JPL Document No. 606-1
s.
MARS SCIENTIFIC MODEL
Claude M. Michaux
Ray L. Newburn, Jr.
Authors
With contributions from
C. F. Capen
C. B. Farmer
E. Haines
R. A. Lyttleton
R. J. Mackin, Jr.
E. D. Miner
E. Monash
M. Neugebauer
R. H. Norton
JET PROPULSION LABORATORY
CALIFORNIA INSTITUTE OF TECHNOLOGY
PASADENA, CALIFORNIA
March 1, 1972
FRONTISPIECE
Photograph of Mars taken by R. B. Leighton of the California institute
of Technology on August 24, 1956, eighteen days before opposition.
The planet was approximately 35.2 million miles from Earth at the
time the photograph was obtained. Mare Cimmerium and Mare
Tyrrhenum dominate the center of the disk, and Syrtis Major is at the
far left. The season is late spring in the southern hemisphere (north is at
the top). The Mt. Wilson 60-inch reflector was used and its aperture was
cut to 21 inches with an off -ax is diaphragm; exposure time was 20
seconds on Kodachrome Type A film. The positive, used in making the
print, was composed by the Jet Propulsion Laboratory. (The repeated
copying of this photograph in the reproduction processes has greatly
decreased the clarity of surface detail and has caused the yellowish
tones of the original positive to appear orange here.)
JPL 606-1 Preface
PREFACE
Our intent in the new edition of this document has been to present a
summary of knowledge about Mars shortly before the time of arrival of
Mariner 9. Emphasis has been given to observational results, with a
limited amount of interpretation where appropriate. Two sections
(5. 3 and 5. 4) have been retained from an earlier edition. The data reflec-
ted in Sections 5. 3 and 5. 4 is nearly 5 years old and should be utilized
solely as background information until later data has been derived. It
is suggested that the interested reader may obtain later data by reference
to the Atmospheres Section of Viking '75 Project, Mars Engineering
Model. Two new Sections, 4. 3 and 5. 5, Secular Change and Atmospheric
Circulation respectively, may be supplied at some future date.
Material in this document has been reviewed extensively both by
Jet Propulsion Laboratory Scientists and by other specialists. Some
errors seem inevitable, however, and the authors will be grateful for
comments, corrections, and criticism from our readers. Each page
shows the date of the latest information and identifies the author or
authors of the material on that page.
Claude M. Michaux
Ray L. Newburn, Jr.
March 1, 1972 page v
Acknowledgments JPL 606-1
ACKNOWLEDGMENTS
Contributions from widely diversified scientific disciplines were
necessary to compile this document. The cooperative spirit of the many
individuals contacted, both on and off the laboratory, is greatly appreci-
ated. Although the document has been almost completely rewritten since
the 1968 edition, something of the spirit and contributions of Mrs. J.
Negus de Wys have come through from that earlier version.
Special thanks are due to James Roth for reviewing the drawings
made after Mariner 6 and 7 TV pictures. Grateful appreciation is
extended to the following investigators who kindly sent Preprints of
important papers which were very useful in the preparation of some of
the sections (indicated in parentheses): D, L. Anderson (2); T. C.
Hanks (2); G. Neugebauer and E. Miner (3. 1); A. B. Binder and J. C.
Jones (3. 2 and 3. 4); R. Goldstein (3. 3); G. C. Pimentel (3. 4); R. E.
Arvidson (3. 5); A. B. Binder (3. 5); W. K. Hartmann (3. 5); G. E.
McGill and D. U. Wise (3. 5); C. H. Thorman and G. G. Goles (3. 5);
A. Woronow and E. A. King (3. 5); W. A. Baum, C. F. Capen, and
L. J. Martin (4. 1 and 4. 2).
We also wish to thank our many colleagues who spent much
valuable time carefully reviewing individual sections of the document.
Their contributions to improve accuracy and clarity have been a great
asset to us.
page vi March 1, 1972
JPl. 606-1 Topical Summary
TOPICAL SUMMARY
The Mars Scientific Model, which is intended to be a source of the most
recent and accurate data for Mars spaceflight program needs, is organized to
provide the user with the means for convenient and expedient location of desired
informiation and also to facilitate updating. The following order of subject
matter appears in each section (or subsection) as applicable: Introduction,
Data Summary, Discussion, Conclusions or Implications, Figures, Tables,
and Bibliography. Some sections also contain a Glossary and/or Appendices.
Each section can be considered a separate entity but can also be used in con-
junction with other correlated sections of the document. Each of the six main
sections contains a detailed Table of Contents, List of Illustrations, and List of
Tables for the material contained within the respective section. This Topical
Summary identifies the primary content of material contained within the appro-
priate section or subsection.
SECTION TITLE AND CONTENT
1. ORBITAL AND PHYSICAL DATA
Historical review of Mars orbit theories and ephemerides.
Orbital elements, related constants and derived data. Earth-Mars
distances. Rotational elements and derived quantities. Preces-
sion the axis. Physical data summary. Seasons. Calendar
(Norton, 1967). Satellites Phobos and Deimos: orbital and phy-
sical data. Astronomical glossary.
2. INTERIOR
Shape of planet: geometrical, dynamical and optical flattenings.
Gravitational potential and coefficient J . Moments of inertia.
Gravity formula of Clairaut. Hydrostatic flattening approximation
of Radau-Darwin. Density models: historical and recent
(Binder, 1 969 ; Ander son, 1972). Thermal history: background
and recent models (Hanks and Anderson, 1969).
3. SURFACE
3.1 Thermal Properties. Theoretical temperatures derivable from
insolation. Atmospheric effects on Mars. Thermal models
(Leighton and Murray, 1966; Kieffer, 1971). Infrared radiometry.
Temperatures observed from Earth. Temperatures observed
from Mariner 6 and 7 spacecrafts. Microwave radiometry. Disc
brightness temiperature s observed from Earth.
3.2 Ultraviolet, Visible, and Infrared Properties . Photometry
and photometric systems in astronomy. Reflection versus emis-
sion from Mars. Integrated photometric properties of Mars:
total brightness, average color, phase function. Bond albedo,
March 1 , 1972 page vii
Topical Summary 606-1
and geometi-ic albedo. Detailed photometry theory: radiance
coefficient, radiance factor, photometric function, normal albedo.
Results of detailed photometry of Mars: table of regional spectral
albedos (Binder and Jones, 197Z). Folarimetry: theory and
results for Mars. Glossary of photometry and polarimetry.
3.3 Radar Properties . Radar astronomy: basic theory and
observational techniques. Radar cross section and reflectivity.
Angular backscattering and surface roughness (average slopes).
Topography from ranging. Results for Mars. Correlation.
3.4 Chemical and Physical Properties . General physical and
chemical properties expected for the Martian ground. Earth-
based reflectance spectra of dark areas and bright areas, with
interpretations. Goethite stability on Mars. Adsorption of
volatiles: carbon dioxide and water vapor experiments. Perma-
frost speculations. Chemical composition of the polar caps.
Mariner 7 IRS spectra of the South Cap. Ozone observation.
Carbon dioxide clathrate hydrate: speculations.
3. 5 Morphology and Processes . Photographic apport of Mariner
6 and 7. Topography obtained by Earth-based radar and the
Mariner 6 and 7 UVS and IRS. New global (Mercator) maps of
Mars, with lists of names. Regional and polar maps. Morphology
of terrain types observed by the Mariners: cratered, chaotic and
featureless terrains. Crater statistics and analyses, from
Mariner 4 and Mariner 6 and 7 photography separately. Crater
modification processes. Ages of large craters: speculations.
Chaotic terrain: distribution, and speculations on age, origin
and processes. Featureless terrain: speculations of age and
origin of Hellas basin and processes on its floor. South Polar Cap:
photography by Mariner 7. Morphology and processes of the 3
zones recognized: margin or edge, interior, and central region
Thickness and permanence of frost cover. Dark and bright areas:
Meridiani Sinus region boundaries and markings observed by
Mariner 6 and 7. Canals, lineaments and oases in the Mariner
photography.
3.6 Mariner 1969 Photographic Atlas of Mars . Mariner 6 and 7
Television Experiment design. Camera system. Image processing.
Far Encounter pictures: a commented selection. Near Encounter
pictures; a commented selection grouped under 5 categories:
Cratered Terrain, Chaotic Terrain, Featureless Terrain, Atmos-
pheric Hazes, and South Polar Cap. Photoference Data Tables
for all ZOO pictures taken.
4. OBSERVATIONAL PHENOMENA
4. 1 Clouds and Hazes . Violet layer and blue clearings. Blue
Clouds.. White clouds and hazes. Yellow clouds and dust storms.
Gray clouds, and bright spots or flares.
page viii March 1 , 1972
JPIj 6 06-1 ^^^^ Topical Summary
4. 2 Seasonal Activity . Polar caps. Polar hoods. Dark polar
fringe. Seasonal evolution of the caps: boundaries and regres-
sion curves. Seasonal behavior of clouds. Recurrent white
clouds. Major yellow clouds or storms. Whitening areas.
Surface features: wave of darkening and seasonal behavior.
Charts of local seasonal changes in equatorial, northern, and
polar areas. Global colored maps of seasonal activity of surface,
with overlays for clouds, possible frost, and wave darkening.
5. ATMOSPHERE
5. 1 Atmospheric Composition. Observed constituents: carbon
dioxide, carbon monoxide, water vapor. Carbon and oxygen
isotopes, dissociation and ionization products (in upper atmos-
phere), ozone (over polar cap). Assumed constituents and upper
limits. Units used for atmospheric abundances.
5. Z Surface Pressure . Historical results using photometry or
polarimetry; critique of the assumptions. Spectroscopic methods
and results. Spacecraft radio occultation method and results of
Mariner 4, 6 and 7. Mean surface pressure, seasonal and
topography effects.
5.3 Lower Atmosphere. Layers of the lower atmosphere. Phy-
sics of the troposphere, stratosphere, and mesophere. Convec-
tive, radiative, and convective- radiative models. Models I II
and III. ' '
5.4 Upper Atmosphere. Layers of the upper atmosphere. Physics
of the photodissociation region. Physics of the ionosphere
including ionization processes and thermal processes. Prelimin-
ary E-Model. F -Model. F -Model.
6. CIS- MARTIAN MEDIUM, RADIATION
Solar electromagnetic radiation at Mars: total irradiation (solar
constant) and spectral distribution (solar spectral irradiance).
Extreme ultraviolet. X-rays, and radio wave radiation components.
Absorption in the Martian atmosphere. Solar particle environ-
ment at Mars: solar wind (protons, alphas, electrons), solar
cosmic rays or flares (protons). Galactic cosmic rays. Induced
charged particle flux (from atmospheric collisions).
Solar interplanetary magnetic field. Martian magnetosphere,
magnetic moment, and surface magnetic field.
Meteoroid environment: sporadic and stream meteoroids. Fluxes
and velocities.
March 1, 1972 p^^^ ,^
JPL 606-1 Orbital and Physical Data
SECTION 1 CONTENTS
1. ORBITAL AND PHYSICAL DATA
Introduction ,
1. 1 Mars — Orbital and Physical Data [ ' ]
Historical Background i
JPL Ephcmerides o
Orbital Elements * ' ' ,
Mean Orbital Elements of Mars ' ' * ^
Orbital Constants and Derived Tabulated Data [ 7
Period of Revolution Y
Daily motion o
Orbital Velocity y
Distance from Sun ^
Distance from Earth jO
Rotational Elements , ^
Position of Mean North Pole of Mars ' * ' 14
Period of Rotation of Mars ' ic,
Derived Quantities 2 5
Martian Longitude System 2 6
Physical Data Summary ' jy
Seasons of Mars ' tq
Earth-Mars Calendar _' ' 22
Ecjuivalences 2 2
Basis of Zero Points * ' p,?
Mars Leap Years ' ' 2?
3
1. 2 Satellites - Orbital and Physical Data . . . . . 2
Orbital Elements 2 3
Physical Data ' ' 23
Bibliography Pq
Appendix A — Earth-Mars Calendar .........[ A-1
Appendix B — Glossary R-l
Figures
1. The reference (orbit and equator) planes of Earth, Mars, and
Satellites ,-
2. Orbital elements /
3. Distances of Mars from Sun and Earth . . 9
4. Oppositions of Mars from 1877 to 1988 12
5. Apparent angular sizes of disk of Mars at oppositions'from
1965 to 1980 ; ^2
6. Precession of the Martian North Pole 14
7. Longitude of central meridian .',*,'* 17
8. Comparison of the Martian seasons and the terrestrial
seasons ., ,
March 1, 1972 0^^ 1 r- 4. j.
' oec. 1, Contents, page 1
Orbital and Physical Data JPL 606-1
1. (cont'd)
9. The Laplacian plane of a Martian satellite 26
10. Illustration of the notation used by Wilkins (1965) and Cain
(1967) in deriving the orbital elements of the Martian
satellites 26
Tables
1, Opposition dates of Mars 11
Z. Exceptionally close Mars approaches 11
3. Superior conjunctions, Earth-Mars, 1960-2000 13
4. Mars physical data summary 18
5. Earth-Mars seasonal durations 20
6. Orbital elements of the Martian satellites 24
7. Nomenclsture for the system of orbital elements necessary
for calculating a Martian satellite's position at anytime t
after epoch to ^^
8. Satellite orbital and related data 27
9. Physical data of Martian satellites 28
A-1. Earth-Mars Calendar A-1
Sec. 1, Contents, page ii March 1, 1972
J PL 6 06-1 Orbital and Physical Data
ORBITAL ANO I IIYSICAL DATA
INTRODUCTION
This section discusses the Orbital and Physical Data for Mars and the
Martian Satellites Deiinos and Phobos.
The discussions are organized into two subsections; the first, 1.1, deals
with the planet Mars while the second, 1.2, deals with the knowTi data covering
the satellites.
1. 1 MARS -ORBITAL AND PHYSICAL DATA
Historial Background
The planet Mars holds a special place in the origin of celestial n-iechanics,
for it was through the study of its orbital motion, from the positional data
gathered by his master Tycho Brahe, that the young Kepler discovered the
three laws of planetary motion which bear his name. (The first two laws were
announced in his Astronomia Nova published in 1609, while the third law
appeared in 1618. ) It was by generalizing these empirically discovered laws
that Newton, with remarkable insight, was aMe to enunciate in 1687 (in the
Principia) the Law of Universal Gravitation, which forms the basis of celestial
mechanics. At the same time, he also formulated the three laws of n-Lotion,
the basis of modern dynamics. The problem of the motion of two bodies under
the inverse square law of gravitational attraction was solved by Newton (yield-
ing the Keplerian ellipse). However, the real problem in dealing with motion
in the solar system is that of n bodies under mutual attraction. This problem
is far more difficult, and in fact the simpler "three-body problem" has not
been solved in its general form. Only special cases of little value to the solu-
tion of the planetary or lunar motion pj-oblems are completely integrable. The
standard procedure in planetary theory is one of successive approximations of
the orbit by the general perturbations method. The planet's motion (a dis-
turbed ellipse) is treated in terms of "perturbations, " or deviations from
Keplerian elliptic motion around the Sun, under the disturbing attraction of
another planet, and thence of all other planets. The analytical expressions
obtained (called "general perturbations") -which usually give the variation
with time of six variables known as "orbital elements" defining the orbit and
position of the planet — are found to contain two classes of terms: (1) "secular
terms, " which change very slowly but proportionally with time (t), although
their sum remains bounded-- (except for very long term dissipative forces),
and (Z) "periodic terms, " which are series of sine and cosine terms, involving
different short periods, where time (t) is again the independent variable. Need-
less to say, these expressions are complicated because the infinite series are
truncated only after naany terms, so as to furnish more accurate results. The
analytical methods employed for the solution of the lunar and planetary
=:Tt was shown by Lagrange and Laplace that these secular terms actually have
extremely long periods (thousands of years).
October 15, 1971 C. Michaux, JPL Sec. 1, page 1
Orbital and I'hysical Data JP.L 606-1
problems \\'ere de-vised primarily by the rrid rhcintir r:,! a s t r o*iO!Tic'" s of the 18th
century, initially, Luler, Claii-aut, ?ind D'Aienibeil, iollowerl by I.ag'ange and
Laplace. The eplnMiierides, or taljbjs of the plautit's positions at regular time
intervals, derived from planetary Iheofy :i h comp reliensi\"'-ly developed by
I;aplace, however, were? of insufficient accuracy (for inany navigational or
astronomical purposes) when compared with actua] observations. Although
inaccurate, Lindenau's (1811) Tables, or del.abindc's (17 92) Tables for Mars
were definitely supetLor to Flaiiey's (1749) tTi.i-in'i labitns. Cons ecjuently,
planetary theory required improvement, which was accomplished in the 19th
century initially by Hansen, and T>everrier, with refine'nents by Hill and
Newcoinb. This often meant nev,- approaches (Hansen's method of treating
perturbations), or major ov^erhauls (Newcomb's introduction of a uniform
system of astronomical constants, such as planetary masses) (Newcomib, 1895).
The resulting improved tables which appeared for Mars were those of
Leverrier (1861) and the more accurate tables of Newcoinb (1898). In the early
part of the ZOth century, observations showed that Newcomb's orbit for Mars
was still unsatisfactory. Indeed, Newcomb had neglected the second order
perturbations, due mainly to the attractions of Flarth and Jupiter, and also due,
in lesser degree, to Venus and Saturn, The empirical corrections given by
Ross (1917) did not significantly improve the situation. A new theory of Mars
was then constructed by Clemence, who used Hansen's method (which Hill had
found so successful in treating Jupiter and Saturn). The new theory included
the previously neglected second-order terms, and some third-order terms,
when it appeared in its complete form in 1961 (Clemence, 1949, for first-order
theory, and Clemence, 1961, for second-order theory). The great accuracy of
the new theory, about 0.04 arc seconds difference between observed and pre-
dicted longitudes of Mars, warrants the issuance of much more reliable
ephemerides of Mars in future volumes of The American FZphemeris and
Nautical Almanac. At the present time, this publication still calculates the
Mars positions according to Newcomb's old theory, supplemented by Ross'
corrections. However, the U.S. Naval Observatory has issued (until final
adoption of Clemence 's work) provisional ephemerides of Mars for the
1800-2000 period (Duncombe, 1964, for 1800-1O50 period, and Duncombe and
Clemence, I960, for 1950-2000 period).
The last two decades have witnessed three revolutionary developments in
engineering which permitted rapid improvement of nnr knowledge of the solar
system through applied celestial iriechanics. These are
1) High-speed digital computers.
2) Spacecraft.
3) Radar astronomy.
With the new computers, it is now possible to perform (without undue
labor) the step-by-step numerical integration of the differential equations of
(disturbed) motion of any planet along its orbit. I'his was first done for Mars
by Herget to test Clemence's theory, but is now performed routinely, bj derive
new ephemerides of ever-increasing accuracy, usinL? all data available (i.e.,
optical, spacecraft, and radar data). With radio t>-a eking of snacec raft flying
past Venus and Mars, and the r-adar bounce r--r<" '"-•'* ■= ' tt these !)'anets, the
Sec. 1, page 2 C, Mirhaux, T^M. Oct-ber 1-, 1971
JPL 606-1 Orbital and Physical Data
precise measurements of their distances and velocities have become available,
as well as their masses and, with lesser accuracy, their radii. It is now
possible to follow Mars by radar around its entire orbit. Therefore, in a few
years, we can expect another revision of the theory of Mars m^otion through
modern analytical and computational methods of celestial mechanics.
JPL Ephemerides
In recent years, necessity and the data obtained from spacecraft and radar
exploration of the nearby solar system (Moon, Venus, Mars, Mercury),
coupled with the availability of high speed computers, have prompted JPL to
produce its own lunar and planetary ephemerides. The JPL ephemierides are
of superior accuracy to the existing publications (such as those published by the
Amierican or British Nautical Almanac Offices), which give latitude and longitude
to -v-O.l arc seconds. The JPL ephemerides are under continual development,
fully utilizing the almost continuous influx of new data from optical, radar, and
spacecraft observations, as well as the latest or most advanced forms of lunar
and planetary theories of motion. Thus, these ephemerides are of ever-
increasing accuracy. To prevent accuracy degradation of the source data, and
also permit direct use by computer programs, magnetic tape recording is used
at every step. The JPL Ephemeris Tape System, established in 1964, is a
collection of procedures, computer programs, and tape archives of ephemeris
data used in generating the new epheraerides called JPL Ephemeris Tapes.
These ephemeris tapes are prepared in sets of three, which collectively cover
the 1950-2000 period. They tabulate for that period the rectangular coordinates
and velocity vector components of the nine planets and the Moon, as referenced
to the mean equator and equinox of 1950. They also include nutations in
longitude and obliquity, and their rates, plus modified second and fourth
differences of all these quantities to facilitate interpolation. The method of
special perturbations or step-by-step numerical integration of the orbits is used
for calculating the planetary positions and velocities, with relativity effects
(as given by the Schwarzschild metric) taken into account. [Note: The special
perturbations solution consists of determining the orbit as closely as possible
by successive approximations ("integration fits") —through numerical integra-
tion of the second order differential equations of motion of the planets with
selected epoch values to best fit the source positions by the least-squares
procedure. This method is also known as 'fitted integration' of an orbit, to fit
source accuracy. The source positions and theories used ultimately determine
the accuracy of the ephemerides produced.]
More complete description of the JPL Ephemeris Tape System can be
found in the report issued by Peabody et al. (1964). JPL ephemerides approved
for external distribution ("Export Ephemerides") so far have been successively:
DE 3 (Peabody et al. , 1964), DE (Devine, 1967), and DE 69 (O'Handley
et al. , 1969). These ephemerides have been extensively tested with real data.
Experimental and special purpose ephemerides are available only to JPL
internal users.
October 15, 1971 C. Michaux, JPL Sec. 1, page 3
Orbital and Physical Data JPL 606-1
Orbital Elements (See Figs. I and 2)
The elliptical unperturbed (or Keplerian) orbital motion of a planet around
the Sun is completely determined by the six elements defined below (Fig. 2) and
referenced to the ecliptic plane and the vernal equinox of Earth (Fig. 1):
a = semi -major axis (of the ellipse)
e = eccentricity (of the ellipse)
i = inclination angle of orbit to ecliptic
^ = longitude of ascending node {ii) of the orbit on the ecliptic, or
CJ - longitude of perihelion (it), measured fronn equinox along the
ecliptic to node, then along the orbit to perihelion, or w = TO + ijir'
M =
or
mean anomaly of planet at time t (epoch), since perihelion passage
time tg, or M = n(t - tg), as derived from the mean daily motion
n = 2tt/P, where P is the sidereal period of revolution of the
planet*
or
L = mean longitude of planet at time t, that is simply L = C + M
Since planetary orbits are never exactly elliptical because of the gravita-
tional perturbations due to other planets and satellites, it is customary to give
a mean reference orbit, valid for a limited period of time (a year or so),
which sufficiently approximates the actual orbit for most astronomical pur-
poses. Neglecting the irregularly varying "periodic" terms of the perturbations
and utilizing the progressive "secular" terms of the perturbations, formulae for
the "mean elements" of such an orbit, at a certain time, are available in the
form of polynomials in the time variable t, which is the time interval since an
epoch t = 0. While the elements e and n (or P) vary extremely' slowly with
time (a°is considered constant), the elements i, ^, and w are dependent upon
the reference system (ecliptic and vernal equinox) and vary more rapidly. The
mean elements may be referenced to either the instantaneous ecliptic and mean
equinox - "of date" which is continually moving - or to a fixed ecliptic and mean
equinox of a conveniently chosen epoch (such as 1950.0). The latter system is
more suitable to astrodynam^ical applications.
. 2 3 _
*Note: From planetary theory, n is related to a by the relation n a -
k2(l + m), where k is the Gaussian gravitational constant, and m is the
planet's mass in terms of the Sun's mass.
Sec. 1, page 4 C. Michaux, JPL October 15, 1971
JPL 606-1
Orbital and Physical Data
a
(U
ca
X
■'-' C
C
s °
nJ
O 4-1
•I-l T-i
4-> Cfi
to"
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a
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rd
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^ ,
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-J->
■r-(
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K 53 <u
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^2-
1 — 1
■rH >H n
ax rt Q,
(U
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h
October 15, 1971
C. Michaux, JPL
Sec. 1, page 5
Orbital and Physical Data
JPL 60b-l
ECLIPTIC
DESCENDING NODE U
PLANET'S ORBIT PLANE INTERSECTION
WITH CELESTIAL SPHERE
P' PLANET'S POSITION (Projected)
V PERIHELION POSITION (Projected)
7>^ ft ASCENDING NODE
VERNAL EQUINOX
LONGITUDE OF ASCENDING NODE: fl = Tfl
O - CENTER OF ORBIT
S - SUN
P - PLANET
V - PERIHELION
LONGITUDE OF PERIHELION:
(TRUE) LONGITUDE OF PLANET:
SEMI-MAJOR AXIS:
ECCENTRICITY:
INCLINATION (UPON ECLIPTIC):
XJ = tO. + At'
L' = Tft +n»-' +ir' P'
a = Oir
e = OS
Ot
Fig. 2. Orbital elements.
Mean Orbital Elements of Mars (Figs. 1 and 2)
1) Referred to (moving) mean equinox and ecliptic of date (American
Ephemeris,
Explanatory-
Supplement,
1961)
a = 1.5236915 AU
e = 0.09331290 + 0.000092064 t - 0.000000077 t
2
i = 1°51'01.20" - 2.430" t + 0.0454 t
U = 48M7'11.19" + 2775.57" - 0.005" t^ - 0.0192" t^
C5 = 334'"13'05.53" + 6626.73" t + 0.4675" t^ - 0.0043" t^
M = 319°31'45.93" + (53^" + 215490.60") t + 0.6509" t^ + 0.0043" t^
= 319.529425° + 0.5240207666° d + 0.000013553° D^
+ 0.000000025°d3
Sec. 1, page 6
C. Michaux, JPL
October 15, 1971
JPL 606-1 Orbital and Physical Data
Time t or interval (t-to)» where to is the fundamental epoch 1900,
January 0.5 E. T. or Julian Date J. D. 2415020.0, is the fundamental
variable and is measured in Julian centuries of 36525 ephemeris days,
counted since the fundamental epoch.
However, for convenience the submultiples D and d are used:
D = 3.6525 t (in units of 10000 ephemeris days)
d = 10000 D = 36525 t (in ephemeris days)
2) Referred to (fixed) inean equinox and ecliptic of 1950.0 (Sturms,
1970)
a = 1.5236915 AU
e = 0.09335891275 + 0.000091987 t - 0.000000077 t^
i = 1.85000° - 0.00821° t - 0.00002° t^
n = 49.17193° - 0.029470° t - 0.00065° t^
C5 = 285.96668° + 0.73907° t + 0.00047° t^
M = 169.458720° + 0.5240207716° d + 0.0001825972° t^
+ 0.0000011944° t3
Time t, in this set of equations, is counted from the fundamental
epoch 1950, January 1.0 E. T. or Julian Date 2433282.5, and is measured
in Julian centuries. It is important to note that this epoch is not exactly
the sanie instant as that of the reference coordinate system epoch 1950.0,
which corresponds to Julian Date 2433282.423357. The reason for this
small shift is to be able to express the Julian Date in a round number
count from the time reference epoch.
Orbital Constants and Derived Tabulated Data
Period of Revolution
The Martian year (687 days) is somewhat less (by 43 days) than two
terrestrial years, while the mean interval of 780 days between successive
oppositions (mean synodic period) is the longest of all major planets.
The mean daily motion of Mars is slower than that of Earth by
0.461576 degree per day.
P Period of revolution (sidereal) = 686.9804 mean solar days
= 1.88089 tropical years
October 15, 1971 C. Michaux, JPL Sec. 1, page 7
Orbital and Physical Data JPL 606-1
n Mean daily motion, n = 360°/P (sidereal) = 0.5Z4033 degree
per day
S Mean synodic period* (Earth/Mars) = 779.94 mean solar days
Daily motion . The daily motion (degrees /day) is given for every 4th day
of the current and preceding years by the American Ephemeris and Nautical
Almanac.
Orbital velocity
, -1
km sec
Mean 24.1
Perihelion 26.4
Aphelion 22.0
The orbital velocity may be calculated from the heliocentric distance
(radius vector) r, the semi -major axis a, and the masses of planet m and
Sun M by the formula:
f7^ , , , ,2 1 V TT ■ , r and a in AU
V (in AU per day) = k ^/(M + m)(_ - -) Units: ^ relative to Sun M = 1,
where k is the Gaussian gravitational constant (k = 0.01720209895)
mi
. sec
15.0
16.4
13.6
Distance from Sun
(Fig. 3) AU km
mi
Mean (semi-major axis) 1.5236(91) 227,800,000 141,500,000
Perihelion 1.3815 206,500,000 128,300,000
Aphelion 1.6660 249,100,000 154.800,000
(American Ephemeris
and Nautical Almanac)
The heliocentric distance (radius vector) is given for every 4th day of the
current and preceding years by the annual volumes of the American Ephemeris
and Nautical Almanac, and for every 10th day for years 1800 to 1980 by the
Planetary Coordinates tables; and for 1973 to 2000 by the tables and graphs
issued by Souders (1970).
*The synodic period S is derived from the sidereal periods Pm °f Mars and Pg
of Earth (Pg = 365.25636 d) by the relation used for superior planets:
1 1 1
s p p ■
e m
Sec. 1, page 8 C. Michaux, JPL October 15, 1971
JPL 606-1
Orbital and Physical Data
10
Jorl. 1
20 30 10 20
Feb. 1
10 20 XI 10 20 30 10 20 30 10 20 30 10 20 30 10 20 30 10 20 30 10 20 30 10 20 30 10 20 30
Mar. 1 Apr. 1 May 1 June 1 July 1 Aug. 1 Sept. 1 Oct. 1 Nov. 1 Dec. 1
1971 DATE
NOTE: Similar figures were issued by Souders (1970) for the years 1973-2000,
Fig. 3. Distances of Mars from Sun and Earth.
October 15, 1971 C. Michaux, JPL Sec. 1, page 9
Orbital and Physical Data
JPL 606-1
Distance from Earth
(Fig. 3) AU
Mean opposition 0.5236
Minimun-L distance 0.37 28
Maximum distance 2.6657
km
78,350,000
55,810,000
398,900,000
mi
48,695,000
34,670,000
247,900,000
(American Ephemeris
and Nautical Almanac)
The geocentric distance (Earth-Mars separation) is given for every day of
the current and preceding years by the American Ephemeris and Nautical
Almanac, and for every 10th day of years 1973 to 2000 by the tables and graphs
issued by Souders (1970).
Opposition dates and minimum distances (at closest approach) are listed
in Tables 1 and 2 for years 1937 to 1980, Figure 4 illustrates orbital
occurrences of many oppositions from 1877 to 1988. Figure 5 compares
angular sizes of the Martian disk as seen from Earth through a complete
15.8-year (average) cycle of successive oppositions, Superior conjunction dates
and maximum distances are listed in Table 3 for years I960 to 2000, as
established from Souders' (1970) tables using graphic interpolation and from
the American Ephemeris and Nautical Almanac.
Rotational Elements
The solid surface of Mars is visible most of the time, which permitted
early observers to quite accurately determine its rotation period. The current
rotation period data is accurate to within 0.02 second. The usual method of
determination consists in timing transits of conspicuous surface features
(e. g. , Meridiani Sinus) across the central meridian of Mars. It is also
possible to measure the areographic longitude of a surface feature on a photo-
graph taken at a precise time. This sometimes has been applied to old draw-
ings, although the resultant accuracy is questionable. It was done, however, by
Wislicenus (1886), and Bakhuyzen (1897), and provided good results,
24h37m2ZS66, when compared to the modern value, +22.67 seconds, derived by
Ashbrook (1953), utilizing only transit material (1877-1952 period).
The present orientation of Mars' axis of rotation in the celestial sphere —
some 10° from the star Deneb -is known with lesser accuracy than the rotation
period, probably to within 0. 1 ° (according to de Vaucouleurs, 1971, private
communication). The inclination of Mars' equator to its orbit, or the Martian
"obliquity, " is very nearly 25°. See Figure 6.
The precession rate of the Mars' axis, or of its equinoxes, is not too
accurately known. Determinations by Struve (1898) —7.07 arc seconds /year,
and by Lowell (1914) —7.08 arc seconds /year, were questioned by Fish (1964),
who proposed 7.34 arc seconds/year (using de Vaucouleurs' inclination value).
The precession rate is still in question, but Struve's value is most commonly
utilized. The corresponding precessional period is about 183,300 years.
Sec. 1, page 10
C. Michaux, JPL
October 15, 1971
JPL 606-1
Orbital and Physical Data
Table 1, Opposition dates of Mar s. (Data from Slipher, 1962; American
Ephemeris and Nautical Almanac, 1937- 1972; Miner, 1967)
Date of
opposition
Interval from
previous date,
days
1937 May
19
795
1939 Jul
23
810
1941 Oct
10
786
1943 Dec
5
771
1946 Jan
14
764
1948 Feb
17
765
1950 Mar
23
770
1952 May
1
784
1954 Jun
24
806
1956 Sep
11
800
1958 Nov
17
775
1960 Dec
30
776
1963 Feb
4
763
1965 Mar
9
767
1967 Apr
15
777
1969 May
31
781
1971 Aug
10
807
1973 Oct
25
791
1975 Dec
15
771
1978 Jan
ZZ
764
1980 Feb
25
Distance from Earth^
Million mile s
47.3
36. 1
38.2
50. 1
60.4
63.0
60.4
51.9
39.8
35.2
45.4
56.3
62.2
62.0
55.8
44.5
34.9
40.4
52.4
60.8
63.2
Million kilometers
76. 1
58.
61.4
80.7
95,6
101.4
97.2
83
5
64
1
56
6
73
90
6
100
1
99
8
89
8
71
7
56
2
65.
84.
3
97.
8
101.
7
s before or aft
er
At closest approach, which may be aa much as 10 days before or aft
opposition.
Table 2. Exceptionally close Mars approaches
Date of
opposition
Interval from
previous date,
years
Distance from Earth*
Million miles
Million kilometers
1877 Sep 5
1892 Aug 26
1909 Sep 18
1924 Aug 22
1939 Jul 23
1956 Sep 11
1971 Aug 10
1988 Sep 28
15.0
17. I
14.9
14.9
17.2
14.9
17.2
34.8
34.5
36.2
34.5
36.1
35.2
34.9
36.3
56.0
55.5
58.3
55.5
58.0
56.6
56.2
58.4
At closest approach, which may be as much as 10 days before or after
opposition .
October 15, 1971
R. Newburn, JPL
Sec. 1, page 1 1
Orbital and Physical Data
JPL 606-1
MIDWINTER
{N HEMiSPHERt
MIDFALL
N HEMISPHERE)
MIOSPRING
(N HEMISPHERE
MIDSUMMER
N HEMISPHERE)
Fig.
4. Oppositions of Mars from 1877 to 1988. Surrounding calendar
indicates time of occurrence. Broken line is major axis of
Mars orbit connecting aphelion and perihelion. Aphelion
and perihelion of Earth are indicated on Earth orbit.
(After Ley, Von Braun, and Bonc'stell, I960)
D1A^^I4 15 6
(SEC OF
ARC ) END
SPRING SUMMER
25 Z
2i 7
MID ■V'NTER
DATE-
16 6 14 5 14
BEGIN
SPRING END
MAX MID SPRING SPRING
o o
75
80
Fig. 5. Apparent angular sizes of disk of Mars at oppositions from
1965 to 1980. Seasons indicated are for nor the rn hemisphe re.
North and south poles and approximate extent of polar caps
are indicated. (Miner, de Wys, 1967)
Sec. 1, page 12
C. Michaux, JPL
October 15, 1971
JPL 606-1
Orbital and Physical Data
Table 3. Superior conjunctions, Earth-Mars, 1960-2000
Year
Superior Conjunction
Date
Conjunction Distance
Million
kilometers*
Million
miles
Furthest Distance
Date
Furthest Distance
Million
kilometers*
Million
miles
1961
1964
1966
1968
1970
1972
1974
1976
1979
1981
1983
1985
1987
1989
1991
1993
1996
1998
2000
Dec. 14
Feb. 17
Apr. 29
Jun. 21
Aug. 2
Sep. 7
Oct. 16
Nov. 26
Jan. 21
Apr. 2
Jun. 4
Jul. 18
Aug. 27
Oct. 1
Nov. 10
Dec. 27
Mar. 6
May 13
Jul. 2
367
355
367
387
398
400
391
374
358
359
380
395
400
395
381
363
355
372
391
228
221
228
240
247
249
243
232
222
223
236
245
249
245
237
226
221
231
243
Nov. 7
Feb. 20
Jun. 13
Jul. 15
Aug. 9
Sep. 1
Sep. 23
Oct. 24
Dec. 13
(1978)
May 21
Jul. 4
Aug. 1
Aug. 28
Sep. 15
Oct. 13
Nov. 19
Apr. 17
Jun. 21
Jul. 20
371
355
373
389
399
400
3 94
378
359
364
384
396
400
3 97
385
367
357
377
392
231
221
232
242
248
249
245
235
223
226
239
246
249
247
239
228
222
234
244
^Calculations using the value adopted by the I. A. U. in 1968 for
1 A. U. = 149600 X 10° m.
October 15, 1971
C. Michaux, JPL
Sec. 1, page 13
Orbital and Physical Data
JPL 606-1
Position of Mean North Pole of Mars (See Fig. 6)
The following formulae for right ascension (^o) a^nd declination (6q) of
the Mars North Pole, referenced to the mean (Earth) equator and ecliptic,
include the effects of known precession rates for Mars and Earth north poles,
but not the effects of nutation (which may be neglected in view of the present
accuracy of the formulae).
1) With respect to t he (moving) mean equator and ecliptic of date :
a = 316.55° + 0.00675r(t - 1905.0) (de Vaucouleurs, 1964)
° (A. E. , 1968; 1971
6 = 5Z,85° + 0.003480''(t - 1905.0) corrected)
o
at start of year t .
Z) With respect to the (fixed) mean equator and ecliptic of 1950.0 :
= 316.8538° - 0.0996°t
50
6^„ - 53.0066° - 0.0566°t
(Sturms, 1970)
where time interval t is counted from the 1950, January
1,0 E. T. , time reference epoch in Julian centuries.
Formulae 2) were derived from formulae 1) by Sturms (1970)
using the appropriate coordinate transformations explained in his
document, but only include the Martian precession.
AD 47471
Vega
LYRA
Fig. 6. Precession of the Martian North Pole.
Sec. 1, page 14 C. Michaux, JPL October 15, 1971
JPL 606-1 Orbital and Physical Data
3) Rate of precession of Mars North Pole used in these formulae:
\i = 7.07 arc seconds per tropical year (Struve, 1895)
This precessional rate appears to be in need of revision, which
would immediately affect the rate coefficients in the above formulae.
Derived quantities (from formulae Z) above:
1) Martian "obliquity" (inclination of Mars equator to orbit):
I = 24.76883° + O.OlZZO't + 0.00006° t^ (Sturms, 1970)
where time interval t is counted from the 1950, January 1.0 E. T.
time reference epoch (as previously) in Julian centuries,
2) Angle along Mars equator, measured from ascending node on the
mean 1950.0 Earth equator to the Mars autumnal equinox : this
angle is a quantity useful in making coordinate transformations.
A^p = 43.34526°- 0.0918rt - 0.00010°t^ (Sturms, 1970)
where t is measured as in 1) in Julian centuries.
Period of Rotation of Mars
The rotational period of Mars has been known with great accuracy for
sometime because of the visibility of the surface itself. Two periods may be
distinguished:
1) Sidereal period of rotation, which is referred to the (moving)
Martian vernal equinox: (This is the period usually quoted. )
H m <5 <5
P = 24 37 22.6689 ±0r0026 (m.e. ) E. T. (Ashbrook, 1953)
Note: The period in mean solar time is 0.0010 seconds shorter over the
period of transit observations, 1879-1952, considered by Ashbrook in
deriving his value. For more information, see de Vaucouleurs (1964),
2) True or actual period of rotation; i. e. , relative to a fixed direction,
eliminating the precession of the Martian equinox)
P = 24 37 22:6701 ±070026 (m.e. ) E. T. (Ashbrook, 1953 and
de Vaucouleurs, 1964)
It is 0.0012 seconds longer than the so-called "sidereal" period
(assuming jx = 7. 07 arc sec per year).
Derived Quantities
1) Sidereal rate of rotation, or daily angular rotation :
360°
R = — p — = 350.891962 degrees per terrestrial day
October 15, 1971 C, Michaux, JPL Sec. 1, page 15
Orbital and Physical Data
JPL 606-1
This rotation rate was adopted in I960 by the American Ephemeris
Office for computation of its Ephemeris for Physical Observations of
Mars, specifically for the "longitude of the central meridian. "
2) The longitude of the central meridian (L. C. M. or co) at any time t
in the Martian longitude system (see below). L. C. M. is the
Martian hour angle of the Earth measured from the zero meridian.
It is determined by the adopted longitude of the central meridian
'^o(to) 3-t a- chosen epoch to and the rotation rate R, as well as the
orbital positions of Earth and Mars, according to the formula with
light-time correction (see Fig. 7):
(L.C.M. = ) w
= V +
o
R (t
-t ) - kA
o
where
Values Adopted by
American Ephemeris, 1971;
V = value of V (defined
below) at chosen
epoch t
^ o
k = value of light -time
per AU
V = 149. 475 deg
o *
k = 0.00577560 day AU
(= 499.012 sec AU" 1
geocentric distance
of Mars in AU
A
E
R
areocentric right
ascension of Earth
sidereal rate of
rotation of Mars
R = 350. 891962 deg day
- 1
(t-t ) = time in Julian days
since epoch t
t = J, D. 2418322.
° (i.e. , 1909.04 epoch)
Martian Longitude System (Adopted by American Ephemeris, 1909)
Areographic longitudes are measured from 0° to 360° opposite to rotation
direction (that is, clockwise as viewed from Mars North Pole) along the equator
of Mars corresponding to the adopted position of the North Pole (see above) and
an adopted value (in 1909) for the longitude of the central meridian, as follows:
uj = 344.41° at chosen epoch t of 1909- 04 (January 15.5 U. T. ), or
t ° = J.D. 2418322.0. °
o
The zero meridian, or origin of longitudes, with this 1909 convention,
then falls at a point on Mars' surface that is about 3 degrees west of the old
origin of longitudes (as used by Schiaparelli, Marth, etc.), which was the apex
of the wedge dividing Meridian! Sinus (and called Fastigium Aryn).
Sec. 1, page 16
C. Michaux, JPL
October 15, 1971
JPL 606-1
Orbital and Physical Data
MARTIAN VERNAL EQUINOX
^^}^ ^9t^,
SENSE OF ROTATION OF MARS
PRIME MERIDIAN (at time t)
Aw = ui' -u, = RkA
(light-time correction)
CM.
CENTRAL MERIDIAN
TO EARTH
PRIME MERIDIAN (at time t-<<A;
i.e., OS $een hxmi Earth ot time t)
ILLUSTRATK3N OF FORMULA:
L.C.M. = oj • V-WcA-|A,|=w'-RkA
WHERE V = Vg + R(t-t ). '
QUANTITIES ARE DEFINED IN TEXT.
Fig. 7. Longitude of central meridian
3) The Martian hour angle of the Martian vernal equinox , measured
from the zero meridian of the Martian longitude system. It is given
either by:
V = 149.475° + 350.891962° (J. D. -2418322.0) (Melbourne et al.
1968)
where time is counted in days since reference epoch 1904 January 15.5
U. T. ; or by:
V = 148.672501° + 350.891962° (J. E. D. - 2433282. 5) (Sturms, 1970)
where time is counted in (ephemeris) days since reference epoch
1950, January 1.0 E. T.
Physical Data Summary
Physical characteristics (constants) of Mars are identified in Table 4,
which presents the best available values to date, both in absolutes and relative
to Earth.
October 15, 1971
C. Michaux, JPL
Sec. 1, page 17
Orbital and Physical Data
JPL 606-1
Table 4. Mars physical data sumniary.
Characteristic
parameter
Flattening:
Geon^etrical
Mariner Occulta -
tion Expts)
Dynamical
(satellites orbits)
Optical
(telescopic)
External gravita-
tional potential:
Zonal harmonic
coefficient J-,
Radius :
Equatorial
(Mariner Occulta -
tion Expts)
At 88 °N latitude
(Mariner Occulta-
tion Expts)
Polar-l= (calcu-
lated from
equatorial radius
and dynamical
flattening)
Surface area of:
Ellipsoid (calcu-
lated from radii)
Volume:
Ellipsoid (calcu-
lated from radii)
Relative
(Earths i;
1.70
1.73
0.5320
0.5310
0.283
0.150
Absolute value
0.0057 ±0.0012
0.00525
0.012
0.00187
±0.00007
3393 ±2 km
3370 ±5 km
3375 km
1.4418
X 108 km^
1.6282
X 10^^ km-^
Source
Kliore(1971)
Cain (1967a)
Null(1971)
Kliore(1971)
Kliore{1971)
=:=No polar radius measurement is available yet from occultation
experiments .
Sec. 1, page 18
C. Michaux, JPL
October 15, 1971
JPL 606-1
Orbital and Physical Data
Table 4. Mars physical data summary (continued).
Characteristic
parameter
Relative
(Earth=l)
Absolute value
Source
Areocentric gravi-
tational constant
GM(^ (Gravitational
constant x mass)
(Mariner 4)
Mass M(^ (derived
using** G = 6.67 32
X 10-23km3g-l
sec "2)
Reciprocal mass
(relative to Sun = 1)
M„/M^
Density:
Mean (calculated
from mass and
volume)
Surface gravity:
Equatorial (cal-
culated from
Clairaut's
equation and
data of this
section)
Velocity of escape:
Equatorial
(calculated)
0.1074
0.1074
(1/0.1074)
0,715
0.379
0.449
42828. 3Z
±0.13 km^sec-^
6.418 X 1026 g
3098714 ±9
3.945 g cm-3
37 1 cm sec "^
5.024 km sec"^
Anderson,
Efron, and
Wong(1970)
Null(1971)
*>;<Value of G determined by Heyl (1942) and recommended by
Mulholland et al. (1968).
October 15, 1971
C. Michaux, JPL
Sec. 1, page 19
Orbital and Physical Data
JPL 606-1
Seasons of Mars
Mars has seasons comparable to those of Earth because of the nearly
equal tilt of its axis to its orbital plane. However, the seasons, in the average,
are about twice as long, which correspond to the greater length of the Martian
year. Furthermore, they are distinctly unequal in duration, as a result of the
appreciable eccentricity of the Martian orbit. Table 5 gives the durations of
these seasons in terrestrial and Martian day units, and also the areocentric
longitudes of the Sun at the two equinoctial and two solsticial points. Figure 8
compares the occurrence of terrestrial and Martian seasons and shows that
they do not occur at the same heliocentric longitudes. This is because of the
different celestial orientation of the two rotational axes. Mars' axis is about
37° away from that of Earth. It is actually 85° in heliocentric longitude ahead
of Earth's in its (also retrograde) precessional motion around the pole of its
orbit. The orbit pole of Mars is very close to the pole of the ecliptic (only
1.85° away, as orbital inclination shows).
In the course of time, the seasons of Mars change very slowly both in
celestial and orbital positions because of the precessions of the node (or line
of equinoxes) and of the perihelion (or line of apsides). The effects (such as
change in duration of the seasons) are more marked than in the case of the
Earth because of the orbit's fair eccentricity. The Martian seasons go through
an effective precessional cycle of about 52,000 years period, which is
determined by the combination of the period of regression of the node
(~183,000 years) and of the period of advance of the perihelion (~72,000 years,
according to Brouwer and Clemence, 1961).
The areocentric longitude of the Sun, Lq — which indicates seasonal date on
Mars — is found tabulated for every second d!ay through Mars apparition periods
in the American Ephemeris and Nautical Almanac annual volumes (under Mars-
Ephemeris for Physical Observations). The relationship of Lg to the heliocentric
orbital longitude of Mars, commonly denoted by tj, is at present appr^ .ximately:
Lg - ^ -85° (constant varying very slowly with precession).
Table 5. Earth-Mars seasonal durations.
Areocentric
longitude
of the Sun
Season
Duration of the seasons on
Mars
Earth
Northern
Hemisphere
Southern
Hemisphere
Martian
days
Terrestrial
days
Terrestrial
days
- 90°
Spring
Autumn
194
199
92.9
90 - 180°
Summer
Winter
178
183
93.6
180 - 270°
Autumn
Spring
143
147
89.7
270 - 360° or 0°
Winter
Summer
154
158
89.1
669
687
365.3
Sec. 1, page 20
C. Michaux, JPL
October 15, 1971
JPL 606-1
Orbital and Physical Data
MARTIAN VERNAL
EQUINOX
Ls= 180*
(^^265°)
Lj - 275*
(, = 0*)
TERRESTRIAL
VERNAL
EQUINOX
^^^52;;^.
SYMBOLS:
Lj = AREOCENTRIC LONGITUDE OF THE SUN (COUNTED FROMTcf)
V = HELIOCENTRIC LONGITUDE OF MARS (COUNTED FROM T )
TT = PERIHELION
a = APHELION
— Figure 8. Comparison of the Martian seasons and the terrestrial seasons (in 1971).
October 15, 1971
C. Michaux, JPL
Sec. 1 , page 21
Orbital and Physical Data JPL 606-1
Earth-Mars Calendar (See Appendix A)
A Martian calendar is desirable to establish an equivalence of Earth and
Mars dates, and to simplify the presentation and interpretation of secular and
seasonal changes on Mars. [Note by C. Michaux, 1972 : The accompanying
calendar (Table A- 1 ) was developed by R. Norton in 1967 and was based on the
position of the Mars North Pole in use then (A. E. 1905-1968). A revision is
to be issued later.]
Equivalences
The basic units of time for a Mars calendar are the Mars tropical year
and the Mars mean solar day. The variability of the length of Martian seasons
resulting fromi the eccentricity of the orbit, limits the usefulness of a seasonal
calendar. However, seasons, referred to the Northern Hemisphere, have been
included in the calendar as rough approximations, beginning on the days listed
below; no attempt has been made to define "months. "
1) One mean Mars day = 1.02749133 mean Earth solar days.
(= 24^39'^35r2509).
2) One Mars tropical year - 668.592159 Mars solar days.
3) Mars spring begins on day 167.1.
4) Mars summer begins on day 361.3.
5) Mars fall begins on day 538.1.
6) Mars winter begins on day 11.3.
Basis of Zero Points
The year 1000 has been adopted as the year of the Mariner IV flyby. The
start of this year has been defined such that the time of year was approximately
0.25 when Mars passed its vernal equinox prior to the Mariner IV flyby in July
of 1965. This equinox occurred on September 12, 1964, Earth time and has
therefore been designated as day number 167 of the Martian year 1000.
The start of each day on Mars is defined to be at midnight at 0° longitude;
i, e. , each day starts when the subsolar longitude is 180°.
Mars Leap Years
A Mars leap year is defined as follows:
1) If year is odd, it is a leap year (669 Mars solar days).
2) If year is even, it is not a leap year, except
a) if year is divisible by 10, it is a leap year, except
b) if year is divisible by 100 it is not a leap year, except
c) if year is divisible by 500 it is a leap year.
Sec. 1, page 22 R. Norton, JPL April 1 , 1967
JPL 606-1 Orbital and Physical Data
At the conclusion of a Martian 500 -year cycle, the cumulative error
between the calendar date and the astronomical date, based on the tropical year,
amounts to 0.08 Mars solar day (about two hours). This accuracy is slightly
better than that of the present Earth calendar.
In the Earth-Mars calendar, the column labeled 'Consecutive Mars day'
lists a day number for Mars, similar to the Julian day number for Earth.
Consecutive Mars day number 1 was the first day of the Martian year (which
was a leap year). The calendar also contains a column of heliocentric ecliptic
longitude equivalents for each day. Earth-Mars data for 1963 through 1983 are
presented at this time.
An additional column has been included in Table A-1 to give opposition
dates and distances, spacecraft events, and other information which may be of
interest. The calendars also can be used to correlate Martian seasons with
past observations.
Earth dates are for 0^ GMT; Mars dates show day of year and fraction of
day elapsed at 0° longitude on Mars.
1. 2 SATELLITES - ORBITAL AND PHYSICAL DATA
Orbital Elements
The two Martian satellites, Phobos and Deimos, discovered in 1877 by
A. Hall, are very small and revolve quite close to the surface of Mars.
Phobos, in fact, revolves faster than Mars rotates, and is the only natural
satellite in the solar system known to behave in this manner. The orbits of
the two satellites are nearly circular and lie very close to the equatorial plane
of Mars. As demonstrated in celestial mechanics, the perturbations principally
due to the oblateness of Mars, but also to the Sun, cause each orbit to precess
on a fixed "Laplacian, " plane passing through the intersection of the equatorial
and orbital planes of Mars, and lying between them though closer to the Mars
equatorial plane. The relation between these planes is illustrated in Fig. 9.
The most reliable orbital characteristics data, Wilkins (1965), are presented
in Table 6. The orbital elements of each satellite are given, using its fixed
(Laplacian) plane as a reference plane, which itself is referred to the standard
Earth equinox and equator plane of 1950.0 (conversion by Cain, 1967 b). The
original notation of Wilkins is explained in Table 7 and illustrated in Fig. 10.
According to Wilkins' analysis of all existing data, the secular acceleration
in longitude of Phobos, once claimed by Sharpless (1945), is unsubstantiated.
Related orbital data is presented in Table 8.
Physical Data
Physical data for the two Martian satellites are rather inadequate as
shown in Table 9- Recently, the size, shape and albedo of Phobos have been
estimated from several Mariner 7 TV pictures. Phobos is approximately 20 km
across (mean diameter), and elongated in the plane of its orbit, while possessing
an extremely dark surface (Smith, 1970).
The masses of these minuscule satellites are still unknown, as are their
densities.
October 15, 1971 R. Norton, C. Michaux, JPL Sec. I, page 23
Orbital and Physical Data
JPL 606-1
Table 6. Orbital elements of the Martian satellites. [Least-square solution
fitting all optical data, as proposed by Wilkins (1965) and converted to the epoch
1950.0 by Cain ( 1967 b), with minor modification* in A by Cain (1966).]
Elements for equinox
of 1950.0 and epoch
to = J. D. 2433282.5
(almost exactly
1950.0)
Longitude of node
of fixed plane
Inclination of fixed
plane to equator
Argument of node of
orbital plane at
epoch
Mean daily motion of
node of orbital plane
Inclination of orbital
plane to fixed plane
Mean longitude
at epoch
Mean daily motion
in longitude
Longitude of
pericenter at epoch
Mean daily motion
of pericenter
Apparent semi-major
axis at unit distance
Eccentricity of orbit
Symbol
N
A
K.
K
R
L,
'N
R
Phobos
46,9 deg
37.57 deg
177.28 deg
-1
-0,438 deg day
0.90 deg
136.21 deg
1128.8443 deg day
76 deg
0.436 deg day
-1
■ 1
12.926 ±0.001 arc sec
(9375.0 ±0.6 km)
0.018
Deimos
46,40 deg
36.64 deg
25,62 deg
-0,0180 deg day-1
1.80 deg
296.475 deg
285.16192 deg day"^
(236) deg formal,
e =
(0.016) deg day
formal, e =
-1
32.344 ±0.002 arc sec
(23,457.7±1.5 km)
0.0
*Note: Cain (1966) modified slightly the semi-major axis values A to take
into account the latest (Mariner 4) mass value of Mars.
Sec. 1, page 24
C. Michaux, JPL
October 15, 1971
JPL 606- 1
Orbital and Physical Data
Table 7. Nomenclature for the system of orbital elements necessary for
calculating a Marian satellite's position at anytime t after epoch t .
Orbital
element
Definition
N,
K,
M
Longitude (TA) of ascending node (A) of fixed (Laplacian) plane of satellite on Earth
standard equator (for 1950.0) [as measured along this equator from vernal
equinox T .]
Inclination of fixed (Laplacian) plane to Earth standard equator.
Longitude (AC) of ascending node (C) of orbital plane of satellite on its fixed
(Laplacian) plane, [as measured along this plane from ascending node A). It is
given by linear expression: K = K^ + Kj^ (t - tg), where to is the epoch
( J. D. 2433282.5, near 1950.0) and t the current Julian Date; therefore, (t - t )
in ephemeris days and K2 and Kj^ as below: °
Longitude K at epoch tg.
Mean rate of regression of ascending node C or of its longitude K (units: deg day" ).
Kj^ is constant and negative.
Inclination of orbital plane of satellite to Laplacian plane. I is constant.
Mean longitude (TA + AC + CS) of satellite (S), as measured along Earth
equator from its vernal equinox T to node A, then along Laplacian_plane to
node C, and finally along satellite orbit to satellite mean position S. It is given
by quadratic expression: L = L2 + Lj^j (t - t^) + Lj^ (t - to)^, where (t - to)
interval as above, and L^, Lji^ as below:
Longitude L at epoch to-
Mean rate of advance (mean daily motion) of satellite S around its orbit.
(Lf^ is considered constant at least over many years. )
Secular acceleration in longitude L of satellite. (Lj^ is, at most, extremely
small; in fact, according to Wilkins' analysis, Lj^ = 0.)
Longitude (TA + AC + CP) of pericenter (P) of satellite orbit, as measured
like L. It is given by linear expression: P = P2 + Pr (t - to), where (t - to)
as defined previously, and P^, Pr as defined below:
Longitude P at epoch to-
Mean rate of advance of pericenter P around the orbit, or of its
longitude P (units: deg day"l). Pp is constant and positive.
Semi -major axis of satellite orbit at unit distance or 1 AU. (Units used:
seconds of arc. )
Eccentricity of satellite orbit.
October 15, 1971
C. Michaux, JPL,
Sec. 1 , page 25
Orbital and Physical Data
JPL 606-1
MARS' EQUATORIAL PLANE* (Fixed)
SATELLITE lAPL^CIAN PLANE (Fixed)
SATELLITE ORBITAL PLANE (Precessing about
pole of Loplocian plane)
- ^AARS' ORBITAL PLANE (Fixed)
*AS DEFINED BY THE GRAVITATIONAL POTENTIAL OF MARS.
Fig. 9. The Laplacian plane of a Martian satellite (values of angles
are for Phobos).
ORBITAL PLANE
OF SATELLITE
VERNAL EQUINOX
(EARTH'S)
LAPLACIAN PlANi.
OF SATELLITE
EARTH EQUATORIAL PLANE
(CELESTIAL EQUATOR)
Fig. 10. Illustration of the notation used by Wilkins (1965) and Cain (1967) in
deriving the orbital elements of the Martian satellites.
Sec. 1, page 26
C. Michaux, JPL
October 15, 1971
JPL 606-1
Orbital and Physical Data
Table 8. Satellite orbital and related data.
Characteristic
Phobos
Deimos
Distance from center of Mars:
Mars radius = 1
(Req = 3393 km)
2.763
6.913
Kilometers
9375
2345 7
Miles
5825
14575
Distance above surface of Mars:
Kilometers
5982
20064
Miles
3717
12467
Period of revolution:
Sidereal (P = , ^ ,
n (or L^)
7S9'"l3f85
or ,
0.31891
lVl7'^54f87
or ,
1.26244^
Synodic
0.319^
1.265^^
Orbital eccentricity:
0.018 (±0.001)
0.0 (±0.0003)
Rate of regression of nodes:
-159.976° yr"^
-6.574° yr-1
or
-0.438° day-1
or
-0.0180° day-1
(Wilkins, 1965)
October 15, 1971
C. Michaux, JPL
Sec. 1, page 27
Orbital and Physical Data
JPL 606-1
fable 9. Physical data of Martian satellites.
Characteristic
Phobos
Dcimos
Diameters:
18 y 22
10 km
(elongated in
orbital plane)
Estimated
(Smith, 1970;
from Mariner 7
TV pictures)
Geometric albedo:
0.065
(average visual)
(Smith, 1970;
from Mariner 7
TV pictures)
Apparent visual magnitude from Earth:
11.6
12.8
(at mean opposition distance)
(Harris, 1961)
(Harris, 1961)
Color:
0.6
0.6
(magnitude B-V)
(greenish)
(greenish)
(Harris, 1961)
(Harris, 1961)
Sec. 1, page 28
C. Michaux, JPL
October 15, 1971
JPL 606-1 Orbital and Physical Data
BIBLIOGRAPHY
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Devine, C.J., 1967, JPL development ephemeris number 19- JPL Tech Ren
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October 15, 1971 C. Michaux. JPL Sec. 1, page 29
Orbital and Physical Data JPL 606-1
Buncombe, R. L. , 1964, Provisional ephemeris of Mars 1800-1950: U.S. Naval
Observ. Circular No, 95 (82 p.), April 30.
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Fish, F. F. , 1964, Precession of Mar s: Spaceflight, v. 6, no. 6, p. 214-215,
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Harris, D. , 1961, Photometry and colorimetry of planets and satellites, p. 272-
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1968, Constants and related information for astrodynamic calculations,
1968: JPL Tech. Rep. 32-1306 (57 p. ), July 15.
Sec, 1, page 30 C. Michaux, JPL October 15, 1971
JPL 6 06-1 Orbital and Physical Data
Miner, E. D. , 1967, (Pasadena, Calif, , Jet Propulsion Laboratory) : private
communication to R, Newburn.
Newcomb, S. , 1895, The elements of the four inner planets and the fundamental
constants of astronomy: Supplement to The American Ephemeris for
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heliocentric motion of Mars: Astronomical Papers prepared for the use
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Newton, I., 1687, Philosophiae naturalis principia mathematica: (The Mathe-
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Null, G. , 1971, (Pasadena, Calif. , Jet Propulsion Laboratory): private
communication to C. M. Michaux, January.
O'Handley, D. A. , Holdridge, D. B. , Melbourne, W. G. , and Mulholland, J. D. ,
1969, JPL development ephemeris number 69: JPL Tech, Rep. 32-1465
(33 p.), December 15.
Peabody, P, R, , Scott, J. F., and Orozco, E. G. , 1964, Users' description of
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October 15, 1971 C, Michaux, JPL Sec. 1, page 31
Orbital and Physical Data JPL 606-1
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Sec. 1, page 32 C. Michaux, JPL October 15, 1971
JPL 606-1
Orbital and Physical Data
APPENDIX A
EARTH-MARS CALENDAR
The attached Earth-Mars Calendar, Table A-1, correlates the Earth-Mars
periods in three categories of time: Earth calendar dates, Julian Dates, and
Consecutive Mars day(s).
The Consecutive Mars day entries are similar to the Julian day number
for Earth.
Table A-1. Earth-Mars Calendar
-"Tisor,
'Anr. va r day
Cojisecati'
Mars da'
([,-ap
1 J ( J .-■ 1 1 IJ
-1
Mar
Zi
! Z.HHJ^O
T
Apr
1
1 J4i.--lliJ
->
Apr
1 1
1 J M .^ 1 1 1 1
~
Apr
^1
i ^ 4 J .-^ ] T' )
-)
Mai
1
l-iin\un
Mav
1 1
^4'.;.170
-
Ma\
il
j ^ n s 1 ,4 1 )
■^
Ma%
31
1 Z.4 5.Kl^^ij
^
Jim
10
, ^ -M K ^ rj
Jun
10
1
sIj"mmf;r
1 z^]s^\■rl
s
Jun
50
■ -:-)5,i^^(j
->
Jul
10
j .M3>.>30
">
Jul
20
1 ^■lJHi^o
1
Jul
30
! J-l3H^Si)
T
Autt
g
■ iA^HZhl)
T
Aua
19
Z;>.nl70
-1 1
Auk
iu
1 i -1 ^ H Z K l>
s
S.p
H
i diihl'":
'
5,.p
FALL
IS
j J.nh^r..<,
S,.p
2rt
Jish.!'}
^
Oct
g
.M ^h iJii
Oct
IS
Zii
070
66S155
66
i4i
403
06SI63
40
l->i
133
668175
1 '
lb\
SOS
66H1S4
S7
271
60
668194
t9
2SI
335
66S2C4
3 5
291
063
66S214
06
300
79S
668223
80
-:-CliptlC
lontiitad.',
:!_.._._.
-ii .
310, 530
320.262
329. 995
339 727
349.460
359 192
SUMMER
368.925
378.657
388,390
398, 122
407 854
417, 587
427.319
437.052
446.784
456.517
466.249
475.981
485,714
495.446
505. 179
514.911
524.644
534.376
FALL
544. 109
553.841
563.573
573.306
583,038
668233 52
668243 26
668252,99
668262.73
668272.45
668282. 19
668291.92
668301.66
668311. 38
668321 . 12
668330.85
668340. 59
668350.31
668360.05
668369.78
668379. 52
668389.24
668398 98
668408.7 1
6684 18.45
6684^8. 17
668437 91
668447.64
668457.38
668467. 10
668476.84
668486. 57
668496 30
668506.03
119 52
121 04
128, 52
132. 97
137 39
141 . 79
146. IS
150, 55
15
. 92
159-29
163. 66
168 04
FEB 4. APIIELIC OFPO.SI
181
30
185
77
190
27
194
82
199
42
204
Oc>
208
77
213
53
218
36
223
27
228
25
233
31
2 38
45
24J, 69
249.01
254.43
259.94
265 55
271 25
277.05
262. 94
288.91
29-97
31. Mars not clc
physical obs
rly vus
■vation
April 1, 1967
R. Norton, JPI,
Sec. 1, Appendix A, page 1
Orbital and Physical Data
JPL 606-1
Table A-1. Earth-Mars Calendar (continued)
J.D.
Julian da
2-13B410.5
2458420-5
2438430 T
2438440,5
24Jb450 5
24384bO 5
2438470 5
2438480.5
2438490 5
2438500.5
2438510.5
2438520 5
2438530 5
2438540 5
2438550 5
2438560 5
2438570. 5
2438580. 5
2438590.5
2438600 5
2438610.5
2438620.5
2438630-5
2438640.5
2438650.5
2438660.5
2438670.5
2438680.
2438690.
243B700
2438710
2438720
2438730
2438740
2438750.5
2438760.5
5
2438770.
2438780
2438790
2438800
2438810
2438820
2438830
2438840.5
2438850. 5
2438860.5
2438870.5
2438880.5
2438890.5
2438900. 5
2438910.5
2438920.5
2438930.5
2438940.5
2438950.5
2438960 5
2438970 5
2438980.5
2438990.5
2439000.5
2439010.5
2439020.5
2439030.5
2439040.5
2439050.5
2439060.5
2439070 5
2439080.5
2439090.5
2439100.5
2439110.5
North
hc-niLsphvr
3 e a 9 o n
and date
n
WINTER
Jan 6
Jan 16
Jan 26
Feb 5
Fib 15
Ffb 25
Mar 6
Mar 16
SPRING
Mar 26
Apr 5
Apr 15
Apr 25
May 5
May 15
May 25
Jun 4
Jun 14
SUMMER
Jun 24
Jul 4
Jul 14
Jul 24
Aug 3
Aug 13
Aug 23
Sep 2
Sep 12
Sep 22
FALL
Oct
Oct
Oct 22
Nov 1
Nov 11
Nov 21
Dec 1
Dec 11
WINTER
Dec 21
Dec 31
2
12
Jan 10
Jan 20
Jan 30
Feb 9
Feb 19
Mar 1
Mar 11
SPRING
Mar 21
Mar 31
Apr 10
Apr 20
Apr 30
May 10
May 20
May 30
Jun 9
Jun 19
SUMMER
Jun 29
Jul 9
Jul 19
Jul 29
Aug 8
Aug 18
Aug 28
Sep 7
Sep 17
FALL
Sep 27
Oct 7
Oct 17
Oct 27
Nov 6
Nov 16
Nov 26
Dec 6
Dec 16
WINTER
Dec 26
1964
(leap
year)
Nor!!u-ni
lenii sphiTf
Const
Mar
and
lOOO
(leap
year)
FALL
592.771
602.503
612.236
621.968
631.701
641.433
651. 165
660.898
1.630
WINTER
11.363
21.095
30.828
40.560
50.293
60.025
69.757
79. 490
89.222
98 955
108. 687
118.420
128. 152
137 885
147 617
157.349
SPRING
167 08 2
176.814
186.547
196.279
206.012
215.744
225.477
235.209
244.941
254.674
264.406
274. 139
283.871
293.604
303.336
313.068
322.801
332.533
342.266
351.998
SUMMER
361 .731
371.463
381 196
390.928
400.660
410. 393
420. 125
429. 858
439.590
449.323
459.055
468.788
478.520
488.252
497.985
507.717
517-450
527. 182
536.915
FALL
546.647
556 380
566. 112
575.844
585.577
595.309
605.042
614.774
624.507
668515.77
668525.50
668535.23
668544.96
668554.70
668564. 43
668574. 16
668583 89
668593.63
668603.36
668613.09
668622.82
668632.55
668642.29
668652.02
668661 .75
66867 1 48
668681 . 22
66869U 95
668700. 68
668710.41
66B720. 15
668729.88
668739-61
668749. 34
668759.08
668768.81
668778 54
668788.27
668798 01
668807 74
6688 17 47
668827.20
668836.94
668846.67
668856.41
668866. 13
668875.87
668885.60
668895. 34
668905.06
668914.80
668924. 53
668934. 27
668943.99
668953.73
668963.46
668973.20
668982 92
668992. 66
669002. 39
669012. 13
669021.85
669031.59
669041.32
669051.05
669060.78
669070.52
669080.25
669089.98
669099.71
669109 45
669119 18
669128.91
669138.64
669148 38
669158 11
669167-84
669177 57
669187. 30
669197.04
669206 77
661216. 50
rchplii
j.iLjitud.-
301 10
307.29
313. 54
319.83
3 26. 15
332.50
338-65
345. 19
351.51
357.79
4.03
10. 22
16. 34
22. 38
28 35
34.23
40.01
45.71
51.31
56.8 1
62.22
67 54
72 77
77 91
82 96
87. 94
97
67
102
43
107
13
111
77
116
37
120
91
125
42
129
89
134
33
138
75
143
14
147
52
151
90
156
26
160
63
165
00
169
39
173
79
178
21
182
66
187
14
191
66
196
22
200
83
205
50
210
22
215
00
219 8t:
224 78
229. 79
234.87
240.05
245.31
250.66
256. 11
261 65
267 29
273 03
278 H6
284. 77
290.77
296.85
303 00
309 II
315. 48
321 7H
I -Jul I - MarB not clcarlv
physical observa
Sep 12 . Mars dl67 ylOOO . appro.i 25 of Mars vli;oO .
Mars vernal rqiunox prior to Maritur IV Uvb>' a
basis of zero points uf this calendar
Nov 5. Marin
Nov 26 . Mariner IV launched {Alias- Agena vehi cle)
Niv 30. Zond 2 launched; unsuccessful mission
MAR 9 APIIELIC OPPOSITION
Earth-Mars distance at clos
62.0 million ml (99.8 niiUio
Earth days until next opposil
Man
ntr [V Mars f
lakt-r
abovi' both lu
{iti^r
L-bsl OLCurri-tl
isphe
j:, p.itures
L- citation
, henilsphe
Sec. 1, Appendix A, page 2 R. Norton, JPL
April 1, 1967
JPL 606-1
Orbital and Physical Data
Table A-1. Earth-Mars Calendar (continued)
EARTH
MARS
Norlh.-rn
Northi-rn
ll.-liore,itru
J . u
h.niisph.i-c
bomi.spherf
Cons.-iuUve
eclifjlic
Julian day
s.aaon
Year
Year
.season
Mars day
lonnilude,
and <lat.-
and vf-ar day
d, u
WINTER
1966
1000
FALL
2439130 S
.Ian 5
(leap
634 239
669226 23
328, 12
2439140 S
Jan 15
year)
643.972
669235 97
334,46
2439150 S
Jan 25
653.704
669245 70
340,81
2439160 5
2439170 5
Feb 4
Fe b 14
663.436
669255,43
347.15
1001
4. 169
669265, 16
353.46
Oeap
WINTER
2439IBC i
F.b 24
year)
13.901
669274, 90
359.73
2439190, '■>
Mar 6
23.634
669284,03
5,95
2439200 5
Mar 16
SPRING
33.366
669294 36
12, 12
2439210 -■>
Mar 26
43.099
669304,09
18.22
2439220 5
Apr 5
52.831
669313.83
24,24
2439230, 5
Apr 15
62.564
669323,56
30. 17
2439240 5
Apr 25
72.296
669333.29
36.02
2439250, 5
May 5
82.028
669343.02
41.78
2439260, 5
May 15
91.761
669352.76
47.44
2439270,5
May 25
101.493
669362,49
53.01
24392H0, 5
Jun 4
111.226
669372,22
58.49
2439290, 5
Jun 14
SUMMER
120.958
669381.95
63.87
2439300,5
Jun 24
130.691
669391.69
69.16
2439310 5
Jul 4
140,423
669401.42
74.36
2439320 5
Jul 14
150.155
669411. 15
79^47
2439330,5
Jul 24
159. 888
SPRING
669420.88
84,50
2439340, 5
Aug 3
169.620
669430.62
89,45
2439350 5
Aug 13
179.353
669440,35
94.33
2439360,5
Aug 23
189.085
669450,08
99, 14
2439370 5
Sep 2
198,818
669459,81
103.88
2439380, 5
Sep 12
208.550
669469.55
108.56
2439390,5
Sep 22
FALL
218.283
669479,28
113.19
2439400.5
Oct 2
228,015
669489, 01
117,77
2439410.5
Oct 12
237,747
669498.74
122, 31
2439420.5
Oct 22
247.480
669508,48
126,80
2439430,5
Nov 1
257.212
669518.21
131,26
2439440. 5
Nov 11
266.945
669527,94
135.70
2439450.5
Nov 21
276.677
669537,67
1 40 , 1 1
2439460.5
Dec 1
286.410
669547,41
144, 50
2439470.5
Dec 11
296, 142
669557. 14
148,87
2439480.5
Dec 21
WINTER
305,875
669566.87
153, 24
2439490.5
Dec 31
315.607
325.339
669676.60
669586,34
157,61
161,98
2439500.5
Jan 10
1967
2439510.5
Jan 20
335.072
669596.07
166, 35
2439520.5
Jan 30
344.804
669605.80
170.74
2439530.5
Feb 9
354.537
SUMMER
669615.53
175, 15
2439540.5
Feb 19
364.269
669625.27
179,58
2439550.5
Mar 1
374.002
669635.00
184,04
2439560.5
Mar 11
SPRING
383.734
669644,73
188. 53
2439570.5
Mar 21
393,467
669654,46
193.06
2439580.5
Mar 31
403, 199
669664.20
197,63
2439590,5
Apr 10
412.931
669673,93
202,26
2439600,5
Apr 20
422.664
669683,66
206.94
2439610,5
Apr 30
432.396
669693 39
211 ,67
2439620.5
May 10
442.129
669703, 13
216.48
2439630, 5
May 20
451.861
669712,86
221. 35
2439640,5
May 30
461.594
669722,59
226, 30
2439650, 5
Jun 9
471.326
669732,32
231.33
2439660,5
Jun 19
SUMMER
481.059
669742.05
236.44
2439670 5
Jun 29
490,791
669751.79
241.64
2439680, 5
Jul 9
500,523
669761.52
246.92
2439690.5
Jul 19
510.256
669771.25
252.30
2439700,5
Jul 29
519.988
669780.98
257.78
2439710,5
Aug B
529.721
FALL
669790,72
263.35
2439720,5
Aug 18
539.453
669800.45
269 02
2439730 5
Aug 28
549,186
669810. 18
274.78
2439740, 5
Sep 7
558.918
669819.91
280.64
2439750, 5
Sep 17
FALL
568,650
669829.65
286.58
2439760,5
Sep 27
578 383
669839, 38
292,61
2439770,5
Oct 7
588. 115
669849. 11
298,71
24397H0 5
Oct 17
597,848
669858.84
304.88
2439790, 5
Oct 27
607.580
669868.58
311.11
2439800,5
Nov 6
617,313
669878, 3r
317,39
2439810,5
Nov 16
627.045
669888.04
323.71
2439820.5
Nov 26
636.778
669897.77
330,04
2439830,5
Dec 6
646.510
669907.51
3 36.40
2439840,5
Dec 16
WINTER
656.242
669917. 24
342.7-1
2439B50.5
Dec 26
665.975
669926,97
349 08
Jan 2 -Oct 1 . Mars
phytic
lot clearly visihl.' from Earlli
al obbt-rvatioiis .
APR ]b
OPPOSITION
Earth-Marb c
5S 8 niillion
Earth dayh ur
April 1, 1967
R. Norton, JPL Sec. 1, Appendix A, page 3
Orbital and Physical Data
JPL 606-1
Table A-1. Earth-Mars Calendar (continued)
Z439870
2439880,
2439890.
243<"90i).
2439910
2439920
2439930
2439940
2439930.
2439960
2439970.
24399fiO.
2439990
2440000.
2440010.
2440020.
2440030.
2440040
2440030
2440060
2440070
2440080.
2440090.
2440100.
24401 10.
2440130
2440140.
2440150
2440160.
2440170
2440180
2440190
2440200
2440210
2440230.
2440240.
2440250.
2440260.
244C270.
2440280.
244l.'29C.
2440300.
2440310.
2440320.
2440330.
2440340.
2440350
2440360
2440370
2440380
2440390
2440410 5
2440420. 5
2410430 5
2440440. 5
2440450 5
2440460. 5
2440470 5
2440480.5
2440490 5
2440300 5
2440510 5
2440520. 5
2440530 5
2440540 5
2440550.5
2440560.5
2440570.5
>-nn»|)lu-r..
aii-l il.-it,-
WIN I t:R
Far.
Fi-b
Ma
Ma:
24
5
15
SPRING
Mar 25
Apr 4
Apr 14
Apr 24
May 4
May 14
May 24
Jiin 3
Jun 13
SUMMER
Jun 23
Jul 3
Jul 13
Jul 23
Aug 2
Aug 12
Aug 22
Sep
Sep
Sep
FALL
Ott 21
Ocl 31
Nov 10
Nov 20
No\
Dec 10
Dec 20
WINTER
Dec 30
Jan 9
Jan 19
Jan 29
Feb 8
Feb 18
Feb 28
Mar 10
SPRING
Mar 20
Mar 30
Apr 9
Apr 19
Apr 29
May 9
May 19
May 29
Jun 8
Jun 18
SUMMER
Jun 28
Jul 8
J ul 18
Jul 28
Aug 7
Aug 17
Aug 27
Sep 6
Sep 16
FALL
Sep 26
Oct 6
Oct 16
(3ct 26
Nov 5
Nov 15
Nov 2 5
Dec 5
Dec 15
WINTER
1968
(leap
year)
WINTER
16.440
55. 370
65, 102
74 .834
94
104
113
123
133
142
152
162
SPRING
172
181
191
201
21 1
220
230
240
250
259
299
032
764
497
229
962
694
426
891
624
356
089
82;
554
286
018
751
1003
(leap
year)
269.483
279.216
28S.948
298 68 1
308,413
318. 146
327.878
337.610
347. 343
357.075
SUMMER
366.808
376.540
386.273
396.005
405.737
415.470
425.202
434.935
444.667
454.400
464. 132
473.865
483.597
493 329
503.062
512.794
522.527
532 259
FALL
541.992
551.724
561.457
571. 189
580.921
590-654
600.386
610. 1 19
619 851
629. 584
639. 3 16
; 49.049
^f-^iilL
513
10 24b
WINTER
19 97 8
Coi,^,
Mill
6699.16. 44
t-69956. 17
669V65 . 90
669975. 63
669985. 37
669995. 10
670004,83
670014, 56
670024. 30
670034.03
670043.76
670053.49
670063.23
670072.96
670082 69
670092.42
670102. 16
670111 .89
670121.62
670131 . 35
670141 .09
670150.82
670160.55
670170.28
670180 02
670189.75
670 199 48
670209 21
670218 95
6702 28. 68
670238 4 1
670248- 14
670257-88
670267-61
670277- 34
670287.07
670296 80
670306. 54
670316.27
670326.00
670335.73
670345.47
670355. 20
670364.93
670374.66
670384 40
670394 13
670403.86
670413.59
670423-33
670433.05
670442.79
670452.52
670462.26
67047 1 .98
670481 72
670491 .45
670501 19
670510.91
670520.65
670530. 38
670540. 12
670549, 84
670559. 58
6705b9. 31
67(i57w. 05
6 70588.77
670598.51
670608. 24
670627.70
67U6'7 44
cliptii
:i;U,,de,
7.85
14.00
20 . 07
26.07
31.98
37.81
43.54
49. 18
54.72
60. 16
65.52
70.78
75.95
81.04
86.04
90.97
95. S3
100.61
105.34
110.00
114.62
119. 18
123.70
128. 19
132.64
137 06
141 ,47
145.85
150.23
154.60
I 58 . 96
163. 33
167.7 1
172. 11
176.52
180.96
185.42
189.93
194.47
199.06
203.70
208.40
213. 16
217.98
222.87
227.85
232.90
238.03
243.26
248.57
253.98
259.48
265.08
270
78
276
56
282
44
288
41
294
46
300
58
306
77
313
01
319
30
325
63
3 3 1
97
138
32
344
6b
3 50
99
Apr 3 -5<ip -
Mars not cltarly visible
phyaical observations.
Feb 23-Apr H . Mariner-'69 launch opportanity (Atlas-
Ctntaur vehicle) Present plana allow
for the launching of two spacecraft, one
during the early and one during the late
launch period wUhin this opportunity.
MAY 31. OPPOSITION
Earth-Mars distance at close at approach:
44. 6 million ini (71,7 million km)
Earth days until next opposition: 7H 1
Jul ^9-Au^i 15 . EaiTiest-latest Marw;
dates for Feb ^3-Apr
nity . Pre-s.-iit plans 1
flyby for one npatecr
torial flyby for ihe ot
f'i9 Mars llyby
unch opport u -
Sec. 1, Appendix A, page 4 R. Norton, JPL
April 1, 1967
JPL 606-1
Orbital and Physical Data
Table A-1. Earth-Mars Calendar (continued)
i.l40hjij, ^
2.140hjli S
^440fi40, 'i
^44;irt^ii ■.
Z44UrtM-J. T
i44l)»70 S
Z440HriO. S
Z440S<,10 S
Z4 4
<;44n')if.
s
Z4404/I)
Z440O3I)
2440941)
2440'll.
2441)'I7
FALL
o< t
0( t
1 1 ,-1 J:)
1')
Z44';'.'''IJ
F,-l»
K
Z44 !i!lji)
f,
IH
244 1'J!0
T
1-7
2H
244L-'20
S
Ma)
10
2441030
"^
Mri!
SPRINOO
20
244 1040
T
Mai
30
244 DOOt
s
Apr
244 lot,;)
s
Apr
19
244 1070
s
Apr
29
244],0,HO
s
May
9
244 1000
■>
May
19
244 1 100
•y
May
29
244 1 ] 10
S
Jiin
H
244 1 120
s
JllTl
1«
SUMMER
244 1 1 30
s
Ji.n
in
244 1 140
Jul
8
244 1 I'^O
■j
Jul
1«
244 1 lf.O
s
Jul
28
2441 170
s
Auk
7
244 1 IBO
s
Aue
17
2441 luo
s
Auk
27
244 1200
s
S,p
6
244 12 10
s
S,-p
FALL
16
244 1220
s
,S,-p
26
244 1230
s
Oct
6
244 1240
^
Oct
16
244 12S0
s
Oct
26
24,1 12('>0
s
Nov
S
2441270
s
Nov
J5
2 4412«0
s
Nov
25
i;4412,>0
1),-.
T
2441 SiiO
"■
i)l c
WIN 1 f.H
1 s
244 1 ill)
1), c
2S
_ . .
^-!4':7 lu
M;.\ -i
1^]-V>7^'''
M;,V 14
Z-44';7iiJ
M.sv -■!
i-140740
S
Jii.i i
^-44 07^0
s
SUMNU.H
i-14(!7':U
T
.1 , = : Zi
Z44'i77l,i
s
;.i i
7h
-i? J
Hh
iO^
^7
H in
1)7
S7(J
17
irn
127.030
1 36.7 08
14 6.000
106,233
160.060
SFRING
175, 697
16 5.430
195. 162
204.895
214.027
224.360
234.092
243.824
253.557
263.289
273.022
2B2.754
292.487
302.219
3 1 1.952
321 . 684
3 31.416
341. 149
350.88!
360.614
SUMMER
370.346
3B0.079
389-81 1
399 544
409.276
419.008
428.74 1
438.473
448 . 206
457.938
467.671
477.403
487. 136
496.868
506.600
516.333
526.065
535.798
FALL
545.530
555.263
564.995
574.728
584.460
594. 192
603.925
613.657
623.390
633. 122
642.855
652.587
662^
3.
WINTER
12.
22.
32.
41.
51.
61
320
784
517
249
982
714
447
6 7 0t,6rj . »_, 1
(,70676. 37
670680. 10
tj70<,95. 84
670705. 56
6707 15. 30
670725.03
670734.77
670744. 50
670754.23
670763. 96
670773.70
670783 43
670793 16
670802.89
670812.63
670622. 36
670832.09
670841 . 82
670851.55
670861.29
670871.02
670880.75
670890.48
670900.22
670909.95
670919.68
670929.41
670939. 15
670948. 88
670958.61
670968.34
670976.06
670987.80
670997.54
671007.27
671017.01
671026.73
671036.47
671046. 20
671055.94
671065.66
671075.40
671085. 13
671094.87
671104.59
6711 14.33
67 1124.06
671133.80
671143.52
671153.26
671162. 99
671172.73
671182.45
671192. 19
671201.92
671211.66
671221. 38
671231. 12
671240.85
671250. 59
671 260. 31
671270,05
671279.78
67 1289 52
671299. 24
671308, 9«
671318 71
67 1328 45
67 1338 17
07 1)47 01
11141
1 16, <H
120
58
125, O'l
129. 57
134. 01
1 18. 43
142. 82
147. 21
155. 04
160. 3 1
164 68
ri, 46
^7. 88
182. U
186. 81
195. 88
200. 48
205. 14
21)'*. 85
214. 6 1
21 9, 48
224.40
220. 30
234. 47
2 19. 6 1
244. 88
2511. 2i
255. 66
261. 20
266. 8 1
i'H'i
11:
29 1,
M
M)i
4H
M\H
(>H
iU
'M
Hi
24
iJ7
S7
M i
42
3-10
26
H6
60
iSJ
'•2
AUG 10 Pi-:Hiiii:LK: OHi'osi CKJr;
t;a!tli -M,.i- -11^1. ■;.. t ,it , l.>~.. -.[
it 'f icuUmm, UN (Si.. 2 miMit,,, k.
April 1, 1967
R. Norton, JPL Sec. I, Appendix A, page 5
Orbital and Physical Data
JPL 606-1
Table A-1. Earth-Mars Calendar (continued)
;i 11 iii), =
i4-ll,i^0, =
^4-41-1'iU '
^■H 1 ^7U -
Z44UBI), -
^441 J'>U . =
241140',). ■
J44 14 10. =
i4414Z0.'
2441430- '
2441440,-
24414'iO '
244 14b0 '
2441470 '
244I4K0 '
24-11 I'll! ■
244 1 SOU . ■
2441S10.
2441=>2C.
2441'iJO.
2441S4n
2441SS0
244n(/0
244 1S7I
244 15hii
244 15',lil
2.14 U.!)0
2.H1'>10.
2441(i20.
Z441t,i0.
244 Ui40.
244 16S»
2441660
244 1670.
M.,v
13
.M;.y
23
Jim
12
.SI
mmf;i
.l.iri
Jul
22
Jul
12
244 16<J0
2441700
2441710.
244 1720.
2441710.
244 1740
24417S0.
2441760.
A.iK 11
.S,-p
S>,p
ALL
10
20
S.,p
30
Oit
10
o> t
20
o> t
30
Nov
9
F.-b 7
F.-b 17
Fi'b 27
Mar 9
Mar 19
SPRING
Ma
29
1072
(l-up
24417H0, ti
Apr
H
2441790. '->
Apr
IB
244 1M00 ■>
Apr
2B
2441S10 5
May
H
2441H20. 5
May
1«
2441S30 1
May
dH
2441840 5
Juii
7
2441«i0. S
Juii
17
SUMMFR
2441H60. T
Ju.i
27
2441K70. ■>
Jul
7
2441SH0 S
Jul
17
2441H90. S
Jul
27
244 1900. 1
AuB
6
244 1910 6
Aug
16
244 1020 ^
Aug
26
2441O30 =.
.S,.p
5
244 1940 •>
S,.p
FALL
!•>
24419S0. '^
.S.p
2S
244 1OP.0 . T
Oi t
S
244 1970 i
Oit
1^
244|OhO s
Of 1
2S
244 1OM0 ^
Mo-,
4
24421)00 =;
Nov
n
24420 1'.) 1
rjov
24
24421)20 ^
D.-c
-1
24 120)0
D. I
WIN' I'LH
14
24,12,;i,Iil . ^
n,,<
24
129, T,' 4
67 1 l"'i . ->7
68.74
1 10, 300,
r.7 1 tO'i . 31) 7 1. u5
140-03'.
6.7 14 1'.. 1
7'' 01-
15H,77 1
'.7 l-t2'i . 77
84 1 1
SPRING
108, 00 i
1.-7 1435. 5U 80. 07 j
17 8,216
67 1415. 21
1. ''5
1 8 7 , ',' 6 8
r,7 14-4. '16
'8.77
197,70 1
6714'.,!. 70
10 1. 52
207,411
67 1474., 13
108. 21
217, lob
07 M8,l , 16
1 12. 8-1
226 , 89H
6714'J3,89 ■ 117.4) !
236,611
6715(,.3,63 121. "7 I,
246, 363
671511, 16
126.47
2'~t.,ou3
67 1521, o'l
1 10. M
266.828
6715 12,82
1 15. 17
275.560
071512, 55
1 I'l. 78
285.293
671552, 29
144. 17
295.025
671562,02
14 8. 5 5
304.758
671571, 75
152.92
314.490
671581.48
157. 1"
324.223
671591.22
16 1. 60
3 3 3.955
67 1660 95
166. 1
143. ')87
671610.68 , 170.42 j|
153.420
67 16 20. 41
17 4. 82
.SUMMER
S63. 152
671610. 15
179. 25
372.885
671619. 88
181. 70
182. 617
67 1649. 62
188. 10
l',)2. 150
671659 34
19 2. 7 2
402. 0H2
671669.08
197. 29
.111.8 15
671678- 81
201. ''!
421 . 547
67 1688-55 206.58 j|
431.279
671(0)8- 27
211. 11 1
441 .012
671708.01
2 16. 11
450. 7. 14
67 17 17 . 74
220. '17
460.477
671727.48 ' 225.''! ;
470.209
671737.20 , 210.''! j
479.942
67 1746.94
216.0,4 1
489 674
67I75t. .67
2-11. 11
499.407
671766.41
2-16. 50
509 139
671776. 13
251. 87
518 871
671785.87 ' 257. 14 jj
528. 604
67 17''5.60
262. '10
FALL
538.336
671805.34
268. 56
548-069
671815.06 1 274. !1 II
557-801
67 1H24. 80
280. 16
567 534
67 18 14. 53
286. 0''
577 - 266
07 1844. 27
AH. 11
586. '198
67 1653 -'19
2') 8. 2"
5'I6.73I
67 186 1-73
10.1. 17
606.46 1
67 1^7 1-46
1 10. 5''
6 16. 196
67 1883. 20
116. 86
025. '>28
67 1802.92
121. 18
635.66 1
67 1''02 . 66
12'). 51
645.193
671')12.39 i 115. 8(,
655. 126
671022 11
1-12. 21 I
864 . 858
67 I'll 1.85
148. 54 I
loo"^
6. 5')0
"^ 67 104 1-59
15 1. 85
(l.-ap
WINIER
16.121
671"5I . 32
1. 1 1
26 .055
67 1'>61 .05
7. ) 1
15. ,'88
67 1''70 7H
I'.OH
4 5.520
67 1''80. 52 ' I''- 5'.
: 55.251
67 1'"<0.25
25- 57
103
1 1 1
[^IM'OSI IK)?,
Sec. 1, Appendix A, page
R. Norton, JPL
April 1, 1967
JPL 606-1
Orbital and Physical Data
Table A-1. Earth -Mars Calendar (continuc-d)
^ 44 .'.070
.i 4 ■)."■) so
^44.;f;'>it
Z4 4^ K)0.
mil )U.5
^44^140, 5
i4421=.D S
244^1bD 5
^44^170 5
.i44ilK0. T
.!4421')0 ■>
i-l-iiiOO. 5
244^<:i0. 5
i14iiilJ. 5
^44iZJ0 S
i442240, S
^44^^^o 5
2442^60.5
2442i70. 5
HiiiM). 5
244Z290. S
2442300, 5
2442310 S
2442j20. 'j
2442330. 5
2442340, ■>
24423^10 5
2442360, t)
2442370.5
24423H0 S
2442390 ".
2442400. 5
_244^4in JS_
"2442420, 5
2442430, 3
2442440, ■-.
2442450, 5
24424f>0. 3
2442470, 5
24421H0, 5
2442490 5
2442500, 5
2442510 5
2442520, 5
2442530, 5
2442540, 5
2442550 5
2442560,5
2442570 5
2442580 5
2442590, 5
2442600, 5
2442610, 5
2442620,5
2442630,5
2442640 5
2442650, 5
2442660,5
2442670, 5
2442680, 5
2442690, 5
2442700,5
2442710,5
2442720,5
2442730, 5
2442740 5
2442750 5
2442760 5
2 4427 (0,5
K,\Rl-!\
M -XR,^
N..rt],. i-n
Norll,.-rr,
ll.-liui, ntr.c
Ji.Mni sp!,.,i-r
Y.-ar
Y. a.
!,,-un-,r.!„-r,
c:oii-..,.-iiii\'.
■ rhplii
Mar h ilay
lun^^itiirb- ,
,..L<1 .la!.-
and y-ar .!:i-,
il.-u
wiK rf:H
197 4
1005 j WINJ'KR
Jau 3
(leap 1 13i 112
67 2068, 1 1
70. 34
,Tar, M
y.-arl : 142 845
f.72077,84
75. 5(
,T a r, 2 3
152,577
672087 , 57
KO. (,i
FrI, 2
162 310
SPRING
67209?, 30
8 5. 6,4
Feb 12
172,042
672 107 04
90, 57
Frb 22
181 , 774
0721 16,77
05. 44
Mai 4
191 ,507
67 2 126 50
100, 23
Mar 14
201,239
1.721 36, 23
Ui4. 96
S 1>R 1 NG
Mar 24
210.972
672145, 97
109. 64
Apr 3
220,704
67 2 155.70
1 14. 2 5
Apr 13
230,437
672165.43
1 18. 83
Apr 2 3
240, 169
672175. 16
121, J5
May 3
249.902
672184.90
127. 84
May 13
259.634
672194.63
132. 30
May 2 3
269-366
672204.36
136. 72
J.m 2
279,099
672214.09
141 13
Jun 12
288,831
672223.83
145. 52
.s[jmmf:r
Jun li
298,564
672233.56
149, 8')
J.il 2
308.296
672243.29
154 26
Jul 12
318.029
672253.02
158, 63
Jul 22
327.761
672262.76
163. 00
Aay 1
337.493
672272.49
167. 36
Aog 11
347.226
672282.22
171. 77
Aug 21
356.958
SUMMLR
672291.95
176. 18
Aug 3 1
366, 691
672301.69
180. 62
Sep 10
376,423
672311.42
185. 08
Sep 20
386, 156
672321 , 15
189, 58
FALL
Sep 30
395, 888
672330.88
194. 12
Ocl 10
405,621
672340,62
198, 71
Oil 20
415,353
672350.35
203. 34
Ocl 30
425,085
672360,08
208. 03
Nov 9
434,818
672369 81
212, 79
Nov 19
444,550
672379, 55
217, 60
Nov 29
454,283
672389, 28
222. 49
Dec 9
464, 015
672399,01
227. 45
D.-c 19
473,748
672408, 74
232. 50
WINTER
Dec 29
483 480
493 213
672418,48
672428, 21
2i7, 62
242, 84
Jail 8
1975
Jan 18
502,945
672437,94
248. 15
Jan 28
512,677
672447,67
253. 54
F. 1, 7
522,410
672457,41
259.04
Fib 17
532, 142
FALL
67 2467, 14
264.63
Feb 27
541 ,875
672476.87
270. 31
Mar 9
551,607
672486, 60
276.09
Mar 19
561 ,340
672496, 34
281,96
.SPRING
Mar 29
571,072
672506.07
287.92
Apr 8
580,805
672515.80
293.96
Apr IB
590,537
672525.53
300, 07
Apr 28
600,269
672535.27
306. 26
May 8
610.002
672545.00
312. 49
May 18
619.734
672554.73
318. 78
May 28
629.467
672564.46
325. 10
Jun 7
639, 199
572574.20
331. 44
Jun 17
648.932
672583.93
3 37.79
SUMMER
Jun 27
658.664
672593,66
344. 13
Jul 7
668.397
672603,39
350. 46
Jul 17
1006
9. 129
672613. 13
356, 76
WINTER
Jul 27
18.861
672622.86
3.0!
Aus 6
28.594
672632.59
9. 21
Aug 16
38.326
672642.32
15. 34
Aug 26
48,059
672652.05
21. 40
Si-p 5
57.791
672661.79
27, 38
Sop 1 5
67.524
672671.52
33, 28
FALL
Sep 25
77,256
672681.25
39. 08
Oct 5
86.989
672690.98
44. 80
Oct 15
96,721
672700.72
50. 42
Oct 25
106,453
672710.45
55, 94
Nov 4
116, 186
672720. 18
61. 37
Nov 14
125.918
672729.91
66. 70
Nov 24
135,651
672739.65
71.95
Dec 4
145.383
672749, 38
77. in
D.-c 14
155, 1 16
672759 1 1
8.^ 17
WINTFR
D.-c 24
l',,4 , 8-18
67276H Ht
87 H,
_ .
DFX 15, QPH-)SllION
t:arrh-M;.r -^ di
S^.4 milhur. n
ill- at closist approach:
1 1 n\iIlion km).
\t opposition; 77 i ,
April 1, 1967
R. Norton, JPL Sec. 1, Appendix A, page 7
Orbital and Physical Data
JPL 606-1
Table A-1. Earth-Mars Calendar (continued)
i':y\R rn
- --. ■ - ■ , -—- ii-
North. TH
J D
beiiusphor.'
Year
>
Julian dav
season
a.id date
WL\' rr.R
V>T,
2442-K0.5
Jan 3 1 (leap 1
i44i.'IO.S
J,,n IS 1 year) 1
a442800,5 J:in Z) | 1]
.!44Z>ilD 5
F.;b 2 1
J.4428i0,5
Feb 12 ; il
2442830, 5
Feb 22 j
2442840 , 5 Mar 3 |
2442850 S
Mar 13
SPRING
2442860 S
Mar 23
2442. 70. S
Apr 2
2442880. S
Apr 12
2442 "'10. 5
Apr 22
2442900.5
May 2
2442910 5
.May 12
2442920 5
May 22
244. 130. 5
Jun 1
2442940 5 Jun 11
SUMMER
2442950. 5
Jun 2 1
2442960 8
Jul 1
2442970.8
Jul 11
2442980. 8
Jul 21
2442990 5
Jul 31
2443000. 5
Aug 10
2443010.5
Aug 20
2443020 5
Aug 30
2443030. 5
Sep 9
2443040.5
Sep 1 9
FALL
2443050 5 ■ S.-p 20
2443080.5
Oct 9
2443070.5
Oct 19
2443080.5
Oct 29
2443090.5
Nov 8
2443100.5
Nov 18
2443110, 5
Nov 28
2443120.5
Dec 8
2443130.5
Dec 18
WINTER
2443140.5
Dec 28
2443150.5
Jan 7
1977
2443160. 5
Jan 17
2443170.5 Jan 27
2443 180.5 Feb b
2443190.5 K.-b 18
2443200.5 ' Ft?b 26
2443210.5
Mar 8
2443220 5
Mar 18
SPRING
2443230.5
Mar 28
2443240. 5
Apr 7
2443250.5
Apr 17
2443260. 8
Apr 27
2443270.5
May 7
2443280. 8
May 17
2443290, 5
May 27
2443300, 5
Jun 6
2443310. 5
Jun 16
SUMMER
j
2443320.5
Jun 26 ]
2443330.5
Jul 6 i
2(43340.5
Jul 16 ! 1
2443350,5
Jul 26
2443360. 5
Aug 5
2443370. 5
Aug 15
2443380.5
Aug 25
2443390.5
Sep 4
2443400.5
Sep 14
FALL
2443410. 5
S.p 2 4
2443420.5
Oct 4
2443430. 5
Oct 14
2443440.5
Oct 24
2443450.5
Nov 3
2443480.5
Nov 1)
2443470 5
Nov 2 3
2443480.5
Dec 3
2443490.5
Dec 13
WINTER
2443500 5
"■"." J
i_
1007
(leap
MARS
Nurlhern
hemisphere
Consecutive
seasori
Mars day
and y ear day
SPRING
174, 580
67 277H, !)8
184 313
672788-31
194 045
6727^18. 04
203.778
672807,77
2 13 5 10
672817, 51
223.243
67 2827,24
232.975
672836,97
242.708
67 2846. 70
252.440
672856.44
262. 172
672866. 17
271 . 905
672875.90
281.637
672885.63
291 .37
672895.37
301. 102
672905. 10
310.835
672914.83
320.567
672924. 56
330.300
672934. 30
340.032
672944.03
349.764
672953.76
359 497
672963.49
SUMMER
369.229
672973.23
378 962
672982.96
388.694
672992.69
398.427
673002.42
408. 159
673012. 16
417.892
673021.89
427.624
673031.62
437.356
673041. 35
447.089
673051.09
456.821
673060.82
466.554
673070.55
476.286
673080. 28
486.019
673090.02
495.751
673099.75
505.484
673109, 48
515.216
6731 19. 21
52, 948
673128.95
534.681
673138.68
FALL
544.4 13
673148.41
554. 146
673158. 14
563.878
673167.88
573.611
673177.61
583.343
67 3187. 34
595.076
673197.07
602 808
673206.80
612, 540
622.273
632.005
641 .736
651 .470
661 203
.._._,!-_
2.935
WINTER
12.667
22.400
32. 132
41 .865
51.597
61 . 330
71.062
80.795
90.527
100.259
109.992
119 724
129.457
139 189
148 922
158. 654
SPRING
168.367
178. 1 19
187 851
197,584
207, 3 16
673216, 54
673226.27
673236.00
67 3245.73
673255.47
673265.20
llehocenlri
ecliptic
longitude.
deg
92.
J8
96.
9 2
101.
70
106.
41
111.
06
115.
67
120.
22
124.
74
129.
21
133.
66
138.
08
142.
48
146.
86
151.
24
155.
60
159.
97
164.
34
168
73
173
12
177
54
181
99
186
46
190
97
195
53
200
13
204
78
209
49
214
26
219
10
224
01
229
00
234
07
239
22
244
47
249
80
255
23
260
76
266
38
27 2
09
27 7
90
283
80
289
. 78
29 5
. B5
301
. 98
308
. 18
67327
.93
673284.66
673294.40
673304. 13
673313.86
673323.59
673333.33
673343.05
673352.79
673362. 52
673372.26
67 3381.98
673391.72
673401.45
67 HI 1 . 19
673420.91
673430,65
314. 44
320.73
327. 06
333. 40
339. 75
346. 09
352.41
673440
38
88
68
673450
12
93
bl
673459
84
98
39
673469
58
10 3
15
356. 69
4. 92
1 I. 10
17. 21
23. 25
29. 20
35. 07
40. 85
46. 5}
52. 12
57.61
63. 01
68. 32
7S. 54
78. 67
83. 71
Sec. 1, Appendix A, page
R, Norton, JPL
April 1, 1967
JPL 606-1
Orbital and Physical Data
Tabic A-1. Earth-Mars Calendar (continued)
^■;m-.'
1 -,
M.I ■ .: ■'.
) '>
A|l!- .i
Z-4-1 Ji. 1
1 . ■>
A,,t- 1^
lh\ i<.J-
l-hiU.:.
( n
M.iv d
^■i'Hi.-i
J T
jM.i. I^
^-1 i %'.=
1 -1
.VI , . .'..:
Z-t-Ht,'
J ^
.l.:n 1
i44 U,7
> S
Jil.. 1!
sr
N
MKH
^44 57 i;
i4-n74:
^ 4 ■! 5 7 ^
Aiiu I
Ai;u d
Z 4-1 '.?■,'
T
-S.-', 't
^■14 i^ 1 ■
^J^ ■ 2'
^ 4 1 5 M ^ .
:...v .s
J44 '.H i'
Nov 1^
^44iri4i
N, ■. 2ri
^4-1 i ■",'.!
[}■ L H
^■t-l IKf.r
'
WI
■: Ir.H
^44 if^Ti
iJ.'i 2.'
"^44 i:^.^:
I..-I 7
^44 ih'ii
.Inn 1^
i44 i'>i.i
,1,1(1 2"
d-\'-i VH(
s
{■■■■I, h
i44iMj'
P-4. !»■
>i 4 4 1 ■ 1 i (/
1
t-.4, dl
dAM'/ii.
s
M,. r «
Z44 3'JSf,
, T
M ,-i 1- ] H
.SH
1 1 \M ;
^44 5'<M
■s
M.,r 2m
244 S^'7iJ
s
Apr 7
^-^■ii'>H<
^
Apr 17
i-i'ii'>"i;
!i
Api' 27
^444'!(:il
M;. > ,■
.'.44 ill ID
M..> 17
^444:1 ''1
J ■in f;
Z44411411
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SIJMMKH
d-t-\-i<i-,>.:
s
J.n, 26
-M 14'Jt,[,i
■>
.till ..
<;444ii7:j
Jul !<■
^■h^■\(i^^i)
s
.lul 2«.
Z444!)''(>
'^
Ai:|; ^
/444 lUO
=,
A. = K IS
^444 1 Ji>
^
A:l^ 2S
Z444 lan
s
S.'p 4
i444 1 SU
s
S<.n 14
h'ALI,
2444 141.1
.S,-p 2 4
2444 It' :■
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2444 !(.■■!
T
('(1 14
2 4 4 1 ] 7 ' .■
s
Oi ■ 2 [
2 4 44 l^-.i.
77 , ■. -1
2444 J'"i
r.'-.v 1 i
244421. 1;
T
r.'-j',- 2 1
^■i■i■Ui<>
T
[)-■. ;
244422!)
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r, 1
Mr.
r, 7 3^7,4
M
''2'>
.Ml
rnn^'T
^0
(1 "ii
. -) 4 4
(.7 VM17
Si
fi 4 -
.27r,
(,7 Vf 17
27
.on
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(,7i<,)27
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!,(.4
c.7393t.
7 3
T
4 74
67 394 6
4 7
vv]>: rf-,1
1 '1
2 ( 1 r ,
^739^^
20
^4
)4
i> -; 1
67 3'aS
't4
44
4(M
tViVKS
40
=i4
Lib
6 7 39'*^=
13
», 1
Ht-h
67400.1
n-
73
60 1
6741114
59
H3
3 i3
674024
33
■' i
(^■>.,
67 4 34
ijt>
1 i;2
7'*H
67 4 04!
HO
i 12
S)U
674053
S2
! 22
2f,3
674063
26
1 51
9'-^S
674072
99
14 1
T2K
6740H2
73
ISI
4(,n
674092
4S
16 1
1'<J
674102
19
■SPRINi,
17 It
'.12 s
6741 n
'U
IHI)
t..S>j,
674121
66
I'yil
300
6741 il
38
200
122 1
6 7 4 i 4 1
12
2 0'»
hSS
f>74 1 SO
8S
2 17
SH7
674 160
S9
1 ".■'.
4 4
14 -i.
H •■
I4h
2)
IS^.
'ih
1 S6.
■'S
sd
!6S.
IVii.
(/>
IS. 0'^
IS. In
42. ill
6-1, t7
80. ,;.i
I'l^
2iii
21 S
,' 1
7 •,
22 T
-2
•-*■ '
T ^.
2411
'.■^
2'?l
44
4S
2'>H
Id
H4
(,H
2HS
2'M
.■'»7
Un
■1 n.
i2'^
J SS
66
00
IS
Ml.
0^
3S4.
M
•f^OSI ] ION
r!li-M.,r.. .n
r, rn; ("I?, H nulho,, 1.,:,) ,
I. mil n. XI .jppn,il,<.n- 7(
April 1, 1967
R. Norton, JPL Sec. 1, Appendix A, page 9
Orbital and Physical Data
JPL 606-1
Table A-1. Earth-Mars Calendar (continued)
RAR 111
1 -
- ! Northern T
J D
henn apher e
Year
V
Julian day
season
anH date
WINTER
1980
I
2444240.5 Jan 2 ;
(leap
2444250.5
Jan 12
year}
2444260.5
Jan 22
2444270.5
Feb 1
2444280. 5
Feb 11
2444290, 5
Feb 21
2444300.5
Mar 2
2444310 5
Mar 12
SPRING
2444320. 5
Mar ii
2444330.5
Apr 1
2444340.5
Apr 11
2444350,5
Apr 21
2444360.5
May 1
2444370 5
May 1 1
2444 380,5
May 2 1 i
2444390,5 May 3 1 |
2444400.5
Jun 10
2444410. 5
Jun 20
SUMMER
2444420.5
Jun 30
2444430, 5
Jul 10
2444440 5
Jul 20
2444450 5
Jul 30
2444460, 5
Aug 9
2444470.5
Aug 19
2444480-5
Aug 29
2444490.5
Sep 8
2444500.5
Sep 18
FALL
2444510.5
Sep 28
2444520.5
Oct 8
2444530.5
Oct 18
2444540.5
Oct 28
2444550.5
Nov 7
2444560. 5
Nov 17
2444570,5
Nov 27
2444580. 5
Dec 7
2444590.5 Dec 17
WINTER
2444600.5 Dec 27
2444610.5 Jan 6
1981
2444620.5 Jan 16
\
2444630.5
Jan 26
2444640,5
Feb 5
2444650.5
Feb 15
2444660. 5
Feb 25
2444670.5
Mar 7
2444680. 5
Mar 17
SPRING
2444690 5
Mar 27
2444700 5
Apr b
2444710. 5
Apr 16
2444720,5
Apr 26
2444730.5
.May 6
2444740.5
May 16
2444750.5
May 26
2444760 5
Jun 5
2444770 5
Jun 15
SUMMER
2444780.5
Jun 25
2444790.5
Jul 5
2444800 5
Jul 15
2444810,5
Jul 25
2444820 5
Au^ 4
2444H30 5
Aus 14
2444840.5
Aug 24
2444850.5
Sep 3
2444860,5
Sep 1 3
FALL
2444B70, 5
Sep 23
2444BK0 5
Oct 3
2444890 5
O^t 1 i
2444900.5
Gel 2 3
2444910. 5
Nov 2
2444920, 5
Nov 12
2444930. 5
Nov li
2444940, 5
Dec 2
2444950. 5
Dec 12
WIN rE;R
2444960 5
Dec 22
.1
FEB 25 APHELIC OPPOSirlON
Earth-Mars distance at closest appr
63.2 ]nillion mi (10 17 n-.illion km).
Sec. 1, Appendix A, page
10
R. Norton, JPL
April 1, 1967
JPL 606-1
Orbital and Physical Data
Table A-1. Earth-Mars Calendar (continued)
TAR IH
MARS
Northern
Northern
Heliocentric
J.D.
hcminpluTi'
Year
Year
Hennisphere
Consecutive
ecliptic
Julian day
season
and date
season
and year day
Mars day
loniiitiide,
dej.
WINTER
198 2
1009
SPRING
2444970 S
Jan 1
(leap
300 985
674909,98
15n. 92
2444980,5
Jan 11
year)
310. 71B
674919, 71
155. 29
2444990. 5
Jan 21
320.450
674929,45
159. 66
2445000, =>
Jan 31
330, 183
674939, 18
164.03
2445010.5
Feb 10
339.915
674948.91
168. 41
2445020.5
Feb 20
349. 648
674958 64
172. 80
2445030.5
Mar 2
359.380
SUMMER
67496B, 38
177. 22
2445040. 5
Mar 12
SPRING
369. 112
674978, 11
181.66
2445050.5
Mar 22
378,845
674987,84
186. 13
2445060.5
Apr 1
388,577
674997,57
190. 64
2445070, 5
Apr 11
398,310
675007, 30
195. 19
2445080,5
Apr 21
408.042
675017,04
199. 78
2445090,5
May 1
417,775
675026,77
204. 43
2445100.5
May 11
427,507
675036, 50
209. 13
2445110,5
May 21
437,240
675046.23
213. 89
2445120.5
May 31
446,97 2
675055.97
218. 73
2445 130,5
Jun 10
456,704
675065.70
223. 63
2445 140.5
Jun 20
SUMMER
466.437
675075.43
228.61
2445150.5
Jun 30
476. 169
675085. 16
233.67
2445 160.5
Jul 10
485.902
675094,90
238. 81
2445170.5
Jul 20
495.634
675104,63
244.05
2445180.5
Jul 30
505.367
675114.36
249. 37
2445190.5
Aug 9
515.099
675124.09
254. 79
2445200.5
Aug 19
524.832
675133.83
260. 30
2445210.5
Aug 29
534.564
FALL
675143.56
265.91
2445220.5
Sep 8
544.296
675153.29
271. 62
2445230.5
Sep 18
FALL
554.029
675163.02
277. 41
2445240.5
Sep 28
563,761
675172.76
283. 30
2445250.5
Oce 8
573,494
675182.49
289. 28
2445260.5
Oct 18
583,226
675192.22
295. 33
2445270.5
Oct 28
592,959
675201.95
301. 46
2445280.5
Nov 7
602,691
675211.69
307. 66
2445290.5
Nov 17
612,423
675221,42
313.90
2445300.5
Nov 27
622, 156
675231, 15
320. 20
2445310.5
Dec 7
631,888
675240,88
326. 52
2445320.5
Dec 17
WINTER
641,621
675250,62
332.86
2445330. 5
Dec 27
651. 353
661.086
675260.35
675270.08
339. 21
345. 55
2445340. 5
2445350. 5
Jan 6
Jan 16
1983
loio
■!.8(S
^75279.91
351. 87
(leap
WINTER
2445360. 5
Jan 26
year)
11.551
675289.55
358. 16
2445370. 5
Feb 5
21. 283
675299.28
4.40
2445380. 5
Feb 15
31.014
675309.01
10. 59
2445390. 5
Feb 25
40. 747
675318. 74
16.71
2445400. 5
Mar 7
50.479
675328.48
22.75
2445410. 5
Mar 17
SPRING
60. 212
675338. 21
28.71
2445420. 5
Mar 27
69.944
675347.94
34. 59
2445430. 5
Apr 6
79.677
675357.67
40. 38
2445440. 5
Apr 16
89.409
675367.41
46.07
2445450.6
Apr 26
99. 142
675377. 14
51.67
2445460. 5
May 6
108. 874
675386, 87
57. 18
2445470. 5
May 16
118.606
675396.60
62. 59
2445480. 5
May 26
128. 339
675406. 34
67.90
2445490. 5
Jun 5
138.071
675416.07
73. n
2445500. 5
Jun 15
SUMMER
147.804
675425. 80
78. 27
2445510. 5
Jun 25
157.536
SPRING
675435. 53
83. 33
2445520. 5
Jul 5
167.269
675445. 27
88. 30
2445530, 5
Jul 15
177. 001
675455.00
93. 20
2445540. 5
Jul 25
186.734
675464. 73
98. 03
2445550. 5
Aug 4
196. 466
675474.46
102. 80
2445560. 5
Aug 14
206. 198
675484. 20
107. 50
2445570. 5
Aug 24
215.931
675493.93
112. 14
2445580. 5
Sep 3
225.663
675503.66
116.74
2445590. 5
Sep 13
FALL
235. 396
675513.39
121. 28
2445600. 5
Sep 23
245. 128
675523. 13
125. 79
24-45610. 5
Oct 3
254.861
675532. 86
130. 26
2445620. 5
Oct 13
264. 595
675542.59
134. 70
2445630. 5
Oct 23
274. 326
675552. 32
139. 12
2445640. 5
Nov 2
284. 1)58
675562. 05
141. 52
2445650. 5
Nov 12
293. 790
675571. 79
147.90
2445660. 5
Nov 22
ilH. 521
675581. 52
152. 27
2445670. 5
Dec 2
315.255
6755'll. J5
156.64
24456B0. 5
Dec 12
WINTER
322.988
67561111. "K
161. iin
2445690. 5
Dec 22
332. 720
67 56 HI. 7 2
165. IH
April 1, 1967
R, Norton, JPL Sec. 1, Appendix A, page 11
JPL 606-1
Orbital and Physical Data
APPENDIX B
GLOSSARY
Apparent Position
Apsides
Areocentric
Areocentric Declination
of the Earth (D„)
ill
Areocentric Latitude
Areocentric Longitude
of the Sun (L„)
Areographic
= Areodetic [ Latitude
Areographic
= Areodetic j Longitude
Areocentric Right
Ascension of the Earth
In astromony, the position of a celestial body as
seen from the center of the Earth or a planet.
The end points of the major axis of an elliptical
orbit.
Mars-centered. (Derived from 'Ares, " Greek
word for Mars, the god of war. )
On an areocentric sphere, the angular distance of
the Earth from the Martian equator, measured
from to 90°, + to the north and - to the south of
the equator.
The angle at the center of Mars between the
equatorial plane and the straight line drawn to a
given point.
Counted from to 90°, + or - according to north-
ward or southward orientation of the line.
The longitude of the Sun on an areocentric sphere
measured in the Mars orbital plane from its vernal
equinox westward from to 360°.
It is equal to the heliocentric orbital longitude of
Mars measured from its autumnal equinox.
The angle between the equatorial plane of Mars
and the normal to the reference ellipsoid (or
spheroid) at a given point on Mars' surface.
Counted from to 90°, + or - according to north-
ward or southward orientation of the normal.
The angle between the reference meridian (plane)
or prime meridian, and the meridian (plane) at a
given point on Mars' surface.
Counted from to 360° westward; i.e. , clockwise
as viewed from the Mars' north pole, or opposite
to Mars' rotation sense.
On an areocentric sphere, the angular distance in
the Martian equatorial plane, measured from to
360° eastward from the Martian vernal equinox to
the great circle through the Earth and the Mars'
celestial poles.
October 15, 1971
C. Michaux, JPL
Sec. 1, Appendix B, page 1
Orbital and Physical Data
JPL 606-1
Aphelion
Argument
Ascending Node
Astronomical Unit
Barycenter
Celestial Equator
Celestial Latitude
Celestial Longitude
The point on a heliocentric orbit that is furthest
from the Sun.
Arc, angle, angular distance, longitude.
The intersection point of the orbit and reference
great circles on the celestial sphere where the
body (planet or satellite) moves from the southern
into the northern hemisphere.
The fundamental unit of distance used in astromony.
It is defined as "the unit of distance in terms of
which, in Kepler's Third Law n^a^ = k.2 (1+m),
the semimajor axis a of an elliptical orbit must be
expressed in order that the Gaussian constant k
may be exactly 0.01720209895, when the unit of
time is the ephemeris day" (Explan. Suppl. , 1961).
Here n is the angular mean motion of any planet in
radians per day, and m the ratio of the mass of the
planet to the mass of the Sun.
The AU is close to but not identical with the mean
distance Earth-Sun. (For Earth, a - 1.00000003
AU. )
For calculations, a rounded value of 1 AU =
1.496 X 10^ km (based on radar measurements)
has been adopted by the lAU Assembly in Hamburg,
1964, and introduced in all national ephemerides
as of 1968.
Center of mass of a group of bodies.
The great circle of the celestial sphere marking
its intersection with the extended equatorial plane
of the Earth.
The arc along a great circle perpendicular to the
ecliptic between the ecliptic and the body, mea-
sured from to 90° positively to the north and
negatively to the south of the ecliptic.
Or (more simply): the angular distance north or
south of the ecliptic.
The angle at the ecliptic pole, or the arc of the
ecliptic, between the circle of celestial latitude of
the vernal equinox and the circle of celestial lati-
tude of the body, measured from to 360° east-
ward from equinox.
Or (more simply): the angular distance east of the
vernal equinox along the ecliptic.
Note: Celestial latitude and longitude of a body are
its celestial coordinates in the 'ecl.iptic' system.
Sec. 1, Appendix B, page 2 C. Michaux, JPL
October 15, 1971
JPL 606-1
Orbital and Physical Data
Celestial Meridian
The hour circle containing the zenith. Also called
the meridian.
Celestial Poles
Celestial Sphere
Central Meridian
Circle of Celestial
Latitude
Declination
Descending Node
Eccentric Anomaly
Eccentricity (of ellipse)
Ecliptic
Ecliptic Poles
The points of the celestial sphere which are the
intersections with the extended rotation axis of the
Earth (or planet).
The imaginary sphere of infinite radius sur-
rounding the observer and upon which all celestial
bodies (except that of the observer) may be
projected.
The fictitious meridian bisecting a planetary disk
as seen externally.
A great circle of the celestial sphere through the
ecliptic poles for the Earth.
The arc along an hour circle from the celestial
equator to the body in question measured fromi
to 90° positively to the north and negatively to the
south of the equator.
Or: angular distance north or south of the celes-
tial equator.
Same as for ascending node except: from northern
into southern.
The angle described by the radius from center of
the orbital ellipse to the projection of the planet
(or satellite) onto the auxiliary circle of the ellipse
and counted from perihelion (or perifocus) position
in the direction of motion.
The ratio of the distance between the center and
the focus of an ellipse to its semimajor axis.
It is expressed by e = \/l - b^/a^ if a is the semi-
major axis and b the semiminor axis.
The great circle marking the intersection of the
celestial sphere with the (mean) orbital plane of
the Earth's center.
Or; approximately: the annual apparent path of
the Sun's center on the celestial sphere, as seen
from the Earth.
The points of the celestial sphere which are the
intersections with the line perpendicular to the
ecliptic.
October 15, 1971
C. Michaux, JPL
Sec. 1, Appendix B, page 3
Orbital and Physical Data
JPL 606-1
Elongation
Ephemeris
Ephemeris Second
Ephemeris Time
Epoch
Equinoxes or Equinoctial
Points
First Point of Aries T
The difference in geocentric (or planetocentric)
longitudes between the Sun and the planet
observed.
A planet is in "eastern elongation" when it follows
the Sun in its apparent daily motion; it is in
"western elongation" when it precedes it.
A table predicting the positions of celestial bodies
(Sun, Moon, planets) at regular intervals of time.
The fraction 1/31556925.9747 of the tropical year
for 1900 January 0^ 12^ ET (i. e. , at the
fundamental epoch).
This fundamental invariable unit of time was
formally adopted by the Comite International Des
Poids et Mesures in 1957.
Note: The fraction above was determined by the
coefficient of the time variable T in Newcomb's
expression for the right ascension of the
fictitious mean Sun. )
A uniform measure of time based on the invariable
unit known as the ephemeris second.
Ephemeris Time was established not for practical
uses (it is not readily accessible like Universal
Time or mean solar time) but for theoretical uses
as an accurate standard in studying the laws of
motion of celestial bodies.
The initial time adopted as the instant of
reference.
The (2) intersection points of the celestial equator
and the ecliptic. (Definition applying to Earth
case. )
One point is the 'vernal equinox' (or point at which
the Sun appears to cross the equator from south to
north). The other point, opposite to it (or 180°
from it) is the 'autumnal equinox. '
These designations were chosen to indicate the
beginnings of (astronomical) spring and autumn,
respectively, in the northern hemisphere of
Earth. (Since seasons are reversed in the
southern hemisphere, it may be preferable to use
the terms March equinox and September equinox.)
The same as vernal equinox (for Earth).
Sec. 1, Appendix B, page 4 C. Michaux, JPL
October 15, 1971
JPL 606-1
Orbital and Physical Data
Fundamental Epoch
Gaussian Gravitational
Constant
(From which Ephemeris Time is reckoned): the
epoch that Newcomb designated as 1900 January 0,
Greenwich Mean Noon, but which is now properly
designated as 1900 January 0, 12^ ET (Explan.
Suppl., 1961).
The instant to which this designation is assigned
is the instant near the beginning of the calendar
year AD 1900 when the geometric mean longitude
of the Sun, referred to the mean equinox of date,
was 279° 41' 48.04". (Explan. Suppl., 1961).
A fixed and exact constant k = 0.01720209895
serving to define the astronomical unit of distance
through Kepler's Third Law.
Originally used as k^ = G, the Newtonian gravita-
tional constant.
General Perturbations
In celestial mechanics, the analytical expressions
(in the form of infinite series) necessary to
calculate the changes or perturbations in the
orbital elements as a function of time.
Greenwich Mean Time
Greenwich Mean Noon
(or GMN)
Heliocentric
Hour Angle (local and
Greenwich)
Hour Circle
Practically identical with Universal Time.
The same as 12^ UT.
Note: 1900 January O.GMN = 1900 January 0^12^
ET Ne"wcomb's fundamental epoch was kept when
revising the UT system and establishing the ET
system.
Sun-centered (derived from 'helios, ' the Greek
word for Sun).
The angle or arc measured westward along the
celestial equator from the local (or Greenwich)
meridian to the hour circle of the body in question.
A great circle of the celestial sphere through the
celestial poles. Also called circle of right
ascension.
Julian Calendar
Julian Century
Julian Date (of an event
or instant)
A simple continuous calendar reckoning the
number of astronomical days elapsed since a
chosen ancient epoch (January 1, GMN: 4713 BC).
Unit of time equal to exactly 36525 days (mean
solar).
The Julian Day Number followed by the decimal
fraction of the day elapsed since preceding noon.
It may be measured in days UT (Julian Date J. D. )
or in days ephemeris (Julian Date J. E. D. ).
October 15, 1971
C. Michaux, JPL
Sec. 1, Appendix B, page 5
Orbital and Physical Data
JPL 606-1
Julian Day Number
Julian Year
Kepler's Laws (of
planetary motion)
Laplacian Plane
(of a satellite)
Line of Apsides
Line of Nodes
Longitude (of a point)
Martian Vernal Equinox
Mean Anomaly
Mean Distance
Mean Equator
Mean Equinox of Date
The integral number of mean solar days elapsed
since the chosen Julian epoch.
Unit of time equal to exactly 365.25 days (mean
solar).
Three laws applying to the motion of one planet
around the Sun, assuming that it is not disturbed
gravitationally by any other planet.
An invariable plane relative to the planet's
equator, and upon which the precessing plane
of a satellite's orbit maintains a (nearly)
constant inclination.
The line connecting the two apsides; i.e., the
major axis of an elliptical orbit
The intersection line of the orbit plane and a
reference plane or circle (usually either the
ecliptic or an equator) on the celestial sphere.
Arc or angular distance along a great circle in a
reference plane, counted from an adopted
reference point.
On an areocentric sphere, the ascending node of
the orbit of Mars on its equator.
The angle described by the radius vector of the
orbiting body in the interval of time (t-tg) since
perihelion passage, assuming constant mean
angular motion n = Ztt/P, where P is the sidereal
period of revolution.
It is the same as true anomaly but with a
fictitious orbiting body revolving with uniform
angular velocity around the Sun (or primary).
Average distance of a body from its primary.
Equal to the semimajor axis.
The mean position of the plane of the equator as
resulting from the effect of precession, but not
of nutation.
The ascending node of the ecliptic on the mean
equator at a particular date.
It is a fictitious equinox with the effect of
nutation removed.
Sec. 1, Appendix B, page 6 C. Michaux, JPL
October 15, 1971
JPL 606-1
Orbital and Physical Data
Mean Motion
The mean angular velocity of the orbiting body.
It is expressed by n = 2tt/P, if P is the period of
revolution.
Mean Solar Day
Mean Solar Second
Mean Solar Time
(for Earth)
Mean Sun (for Earth)
Meridian Transit
The basic interval of time in the mean solar time
system; i. e. , the interval between two consecutive
transits of the fictitious mean sunover a meridian
(corrected for the motion of the pole).
The fraction 1/86400 of the mean solar day. (The
mean solar second was the fundamental unit of
time before the adoption of the ephemeris second
in 1957.)
(At any place) the local hour angle of the fictitious
mean sun plus 12 hours (to start from midnight).
A fictitious sun moving eastward along the
celestial equator at a uniform rate such that it
completes a revolution in the same time as that
of the actual Sun along the ecliptic.
The passage of a celestial body across a
celestial meridian.
Nadir
Nutation
Obliquity of the Ecliptic
Opposition
Orbit
Pericenter
The point of the celestial sphere exactly opposite
(180°) from the zenith.
The somewhat irregular circular motion of the
true pole of a planet's equator about the mean
pole. For Earth, the nutation period is of
18.6 years, while its angular amplitude (constant
of nutation) is of 9-210 arc sec. For Mars, it is
still not known but appears to be even smaller.
Nutation is essentially the short-period periodic
portion of the precessional motion of the pole and
depends principally (in the case of the Earth) on
the periodic motion of the Moon with small
contributions from the Sun and other planets.
The angle between the planes of the celestial
equator and of the ecliptic.
Time when the apparent geocentric (planeto-
centric) longitudes of another planet and the Sun
differ by exactly 180°.
The path of a celestial body under the gravita-
tional attraction of another body or bodies.
The point on an orbit that is nearest to the center
of attraction.
October 15, 1971
C. Michaux, JPL
Sec. 1, Appendix B, page 7
Orbital and Physical Data
JPL 606-1
Perifocus
Perihelion
Periodic Perturbations
Perturbations
(of an orbit)
Precession (of the
equinoxes of a planet)
Quadrature
Retrograde Sense
Revolution
Right Ascension
Rotation
Secular Perturbations
Semimajor Axis (of
ellipse)
Same as pericenter.
The point on a heliocentric orbit that is closest
to the Sun.
Perturbations which are of short period compared
to secular perturbations.
Deviations from Keplerian or true elliptical
orbital motion around the Sun or primary.
The very slow conical motion of the rotational
axis of a planet about the normal to its orbital
plane, as caused by the attraction of the Sun, the
other planets, and especially any large satellite
upon the equatorial bulge of the planet. (Pre-
cession is commonly exhibited by a spinning top
under the torque produced by gravity. )
The position of a planet (or Moon) at 90° to the
Sun as measured from the center of the Earth.
The opposite of direct sense of rotation or
revolution; that is clockwise as seen from the
north pole.
The orbital motion of a celestial body (planet or
satellite) around another more massive body (Sun
or primary).
The arc along the celestial equator measured
eastward from the vernal equinox to the hour
circle of the body, from to 24 hours or
0° to 360°.
Note: Right ascension and declination of a body
are its celestial coordinates in the 'equatorial'
system.
The spinning of a celestial body (Sun, planet or
satellite) about an axis passing through it.
Perturbations which are very slowly changing the
orbit of a planet (or satellite).
It was shown by Laplace and Lagrange that such
perturbations actually are periodic with extremely
long periods.
One -half the longest diameter of an ellipse.
Also called "mean distance. "
Sec. 1, Appendix B, page 8
October 15, 1971
JPL 606-1
Orbital and Physical Data
Semiminor Axis
(of ellipse)
Sidereal
Solstices or Solsticial
Points
Special Perturbations
Superior Conjunction
Synodic Period
(of revolution of two
planets or satellites)
Tropical Year
True Anomaly
True Equinox of Date
Universal Time
Vernal Equinox (T)
Zenith (astronomical)
One -half the shortest diameter of an ellipse.
With reference to the stars.
The (2) points on the ecliptic (Earth case) which
are 90° from the equinoxes. One is the summer
solstice, the other the winter solstice.
In celestial mechanics, the perturbations in the
orbital elements at successive time intervals as
calculated through a stepwise numerical integra-
tion of the equations of motion.
Time when the apparent geocentric (planetocentric)
longitudes of another planet and the Sun are
exactly the same, with the Sun between the two
planets (except for latitude differences).
The time interval between two successive
heliocentric or planetocentric conjunctions in
longitude.
The interval between two successive returns of
the (fictitious mean) Sun to the (mean) vernal
equinox.
Or: the interval during which the Sun's mean
longitude (or right ascension), referred to the
mean vernal equinox, increases by 360°
(Explan. Suppl., 1961).
The angle described by the radius vector from
the Sun (or primary) to the planet (or satellite)
counted from perihelion (or perifocus) position
in the direction of motion.
The intersection of the true equator (as affected
by both precession and nutation) with the true
ecliptic.
The miean solar time referred to the Greenwich
meridian.
The ascending node of the ecliptic on the equator.
The point of the celestial sphere directly
(vertically) above the observer.
October 15, 1971
C. Michaux, JPL
Sec. 1, Appendix B, page 9
J^^L 606-1 I^.^ri^^
SECTION 2 CONTENTS
2. INTERIOR
Data SiiiriiTiary
Discussion
2. 1 The Theory of a Rotating Planet '. . .
FlattenintT
Geometric Relationships
Dynamical Flattening .* ' ' * 2
Optical Flattening ' * o
External Gravitational Potential . . . 4
Potential of Gravity ' r
Determination of Coefficient J2 From Satellite
Orljital Precession /
Determination of the (Polar) Moment of Inertia . . . 7
Gravity Relation o
Hydrostatic Flattening '....'.' 9
2. 2 Density Models n
IntrodTiction p
Historical Density Models '..'.'' 10
Homogeneous Models * ' ' 11
Chemically Differentiated Models 11
Hybrid Models , ,
Recent Density Models I 1
Binder's (1969) Models 14
Anderson's (1972) Models 15
Implications of Anderson' s (1972) Chondritic Models' ." . . . . '. 18
Planet formation 10
Magnetic field * ' " jg
2. 3 Thermal History Models * ' ] ,q
Introduction ,„
Some Planetary History 2
Recent Thermal Models. ] ' 2]
Anderson and Phinney (1967) . . . . 21
Hanks and Anderson (1969) . . . ?!
Bibliograpliy
12
Figures
1(a). Density models of Mars' interior
1(b), Theoretical interior density profiles for Earth and Mar's* . . . . 12
2(a). Density, gravity, and pressure profiles of a model
representative of the models that satisfy the moment
of inertia and temperature limits and that have mantle
composition similar to the composition of the mantle
of Earth , .
lb
March 1, 1972 o^^ 7 r^ 4. ^
bee. 2, Contents, page 1
Interior
JPL 606-1
(cont'd)
2(b). Density, gravity, and pressure profiles of a model
representative of the models that satisfy the moment
of inertia and temperature limits and that have a mantle
composition intermediate between the composition of the
terrestrial mantle and the composition of olivine in
meteorites
3 Schematic binary phase diagram for the system Fe-FeS 1 .^
4 Possible density models of Mars satisfying proper mean density
and moment of inertia of the planet, calculated by Anderson. . . 19
5 Temperature profiles for Mars ID and Mars IID 23
6 Thermal history profiles for Mars model (Mars II'D) which
assumed radioactive abundances 0.8 of the "terrestrial"
and total accretion time 300,000 years 23
Sec. 2, Contents, page ii March 1, 1972
JPL 606-1
Interior
2. INTERIOR
DATA SUMMARY
Gravitational Coefficient J^
Flattening:
Dynamical (preferred
value) f ,
Optical f^
Polar moment of inertia
factor (hydrostatic) C/MR^
eq
0.00195
0.00525
0.012
0.377
(Wilkins, 1967; Cain, 196 7)
(Cain, 196 7)
(de Vaucouleurs, 1964)
Interior models
See Figs. 1 and 2, etc.
DISCUSSION
2. 1 The Theory of a Rotating Planet
Flattening
The flattening of a planet is the difference between the equatorial and
polar radii divided by the equatorial radius. (For a more detailed description,
see MacDonald, 1962.)
Geometric Relationships
Any rotating, self -gravitating, fluid body will assume a shape such that
its surface is everywhere normal to the resultant of gravity and centrifugal
force (Sterne, I960), This equipotential surface will approximate a spheroid
for slow rotation (the case for planets) (Jardetzky, 1958). Real bodies may
have sufficient strength in their mantles to allow their surfaces to depart some-
what from an equipotential surface or to support internal inhomogeneities.
The simple geometric relationships for a spheroid of revolution about an
axis are given below, where
f
e
R
R
eq
R
P
R
m
flattening
eccentricity
radius at geocentric latitude (j>
equatorial radius
polar radius
mean radius
November 15, 1971
C. Michaux, R. Newburn, JPL
Sec. 2, page 1
Interior JPL 606-1
By definition
R - R R
f - eg P 3 1 P
R R
eq eq
and
\ eg P/
1/2 . . .1/2
r"^
Thus
and
e = ^ ^^ n ^^ = (1 -^
%q \ R^
®g,
e^ = f(2 - f)
i = I - y/l
or, ignoring 0(f2),
>2
'=T
At any planetocentric latitude cf)
/ 2 \'/'
° \ 1 - e2 cos ^ /
or, ignoring 0(f2),
R := R (1 - f sin2<j>)
eq
and mean radius
_('4)
Dynamical Flattening
R = R
m eq
Dynamical flattening is a flattening given by the gravity potential of a
primary body; the equipotential surfaces described by this potential may or may
not refer to the material surface of that body. The dynamical flattening of Mars
Sec. 2, page 2 R. Newburn, JPL November 1, 1967
JPL 606-1 Interior
is derived from consideration of the orbital perturbations of Phobos and
Deimos. The theory of this motion, vvith improved data from Mariner 4, yields
a value of fj = 0,005 25 ±0.00001 (Cain, 1967). A quantitative discussion of
these ideas is given in paragraphs beginning on page 4.
The dynamical flattening of Mars is well determined and corresponds to a
value theoretically reasonable for a planet with a very small amount of com-
pr(u-3sion toward its center. Spectroscopic measurements agree with this value.
Spectroscopic determinations also show no more CO2 at the polar caps than at
the equiitor. If the true surface were that implied by the optical flattening
value, CO2 pressure at the caps should be several times that at the equator
(Hanselman, 1965). The figure as determined by Mariners 4, 6, and 7 agrees
nicely with the dynamical value (see Sections 1 and 5. 2). For these reasons
the mean surface of Mars is suggested to be nearly that indicated by the
dynamical flattening.
Optical Flattening
Optical flattening of the apparent planetary surface is measured either
directly on the optical image with a micrometer or heliometer, or indirectly
using photographic images or photoelectric scans. Most results are between
0.010 and 0.015, with recent values near 0.012, apparently corresponding
to a surface far from an equipotential. However, optical flattening determina-
tions are far less accurate than those obtained for dynamical flattening. Optical
measurements are difficult to make and are subject to many sources of
systematic error and personal evaluation, for the following reasons:
1) Obscuration, by the Martian atmosphere, of the planetary surface
(edge of the disk).
2) Turbulence in the Earth's atmosphere.
3) Preferential placement of crosswire relative to a bright or dark
area.
4) The gibbous nature of the Martian disk; noncircularity of
the disk due to phase angle and oblateness of the planet.
5) Possible exaggerated equatorial diameter measurements caused by
dust particles in the atmosphere during the perihelic oppositions
when such measurements are characteristically made.
6) Mars seldom exhibits a true polar diameter because of axial tilt.
7) Photographic images also lack sharpness at image edge,
8) Photoelectric scans suffer from finite slit width, instrumental
scattering, and orientation problems.
November 15, 1971 R. Newburn, JPL Sec. 2, page 3
Interior JPL 606-1
External Gravitational Potential (of a spheroid)
It is well known that the external gravitational potential V(r,6A). of a
planet is given by an expansion in terms of "spherical harmonics, " which are
functions of latitude ^ and longitude X; (see Jeffreys, 1970). If the planet is
axially (i. e. , rotationally) symmetric (a spheroid), the spherical harmonics
reduce to "zonal" harmonics, which are functions of latitude <^ only (through
Legendre polynomials P^(sin 6). In its simplest form, the potential V at a
point (r, (b) external to such a planetary body is given (keeping only the two first
terms) by
,, GM ^ GM „ 2 J ,^ . 2, ,.
V = + R . Jt (3 sm 9 - 1)
r T 3 eq 2 ^
2r-^ ^
where (r,4>) are the planetocentric coordinates of distance and latitude of the
point:
G = the absolute gravitational constant
M = the mass of the planet
R = its equatorial radius
eq ^
J^ = a dimensionless coefficient (to be determined)
This coefficient Jt of the second harmonic is the most important in
defining the shape or figure of the planetary body (or rather its approximating
equipotential spheroid). If higher harmonics were considered, their coefficients
(J4, . . . ), which are much smaller, would describe the deviations of the actual
equipotential surface from the spheroid.
It can be shown theoretically that coefficient Jt is related to the difference
between the principal moments of inertia C (about the polar axis) and A (about
an equatorial axis) of the body by an important relation independent of internal
structure:
J
2 ~ 2
MR^q
These coefficients J^, J41 etc. may be measured observationally with
great precision from study of the orbital motions (perturbations) of close
satellites and orbiters (such as planned for Mariner 9). From Earth, observa-
tions of the Martian satellites Phobos and Deimos are accurate enough to yield
a reliable value of Jt, but not J^. It is also possible to obtain J2 from per-
turbations of a flyby trajectory.
Sec. 2, page 4 C. Michaux, JPL November 15, 1971
JPL 606-1 Interior
Potential of Gravity (at surface)
The potential of gravity at any point on the surface of the approximating
equipotential of the body is the sum of gravitational potential and rotational
potential (due to centripetal acceleration):
U = V :r- r cos 4>
where cu is the angular rate of rotation of planet, and (r, cj)) are now the co-
ordinates of the surface point.
Note: In the case of the Earth, the surface of constant (geo) potential U, or
equipotential Uq, coinciding with the mean sea level, is by definition the so-
called "geoid. " This irregular geoid is approximated in (physical) geodesy by
the more regular "normal spheroid" of revolution, or even by the "reference
ellipsoid" of revolution (used in geometrical geodesy), which are mathemati-
cally tractable surfaces, -'■
From the expression for the potential of gravity U, the flattening of the
approximating spheroid or ellipsoid can be derived to first order terms by
R - R - , w^R^
-^ E = f, =lj, +1. 23.
often written
2
R 2 2 GM
eq '
f 3 ^ ^ 1
^d == I J2 + Y "^
where
that is
R and R = the equatorial and polar radii of the spheroid or
^ ^ ellipsoid
m = the ratio of centrifugal to gravitational acceleration
at the equator
u)^R
m = 2^
GM/R^
eq
='=See, for example, Handbook of Geophysics and Space Environments (section
Geodesy), edited by S. L. Valley, at Mc Graw-Hill, N.Y. (1966).
November 15, 1971 C. Michaux, JPL Sec. 2, page 5
Interior JPL 606-1
Using the mean density of the planet
•■p = M/4^R^
3 tn
and its period of rotation
P =-'-^
this ratio is sometimes taken to be equal to
3rT
m =
GpP^
which is an approximation valid when flattening is small (i. e. , R^^ == '^eq^*
The flattening f , is usually referred to as the "dynamical flattening. "
Determination of Coefficien t J^ From Satellite Orbital Precession
An astronomically observed value of gravitational coefficient J2 can be
deduced for a planet from the period of precession of the orbital pole (or line of
nodes) of a close satellite's motion, when perturbed by the planet's oblateness:
J - 2 a p
where
2 " 3 r2 ^^
eq
a = semi-major axis of satellite orbit
p = period of revolution of satellite
R = equatorial radius of planet
eq ^
u = period of precession of satellite orbit
The estimate of J-, does not depend on the value of mass M nor on any
assumption on the internal mass distribution of the planet.
is:
The value of J2Req obtained from the best available satellite orbital data
J^r2 = (1.013 ±0.003) X 10"^^ AU^ (Wilkins, 1967)
2 eq
Sec. 2, page 6 C. Michaux, JPL November 15, 1''7 1
JPL 606-1 Interior
Using this value and the equatorial radius Req ~ 3393.4 km obtained from
Mariner 4 data, Cain derived the coefficient J2 = 0.00195 ±0,00002
(Cain, 1967).*
Determination of the (Polar) Moment of Inertia
The moment of inertia of a planet can only be determined from measure-
ments nnade externally to its surface (and not from gravity measurements at
its surface). There are two well-known methods for determining the principal
or polar moment C:
1) Method A utilizes the "mechanical ellipticity" of the planet, a
quantity defined by H = (C-A)/C, and the coefficient
Jz
(C-A)
mr2
eq
Indeed, their ratio
h
H MR^
eq
yields C imnnediately.
This method works well for Earth, because H is well determined
from the rate of precession of its axis of rotation, resulting from
the torques exerted upon it by the Moon and the Sun. In the case
of Mars, the mechanical ellipticity H is not well determined.
Because the tiny satellites do not exert a significant torque, and the
Sun produces very little torque, H can only be approximated from
an assumed internal density distribution, which is unsatisfactory.
(See, for example, Lowell's 1914 estimate: H = 0.005). Hopefully,
in the near future, H will be determined by spacecraft.
2) Method B combines the observed coefficient J2> or equivalently the
dynamical flattening
^d " 2 J^+T"^
*This value of J2 may be compared with that obtained by Null (1971) from the
perturbations of Mariner 4 flyby trajectory: J2 = 0.00 187 ±0.00007. Thus,
J2 = 0.0019. A more final value of J2 (at least to the third significant digit)
should be easily obtainable by the Mariner 9 orbiter. It was practically im-
possible to derive a reliable value of J2 from the Mariners 6 and 7 flyby
trajectories perturbations because of additional, nongravitational forces (N2
gas leaks from IRS experiment, etc.), during encounter.
November 15, 1971 C. Michaux, JPL Sec. 2, page 7
Interior JPL 606-1
and the Radau-Darwin relation, which assumes hydrostatic
equilibrium; that is
£ - f
^d ^h
Indeed, this Radau approximation relates
—, m and f, (~f )
MR
eq
This method has been applied to Mars to yield an approximate
value of
^ = 0,377
MR^
eq
usine the Mariner 4 data for M and R (which both enter in m also),
o eq
or
C = 2.79X10^3 g_cm2
By substituting this result for C into
(C-A) _
MR'^ ^
eq
one may also obtain an approximate value of the equatorial
moment A.
Gravity Relation (Clairaut's Formula)
Gravity or gravitational acceleration at the surface of the planet, ap-
proximated by the equipotential ellipsoid, is normal to this ellipsoid and of a
magnitude given by the gradient of the (total) potential:
g = - grad U
From the expression for U, one obtains gravity at the equator (where
r = Rgq and 4> = 0):
GM ,1 , r 3 ^
(1 + f , - — m)
'« r2 d 2
eq
Sec. 2, page 8 C. Michaux, JPL November 15, 1971
JPL 606-1
Interior
The gravity at any other latitude cj) is given in terms of g (to the accuracy
of the first power of f ,) by Clairaut's formula;
) sin ()3
1 + (tt- m - f J
(Note: Planetographic latitude c|>'may be used instead of planetocentric latitude
<^, with negligible loss in accuracy. )
Again, Clairaut's formula, based on fj or Jt, is also independent of any
assumption on the internal structure of the planet 7such as hydrostatic
equilibrium).
Hydrostatic Flattening (Radau's Approximation)
The shape or figure of equilibrium of a rotating fluid planet, where
density distributes itself in concentric layers, is provided by the solution to
Clairaut's equation (see Jeffreys, 1970). This is well expressed for our pur-
poses by the hydrostatic flattening approximation (neglecting terms of order f2);
'5 15 C
m
1
MR^
eq
This relation between hydrostatic flattening fj^, ratio m, mass M, equa-
torial radius Rgg. and principal moment of inertia C is known as the Radau-
Darwin approximation. It is often written;
MR'^
eq
Z
3
i-l,/i.--i
If an observed or true C/MRA value were available for Mars (such as
given by J2/H method), one could determine a hydrostatic flattening i-^ and
connpare it to the dynamical flattening f^.
A hydrostatic J 2 can be derived from fh by substitution in the J2 expres-
sion in fj. Comparison of this hydrostatic J2 with the observed J2 may reveal
a difference. This difference would indicate the extent of deviation of the
actual planet from the ideal (hydrostatic) planet. This would provide evidence
on the anelasticity of the planet's interior, and give an idea of the magnitude
of the stress differences supported by a portion of (the mantle) or possibly the
whole body (in the case of a small planet like Mars). (Note: Such stresses are
likely to arise during the thermal history of the planet, )
2. Z Density Models
Introduction
At the present time, and until new measurements are secured by orbiters
and/or landers, the internal structure of Mars can only be inferred from the
November 15, 1971
C. Michaux, JPL
Sec. 2, page 9
Interior JPL 606-1
very few available pUinetary constants (the so-called "observables '). These
available constants are the mass, the radii, or simply the mean radius, and
the polar (or mean) moment of inertia. - The Mariner 4 flyby provided a very
accurate value for the mass and a reasonable value for the mioment of inertia.
The three Mariner flybys established an accurate value for the equatorial
radius. However, the polar radius and geometrical flattening are still not
accurately known.
The two basic assumptions underlying all internal models of Mars have
been: (1) analogy with the Earth, either in structure or composition; and
(2) hydrostatic equilibrium. Other assumptions usually pertain to the possible
discontinuities; that is, either changes of state (under increasing pressure or
temperature) or composition, and the equations of state (relating pressure p
to density p )*=- to adopt for each layer (crust, mantle, core).
Use is made of experimental and theoretical information obtained for the
Earth, from both geophysics (seismic studies, high-pressure physics, etc.)
and geochemistry (composition of rocks, ... ), and data from meteoritics.
The general method for constructing an internal density distribution
model for a planet is to integrate the hydrostatic equation inward, using the
assumed equations of state for each layer, and attempt, by varying the thick-
ness of each, to approximately satisfy the boundary conditions of mass, radius,
and moment of inertia. The planet is generally assumed spherical in shape,
except for the utilization of dynamical flattening or coefficient J? to derive the
moment of inertia. The product of the calculations, the model(s), is therefore
a density (p) versus radial depth (r) distribution, or a density profile p (r) of
the planet's interior.
Historical Density Models
Prior to late 1965, the exact value of the radius was still unknown.
Investigators who used lower values of radius (as low as 3310 km) concluded
that Mars was nearly homogeneous (i.e. , without a chemically distinct core
of appreciable size), such as the models of Urey (195Z), MacDonald (1962), and
Kovach and Anderson (1965).
Investigators who used the higher values of radius (as high as 3423 km)
concluded that Mars was chemically differentiated and composed of an iron-rich
(Fe-Si-Ni) core similar to Earth's, but smaller. Chemically differentiated
models were produced by Jeffreys (1937), Bullen (1949), and Kozlovskaya
(1966), A special or "hybrid" type of model, chemically homogeneous, but
*The polar moment of inertia is derivable from the observed coefficient J2,
or equivalently the dynamical flattening fd and the rotation rate w, by means
of Radau's relation assuming hydrostatic equilibrium.
**Other mechanical parameters related to density p (or pressure p) are also
involved, namely, the (adiabatic) incompressibility k, and the rigidity \i .
All these parameters depend, of course, on temperature and chemical
composition. Finally, the gravity intensity g, dependent on depth r, has to
be taken into account.
Sec. 2, page 10 C. Michaux, JPF November 15, 1971
JPL 606-1 Interior
physically differentiated into core and mantle (with the core a dense
high-pressure phase of the same silicates constituting the mantle), was
ingeniously conceived by Ramsey (1948) and worked out by Lyttleton (1963).
These models are very briefly reviewed by category.
Homogeneo us Models
Urey (1952) argued, by comparing the terrestrial planets, that Mars must
be roughly homogeneous (or at least much less differentiated than Earth), con-
sisting of a mixture of iron and silicate phases uniformly dispersed. He esti-
mated from a density pressure relationship that Mars must contain about 30''o
iron phase.
MacDonald (1962) calculated numerous possible models of Mars' interior
by varying the structural assumptions of radius, flattening, and mantle phase
changes. He showed that core radius and surface density depend critically
upon the value of (hydrostatic) flattening and on planet radius. Thus, a slight
decrease in flattening leads to a considerable increase of core radius (and
mass); while a small decrease in planet radius leads to unduly large surface
densities. Adopting a mean radius of 3345 km and a flattening of 0.0050 for
Mars, MacDonald finally proposed two alternative models, which are both
core models. The first, MI, is a small core model, where its mass amounts to
about 1% of total planetary mass and radius is about 700 km, and was derived
by assuming for the mantle a phase transition creating an abrupt discontinuity
in its overall (mantle) density distribution. The second model, MH, is a larger
core model, where core mass is about 10% (exactly, 9.3%) of total mass and its
radiu<^ about 1250 km; but it has no phase transition assumption for the mantle.
The surface densities obtained were quite high, from 3.5 to 3.8 g cm"^. The
metallic core, composed mostly of Fe and Si, has a (compressed) density of
9 g cm-3. MacDonald favored the small core model, ML Both models are
shown in Fig. 1.
Kovach and Anderson (1965), starting from the chemically differentiated
Earth (iron-rich core and silicate mantle) as a model, investigated various
Mars models, in which core and mantle material were mixed in assigned pro-
portions. They were able to obtain a common overall (planetary) composition
only by utilizing a very low value (3310 km) for the Mars' radius. Thus, they
concluded that with the hypothesis of chemical differentiation, it is impossible
to have a Mars similar to Earth in overall chemical composition, if its radius
is much greater than 3310 knn.
Chemi cally Differentiated Models
Jeffreys (1937), who did the initial work on interiors of terrestrial planets,
actually proposed two models for Mars, illustrated in Fig. 1. One model, JI,
resembled the Earth with three layers, while the other, JII, had only two
layers, and no metallic core. Jeffreys preferred the first model, JI.
November 15, 1971 C. Michaux, JPL Sec. 2, page 11
Interior
JPL 606-1
p
10
? 8
cr
a.'
6
4
■---J:^
^i^/
=.:b
\
\
4 'J
;,2;
*
'/
t
Ki
t I
3,2'!
'1* -^°
^
2000 ^ ^'5 4000
Oislance From Center, r, km
Jeffreys (1937)
6000
Mars I
10,3
5./
1" 1
4.23
_B_II
ui MaiR
ll&BM
3,?9
' -|387
3.361
2000
Distance From Center, r, km
Bullen (1949)
3.39 mm 14000
Mars
Ml
Mar<-
M.li
>>^^
t:
2000 4000
Distance From Center, r, km
MocDonald (1962)
Fig. 1(a). Density models of Mars' interior
3000 dL 4000
RADIUS, km
Fig. 1(b). Theoretical interior density profiles for Earth and Mars.
(Jeffreys, 1937, Bullen, 1949, Lyttleton, 1963).
Sec. 2, page 12
C. Michaux, JPL
November 15, 197 1
JPL 606-1 Interior
Bullen (1949) also calculated two models. One model, BI, with a
differentiated Fe-Ni core, and the other model, BII, which was a modification
of Ramsey's Earth model by using a quadratic law of incompressibility (k)
versus pressure (p), instead of the usual linear law. Bullen preferred the first
model, BI. See Fig. 1.
Kozlovskaya (1966) also (like many others up to Kovach and Anderson,
1965) attempted to prove the hypothesis of a common overall chemical composi-
tion for the terrestrial planets (Mars, Venus, and Earth). Starting with an
Earth model of the "hybrid ' type (following Ramsey, with a core of silicates
in metallic state), he found that Mars on the whole must be somewhat denser
than Earth. 'Mars material'' could be obtained by adding 5 to 8% iron to
Ea rth material .
Hybrid Models
Ramsey (1948), at Jeffreys' suggestion, undertook the investigation of a
two-layer (or two-zone) model with identical chemical (silicates) composition
throughout, but with differentiation into a core of silicates in a high-pressure
"metallic" state and a mantle of the same silicates in their normal (low-
pressure) molecular state. The calculations to confirm his model were not
completed. Ramsey's intent was to explore the broader hypothesis that all
four terrestrial planets were of the same basic (silicate) composition.
Lyttleton (1963 and 1965 a, b) calculated a series of two -zone models of
Mars, from the known mass, on the hypothesis that its composition is similar to
that of the Earth, with the interior zone representing a phase -change produced by
both pressure and temperature effects, thereby extending the Ramsey hypothesis.
Because of the much lower central pressure in Mars, which on almost any
model must be less than 0.3 X 10^2 dyn cm-2 (about 1/10 that for the Earth),
Mars can consist of two zones only — an inner one of solid material in the same
high-pressure phase as the present mantle of the Earth (below 413-km depth),
and an outer one of solid material in the same form as the present outer shell
of the Earth (above 413-km depth). According to this theory. Mars would be
entirely solid without a liquid metallic core. Consequently, despite the closely
similar angular velocity and obliquity of Mars and Earth, this theory predicted
the absence of any main Martian magnetic field. This prediction was subse-
quently confirmed by Mariner 4 data obtained in July 1965.
The internal temperature is regarded as arising from release of radio-
actively produced energy, and, because of its somewhat smaller size, Mars
might be expected to be at slightly lower temperatures (at comparable depths)
than in the Earth. The pressure at which the phase-change (corresponding to
the 20-degree discontinuity in the Earth) occurs is known to be highly sensitive
to temperature (Ringwood, 1962). For this reason, the interface-pressure in
Mars is not at present precisely determinable. If the interface pressure were
known, a unique structure for the planet would emerge solely from the mass.
To comply with the "best" observed radius - 3933 km (given by Mariner 4,
6, and 7) — the interface pressure would be about 0,07 X 10l2 clyn cmi"2 (com-
pared with 0.14 X 10l2 dyn cm"2 in the Earth), and the resvilting structure
would have about 60% of the total mass in the central region in mantle form
R, A, Lyttleton, R. Newburn,
November 15, 1971 C. Michaux, JPL Sec. 2, page 13
Interior JPL 606-1
and 40"'o in the outer shell, with the interface occurring at a depth of about
560 km. This configuration is also consistent (within the limits of error) with
the dynamical flattening derived from the satellite motions. However, closer
limits to the value of this quantity and to the radius would supply a more
stringent test of the theory.
The theory further suggests that with rising internal temperature (with
radioactive energy release), the depth-level of the phase-discontinuity will
increase to provide the requisite higher pressure, and, as a result, the planet
will have undergone slight expansion, which may still be occurring. The
amount of expansion is uncertain but could well be of the order of 10 km in
radius during the life of the planet, with a consequent increase of surface area
in the order of 10° km^. This might well result in rifting of the extreme outer
layers of Mars, although this might be disguised by effects of subsequent (long-
term) surface modifications by weathering and erosion processes.
Recent Density Models
Binder's (1969) Models
The upper n:iantle of the Earth is currently considered to be composed of
olivine 1= chrysolite (a synonym), or peridot (the gem variety of olivine)],
represented by the miineral series (Mg, Fe)-, SiO^,
Olivine is a solid solution of forsterite (pure Mg^SiO. ) and fayalite (pure
Fe2Si04). The composition of olivine, which therefore miay vary between these
two extremies, is usually represented as mole percent of these extreme con-
stituents abbreviated as Fo and Fa. Thus, for example, Fooq denotes olivine
composed of 80% forsterite and 20% fayalite.
Recent findings on the comiposition and structure of the Earth's upper
mantle by Anderson (1967) have shown that
1) There are actually two (and not one) density discontinuities, be-
ginning at depths of 365 and 620 km, which can be correlated with
changes in the packing of olivine to a spinel and then a post-spinel
structure. These packing transitions (and sudden increases in
density) are dependent on pressure, temperature, and the Fo/Fa
ratio.
2) The composition of the mantle slowly changes fromi Foqq near the
surface to Fo^q at about 700 km, below which depth it apparently
remains constant. This finding has been substantiated by Press
(1968).
These compositional and structural data on the terrestrial mantle were
applied directly to the Martian miantle in the nnodels of Mars' interior con-
structed by Binder (1969). Therefore, the mantle of Mars was assumied by
Binder to have three possible regions (olivine region, spinel region, and post-
spinel region), each of constant composition characterized by the Fo/Fa ratio
(decreasing or stationary at the discontinuity). The mantle was also assumied
R, A. Lyttleton, R. Newburn,
Sec. 2, page 14 C. Michaux, JPL November 15, 1971
^ — -
JPL 606-1 Interior
to be "pure, or nearly pure, silicates and that all the free iron or iron-nickel
has been differentiated into a nearly silica-free core" (Binder, 1969). The
Fe-Ni core material equation of state (pressure-density function p = f(p), where
the incompressibility k enters) was derived from data on the density of Earth's
core (Jacobs, 1963), allowing a 3% decrease in density to account for a solid
rather than liquid Martian core. Finally, the Martian crust was assumed to
have a density (2.8 g cm"^) and thickness (20 km) similar to the Earth's crust.
Important constraints imposed on the models were
1) Internal (core) temperature limits of 700° and 1500°C. The lower
limit is derived from lunar analogy. The insignificant magnetic
field results (Ness et al. , 1967) are taken as evidence that the
Moon's interior is not above 700''C. The upper limit is derived
from the assumption that lack of an appreciable magnetic field for
Mars implies the lack of a fluid-conducting core, and therefore a
temperature of its postulated Fe or Fe-Ni core below the melting
point of these metals (Fe mp: IBBO'C uncompressed).
2) Accurate modern values of the Mars radius (Rg = 33 94 ±5 km)
and dynamical flattening (f^ = 0.00525 given by Cain, 1967), both
provided by the Mariner 4 flyby. These also yield a moment of
inertia factor C/MR^ of 0.377 ±0.002, if the Mariner 4 mass of
Mars (M = 6.423 • lO^o g) ig used.
Binder's computational model results for Mars are shown in Fig. 2. The
mantle compositions varying between Fo^g and Foyg are similar to those of
Earth's mantle (averaging Fo^5). However, their structure is somewhat dif-
ferent, with the phase transitions occurring at different depths than those for
Earth. * Binder also evolved another model where the mantle composition is
intermediate between Earth's and meteoritic olivine. These models have a
metallic (Fe, or Fe-Ni) solid core of compressed density about 8.5 g cm"^,
and a radius ranging from 790 to 950 km, for the most probable models, thus
representing a fractional planetary mass ranging from 2.7 to 4.9%. (Note: The
actual range for all models was core radius 680 to 1050 km, and core frac-
tional mass 1.7 to 6.3%.)
Anderson's (1972) Models
Anderson (1972) reconsidered the question of differentiation for Mars
after the Mariner 4 data on mass and radius were fully confirmed by Mariners 6
and 7. Although the Binder (1969) models, based on the Mariner 4 data, had
indicated the likelihood of a sizable metallic (Fe-Si-Ni) core, the thermal
history models of Anderson and Phinney (1967) and Hanks and Anderson (1969)
seemed to thermally restrict the development of such a core which requires
high melting points (of Fe, Si, Ni alloys) to be reached at some time. The re-
quirement (taken as a constraint) of a lower melting point material for the
Martian core than the usual Fe-Si-Ni mixture (assumed for cores of all ter-
restrial planets) can be satisfied by replacing the light element Si with sulfur.
=:=Because of the slower increase in pressure with depth for Mars.
November 15, 1971 C. Michaux, JPL Sec. 2, page 15
Interior
JPL 606-1
J
<^
1/MR*
= 0.3T7
TEMP
-. Z 8 •/.
.2T
" 5
1500
20O0
DEPlM . KM
Fig. 2(a). Density (p), gravity (g), and pressure (p) profiles of a model
representative of the models that satisfy the moment of inertia and
temperature limits and that have mantle composition similar to
the connposition of the nnantle of the Earth (Binder, 1969).
"^
^
^ ^^
X
-
/
3
^N.
^ m -^-^ ^''^
>■ 4
\
Z
^ ^
\
O
^
\
/
\
X
\
\
\
X
^^ 1/MR'iO 377
\
^ «(■'"?' *''"*'"
\
y ^^*"\
\
/
\
\
/ , , . ,
1 N
1500 20O0
DEPTH, KM
Fig. 2(b). Density (p), gravity (g), and pressure (p) profiles of a model
representative of the models that satisfy the mioment of inertia and
temperature limits and that have a mantle composition interme-
diate between the composition of the terrestrial mantle and
the composition of olivine in meteorites (Binder, 1969).
Sec. 2, page 16
C. Michaux, JPL
November 15, 1971
JPL 606-1 Interior
which is cosmically abundant yet produces an alloy (Fe-S-Ni) that melts at a
lower temperature. The possibility of sulfur as a major element in the core
has been discussed for the Earth by Murthy and Hall (1970), and for both Earth
and Mars by Anderson et al. (1971), who concluded that sulfur should be abun-
dant in the core of small, relatively cold planets like Mars.* An examination
of the binary phase diagram for the system Fe-FeS (see Fig. 3) indicates that
the eutectic Fe-FeS melts at a low temperature, near 990''C, and is relatively
insensitive to pressure, according to the Brett and Bell (1969) experiments
up to 30 kb. Thus, by assuming an FeS or S rich meteoritic (chondritic) com-
position for Mars initially, a core can start forming at temperatures exceed-
ing this low eutectic temperature. However, core formation cannot proceed to
completion (total differentiation), or nearly so, unless most of the planet's
interior exceeds the liquidus temperature for some time.
On this basis, Anderson (1972) constructed meteoritic (chondritic) models
for Mars interior; two of which, in particular, gave a sizable Fe-S-Ni core--
of radius nearly half (R(,/R~0.45) that of the planet--containing about 12% of
the total mass of Mars and a mantle rich in (the remaining) Fe or FeO. The
total Fe content of Mars would be about 25-28%, close to that of chondrites, **
but less than Earth's 35% Fe, and this regardless of assumptions on overall
composition or distribution of Fe in the Mars models. The implied FeO con-
tent of the mantle is 21-24% by mass. Thus, they concluded from their chon-
dritic models (that these suggest) that Mars: (1) should have a smaller and
lighter (less dense) core than Earth, but definitely richer in sulfur and poorer
in silicon; and (2) would have a denser mantle rich in FeO or Fe (and Ni).
The details of Anderson's (1972) two main representative meteoritic
models are as follows:
1) One model ("Model II") had the starting composition of ordinary
chondrites (11.7% Fe, 1.3% Ni, 5.9% FeS; i.e., 19% potential core
forming material). Upon heating, most of the sulfur (FeS) melts
and settles into a core, which collects about 63% of the available
potential core-forming material, while much Fe + Ni remains solid
in the mantle nnixed with the silicates. The resulting density of
the core is 5.85 g cm'^, and that of the mantle is 3.54 g cm'^ (at
zero-pressure). This model is incompletely differentiated.
2) An alternative model ("Model I"), composed of a mixture of 75%
carbonaceous chondrites Type III and 25% ordinary chondrites, and
accommodating higher temperatures, yielded complete separation
of core and mantle, which have densities of 5.78 and 3.49 g cm"-^,
respectively. The mantle here has high FeO, from the Type in
carbonaceous chondrite. The model is fully differentiated.
*The idea is not new; Murthy and Hall mention that Fish et al. (I960), and
Urey (1966), had "pointed out that among the major phases in meteorites, an
iron-sulfur melt would be the first to be produced on heating. "
♦♦Ordinary chondrites have 17% free Fe and 5% FeS; carbonaceous chondrites
have little free Fe but 7-25% FeS.
November 15, 1971 C. Michaux, JPL Sec. 2, page 17
Interior
JPL 606-1
1600°
(at 30kb)
'1534°
(atOkb)
CC
CC
UJ
Q.
'900'
earth's core
Fe + Liquid
•1190°
IRON RICH MANTLE
SULPHUR RICH CORE
Eutectic Temp
Liquid + FeS
Fe + FeS (Solid)
Fe
FeS
Fig. 3.
COMPOSITION (wt % Fe)
Schematic binary phase diagram for the system
Fe-FeS (Anderson, 1972).
Figure 4 shows a wide range of possible Mars models (not necessarily of
meteoritic composition) which are consistent with the presently known values of
radius R, mass M, and moment of inertia C. It can be seen that the more
plausible models, those of chondritic compositions (such as Models I and II)
only occupy a narrow (hachured) region in the plot of density versus fractional
radius. Note that all models have about the same mantle density, while core
density and size (fractional radius) have wide, inverse variations.
Implications of Anderson's (1972) Chondritic Models
Planet formation . The greater ability of Mars to retain sulfur (than
Earth) would be explainable in terms of lower accretional energies and tem-
peratures during its growth, which was probably slower than the Earth's. While
the Martian core is rich in sulfur, it is poor in silicon, according to the models.
High accretional temperatures and rapid solidification would be necessary for
a Martian core to be rich in silicon. (See Anderson et al. , 1971).
Magnetic field . The observations of a negligible magnetic field (by
Mariner 4) and a rapid rotation rate for Mars have usually been interpreted as
evidence against the presence of a molten iron-rich core, as in Earth, where
a dynamo effect takes place generating the field. The exact mechanism driving
the geomagnetic dynamo has not been ascertained yet, but recent views (e. g. ,
Malkus, 1968) attribute the driving force to the differential precessional torque
Sec. 2, page 18
C. Michaux, JPL
November 15, 1971
JPL 606-1
Interior
ro
o
\
C7>
CO
Q
4 —
MARS
(^moluiclmis
Fe
Earth's Core
-f-i-4f-,
- CORE
Eutectic mix J
P^30 kb
P= kb
FeS
MANTLE
Earth's Mantle
Fig. 4.
2
0.4
Fractional
0.6
Radius
8
Possible density models of Mars satisfying proper mean
density and moment of inertia of the planet,
calculated by Anderson (1972).
acted upon the Earth's core and mantle, primarily by the large, close Moon,
and to a lesser extent by the Sun, In the case of Mars, because there is no
large Moon, and the Sun is further away, an effective torque may not be
present. Other reasons limiting the dynamo action were also advanced by
Anderson (1972): the smaller size of the (proposed) Martian core, and, if
sulfur is abundant, its high resistivity which would lower the magnetic Reynolds
number; and, possibly, the low viscosity of the core (under lower prevailing
temperatures and pressures), and the shielding effect of an iron-rich mantle.
Anderson (1972) considers it is "unlikely that Mars ever had a substantial mag-
netic field of internal origin. "
2. 3 Thernnal History Models
Introduction
Calculations of the thermal history of the interior of a planet can aid
greatly in understanding the evolution of the whole planet, from its formation
down to the present. * The study of the thermal history of the terrestrial planets
*At least they permit interesting speculations; see for example Fanale (1971).
November 15, 1971 C. Michaux, JPL Sec. 2, page 19
Interior JPL 606-1
was revived some ten years ago when accretional theories^:- of planet formation
at relatively low temperatures had gained wide acceptance. New thermal
models were constructed first for the meteorites and asteroids (Allan and
Jacobs, 1956), then for the Earth (Lubimova, 1958), the Moon (Levin, 196Z),
and the terrestrial planets (MacDonald, 1959, 1962). The method for construct-
ing such models has been reviewed by MacDonald (1959). Because many
assumptions are required, the calculations are highly speculative. The initial
conditions prevailing upon accretion, as well as the subsequent mechanisms of
heat generation, all have to be prescribed from meagre data and uncertain dis-
cussions. Extending the calculations for the Earth and using conclusions de-
rived therefrom, similar miodels were constructed for the other terrestrial
planets. However, it was soon realized that the Earth analogy should not be
carried too far, especially in the case of Mars, because of its much smaller
mass. Later, to meet certain geochemical constraints, a challenging 'hot
origin' theory was proposed (for the Earth especially) by Ringwood (1966) which
assumies: rapid accretion followed imimediately by core formation or differen-
tiation. This 'hot origin' hypothesis was not widely accepted. Models were
then developed by Anderson et al. to study the sensitive interplay of accretion
and radioactive abundances upon differentiation, for both Earth and Mars. The
Anderson miodels appear to point toward this hypothesis for Earth, and perhaps
for Mars, unless a sulfurized core is assumed (Anderson, 1972).
Some Planetary History
Before reviewing thermal history models, it is instructive and helpful to
briefly describe the origin and evolution of a typical Earth-like planet, as com-
monly understood. This evolution, it will be seen, is chiefly governed by the
generation and release of thermal energies at various times, and the possibility
of chemical differentiation.
It is now generally believed that most bodies of the solar system origi-
nated through the gravitational agglomeration, "accretion, " of a primordial
mixture of dust and gases under relatively low temperatures (100° - 1000°K)
and in relatively short periods of time (1 - 100 million years). This major
early accretional phase produced the so-called ''protoplancts, " which were
still, relatively cold, undifferentiated bodies. The protoplanets ' "initial" tem-
perature was primarily dependent on the accretion rate (or total accretion
time), being much higher if the rate were faster. Other lesser heat sources,
besides loss of gravitational energies, also contributed to temperature; namely,
adiabatic self-compression, and the decay of short-lived radioisotopes (such
as I^^^, with half-lives < 10^ years). Subsequently, the temperature of the
protoplanet continued to rise internally, due to the decay of long-lived radio-
isotopes (K^^, Th232^ U^35 a,nd U238)^ and the difficulty for the generated heat
to flow toward the surface and escape into space (mostly rocky material having
low thermal conductivity and high opacity). After some time, the internal tem-
perature became so great as to reach the melting point of iron (and similar
=:'These accretional theories were developed through the work on the origin of
the solar system by such men as Urey (1952), Alfven (1954), Hoyle (I960),
Cameron (1962), Fowler et al. (1962), etc.
Sec. 2, page 20 C. Michaux, JPL November 15, 1971
JPL 606-1 Interior
metals), and "differentiation" began. At this point, molten iron descended by
gravity toward the center of the planet to form a heavy metallic "core, " while
the rocky silicate material ascended to form the "mantle. " This separation or
core formation converted more gravitational energy into heat, accelerating the
differentiation process to completion in a shorter time. This resulted in a hot,
differentiated planet.
The so-called "early thermal history" of a planet commences with accre-
tion and ends just before differentiation or the beginning of core formation. The
"later thermal history" starts with core formation and covers the total differ-
entiation stages of the planet. This evolution reflects the classical views.
Recent Thermal Models
Anderson and Phinney (1967)
Anderson and Phinney made early thermal history calculations for undif-
ferentiated planets of the mass of Earth and Mars, starting (initially) with a
surface temperature T(0, 0) of either 400°, 330°, or 100°K, and assuming
homogeneous distribution of the long-lived radioactive elements throughout the
protoplanet, with average abundances called "terrestrial"* by geophysicists
(instead of the chondritic abundances previously used). The calculated profiles
showed that for the case T(0, 0) = 400°K, the melting point of iron is reached
at the center of both planets, some 2,109 years after accretion for the Earth,
but 4.46 X 109 years after accretion (i. e. , the present time) for Mars. For
the case where T(0, 0) = 330°K, no melting of iron occurs for Mars (regard-
less of the lattice conductivity postulated), while there is still some iron melt-
ing at Earth's center. For the very low temperature T(0, 0) = 100 °K, neither
Mars nor Earth can reach the melting point of iron, and both remain undif-
ferentiated. Their conclusion was that "it is fairly easy to differentiate an
Earth, but difficult to differentiate a Mars" if the estimates of initial tempera-
tures and radioactive abundances are close to true conditions.
Hanks and Anderson (1969)
Hanks and Anderson calculated early thermal histories of both Earth and
Mars for comparison purposes, assuming their overall compositions and accre-
tional histories to be roughly similar. Their main objective was to determine
the early thermal conditions for Earth permitting core formation and large-
scale differentiation to occur within the first 1.1 billion years following accre-
tion, in order to antedate the oldest known terrestrial rock (3.4 billion years
old). The use of Mars as a comparison planet (rather than Venus) was of
=:=Specifically, they took the ratios K/U = 10"^ and Th/U = 3.7, determined
from terrestrial rocks by Wasserburg et al, (1964), then "calibrated" these
ratios to the present U content in Earth's crust and nnantle, 4.5 X 10-8 g g~ ,
as proposed by MacDonald (1964). Finally, they reduced the values obtained
to account for dilution by mixing with the present core assumed free of
radioactivity.
November 15, 1971 C. Michaux, JPL Sec. 2, page 21
Interior JPL 606-1
interest due to the fact that its mass is much smaller than Earth's, and
therefore the gravitational energies released (converted to heat) upon accretion
are also much smaller. In addition, since astronomical data showed Mars
appeared to be nearly homogeneous or undifferentiated, it would provide a
sensitive test for the hypothetical onset of differentiation (melting of iron and/or
silicates) in a smaller body.
They first constructed two Mars models patterned after two Earth models,
which differed in their radioactive abundances and total accretion times (ta,cc)-
One model. Mars ID, composed of the usual chondritic abundances and t^^^, -
500,000 years, and a second model. Mars IID, with the so-called "terrestrial"
abundances (of Wasserburg et al. 1964), scaled down by MacDonald (1964) and
reduced by Anderson and Phinney (1967) and tg^^c = 300,000 years. The result-
ing thermal profiles showed that the Mars ID model produced molten iron in
2.1 billion years and the Mars IID model in 3,0 billion years after accretion,
see Fig, 5, They concluded that the assumed radioactive abundances were too
high for Mars, if it is still undifferentiated today. To prevent melting of iron
in Mars, they then reduced the terrestrial abundances to 0,8 of their original
values and proceeded to recalculate profiles for a Mars II'D model, whereby
it would take 4.7 billion years after accretion to produce molten iron. See
Fig, 6 (shown originally without the FeS liquidus and Fe-FeS eutectic dashed
lines).
Unfortunately, at the time, the authors did not consider the alternative
possibility of a differentiated Mars and further investigate the thermal condi-
tions leading to the formation of a Martian core of lighter composition than
Earth's core. Later, however, Anderson (1972), in the light of his Fe-S-Ni
core model for Mars, added the dashed lines for the low-melting Fe-FeS
eutectic and for the FeS liquidus data to Fig. 6, and stated that "much of the
interior of Mars is between the eutectic and liquidus temperatures, which
suggests partial, rather than total, melting and incomplete separation of
potential core forming material. " By considerably lowering the temperature
necessary for core formation in Mars interior, the Anderson (1972) scheme
involving svilfur also reduces the time required for the onset of differentiation -
to perhaps 3 X 109 years in Mars II'D (see Fig. 6). According to Anderson
(1972), the formation of a sulfurized Martian core would release only a
"negligible amount of heat" (primarily because of the small mass of the planet).
The differentiation process would then be much slower than for the Earth.
The thermal models presented above implicitly presuppose homogeneous
accretion; i. e. , accretion of primordial matter into an essentially homogeneous
protoplanet with differentiation subsequently taking place. Recently, a new class
of models was proposed by Turekian and Clark (1969), based upon inhomogene-
ous accretion. As the primitive nebula cooled down, successive condensations
of elements and components (in roughly the order of increasing vapor pressure
and decreasing density) formed the protoplanet presumed to be initially layered
with a central core of Fe and Ni. Such a model has in fact been presented for
the Earth by Clark, Turekian, and Grossman (1972), and also for the Moon by
Hanks and Anderson (Hanks, 1972). The application of the inhomogeneous
accretion concept to Mars would produce direct core formation (a small core)
without need of resorting to the melting-by-sulfur differentiation concept.
Sec. 2, page 22 C. Michaux, JPL March 1, 1972
JPL 606-1
Interior
3000
^ 2000
lOOO
1 I i I i r
Iron Melting Curve (Strong. 1959)
0-
Mors ID
(Chondritic Abundonces)
'occ "5" 10' yeors
_1_
_!_
_L
_l_
JL
_1_
3000
2000
1000
02 04 06 08 10
Fraction of Mors Radius
02 04 06 06 LO
~i 1 1 T r
Iron Melting Curve (Strong. 1959)
(Terrestrial obundonces)
•act "^'lO' yeors
Fig, 5. Temperature profiles as a function of time (in billion years) for
thermal history models Mars ID ("chondritic" radioactive abundances
and t
ace
5 X 10^ years) and Mars IID ("terrestrial" radioactive
abundances and t
ace
3 X 105 years),
(Hanks and Anderson, 1969)
3000
2000
o
a>
CL
E
^ 1000
MARS n'D tQccr^xlO^ yeors
(0.8 Terrestrial obundonces)
_FeS liquidus\
ZJ
2^
Ve-FeS
b^
Iron Melting Curve
(Strong, 1959)
0.2
0.4 0.6
Fractional Radius, r/R
0.8
.0
Fig. 6, Temperature profiles as a function of time (in billion years) for
thermal history model Mars II'D (reduction to 0.8 of the "terrestrial"
radioactive abundances and same t^^c = 3 X 10^ years as in Mars IID).
(Hanks and Anderson, 1969)
Note : Dashed lines for FeS liquidus and Fe-FeS eutectic were added later
by Anderson (1972), in consideration of his new differentiated internal
density models. No revised profiles have yet been calculated to take
into account differentiation possibilities.
March 1, 1972
C. Michaux, JPL
Sec. 2, page 23
Interior JPL 60b- 1
BIBLIOGRAPHY
Alfven, H. , 1964, On the origin of the solar system, Oxford U. Press.
Allan, D. W. , and Jacobs, J. A. , 1956, The melting of asteroids and the origin
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Anderson, D. L. , 1972, Internal constitution of Mars: J. Geophys. Res. , v. 77
(5), p. 789-795, February 10.
Anderson, D. L. , 1967, Phase changes in the upper mantle: Science, v. 157,
no, 3793, p. 1165-1173, Septembers.
Anderson, D. L. , and Phinney, R. A. , 1967 (584 p. ), Early thermal history of
the terrestrial planets, Chapter 3 in Mantles of the Earth and terrestrial
planets; Runcorn, S. K. , Editor : New York, Interscience Publishers,
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Anderson, D. L. , Sammis, C. , and Jordan, T. , 1971, Composition and evolution
of the mantle and core: Science, v. 171, no. 3976, p. 1103-1112, March 19-
Binder, A. B., 1969, Internal structure of Mars: J. Geophys. Res. , v. 74 (12),
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Cain, D.L., 1967, The implications of a new Mars mass and radius, p. 7 -9 in
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Cameron, A. G. W. , 1962, The formation of the Sun and planets : Icarus, v. 1 (1),
p. 13-69, May.
Clark, S. P. , Jr. , Turekian, K. K. , and Grossman, L. , 1972, Model for the early
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de Vaucouleurs, G. , 1964, Geometric and photometric parameters of the ter-
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Fanale, F. P. , 1971, History of Martian volatile s: implications for organic
synthesis: Icarus, v. 15 (2), p. 279-303.
Sec. 2, page 24 C. Michaux, JPL March 1, 1972
JPL 606-1 Interior
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Fowler, W. A., Greenstein, J. L. , and Hoyle, F. , 1962, Nucleosynthesis during
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communication to C. Michaux.
Hanks, T. C. , and Anderson, D. L. , 1969, The early thermal history of the
Earth: Phys. Earth Planet. Inter. , v. 2 (1), p. 19-29, April.
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distribution: Abco/Rad.
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York, McGraw-Hill Book Co. , 2nd edition.
Hoyle, F. , I960, On the origin of the solar nebula: Quart. J. Roy. Astron. Soc. ,
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September.
March 1, 1972 C. Michaux, JPL Sec. 2, page 25
Interior JPL 60b- 1
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Soc. , V. 130 (1), p. 95-96.
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Sec. 2, page 26 C. Michaux, JPL November 15, 1971
J PL 606-1 Interior
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November 15, 1971 C. Michaux, JPL Sec. 2, page 27
JPL 606-1 Surface
SECTION 3 CONTENTS
SURFACE
3. 1 Thermal Properties
Introduction l
3. 1. 1 Theoretical Temperatures 1
Thermal Models — General 2
Atmospheric Effects 2
Leighton-Mur ray Thermal Model 3
Kieffer Thermal Model 5
3. 1. 2 Infrared Radiometry 6
Infrared Radiometry From Earth 10
Infrared Radiometry From Spacecraft 13
3.1.3 Microwave Radiometry 21
Observations 22
Bibliography 26
Figures
1. Theoretical curves of diurnal variation of surface tem-
perature near the equator at southern spring equinox 4
2. Theoretical curves of annual variation of average tem-
perature of the disk as viewed from the Sun for various
depths 4
3. Average annual surface temperature as a function of
latitude 5
4. Diurnal variation curves of surface temperature at
perihelion for latitudes 0°, ±30°, ±45° and ±60° 7
5. Diurnal variation curves of surface temperature at
aphelion for latitudes 0°, ±30°, ±45° and ±60° 8
6. Earth: atnnospheric transmission 9
7. Radiometric drift-curves across Mars on July 20, UT 11
8. Equatorial diurnal temperature variation 12
9. Brightness temperature map of Mars 12
10. Theoretical temiperature distribution with latitude and
local time for Mars at the equinox 13
11. Variation of temperature with latitude 14
12. Surface kinetic temperatures versus local time, obtained
by the Mariner 6 IRR swaths 15
13. Surface kinetic temperatures versus local time, obtained
by the Mariner 7 IRR swaths. 1£,
14. Surface kinetic temperatures along 2 swaths of Mariner 6
between latitudes 5 °N to 20 °S 17
15. Surface kinetic temperatures along 1 swath of Mariner 6
between latitudes 20 °N to 15 °S 18
March 1, 1972 Sec. 3, Contents, page i
Tables
Surface JPL 606-1
3. 1 (cont'd)
16. Surface kinetic temperatures along 2 swaths of Mariner 7
between latitudes 10 °S and 45 °S 19
17. The radiowave spectrum of Mars, as plotted by Epstein
(1971) after initial evaluation of all available derivations .... 23
I. Disk-averaged brightness temperatures of Mars 24
3. 2 Ultraviolet, Visible, and Infrared Photometric Properties
Data Summary 1
Disciission 1
3. 2. 1 Photometric Nomenclature and Theory 1
3. 2. 2 Reflection Versus Emission on Mars 4
3. 2. 3 Integrated Photometric Properties 6
Brightness, Opposition Effect, and Color 6
Phase Function 8
Geometric Albedo 9
Bond Albedo and Phase Integral 12
Bolometric Bond Albedo 13
3. 2. 4 Detailed Photometric Properties 13
Radiance Coefficient 14
Radiance Factor, Photometric Function, Normial Albedo, etc 14
Empirical Photometric Behavior 15
Results of Detailed Martian Photometry 16
3. 2.5 General Photometric Conclusions 19
3. 2.6 Polarimetric Nomenclature and Results 20
Introduction to Polarimetry 20
Observations 21
Bibliography 22
Appendix A — Martian Albedos A- 1
Appendix B — Glossary of Photometric and Polarimetric Terminology ... B-1
Figures
1. Reflected solar radiation versus emitted radiation
from Mars 5
2. Photometric coordinates 6
3. Local photometric geometry 15
4. Reference albedo (the average of the albedo at 1.04 ^J. and
1.24 (J. at 10.3 ° phase) 17
5. Spectral radiance factor at 5 ° phase for seven Martian
areas '■°
A-1. Martian areas studied by Binder and Jones A-2
Tables
1. Effective wavelengths in the UBV system 3
2. Infrared geometric albedos 11
A-1. Martian albedo A-3
Sec. 3, Contents, page ii March 1, 1972
JPL 606-1 Surface
3. 3 Radar Properties
Introduction 1
3. 3. 1 Basic Concepts of Radar Astronoiny 1
3. 3. 2 Cross -Section and Reflectivity 2
Fundamental Concepts 2
Observation Techniques ^
Early Radar Observations g
Recent Radar Ol^servations 10
3. 3. 3 Angular Scattering and Roughness 14
Fundamental Concepts 14
Backscattering Model Fitting 1 £,
Power Spectrum. Frequency Offset Measurement 16
Height Profile Differentiation 17
Experimental Results 17
Interpretation 17
3. 3. 4 Topography 20
Fundamental Concepts 20
Experimental Results 20
3.3.5 Topography — Cross -Section — Roughness Correlation 23
Bibliography 25
Appendix — Mars Radar Observations in 1971: Topography and
Radar Cross -Sections A- 1
Figures
1. The system of constant delay rings and doppler shift
strips on the disk 5
2. Samples of range-gated freqiiency power spectra set 7
3. Excursions in Martian latitude of the subradar point for
the 1963 to 1971 apparitions of Mars 8
4. The variation of the radar cross-section near the 22 °N
latitude as a function of the longitude of the centrnl
meridian of the visible disk as uL-'ained at 70-cm
wavelength 10
5. Radar cross -section variation with longitude near the 21 "N
parallel of latitude, as obtained at 12.5-cm wavelength H
6. Relative radar cross-section variation with longitude as
inferred from CW echo-Doppler spectrograms at 3.8 cm
wavelength 13
7. Relative radar cross-section variation with longitude near
four latitudes, as inferred from phase-coded (ranging)
measurements at 3.8 cm wavelength (bandwidth 1 kHz) 13
8. Relative radar cross-section variation with longitude near
several latitudes 3 °- 12 °N 15
9. Frequency power spectrum for Mars converted by Besscl
transformation for comparison with the Moon, Mercury,
and Venus Ig
10. Average doppler spectrogram and angular backscattering
curves at several wavelengths I9
11. Topography variation with longitude in four latitude steps
from 3 ° to 22 °N 22
March 1, 1972 Sec. 3, Contents, page iii
Surface JPL 606-1
3. 3 (cont'd)
IZ. Topography variation with longitude in latitudes 3°-12°N . . . . 23
A- 1(a). Topography variation with longitude near 16.5 °S latitude A-2
A-l(b). Topography variation with longitude near 16°S latitude A-2
A- 1(c). Topography variation with longitude near 15. °S latitude A-3
A-2. Topography and radar cross-section variation with
longitude near 14°S latitude A- 5
Tables
1. Doppler spread, or "limb-to-limb bandwidth" B of the
Martian echo, as a function of operating freqxiency
f (=f ' ) or wavelength Kq( \ q) 5
2. Comparison of radar cross -sections obtained by
various observers 1 "^
3. 4 Chemical and Physical Properties
Introduction
Physical Properties of the Ground Surface Material 1
Granularity ^
Density -^
Chemical Properties of the Ground: General Aspects 3
3.4. 1 Composition Inferred From Reflectance Spectrophotometry 4
3. 4. 2 Reflectance Spectra of the Bright and Dark Areas 5
Description and Interpretation of the Martian Spectra 7
Water of Hydration ^ 'J
Laboratory Simulation Experiments 10
Exotic Interpretations 1 ^
Distribution of Martian Surface Materials: Preliminary
Results 13
Stability of Goethite on Mars 14
3. 4. 3 Adsorption of Volatiles: CO2 and H^O 15
CO2 Adsorption in a Bright Area: Experiments and Calculations . . 16
H2O Adsorption: Experiments 16
3. 4. 4 Martian Permiafrost: Speculations 17
Permafrost From' Atmospheric Water: Calculations 18
Possible Frost-Heaving Caused by Atmospheric Water 18
3. 4. 5 Liquid Water • 18
Chemical Properties of the Polar Cap Deposit 21
3. 4. 6 South Polar Cap: Mariner 7 IRS Results 22
3. 4. 7 Near-Infrared Reflection Spectrk of CO2-H2O Frosts 22
3. 4. 8 Possible O3 Adsorption by the Polar Cap: Mariner 7 UVS Results . 24
3. 4. 9 Speculations on the Composition and Structure of the Polar Caps . . 24
Bibliography
Figures
1. Spectral geometric albedos of a typical bright area (Arabia)
and dark area (Syrtis Major) for the 0.3 to 2.5 \i spectral
region according to the data of McCord and Westphal and
other investigators o
2. Seasonal changes in dark areas 6
Sec. 3, Contents, page iv March 1, 1972
JPL 606-1 Surface
3. 4 (cont'd)
3. Ternary diagram (Ca, Mg, Fe++) Si03 showing the
compositional variations of pyroxenes 8
4. Simulation experiments with oxidized basalts for the
spectral geometric albedo of Martian bright and
dark areas , 22
5. Equilibrium vapor pressure curve for the goethite-
hematite system 15
6. Mean annual temperature as a function of latitude, with
indication of condensation temperatures of water vapor
for three atmospheric abundances 19
7. Depth of top surface of H2O permafrost, as a function of
latitude 19
8. Phase relationships of CO2 and H2O 20
9. Near-infrared spectra of the South Polar Cap by
Mariner 7 and comparison laboratory spectrum of solid
CO2 at 77 °K 23
10. Ratio of reflectance of polar cap to reflectance of a
desert region 25
11. Phase diagram of carbon dioxide hydrate 26
Tables
1. Estimates of average physical properties of Martian
surface material 2
3. 5 Morphology and Processes
Introduction j
3. 5. 1 Topography 2
Spectroscopic Methods 2
Infrared 2
Ultraviolet 2
Summary of Present Topographic Inforrriation 3
Interpretation g
3. 5. 2 New Mars Maps y
Mariner Mars 1969 Chart ] 9
International Planetary Patrol Photographic Maps of
Mars 1969 and 1971 U
Mars 1969 [ 1 1
Mars 1971 .......[.. 11
Mariner Mars 1971 Planning Charts 11
Charts for the South Polar Region and Cap 15
Mariner Mars 1969. I9
Mariner Mars 1969 Meridiani Sinus Region Map 1 9
Mariner Mars 1969 South Polar Region Map .' 20
3. 5. 3 Types of Terrains 21
Cratered Terrain 2 1
Mariner 4 Photography 21
Mariner 4 Crater Statistics and Analyses 23
Leighton et al. Analysis 23
Chapman et al. Analysis 25
March 1, 1972 Sec. 3, Contents, page v
Stirface JPL 606-1
3. 5 (cont'd)
Mariner 6 and 7 Photography 25
Mariner 6 and 7 Crater Statistics and Analyses 26
Murray et al. Analysis 26
Woronow and King Analysis 28
McGill and Wise Analysis 31
Crater Modification Processes 35
Age of the Large Craters 39
Chaotic Terrain 41
Distribution 41
Relative Age 41
Origin and Possible Processes 42
Featureless Terrain 43
Origin and Age of the Hellas Basin 44
Origin and Age of the Featureless Floor of Hellas 45
3.5.4 South Polar Cap 45
Mariner 7 Photography and Observations 47
Morphology 47
Processes 50
Marginal Zone 5
Polar Cap Interior 50
Central Polar Region 5
Thickness of Frost Cover 51
Permanence of Frost or Ice 51
3. 5. 5 Dark and Bright Areas: Boundaries and Markings 52
3. 5. 6 Canals and Lineaments 53
Canals 53
Lineaments 5 3
Oases 54
Bibliography 57
Figures
1. Mariner 6 UVS and IRS surface pressures and derived
altitudes 4
2. Mariner 7 UVS and IRS surface pressures and derived
altitudes 5
3. Mariner 6 UVS and IRS surface pressures and derived
altitudes 6
4. NASA Mars Chart 1969 10
5. 1969 Mars Patrol Photographic Map 12
6. 1971 Mars Patrol Photographic Map 14
7. Mariner Mars 1971 Planning Chart 16
8. Mariner Mars 1971 Planning Charts of South Polar Regions ... 18
9. Mariner Mars 1969 Photomap 20
10. Mariner Mars 1969 Meridiani Sinus Region Map 21
11. Mariner Mars 1969 South Polar Region Map 22
12. Cumulative size -frequency distribution of craters recognized
in Mariner 4 pictures 4N7-12 24
13. Cumulative size-frequency distribution of craters in
Deucalionis Regio 27
14. Plots of crater abundances for individual wide-angle and
narrow-angle frames 28
Sec. 3, Contents, page vi March 1, 1972
JPL 606-1 Surface
3. 5 (cont'd)
15. The Deucalionis Regio crater abundances of Fig. 13
compared with those of the lunar maria and the uplands 29
16. Cumulative size-frequency probability distributions of
craters found in wide-angle frames 30
17. Cumulative size -frequency probability distribution of
craters found in six narrow-angle frames 3
18. Size-frequency distribution of Martian craters in
four regions 33
19- Average degradation numbers for small craters and
large craters in four Martian regions 34
2 0. Summary plots contrasting distribution of small craters
among degradation classes in four Martian regions 35
21. Model explaining differences in degradation - density
curves for small craters from four Martian regions 36
22. Threshold drag velocities plotted over a range of particle
sizes for Mars and Earth 38
23. Lowest threshold wind velocities for Mars and Earth 39
24. Settling velocities over a range of particle sizes for Mars
and Earth 40
25. Interpretive map of chaotic-terrain distribution constructed
from Mariner 1969 photos 42
26. Diagram of the Hellespontus to Hellas transition zone as
viewed in 7N27 44
27. Mariner 6 far-encounter views of South Polar Cap,
enlarged to a common scale 46
28. Sketch of South Polar Cap: interior and central region,
morphological features appearing in 7N17 49
29. Rose diagram showing the azimuthal distribution of 868
lineaments mapped from Mariner 4 photographs 4N3-15 55
30. Rose diagrams of Martian lineaments at several latitudes 56
Tables
1. List of names used on the International Planetary Patrol
Photographic Map of Mars 1971 13
2. List of names used on the Mariner Mars 1971 Planning
Chart 17
3. Summary compilation of Martian crater data from
Mariner 4 pictures 4N3-16 23
4. Cr&ter percentages by class at several diameter intervals
for Mars 25
5. Classification of Martian craters by degradation number 32
3. 6 Mariner 1696 Photographic Atlas of Mars
Introduction 1
Television Experiment Design 1
Camera System 4
Image Processing 4
3. 6. 1 Far Encounter 6
Introduction 6
March 1, 1972 Sec. 3, Contents, page vii
Surface JPL 606-1
3. 6 (cont'd)
Mariner 7 — First Series 11
Mariner 6 — First Series and Mariner 7 — Second Series 14
Mariner 7 — Third Series and Mariner 6 — Second Series 22
3. 6. 2 Near Encounter 37
Cratered Terrain 41
Mosaic of Seven Camera A Frames 6N9 Through 6N23 41
Chaotic Terrain 52
Mosaic of Four Camera A Frames 6N1 Through 6N7 52
Lea: iix les s Terrain. 62
Mosaic of Frames 7N21 Through 7N31 62
Atmospheric Haze 69
Mosaic of Frames 7N1 and 7N3 69
Dark and Light Areas 72
Mosaic of Three A Frames 7N5, 7N7, and 7N9 72
South Polar Cap 76
Mosaic of Frames 7N1 1 Through 7N19 76
Bibliography 93
Figures
1. Spectral transmission, Mariner 6 camera A filters 2
2. Spectral transmission, Mariner 7 camera A filters 3
3. Transniission characteristics as a function of wavelength
for the camera B filter 3
4. The globe of Mars 7
5. Far Encounter Frame 7 F2 12
6. Far Encounter Frame 7F16 12
7. Far Encounter Frame 7F28 13
8. Far Encounter Frame 7F33 13
9. Far Encounter Frame 7F40 15
10. Far Encounter Frame 7F44 16
11. Far Encounter Frame 7F48 17
12. Far Encounter Frame 7F52 18
13. Far Encounter Frame 7F59 19
14. Far Encounter Frame 7F67 2
15. Far Encounter Frame 6F32 21
16. Far Encounter Frame 7F70 23
17. Far Encounter Frame 6F34 24
18. Far Encounter Frame 7F74 25
19. Far Encounter Frame 7F76 26
20. Far Encounter Frame 7F76 modified to show the positions
of Mariner 4 pictures 4N7 through 4N14 27
21. Far Encounter Frame 7F78 28
22. Far Encounter Frame 6F38 29
23. Far Encounter Frame 7F80 30
24. Far Encounter Frame 7F83 . 31
25. Far Encounter Frame 6F46 32
26. Far Encounter Frame 7F91 33
27. Far Encounter Frame 7F91 (magnified portion showing
Phobos) 34
28. Far Encounter Framie 7F93 35
Sec. 3, Contents, page viii March 1, 1972
JPL 606-1 Surface
3. 6 (cont'd)
29. Far Encounter Frame 6F49 36
30. Mariner 6 picture locations on a painted globe of Mars 38
31. Mariner 7 picture locations on a painted globe of Mars 39
32. Mosaic 6N9 through 6N23 42
33. Near Encounter Frame 6N11 43
34(a). Near Encounter Frame 6N13 (Max-D version). . 44
34(b), Near Encounter Frame 6N13 (photometric version) 45
35. Near Encounter Frame 6N17 46
36. Near Encounter Frame 6N18 47
37. Near Encounter Frame 6N19 48
38. Near Encounter Frame 6N20 . . . ,, 49
39. Near Encounter Frame 6N21 50
40. Near Encounter Framie 6N22 51
41. Mosaic 6N1 through 6N7 53
42. Near Encounter Frame 6N3 54
43. Near Encounter Frame 6N5 . , . 55
44. Distributions of light and dark markings and chaotic
terrain in equatorial region photographs 6N5, 7, and 9 56
45. Near Encounter Frame 6N6 57
46. , Near Encounter Frame 6N7 58
47. Near Encounter Frame 6N8 59
48. Near Encounter Frame 6N14 60
49. Near Encounter Frame 6N15 61
50. Mosaic 7N21 through 7N31 63
51. Near Encounter Frame 7N25 64
52. Near Encounter Fraine 7N26 65
53. Near Encounter Frame 7N27 66
54. Near Encounter Frame 7N28 67
55. Near Encounter Frame 7N29 68
56. Mosaic 7N1 and 7N3 69
57. Near Encounter Frame 7N1 70
58. Near Encounter Frame 7N2 71
59. Mosaic 7N5, 7N7, and 7N9 72
60. Near Encounter Frame 7N5 73
61. Near Encounter Frame 7N6 ■ 74
62. Near Encounter Frame 7N7 75
63. Mosaic 7N11 through 7N19 (photometric version) 76
64. Mosaic 7N11 through 7N19 (Max-D version) 77
65(a). Near Encounter Frame 7NI1 (Max-D version) 78
65(b). Near Encounter Frame i7Nll (photometric version) 79
66. Near Encounter Frame 7N12 80
67(a). Near Encounter Frame 7N13 (Max-D version) 81
67(b). Near Encounter Frame 7N13 (photometric version) 82
68. Near Encounter Frame 7N14 ; 83
69(a). Near Encounter Frame 7N15 (Max-D version) 84
69(b). Near Encounter Frame 7N15 (photometric version) 85
70. Near Encounter Frame 7N16 86
71(a). Near Encounter Frame 7N17 (Max-D version) 87
71(b). Near Encounter Frame 7N17 (photomietric version) 88
72, Near Encounter Frame 7N18 89
March 1, 1972 Sec. 3, Contents, page ix
Surface JPL 606-1
3. 6 (cont'd)
73(a). Near Encounter Frame 7N19 (Max-D version) 90
73(b), Near Encounter Frame 7N19 (photometric version) 91
74. Near Encounter Frame 7N20 92
Tables
1. Characteristics of the Mariner 6 and 7 camera optics 2
2. Far encounter photoreference data 8
3. Near encounter photoreference data 40
Sec. 3, Contents, page x March 1, 1972
JPL 606-1 Thermal Properties
3. 1 THERMAL PROPERTIES
INTRODUCTION
This section discusses Mars' surface and subsurface temperatures,
their spatial and temporal variations, and the thermal properties derivable
therefrom.* Observed temperatures have been obtained by remote sensing
(radiometry) and are treated under two separate headings: (1) Infrared
Radiometry, which provides surface temperatures exclusively; and (2) Micro-
wave Radiometry, which provides subsurface temperatures. The theoretical
temperatures expected are derivable for Mars from physical theory and avail-
able astronomical and physical data. These theoretical temperatures are
treated first as background data for understanding the observed temperatures
and their implications in terms of the thermal properties of the Martian
surface.
The surface thermal environment of Mars differs from that of Earth in
two ways: (1) the mean surface temperature of Mars is much lower because of
the greater solar distance; and (2) diurnal thermal amplitudes are much larger
because of the thin, dry atmosphere and lack of heat-storing oceans. Mars
resembles Earth in its diurnal and annual thermal rhythms, but Martian
seasonal periods are nearly twice as long and display much more north-south
asymmetry.
3.1.1 THEORETICAL TEMPERATURES
The orbital and mechanical data for the planet Mars are well known (see
Section 1) and the atmosphere is thin enough to be neglected in a first approxi-
mation, so that it is a relatively easy matter to calculate representative simpli-
fied thermal models (of surface and subsurface temperature variation and
distribution), in the same general manner as for the Moon (insolation and heat
conduction of an airless, smooth, homogeneous body). The temperatures
closely follow the insolation in its diurnal and annual rhythms, at any given
location, while the differences between contiguous areas depend upon their
albedo, local slopes, and thermophysical parameters such as infrared
emissivity and thermal inertia. Such models are meaningful only if the values
of these parameters are correctly chosen, a choice v/hich is itself guided by a
comparison of observations with previous modeling results.
If the atmosphere is included, by taking into account some of its effects
on surface temperature, such as infrared back-radiation (greenhouse effect),
condensation/sublimation of volatiles (CO2 and H2O), turbulent convection in
the lowest atmospheric layers, etc. , then more nneaningful models can be
derived, but the calculations become more complicated.
Thus, thermal models with increasing degrees of sophistication can be
constructed for Mars surface environment. Only the highlights of two such
^Temperatures above the surface (atmospheric temperatures) are not treated
here; they are discussed in Sections 5.3 and 5.4 which cover the Lower and
Upper Atmosphere, respectively.
February 15, 1972 C. Michaux, E. Miner, JPL Sec. 3. 1, page 1
Thermal Properties JPL 606-1
models involving the atmosphere are given here; those of Leighton and Murray
(1966) and of Kieffer (Neugebauer, et al. , 1971).
It is necessary, however, to caution the reader that while such thermal
models (which supply both spatial and temporal temperature distributions for
the Martian globe over one Martian year) are most useful in that they supple-
ment the meager observational information obtained so far, their theory and
calculations still rely upon a number of simplifying assumptions. Also, the
choice of values for the thermophysical or other parameters may not be
adequate. Therefore, the models represent at best only rough approximations
of the true distributions. The assumptions made and values of parameters
chosen will be indicated where applicable.
Thermal Models ~ General
All thermal models of the Martian surface have as their basis the plane -
parallel, homogeneous, one -dimensional, partial differential equation for sub-
surface heat conduction: 9T/9x = K/pc • S^T/Sx^, where T is the absolute
temperature, x is the vertical depth measured downward from the surface, K is
the thermal conductivity, p the density, and c the specific heat. The boundary
conditions are that (1) the net thermal energy flux at the surface (x = 0) must be
equal at all times to the difference between the absorbed energy and the
radiated energy, and (2) the temperature is constant at some depth x = £ . The
absorbed energy, neglecting atmospheric effects, is due to solar insolation and
is proportional to (1-A), where A is the Bolometric Bond Albedo (see
Section 3.Z). The radiated energy is ecrT^, where e is the surface emissivity,
cr is the Stefan-Boltzmann constant, and T is the surface absolute temperature.
The solutions obtained from the diurnal surface temperature variation
arc found to depend on the "thermal inertia, "I - (Kpc)-^' ^. A very low thermal
inertia of 0.001 cal cm"'^ sec"-^''^ deg K"l, as may be found on the Moon, is
indicative of dust or powdered rock in a vacuum. Terrestrial rocks have
thermal inertias near 0.05. As thermal inertia increases, the diurnal temper-
ature amplitude decreases, and maximum temperatures occur later in the
afternoon. A more detailed discussion is given by Sinton and Strong (1960a).
Atniospheric Effects
The Martian atmosphere slightly modifies the surface heat balance and
surface temperature expected for an idealized planet without an atmosphere.
Several effects are recognized:
1) Effective Thermal Conductivity. The comparatively extensive
literature on the thermal properties of rock powders in a vacuum
is largely inapplicable to Mars, as the conductivity of the Martian
atmosphere alone, approximately Z. Z X 10-5 (pure CO2 at ZOO°K),
exceeds the conductivity of the lunar surface. In a vacuum,
radiative heat transfer can contribute to the effective conductivity.
The size of the contribution is proportional to T^ (Watson, 1964).
Fountain and West (1970) have shown that under simulated Martian
conditions this radiative effect is small in comparison with pure
conduction. Therefore, temperature -independent conductivity is
appropriate for the Martian conditions.
Sec. 3. 1, page Z C. Michaux, E. Miner, JPL February 15, 197Z
JPL 606-1 Thermal Properties
Z) Infrared Back-Radiation (CO2 greenhouse effect) . While a
predominantly CO2 atmosphere is practically transparent to
incoming solar radiation, the outgoing thermal radiation undergoes
selective absorption. For blackbody radiation in the 200-300° K
range, the 1 5 |jl band of CO^ is the prime contributor to this
Martian atmospheric thermal opacity. Some of this energy is
reradiated back to the surface, producing a weak greenhouse effect.
3) Turbulent Heat Conduction . Turbulent heat transfer across the
boundary layer has been considered by Gierasch and Goody (1968).
The primary effect is to decrease the surface -temperature diurnal
variation by a few degrees, corresponding to an apparent increase
of thermal inertia.
4) Condensation/Sublimation of C02- The condensation temperature
of CO2 is dependent on its pressure, but at Martian surface
pressures (see Section 5. 2) it falls in the range of 145-150°K.
When surface temperatures on Mars reach these values, CO2
condenses, releasing its latent heat and inhibiting any further
temperature drop. Similarly, deposits of CO2 ice inhibit tempera-
ture rises until sublimation is complete. The rate of mass forma-
tion (and the subsequent sublimation) of CO2 ice per -unit-area at
the surface of Mars has been calculated by Leighton and Murray
(1966).
5) Screening and Blanketing by Clouds or Aerosol (Dust) Layers .
While these effects have not been studied in detail for the Mars
environment, they should generally decrease the surface tempera-
ture in daytime (by absorption of solar radiation) and increase it at
night (by back-radiation). The effects are usually local in extent
and certainly highly variable. Gierasch (1971) has made prelimi-
nary calculations on the formation of CO2 clouds in the Martian
atmosphere and their radiative effects. Aleshin and Fedoseeva
(I9VO) calculated that a dust-laden CO2 atmosphere model at 5-mb
pressure v/ould cause the predawn temperatures to increase by as
much as 30° K.
Leighton-Murray Thermal Model
Leighton and Murray (I966) were the first to construct a thermal model
of the Martian surface including the effects of CO2 condensation and sublima-
tion. They incorrectly assumed that the CO2 greenhouse effect could be
accounted for by reducing the apparent infrared surface emissivity by
10 percent. An error was also made in the placement of the aphelion, relative
to northern summer solstice. In spite of these errors, the model provides a
fair idea of the gross thermal conditions at the surface of Mars and in its sub-
surface. Figure 1 depicts their results in terms of diurnal surface temperature
curves for the equator near the autumnal equinox of Mars. The variation of
the average temperature of the disk, as viewed from the Sun, is shown as a
function of season and depth in Fig. 2. The results of calculations of mean
annual temperatures as a function of latitude are shown in Fig. 3. Condensa-
tion temperatures of water vapor are also indicated in Fig. 3.
February 15, 1972 C. Michaux, E. Miner, JPL Sec. 3. 1, page 3
Thermal Properties
JPL 606-1
JOO
■ IS
^
l»0
- f \
UJ
Q.
Z
Ui
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UI
o
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5 200
-
V
Ill \\^-~^°°"
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■~~.^,^/| \,,^^0 006
^^ — .«^__^^ 1 colcm-2s«c-"2dejK-'
150
1 1 1 1 .
12
TIME (hours)
Fig. 1. Theoretical curves of diurnal variation of surface temperature near the
equator at southern spring equinox (approximate seasonal date). Sinton and
Strong (1960a) observations are compared. (Leighton and Murray, 1966)
3 C 9 12
MARTIAN MONTHS FROM N, SUMMER SOLSTICE
Fig. 2. Theoretical curve of annual variation of average temperature of the
disk as viewed from the Sun for various depths in centimeters (using the
CO^ condensation model). (Leighton and Murray, 1966)
Sec. 3. 1, page 4
C. Michaux, E. Miner, JPL
February 15, 1972
JPL 606-1
Thermal Properties
o
UJ
cr
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cr
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200
HjO
100
VAPOR
PRECI-
PITABLE
Condensation temperatures of water
vapor are indicated for 3 abundances
30 60
LATITUDE (deg)
90
Fig. 3. Average annual surface temperature as a function of latitude (using
the CO2 condensation model). (Leighton and Murray, 1966)
The best fit with the Sinton and Strong (I960) radiometric data was
obtained with the following parametric values:
,1/2
Thermal Inertia, (Kpc)'
Bolometric Bond Albedo, A
Surface Emissivity (8-12 p.), sjq
Kieffer Thermal Model
= 0.009 cal cm-2 sec'l/^ deg K"!
- 0.15 for Martian dark areas
= 0.65 for polar ice caps
= 0.85 for bare ground and polar ice
caps
H. Kieffer has constructed thernnal models similar in nature to the
Leighton-Murray model for the purpose of interpreting Mariner 1969 IR
Radiometry data (see Neugebauer, et al. , 1971). These homogeneous ground
conduction models assume (as in the L-M model) temperature -independent
conductivity and specific heat. The orbital geometry error of the L-M model
was corrected, and the greenhouse effect was accounted for by using a constant
atmospheric back-radiation equal to 0.01 of the local noon solar flux. Turbu-
lent heat transfer across the boundary layer was not included in the models, but
was estimated to correspond to a 10 percent increase in the apparent thermal
inertia. CO2 condensation and sublimation was included in the models.
February 15, 1972
C. Michaux, E. Miner, JPL
Sec. 3. 1 , page 5
Thermal Properties JPL 606-1
Figures 4 and 5 show the resulting diurnal surface temperature variation at
perihelion (Fig. 4), and at aphelion (Fig. 5). Perihelion (r = 1.381 AU,
r] = 335°) corresponds closely to northern winter solstice, while aphelion
(r = 1.666 AU, r\= 155°) corresponds to northern summer solstice. Curves
are depicted for latitudes 0°, ±30°, ±45°, and ±60°. Curves for the equinoxes
do not differ much from being intermediate between the extreme positions.
The parametric values adopted were:
Thermal Inertia, (Kpc) = 0.004 to 0.010 cal cm'^ sec "^ /^
deg K-1
Bolometric Bond Albedo, A = 0.20 to 0.40
Surface Emissivity (8-12 |jl), €10 =0.90
3.1.2 INFRARED RADIOMETRY
The relationship between incident solar radiation, reflected solar
radiation, and emitted thermal radiation from Mars is discussed in
Section 3. 2. 2. It can be summarized by stating that the am.ount of incident
solar energy absorbed by the surface of Mars, must be balanced by an equal
amount of thermal energy radiated from Mars. It is this emitted thermal
energy that can be measured using infrared and microwave radiometry. For
temperatures between 150° and 300°K, most of the energy is emitted in the
5 to 30 |j. wavelength range. Infrared radiometric measurements of the
Martian surface are made at these wavelengths. If Mars were an ideal
blackbody radiator, the radiometrically observed energies could be directly
interpreted as surface temperatures by means of Planck's blackbody radiation
equation. However, Mars surface does not radiate precisely like a blackbody,
because of selective absorption of the thermal radiation by the Martian surface
and atmosphere. This results in a slight redistribution of the energy. At a
given wavelength, the observed brightness temperature can be converted to an
actual surface temperature, only if the effective emissivity is known or
assumed.
The atmosphere of Mars is generally transparent to visible and infrared
light, except for absorption at selected wavelengths due to carbon dioxide and
small amounts of water and other trace constituents (see Section 5). Martian
polar regions are obscured in wintertime by the polar hoods, and other regions
of the planet are occasionally obscured by dust or clouds. Observations of the
Martian surface from Earth are hindered primarily by absorption in the Earth's
ovm atmosphere (see Fig. 6). Earth-based infrared radiometric observations
can be made only through the atmospheric windows, such as the 8-14 micron
window. Other disadvantages in observing Mars from Earth are those
encountered at all wavelengths: (1) the planet is always at great distance, with
the result that very small thermal fluxes are received and only low areal
resolution is obtainable over the disk; (2) the relative orbital geometry of
Mars and Earth restricts the phase angle range observable to ±47°, thereby
disclosing little of the nightside; (3) the inclination of the Mars equatorial plane
with respect to the ecliptic can be great enough to conceal a large sector of
latitudes centered on the more distant pole (at the time of the observation).
Sec. 3. 1, page 6 C Michaux, E. Miner, JPL February 15, 1972
JPL 606-1
Thermal Properties
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February 15, 1972
C. Michaux, E. Miner, JPL
Sec. 3. 1, page 7
Thermal Properties
JPL 606-1
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Sec. 3. 1, page
C. Michaux, E. Miner, JPL
February 15, 1972
JPL 606-1
Thermal Properties
UJt/>
100
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WAVELENGTH (MICRONS)
10 50
100
500
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1,000 2.000
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T
Fig. 6. Earth; atmospheric transmission (Murray and Westphal, 1965).
In the particular case of Mars, the concealed polar sector may extend down to
the 74° latitude circle (i.e., 26° from the pole itself). Spacecraft astronomy
meets with fewer hindrances; complete global coverage is possible over
dayside or nightside (subject only to instrument pointing constraints and time
available for observations), and the areal resolution and accuracy attainable
are far superior.
To date, only two sets of modern reliable, infrared radiometric
information on Mars exist. * Together, these measurements cover only six
days of the southern spring season. The main difference between them lies in
their areal resolution and surface coverage, as follows:
1) Low-resolution (500 km) but extensive (within ±70° latitude)
measurements from Earth's largest telescope (200-inch aperture)
during four days (end of July 1954) by Sinton and Strong (1960a, b).
Only dayside measurements from 7 a.m. to 3 p.m. Martian local
time.
2) High-resolution (50 km) but limited (2 percent of Mars surface)
measurements from Mariner 6 and 7 spacecraft during two days
(July 31 and August 5, 1969). Tracks crossed the equatorial and
southern Martian regions, and included the South Polar Cap and
early nightside sampling.
^Earlier radiometric information, obtained at Lowell Observatory by Coblentz
and Lampland between 1922 and 1942 (see Coblentz and Lampland, 1927),
and reduced by Gifford (1956), is extensive, covering three Martian seasons.
However, the temperature measurements are of limited accuracy. These
measurements are not considered here, nor are those of Pettit and Nicholson
(1924) at Mt. Wilson Observatory, belonging rather to the category of
'historical observations'. Their value lies more in the preliminary observa-
tional insight on Martian climatology which they provided at the time.
February 15, 1972
C. Michaux, E. Miner, JPL
Sec. 3. 1, page 9
Thermal Properties JPL 606-1
Infrared Radiometry From Earth
The history of ground-based infrared radiometry of planets is given by
Sinton and Strong (1960b), along with their data for Venus and Mars (July 1954)
and description of their observing equipment. Examples of their scan data are
given in Fig. 7, Of the 33 scans obtained, only 6 equatorial scans were reduced
by Sinton and Strong (1960a). The entire set was analyzed by Morrison (1968).
He reduced the same equatorial scans as Sinton and Strong and obtained kinetic
temperatures as depicted in Fig. 8. The Tj values (Fig. 8) are derived assum-
ing an isotropic emissivity of 0.95, while T2 values are obtained with a lunar
variation of emissivity with direction (Sinton, 196Z), normalized to a mean of
0.95. The model curves shown are for a homogeneous planet with no atmos-
phere. Morrison grouped all of the data into regions 10° X 10° in latitude and
local time, and constructed the brightness temperature contours shown in Fig, 9.
The season on Mars was early southern spring (heliocentric longitude T) ~290°),
Morrison also grouped the data according to bright and dark areas.
Adopting a bolometric albedo A = 0.25 for the bright areas, and an emissivity, e ,
of 0.93, he concluded that the thermal inertia was 0.004 cal cm"^ sec"-^'^
deg K-1. A similar calculation for the dark areas (using A = 0.15 and e = 0.93)
led to a thermal inertia of 0.006. Using these values, Morrison proceeded to
compute the expected temperature distribution with latitude and local time for
the two types of areas when Mars is at the northern autumnal equinox ( V - 265°).
The results are depicted in Fig. 10. He concluded that dark areas are always
hotter than bright areas, although during midmorning this difference is very
small. The greatest temperature differences, about 15° K, develop near
sunset and persist throughout the night.
Grouping the north- south scan data for bright areas into three intervals
of local time, Morrison plotted the average brightness temperature versus
latitude for tj ~290°. These are shown in Fig. 11, where for comparison the
curve of derived peak kinetic temperatures for bright areas has been added.
The fit is fair for northern latitudes, but the data drops much more rapidly in
the southern hemisphere than the theoretical curve. The difference is probably
due to seasonal effects, including the presence of a large southern polar cap.
These were not included in Morrison's nnodel.
To summarize the values of the assumed or derived thermal parameters
obtained by Morrison from Sinton and Strong's data:
Bolometric Albedo, A = 0.25 for bright areas
= 0.15 for dark areas
Surface Emissivity (8-13 |Ji), €10 = 0-93
Thermal Inertia, (Kpc)!/^ = 0.004 cal cm' sec" deg K
for bright areas
= 0. 006 for dark areas
Specific Heat Capacity, c = cal g deg K
Sec. 3.1, page 10 C. Michaux, E. Miner, JPL February 15, 1972
JPL 606-1
Thermal Properties
#10
#13
#11
4 37UT
CM 28*
4.'26UT
CM 22"
4=33 UT
CM24*
NOTE: Scan #2, which crossed a large yellow cloud, is relatively flat and
cold. Circle represents scanning aperture (1.5 arc sec). Abbreviation
CM. means central meridian and UT universal time.
Fig. 7. Radiometric drift-curves across Mars on July 20, 1954 ( tj = 287. 5°).
Numbers on the scans match w^ith those on chart and give
positions at which photographs were taken.
(Sinton and Strong, 1960a).
February 15, 1972
C. Michaux, E. Miner, JPL
Sec. 3. 1, page 1 1
Thermal Properties
JPL 606-1
(Values of fhermal
inerMa indicated)
-80 -60 -40 -20 20
SOLAR HOUR ANGLE (dsg)
40
Fig
8. Equatorial diurnal temperature variation
m Mars in early southern spring (t]~Z90°).
(Morrison, 1968)
10 12
LOCAL TIME (hours)
Fig. 9. Brightness temperature map of Mars in early southern spring (t)~290°),
All data from bright areas are included. (The dashed isotherms are known
with less accuracy than the solid lines. Contour interval is ICK.)
(Morrison, 1968)
Sec. 3. 1, page 12
C. Michaux, E. Miner, JPL
February 15, 1972
JPL 606-1
Thermal Properties
LOCAL TIME - BRIGHT AREA
10
12 14 16
LOCAL TIME - DARK AREA
18
20
22
Fig, 10. Theoretical temperature distribution with latitude and local time
for Mars at the northern autumn or southern spring equinox ( T] = 265°).
The upper map was computed for bright areas: albedo, 0.25 and
thermal inertia, 0.004 cal cm"^ sec"!'^ deg"-'-. The lower map
represents dark areas: albedo, 0.15 and thermal inertia,
0.006 cal cm~2 sec"l/2 deg"l.
(Morrison, 1968)
Density, p
Thermal Conductivity, K
Infrared Radiometry From Spacecraft
= 2 g cm
= 5 X 10-5 ^^i cm-1 sec"^ deg K" 1
for bright areas
= 12 X 10-5 fQj- dark areas
Mariners 6 and 7 each carried a two-channel infrared radiometer
designed to measure the thermal emission of Mars surface (and thereby the
equivalent blackbody temperature) at an areal resolution of about 50 km. The
two spectral channels (8.1 to 12.5 |jl and 17.9 to 25.1 \i) were chosen to avoid
interference from atmospheric CO2 emission and to correspond approximately
to the peaks of the blackbody curves for 300 °K and 150°K, respectively. A
description of the instrument has been given by Chase (1969).
The final results of the IRR experiment have only recently been
published (Neugebauer et al, , 1971). Measurements made during the approach
phase of the flyby did not exceed Earth-based data in resolution and suffer
from insufficient data regarding the precise pointing angle of the instrument.
The measurements were used primarily as a check on the prelaunch instrument
February 15, 1972
C, Michaux, E, Miner, JPL
Sec. 3. 1 , page 1 3
Thermal Properties
JPL 606-1
300
280
260
5f 240
220
2 00
180 —
J_
O 10.7<LT<11.7
• 11.7<LT<13.0
A 13.0 < LT < 14.0
I
±
-60 -50 -40 -30 -20 -10 10
LATITUDE
20 30 40 50 60
Fig. 11. Variation of temperature with latitude on Mars in early southern
spring (t]~290''). The data points are average brightness temperatures
with indicated standard deviations in the mean. The solid curve is the
theoretical peak thermometric temperature for an albedo of 0.25 and a
thermal inertia of 0. 004 cal cm"'^ sec
1/2 d
eg"l. This curve should
be above the data points by less than 10°K. (Morrison, 1968)
radiometric calibration. Results of the near-encounter phase are given in
terms of plots of kinetic temperatures versus local Mars time, as shown in
Figs. 12 and 13. The actual scan traces, across Mars, are shown in Figs. 14,
15, and 16, where the 10 fi kinetic temperatures and the difference between 10
and 20 [i temperatures have also been plotted as reference data. The Meridiani
Sinus region was the only area viewed (during near -encounter phases) by both
spacecraft. Mariner 6 prinnarily scanned the equatorial region of Mars while
Mariner 7 scanned from Meridiani Sinus to the South Polar Cap and then north-
ward through Hellas. Only about two percent of Mars was covered by these
scans. The derived tem.peratures were estimated to have a relative accuracy
of 0.5 °K and an absolute accuracy of 2''K, except at the South Polar Cap, where
depressed temperatures and field-of-view corrections reduced the absolute
accuracy to about 5''K.
In order to reduce brightness temperatures to kinetic temperatures, a
value for the surface emissivity had to be assumed. Based on the best 8-12 [i
laboratory measurements of the emissivity of terrestrial materials (Hovis and
Callahan, 1966) a mean Martian emissivity of e^g - 0-90 was adopted. By
minimizing the differences AT = T20 - Tjq over all the data, a long wavelength
channel emissivity ^20 = 0-88 ±0.03 was derived, along with a = 0.10 ±0.07 for
the exponent a characterizing the cos"' 6 angular dependence law. The only
Sec. 3. 1, page 14
C. Michaux, E. Miner, JPL
February 15, 1972
JPL 606-1
Thermal Properties
<
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February 15, 1972
C, Michaux, E. Miner, JPL
Sec. 3. 1 , page 1 5
Thermal Properties
JPL 606-1
— rx
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C. Michaux, E. Miner, JPL
February 15, 1972
JPL 606-1
Thermal Properties
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C. Michaux, E. Miner, JPL
Sec. 3. 1, page 17
Thermal Properties
JPL 606-1
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Sec. 3. 1, page II
C. Michaux, E. Miner, JPL
February 15, 1972
JPL 606-1
Thermal Properties
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February 15, 1972
C. Michaux, E, Miner, JPL
Sec. 3. 1 , page 1 9
Thermal Properties
JPL 606-1
instance where a could be directly determined from the observations was for a
single area near Margaritifer Sinus, which was viewed at angles of 30° and 56°
from, the vertical. The value obtained was a = 0.17 ±0.05, which is in fair
agreement with the previously derived value.
Comparison of the observational data with Kieffer's model (Section 3.1.1)
yielded a mean surface thermal inertia of 0.006 cal cm-2 sec"-^''^ deg K-1, with
bolometric Bond albedo A = 0.2 to 0.4. These values, which apply to the daytime
measurements lead to temperatures for the nightside which are several degrees
lower than observations. The discrepancy may be due to actual horizontal
variations of the thermal inertia or alternatively, to vertical inhomogeneity.
A further complexity is added to the data analysis by the local topography
of Mars. Although Mars surface temperature is unlikely to be strongly
dependent on altitude (the tenuous Martian atmosphere is probably thermally
decoupled from the surface in a first approximation), any slope which deviates
from that of a perfect sphere causes the solar insolation, at that point, to differ
from the insolation assumed for the models. In the late afternoon or early
evening, such slopes could cause temperature differences of more than ±1°K
for each degree of slope. These differences cannot be modeled until a detailed
topographic map of Mars is available.
The IRR investigators concluded that the surface of Mars appears to be
strongly nonhomogeneous in its thermal properties, on scales ranging from
those of the classical light and dark areas to the limit of resolution of the
radiometers. The derived or assumed thermal parameters are summarized
below:
Bolometric Albedo, A
Surface Emissivity, e^g
«20
1/2
Thermal Inertia, (Kpc)
Specific Heat Capacity, c
Density, p
Thermal Conductivity, K
= 0.2 to 0.4
- ^. 0.10 ±0.07.
= 0.90 cos y
0.10 ±0.07
0.88 cos
6
-2 -1/2
= 0.006 to 0.010 cal cm sec
deg K-1
= 0.16 cal g' deg K'
-3
= 1.3 to 2.0 g cm
-4 -1 -1
= 1 to 3 X 10 cal cm sec
deg K-1
For the probable materials on Mars' surface, thermal conductivity
depends primarily on particle size, only slightly on porosity, and is not
dependent on composition. Using the Wechsler and Glaser (1965) conductivity-
particle size relation for silicate powders under 6 mb pressure, one may
derive (for the K values above) mean particle sizes (diameters) of about 200 \i
for I = 0.006 and 1 mm for I = 0.010.
The low thermal emission of the polar cap could be measured with
reasonable signal-to-noise ratios only in the 18 to 25 [i range. The continuous
Sec. 3. 1 , page 20
C. Michaux, E. Miner, JPL
February 15, 1972
JPL 606-1 Thermal Properties
frost cover started south at about -62° latitude, and the viewing geometry
became extreme by the end of the polar swath. The temperatures (T20)
measured over the cap proper have a greater uncertainty than those measured
elsewhere (on bare or unfrosted ground). This was attributed to two main
causes: (1) effects of the extended field of view, and (Z) thermal offsets of the
internal calibration plate (arising probably when viewing "empty space"). The
first effect was estimated at no more than 2°K error (including noise and gain
errors, etc. ), by making use of a model of the cap at 148°K, the uncorrected
observed minimum temperature. The thermal offset effect was determined to
cause a 5°K uncertainty at the most. It was concluded that the temperature of
the South Polar Cap was 148 + ^°K. -I'
3. 1. 3 MICROWAVE RADIOMETRY
Planck's blackbody radiation equation shows that for temperatures of
150° to 300°K (the range of Martian surface temperatures) most of the energy
emitted is in the 5 to 30 (x wavelength region. At the much longer radio wave-
lengths, where the Rayleigh- Jeans approximation to Planck's law is valid, the
emitted thermal energy is proportional to T/\^, where T is the absolute
temperature of the source, and \ is the wavelength. Although the available
energy-per-unit wavelength is several orders of magnitude smaller in the
microwave region than in the infrared region of the spectrum, measurements
of the microwave spectrum at reasonably high signal-to-noise ratios are
possible using large antennas and more sensitive signal detectors. However,
since high angular resolution is difficult to obtain in the microwave region
only integrated-disk measurements of thermal emission from Mars have been
made. (To date, none of the spacecraft sent to Mars have carried microwave
radiometers. ) The strength of any microwave signal received is directly
proportional to the solid angle subtended at Earth by the disk of Mars. For
this reason, precision measurements have been limited to periods of near-
opposition. The diurnal variation of microwave temperatures can only be
obtained by measuring the disk at different phases due to the low spatial
resolution.
As in the case of infrared radiometry, only the brightness temperatures
can be directly measured. The actual kinetic temperatures can be deduced only
if the appropriate emissivity is known. However, if the diffuse radar
reflectivity of Mars is known at the same microwave wavelength, Kirchhoff's
Law states that the emissivity will be equal to the absorptivity (and
absorptivity = 1 - reflectivity).
Temperatures observed in the infrared refer in particular to the upper
few millimeters of the surface of Mars. Microwave measurements refer to the
subsurface, at depths approximately equivalent to ten times the wavelength,
assuming there are no sharp boundary layers in the subsurface. The discussion
of thermal models in Section 3. 1. 1 makes it apparent that diurnal thermal
*A mean infrared emissivity ^20 of 0-90 in the 18 to 25 jj. range was also
assumed for the polar cap. This admittedly arbitrary value was selected
because of a lack of laboratory measurements on CO2 frosts in the mid-
infrared region.
February 15, 1972 C, Michaux, E. Miner, JPL Sec. 3, 1, page 21
Thermal Properties JPL 606-1
amplitude must decrease with increasing depth, and thus with increasing
wavelength until the depth, H, is reached where the amplitude has decreased
to zero.
For further details on the theory of radio emission from a planetary-
surface, see Piddington and Minnett (1949)- They assumed an airless planet
and a vertically homogeneous subsurface, with thermial and electrical properties
independent of temperature.
Observations
A critical evaluation of all available radio observations of Mars has
recently been made by Epstein (1971), taking into account calibration
uncertainties and sensitivity of equipment used. (In some cases he recalculated
the uncertainties.) By utilizing only the most accurate and reliable observations
to date (2Z of the 32 data sets examined), Epstein arrived at the plot presented
in Fig. 17 of the normalized temperatures T^ versus wavelength \ .* The error
bars plotted are the total uncertainty (1-a) including the absolute calibration
uncertainties.
The disk-averaged brightness temperatures used for the plot are listed
in Table 1, both as observed (Tb) and normalized (Tb = Tb • C) to the mean
solar distance Tq = 1.524 AU. The normalization factor C (discussed below)
was taken as follows:
1/2
For observations at X < 1 cm, C = (r/r )
o
1/4
For observations at 1 cm < \ < 10 cm, C =(r/r^)
For observations at \ > 10 cm, C =1
The wavelength dependence of this normalization factor, C, was proposed
by Morrison, Sagan and Pollack (1969) and discussed by Epstein et al. (1970).
The argument is as follows: As depth beneath the surface increases, the sub-
surface layers are found to be affected less and less by the insolation variations
(diurnal and even annual), so that at a depth corresponding to the origin of very
long centimeter waves, the temperature remains essentially constant. There-
fore, at such depths and corresponding wavelengths, one may take C = 1. At
intermediate depths and shorter centimeter wavelengths, the temperature may
be "scaled approximately as the average of dayside and nightside heliocentric
distance corrections, " that is between C = (r/ro)y 2 and C = 1 respectively. A
good approximation to this average is C := (r/ro)^'^. Of course, as Epstein
(1971) noted, the wavelength intervals and correction forms are somewhat
arbitrary; but the corrections are small for Mars eccentricity (less than ±5 /o
for C = (t/to)^'^).
*The full details of the evaluation as well as the rejected observations may be
found in Epstein's paper.
Sec. 3. 1, page 22 C. Michaux, E. Miner, JPL February 15, 1972
JPL 606-1
Thermal Properties
300
1 10
WAVELENGTH (cm)
100
Fig. 17.
The radiowave spectrum of Mars, as plotted by Epstein (1971)
after initial evaluation of all available observations.
The present picture of the Martian microwave spectrum cannot be
considered to be a very accurate one. As usual in radioastronomy measure-
ments (see the lunar case, for example), the measurements of the continuous
spectrum are plagued with large error bars associated with individual data
points. The sources of these errors are well known: calibration procedures
(5-15%), antenna pointing (1-3%), atmospheric attenuation (H2O) correction
(1-3%), etc. ; but the determination of the respective amounts of those errors,
or uncertainties, is never an easy matter. In the case of the Martian spectrum,
calibration has been the dominant problem. At long wavelengths, the Martian
flux is weak and difficult to measure (low signal-to-noise ratio); however, good
calibration is available from the standard radio sources (usually radio
galaxies) with strong, well-measured fluxes. Here the calibration uncertainty
may be -4-5%. At short wavelengths (< 2 cm), these radio sources are very
weak and have not been accurately measured; therefore, one resorts to the
Moon, Sun, or planets for calibration. However, their fluxes are not always
stable and well defined, and the calibration uncertainty is large (10-15%). The
flux equivalences of all these comparison sources have not always been reliably
established. As a result, the Martian spectrum suffers to some unknown
extent from internal inconsistency due to the variety of calibration procedures
over the 0.1 to 20-cm wavelength range.
February 15, 1972
C. Michaux, E. Miner, JPL
Sec. 3. 1 , page 23
Thermal Properties
JPL 606- 1
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C. Michaux, E. Miner, JPL
February 15, 1972
JPL 606-1 Thermal Properties
No statistically significant variations of brightness temperature with
Martian phase angle have been detected, even at millimeter ■wavelengths. At
the depths in the Martian subsurface where the radio emission originates, it
appears the temperature remains essentially constant (see Dent et al. , 1965).
Despite the large scatter in the data, the general trend of T)-, (see
Table 1) appears to be flat over the entire range of wavelengths used. Due to
the uncertainty in the absolute calibration of the measurements it is not possible
to discern whether or not slight spectral slopes exist, as various theoretical
models predict. The mean surface temperature in the microwave is about
190 ±10°K. This corresponds (assuming an emissivity « = O.9O) to a kinetic
temperature of about 210°K, which is in fair agreement with the mean disk
temperature obtained from infrared observations. Some investigators have
interpreted the lower temperatures at 1 mm as a real discrepancy, and have
attempted to formulate models of the Martian subsurface to account for it (for
example, Sagan and Veverka, 1971).
February 15, 1972 C. Michaux, E. Miner, JPL Sec. 3.1, page 25
Thermal Properties JPL 606-1
BI13LIOGRAPHY
Aleshin, V. I. , and Fedoseeva, T. N. , 1969, The diurnal temperature variation in
the aerosol-gaseous atmosphere and surface layer of Mars: Astron. Zh.
V.46, no. 5, p. 1095-1103, September-October ; translation in Soviet
Astronomy-AJ, v. 13, no. 5, p. 858-864 (1970), March -April.
Alsop, L. E. , and Giordmaine, J. A. , 1961, The observation of three centimete r
radiation from astronomical objects with a ruby maser: Columbia
Radiation Laboratory Special Technical Report, June 1.
Chase, S. C. , Jr. , 1969, Infrared radiometer for the 1969 Mariner mission to
Mars: Appl. Opt. , v. 3, no. 8, p. 639-643, March.
Coblentz, W. W, , and Lampland, C, O. , 1927, Further radiometric measure -
ments and tem.perature estimates of the planet Mars, 1926: Scientific
Papers of the National Bureau of Standards, v. 22, no. 553, p. 237-276.
Davies.R. D., and Williams, D., 1966, Observations of the continuum emission
from Venus, Mars, Jupiter, and Saturn at 21.2 cm wavelength: Planet.
Space Sci. , v. 14, no. 1, p. 15-32, January.
Dent, W. A., Klein, M. J. , and Aller, H. D. , 1965, Measurements of Mars at
X3.75 cm from February to June, 1965: Astrophys. J. , v. 142, no. 4,
p. 1685-1688, November 15.
Drake, F.D., 1970, Private communication. See Epstein (1971).
Epstein, E. E. , 1971, Mars: a possible discrepancy between the radio spectrum
and elementary theory: Icarus, v. 14, no. 2, p. 214-221, April.
Epstein, E, E. , Dworetsky, M. M. , Montgomery, J. W. , Fogarty, W. G. , and
Schorn,R.A., 1970, Mars, Jupiter, Saturn, and Uranus: 3. 3-mm bright-
ness temperatures and a search for variations with time or phase angle:
Icarus, v, 13, no. 2, p. 276-281, September.
Fountain, J. A. , and West, E. A. , 1970, Thermal conductivity of particulate
basalt as a function of density in simulated lunar and Martian environ-
ments: J. Geophys. Res. , v. 75, no. 20, p. 4063-4069, July 10.
Gierasch,P. , 1971, Dissipation in atmospheres : The thermal structure of the
Martian lower atmosphere with and without viscous dissipation:
J. Atmosph. Sci. , v. 28, no. 3, p. 315-324, April.
Gierasch, P, and Goody, R. , 1968, A study of the thermal and dynamical
structure of the Martian atmosphere: Planet. Space Sci., v. 16, no. 5,
p. 615-646, May.
Gifford, F. , Jr. , 1956, The surface-temperature climate of Mars: Astrophys. J.
v. 123, no. 1, p. 154-161, January.
Sec. 3. 1, page 26 C. Michaux, E. Miner, JPL February 15, 1972
JPL 606-1 Thermal Properties
Giordmaine, J. A. , Alsop, L. E. , Townes, C. H. , and Mayer, C. H. , 1959,
Observations of Jupiter and Mars at 3 -cm wavelength: Astronom. J. ,
V.64, no. 8, p. 332-333, October.
Hobbs, R. W. , and Knapp,S.L. , 1971, Planetary temperatures at 9.55-m
wavelength: Icarus, v. 14, no. Z, p. Z04-209, April.
Hobbs, R.W., McCullough, T. P. , and Waak, J. A. , 1968, Measurements of Ma rs
at 1.55-cmand 0.95 -cm wavelengths : Icarus, v. 9, no. 2, p. 360-363,
September.
Hovis, W. A. , Jr. , and Callahan, W. R. , 1966, Infrared reflectance spectra of
igneous rocks, tuffs, and red sandstone from 0.5 to 22 \i: J. Opt. Soc. Am. ,
V. 56, no. 5, p. 639-643.
Hughes, M. P., 1966, Planetary observations at a wavelength of 6 cm: Planet.
Space Sci. , v. 14, no. 10, p. 1017-1022, October.
Kellermann, K. I. , 1965, Radio observations of Mars: Nature, v. 206, no. 4988,
p. 1034-1035, June 5.
Kieffer, H. , 1972: (Los Angeles, Calif., University of California) private
communication to C. Michaux, March.
Klein, M. J. , 1971, Mars: measurements of its brightness temperature at 1.85
and 3.75 cm wavelength: Icarus, v. 14, no. 2, p. 210-213, April.
Kostenko, V. I. , Pavlov, A. V. , Sholomitsky, G. B. , Slysh, V. I. , Soglasnova, V. A. ,
and Zabolotny, V. F. , 1970, The brightness temperatures of planets in the
wavelength range centered at 1.4 mm: Paper and preprint presented at
the XIV General Assembly of the International Astronomical Union,
Brighton, England, August 18-27.
Kuzmin, A.D., Losovsky, B. Ya. , and Vetukhnovskaya, Yu. N. , 1971, Measui-e-
ments of Mars radio emission at 8.22 mm and evaluation of thermal and
electrical properties of its surface: Icarus, v. 14, no. 2, p. 192-195,
April.
Leighton, R. B. , and Murray, B. C. , 1966, Behavior of carbon dioxide and other
volatiles on Mars: Science, v. 153, no. 3732, p. 136-144, July 8.
Low, F, J. , and Davidson, A. W. , 1965, Lunar observations at a wavekmgth of
1 millimeter: Astrophys. J. , v. 142, no. 3, p. 1278-1282, October.
Mayer, C. H. , and McCullough, T. P. , 1971, Microwave radiation of Uranus and
Neptune: Icarus, v. 14, no. 2, p. 187-191, April.
Morrison, D., I968, Martian surface temperatures : Smithsonian Astrophysical
Observatory Special Report No. 284 (40 p. ).
Morrison, D., Sagan, C. , and Pollack, J. B. , 1969, Martian temperatures and
thermal properties: Icarus, v. 11, no. 1, p. 36-45, July.
February 15, 1972 C, Michaux, E. Miner, JPL Sec. 3.1, page 27
Thermal Properties JPL 606-1
Muhleman, D. O. , 1971, Lecture at University of California, Los Angeles,
March 2.
Muhleman, D. O. , 1971, Private communication. See Epstein (197 1).
Muhleman, D. O. , and Sato, T. , 1965, Observations of Ma rs at 12. 5-cm
wavelength: Radio Sci. , v. 69D, no. 12, p. 1280, December.
Murray, B.C., and Westphal, J. A. , 1965, Infrared astronomy: Scientific
American, v. 213, no, 2, p. 20-29, August,
Neugebauer, G, , Munch, G. , Kieffer, H. , Chase, S. C. , Jr. , and Miner, E. ,
1971, Mariner 1969 infrared radiometer results: temperatures and
thermal properties of the Martian surface: Astronom. J. , v. 76, no. 8,
p. 719-728, and 747-749, October.
Pauliny-Toth, I. I. K. , and Kellermann, K. I. , 1970, Millimeter -wavelength
measurements of Uranus and Neptune: Astrophys. Lett. , v. 6, p, 185-187.
Pettit, E. , and Nicholson, S. B. , 1924, Measurements of the radiation from the
planet Mars: Popular Astronomy, v. 32, p. 601-608.
Piddington, J. H. , and Minnett, H. C. , 1949, Microwave thermal radiation from
the Moon: Austral. J. Sci. Res . , Series A: Phys.Sci., v. 2, no. 1,
p, 63 -77, March.
Sagan, C, and Veverka, J. , 1971, The microwave spectrum of Mars: an
analysis: Icarus, v. 14, no. 2, p. 222-234, April.
Sinton, W.M, , 1962, Temperatures on the lunar surface, p. 407-428, in
Physics and astronomy of the Moon; Kopal, Z. , Editor ; Academic
Press, New York, (538 p).
Sinton, W. M. , and Strong, J, , 1960a, Radiometric observations of Mars:
Astrophys. J, , v. 131, no. 2, p. 459-469, March.
Sinton, W. M. , and Strong, J. , 1960b, Observations of the infrared emission of
planets and determination of their temperatures: Baltimore, Md. , The
John Hopkins University Laboratory of Astrophysics and Physical
Meteorology, Contract Nonr 248 (01), Progress Report, April 15,
Stankevich, K, S. : Observations of Mars and Venus at 11.1 cm: Austral. J. Phys.
v. 23, p. 111-112, (1970).
Watson, K. , 1964, The thermal conductivity measurements of selected silicate
powders in vacuum from 150° -350°K: California Institute of Technology,
Pasadena, California, Ph.D. Thesis.
Wechsler, A. E. , and Glaser, P. E. , 1965, Pressure effects on postulated lunar
materials: Icarus, v. 4, p. 335-352.
Sec. 3. 1, page 28 C Michaux, E. Miner, JPL February 15, 1972
JPL 606-1 Ultraviolet, Visible, and Infrared Photometric Properties
3.2 ULTRAVIOLET, VISIBLE, AND INFRARED PHOTOMETRIC PROPERTIES
DATA SUMMARY
Mars brightness at zero phase (See pages 6-8.)
U(l. 0) = +0,25 h(l, 0) = -2.38
B(l, 0) = -0.31 g(l, 0) = -2.37
V(l, 0) = -1.61 e(l, 0) - -2.34
Phase Coefficients, mag/deg (See pages 8-9.)
U = 0.019 h = 0,014
B = 0.018 g = 0,013
V = 0.017 e = 0,013
Geometric Albedo (See pages 9-11.)
p, „ = 0.28
■^1.5 [JL
p. _^ = 0.28
^1.75 |JL
p- „ = 0.23
^2,0 \i
P-? jc = 0-29
^^2.25 ^JL
p, ^ = 0.26
2,5 \i
PX2500
0.12
Py
= 0.16
P\2750
0.07
Ph
= 0.33
P\3000
0.06
Pg
= 0.33
Pu =
0.06
Pe
= 0.32
Pb "
0.09
Pi,
.25 fx
= 0,29
Bolometric Bond Albedo (See page 13. )
)ol
Normal Albedo (See pages 14-15.)
A^^, = 0,26
Detailed tables are given in Appendix A. Exact definitions, more detailed
tables, and sources of the quantities above (and others) are given in the main
body of text.
DISCUSSION
3. 2. 1 Photometric Nomenclature and Theory
The photometric properties of any celestial body may be divided conve-
niently into two categdries: integrated photometric properties, which are stud-
ies of the entire body as a unit; and detailed photometric properties, which are
studies of a body on a point-by-point basis. In practice, any remote sensing
technique will integrate over a considerable area, but hopefully not so great an
area that all details are lost. The Moon is a unique body in being near enough
October 1, 1971 R. Newburn, JPL Sec. 3.2, page 1
Ultravioletj Visible, and Infrared Photometric Properties JPL 606-1
to the Earth that detailed photometry was attempted nearly 200 years ago, by
Scliroeter. Mars exhibits a disk two orders of magnitude smaller than tlie Moon
in angular size, even under favorable conditions, and photometric interpreta-
tions are further complicated by an atmosphere whose effects are difficult to
separate from surface properties. Nevertheless, attempts to separate bright
and dark areas of Mars' surface in visual photometric studies were made as
early as 1909, while visual integrated photometry goes back to at least 1864.
Extensive vis\ial and photographic work was done in the period 1920-1939, espe-
cially in Russia and France. Much of this work was important in early attempts
to gain a qualitative understanding of Mars, but very little of it meets modern
quantitati\c requirements. A comprehensive discussion of this period has been
produced l)y de Vaucouleurs (1954). Modern quantitative photometry of Mars is
performed using various types of electronic detectors (photomultipliers, PbS
cells, etc. ), but it still presents certain problems. The general nature of
tliese problems will be indicated in succeeding paragraphs.
Care must be taken to differentiate between photometric and radiometric
data. Photometric data refers to the response of some particular detector
system. The data are convolved with the spectral response of the filters, detec-
tors, etc. By pure definition, "photometric data" are visible data received by
the human eye, and "physical photometric units" assume an international spec-
tral luminous efficiency curve for the eye. In astronomy, the word "photometric"
is used in the broader sense.
In any astronomical photometric systemi, values are usually given in
magnitudes, based on a logarithmic scale o f anc ient origin. One magnitude
difference is exactly the fifth root of 100 ( ylOO) ratio in flux; that is, approxi-
mately 2.512. Five magnitudes is exactly a factor of 100, 10 magnitudes a
factor of 10,000, and etc. The simple formula relating magnitudes m and m}-,
and fluxes f and f^-^ is
log 7^ = -0.4 (m - m, )
The magnitude scale is an inverse system; that is, the brighter the object, the
smaller its numerical magnitude. For example, an object brighter than magni-
tude one may be magnitude zero or have a negative magnitude.
The most common photometric system in use today is the ultraviolet,
blue, and visual system (UBV), sometimes with red and (near) infrared (IR) and
even longer (JKLMN) wavelength measurements added. The exact system is
defined by a set of magnitudes for a group of standard reference stars to which
observations with any local system must be transformed. Each passband is
approximately fixed in any given photometer by a standard detector-filter com-
bination. Colors are given by nnagnitude differences between passbands in
order of increasing wavelength; e.g., U-B, B-V, V-K. The "effective wave-
length" of each passband (the mean wavelength integrated over the passband) is
given in Table 1. Zeros of the system have been chosen so that U-B, B-V, and
all other "colors" are 0.00 for an unreddened star of spectral type AO V, and so
that passband V agrees in zero point with an older "classic" photometric sys-
tem. Detailed response curves are given by Johnson and Mitchell (1962), Low
and Johnson (1964), and (most accurately) for U, B, and V by Azusienis and
Straizys (1966).
Sec. 3.2, page 2 R. Newburn, JPL October 1, 1971
JPL 606-1
Ultraviolet, Visible, and Infrared Photometric Properties
Table 1. Effective wavelengths in the UBV system. (Johnson, 1966)
Passband
Wavelength, \x
U
0.36
B
0.44
V
0.55
R
0.70
I
0.90
J
1.25
K
2.20
L
3.4
M
5.0
N
10.2
The UBV system was designed for stellar work. It has broad passbands,
making it easy to work with faint objects, and so long as the objects in question
have energy distributions similar in form (a blackbody curve), no problems
arise. Even in stellar work, a slight variation in the ultraviolet passband from
photometer to photometer can cause discrepancies (because of departures from
a blackbody at the Balmer discontinuity), as can interstellar reddening. Even
differences in observatory altitudes can make data reductions difficult because
of varying atmospheric opacity in the U-passband. Transformations from a
"local" photometric system to the standard system can be multivalued in the
worst casesj causing as much as 0.1-0.2 magnitude differences in supposedly
"precise", results. By the time Mars* albedo (reflectivity) has been convolved
with the solar spectral irradiance, its spectral radiant exitance looks some-
thing like an early B-star with 1.5 magnitudes of interstellar reddening, and
inevitably there have been systematic differences of as much as 20% in various
UBV measurements of Mars (Young, 1970). A general discussion of broadband
photometric systems has been produced by Johnson (1963).
In an attempt to fill basic gaps in photometric data and to overcome
some of the difficulties of the UBV system, NASA sponsored a program for
multicolor photometry of the planets by Harvard College Observatory personnel
during the period 1962-1965, This work utilized 10 narrowband (interference)
filters, between 3,150 A and 10,600 A, in addition to the UBV system (Young
and Irvine, 1967). McCord and coworkers have used up to 52 interference
filters^and a double-beam photometer to study discrete Martian regions between
3,010 A and 25,230 A (McCord and Westphal, 1971; McCord, 1968). Still
another modern technique utilizes a scanning spectrometer as a monochromator .
Younkin (1966) used such an instrument with a bandpass of 50 A in his Mars work.
October 1, 1971
R. Newburn, JPL
Sec. 3. 2, page 3
Ultraviolet, Visible, and Infrared Photometric Properties JPL 6O6-I
Any photometric system may be converted to a radiometric one by
absolute calibration of the photometric system response curve. Either photo-
metric or radiometric values can also be cited in absolute physical units.
General discussions of the absolute calibration problem have been offered by
Code (1960), Willstrop (I960), and Oke (1965). Numerical values for absolute
calibration of the UBV system in all the passbands U through N have been given
by Johnson (1966), but these values may be in error by as nauch as 10%
(Azusienis and Straizys, 1966; Young, 1969). The best available calibrations
for the U-, B-, and V-passbands have been given by Azusienis and Straizys
(1966). For planetary studies, measurement of the absolute spectral exitance
may not be required. Mars only emits (as opposed to reflects) significantly in
the far infrared. Therefore, various ratios of "flux out" to "flux in" are suffi-
cient to describe the planet's behavior in the shorter wavelengths. Comparison
with the Sun is usually made through one or more internnediary standard stars
which have been intensively studied and which in turn have been (indirectly)
compared to the Sun.
3. 2. 2 Reflection Versus Emission on Mars
The ratio of reflected to emitted flux, as a function of wavelength at any
given location on Mars, is a function of (1) the incident radiant energy (see
Section 6); (2) the atmosphere (see Sections 5. 3, 5. 4, and 6); (3) the local nor-
mal albedo (see paragraphs following); (4) the local geometry and emissivity
(see Section 3. 1); and (5) the local temperature (which, in turn, is determined
by a complex interaction of local transport properties and the past history of
the radiant field).
The spectral irradiance of the Sun at Mars is better known than the
other quantities in question, although there are still uncertainties of at least
1-2% (see Section 6). This model assumes the values given by Thekaekara
(1970), which integrate to a total electromagnetic irradiation of 582.7 W m" at
Mars' mean solar distance.
The Martian normal albedo (precisely defined in later paragraphs) is
poorly determined in the infrared beyond 2.5 \i. However, between 1.0 jjl and
2.5 \x the albedo appears to be about 0.4 for bright areas and 0.15 for dark
areas (McCord, Elias, and Westphal, 1971). Sinton (1967) has presented evi-
dence for a large absorption feature at 3.0 \i, where the albedo drops to
perhaps 0.1 in both regions and then rises back to the previous level at about
4.0 |j.. The work of Beer, Norton, and Martonchik (1971) confirms this general
behavior. There are no published data beyond 4.0 fji.
The brightness temperature of a surface is the temperature that a
blackbody (a body with relative emissivity of unity) would have to possess in
order to emit the same power, at the wavelength of observation, as is actually
measured. Remote observations of Mars in the "radiometric window" (8-14 \x)
and at microwave frequencies are often reported as brightness tennperatures,
since these are directly related to what is actually measured (see Section 3. 1).
Brightness temperature can be converted to (kinetic) temperature, only if the
emissivity is known or assumed.
Sec. 3.2, page 4 R. Newburn, JPL October 1, 1971
JPL 606-1
Ultraviolet, Visible, and Infrared Photometric Properties
Figure 1 is a plot of curves of spectral radiant exitance (power emitted
per-unit-area and wavelength) versus wavelength for blackbodies of 150 °K,
ZOO°K, 250 °K, and 300 °K, and also curves of the mean solar spectral irradi-
ance times the normal albedo (for typical bright and dark areas). Allowing for
the (roughly) cosine effect of nonorthogonal illumination would further reduce
the reflected contribution, of course, making the later set of curves the highest
possible for the reflected contributions. The normal albedo is assumed con-
stant at the 4 |j. value for all longer wavelengths. Somewhere beyond 4 |jl the
albedo presumably drops, so the assumed values may be too high. It is clear
that at wavelengths shorter than about 3.5 |i,, reflected energy dominates emitted
energy. Beyond 9 [x, emitted energy even dominates reflected energy from the
cold, highly reflective, polar caps. The curves do not allow for effects of
local geometry or non-isotropy in reflection or emission.
<
10
(BRIGHTNESS
TEMPERATURE)
300" K
250° K
200° K
150°K
■• 1.0
■" BRIGHT (0.4)
•] DARK (0.15)
(NORMAL
ALBEDO)
SPECTRAL IRRADIANCE x
NORMAL ALBEDO
SPECTRAL RADIANT
EXITANCE
3 4 5 6 7 8
WAVELENGTH - MICRONS
Fig. I. Reflected solar radiation versus emitted radiation from Mars.
October 1, 1971
R. Newburn, JPL
Sec. 3. 2, page 5
ultraviolet, Visible, and Infrared Photometric Properties
JPL 606-1
3. Z. 3 Integrated Photometric Properties
As described previously, Martian photometric properties can be divided
into integrated and detailed. The integrated properties are considered first.
These include the total brightness, the average color, the variation of brightness
and color with phase (the phase function), the Bond (or spherical or Russell-
Bond) albedo, and the geometric albedo. All of these vary somewhat with
Martian rotation, as various features rotate into the observable hemisphere, and
with Martian season, as polar caps grow and shrink and the surface undergoes
the "wave of darkening" (see Section 4. 2). At shorter wavelengths, additional
variable effects (presumably atmospheric) are noted.
Brightness, Opposition Effect, and Color
The body-centered angle between the source of illumiination (the Sun) and
the observer (detector) is called the phase angle o (see Fig. 2). Where there
is need to discriminate, the phase angle is negative before zero phase (opposi-
tion). In fact, because the orbits of Earth and Mars are not coplanar, Mars
could reach true zero phase, as seen from Earth, only if opposition were to
occur when the planet was at an orbital node. The brightness is a nonlinear
function of phase. Near zero phase, the planet appears slightly brighter than
would be predicted by extrapolation to zero phase of a linear fit to data in the
10° to 30° phase interval. This "enhanced" brightness for phase angles |c|<10°
is called the opposition effect. Actually, the phase curve has a slight continu-
ous upward curvature out to 40° phase, and there is argument whether any
effect remains, if a quadratic or cubic fit is used rather than a linear one.
(The Moon, on the other hand, has an unquestionably real opposition effect. )
5 5UBSOLAR POINT
O SUBOBSERVER POINT
P GENERAL SURFACE POINT
( ANGLE OF EMERGENCE
; ANGLE OF INCIDENCE
a PHASE ANGLE
i LUMINANCE LONGITUDE
X LUMINANCE LATITUDE
Fig. 2. Photomietric coordinates.
Sec. 3. 2, page 6
R. Newburn, JPL
October 1, 1971
JPL 606-1 Ultraviolet, Visible, and Infrared Photometric Properties
Ignoring rotational, seasonal, or secular effects, the mean magnitude
of a planet is given by the expression
M = M(l, 0) + 5 log(r d) + AM (a)
where
M(l, 0) = mean nnagnitude at zero phase and at unit distance from
Earth and Sun, whether this configuration is physically-
possible or not. It is a form of absolute magnitude for
solar system objects.
r = distance from Earth, in AU
d = distance from Sun, in AU
AM (o-) = correction for phase angle o-
Because Mars rarely ever approaches true zero phase, it is useful for some
purposes to use M£(l, 0) which is the same as M(l, 0) except it is calculated
from a linear extrapolation to zero phase from phase data taken for a > 10°
and therefore does not include any curvature of the phase function or opposition
effect.
The value of Mg{l, 0) will appear to vary v/ith rotation, season, etc. ,
depending upon the passband. In general, the variation is greater in the red
regions of the spectrum, where more surface detail is apparent, than in the
blue regions. In the U- and B-passbands, the variation may be in the order of
0.10 to 0,15 magnitudes (Young, 1970). Typically, in the V-passband there miay
be 0.15 magnitude variation as a function of central meridian longitude (Young,
1970; de Vaucouleurs, 1970; Irvine et al. , 1968a; Irvine et al, , 1968b). In the
red, the variation increases to about 0,3 magnitudes at X.7300, and perhaps 0.35
magnitudes at )^10600 (Irvine et al, , 1968a; Irvine et al. , 1968b),
Any attempt to quote values of Mg(l, 0), or worse M(l, 0), m>ust neces-
sarily be restricted to mean values of inhomogeneous data until such time as a
program such as Young's, to quantitatively disentangle all of the various effects,
is con^pleted. The review article by de Vaucouleurs (1970) effectively covers
work at three observatories during the Martian oppositions of 1952, 1954, and
1958. The Harvard work at Le Houga (Irvine et al. , 1968a) and Boyden (Irvine
et al., 1968b) covers the period from May 1963 through July 1965, The
Le Houga results and those reported by de Vaucouleurs are in good agreement,
while the Boyden results are consistently fainter in the red regions, and,
according to the authors, reflect the presence of more dark markings, as seen
from the Boyden station (although it is not obvious that this is the case). This
review arbitrarily takes the mean of the extensive Le Houga and Boyden results
for use here as follows:
U£(1,0) = +0.35 yg(l,0) = -1.51 g£(l,0) = -2.27
B^(1,0) = -0.21 h£(l,0) = -2.28 0^(1,0) = -2.24
October 1, 1971 R. Newburn, JPL Sec. 3,2, page 7
Ultraviolet, Visible, and Infrared Photometric Properties JPL 606-1
The h, g, and e passbands of the Harvard photometry have effective wavelengths
of 7,297 A, 8,595 A, and 10,635 A, respectively.
The Harvard photometry did not disclose any opposition effect. The
Le Houga observations included only three sets of observations under 10° phase
angle, but the Boyden observations included nine sets under 5° and 18 sets
under 10°. O'Leary worked in 1967 and 1969 when the phase angle reached the
very favorable small values of 1.2° and 1.3°, respectively (O'Leary, 1967;
O'Leary and Jackel, 1970), and generally obtained results which qualitatively
supported earlier suggestions of an opposition effect by de Vaucouleurs (1959)
and others. The 1969 effect reported was only about half the size of the 1967
effect (and the data appear to be of greater precision). It is suggested here
that at visual wavelengths an effect of 0.1 magnitudes should be "added"; that is,
M(1,0) = M£(1,0) -0.10
Perhaps half of this is "true" opposition effect, and the remainder is correction
for curvature in the phase curve. Martian photonnetry is in a sufficiently
primitive state that further sophistication seems unjustified at the present time.
As previously noted, colors are given by magnitude differences between
two passbands. The colors of Mars compared to those of the Sun (as given by
Irvine et al. , 1968a) are
Color Mars Sun
U-B
0.56
0.14
B-V
1.30
0.65
V-h
0.77
V-g
0.76
V-e
0.73
These figures indicate Mars to be a body much redder than the Sun, which is no
great surprise.
Phase Function
The phase function of a body, commonly written (^(o), describes the flux
reflected as a function of phase angle, normalized to zero (full) phase; i. e. ,
<t)(a) =
F(a)
As always, practical measurements are made in a particular photometric sys-
tem and, when properly calibrated, can sometimes be transformed to other sys-
tems, with only a small loss in accuracy. Plots of <^(<y) for the Moon, Venus,
and Mercury often appear with a a polar coordinate. Mars, however, can never
be observed from Earth at a phase angle greater than 48°, and, except for the
opposition effect, this phase variation can be roughly represented by a linear
expression
AM(o') = \i.a
Sec. 3.2, page 8 R. Newburn, JPL October 1, 1971
JPL 606-1
Ultraviolet, Visible, and Infrared Photometric Properties
where the coefficient \i is often called the phase coefficient. In this form,
AM(q) must be used with the "linear" absolute magnitudes Mjg(l, 0) to derive
Martian brightness. In theory, additional coefficients and negative powers of
a could be included in the expression for AM(o) to represent the opposition
effect, but, as noted in previous paragraphs, there is still too much uncertainty
about the amount or even the existence of the effect. Ultimately, it is hoped
that spacecraft data can be used to extend the phase function to phase angles
greater than 48°, but such data are not yet available. The phase coefficients
for the various passbands being used in this review are taken from Irvine et al.
(1968b), as follows:
Passband
^^
mag/de
■■£
Pai
ssband
KL
mag /deg
U
0.019
h
O.OM
B
0.018
g
0.013
V
0.017
e
0.013
Thus, there is a slight change in color with phase, in that Mars appears slightly
redder with increasing phase. It must be noted that a comprehensive study of
all Martian photometric data by Young* shows a distinct change of B - V with
phase but no evidence of any change in U - B.
Geometric Albedo
The geometric albedo p of a body is the ratio of its mean luminance at
full phase (a = 0) to the luminance of a perfectly diffusing ("intrinsically white")
plane surface at the same point and perpendicular to the source of illumination
(the Sun). A perfect diffuser is one which scatters 100% of the power incident
upon it (absorbing none) according to the Lambert law of cosines. Its luminance
appears the same from any angle, and is proportional to the cosine of the angle
of incidence. The geometric albedo can be calculated from the expression
log p = 0.4 [M - M(l, 0)] -2 log R + 16.350
where
M = apparent magnitude of the Sun at one AU in the
photometric band in use
M(l, 0) - planetary absolute magnitude in the photometric band in
use (as in previous paragraphs)
R = mean radius of body (in kilometers)
Details of the derivation can be found in Chapter VI of Sharonov (1964) or, in
less detail, in Harris (1961).
='=Private comrrainication (work to be puljlisherl).
October 1, 1971
R. Newburn, JPL
Sec. 3. 2, page 9
)wn
Ultraviolet, Visible, and Infrared Photometric Properties JPL 606-1
Geometric albedo is a widely used concept, but it does suffer observa-
tional and theoretical inconveniences. The magnitude of the Sun must be kno\
in the relevant photomietric systemi, as the Sun is so bright relative to other
objects, there are usually certain inaccuracies in measurement. The zero-
phase magnitude of Mars is uncertain because of the opposition effect, as pre-
viously discussed. Real bodies in space, such as Mars, are generally more or
less spherical. Comparing a planet's hemisphere to a perfectly diffusing plane
is rather artificial. Comparison to an intrinsically white body of the same
shape would be more illustrative. A perfect Lambert surface would then have
an albedo of unity, whereas its geometric albedo is only two-thirds. This should
not imply a body can have a geometric albedo no larger than two-thirds. A body
with very strong backscattering characteristics (something approaching a per-
fect retrodirective reflector), can have an indefinitely large geometric albedo.
A V-magnitude for the Sun of -26.8 is adopted here. Various determina-
tions generally report internal errors of perhaps 0.02 magnitude, but differ
from each other by far greater amounts. A realistic total probable error may
be as much as 0.1 magnitude. Solar magnitudes in the other passbands are
obtained by applying the colors previously given. A Martian mean radius of
3383 km is used, and an opposition effect of 0.1 magnitude is arbitrarily applied
to all Mj?(l,0) values. The resulting geometric albedos are
PU = 0-06 p^ = 0.33
'B
0.09 p = 0.33
p,, = 0.16 p = 0,32
'V
e
Observations of ultraviolet flux from Mars have been made by Evans (1965)
and by Broadfoot and Wallace (1970), using sounding rockets. These results are
difficult to calibrate absolutely, as they each refer to one point in time. Evans'
work was done with an objective grating, while Broadfoot and Wallace used a
slit spectrometer that accepted 42% of the disk area. For purposes of this
review, the curve of Evans is "calibrated" by matching it at 3500 A to the
Harvard photometry quoted previously. On this basis, crude geometric albedos
can be assigned at a few additional wavelengths, as follows:
P\3000 " °'°^
PX2750 = 0-0^^
P\2500 " 0-12
Earth and Hord (1971) have reported on spectra taken with the Mariner
ultraviolet spectrometers. They show an ultrdviolet reflectance for a "desert
region, " which is said to be Meridiani Sinus, actually a dark area (the distinc-
tion is not very significant at short wavelengths). In any case, their results
refer to a specific point and are imipossible to transform to geometric albedos
(see succeeding paragraphs on detailed photometry).
Sec. 3.2, page 10 R. Newburn, JPL October 1, 1971
JPL 606-1
Ultraviolet, Visible, and Infrared Photometric Properties
Observations of Mars at wavelengths longer than one micron all are
fairly recent and are primarily detailed photometry, rather than integrated
photometry. McCord and Westphal (1971) have presented integrated disk results
for one Martian longitude, centered on Amazonis, with bright areas dominant.
Their values, given in Table 2, are taken from Figs. 7 and 8 of McCord and
Westphal and refer to a phase angle of 5°. The "corrected values" were de-
rived, by assuming a phase coefficient of 0.013 magnitude per degree and an
opposition effect of 0.1 magnitude, in an attempt to make them true geometric
albedos. The resulting values at 0.85 [jl and 1.05 \i are slightly smaller than
those of the Harvard photomietry, but neither bandwidths nor wavelengths are
identical, and the absolute accuracy is probably no better than 10% in any event.
There is no absolute integrated photometry beyond 2.5 ^j.. The works of
Sinton (1967) and of Beer, Norton, and Martonchik (1971) show evidence of a
large decrease in albedo near 3 \j. and then a rise back to roughly the same level
as at shorter wavelengths, but all of this work is on relative scales, and Sinton's
work does not refer to the entire disk. In addition, at these wavelengths there
are no phase coefficients to be used for correction to zero phase. Therefore,
it can only be stated that the geometric albedo appears to drop under 0.1 near
3 |j. and to rise back over 0.2 from 3.6 fj. to 4,0 |j.. Beer=:' and Norton recently
have acquired data beyond 4.0 fx, but no results are available at this time.
Table 2. Infrared geometric albedos.
Wavelength
(f^)
0.855
1.053
1.25
1.50
1.75
2.00
2.25
2.50
Geometric Albedo
McCord and Westphal
0.27
0.27
0.25
0.24
0.24
0.20
0.25
0.22
"Corrected Values'
0.31
0.31^
0.29
0.28
0.28
0.23
0.29
0.26
=i=Private communication.
October 1, 1971
R. Newburn, JPL
Sec. 3.2, page 1 1
Ultraviolet, Visible, and Infrared Photometric Properties JPL 606-1
Bond Albedo and Phase Integral
The Bond albedo A, often called spherical albedo and sometimes the
Russell-Bond albedo, is a quantity with clear physical meaning. It is the ratio
of the power (flux) reflected in all directions by a body to the power incident
upon it in a collimated beam. It is the fraction of flux incident upon a body that
is not absorbed. The geometric albedo is a measure of the flux returned at
zero phase. There is a quantity q called the phase integral, that is a measure
of scattering at other phases, which relates A and p. It is the ratio of flux
scattered in all directions to that scattered at zero phase, per unit solid angle.
It can be shown (Sharonov, 1964; Harris, 1961) that
A = pq
where
.'0
sin a do
and 4'{o) is the phase function (converted into intensity units rather than the
logarithmic magnitude units) as in previous paragraphs. Phase coefficients
vary slowly with color, and therefore q must be a (weak) function of color.
Worse, Mars cannot be observed froni Earth at o > 48°. It has been customary
to use Russell's Rule, a 55-year-old empirical relationship Russell found to
hold true within ±5% for Mercury, Venus, and the Moon, which states that
q ~ 2.2 (t>(50°)
Veverka (1971) has recently shown that this empirical law is the direct conse-
quence of rapid decrease in brightness with increased phase. Using a two-point
Gaussian quadrature on the phase integral, he showed that
q - 2(1+6) c}5(55°)
_ tt)(125°)
where 6 = — — — and is a snnall quantity (~0.1). Using actual values of & and
9(55 )
c|> for Mercury, Venus, and the Moon, the two formulas, Gaussian and Russell's,
agree quite well.
Irvine et al. (1968b) derived the following values for the phase integral
q, using a version of Russell's Rule, with a constant 2.17 based upon modern
data by Harris (1961):
Passban
id
Ph;
a.se Inte
[g
ral
P,
assba
nd
Ph;
a.se Integral
U
0.92
h
1.12
B
0,94
g
1,17
V
1,01
e
1,20
Sec. i.2, page 12 R. Newburn, JPL October 1, 1971
JPL 606-1 Ultraviolet, Visible, and Infrared Photometric Properties
Corresponding values for the Bond albedo A, using the geometric albedos given
on page 10, are as follows:
Passband
B^
ond Albe
ido
R
as sbi
and
Bon
d Albedo
U
0.05^
h
0.37
B
0.08-^
g
0.39
V
0.16
e
0.38
Bolomietric
Bon
id.
Albe
:do
The quantity of importance in energy balance studies is the bolometric
Bond albedo, the average albedo weighted by the solar flux Fq(\).
t
^bol "T^Fq(M d\
A constant value of q = 1.2 has been used with the infrared geometric albedos
previously quoted, the values above, and the solar flux values of Section 6 to cal-
culate this integral. A numerical integration from 0.3 \i to 2.5 jjl gives a value
of 0.266. The neglected wavelengths less than 0.3 \i. include 1.2% of the solar
flux and are of low albedo. The wavelengths greater than 2.5 \x include 3.7% of
the solar flux and are low at 3.0 \x and high at 4.0 [jl. Assuming zero albedo for
these regions would only reduce the value to 0.253. Therefore,
A, , = 0.26
bol
is adopted as the best current value for the bolometric Bond albedo. This is
slightly larger than the value of Irvine et al. (1968b), because 0.10 magnitude
of opposition effect has been included, the Sun is assumed 0.01 magnitude fainter,
and new infrared data have been used. The potential sources of systematic
error are such that assignment of a "probable error" makes little sense.
3. 2. 4 Detailed Photometric Properties
Detailed photometric properties of Mars include brightness and color of
localized areas as a function of phase, the normal albedo, the photometric
function, the radiance factor, and the radiance coefficient. The disk of Mars
is so small, even near opposition, that excellent seeing is required for rela-
tively crude detailed photometry by Earth-based telescopes. Results to date
have been limited to a few of the larger dark and bright Martian features and
within a limited range of phase angles. Hence, many of these listed detailed
properties are simply unknown. They are described here briefly for complete-
ness. Hopefully, the current program of spaceflights to Mars may provide
additional answers. '
October 1, 1971 R. Newburn, JPL Sec. 3. 2, page 13
Ultraviolet, Visible, and Infrared Photometric Properties JPL 606-1
Radiance Coefficient
The radiance coefficient r is defined as the ratio of the radiance observed
to that of a white Lambert surface (plane) at the same inclination. The radiance
factor p is defined as the ratio of the radiance observed to that of a white
Lambert surface at zero inclination to the incident illumination. Obviously,
p = r cos i
In practical observing, the detector will have an irregular passband and the
quantities measured are luminances, so they are somietimes assigned the more
appropriate adjective, luminance, becoming luminance coefficients and lumi-
nance factors.
Radiance Factor, Photometric Function, Normial Albedo, etc.
The radiance factor of an element P on a sphere can be written
P = pQ f(i.«-0')
where f(i, t , o) is called the photometric function and is normalized, so that
f(0, 0, 0) = 1, by Pq which is called the normal albedo. The quantities i, £,
and a are the angle of incidence, angle of emergence, and phase angle as pre-
viously shown in Fig. 2. The normal albedo then is the exact equivalent in
detailed photometry to geometric albedo in integrated photomietry; that is, the
luminance at zero phase, comipared with an intrinsically white plane La.mbert
surface perpendicular to the source of illumination. It differs only in referring
to a point at the center of the disk, rather than to the mean of the entire hemi-
sphere. There is an unfortunate tendency on the part of many recent authors
of papers on detailed photometry to use the term geometric albedo for any
reflectivity measured at zero phase or reduced to zero phase. This is NOT
correct for any body which exhibits detectable limb darkening; i. e. , Mars.
There are additional ternns in common use in detailed photometry. The
phase plane (occasionally phase-angle plane) is the Sun-object-observer plane
(sec Figs. 2 and 3), and the luminance equator (or radiance equator) is the
intersection of the phase plane with the surface under study. The luminance
longitude is the angle of observation (reflection angle) projected into the phase
plane. Usuallv, the angle is taken to be negative if it is on the subsolar point
side of the subobserver point.
When observing Mars fronn the Earth, it is perfectly acceptable to con-
sider the phase plane as passing through the center of Mars for all observations
(plane SCO in Fig. 2), since the error involved is, at most, a few arc seconds.
The luminance equator then becomes the great circle where the phase plane
intersects Mars. The relationship between luminance latitude X and longitude S.
and the angles previously used (see Fig. 2) is (Harris, 1961)
cos « = cos \ cos 2
cos i = cos Xcos a - a)
Sec. 3. 2, page 14 R. Newburn, JPL October 1, 1971
JPL 606-1
Ultraviolet, Visible, and Infrared Photometric Properties
However, when Mars is observed from a flyby, an orbiter, or a spacecraft on
the Martian surface, it is absolutely necessary to use the proper (local) defin-
itions as shown in Fig. 3.
PHASE
PLANE
SOLAR VECTOR
J> OBSERVATION VECTOR
P OBSERVED POINT
N NORMAL TO SURFACE AT P
X LUMINANCE LATITUDE
i ANGLE OF INCIDENCE
e ANGLE OF OBSERVATION (EMERGENCE)
a PHASE ANGLE
i LUMINANCE LONGITUDE (SHOWN
NEGATIVE IN CASE DRAWN)
Fig, 3. Local photometric geometry, (after Holt and Rennilson, 1968)
Empirical Photometric Behavior
Most detailed Martian photometry has been indirectly absolute, in the
sense that various regions on the disk have been ratioed to one area near the
disk center, and that area alone has been compared to the Sun through inter-
mediary standard stars. Modern detectors, amplifiers, and recorders are
sufficiently linear and stable in operation, such that system drift is no longer a
severe problem. The extinction caused by the Earth's atmosphere does vary
strongly with zenith distance, with wavelength, and occasionally in time, how-
ever, and continuous reference to standards is necessary to remove these effects
in order to obtain absolute results.
The amount of data required for full evaluation of the photometric func-
tion for each physiographic unit is quite large and has never been totally obtained.
It is not possible to obtain all of the required data from observations on Earth,
It has sometimes been assumed that Mars is a Lamibert surface (with p = p^
cos i), but this is really unacceptable, A Lambert surface has q = 1,5, and the
phase integral for Mars is clearly much smaller (as stated previously). Young
(1969) suggested the use of Minnaert's empirical law for the photometric func-
tion. It is a relatively simple law and satisfies the reciprocity principle (stating
that light rays should be reversible "through the system" unless there is a
October 1, 1971
R. Newburn, JPL
Sec, 3. 2, page 15
Ultraviolet, Visible, and Infrared Photometric Properties JPL 606-1
significant amount of polarization). Minnaert's law can provide an excellent fit
to lunar data (Harris, 1961), even though the Moon does exhibit some polariza-
tion. The law has the form
,-/ \ r • 1 K(o)
p cos (. = C(o) [cos 1 cos £ J
where C and K are functions only of phase angle a for a given passband, and the
other quantities are as given previously. For K equal to unity and constant C,
this reduces to the Lambert law. Qualitatively, the K is a measure of limb-
darkening, and the C is a measure of reflectivity. When a = i = e = o, then
C reduces to the normal albedo. The law is commonly written as a function of
luminance, but this differs from the luminance factor only by the amount of the
solar illuminance, a constant at any given distance from the Sun.
Results of Detailed Martian Photometry
Young and Collins (1971) used the Minnaert law to reduce data from the
far-encounter pictures of Mariners 6 and 7. The data, all taken near a phase
angle of 22° , resulted in quite sensibly linear plots of log (pcos e) versus log
(cos i cos € ) for each of a number of individual Martian features, indicating the
Minnaert law to be a good one. There were differences in the results of the two
Mariners, however, indicating calibration difficulties. Reference should be
made to the Young and Collins (1971) article for the quantitative results.
Binder and Jones (1972) carried out an extensive program of Martian
photometry during the 1969 opposition. Lacking sufficient data to derive a com-
plete photometric function without some mathematical framework, they too
chose to use the Minnaert formalism to display their results, which give the
most extensive detailed coverage of the Martian surface yet attempted by
modern means. They utilized a 10-channel spectrophotometer with medium-
width passbands, centered upon wavelengths from 0.595 (jl to 2.270 [x. Results
were obtained for four phase angles (7.2°, 10.3°, 17.7°, and 18.5°) and for
156 different points on the disk. Binder and Jones' data fit into the Minnaert
law very nicely and give evidence that the Minnaert parameter K varies linearly
with phase angle and with wavelength over a wide range of values of these quan-
tities, making extrapolations beyond the observations conceivable, though not
desirable. Having no observations near zero phase, it was not possible for
them to discuss any opposition effect, of course. In fact, lack of opportunity
to obtain sufficient reference to standards made it impossible to obtain a normal
albedo map. Good data were obtained at 10.3° phase angle. Figure 4 is a plot
of "reference albedo" (the average of the albedo at 1.04 (x and 1.24 [jl at 10.3°
phase) taken from Binder and Jones (1972). They suggest that the true normal
albedos should be 10-15% larger. Binder and Jones data seem sufficiently
important that the "spectral albedos" (luminance factors at 10.3° phase) for each
of their 10 passbands are listed in Appendix A for 150 observed points.
Binder and Jones (1972) found that their reference albedos tended to
cluster around two values, 17% and 35%. Ratios of albedos at different wave-
lengths, a direct measure of color, also tended to fall into two groups, one
associated with each albedo group. This is a quantitative measure of the obvious
division of Mars into bright and dark areas. Binder and Jones (1972) also made
Sec. 3.2, page 16 R. Newburn, JPL October 1, 1971
JPL 606-1
Ultraviolet, Visible, and Infrared Photometric Properties
s s
T3
C
(T)
00
0)
C
Ml
W
o
z
o
> Vi
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8
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October 1, 1971
R. Newburn, JPL
Sec. 3. 2, page 17
Ultraviolet, Visible, and Infrared Photometric Properties
JPL 606-1
South Polar Cap measurements in six areas. These showed infrared albedos as
low as the dark areas. They suggest these may have been caused by the large
values of i and e at which they had to observe, or perhaps by "contamination"
from underlying or surrounding dark areas.
McCord and Westphal (1971) intercompared seven areas of Mars at
52 wavelengths from 0.30 [x to Z.52 |a. Their work is well standardized but
covers an insufficient range of phase angles to give any information on the photo-
metric function. Their results, showing "geomietric albedo" (actually the radi-
ance factor at 5° phase) versus wavelength, shows the full range of Martian
reflectivity from darkest Syrtis Major to brightest Arabia, and is reproduced
here as Fig. 5.
m
II
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WAVELENGTH (/a)
1.0
I.I
Fig. 5. Spectral radiance factor at 5° phase for seven Martian areas
(from McCord and Westphal, 1971). The insert at the upper left
shows the ultraviolet parts of the curves on an expanded scale.
Sec. 3. 2, page 18
R. Newburn, JPL
October 1, 1971
JPL 606-1 Ultraviolet, Visible, and Infrared Photometric Properties
O'Leary and Jackel (1970) made scans of Mars at very small phase
angles (1.3° < a < 3.5°) in 1969. These seemed to indicate a greater opposition
effect for dark Syrtis Major than for bright Arabia.
Barth and Hord (1971), reporting on Mariner Ultraviolet Spectrometer
results, were unable to present useful ultraviolet albedos because they lac^-ed
comparison spectra of the Sun. They were able to comipare the South Polar Cap
to the Argyre I desert region under virtually identicjal photometric conditions,
finding the polar c^ap three times as bright at 2000 A and 28C0 A and about twice
as bright at 2500 A. Barth and Hord suggest that the Martian surface has an
extremely low normal albedo in the ultraviolet except for the polar caps, and
that the observed radiation at these wavelengths is mostly scattered by the
atmosphere (except over the poles). With this assumption, there should be a
direct correspondence between ultraviolet intensity and atmospheric pressure,
and, therefore, topography. In fact, a good correlation exists between ultra-
violet topography and the Mariner infrared spectronneter topographic results.
3. 2. 5 General Photometric Conclusions
It is very difficult to do accurate detailed absolute photometry of an
angularly small body such as Mars. A large amount of pioneering in visual and
photographic photonnetry (not discussed here) served to provide a qualitative
indication of many interesting Martian problems, but it is generally worth 'ess
quantitatively. Even the best modern photoelectric work, as discussed ir pre-
ceding paragraphs, has not been totally adequate, mostly because no inoivndual
worker has taken data over a sufficient time span to obtain adequate phase -angle
coverage, even for a "static" planet. There is no acceptable quantitative mea-
sure of the many supposed dynamic effects, such as the seasonal "wave of
darkening. " The many classic attempts, to compare changes in dark to bright
areas as a function of season, have not proven which area is changing or whether
both areas may change. Furthermore, such studies often did not prove whether
the changes were true seasonal changes or photometric anomalies. Hopefully,
the combined efforts of spacecraft and modern ground-based systems will
rectify our lack of knowledge, although a very long-lived spacecraft will be
needed.
Both integrated and detailed photometry of Mars have been presented in
this section, with little reference to surface or atmosphere. Young (1969) has
estimated the Martian atmospheric extinction at about 3% near X6odo, and
Barth and Hord (1971) find it rising to 10% at \3050. The photometry in visible
and infrared wavelengths therefore refers to the surface, except possibly near
the limbs. Martian photometry primarily refers to the atmosphere only in Lhe
ultraviolet, where surface albedo becomes very small while extinction rises.
One result is conclusive, Mars is a very RED planet in visil^le wave-
lengths. Even the dark areas which are often described as predominantly green
or greenish-grey, are in fact red, having much higher reflectivity at wave-
lengths of 6OOOA and longer than at wavelengths of 5000 A and less. The dark
areas are less red than the bright areas, but are nonetheless red. They appear
"greenish" to the eye, principally because of visual effects such as the "color-
brightness" effect, which makes "areas darker than the average brightness
October 1, 1971 R. Newburn, JPL Sec. 3.2, page 19
Ultraviolet, Visible, and Infrared Photometric Properties JPL 606-1
appear in a contrast color to the illuminant" (Schmidt, 1959). Astronomical
photographic exposvires are relatively long, and there is reciprocity failure of
the red layer of color film, causing it to lose sensitivity relative to the blue and
green layers. The result of this is to flatten the apparent spectral reflectivity
curve of a dark area. A bright area is also affected, but it retains enough red
to simply make it appear more yellow than is realistic. In many color prints,
the dark areas appear sufficiently dark that they show little color of any kind
(often due to contrast exaggeration through repeated re-copying).
No attempt has been made in this section to discuss the several subtle
spectral features caused by the Martian surface materials or to interpret the
reported photometric results. These topics are covered in Sections 3. 4 and
3. 5. Martian infrared and radio emission properties are covered in Section 3. 1.
3. 2. 6 Polarimetric Nomenclature and Results
Introduction to Po ' arimetry
Astronomical polarimetry has not been a popular field of research, and
Martian polarimetry has been no exception. It might best be described as a
field of unrealized over expectations. Polarimetry promises (and delivers) more
than simple total intensity photometry, but cannot be expected to produce chem-
ical anal-ses of solids. Polarimetry can offer insights into particle size and
can furnish indices of refraction (when a sufficient range of phase angles is
covered). It can "cast the deciding vote in a close race, " but cannot by itself
furnish unambiguous statements, for example, about composition.
Planetary oolarimetry was pioneered in France in the early 1920's by
Lvc'-, and has been carried on assiduously in that country by his pupil Dollfus
and coworkers, l.i fact, the majority of all polarization studies of Mars ever
carried out anywhere have been those of Dollfus and Focas in France. Other
work has included that of Morozhenko in Russia, and Gehrels, Hall, and
Ingersoll in the United States.
A complete description of monochromatic polarized lightrequires four
parameters, those of Stokes being commonly chosen (see Shurcliff, 1962).
However, in common astronomical practice at optical wavelengths, all four
parameters are not measured. Normally, the maximum intensity I^ax (^^ the
plane containing the electric vector and the direction of propagation), the mini-
mum intrnsity I^i„ (perpendicular to that plane), and. sometimes, the orienta-
tion of the electric vector (or magnetic vector) are measured. These measure-
ments are used to derive a quantity called the degree of polarization V, defined
as follows:
I - I .
max min
I + I .
max mm
or, more commonly, the percent of polarization P:
P = 100 V
Sec. 3. 2, page
20 R. Newburn, JPL October 1, 1971
JPL 606-1 Ultraviolet, Visible, and Infrared Photometric Properties
and, sometimes, the permil of polarization, written %o and defined as 10 P.
Some of the techniques used for these measurements have been described by
Hiltner (1962) and Gehrels and Teska (I960). A typical paper may then include
as little as the percent of polarization, as a function of phase angle in some
specified passband and location on the source, or as much as the percent of
polarization, the total intensity (Imax + Imin)' and the orientation of the electric
vector at many wavelengths and phase angles for many different locations on the
object under study. When the component of the electric vector perpendicular to
the phase plane is the larger, polarization is said to be positive. When the
component in the phase plane is larger, the polarization is negative. The theory
has been developed by Fymat and Abhyankar (1970) for an accurate inter fer-
ometric determination of the complete Stokes vector (all four Stokes parameters)
at very high spectral resolution. A polarimetric observing program of Venus
is now being attempted by scientists at the University of Arizona.
Unfortunately, there is no complete "direct" theory of the interaction
of light with a complex real surface; only a theory of interaction with smooth,
homogeneous surfaces such as optical elements. The "inverse" problem,
observing polarized radiation and attempting to derive the nature of the surface
from which it was scattered, has not been solved. Attempts to interpret obser-
vations of Mars have been strictly empirical, a process of comparing planetary
observations to observations of laboratory samples.
At shorter wavelengths, there is a contribution to polarization from the
atmosphere. Much early work assumed a pure Rayleigh atmosphere (a pure
molecular scattering atmosphere) and made attempts to derive the total atmo-
spheric abundance. These attempts failed for many reasons (see Sec. 5. 2;
Chamberlain and Hunten 1965; Coulson, 1969). More recent work has been
largely at wavelengths near 6000 K. or longer, in an attempt to model surface
particle size and composition (see Section 3. 4).
Observations
The general polarimetric behavior of Mars can be summarized fairly
simply. At small phase angles, polarization is negative, reaching about -1.0%
near 12° phase. The inversion angle (the phase angle at which polarization goes
from negative to positive) is between 24° and 30°, and polarization reaches
+ 2% between 40° and 45° and apparently continues rising (DoUfus and Focas,
1969). In general, the dark surface areas are more highly polarized than the
bright areas, following the "Umov effect, " which states that, for materials in
general, the degree of polarization is an inverse function of normal albedo.
The dark markings appear to show a decrease in polarization during Spring and
early Summer (Focas, 1961). Clouds sometimes appear to cause polarization
effects (see Section 4. 1 and Dollfus, 1961), A more complete review of the
observational data has been given by Pollack and Sagan (1969).
October 1, 1971 R. Newburn, JPL Sec. 3,2, page 21
Ultraviolet, Visible, and Infrared Photometric Properties JPL 606-1
BIBLIOGRAPHY
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Binder, A. B., and Jones, J. C. , 1972, Spectrophotometric studies of the photo-
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Sec, 3, 2, page 22 R. Newburn, JPL October 1, 1971
JPL 606-1 Ultraviolet, Visible, and Infrared Photometric Properties
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Sec. 3.2, page 24 R. Newburn, JPL October 1, 1971
JPL 606-1
Ultraviolet, Visible, and Infrared Photometric Properties
APPENDIX A
MARTIAN ALBEDOS
, in I-^^^ appendix identifies the albedos of 150 Martian mare and desert areas
at 10 different wavelengths and at 10.3" phase angle, as determined by Binder
and Jones (1972) Figure A-1 indicates the areas observed. The numbers con-
,u' H f'" \^' . ^'''' "°^ consecutively complete, as some were deleted during
the data reduction process. Detailed results for each number-identified area
are contained in Table A-1. The entries for each numbered area identify lati-
tude and longitude m degrees, and albedo in percent. Entires prefixed with D
TlheTo fnlrlVn"} '" ^'f ' '°' '"'''"'" ^""^ longitude, and in percent for
util^^eH T^ !f- f ^^'.^ ^""^^ ^''^""^ ^^^ correlated to the 10 wavelengths
nhir V ;• ^°°^dinates change with wavelength because of terrestrial atmos-
pheric dispersion. In a few cases, this caused an anomalous albedo curve by
moving the infrared observations onto a different physiographic unit. Most
L°ducw'\rwterr"''"^^ "^^^ ^^^^^^^' '' ^^^^^^ ^"^ J-- ^-^"g ^^^ d-ta
We are particularly grateful to the authors (A. B. Binder and J. C. Jones)
for sending us this important material prior to publication. '
October 1, 1971
R. Newburn, JPL Sec. 3. 2, Appendix A. page 1
Ultraviolet, Visible, and Infrared Photometric Properties
JPL 606-1
C
o
H)
•V
C
n)
(U
a
>■
XI
•H
;3
en
a
• H
U
X*0 OO'OZ- 0O*O»- 00*09-
X-0» 00-02
(S33aS30) aOniliVI 3l81N3303aV
Sec. 3.2, Appendix A, page 2 R. Newburn, JPL
October 1, 1971
JPL 606-1 Ultraviolet, Visible, and Infrared Photometric Propertiei
Table A-1. Martian albedo
WAVELENGTH
•60 •6'' -70 .89 1.04 1.24 1.61 1.74 2.14
2.27
NUMBER LATITUDE
2 DLATITUDE
9.0
2.0
9.7
2.0
10.6
2 .0
11.9
2 .
12.5
2.0
36.9
3.
13.1
2.0
36.6
3.0
31.9
.5
13.5
2.0
36.3
3.0
33.2
.4
13.7
13.9
13.9
LONCilTUDE
DLONGITUDE
39,0
3.0
33.6
3.0
38.1
3.0
37.3
3.0
2 .
36.2
3,0
36.1
,5
2.
36.1
2.0
36.
ALBEDO
DALBEDO
18.2
.6
23,2
.5
27.7
.4
29.8
.4
31.8
.4
3 .
30 .4
.7
3 .
31.8
1.0
NUMBEF
3
i LATITUDE
DLATITUDE
LONGITUDE
DLONGITUDE
12.0
2.0
28,0
3.0
12.7
2.0
27.5
3.0
13.6
2.0
27.0
3.0
14.9
2.0
26.1
3.0
15.5
2.0
25.7
3 ,
16.1
2.0
25.3
3.0
31.2
.8
16.6
2.0
25.0
3.U
33.1
.9
16,8
2,0
24,9
3.0
35,7
1,0
16.9
2.0
24.8
17.0
2.0
24.7
ALBEDO
DALBEDO
17.3
.5
22,4
.7
27.1
.6
29.1
.6
31.2
.7
3 ,
30.9
1.1
3.
32.9
1.7
NUMBER
4
LATITUDE
DLATITUDE
LONGITUDE
DLONGITUDE
ALStDO
DALBEDO
16. 3
2.0
15.9
3.0
17.6
.8
17.6
2.0
15.4
3,0
22.0
1.0
13.5
2.0
14.7
3.0
26.7
1.2
19.8
2.0
13.7
3.0
29.0
1.3
20.5
2.0
13.2
3.0
31.9
1.5
21.1
2.0
12.7
3.0
31.8
1.5
21.6
2.0
12.3
3.0
33.6
1.6
21,8
2,
12,2
3,
36.8
1 .8
21.9
2.0
12.0
3.0
31.3
1.8
22.0
2.0
12.
3.
32.5
1.9
NUMBER
LATITUDE
-12.7
-12.0
-11.0
-9.7
-9.0
-8. 4
-7.9
-7.7
2.0
6.5
2.0
23.1
1.3
-7.6
2.
6.4
2.0
13.3
1.4
5
DLATITUDE
LONGITUDE
DLONGITUDE
2.0
8.9
2.0
2.0
8,5
2.0
2.0
8.1
2.0
2.0
7.4
2.0
2.0
7.1
2.0
2.0
6.8
2.
2.0
6.5
2 .
-7.5
2.0
6.3
ALBEDO
DALBEDO
14.6
.9
16.3
1.0
19.4
1.2
20.1
1.1
21.0
1.1
20.3
1.2
21.4
1.2
2 .
19.1
1.4
NUMBER
8
LATITUDE
DLATITUDE
41.0
3.7
41,9
3,7
43.2
3.3
45.0
3 . 9
45.9
3.9
46.7
4.0
22.2
2.0
47.4
4.0
21,6
2.
47.6
47.9
48.0
LONGITUDE
DLONGITUDE
26.9
2.0
26.2
2.0
25.2
2.0
23.7
2.0
22 ."9
2.0
4 .
21.3
2.0
16.5
1.5
4 ,
21.1
2.0
13.6
1.3
4.1
21.0
2.0
13.2
1.5
ALBEDO
DALBEDO
14.2
1.0
15.9
1,0
18.3
1.3
17.7
1.4
17.8
1.4
16.8
1.4
15.8
1.4
NUMBER
9
LATITUDE
DLATITUDE
41.0
3.7
42,0
3,7
43.2
3.8
45.0
3 . 9
46.0
?, 9
46.8
4.0
22.2
2.0
16.3
1.4
47.5
4.0
21.5
2.0
15.7
1.4
47.3
48.0
48.1
LONGITUDE
DLONGITUDE
ALBEDO
DALBEDO
26.9
2.0
13.4
1.0
26,2
2.0
15.6
1,0
25.2
2.0
18.2
1.3
23.7
2,0
17.4
1.4
0.7
22.9
2.0
17.5
1.4
4.
21.3
2.0
16.5
1.5
4.1
21.0
2.0
13.4
1.4
4.1
20.9
2.0
13.4
1.6
NUMBER
10
LATITUDE
DLATITUDE
LONGITUDE
DLONGITUDE
8.9
2.0
39.0
3.0
9.8
2.0
38.7
3,0
10.6
2.0
38.3
3.
12.3
2.0
37.7
3 . n
13.0
2.0
37.4
3.0
31.2
.6
13.7
2.0
37.1
3.0
31.3
.'7
14.2
2.0
36.9
3.0
32.9
.7
14.4
2.0
36,8
14.6
2.
36.7
14.7
2.0
36.7
ALBEDO
DALBEDO
17.9
.6
22,5
.5
27.1
.6
w . u
29.4
.5
3,
34.7
.9
3.
29.2
.9
3,0
31,2
1,0
NUMBER
11
LATITUDE
DLATITUDE
LONGITUDE
DLONGITUDE
ALBEDO
DALBEDO
27.9
2.0
31.8
2.0
14.3
.8
28,8
2.0
31,3
2,0
16,5
,8
29.9
2.0
30.7
2.0
18.5
.8
31.5
2.0
29.8
2.0
18.0
.9
32.4
2.0
29.3
2.0
18.3
1.0
33.1
2.0
28.8
2.0
17.1
1.1
33,7
2.0
28,5
2.0
16.3
1.3
33.9
2.0
28.3
2.0
17.1
1.4
34,1
2.0
28.2
2.0
12.9
1.2
34,3
2.0
28.1
2.0
15.0
1.3
NUMBER
12
LATITUDE
DLATITUDE
28.0
2,
26,9
2,0
30.1
2.0
30.9
2.0
19.2
.8
31.7
2.0
30.
2.0
19.0
.9
32.6
2.0
29.5
2.0
19.1
.9
33,3
34.0
34.2
34. 4
34.5
LONGITUDE
DLONGITUDE
ALBEDO
DALBEDO
32.
2.0
14.2
1.1
31.5
2,0
17,0
,8
2.0
29.0
2.0
17.9
.9
2.0
28.6
2.0
17.0
.9
2,0
28.5
2,0
17,9
1.
2.0
28. 4
2.0
14.2
1.0
2.0
28, 3
2.0
14,3
. 9
October 1, 1971 R. Newburn, JPL Sec. 3.2, Appendix A. page 3
Ultraviolet, Visible, and Infrared Photometric Properties JPL 606-1
Table A-1. Martian albedo (continued)
WAVELENGTH
.60 .64 .70 .89 1.04 1,24 1.61 1.74 2.14 2.27
MUM8ER
13
LATITUDE
DLATITUDE
LONGITUDE
DLONGITUDE
ALBEDO
DALBEDO
-11.1 -
3.1
27.1
3,0
13.3
1.1
10.2
3.1
26,9
3.,0
15.2
.s
-9.1
3.1
26.6
3,0
16.7
1.1
-7.5
3.1
26.2
3,0
16,5
,3
-6,7
3.1
26.0
3.
16.7
,8
-6.1
3.1
25.8
3.0
16.6
.7
-5,5
3.1
25./
3.0
17.3
.6
-5.3
3.1
25,6
3.0
18.3
.8
-5.1
3,1
25.6
3.0
15.5
,8
-5.0
3.1
25.5
3.0
15.5
.9
NUMBER
14
LATITUDE
DLATITUDE
LONGITUDE
DLONGITUDE
ALBEDO
DALBEDO
-24.9 -
3.5
24.9
3.9
12.6
1.2
23.9 -
3.5
24.9
3,9
15.2
1.2
22.7 -
3.4
24.8
3.8
16.9
1.3
20.9 -
3,4
24.6
3,8
16.8
1.3
20.0 -
3.3
24.5
3.8
16.7
1.2
19.2 -
3.3
24.5
3.7
16.9
1.2
18.6 -
3.3
24 .4
3.7
17.6
1.2
16.4 -
3.3
24.4
3.7
18,3
1,3
18,2 -
3,3
24.3
3.7
15.5
1.2
13.1
3.3
24.3
3.7
15.9
1.2
NUMBER
16
LATITUDE
DLATITUDE
LONGITUDE
DLONGITUDE
ALBEDO
8,9
2.0
39,0
3,0
17,6
9,9
2.0
33.7
3.0
22.5
11.0
2.0
38.4
3.0
26.8
12.7
2,0
37,9
3.0
28.8
13,6
2.0
37,7
3,0
30,8
.9
14.3
2.0
37.5
3.
31.0
.9
15.0
2.0
37.3
3.0
32.9
1.0
15,2
2,0
37.2
3.0
35,5
1 ■ 2
15.4
2,0
57.2
3.0
29.8
1 . 1
15.5
2.0
37.1
3.0
31.4
1,3
DALBEDO
.7
.8
,7
. 9
NUMBER
17
LATITUDE
DLATITUDE
LONGITUDE
DLONGITUDE
ALBEDO
DALBEDO
-19.1 •
3.3
38.1
3.5
12.3
L.
•18,0 ■
3.3
38.0
3.5
15.5
1.0
■16,7 ■
3,2
37.9
3.4
17.1
1.1
■14.8 •
3,2
37.7
3.4
16,8
.9
■13,9 •
3.2
37,6
3,4
17,3
,a
■13.1 ■
3.2
37.6
3.4
17.5
.6
■12.4 ■
3.2
37.5
3.4
18,4
.8
-12,2 ■
3,2
37,5
3,4
19.5
.9
■12.0 •
3.1
37.4
3,3
15.9
.9
•11.9
3.1
37.4
3.3
17.0
1.3
NUMBER
19
LATITUDE
DLATITUDE
LONGITUDE
22.9
3.1
76.1
24.0
3.2
76.1
25.3
3.2
76,1
27.3
3.2
76.1
23.3
3.3
76.1
29.2
3.3
76.1
3.4
37.3
1.6
30.0
3,3
76.1
3.4
38.8
1.8
30,2
3,3
76.2
3.4
41,7
2,3
50.5
3.3
76,2
3 , 4
30.6
3,3
76,2
3, 4
DLONGITUDE
ALBEDO
DALBEDO
3.3
20.3
.8
3,3
25.7
.8
3.3
31.2
1.1
3 . 4
34.5
1.3
3,4
36.9
1.5
35,6
2.2
36,0
2.0
NUMBER
21
LATITUDE
DLATITUDE
LONGITUDE
DLONGITUDE
ALBEDO
DALBEDO
8.9
2.0
39,0
3.0
18,2
.7
10.0
2.0
38.8
3.0
23,3
.9
11.3
2.0
38.6
3.0
27,9
1,2
13,3
2.0
38,2
3.0
30,2
1.1
14.3
2.0
38.0
3,0
32.0
1.2
15.1
2.0
37.8
3.0
32.2
1.3
15.9
2.0
37.6
3.0
33.2
1.4
16,1
2,0
37.6
3.0
35.7
1.5
16.4
2,0
37.5
3.0
30,2
1,5
16.5
2.0
37,5
3.0
30,3
1,5
NUMBER
22
LATITUDE
DLATITUDE
LONGITUDE
DLONGITUDE
ALBEDO
DALBEDO
23,0
3,1
247.0
3,4
17.5
.7
23.5
3.1
246.0
3.5
21.1
.9
24.1
3.1
244.7
3.5
24,6
1,1
25,0
3.1
242.9
3.6
26,6
1,3
25.4
3, -2
241.9
3.6
28.0
1.4
25.8
3,2
241,1
3.7
28,2
1.5
26.1
3.2
240,3
3.7
29.0
1,6
26,2
3.2
240.1
3,7
29,6
1.7
26.3
3.2
239,3
3,7
26.7
1,7
26.4
3,2
239.7
3.7
27.2
2.0
NUMBER
23
LATITUDE
DLATITUDE
LONGITUDE
15.0
3.0
240.0
15.5
3.0
239.0
16.1
3.0
237.8
16,9
3,0
235,9
17.3
3,0
235.0
17,7
3.0
234.1
18,0
3,0
233,4
18.1
3.1
233.1
18,2
3.1
232.9
4.0
27.1
2.0
18.2
3.1
232.8
4.0
28,3
2.2
DLONGITUDE
ALBEDO
DALBEDO
3.6
17.0
.9
3.6
20.5
1.1
3,7
24,2
1.4
3.8
25,9
1,6
3,9
27,6
1,6
3,9
27,9
1,9
4 .
29,0
2.0
4 .
29.9
2.2
NUMBER
24
i LATITUDE
DLATITUDE
LONGITUDE
DLONGITUDE
ALBEDO
DALBEDO
14.9
3.0
253.0
3.2
17.5
.5
15.4
3.0
252.1
3.3
20.9
.7
16,0
3,0
251.1
3,3
24,2
,8
16,9
3,0
249,5
3,3
25,4
1,0
17.3
3,0
248,7
3.4
26.5
1.1
17,7
3,0
248,0
3.4
26.3
1,1
18.0
3.0
247.4
3.4
26.9
1.1
18.1
3.1
247,2
3.4
27.8
1.3
18.2
3,1
247.0
3,4
24,2
1,2
18.3
3.1
246.9
3.4
25.1
1,2
Sec. 3.2, Appendix A, page 4 R. Newburn, JPL October 1. 1971
JPL 606-1
Ultraviolet, Visible, and Infrared Photometric Properties
Table A-1. Martian albedo (continued)
,60
.6'i
70
WAVC-LEKGTH
.69 1.04 1.24 1.61 1.74 2.14 2.27
MUMBER
LATITUDE
16.9
:.7.5
IB. 2
19.2
19.7
20.1
20.5
20. 6
20 .7
20 . 3
25
DUATITUDS
2.
2.0
2.0
2.U
2.0
2.0
2,0
2.0
2,
2.
LU\'C1TUD£
265.1
2£4.4
263 . i
282.0
281 .6
281.1
280.6
280.4
230.3
2 e ' 2
DUONGITUDE
2.0
2.0
2.0
2. 3
2.0
2.0
2.0
2.0
2 .
2 .
Aiasoo
1C.9
12,7
14.2
13.6
14.1
14.3
15.6
16,3
13.3
13 . 7
DALfitDO
••■'
.3
. 3
.3
.2
.3
.2
,5
. 4
.7
NUMBER
LAT iTUDE
-<.9
-4.4
-5.8
-2. a
-2.4
-2.0
-1.7
-1.5
-1.4
-1.4
26
DLATITUJE
3.0
3.0
3.0
3.0
3.0
3.0
3.0
3.0
3.
3.
LONGITUOE
252.0
251.2
250.3
249.0
248.4
247.8
24 7.3
247.1
246.9
246 . 8
DUONGITUDE
3.4
3.4
3.4
3 . 4
3.5
5.5
3.5
3.5
3.5
3 ,5
Al aSDO
11.8
1^.7
15.9
14.7
15.0
15.0
15.5
16.1
13.6
14.6
DALBEDO
.6
.6
.7
.7
.6
.7
.7
. 7
.8
1.0
NUMBER
LATiTUDE
-11.0
-10.4
-9.7
-8.G
-8.3
-7.8
-7.5
-7.3
-7.2
-7 .2
27
Dl,ATiTUjE
3.0
3,0
3.0
3.0
3.0
3.0
3.0
3.0
3.
3.
LONGITUDE
262.0
261.4
260 .6
259.4
258.3
258.3
257.9
257.7
257 .6
257 .5
DUONGITUDE
3.3
3.3
3.3
3.3
3.3
3.3
3.3
3.3
3.3
3. 3
ALBEDO
11.6
13.8
15.5
15.2
15.3
15.5
16.2
16,6
14.9
16. 4
DALDL-DO
.4
. p
.5
.6
.5
.5
• 6
,6
.7
1.2
NUMBER
LATITUDE
-15.0
-14.4
-13. a
•12.8
-12.3
-11.8
-11.5
-11.4
-11.2
-11 . 2
28
DUATITUDE
3.0
3,
3.0
3.0
3.0
3.0
3.
3.
3.
3.
LONGITUDE
249.0
2'. 3. 4
247.5
246.4
245.7
245.2
244.8
244.6
244.4
244. 4
DUONGITUDE
3.3
3.3
3.8
3.9
3.9
3.9
3.9
3.9
3.9
3.9
ALBEDO
11.9
13,6
15.5
14,6
15.0
15.3
16.1
16.6
14.2
16.7
DALBEDO
.8
.9
1. 1
1.1
1.1
1. 1
1.2
1.4
1.1
1.7
NUMBER
LATITUDE
-29.9
-29.2
-20.4
-27.2
-26.5
-26,
-25.6
-25.4
-25.3
-25.2
29
DLAT ITUDE
3.8
3.7
3.7
3.7
3.6
3.6
3.6
3.6
3.6
3.6
LONGITUDE
270.3
270,2
269.4
268.3
267.8
267.3
266.9
266.8
266.7
266 . 6
DLONGITUOE
4.0
".0
4.0
4.0
4.0
4.0
4,0
4,0
4.0
4.
ALBEDO
13.6
18.3
23.5
25.7
27.1
27.3
27.9
29,6
26.7
29, 6
DALBEDO
1.3
1.7
2.0
2.2
2,3
2.3
2.2
2,5
2.5
2.6
NUMBER
LATITUDE
21.9
22,5
23.3
24.4
24.9
25.4
25,8
25,9
26.
26. 1
30
DLATITUDE
3.1
3.1
3.1
3.1
3.1
3.2
3.2
3,2
3.2
3.2
LONGITUDE
273.9
273.3
272.4
271.2
270.6
270.1
269,6
269,4
269.3
269 ,2
DLONGITUDE
3.1
3.2
3.2
3.2
3.2
3.2
3,3
3,3
3.5
3.3
ALBEDO
14.0
16.7
19.6
2U.0
20.9
21.0
21.8
23,6
20 .3
22. 4
DALBEDO
.6
.4
.6
.7
.9
.9
.9
1.0
1.
1.5
NUMBER
LATITUDE
32.0
32.6
33.4
34.6
35.2
35.7
36.2
36.3
36.5
36 . 6
31
DLATITUDE
3.3
3.3
3.3
3.4
3.4
3.4
3.4
3 .4
3 .4
3 . 4
LONGITUDE
280.9
280.2
279.3
278.1
277.4
276.9
276.4
276.2
276.
275.9
DUONGITUDE
3.3
3.3
3.3
3.4
3.4
3.4
3.5
3.5
3,5
3.5
AUBEDO
17.5
21.1
25.1
26.3
27.5
.27.5
28.3
30 .1
25,3
27 . 4
DAUBEDO
.8
.8
.9
1.1
1.1
1.2
1.3
1.3
1.4
1.3
NUMBER
LATITUDE
15.0
15.6
16.4
17.5
18.0
18.5
18.9
19.
19.2
19 . 3
32
DuATITUDE
2.0
2,0
2.0
2.0
2.0
2.0
2.0
2,
2 .
2 .
LONGITUDE
DUONGITUDE
287.0
2.0
286.5
2,0
285.8
2.0
284.8
2.0
284.3
2.0
283.3
2.Q
283.5
2.0
233.3
2.0
233.2
2 .
283.1
2 .
AUBEDO
10.6
12.6
14.2
13.6
14.1
14.5
15,9
17.0
14 .
14. 9
DAUBEDO
.5
.5
.3
.4
.4
.3
.4
.4
.6
.8
NUMBER
UATITUDE
7.0
7.6
8.4
9.5
10.0
10.5
10.9
11.1
11.2
11.3
33
DUATITUDE
2.0
2.0
2.0
2.0
2.0
2.0
2.0
2 .
2 ,
2 .
UONGITUDE
236.0 ,
235.5
234.3
233.9
283.4
235.0
282.7
232.5
232.4
282! 4
DUONGITUDE
2.0
2.0
2.0
2.0
2.0
2.0
2.0
2.0
2.
2 .
AUBEDO
DALU'EDO
9.6
.3
11.9
.2
13.0
.2
11.6
.2
11.5
.2
11.4
.2
12.2
.3
12.3
.2
10.7
.3
ll!c
. 6
October I, 1971
R. Newburn, JPL Sec. 3. 2, Appendix A, page 5
Ultraviolet, Visible, and Infrared Photometric Properties JPL 606-1
Table A-1. Martian albedo (continued)
.60
WAVELENGTH
,70 .69 1.04 1.24 1.61 1.74 2.14 2.27
NUMBEH
LATITliCE
-33.9 -33,0 -32.0 -30.5 -29.8 -
29.2 -28.6 -23,5 -28.3 -28.2
34
OLaTI 'UDE
L Q N' G ! T U D E
4.0 '5,0 3.9 3,8 3.8
300. 5 299.8 298.9 297.7 297.1 2
3.7 3.7 3,7 3.7 3.7
96.6 296.2 296,0 295.9 295.8
Dl,0''!;ITUDE
4.1 4,0 4.0 3,9 3.8
3.8 3.8 3,' 3.7 3,7
ALBLLiO
13.4 15.8 17.9 17.1 17.7
17.3 17.7 18.9 15.3 17.2
DALBEDO
1.2 1,3 1.3 1.2 1.2
1.1 1.1 1.2 1.1 1.2
NUMBER
LAT ; TUDE
45.0 45,8 46.8 48.2 48.9
49.6 50.1 50.3 50.5 50.6
35
D LATITUDE
LONGITUDE
20 2,0 2.0 2.0 2.0 2.0 2.0 2.0 2.0 if.U
276.1 275.2 274,1 272.3 271.3 270.5 269.7 269.4 269.2 269.1
DLC'JG I TUDE
3,9 3,9 4,0 4.1 4.2
4.3 4.3 4.3 4.4 'i .<*
AL Ec DO
13.1 l^'.y 16.6 15.7 15.6
14.5 13.9 14.7 12.4 12.5
DALBEDO
1.0 1.1 1.2 1.3 1.3
1.3 1.2 1.4 1.2 1.5
NUMBER
LATITUDE
24,9 25,6 26.5 27.7 28.3
28.9 29,3 29.5 29.7 29.7
36
DLATHUIJE
LONGITUDE
3.1 3,2 3.2 3.2 3.2
304.1 303,7 303.1 302.3 301, - :
3.2 3.2 3,2 3.2 3.2
501.5 301.2 301.1 300.9 300.9
DLONGITUDE
ALBEDO
DALBEDO
3.2 3,2 3.2 3.3 3.3
18.5 22.9 27.0 28.5 30.5
.7 .8 .9 1.0 1.1
3.3 3.3 3.3 3.3 3.3
30.2 31.2 33.0 28,5 30.2
1.1 1.2 1.2 1.0 1.5
NUMBER
37
LATPUDE
PLATITUDE
LONGITUDE
32.0 32,7 33.6 35.0 35.6
2.0 2,0 2.0 2.0 2.0
296.3 295,8 295.2 294.3 293.8
36.2 36,7 36,9 37.1 37.2
2.0 2.0 2.0 2.0 2.0
293.4 293.0 292.9 292.7 292.7
3.4 3.4 3.4 3.4 3.5
32.0 32.9 35,0 30.2 30.8
1.4 1.4 1.4 1.4 1.7
DLONGITUDE
ALBEDO
DALBEDO
3.3 3,3 3.3 3.4 3.4
19.1 23.4 28.1 30.1 31.7
.7 ,8 1.0 1.1 1.3
NUMBER
LATI'l'UDE
12.9 13,6 14.5 15.8 16.4
17.0 17.5 17,6 17.3 17.9
^-.— ~i*^ "^rt ■TA
38
PLATITUDE
LONGITUDE
3.0 3,0 3.0 3.0 3.0
328.0 327.6 327.1 326.5 326.1
3,0 3.0 3.0 3.0 3.0
325.9 325.6 325.5 325.4 325.4
3.5 3.5 3.5 3.5 3.5
DLONGITUDE
3.6 3,6 3.6 3.5 3.5
ALBEDO
DALBEDO
20.0 25,1 30.0 32.3 34.7
1.1 1,3 1.4 1.5 1.6
34.7 36.1 38.4 33.5 34.4
1,7 1.7 1.8 1.6 1.7
NUMBER
LATITUDE
6.0 6,7 7.5 8.8 9.4
9.9 10.4 10,6 10 .7 10.8
2.0 2.0 2.0 2.0 2.0
288.6 288.4 288.3 288.2 288.1
39
PLATITUDE
LONGITUDE
2.0 2,0 2.0 2.0 2.0
291.0 290,6 290.1 289.3 289.0
DLONGITUDE
ALBEDO
2.0 2,0 2.0 2,0 2,0
1C.4 12.2 13,7 12.4 12.2
2.0 2.0 2.0 2.0 2.0
12.2 13.1 14,0 11.4 11.5
.5 .4 .5 .2 .4
DALBEDO
,6 .4 .3 .4 .4
NUMBER
LATITUDE
3,9 4.7 5.7 7.2 7.9
8.6 9.1 9.3 9.5 9.6
2.0 2.0 2,0 2.0 2.0
333.4 333.3 333.2 333.1 333.1
3.3 3.3 3.3 3.3 3.3
34.9 36.4 38.5 33.6 35.0
1.3 1.3 1.4 1.2 1.5
40
DLATITUUE
LONGITUDE
2.0 2.0 2.0 2.0 2.0
335.0 334.7 334.4 333.9 333.6
DLONGITUDE
ALBEDO
DALBEDO
3.4 3.4 3.4 3.4 3.4
18.7 23.9 29.4 32.1 34.5
.8 1,0 1.1 1.2 1.2
NUMBER
LATITUDE
15.9 16.8 17.8 19.4 20.1
20.8 21.4 21.6 21.8 21.9
41
DLATITUDE
LONGITUDE
3.0 3.0 3.0 3,1 3.1
300.9 300,6 300.2 299.7 299.4
3.1 3.1 3.1 3.1 3.1
299.1 298.9 298.8 298.8 298.7
2.0' 2.0 2.0 2.0 2,0
35.3 36.9 38.8 32.5 34.2
.7 .8 .8 .7 1.1
DLONGITUDE
ALBEDO
DALBEDO
2.0 2.0 2.0 2.0 2.0
17.9 24.1 29.5 33.1 35.0
.5 .5 .7 .8 ,7
NUMBER
LATITUDE
21.1 22.1 23.4 25.3 26.3
27.1 27.8 28.1 28.3 28.5
42
DLATITUDE
LONGITUDE
3.1 3,1 3.1 3.1 3.2
337.0 336,9 336.9 336.9 336.9
3.2 3.2 3.2 3.2 3.2
336.9 336.9 336.9 336.9 336.9
DLONGITUDE
3.2 3,2 3.3 3.3 3.3
3.3 3.3 3,3 3,4 3.4
ALBEDO
18.5 23.7 28.6 31.3 34,7
35.8 38 .7 45,3 36.1 37 .5
DALBEDO
.7 ,7 1,0 1.0 1.1
1.2 1.3 1.7 1.4 1.6
Sec. 3.2. Appendix A. page 6 R. Newburn, JPL October 1, 1971
JPL 606-1
Ultraviolet, Visible, and Infrared Photometric Properties
Table A-1. Martian albedo (continued)
.60
.64
.7u
WAVELE-\GTH
.69 1.01 1.24 1.61 1.74 2.14 2.27
NUMBER
LiTITUDE
35.0
36,1
57.6
39.8
40 .9
41.9
42.7
43.0
43. 3
43.4
43
DLATITUDE
3.4
3,4
3.5
3.5
3.6
3.6
3.7
3.7
3.7
3 . 7
LC.^GITUUE
337.1
337.2
337 ,3
337.6
337 ,7
337.8
337.9
338.0
336.0
336. 1
DLONGITUDE
3.5
3.6
3.6
3.7
3,7
3.8
3.6
3.8
3.9
3 . 9
ALBEDO
17.3
21.4
26.2
26.7
28.9
29.2
30 .8
35.6
28.0
29. 4
DALBEDO
.9
1.2
1.6
2.1
2.2
2.3
2.4
2.9
2.4
2.3
NUMBER
LATITUDE
14.0
14.9
16.0
17,6
18.3
19.0
19.6
19.8
20 .
20 . 1
44
DLATITUDE
3.0
3.0
3.0
3.0
3.0
3.0
3.0
3.1
3.1
3. 1
LONGITUDE
221.0
220.7
220.5
220.0
219.8
219.6
219.4
219.4
219.3
219.3
DLONGITUDE
3.0
3.0
3.0
3,0
3.0
3.0
3.1
3.1
3.1
3.1
ALBEDO
18.8
24.0
29,5
32.4
35.0
35.6
37.6
40,4
34.9
.
DALBEDO
.6
.5
.6
.6
.8
.8
1.0
1.1
1.2
0.0
NUMBER
LATITUDE
11.9
12.8
14.0
15.6
16.4
17.1
17.7
17.9
18.1
18.2
46
DLATITUDE
3.0
3,0
3.0
3.0
3.0
3.Q
3.0
3.0
3.0
3.0
LONGITUDE
202.0
201.7
201.3
200.7
200.4
200.1
199.9
199,8
199,7
199. 7
DLONGITUDE
3.2
3.2
3.3
3.3
3.3
3.3
3.3
3,3
3.3
3 . 3
ALBEDO
19.0
23.6
28.6
30.2
32.0
31.9
33.1
35.8
30.
.'
DALBEDO
.6
.8
1.0
1.0
1.1
1.1
1.2
1.3
1.4
0.0
NUMBER
LATITUDE
-23.0
-21,9
-20.6
-18.7
-17.7
-16.9
-16.2
-16.0
-15.8
-15.7
47
DLATITUDE
2.0
2.0
2.0
2.0
2.0
2.0
2.0
2.0
2.0
2.
LONGITUDE
224,1
223.9
223.7
223,3
223.1
223.0
222.9
222.8
222.8
222. 8
DLOfvGITUDE
3,5
3,5
3,5
3.4
3.4
3.3
3.3
3.3
3.3
3.3
ALBEDO
12.4
14.5
16.9
16.0
16.5
16.7
17,3
16,8
15.3
.
DALBEDO
.7
.3
.8
.7
.8
.8
.9
,9
.9
0.0
NUMBER
LATITUDE
13.9
14,9
16.2
17.9
18.8
19.6
20.3
20.5
2C .7
20 .9
49
DLATITUDE
3.0
3,0
3.0
3.0
3.0
3.0
3. i
3.1
3.1
3 . 1
LONGITUDE
241.0
240,9
240.8
240.6
240.5
240 .4
240 ."^
240.4
240 .3
2 4 0,3
DLONGITUDE
3.1
3.1
3.1
3.1
3.1
3.1
3.1
3.1
3.1
3.1
ALBEDO
17.1
21.1
24.8
26.2
27.8
28.0
28.7
31.3
25.4
,
DALBEDO
.5
.7
1.0
1.0
1.2
1.2
1.5
1.7
1.9
0.0
NUMBER
LATITUDE
4.1
4.5
5.0
5.7
6.0
6.3
6.6
6.7
6. 7
6.8
51
DLATITUDE
3.0
3.0
3.0
3.0
3.0
3.0
3.0
3.0
3.
3.
LONGITUDE
136.0
134.9
133.7
131.8
130.8
130.0
129.3
129.0
128.6
128 . 6
DLONGITUDE
3.4
3.4
3.5
3.6
3.6
3.6
3.7
3.7
3.7
3.7
ALBEDO
19.1
24,0
28.2
31.1
32.6
32.9
34.5
36.5
31.7
33 . 5
DALBEDO
.8
1.0
1.3
1.5
1.7
1.8
2.0
2.2
2.0
2.3
NUMBER
LATITUDE
17.8
18.2
18.7
19.3
19.7
19.9
20.2
20.3
20.3
20 , 4
52
DLATITUDE
3.0
3,0
3.0
3,0
3.0
3.1
3.1
3,1
3. 1
3 . 1
LONGITUDE
131.0
129.8
128.3
126.1
125.0
124.0
123.1
122.8
122.5
122 . 4
DLONGITUDE
3.6
3.7
3.8
3,9
4.0
4.0
4.1
4.1
4 . 1
4 . 1
ALBEDO
19.5
24,2
29.2
31.6
33.3
33.3
34.9
36,6
3 ' .3
33 . 4
DALBEDO
1.0
1.5
1.8
2.2
2.5
2.6
2.9
3.2
2.8
3,3
NUMBER
LATITUDE
-13.9
-13.5
-13.0
-12.2
-11. S
-11.5
-11 , 2
-11.1
■■ j 1 .0
-11.0
53
DLATITUDE
3.3
0.3
3.2
3.2
3.2
V '^
3. 2
'i
LONGITUDE
137.0
136,0
134.9
133.2
132.4
131.6
131.0
130 .8
130 6
1 3 P .5
DLONGITUDE
3.7
3.8
3.8
3.8
3.9
3.9
3.9
3 . V
3.9
3 9
ALBEDO
19.1
23.5
28.5
30.7
32.6
33.1
34.8
3 7,
3?,
34.1
DALBEDO
1.2
1.6
1.9
2.1
2 . 4
2.5
• , /
1 . r.
NUMBER
LATITUDE
-45.0
-44.2
-43.2
-41.9
-41.3
-40.7
-40 .2
-40.1
- 59 . i
- 7; 9 . •'
55
DLATITUDE
5.2
5.1
5.0
4.8
4 .7
4 .7
4 .6
4 . 6
*t , o
A
LONGITUDE
153,8
152.9
151.8
150.1
149.3
148.6
148.0
14/. 8
147.6
1 4 7 .' 5
DLONGITUDE
5.3
5.2
5.2
5.1
5.0
5.0
4 .9
4.9
4.9
4. 9
ALBEDO
13.3
15,0
17.1
17.0
17.5
17.5
18.5
19.9
15.5
15.8
DALBEDO
2.4
2.6
2.9
2.7
2.8
2.7
2.8
3.2
2.6
2.6
October 1, 1971
R. Newburn, JPL Sec. 3. 2, i^ppendix A, page 7
Ultraviolet, Visible, and Infrared Photometric Properties
JPL 606-1
Table A-1. Martian albedo (continued)
WAVELENGTH
.60
NUMBER
56
NUMBER
57
NUMBER
61
NUMBER
63
NUMBER
65
NUMBER
66
NUMBER
67
NUMBER
70
NUMBER
71
LATITUDE
DLATITUDE
LONGI TUDE
DLONGITUDE
ALBEDO
DALBEDO
LATITUDE
DLATITUDE
LONGITUDE
DLONGITUDE
ALBEDO
DALBEDO
LATITUDE
DLATITUDE
LONGITUDE
DLONGITUDE
ALBEDO
DALBEDO
LATITUDE
DLATITUDE
LONGITUDE
DLONGITUDE
ALBEDO
DALBEDO
LATITUDE
DLATITUDE
LONGITUDE
DLONGITUDE
ALBEDO
DALBEDO
LATITUDE
DLATITUDE
LONGITUDE
DLONGITUDE
ALBEDO
DALBEDO
LATITUDE
DLATITUDE
LONGITUDE
DLONGITUDE
ALBEDO
DALBEDO
LATITUDE
DLATITUDE
LONGITUDE
DLONGITUDE
ALBEDO
DALBEDO
LATITUDE
DLATITUDE
LONGITUDE
DLONGITUDE
ALBEDO
DALBEDO
9
3
158
3
IS
25
3
157
3
19
.64
10 .4
3,0
157.2
3.1
23.2
.5
.70
,89 1.04 1.24 1.61 1.74 2.14 2.27
26,
3,
157,
3,
23,
11.
3,
156,
3 ,
28,
27,
3,
155,
3
26,
1
12,
3
11.8
3.0
154.8 154
3.1 3
31.4
.7
33
27.9
3.2
154.
3,
31.
1,
28.
3,
153,
3,
33,
1,
12.
3.
153,
3,
34,
28.
3,
152,
3,
33,
1,
12
3
152
3
35
-4.0
3.1
165.0
3.1
17.5
.5
-3,5 -2.8 -1.9 -1,4 -1
3.1 3.1 3.1 3.1 3
164,3 163.5 162.2 161.6 161
3,1 3.1 3.1 3.1 3
23,1 27.4 29.9 32.2 32
,6 ,8 .7 .7
-34.9 -34,2 -33.3 -32.1 -31.5
4.2 4,2 4.1 4.0 4.0
165.7 165,0 164.2 163.0 162.4
4.3 4,2 4.2 4.1 4.1
12.4 15,1 16.4 15.3 16.3
1.3 1,5 1.6 1.5 1.5
9.9 10,5 11.1 12.0 12.4
3.0 3,0 3.0 3.0 3.0
158.0 157,3 156.4 155.1 154.5
3.2 3,2 3.2 3.3 3.3
18.0 23,8 28.4 30.6 32.8
.5 ,7 .9 1.0 1.1
29,
3,
151.
3,
35,
1
3
160
3
34
-30
3
161.8 161
-30
3
4,
16,
1.
12.8
3.0
4
17
1
13,
3,
1
2
9
4
7
5
7
,0
,6
,1
,3
,8
,5
.9
.4
.1
.1
,5
2
13.
3.
152.
3.
39.
1.
29.
3,
151,
3,
38,
1.
3
160
3
37
-30
3
161
4
17
1
13.
3.
152.
3.
33.
1.
29.
3.
151,
3,
33,
1,
13
3
152
3
34
1
29,
3,
151
3,
34.0
1.8
-.5
3,0 3.
160.3 160,
3.2 3,
31.2
.9
32
1
13.3
3.0
-30,
3,
161,
4,
15
1
13
3
153.9 153.4 153.2 153
3.3 3.3 3.3 3
32.9 34.3 37.9 31
1.3 1.4 1.6 1
19.0
3.0
19,6
3,0
20.
3
178.0 177,4 176
3.0
17.8
.6
3.1
23.8
.7
3
28
3 21.3
1 3.1
6 175.5 174
1' 3.1 3
7 31.5
6 .7
21.8
3.1
33
32.5 33.2
3,3 3,3
170.7 170,0 169.0
3.3 3,3 3.4
24,1 28.8
1.0 1.3
31.9
3.2
18.5
.8
34,
3,
34.3
3.3
167.6 166
3.4 3
31.9
1.5
34
1
22.
3,
174,
3.
34
35,
3,
166,
3,
34,
■ 1.
22.
3.
173,
3
36
22.8
3.1
173.8 173
3.1 3
38.9 32
1.0 1
22
3
35,
3,
35.7
3,3
165.6 165
3.5 3
36.8
1.9
39
2
36.0
3.4
165
3
34
2
30,
3,
161,
4,
15,
1,
13
3
152
3
33
1
23.
3.
173.
3.
32.
1.
36.
3,
165,
3,
34,
2,
4
2
2
4
1
1
9
1
4
5
,5
,0
,9
,3
,6
.8
.0
1
5
1
9
1
1
4
2
5
2
2
8
3
182
3
17
8,6 9.3 10.4 10.9 11.3 11.7 11.9 12.0 12.0
3,0 3.0 3.0 3.0 3.0 3.0 3.0 3,0 3.0
181,4 180.8 179.8 179.3 178.8 178.4 178.3 178.2 178.1
3,0 3.0 3.0 3,0 3.0 3,0 3.0 3.0 3.0
22.8 27.6 29,7 32,3 32,8 34.1 37.7 32,2 31.7
,2 ,5 ,4 .3 .3 .3 .6 .6 .9
•7,4 -6.6 -5,
3,1 3.1 3,
184.0 183.5 182.8 181,
3.1 3,1 3.1 3.
20,5 24.6 26,
,6 ,7
-8.0
3.1
15.8
.7
5
-5,
,0
-4.5
-4.
,1
-3,
,9
-3.8
-3,
7
1
3.
,1
3.1
3,
.1
3,
,1
3.1
3,
,1
9
181,
,4
181.0
180,
,6
180
.5
180.4
180.
,3
1
3
.1
3.1
3,
,1
3
.1
3.1
3
, 1
9
29
.2
29.7
31
,5
34
.4
29.2
28
.7
6
.7
.6
.7
.8
1.0
1
.2
Sec. 3.2, Appendix A, page 8 R. Newburn, JPL
October 1, 1971
JPL 606-1
Ultraviolet, Visible, and Infrared Photometric Properties
Table A-1. Martian albedo (continued)
WAVELENGTH
.64
,70
.59
1.04 1.24 1.61 1. 74 2.14 2. 27
NUMBER
72
NUMBER
73
NUMBER
74
NUMBER
75
NUMBER
76
NUMBER
77
NUMBER
78
NUMBER
79
NUKBER
80
L A T : - U c
DuAT ITUCE
LONGITODS
DLO\uITUDE
A1,8ED0
DALb^rDO
lat:tuce
DLAT i TUDt
LCNGiTUCt
D'-0,\GITUDE
ALBEDO
DALStDO
LATITUDE
DLATITUDE
LONGITUDE
DLONGITUDE
ALBEDO
DALBEDO
LATITUDE
DUATITUDE
LONGITUDE
DLONGITUDE
AU8ED0
DALBEDO
LATITUDE
DLATITUOE
LONGITUDE
DLONGITUDE
ALBEDO
DALBEDO
LATITUDE
DLATITUDE
LONGITUDE
DLONGITUDE
ALBEDO
DALBEDO
LATITUDE
DLATITUDE
LONGITUDE
DLONGITUDE
ALBEDO
DALBEDO
LATITUDE
DLATITUDE
LONGITUDE
DLONGITUDE
ALBEDO
DALBEDO
LATITUDE
DLATITUDE
LONGITUDE
DLONGITUDE
ALBEDO
DALBEDO
•26.
3 .
L61.
5 .
12,
11.
1.
-25.
3.
166 .
3.
12.
1.
23.
3.
155,
3.
19,
1.
15,
3,
190,
2.
18.
5.0
3.0
195.0
3.0
18.2
.5
5
2
204
3
17
-17,
3.
196,
3,
11,
-25
3
l3 G
o
14
1
-34
3
193.2 197
4.3 4
,3
.3
13
1
-.8
2.0
215.1
3.3
15.5
.9
-25
3
166
3
14
1
23
3
154
3
25
1
15
3
139
2
23
5
3
194
3
22
5
2
203
3
22
2
214
3
20
1
-16
3
196
3
14
-24,
3.
179,
3.
16.
1 ,
-33,
3.
196,
4 ,
14,
-23.
o ,
176,
3,
10,
1.4
- 3 1 . !3
3.0
195 .4
4.0
14.1
1.2
-22,
3,
176.
3,
17
1.
-30 .
3.
194 ,
4,
14,
5 -21.
3,
173,
3,
17,
1 ,
-30,
3,
194.
3,
14,
-21.4
3.0
17 7.7 177,
3.5 3,
19.0
1.2
-21,
3,
20
1
-29,
3.
193,
3,
15,
1.2
1.2
1.2
-29.4
3.0
193.7
3.9
17.0
1.4
•21.
3,
177,
3,
16.
1 ,
-29.
3.
193.
3,
13,
-21.
3,
177.
3,
16,
1 ,
1.1
-29.2
3.
193.5
3.9
14.6
1.2
1 -24.2 -22.9 -22.3 -21.8 -21.3 -21,2 -21.0 -20.9
Q 3.0 3.0 3.0 3.0 3.0 3.0 3.0 3.0
2 165.7 164.9 164.5 164,1 163.8 163.7 163.6 163.5
9 3.9 3.9 3.9 3.8 3.S 3.8 3.8 3.8
15.5 15.1 15.9 16.3 17.2 18.9 14,7 15.0
1 1.2 1.2 1.3 1.4 1.4 1,6 1,5 1.5
24,
3,
153,
3,
29,
2,
16,
3,
133,
2,
27,
25,
3,
151,
3.
33
2
25,
3,
150.
4,
35
2.
26,
3,
150,
4,
36,
2,
17.6 18.2
3,0 3.0
188,0 137.6 187
2.0 2.0 2
29.6 32.5
.5 .5
18,
3,
32
26,
3,
149,
4,
38.
3,
19,
3.
186,
2,
34,
26
3
149
4
42
3
27.
3.
149,
4,
34.
3,
19.3
3.0
186.3 186,
2.0 2,
37.9
.7
19,
3.
31
1
7.7 8.2 8.8 9.2 9
3.0 3.0 3,0 3.0 3
194.0 193.2 192.8 192.5 192.2 192
3.0 3.0 3.0 3.0 3.0 3
29.1 31.5 31.9 33.6
.4 .4 .4 .4
6.5
3.0
27.3
.5
37
9
3
192
3
30
27,
3.
149,
4,
36,
3,
19.5
3.0
186.6
2.0
33.1
1.3
9,
3,
191,
3
31
1
6.5
7.7 8.3
8.9
9.3 9.5
9.6
9.7
2.0 2.0 2.0 2.0 2.0 2,0 2,0 2.0
6 203,1 202,3 201.9 201.6 201.3 201.2 201.1 201,1
1 3,0 3,0 3,0 3,0 3,0 3.0 3.0 3.0
3 26.4 27.7 29,8- 30,0 31,2 34.1 27.3 28.5
6 .5 .5 .6 . .5 .7 ,6 .9 1.4
,8
2.0
2.7
3.2
3.7
3,8
4,0
4.1
2.0 2.0 2.0 2.0 2,0 2,0 2,0 2,0
6 214.1 2l3,3 212,9 212,6 212,3 212,2 212,1 212.1
3 3,2 3,2 3.2 3.2 3.2 3.2 3.2 3.2
5 24,3 25.4 27,5 27,6 28,5 31,3 ^5.5 27.0
.9 ,9 1.0 ,91,0 1,C 1,0 1.3
2 -15.3 -14.0 -13.3 -12.8 -12.3 -12.1 -il.9 -n,9
3.0 3.0
195,9 195,1
3,0 3,0
17.2 17.3
.8 .9
3.0 3.0
194,8 194.4
3,0
18.6
1.3
3.0
19.4
1.6
3.0 3.0
194.2 194.1
3.0 3.0
21.0 22,7
1.7 2,1
3.0 3.0
194,0 193.9
3.0 3.0
18,5 20,3
1.7 2.2
October 1, 1971
R. Newburn, JPL Sec. 3. 2, Appendix A, page 9
Ultraviolet, Visible, and Infrared Photometric Properties
JPL 606-1
Table A-1. Martian albedo (continued)
WAVELENGTH
.60
.64
,70
.09
1.04 1.24 1.61 1.74 2.14 2.27
NUMBER
LATITUDE
-21.9
'21,0
-20 .0
-18.6
-17.9
-17.3
-16.6
-16.6
-16.4
-16.3
61
DLATPUDE
3.
3,0
3,0
3.0
3.0
3.0
3.
3,0
3.
3.0
LONGITUDE
208.4
207.3
207.2
206.4
205.9
205.6
205.3
205,1
205.0
205.0
DLOK-GITUDE
3.6
3.6
3.5
3.5
3.4
3.4
3.4
3,4
3.4
3.4
ALBEDO
11.1
13,0
14.5
13.7
14.4
14.7
15. p
16,8
14.3
14.3
DAI ■ii:[ T
.8
,'3
.9
.8
. 7
.8
.9
.9
,6
.8
NUMBER
LATITUDE
29,9
30.6
31.5
32.9
33.5
34.1
34.6
34.8
34.9
35.0
62
DLATITUDE
3.2
3.2
3.2
3.3
3.3
3.3
3.3
3.3
3.3
3.3
LONGITUDE
185.6
185.3
184.6
183.5
183.0
182.5
182.1
182,0
181.8
181.8
DLONGITUDE
3.3
3,3
3.3
3.4
3.4
3.4
3.4
3,4
3.4
3.4
ALBEDO
19.1
24.5
29.7
31.9
34.5
35,1
37.0
40 ,
33.0
35.2
DALBEDO
1.0
1.0
1.2
1.5
1.6
1.7
1.8
2.1
1.7
1.9
NUMBER
LATITUDE
44.9
45,8
46.9
48.5
49.3
50.0
50 .6
50.8
51.0
51,1
85
DLATITUDE
2.0
2.0
2.0
2.0
2.0
2.0
2.0
2.0
2.0
2.0
LONGITUDE
185.7
165.0
184.0
182.5
181.7
181.0
180 .3
180,1
179.9
179.8
DLONGITUDE
2.0
2.0
2.0
2.0
2.0
2.0
2.0
2.0
2.0
2.0
ALBEDO
16.5
21.1
25.4
25.9
27.9
28,1
29.7
32.6
27.4
27.2
DALBEDO
1.0
1.3
1.5
1.8
2.0
2.1
2.4
2,7
2.3
2.6
NUMBER
LATITUDE
'20.0
20 .8
21.7
23.1
23.8
24 .4
24 .9
25,0
25.2
25.3
66
DLATITUDE
2.0
2.0
2.0
2.0
2.0
■2.0
2,0
2,0
2.0
2.0
LONGITUDE
196.0
195.6
195.2
194.4
194.1
193.7
193.4
193.4
193.3
193.2
DLONGITUDE
2.0
2,0
2.0
2.0
2.0
2.0
2.0
2,0
2,0
2.0
ALBEDO
15.6
20.4
24.1
25.7
28.1
28.7
30 .6
33.3
27,9
28,6
DALBEDO
.6
.8
.9
1 .2
1 .4
1.3
1.3
1.5
1.2
1.3
NUMBER
LATITUDE
42.0
42.9
44.1
45,8
46.6
47.4
48.0
48,3
48.5
48.6
87
DLATITUDE
3.6
3.6
3.6
3,7
3.8
3.8
3.8
3,8
3.9
3.9
LONGITUDE
204.2
203.8
203.2
202,4
202.0
201.6
201.3
201.2
201.0
201.0
DLONGITUDE
3.6
3.6
3.6
3.7
3.8
3,8
3.8
3,9
3.9
3,9
ALBEDO
18.3
23,5
27.7
28,9
30.7
31,0
32.5
35.5
29,0
30,2
DALBEDO
1.2
1.5
1.9
2.2
2.5
2.7
2.9
3.3
2.7
3.6
NUMBER
LATITUDE
26.0
26.8
27.8
29.3
30,1
30.8
31.3
31,5
31.7
31,8
89
DLATITUDE
2.0
2.0
2.0
2,0
2.0
2.0
2,0
2,0
2.0
2.0
LONGITUDE
210.1
209.7
209.4
208.8
208.5
208.3
208.1
208,0
207.9
207,9
DLONGITUDE
2.0
2.0
2.0
2,0
2.0
2.0
2.0
2,0
2.0
2.0
ALBEDO
19.1
24,0
29.3
31.3
33,5
33.7
35.1
37.5
31,8
33.4
DALBEDO
.5
.6
.7
.8
.9
1.0
1.0
1.0
1.4
1,5
NUMBER
LATITUDE
22.8
23,6
24.6
26,1
26.8
27.4
28.0
28.2
23.3
28.4
90
DLATITUDE
3,1
3.1
3.1
3.1
3.1
3.2
3.2
3,2
3,2
3.2
LONGITUDE
184.9
184.4
183.8
182.9
182.4
182.0
181.6
181.5
181.4
181.3
DLONGITUDE
3.3
3,3
3.4
3.4
3.4
3,5
3.5
3,5
3,5
3.5
ALBEDO
19.4
24.1
29.1
32.1
34.9
35.0
36.4
39,3
34.2
36,2
DALBEDO
,8
.9
1.2
1.4
1.5
1.6
1.7
2,0
1.9
2.0
NUMBER
LATITUDE
19.9
20.8
21.8
23.3
24.0
24,7
25,3
25,4
25.6
25.7
91
DLATITUDE
2.0
2.0
2.0
2.0
2.0
2.0
2.0
2.
2.0
2.0
LONGITUDE
212.1
211,8
211.5
211.0
210.7
210.5
210.3
210.3
210.2
210.2
DLONGITUDE
2.0
2,0
2.0
2,0
2.0
2.0
2.0
2.0
2.0
2.0
ALBEDO
19.1
23.4
28.7
30.9
33.5
33.6
34.8
36.8
31.9
34.2
DALBEDO
.7
.5
.5
.6
.7
.8
.9
,9
.8
1.0
NUMBER
LATITUDE
32.0
32,9
34.0
35.6
36.4
37.2
37.6
38,0
38.2
36.3
92
DLATITUDE
2.0
2,0
2.0
2,0
2.0
2.0
2.0
2,0
2.0
2.0
LONGITUDE
208.3
208,0
207.6
207.0
206.7
206.4
206.1
206,1
206,0
205.9
DLONGITUDE
2.0
2.0
2.0
2.0
2.0
2.0
2.0
2.0
2.
2 .
ALBEDO
18.9
24.3
29.2
31.3
33.6
33.6
35.1
37.2
32.1
33.9
DALBEDO
.6
.9
1.0
1.1
1.2
1.3
1.3
1.5
1.6
1.6
Sec. 3. 2, Appendix A, page 10 R. Newburn, JPL
October 1, 1971
JPL 606-1
Ultraviolet, Visible, and Infrared Photometric Properties
Table A-1, Martian albedo (continued)
WAVELENGTH
NUMBER
93
NUMBER
94
NUMBER
95
NUMBER
96
NUMBER
97
NUMBER
98
NUMBER
99
NUMBER
100
NUMBER
LATITUDE
PLATITUDE
LONGITUDE
DLONGITUDE
ALBEDO
DALBEDO
LATITUDE
DLATITUDE
LONGITUDE
DLONGITUDE
ALBEDO
DALBEDO
LATITUDE
DLATITUDE
LONGITUDE
DLONGITUDE
ALSEDO
DALBEDO
LATITUDE
DLATITUDE
LONGITUDE
DLONGITUDE
ALBEDO
DALBEDO
LATITUDE
DLATITUDE
LONGITUDE
DLONGITUDE
ALBEDO
DALBEDO
LATITUDE
DLATITUDE
LONGITUDE
DLONGITUDE
ALBEDO
DALBEDO
LATITUDE
DLATITUDE
LONGITUDE
DLONGITUDE
ALBEDO
DALBEDO
LATITUDE
DLATITUDE
LONGITUDE
DLONGITUDE
ALBEDO
DALBEDO
LATITUDE
DLATITUDE
LONGITUDE
DLONGITUDE
ALBEDO
DALBEDO
.60
25.0
2.
.64
25.9
2.0
.70
26.9
2.0
1.04 1.24 1.61 1.74 2.14 2.27
203.0 202.7 202.2
2,0
2.0
2.0
28.5 29.3 30.0 30 .6
2.0 2.0 2,0 2.U
201.6 201.3 201.0 200.7
2.0 2.0 2.0 2.0
30 ,8
2.0
200 .7
2 ,
31 .
2.
200.6
2.
18.7 23,6 23.5 31.2 33.9 34.0 Si,
31.1
2.
2 0.5
2.
37.5 62 .6 33.8
.5
.7
.6
.9
.9
1 .
1.2
1.3
27,0 27,9 29. 3U ,
31.4 32.1 32.7 32.9 33.1 33.2
2.0 2.0 2.0 2.0 2.0 2.0 2.0 2.0 2.0 2.0
217.1 216.8 216.5 216.1 215.9 215.7 215.6 215.5 215.4 215.4
2.0 2,0 2.0 2.0 2.0 2.0 2.0 2.0 2.0 2.0
19.6 24,8 29.9 32.0 34.0 33.8 35.0 37.0 31.4 32.6
.6 .8 .8 1.0 1.0 1,1 1.2 1.3 1.1 1,6
-19.9 -18,9 -17.7 -16.0 -15.2 -14.4 -13.8 -13.6 -li,4 -13.3
3.0 3,0 3.0 3.0 3.0 3.0 3.0 3.0 3.0 3.0
219.3 218.9 218.6 216.0 217.7 217.5 217.3 217.3 217.2 217.2
3.5 3.5 3.4 3.4 3.3 3.3 3.3 3,3 3.3 3.3
11,4 12.8 14.3 13,2 13,5 13.4 14.0 14,6 11,6 12.9
.6 ,6 ,8 .6 .6 .6 .7 ,7 .7 .9
6.1
7.0
8.2
9.8 10.6 11.3 11.9 12.1 12.3 12.4
2.0 2,0 2,0 2.0 2.0 2.0 2.0 2.0 2.0 2.0
225.0 224,8 224.6 224.2 224.1 223.9 223.8 223.7 223.7 223.7
3.1 3,1 3.1 3.1 3.1 3.1 3.1 3.1 3.1 3.1
16,7 20,2 23.8 25.3 27.0 26.8 27.5 29.3 25,1 25,5
.6 ,7 .8 1.0 1.0 1.0 1.2 1.2 1.1 .9
14
3
231
3
14
15.8
3,0
17.0
3,0
18.6
3.0
19.5
3,0
20,
3.
20 .8
3.x
21.
3.1
21.3
3.1
21.4
3.1
230.9 230.7 230.5 230.4 230.3 230.2 230.1 230.1 230.1
3.2
17.1
.7
3.2
19.2
.7
3.2
.8
3.2
19.1
.7
3.
18,
2
1
. 7
3.2
17.5
3.2
l6,5
3.2
15.1
1.0
3. 2
14.9
1.3
33,1 34.1 35.4 37.2 38,2 39.0 39.7 40.0 40.2 40.3
3.3 3,3 3.3 3.4 3.4 3.4 3.5 3,5 3.5 3 5
231.9 231.9 231.8 231.8 231.8 231.6 231.6 231.8 231.8 231.6
3.4 3.5 3.5 3,6 3,6 3.6 ,3.6 3.6 3.7 3 7
15.5 17.6 20,5 19,2 19 . 6 18.5 17.5 18.6 14.5 14.0
.9 1,2 1.2 1.3 1.3 1.3 1.3 1.4 1.4 1.9
49.9 51.2 52.7 55.1 56.3 57.4 58.4 58.7 59.0 59.2
3.9 4.0 4.1 4,3 4.4 4.5 4.6 4.6 4.6 4.6
227,4 227.4 227,5 227,6 227,7 227.8 227.9 227.9 227.9 228.0
4,0 4.1 4.2 4.4 '4,5 4,6 4,7 4,7 4.7 4.8
15.9 18.3 20.5 19.8 20.1 19.3 18.8 19,8 17,2 17,1
1.4
1.8
2.3
2,4
2.6'
2.7
2.1
3.1
2.9
3.3
26.0 27.0 23,2 30,0 31.0 31.8 52.5 32.7 32.9 33.0
2,0 2.0 2.0 2.0 2.0 2,0 2.0 2,0 2.0 2.0
210.0 209.7 209.4 208.9 208.7 208.4 206.2 208,2 20,8.1 206 1
2.0 2,0 2,0 2,0 2.0 2.0 2.0 2.0 2.0 2,0
19.1 24.0 29.4 32,3 34.4 34.6 35,7 37.9 32.9 33.0
.6 ,5 ,8 ,8 1,0 1.0 1.1 1.3 1.6 1.5
9,0 10.0 11,1 12,8 13,7 14.5 15.1 15,3 15.5 15.6
3.0
3.0
3.0
3.0
3.0
3.0
3.0
3.0
3.0
3.C
191.0 190.7 190.4 189.9 189.6 189.3 169.1 189.0 189.0 1S3.9
3.3 3.3
18,3 21.7
.7 .9
3,
27,
1,
3.4
30. 2
1.2
3. 4
31.9
1.3
3.4
32.5
1.3
3.4
33.7
1.4
3. 4
36 . 2
1.5
01 .
1.
3,
3?.
October 1, 1971
R. Newburn, JPL Sec. 3.2, Appendix A, page 11
Ultraviolet, Visible, and Infrared Photometric Properties
JPL 606-1
Table A-1. Martian albedo (continued)
ftAVtieiVCTH
.6u
.64
.70
.89 1.04 1,24 1.61 1.74 2.14 2.2?
NUMBER
LATITUDE
-18.0
-16.9
-15.5
-13.6
-12.6
-11.8
-11.1
-10.9
-10.6
-10.5
102
OUATiTUDE
3.n
3.0
3.0
3.
3.0
3.0
3.0
3,0
3.0
3.0
LONGITUDE
227,0
22 6,7
226.5
226,1
225.9
225.7
225.6
225,6
225,5
225.5
DUONGITUDt
3.4
3.4
3.3
3 , 3
3.3
3.3
3.2
3.2
3,2
3.2
ALBEDO
11.2
12.0
14.3
13.7
13.2
13.3
13.6
14.6
11.2
11.5
DALBEDO
.6
.6
.6
.5
.5
.5
.5
.5
,7
1.1
NUMBER
LATITUDE
14.9
15.5
16.2
17.2
17.7
18.2
18.6
10.7
18.8
18.9
111
DLATITUDE
3.0
3.0
3.0
3.0
3.0
3.0
3.0
3.0
3.0
3.
LOMGITUDE
110.0
109,2
108.2
106.8
106.1
105.5
105,0
104.3
104,6
104.5
DLONGITUDE
3.5
3.4
3 .4
3.4
3.4
3.3
3.3
3,3
3.3
3.3
ALBEDO
16.5
21.4
25.9
29.1
31.4
32.0
33.7
34,8
31.1
31.8
DALBEDO
1.2
1.5
1.3
1.9
2.1
2.1
2.2
2.3
2.1
2.3
NUMBER
LATITUDE
12.0
12,5
13.2
14.1
14.5
14.9
15.3
15.4
15.5
15.6
112
DUATITUUE
3.0
3.0
3.0
3.0
3.0
3.0
3.0
3,0
3.0
3.0
LONGITUDE
83.0
82.3
81.4
80. 1
79.4
78.8
78,3
73,1
78.0
77.9
DLONGITUDE
3,0
3.0
3.0
3.0
3.0
3.0
3.0
3,0
3.0
3.0
ALBEDO
16.2
21,2
25.4
27.7
29.7
29,9
31.4
32,1
28.1
30.3
DALBEDO
.4
.4
.6
.6
.8
.7
.8
1.0
1.7
1.1
NUMBER
LATITUDE
11.1
U.7
12.5
13.6
14.1
14.6
15.0
15.2
15.3
15.4
120
DLATITUDE
3.0
3.0
3.0
3.0
3.0
3.0
3.0
3.0
3.0
3.
LONGITUDE
116.0
117.4
116.6
115.5
115.0
114.5
114.1
114,0
113.8
113.8
DLONGITUDE
3.4
3.4
3.4
3.3
3.3
3.3
3.3
3.3
3.3
3.3
ALBEDO
17.8
23.7
28.4
31.2
33.5
34,2
35.8
37.7
31.9
33.4
DALBEDO
1.1
l.-?
1,7
1.8
1.9
2.0
2,0
2.2
2.1
2.4
NUMBER
LATITUDE
12.0
12.6
13.3
14.3
14.8
15.3
15,7
15,8
15.9
16.0
121
DLATITUDE
3.0
3.0
3.0
3.0
3.0
3.0
3,0
3.0
3.0
3.0
LONGITUDE
91.0
90.4
89.7
38.6
88.0
87.6
87,1
87.0
86.9
86.8
DLONGITUDE
3.0
3,0
3.0
3.0
3.0
3.0
3.0
3,0
3.0
3.0
ALBEDO
17.3
23.2
28.1
31.0
33.0
33.6
34.6
36.6
31.4
33.3
DALBEDO
.5
,4
,5
.5
.6
.6
.7
.7
.8
1.4
NUMBER
LATITUDE
12.0
12.6
13.3
14.3
14.8
15.2
15.6
15.7
15.8
15.9
122
DLATITUDE
3.0
3.0
3.0
3.0
3.0
3.0
3.0
3,0
3.0
3.0
LONGITUDE
83.0
62.4
81.6
80.5
79.9
79.4
79.0
78,9
78.7
78.7
DLONGITUDE
3.0
3.0
3.1
3.1
3.1
3.1
3.1
3.1
3,1
3.1
ALBEDO
16.2
21.0
25.9
28.6
30.9
31,4
32.9
34.6
29.5
31.8
DALBEDO
.3
.5
.5
.6
.6
.7
.8
,9
1.0
1.1
NUMBER
LATITUDE
39.0
39.7
40.6
41.8
42.4
43.0
43.5
43.6
43.8
43.9
125
DLATITUDE
3.4
3.4
3.5
3.5
3.5
3.5
3.6
3.6
3.6
3.6
LONGITUDE
92.1
91.4
90.5
69.2
88.5
87.8
87,3
87,1
86.9
86.8
DLONGITUDE
3.4
3.4
3.5
3.5
3.5
3.6
3.6
3.6
3.6
3.6
ALBEDO
17.6
23,8
28.3
30.7
32.6
33.0
34.1
36,7
30.7
31.6
DALBEDO
1.1
1.1
1.7
2.1
2.3
2.5
2.7
3.0
2.6
2.8
NUMBER
LATITUDE
28.0
28.6
29.4
30.6
31.2
31.7
32.1
32.3
32.4
32.5
126
DLATITUDE
3.1
3.2
3,2
3.2
3.2
3.2
3.2
3.2
3.2
3.2
LONGITUDE
93.9
93.3
92.5
91.4
90.9
90.3
89,9
89.8
89.6
89.5
DLONGITUDE
3.1
3.2
3.2
3.2
3.2
3.2
3.2
3.2
3.3
3.3
ALBEDO
17.7
23.1
28.2
31.0
33,6
34.1
35.5
38.7
32.4
35,4
DALBEDO
.8
.9
1.1
1.2
1.4
1.5
1.6
1.9
1.7
1.8
NUMBER
LATITUDE
15.0
15.6
16.4
17.5
18.1
18.6
19,0
19.1
19.3
19,3
127
DLATITUDE
3.0
3.0
3.0
3.0
3.0
3.0
3,0
3.0
3.0
3.0
LONGITUDE
96.0
95.4
94.8
93.8
93.4
92.9
92.6
92.4
92.3
92,5
DLONGITUDE
3.0
3.0
3.0
3.0
3.0
3.0
3.0
3.0
3.0
3.0
ALBEDO
18.5
23,5
29.0
31.6
34.2
34.9
36,1
39,0
33.8
34.1
DALBEDO
.5
.4
.6
.6
.7
.7
,8
.9
.9
1.2
Sec. 3. Z, Appendix A, page IZ R. Newburn, JPL
October 1, 1971
JPL 606-1
Ultraviolet, Visible, and Infrared Photometric Properties
Table A-1, Martian albedo (continued)
WAVELENGTH
(
.60
.64
.70
.89
1. 04
1.24
1.61
1.74
2.14
2.27
NUMBER
LATITUDE
3.0
3.7
4.5
5.6
6.2
6.7
7.1
7,3
7.4
7.5
128
DLATITUDE
3.0
3.0
3.0
3.0
3.0
3.0
3.0
3.0
3.0
3,0
LONGITUDE
96.9
96.5
95.9
95.0
94.6
94.2
93.9
93.8
93,6
93.6
DLONGITUDE
3.0
3,0
3.0
3.0
3.0
3.0
3.0
3.
3.
3.0
ALBEDO
18.0
23.4
29.2
31.5
33.8
34.3
35.4
37,8
31.7
34 .4
DALBEDO
.5
.4
.4
.6
.5
.5
.6
.6
,6
,6
NUMBER
LATITUDE
-10 .0
-9,3
-8.4
-7.2
-6.6
-6.1
-5.6
-5.5
-5.3
-5.3
129
DLATITUDE
3.2
^.2
3.2
3.2
3.1
3.1
3.1
3.1
3.1
3.1
LONGITUDE
96. V
96,4
95.8
95.0
94.6
94,2
93.9
93.8
93.7
93,6
DUONG iVUDE
3.2
3.2
3.2
3.2
3.2
3.2
3.2
3,1
3.1
3.1
ALBEDO
17.6
23,1
27.9
30.6
33.1
33.4
34.8
37.2
31.7
31.8
DALBEDO
.6
.7
.8
.8
.9
.9
,9
1.0
1.1
1.3
NUMBER
LATITUDE
-21.0
-20.2
-19.3
-13.0
-17.3
-16.7
-16.2
-16.1
-15.9
-15. a
130
DLATITUDE
3.5
3.5
3.5
3.4
3.4
3.4
3.4
3.4
3.4
3.4
LONGITUDE
96.1
95.6
95.0
94.2
93.8
93.5
93.2
93.1
93.0
92.9
DLONGITUDE
3.5
3,5
3.5
3.5
3.4
3.4
3.4
.4
3.4
3.4
ALBEDO
15.8
21.0
25.4
27.1
29.3
29.5
30 .8
3o,3
28.8
28.0
DALBEDO
.9
1.1
1.3
1.3
1.4
1.4
1.4
1.6
1.4
1.6
NUMBER
LATITUDE
34.0
34,7
35,6
36.9
37.6
38.2
38.7
38,9
39.0
39.1
132
DLATITUDE
3.3
3,3
3.3
3.3
3.4
3.4
3.4
3.4
3.4
3.4
LONGITUDE
99.7
99,1
98.4
97.4
96.8
96.3
95.9
95.8
95.6
95.5
DLONGITUDE
3.3
3.3
3.3
3.4
3.4
3.4
3.4
3.4
3.4
3.4
ALBEDO
19.5
24,4
29.8
33.5
35.4
36.3
38.1
41.5
34.6
35,4
DALBEDO
.9
1.2
1.5
1.9
2.0
2.1
2.4
2.7
2.3
2.8
NUMBER
LATITUDE
3.0
3.7
4.6
5.8
6.4
7.0
7.4
7.6
7.7
7.8
133
DLATITUDE
3.0
3.0
3.0
3.0
3.0
3.0
3.0
3.0
3.0
3,
LONGITUDE
98.0
97,6
97.1
96.3
95.9
95.6
95.3
95.2
95.1
95.1
DLONGITUDE
3.1
3.1
3.0
3.0
3.0
3.1
3.1
3.1
3.1
3.1
ALBEDO
18.7
24,2
29.7
32.1
34,4
34.7
36.0
38.7
33.4
33.6
DALBEDO
.5
.6
.5
.6
.5
.5
.6
.7
.7
.9
NUMBER
LATITUDE
35.9
36.7
37.7
39.2
39.9
40.6
41.1
41,3
41.5
41 .6
135
DLATITUDE
3.3
3.3
3.4
3.4
3.4
3.5
3.5
3.5
3.5
3.5
LONGITUDE
120.3
119.9
119.5
116.8
118.5
118.2
117.9
117.8
117.7
117.7
DLONGITUDE
3.4
3.4
3.4
3.5
3.5
3.5
3.5
3.5
3.5
3.5
ALBEDO
19.3
24.6
30.2
32.0
34.6
34.7
36.2
39.3
32.7
34.7
DALBEDO
1.3
1.6
2.0
2.3
2.5
2,6
2.9
3.1
2.9
3.0
NUMBER
LATITUDE
25.0
25.7
26.7
28.0
28.7
29.3
29.9
30.0
30.2
30 .3
136
DLATITUDE
3.1
3.1
3.1
3,1
3.2
3.2
3.2
3.2
3.2
3.2
LONGITUDE
117.1
116.7
116.2
115.6
115.2
114.9
114.7
114,6
114.5
114.4
DLONGITUDE
3.1
3.1
3.1
3.2
3.2
3.2
3.2
3.2
3.2
3.2
ALBEDO
. 18.7
24,1
29.3
31.6
34.1.
34.4
36.1
39,5
31.9
34.1
DALBEDO
.6
.9
1.2
1.3
1.5
1.5
1.6
1.8
1.8
1.7
NUMBER
LATITUDE
10.1
10.8
11.7
13.1
13.7
14,3
14.8
i5;6
15.2
15.2
137
DLATITUDE
3.0
3.0
3.0
3.0
3.0
3.0
3.0
3.0
3.0
3.0
LONGITUDE
119.0
118.6
118.2
117.6
117.2
117.0
116.7
116,6
116.6
11.6.5
DLONGITUDE
3.0
3.0
3.0
3.0
3.0
3.0
3.0
3.0
3.0
3.0
ALBEDO
18.8
24,3
29,2
31.6
34.2
34.3
35,8
38,6
31.4
33.6
DALBEDO
.6
.9
.9
.9
1.0
1.0
1.0
1.2
1.3
1.2
NUMBER
LATITUDE
.0
.8
1.8
3.1
3.3
4.4
4.9
5.1
5.2
5.3
138
DLATITUDE
3.1
3.0
3.0
3.0
3.0
3.0
3.0
3.0
o.O
o.
LONGITUDE
120.0
119.7
119.2
118.6
lie. 2
118.0
117.7
117,6
117.6
117.5
DLONGITUDE
3.1
3.1
3.1
3.1
3.1
3.1
3.0
3.0
3.
3.0
ALBEDO
13.8
24,0
28.8
30 .9
33.2
33.0
34,0
38.1
32.2
33,9
DALBEDO
.7
.7
,9
.8
.8
.9
.9
1.0
1.2
1.6
October 1, 1971
R. Newburn, JPL Sec. 3.2, Appendix A, page 13
Ultraviolet, Visible, and Infrared Photometric Properties
JPL 606-1
Table A-1. Martian albedo (continued)
WAVELENGTH
.60
.6 4
.7a
.89 1.04 1.24 1.61 1.74 2.14
:.27
MUM3ER
LATITUDE
-CO
-7.2
-6.2
-4.8
-4,1
-3.4
-2.9
-2,7
-2.6
-2.5
139
DLATITUDE
3.2
3.2
3.1
3.1
3,1
3.1
3.1
3.1
3.1
3.1
LONGITUDE
12D.0
119.6
119.2
118.5
118.2
117.9
117.7
117.6
117.5
117.5
DLONGITUDE
3.2
3.2
3.2
3.1
3,1
3.1
3.1
3.1
3.1
3.1
ALBEDO
18.3
22.7
27.9
30.1
32.1
32.5
33.9
36.7
31.1
31.4
DAL6ED0
.8
' .3
1.0
.9
1.0
1.0
1.0
1,1
1.1
1.1
NUMBER
LATITUDE
-24.0
-23.0
-21.9
-20 .2
-19.4
-18.7
-18,1
-17,9
-17.7
-17.6
140
DLATITUDE
3.7
3.6
3.6
3.5
3,5
3.5
3.4
3.4
3.4
3.4
LONGITUDE
121.1
120.7
120.1
119.4
119.0
118.7
113.4
118.3
118.2
118.2
DLONGITUDE
3.7
3.7
3.6
3.5
3.5
3.5
3.5
3.4
3.4
3.4
ALSiDO
17.5
22.3
26.9
29.1
31.2
31.6
33,1
36.2
30.0
30.5
DAlBEDO
1.3
1.6
1.8
1.8
1.8
1.9
1.9
2.1
1.8
1.9
NUMBER
LATITUDE
10.0
10.8
11.7
13.1
13.8
14 . 4
14.9
15.1
15.3
15.3
141
PLATITUDE
3.0
3.0
3.0
3.0
3.0
3.0
3.0
3.0
3.0
3.0
LONGITUDE
111.0
110.6
110.2
109.6
109.3
109.0
103.7
108.7
108.6
106.5
DLONGITUDE
3.0
3.0
3.0
3.0
3,0
3.0
3.0
3,0
3.0
3.0
ALBEDO
19.1
24.3
30.1
32,5
35.0
35.6
37.0
40,2
34,3
35.7
DALBEDO
.4
.5
.5
,5
.5
.6
.6
.8
,8
1.3
NUMBER
LATITUDE
-10.1
-9.2
-8.2
-6.8
-6,1
-5.5
-4.9
-4.7
-4,6
-4.5
142
DLATITUDE
3.2
3.2
3.2
3.1
3.1
3..1
3.1
3,1
3.1
3.1
LONGITUDE
94.1
93.8
93,4
92.8
92.5
92.3
92.1
92,0
91.9
91.9
DLONGITUDE
3.4
3.4
3.4
3.3
3.3
3.3
3.3
3,3
3.3
3.3
ALBEDO
18.0
23.4
27.7
30.3
32.4
33.1
35.3
37,3
32.0
32.8
DALBEDO
•8
1.1
1.1
1.2
1,3
1,2
1.3
1,5
1.6
1.6
NUMBER
LATITUDE
35.9
36.7
37.8
39.3
40.1
40 .7
41.3
41,5
41.7
41.8
143
DLATITUDE
3.3
3.3
3.4
3.4
3.4
3,5
3,5
3.5
3.5
3.5
LONGITUDE
89.8
89.1
88,2
86.8
86.1
85,4
84.8
84,6
84.4
84.3
DUONGITUDE
3.6
3.7
3,7
3.8
3.9
3.9
4.0
4.0
4,0
4.0
ALBEDO
19,0
25.5
30.7
33.6
35.5
36.4
39.1
42.5
35,0
36.9
DALBEDO
1.2
1.8
2.2
2.7
2.9
3.1
3.5
4.0
3.4
3.7
NUMBER
LATITUDE
12.8
13.5
14.5
15.9
16.6
17.2
17.7
17,9
18.1
18,1
144
DLATITUDE
3.0
3.0
3.0
3.0
3.0
3.0
3.0
3,0
3.0
3.0
LONGITUDE
91.0
90.6
90,0
89.2
88.8
88.4
88.1
88,0
87.9
87.8
DLONGITUDE
3.3
3.3
3,3
3.3
3.4
3.4
3.4
3,4
3.4
3.4
ALBEDO
19.6
24.9
29.9
32.6
35.0
35.5
36.7
39,9
33.5
35.0
DALBEDO
.8
.9
1.1
1,3
1.5
1,5
1.6
1,8
1.7
1.8
NUMBER
LATITUDE
6.1
6.9
8.0
9.5
10,2
10,9
11.5
11,7
11,9
12,0
145
DLATITUDE
3.0
3.0
3.0
3.0
3,0
3.0
3.0
3,0
3,0
3.0
LONGITUDE
112.9
112.7
112.3
111.8
111,5
111,3
111.1
111,0
111,0
110.9
DLONGITUDE
3.0
3.0
3,0
3,0
■ 3.0
3,0
3.0
3,0
3,0
3.0
ALBEDO
18.5
23.8
29.4
32,3
34.9.
. 35.3
36.9
39,5
34.
34. 8
DALBEDO
,5
.4
,5
,5
.5
.6
.5
,7
.8
.9
NUMBER
LATITUDE
24.8
27.7
28.9
30,6
31.4
32.2
32.8
33,1
33.3
33.4
147
DUATITUDE
3.1
3.1
3,2
3.2
3.2
3.2
3.2
3,2
3.2
3.2
LONGITUDE
147.1
147.1
147.0
146.9
146.9
146.9
146.9
146,9
146.9
146.9
DLONGITUDE
3.5
3.5
3.5
3,6
3,6
3.6
3.6
3,6
3.6
3.6
ALBEDO
19.1
23.6
28.3
29.7
31.8
31.7
33.4
36,3
51.1
33.0
DALBEDO
1.3
1.7
2,0
2,3
2,5
2.6
2,8
3.2
2.9
3.0
NUMBER
LATITUDE
9.9
10.9
12.0
13,6
14.5
15.2
15.8
16.0
16.2
16.5
148
DLATITUDE
3.0
3.0
3.0
3.0
3.0
3,0
3.0
3.0
3.0
3.0
LONGITUDE
146.0
145.8
145,6
145.3
145.1
145,0
144.9
144 ,9
144,8
144 .8
DLONGITUDE
3.3
3.3
3,3
3,3
3.3
3,3
3.3
3,3
3.3
5.5
ALBEDO
19.4
24.2
28.9
30 ,9
33.6
33,6
35.1
38,0
33.0
35.4
DALBEDO
1.1
1.4
1,5
1,7
1.8
1,9
2.0
2.2
2,1
2.5
Sec. 3.2, Appendix A, page 14 R. Newburn, JPL
October 1, 1971
JPL 606-1
Ultraviolet, Visible, and Infrared Photometric Properties
Table A-1. Martian albedo (continued)
WAVELENGTH
.60
.64
.70
.89
1.0 4
1.24
1.6:
1.74
2.14
2.
-■7
NUMBER
LATITUDE
-5.3
-4,8
-3.6
-1.9
-1.0
- . 2
.4
.6
.3
. 9
■149
DLATI TUDE
3.1
3, 1
5. 1
3. '
3.1
3. 1
v^ • J
3. C
f\
•V
LONGITUDE
147.1
I'i^.a
1 4 6 . v
146.0
14^. a
14S.6
145. ■>
1 4 D . 4
14^14
14i
3
DLONGiTUDE
5.4
3,4
3.4
3.4
3.3
3.3
3.3
3.3
3.3
3
3
ALBEDO
19.2
23.8
28.9
31.3
34.3
34.7
36. 3
39.
32.6
34
.6
DALBEDO
1.2
1.5
1.7
1.9
2.0
2.1
2.1
2.5
2.1
2
4
NUMBER
LATITUDE
-21.9
-20.8
-19.4
-17.4
-16.5
-15.6
-14.9
-14.7
-14.4
-1-
. 3
150
DLATITUDE
3.6
3,5
3.5
3.4
3.4
3.4
3.3
3,3
3 . 3
o
.3
LONGITUDE
144.1
143,7
143.2
142.6
142.3
142.0
141 .b
141,7
141.6
141
. 6
DLONGITUDE
3.3
3,8
3.7
3.6
3.6
3.5
5.5
3,5
3 . 5
3
.5
ALBEDO
18.8
■2 4.0
28.3
31.2
33.6
33.9
35.2
5 7.8
33.
33
. 4
DALGEDO
1.7
1.9
2.1
2.2
2.3
2.2
2.4
2.5
2.5
2
6
NUMBER
LATITUDE
-15.0
-13.9
-12.6
-10.7
-9.7
-3.9
-8.2
-a.o
-7.8
-7
7
151
DLATITUDE
3.3
3.3
3.3
3.2
3.2
3,2
3.2
3.2
3.2
3
2
LONGITUDE
129.3
129.0
128.8
128.4
128.2
128.
127.9
127.9
127.3
127
8
DLONGITUDE
3.3
3,3
3.3
3.2
3.2
3.2
3.2
3.2
3.2
3
2
ALBEDO
17.7
22.6
27.6
30,3
32.7
32.7
33.9
35.6
31.1
51
2
DALBEDO
1.1
1.0
1.2
1.2
1.3
1.2
1.2
1.3
1.4
1
4
NUMBER
LATITUDE
11.0
12.0
13.2
15.0
15.8
16.6
17.3
17,5
17.7
17
8
152
DLATITUDE
3.0
3.0
3.0
3.0
3.0
3.0
3.0
5,0
3.0
3
LONGITUDE
122.0
121.8
121.5
121.2
121,0
120,8
120.7
120,6
120.6
120
5
DLONGITUDE
3,0
3.0
3.0
3.0
3,0
3.0
3.0
3.0
3,0
3
ALBEDO
19.0
24.3
29.6
32.6
35.2
35.5
36.9
39.0
34,0
33
7
DALBEDO
.4
.5
.5
.6
.7
.7
.7
.8
1.0
^
3
NUMBER
LATITUDE
10.1
11.1
12.3
14.1
15.0
■15. S
16.5
16.7
16.9
17
153
DLATITUDE
3.0
3,0
3.0
3.0
3.0
3.0
3.0
3.0
3.0
3
LONGITUDE
136.0
135,9
135.7
135,4
135.3
135.2
135.1
135.0
135.0
135
DLONGITUDE
3.0
3,0
3.0
3.0
3.0
3.0
3.0
3.0
3,1
3
1
ALBEDO
18.4
22.6
27.1
30.2
32.4
32.5
33.4
35.7
31.2
31
9
DALBEDO
.7
,7
.9
1.0
1.0
1,1
1,1
1.1
1.2
1
7
NUMBER
LATITUDE
6.9
7,9
9.1
10.9
11.8
12,6
13.2
13,4
13.7
13
6
154
DLATITUDE
3.0
3,0
3.0
3.0
3.0
3.0
3.0
3.0
3.0
3
LONGITUDE
104.0
103,8
103.5
103.0
102.7
102.5
102,5
102.2
102.2
102
1
DLONGITUDE
3.3
3,3
3.3
3.3
3.3
3.3
3.3
3.3
3.3
3
3
ALBEDO
19.3
24,7
29.5
33.2
35.3
35.8
37.1
38.8
34.6
35
DALBEDO
.8
,9
1.0
1.2
1.3
1.3
1.5
1.5
1,5
1
6
NUMBER
LATITUDE
12.0
12,7
13.5
14.7
15.2
15.8
16.2
16,4
16,5
16
6
155
DLATITUDE
3.0
3,0
3,0
3.0
3.0
3.0
3.0
3.0
3,0
3
LONGITUDE
90.0
89,5
88.9
88,1
87.6
87.:;
86.9
86.8
86,7
6 6
7
DLONGITUDE
3.0
3.0
3.0
3.0
3.0-
3,0
3.0
3.0
3.0
3
ALBEDO
18.3
23.1
27.9
30.2
32.3
32.6
34.0
36.2
30.6
32
4
DALBEDO
.6
.4
.5
.6
.6
.6
.7
.7
.9
1
NUMBER
LATITUDE
10.1
10.8
11.6
12.7
13.3
13.8
14.2
14.4
14.5
14
6
156
DLATITUDE
3.0
3,0
3.0
3.0
3.0
3.0
3.0
3.0
3.0
3
LONGITUDE
77.0
76.5
75.9
74.9
74.5
74.
73.7
73,6
73.4
73
4
DLONGITUDE
3.1
3.1
3.1
3.1
3.1
3.2
3.2
3.2
3.2
3
2
ALBEDO
16.6
21.5
25. a
27.8
29.9
30.2
•31.3
33.7
26.3
3n
DALBEDO
.6
.5
.6
.7
.8
.9
.9
1.0
.9
1
7
NUMBER
LATITUDE
12.0
12.6
13.4
14.5
15.1
15.5
16.0
16.1
16.2
16
3
157
DLATITUDE
3.0
3.0
3.0
3.0
3.0
3.0
3.
3.0
3.0
3
LONGITUDE
62.0
61.4
60.6
59.5
5S.9
53.4
57.9
57.8
5 7.6
57
5
DLONGITUDE
3.5
3.5
3.6
3.6
3.6
3.7
3.7
3.7
3.7
'7
ALBEDO
14.1
13.7
21.4
23.0
24.8
25.3
26.4
26.7
24.3
26
7
DALBEDO
.3
.9
1.1
1.3
1 .4
1.6
1.7
2.0
1.6
2
October 1, 1971
R. Newburn, JPL Sec. 3.2, Appendix A, page 15
Ultraviolet, Visible, and Infrared Photometric Properties
JPL 606-1
Table A-1. Martian albedo (continued)
WAVc;Lc\-GTH
,60
.64
,70
.39 1.04 1,24 1.6:
1.74 2.1.
2.27
,^JUMBER
LATI'UDE
-26.9
-::8.0
-26.9
-25.3
-24,6
-23.9
-23.3
-23.1
-23 , C
-22,9
159
OLA"' ITU i;::
3.9
3,9
5.S
3.7
3,7
3.7
3,.'.
3.6
3.6
3, 6
l.OXG!TuaE
100.4
5 9.9
99.3
9 3, 4
97 .9
97.5
97, i
97,1
97.0
97.0
DLO;.G:T;jDt
3.9
3.9
3.3
3.7
3,7
3.7
3. 6
3.6
3.6
3.6
AU3-D0
17.4
21,3
26.3
23.3
30 .1
30 .5
31.6
34 ,3
29.5
30.4
D A L b t D
1.5
1.3
2.0
2. 1
2.1
2.1
2 .2
?,5
2.1
2,3
NUMBER
LATITUDE
-4,0
-3,5
-2.4
-1.1
- .4
.1
.
.3
.9
1.0
160
DLATITUDb
3.1
3.1
3.1
3.1
3.1
5.1
3.0
3,0
3.0
3,0
LONGITUDE
93.9
95 ,5
96.0
97.3
96.9
96.6
96.3
96,2
96.1
96.1
DLO.NGITUDE
3.1
5.1
3.1
3.1
3,1
3.1
3,0
3.0
3.0
3.0
ALBEDO
17.5
22,7
27.6
30.1
32.2
32,4
33.7
35,7
30,5
31.4
DALBEDO
.6
.5
.7
.8
.7
.7
. B
,8
.9
,7
NUMBER
LATITUDE
29,0
29.3
30.7
32.1
32,7
33.3
33.9
34,0
34,2
34.3
161
DLATiTUJE
3.2
3.2
3.2
3.2
3.2
3.2
3.3
3,3
3,3
3,3
LONGITUDE
97.0
96.5
96.0
95.1
94.7
94.3
94.11
93.9
93,7
93.7
DLONJGITUDE
3.2
3.2
3.2
5.2
3.2
3.2
3.3
3.3
3,3
3.3
ALBEDO
13,9
23.7
29,2
31. S
"34.5
34.9
37.0
39.7
34,2
35.7
DALBEDO
,S
.9
1.2
1.4
1,6
1.7
1.9
2.2
1,9
2.0
NUMBER
LATITUDE
13,0
13.3
14.7
16.1
16,7
17.3
17.8
18.0
18.2
16.3
163
DLATITUDE
3,0
3.0
3.0
3.0
3,0
3.0
3.0
3,0
3.0
3.0
LONGITUDE
114.0
113.6
113.2
112.6
112.3
112.
111.7
111,7
111.6
111.5
DLONGITUDE
3.1
3.1
3,1
3.1
3.1
3.1
3.1
3.1
3.1
3.1
ALBEDO
18,9
24,2
28.8
31.5
33,9
34.4
35.9
39.0
32.9
34.7
DALBEDO
.7
.9
1.1
1.3
1,3
1.4
1.5
1.7
1.6
1.9
NUMBER
LATITUDE
13.0
13.8
14.8
16.2
16,9
17.5
18.0
16,2
18.3
16.4
164
DLATITUDE
3.0
3,0
3.0
3.0
■ 3.0
3.0
3.0
3.0
3,0
3,0
LONGITUDE
130,0
129.7
129.3
126.7
128,4
128.2
127.9
127,9
127.8
127.6
DLONGITUDE
3.5
3.5
3.5
3.5
3.5
3.5
3.5
3,5
3,5
3.5
ALBEDO
19.1
23.6
29.2
31.8
34.2
34.5
36.6
39,8
33.9
33.8
DALBEDO
1.5
1.8
2.2
2.4
2,5
2.6
2,8
3,2
2,8
3.0
NUMBER
LATITUDE
-15.0
-14,2
-13.3
-12.0
-11.3
-10.7
-10.2
-10,1
-9.9
-9,8
170
DLATITUDE
2.0
2.0
2.0
2.0
2.0
2.0
2.C
2.0
2,0
2.0
LONGITUDE
340.0
339.6
339.1
338.4
338.0
337.6
337.4
337.3
337,2
337.1
DLONGITUDE
3.7
3.7
3.6
3.6
3.6
3.6
3.6
3.6
3,6
3.6
ALBEDO
16.1
19.7
22.7
22.2
22.6
22.4
22.5
23.7
18,4
17.9
DALBEDO
1.1
1.2
1.4
1.5
1.5
1.5
1.5
1.7
1.6
1.8
NUMBER
LATITUDE
-17.0
-16.2
-15.2
-13.8
-13.1
-12.5
-12.0
-11.8
-11.7
-11.6
171
DLATITUDE
2.0
2.0
2.0
2.0
2.0
2.0
2,0
2.0
2.0
2.0
LONGITUDE
350.1
349,7
349.2
343,5
348.1
347.8
347.5
347,4
347.4
347.3
DLONGITUDE
3.5
3.5
3.5
3,5
3.5
3.4
3.4
3,4
3.4
3.4
ALBEDO
15.0
,13.2
21.0
20. 1-
20,7-
20.0
20.1
21.0
17.1
17.2
DALBEDO
1.0
1.1
1.1
1,1
1.1
1.1
1.1
1,3
1.2
1.5
NUMBER
LATITUDE
-4.9
-4.2
-3.2
-1.9
-1.2
- ,6
-.1
.0
.2
.3
172
DLATITUDE
2.0
2.0
2.0
2.0
2.0
2.0
2.0
2.0
2.0
2.0
LONGITUDE
355.9
355,5
355.0
354,2
353.9
353.5
353.2
353.2
353.3
3 5 3.0
DLONGITUDE
2.0
2.0
2.0
2.0
2.
2.0
2.
2.0
2. D
2,
ALBEDO
13.5
15.2
17.2
16.1
16.4
16.3
16.3
16.5
14,7
14.9
DALBEDO
.5
,6
.6
.5
.5
.5
. N
,5
1.1
,8
NUMBER
LATITUDE
9.0
9.7
1C.6
11.8
12.4
13.0
13.5
13.6
13,6
13.9
173
DLATITUDE
3.0
3.0
3.0
3.0
5.0
3.0
5.0
3,0
3,0
3 .0
LONGITUDE
341,0
340.5
339.9
338.9
338.4
333.0
337.6
337.5
33-', 3
337.3
DLONGITUDE
3.3
3.3
3.4
3.4
3. 4
3.4
3.4
3.4
3,4
3.4
ALBEDO
13,7
23.4
28.7
30.9
33.0
32.9
33,6
36.3
30.6
32,0
DALbEDO
.a
.9
1.2
1.4
1.6
1.6
1.7
1.9
1,3
1.8
Sec. 3. 2, Appendix A, page 16 R. Newburn, JPL
October 1, 1971
JPL 606-1
Ultraviolet, Visible, and Infrared Photometric Properties
Table A-1. Martian albedo (continued)
WAVtLENGTH
.60
.64
.70
.89 1.04 1.24 1.61 1.74 2.14 2.2?
NUMBER
LATITUDE
-33.8
-32,7
-31.5
-29.7
-28.8
-28.1
-27,
, 4
-27,
,2
-27.0
-26.9
174
DLATITUDE
4.3
^^.2
4.1
4.0
4,0
3.9
3
,9
3
.9
3.8
3.8
LONGITUDE
347.5
347.2
346.9
346.4
346.2
346.0
345,
,8
345,
,7
345.6
345.6
DLON'GiTUDE
4.6
""'.S
4 .4
4.3
4.2
4.2
4,
,2
4 ,
,2
4.1
4,1
ALBEDO
16.1
■IB, 6
20.8
20 ,7
21.2
21.3
21,
,7
23,
,1
17.8
20.0
DALBEDO
2.3
2.5
2.7
2.5
2.5
2.5
2,
,5
2,
, 7
2.2
2.4
NUMBER
LATITUDE
, 1
.9
1.8
0.1
3.8
4.4
4,
,9
5.
,0
5.2
5,3
175
DLATITUDE
2.0
2.0
2.0
2.0
2.0
2.0
2,
,0
2
.0
2.0
2.0
LOi\GlTUD£
343.0
342,5
342.0
341.2
340.3
340.4
340.
.1
340 ,
,0
339.9
339.8
DLONGITUDE
3.0
3,0
3.0
3.0
3.0
3.0
3,
.0
3
.0
3.0
3.0
ALBEDO
16.9
21.7
26.1
28.5
31.2
31.4
32,
, 4
35,
,2
28.8
29.3
DALBEDO
.9
1.0
1.1
1.4
1.5
1.6
1
,6
1
,7
1.6
1.8
NUMBER
LATITUDE
25.9
26.9
28.0
29.7
30.5
31.2
31,
,9
32,
,1
32.3
32.4
180
DLATITUDE
2.0
2.0
2.0
2.0
2.0
2.0
2,
.0
2
,0
2.0
2.0
LO\'GITUDE
36.0
35.9
35.6
35.3
35.1
35.0
34
.9
34
.8
34.8
34,8
DLONGITUDE
3.0
3,0
3.0
3.0
3.0
3.0
3
, U
3
.
3.0
3.0
ALBEDO
14.4
16.1
19.2
17.9
18.5
17.4
16
.9
17
.7
14.1
15.7
DALBEDO
,9
1.1
1.3
1.2
1.3
1.3
1.
,3
1
.5
1.8
1.7
NUMBER
LATITUDE
-12.0
-11.0
-9.8
-8.1
-7.3
-6.6
-6 ,
.0
-5
,8
-5.5
-5.4
ISl
DLATITUDE
3.3
3.3
3.2
3.2
3,2
3.2
3,
,2
3,
,1
3.1
3.1
LONGITUDE
20.2
,19.8
19.5
18,9
18.6
18.4
18,
.2
18
.1
18.0
18.0
DLONGITUDE
3.3
3,3
3.2
3.2
3.2
3.2
3,
,2
3
.1
3.1
3.1
ALBEDO
12.4
14,0
16.0
15.0
15.6
15,5
15,
,5
16
,2
12.5
14,3
DALBEDO
.6
.3
.7
.7
.7
.7
,7
.8
.5
.7
NUMBER
LATITUDE
-22.0
-20,9
-19.6
-17.7
-16.8
-16.0
-15,
.3
-15
.1
-14.8
-14.7
183
DLATITUDE
3.6
3,6
3.5
3.4
3.4
3.4
3
.4
3
,4
3.4
3.4
LONGITUDE
20.0
19.7
19.3
18.8
18,5
18.3
18,
,1
18
,0
18.0
17.9
DLONGITUDE
3.6
3,6
3.5
3.4
3.4
3.4
3
,4
3
.4
3.4
3,4
ALBEDO
11.6
13.2
14.0
13.4
14,3
14.0
13
.8
14
.9
11.3
11.8
DALBEDO
,9
.9
.9
.8
.9
.8
.8
1
.0
1.0
1.5
NUMBER
LATITUDE
4.1
5,1
6.3
8.0
8.8
9,6
10,
,2
10
.4
10.6
10.8
184
DLATITUDE
3.0
3,0
3.0
3.0
3.0
3.0
3,
,
3 ,
,0
3,0
3.0
LONGITUDE
45.0
44.8
44.5
44 ,1
43.9
43.7
43
.5
43
.5
43.5
43.4
DLONGITUDE
3.3
3.3
3.3
3.3
3.3
3.3
3,
,3
3
.3
3.3
3.3
ALBEDO
15.8
20,0
22.5
23.3
25.4
25.3
26,
,0
27
.8
23.5
22,7
DALBEDO
1.3
1.5
1.8
1.7
1.9
2.0
2
.1
2
,4
2.4
2.4
NUMBER
LATITUDE
-6.0
-4.9
-3.7
-1.8
-.9
-.1
,6
,8
1.0
1.1
185
DLATITUDE
3.2
3,1
3.1
3.1
3.1
3.1
3,
.1
3,
.1
3.1
3.1
LONGITUDE
50.0
49.7
49,3
48.8
48.5
48.3
48
.1
48
.1
48,0
4S.0
DLONGITUDE
3.6
3,5
3.5
3.5
3.4
3,4
3,
,4
3
.4
3,4
3.4
ALBEDO
15.9
18,9
22.4
23,9
26. !■
25.7
26
,7
28
.1
24.4
24.9
DALBEDO
1.5
2.0
2.1
2.3
2,3
2.5
2,
,6
2 ,
,9
2,8
2.7
NUMBER
LATITUDE
-28.1
-26,6
-25.2
-23.0
-21.9
-20.9
-20.
.1
-19,
,9
-19.6
-19.5
166
DLATITUDE
3.9
3,B
3.8
3.7
3.6
3.6
3,
,5
3,
,5
3.5
3.5
LONGITUDE
45.8
45,3
44.6
43.7
43.3
4.3.0
42,
,7
42.
,6
42.5
42.5
DLONGITUDE
4.3
4,2
4.1
3.9
3.9
3.8
3,
,8
3,
,8
3.7
3.7
ALBEDO
14.7
15. a
18.6
16.3
18.7
18.5
18,
,'9
19,
, 1
16.4
16,7
DALBEDO
1.9
1.9
2.2
2.0
2.0
1.9
1.
9
2,
,1
1.9
2.3
NUMBER
LATITUDE
10.6
11.0
11.5
l-^.l
12.4
12.7
13.
13.
13.1
13.1
187
DLATITUDE
2.0
2.0
2.0
2.0
2.0
2.0
2.
2.
2.0
2.0
LONGITUDE
287.0
285.9
284.4
282.2
281.1
280.1
279,
,2
276.
9
278.5
273 .4
DLONGITUDE
2.0
2.0
2.0
2.0
2.0
2,0
2,
2 ,
2.0
2. G
ALBEDO
10.1
12.3
13.5
12.4
12.6
12.5
13,
1
13,
9
11.4
11.7
DALBEDO
1.0
1.3
1.4
1.4
■ 1.5
1.6
1.
a
2.
1.7
1.9
October 1, 1971
R. Newburn, JPL Sec. 3. 2, Appendix A, page 17
Ultraviolet, Visible, and Infrared Photometric Properties
JPL 606-1
Table A-1. Martian albedo (continued)
WAVELENGTH
.60
.64
70
,89 1.04 1.24 1.61 1.7< 2.14 2.27
NUMBER LATITUDE
192 DLATITUDE
LONGITUDE
DLONGITUDE
ALBEDO
DALBEDO
-3.8 -3.1 -2.3 -1.0 -.4 .1 .6 ,7 .9 1.0
3.1 3,1 3.1 3.1 3.1 3.1 3.1 3.1 3.1 3.1
317.9 317.4 316.8 315.8 315.3 314.9 314.5 314.4 314.3 314.2
3.G 3.8 3.9 3.9 3.9 3.9 3.9 3.9 3.9 3.9
16.3 22,6 26.5 29,2 31.0 31.1 31.0 33.9 27.8 27.3
1.2
1.7
2.0
2.4
2.6
2.7
2.8
3.2
2,7
2.7
■NUMBER LATITUDE
194
DLATITUDE
LONGITUDE
DLONGITUDE
ALBEDO
DALBEDO
-17.9 -17,1 -16.1 -14.6 -13.9 -13.3 -12.7 -12.6 -12.4 -12,
2.0 2.0 2.0 2.0 2.0
331.9 331.6 331.1 330.4 330.1
3.3 3.7 3.7 3.7 3.7
14.0 17,0 20.2 20.0 20.4
1.0 1.2 1.4 1.4 1.4
2.0 2.0 2.
329.8 329.5 329,
3.7 3.7 3,
19.9 19.3 20 ,
1.4 1.3 1,
2.0 2.0
4 329.4 329.3
7 3.7 3.7
5 16.2 16.2
5 1.2 1.4
NUMBER
195
LATITUDE
DLATITUDE
LONGITUDE
DLONGITUDE
ALBEDO
DALBEDO
-41.9 -40.6 -39.1 -37,0 -36.0 -35.2 -34.4 -34.2 -33.9 -33.8
5.1 4.9 4.8 4.6 4.5 4.4 4.3 4.3 4.3 4.3
343.5 343.2 342.8 342,2 341.9 341.7 341.5 341.4 341.3 341.3
5.2 5,1 4.9 4.7 4.6 4.5 4.5 4.5 4.4 4.4
16,5 18,9 21.7 21.5 22.2 21.6 21.5 22,7 16.8 16.7
3.3
3.5
3.7
3.4
3.4
3.3
3.2
3.4
2.9
3.1
NUMBER
197
LATITUDE
DLATITUDE
LONGITUDE
DLONGITUDE
ALBEDO
DALBEDO
17,
4,
7,
4,
18,
18,5
4.1
7.1
4,1
22.9
.7
19,
4,
6,
4,
28,
20.3 20.8 21.2 21.6 21,7 21.9
4.1 4.1
4.7 4.0
4.2 4.2
31.0 33.5
1.1 1.2
4.2
3.4
4.2
34.0
1.3
4 .2
2.8
4.3
35.8
1.4
4.2 4.2
2.6 2.4
4.3 4,3
38,9 33.3
1.6 1.5
21
4
2
4
34
1
NUMBER
198
LATITUDE
DLATITUDE
LONGITUDE
DLONGITUDE
ALBEDO
DALBEDO
■17,
4,
24,
4,
13
•16,4
4.3
24.1
4.4
16.6
.8
■15.6
4.3
23.2
4.3
18.8
.9
-14,5 -13.9 -13.5 -13.0 -12.9 -12.8
4,3
21.9
4,3
17.9
.9
4.2
21.2
4.3
16.8
.9
4.2
20.7
4.2
18.7
.9
4.2
20 .2
4.2
19.7
.9
4 ,2
20 .0
4.2
21,2
,9
4.2
19.9
4.2
17.2
1.0
■12.7
4.2
19.8
4.2
17.8
1.1
NUMBER
199
NUMBER
200
LATITUDE
DLATITUDE
LONGITUDE
DLONGITUDE
AL8ED0
DALBEDO
LATITUDE
DLATITUDE
LONGITUDE
DLONGITUDE
ALBEDO
DALBEDO
'17.0 -16.3 -15.5 -14,4 -13,8 -13.4 -12.9 -12.6 -12.7 -12.6
4
25
4
13
4
35
4
17
7.8
4,3
24.4
4.4
16.0
.8
4.0
34,3
4.2
22.3
.6
4,3
23.5
4.3
17,9
.8
4.3
22.2
4.3
17.3
.7
4,2
21.5
4,3
17,8
.7
4,
21
4
17.8
.8
4.
20 ,
4,
18,
4 ,
20
4
20,
4.2
20.2
4 .2
16.9
.9
4.0
33,5
4,2
26.3
.7
4,
32,
4,
28,
4.
31.
4,
29,
4,
31,
4,
-29,
4,
30,
4 ,
30 .8
.6
4,
30,
4,
33,
4.0
30.4
4 .1
28. ij
,9
4
20
4
17
6.5 9.2 10.3 10.8 11.3 11.7 11.8 12.0 12.0
4 ,
30,
4
2?!
1
NUMBER
202
LATITUDE
DLATITUDE
LONGITUDE
DLONGITUDE
ALBEDO
DALBEDO
15.0 15.7 16,7 13.0 18.7 19.3 19.8 20,
4,0 4,0 4.1 4,1 4.1 4.1 4.1 4,
313,1 312,7 312,3 311,6 311.3 311.0 3i0.8 310,
4.0 4,0 4.1 4.1 4.1 4-.1 4.1 4,
19.8 25.1 30.4 33.5 35.9 36.5 38.5 40,
.4
.5
.6
.7
.8
20 .2 20.3
1 4.1 4.1
7 310.6 310.6
1 4.1 4.1
6 35 .1 36.5
.9
,8
NUMBER LATITUDE 33.0 33.8 34.9 36.4 37.2 37,9 38.4 38.6 33. tt 33.9
203 DLATITUDE 4.4 4.5 4.5 4.6 4.6 4.6 4.7 4.7 4.7 4.7
LONGITUDE 308.7 308,3 307.8 307.0 306.5 306.2 305.8 305.7 305.6 3 05.'"'
DLONGITUDE 4.4 4.5 4.5 4.6 4.6 4.7 4.7 <! , 7 4.7 4.7
ALBEDO 18.6 23.5 28.2 29.9 32,0 32.1 33.3 35.5 30,1 .'5 1 . 4
DALBEDO .9 1.2 1.6 1.6 2.0 2.1 2.3 2.5 2.2 2.3
Sec. 3. 2, Appendix A, page 18
R. Newburn, JPL
October 1, 1971
JPL 606-1
Ultraviolet, Visible, and Infrared Photometric Properties
Table A-1. Martian albedo (continued)
WAVELENGTH
.60
.64
NUMBER
204
NUMBER
206
NUMBER
207
NUMBER
208
NUMBER
209
NUMBER
210
LATITUDE
DLAT ITUDE
LONGITUDE
DLOMGITUDE
ALBEDO
DALBEDO
LATITUDE
DLATITUDE
LONGITUDE
DLOMGITUDE
ALBEDO
DALBEDO
LATITUDE
DLATITUDE
LONGITUDE
DUONGITUDE
ALBEDO
DALBEDO
LATITUDE
DLATITUDE
LO.-^GITUDE
DLONGITUDE
ALBEDO
DALBEDO
LATITUDE
DLATITUDE
LONGITUDE
DLONGITUDE
ALBEDO
DALBEDO
LATITUDE
DLATITUDE
LONGITUDE
DLONGITUDE
ALBEDO
DALBEDO
4
335
4
18
10
9.6
4,0 4
334,7 334
4.3
23,7
.9
4
28
1
70
6
3
3
6
1
.89 1.04 1.24 1.61
. • 74 2.14 ?. ?7
12,
4
333
4,
31
1,
12
4
333
4
34 .0
1.4
13,
4 .
333,
4 ,
34,
1 ,
13,
4
335
4 ,
36
1 ,
14.
4 ,
333 ,
4 ,
39,
1 ,
4
333
4
33
1
14.3
-= .
= ^? . 9
4 .3
34 .5
1 .5
23.0 23,8 24.8 26.3 27.0 27.7 2S . 3 26.4 28.6
4.2 4,2 4.2 A, 2 4.2 4.3 4,3 4. J 4.3
328.0 327.3 327.5 327.0 326.8 326.6 326.4 326.3 326.3
4.3 4.3 4.3 4.3 4.4 4.4 4.4 4.4 4.4
19.3 24.3 29.6 32.3 34.9 35.4 37,1 40.6 34.5
2,S
526 .
.7
.9
1.2
1.3
1.4
1.5
1.6
l.S
1.5
14.9 15.7 16.6 13.0 13.7 19.3 19.9 20.0 20,2 20.3
4.0 4,0 4.1 4.1 4.1 4.1 4.1 4.1 4.1 4.1
287.0 286.5 265.9 285.1 284.6 284.2 283.9 263.7 2S3.6 263.6
4.5 4,6 4.6 4,6 4.7 4.7 4.7 4.7 4.8 4.8
10.9 13.0 15,0 13.9 14.6 14.8 15.6 16,7 13.3 15,9
.6 ,8 1.0 1.1 1.2 1.2 1.2 1.3 1.1 1.1
14.9 15.7 16.7 18.1
19.4 20,0 20.1 20,3
4.0 4.0 4.1 4.1 4.1 4.1 4.1 4.1 4,1 4.1
287.0 286,5 286.0 285.1 284.6 284.2 283.9 28', 3 233.6 263,6
4.5 4,6 4.6 4,7 4.74.7 4.7 43 4,8 4.3
11.2 13,4 15.4 15.0 16.0 16.4 17.4 lti.3 15.0 16,6
.7 .8 1,1 1.1 1.2 1.3 1.2 1,3 1.2 1.2
2.1
4.0
289.0
4,5
12.0
.8
2.9
4.0
3.8 5.3
4.0
4.0
6.0
4.0
6.6
4.0
7.1
4.0
7.3
4.
268,6 288.2 287.6 287.2 286.9 266.7 286.6
4.5
14.0
.9
4.5
15.9
1.0
4.6
15.1
1.0
4.6
15.8
1.1
4.6
16.1
1.2
4.6
16,3
1. 2
4 .6
17.6
1.3
7. 4
4.0
286.5
4.6
14.5
1.2
7,5
4 ,
236,5
4 .6
15.4
1,5
28.0 28.8 29.9 31.5 32.3 33.0 33.6 33.8 34.0 34.1
4.3 4.3 4.3 4.4 4.4 4.4 4.5 4,5 4.5 4,5
328.1 327,9 327.6 327.2 327.0 326.3 326.6 326.6 326,5 326.5
4.4 4.4 4.4 4.5 4.5 4.5 4.5 4,5 4,5 4.5
18.9 24,5 29.8 32.9 35.4 35.8 37.8 40,7 35,1 36,0
.8 1,1 1.4 1.6 1.8 1.9 1.9 2,2 2,0 2.C
October 1, 1971
R. Newburn, JPL Sec. 3. Z, Appendix A, page 19
JPL 606-1
Ultraviolet, Visible, and Infrared Photometric Properties
APPENDIX B
GLOSSARY OF PHOTOMETRIC AND POLARIMETRIC TERMINOLOGY
Albedo
Blackbody
Bolometric
Bond (sometimes
Russell-Bond or
spherical) albedo
Brightness
temperature
Degree of
polarization
Detailed
photometry
Geometric
albedo
See geometric albedo. Bond albedo, normal albedo.
A body in complete thermodynamic equilibrium with its
surroundings, which in turn implies detailed balancing
of all associated atomic and molecula.r processes. A
blackbody is completely defined by one parameter, the
equilibrium temperature, and radiates according to
Planck's law.
An adjective implying radiometric data (rather than
photometric) integrated over all wavelengths. A bolo-
metric magnitude is thus a measure of total power, and
a bolometric albedo is a miean albedo over all wave-
lengths, unaffected by atmospheric or photometric
system absorptions and response.
The ratio of the power (flux) reflected in all directions
by a body to the power incident upon it in a coUimated
beam. It is the fraction of incident solar flux that is
NOT absorbed.
The temperature that a blackbody would have to have in
order to emit the same power (flux) that is actually
measured at the wavelength in question.
By definition, the difference Imax less Imin divided by
the sum Imax plus I^in- ■'Wiax ^^ *^® miaximum intensity
of the polarized beam which lies in the plane containing
the electric vector and the direction of propagation. I
is perpendicular to that plane.
Point by point photometry of an extended source.
The ratio of mean luminance of a body at full phase
(phase angle zero) to the luminance of an "intrinsically
white" plane surface normal to the source of illumiin-
ation (Sun). An "intrinsically white" surface scatters
all of the power incident upon it (absorbing none) and
does so according to Lambert's cosine law. Such a
surface is also called a perfectly diffusing surface.
mm
October 1, 1971
R. Newburn, JPL Sec. 3. 2, Appendix B, page 1
Ultraviolet, Visible, and Infrared Photometric Properties
JPL 606-1
Illuminance
Integrated
photometry
Inversion angle
Irradiance
Lambert surface
Lambert's
cosine law
The photometric equivalent of irradiance, the power-per-
unit solid angle and unit-projected area falling upon a
surface within the passband of a photometric system. In
the narrowest sense, this refers only to the passband of
the human eye, but in astronomy, the broader sense is
often used, applying the terminology to any defined
photometric system.
Photometric study of an entire body as a unit, as opposed
to detailed photometry.
The phase angle at which the degree of polarization
changes from negative to positive.
Radiometric term for power-per-unit solid angle and
unit-projected area falling upon a surface.
A surface which has the same radiance when viewed
from any angle.
A surface radiating (or reflecting or transmitting) an
amount of flux per unit area and unit solid angle pro-
portional to the cosine of the angle between the surface
normal and the direction of observation is said to follow
Lambert's cosine law. Such a surface is a Lambert
surface.
Luminance
Luminance
equator
Luminance
latitude
The photonnetric equivalent of radiance, the power-per-
unit solid angle and unit-projected area leaving a surface
within the passband of a photometric system. In the
narrowest sense, this refers only to the passband of the
human eye; but, in astrononny, the broader sense is
often used, applying the terminology to any defined photo-
metric system.
The intersection of the phase plane with the surface
under study.
The angle between the phase plane and the normal to the
surface at the point of observation. See Figs. 3 and 5 in
the text.
Lunninance
longitude
Magnitude
The angle of observation (reflection angle) projected into
the phase plane. See Figs. 3 and 5 in the text.
A logarithmic unit of electromagnetic flux, ancient in
origin, used in astronomy. In nnodern usage, one mag-
nitude is equivalent to a ratio of Z.512 in flux. See the
text for a complete definition of the system.
Sec. 3. 2, Appendix B, page 2 R. Newburn, JPL
October 1, 1971
JPL 606-1
Ultraviolet, Visible, and Infrared Photometric Properties
Normal albedo
The ratio of luminance of a point at zero phase to the
luminance of an intr-nsically white plane Lambert sur-
face normal to the illumination. This is the photometric
equivalent in detailed photometry to geometric albedo in
integrated photometry. See geometric albedo; also the
main text.
Opposition effect
Percent
polarization
Phase angle
Phase coefficient
Phase functior
Phase integral
Phase plane
Photometric data
Photometric
function
An enhanced brightness occurring for phase angles
|q|<10°. For example, in the B-passband, linear extra-
polation to zero phase would give Bjg(O) = -0.21 as Mars'
magnitude, where, in fact, it is B --0.31.
o
Just 100 times the degree of polarization. Sometimes
permil polarization, written %„ and equal to 1000 times
the degree of polarization is also used.
Astronomically, the object (body) centered angle between
the source of illumination (the Sun) and the observer
(detector). See Fig. 3. In local photometry, the angle
between source (Sun) and detector measured at the
observed point. See Fig. 5.
The phase function of a body is often presented in the
form of a polynomial expansion with the phase angle as
the argument. The linear coefficient is often called the
phase coefficient (and over the limited range of phase
angle available for Mars it is an adequate representation
without additional terms).
The ratio of power (flux) scattered at phase angle o to
that scattered at zero phase. The phase function is often
given as a cubic polynomial, in tabular form, or as a
polar graph.
The ratio of power (flux) scattered in all directions to
that scattered at zero phase, per unit solid angle. The
phase integral multiplied by the geometric albedo equals
the Bond albedo.
The Sun - object - observer plane, the plane containing
the phase angle. See Figs. 3 and 5.
Flux data convolved with the response of a particular
detector system. In the strictest sense, photometric data
is that received by a "standard" human eye, but, astro-
nomically, the term is applied to any calibrated combi-
nation of filters, detectors, etc.
The ratio of the radiance factor to the normal albedo for
a point on a sphere. The photometric function is a func-
tion of three parameters: the phase angle, angle of inci-
dence, and angle of observation; or, alternately, the phase
angle, luminance longitude, and luminance latitude.
October 1, 1971
R. Newburn, JPL
Sec. 3. 2, Appendix B, page 3
Ultraviolet, Visible, and Infrared Photometric Properties
JPL 606-1
Radiance
Radiance
(luminance)
coefficient
Radiance
(luminance)
factor
Radiometric data
Solar constant
Radiometric term for power-per-unit solid angle and unit-
projected area leaving a surface .
The ratio of radiance (luminance) observed to that of a
white plane Lambert surface at the same inclination to
the source of illumination.
The ratio of radiance (luminance) observed to that of a
white plane Lambert surface normal to the source of
illumination.
Flux data given in absolute units, deconvolved of any
photometric system response curve. These units can
be either astronomical (magnitude) or physical (watts)
and may still refer to a particular region of the spectrum,
rather than be integrated over all wavelengths (bolo-
metric data).
The irradiance (power-per-unit area) from the Sun at a
distance of one astronomical unit.
Spectral
irradiance
Spectral radiant
exitance
Stokes parameters
Radiometric ternn for power incident per unit area and
wavelength upon a surface.
Radiometric term for power-per-unit area and wavelength
coming from a surface.
Four parameters which give a complete description of
polarized light. See van de Hulst (1957) or Shurcliff
(1962) for details.
Sec. 3. 2, Appendix B, page 4 R. Newburn, JPL
October 1, 1971
JPL 606-1 Radar Properties
3.3 RADAR PROPERTIES
INTRODUCTION
Earth-based planetary radar observations have substantially improved
our knowledge of the orbits of the terrestrial planets, refining their ephemer-
ides, and providing data of importance in basic research, as well as in the
planning of missions to these planets. Radar observations have helped to pro-
vide a more reliable value for the radius of Mars, a value confirmed by the
Mariner flybys (see Section 1). Radar studies have provided extensive informa-
tion on Martian topography in equatorial latitudes, as well as on the dielectric
constant and roughness of the Martian surface.
Radar astronomy is simple in basic concept, but it is complex in execu-
tion and theory. The paragraphs immediately following discuss the concepts (it
is hoped) in sufficient depth to understand the observational results. No attempt
is made to offer detailed derivations of formulas, as these are available in
many texts . ''
3. 3. 1 BASIC CONCEPTS OF RADAR ASTRONOMY
In radar astronomy, a signal is transmitted with the highest available
power and the narrowest possible beamwidth in order to concentrate as much
power as possible on the target. The flux decreases proportionate to the square
of the distance from the transmitter. The angular width of the transmitted beam
is much greater than any planetary target, so only a small fraction of the trans-
mitted power strikes the target and is scattered back toward Earth. The target
returns also grow weaker as the square of the distance and are received by the
largest possible collecting area (antenna). The basic radar equation relating
power transmitted P^ to power (signal) received Pg is
P G^A 0-
s t r
P ~ 2 Z
where
G is gain of the transmitting antenna (compared to an isotropic
radiator)
A is area of the receiving antenna
0- is the target scattering cross-section
r is the target distance
-23
Typically, for the largest radar facilities this ratio is about 10 for Mars,
even when the planet is near a perihelic opposition and as near as it ever comes
*See (for example) Radar Astronomy by Evans and Hagfors (1968).
November 15, 1971 C. M. Michaux, JPL Sec. 3.3, page 1
Radar Properties JPL6O6-I
to Earth. Thus, in radar astronomy, as in radio astronomy, the reception of
weak signals is the general rule, and improvements in the signal-to-noise ratio
have come about through larger antennas, increased transmitter power, and
more sensitive (lower noise) receivers.
The transmitted frequency is limited to wavelengths longer than about
1 cm, due to increasing absorption in the Earth's atmosphere caused by water
vapor and oxygen at the shorter wavelengths. At wavelengths greater than about
20 m, signals are scattered by the Earth's ionosphere. Within the available
frequency range (1 cmi to 20 m), the choice is usually at the shorter wavelength
end, because antenna gain increases inversely as the square of the wavelength
for a given antenna area, at least to a point where imperfections in the surface
of the antenna begin to approach the wavelength transmitted.
The signal may be transmitted continuously (CW) at powers up to sev-
eral hundred kilowatts, or it may be pulsed with a peak power of many mega-
watts. Either type of transmission may be additionally modulated or coded; to
obtain range resolution with continuous wave (CW) radar, and to remove range
ambiguities with pulsed radar.
The signal-to-noise ratio may be improved by time integrations of the
received signal. Since the round-trip time for a signal transmitted to Mars is
more than six minutes, even when Mars is closest to Earth, it is possible to
transmit for at least six minutes and then listen (receive) for six minutes with-
out confusion, even using only one antenna. For Mars, there is a complication
introduced by the relatively rapid rotation of the planet. Long integration times
inevitably result in a loss of surface resolution.
All results discussed in this section are the product of monostatic radar
experiments, experiments in which the angle of incidence of the transmitted
signal is the same as the angle of reflection (or backscattering) of the received
signal. The angular size of Earth as seen from Mars is so small (< 1 arc min
at closest approach) that even experiments using different transmitting and
receiving antennas are still essentially monostatic. A true bistatic experiment
could be carried out by transmitting from Earth and receiving from a nonter-
restrial site or spacecraft.
3.3.2 CROSS -SECTION AND REFLECTIVITY
Fundamental Concepts
The target or total radar cross-section cr is usually defined as 4n timies
the ratio of the power-per -unit solid angle scattered back toward the trans-
mitter, to the power -per-unit area (power density) striking the target (Westman,
1956). ■■■
=:=This classical definition assumes a monostatic radar systemi is used. With a
bistatic radar, the radar cross-section cr, of the illumiinated target is a mea-
sure of the energy scattered in the direction of the receiver (Skolnik, 1962).
Sec. 3.3, page 2 C. M. Michaux, JPL November 15, 1971
JPL 606-1 Radar Properties
The radar cross- section is a characteristic of the target and is a
measure of its size as seen by the radar (Skolnik, 1962), and has the dimen-
sions of area (m ). '■'■'■ The radar cross -section depends not only upon the actual
size and orientation (the geometrical cross - section), but also on the reflectivity
and the roughness oi the target.
The radar cross - section cr of a spherical body, such as a planet in first
approximation, is expressed mathematically in two different ways, relative to
its geometrical cross- section ttR (where R is the planetary radius):
2
1) first formulation: cr = p gnR
_ o
2) second formulation: cr = pGirR
where the dimensionle s s parameters p^ , g, p, and G are the "reflectivity
under normal incidence," the "directivity," the "spherical (or Bond) reflec-
tivity, " and the "gain in backscattering, " respectively. Only the first formu-
lation is useful in describing inonostatic measurements of Mars, and it will be
considered in more detail.
The reflectivity p^ is the classical Fresnel power reflection coefficient
for a plane homogeneous surface when a plane wave is incident upon it normally:
Po
1 +
where the symbols are
s = electrical conductivity
e = permittivity \ with subscript ^ referring to free space values;
> therefore, eA^ and |J./p are relative permittivity
fj. - permeability ) (dielectric constant) and relative permeability
oj - angular frequency - related to frequency f or free space wave-
length X , respectively, by co = Z-rrf =: 2ttcAq,
where c = velocity or light
i = imaginary unit
Thus, it is seen that generally p^ depends upon wavelength and upon the electro-
magnetic characteristics of the reflecting surface material. K the material is
a perfect dielectric (s = and |a = M-q), then p becomes
-The radar cross -section often is expressed in fractional terms of the ratio
(t/ttR where R is the radius of the target, or as a percentage, throughout
this section.
November 15, 1971 C. M. Michaux, JPL Sec. 3.3, page 3
Radar Properties
JPL 606-1
1 + A /—
Therefore, for a perfect dielectric, p^ is independent of wavelength \ and only
depends upon the dielectric constant k - c/^^ • Generally, rocks and soils in a
very dry environment have a very low conductivity s. Therefore, provided the
relative permeability [jl/Vq ^^ close to unity (when only minute amounts of mag-
netic material are present), this last equation can be used with fair approxima-
tion. ='■ Its application seems to have met with success in the case of the xMoon.
It may be less valid for Mars, which is closer to the asteroidal belt and possibly
has larger mounts of magnetic meteoritic inaterial on its surface. However,
for lack of firm knowledge, this equation will be used later to derive an approxi-
mate value of the effective dielectric constant.
The directivity g is determined by the ability of the actual (rough spheri-
cal) surface to backscatter favorably toward the illuminating source (monostatic
radar). The actual returns are compared with those obtained from a perfectly
smooth spherical (isotropic) surface, for which g = 1. Directivity g should not
be confused with gain (Rea et al., 1964). For an arbitrary rough sphere, g can
be calculated, pruvi led the statistics of height deviations and surface slopes are
known. For example, tlie case of a smooth undulating sphere was treated by
Hagfors (1964), who gave the approximation: g = 1 + a^ , where a is the rms
surface olope {a is sinall; e.g. , a. ~ 0.1 at X 1 m).
The distance between Earth and Mars constantly changes as the two
planetr mo' c around the Sun. The frequency transmitted from Earth therefore
returns at • sLghtly (Doppler) shifted central frequency f. The amount of this
shift is accurately known and is usually automatically compensated for by means
of a computer-controlled adjustment of the receiver. Far miore imiportant is
the Doppler spread introduced by the rotation of Mars. Because the rotating
target is a sphere (approximately), its illuminated hemisphere returns the wave
in the form of fine semiannular slices of its surface, with constant frequency f"
aligned parallel to the plane formed by the axis of apparent rotation and the line-
of-sight. To the radar observer, the disk to be explored may be divided into a
series of parallel lines or narrow strips of constant Doppler shift Af, symmetri-
cal about the central frequency f line or strip, which has zero shift (Af = 0)
and represents the apparent rotation axis (see Fig. 1). The spread Af, corre-
sponding to each line, is proportional to its distance from the f line (axis), and
the maximum spread Afj^^^^ occurs for returns from the two most distant lines,
tangent to the limbs. This is given by Afj^-^^j^ - Zcu' • R/^q' where lo' is the
apparent rotation ra^e (rad^sec), R the radius of the planet, and Xq = c/f the
center v/avelength.
Af
max
Because of the fast rotation of Mars, the rotational Doppler effect (i.e. ,
at the limbs) is much larger than the orbital Doppler effect. The total
'Pq may depend upon X, since longer wavelengths penetrate to greater depth
where denser soils will have higher dielectric constant.
Sec . 3.3, page 4
CM. Michaux, JPL
November 15, 1971
JPL 606-1
Radar Properties
AXIS OF ROTATION
APPROACHING
SIDE
MOON
RECEDING
APPROACHING
DECREASED NORMAL
FREQUENCY FREQUENCY
INCREASED
FREQUENCY
Fig. 1. The system of constant delay rings and doppler shift strips
on the disk (Evans and Hagfors, 1968).
(Only one ring and one strip are indicated. )
limb-to-limb bandwidth B (= 2Afj^ajj.) is given in Table 1 for several commonly
used radar frequencies.
Table 1. Doppler spread, or "limb-to-limb bandwidth" B of the Martian echo,
as a function of operating frequency f (=f') or wavelength X ( = \' ).
f, MHz
\' ^"^
B, Hz
7,840
3.8
-23,500
2,388
12.5
7,670
700
43
2,200
430
70
1,280
November 15, 1971
CM. Michaux, JPL
Sec . 3.3, page 5
Radar Properties JPL 606-1
An additional contouring of the planet is introduced by the finite velocity
of light. When a short pulse of duration or width t transmitted to a planet, and
echoes are received at a later time T, those echoes must come from an annulus
(or ring) whose edge-on width is ct/2. Furthermore, if the first echoes (from
the subradar point) arrive at a time T^, and the echo of interest arrives at a
time t = T - Tq later, then the distance between the plane of the annulus and
the subradar point is c t/Z. These rings are also plotted in Fig. 1. The area
of the annulus is independent of its location (or delay time t) and is equal to
ttRct, determined only by pulse duration t, with R the radius of the planet, and
c the speed of light. The angle of incidence i (equal to the angle of back-
scattering 4)) is practically constant over the annulus area (if the pulses are very
short, order of fisec), and is related to the delay time t by the interchangeable
relations (with unique correspondence):
= 4) = arc cos
(' ■ i)
or
2R
(I - cos i)
It can be seen that the combination of Doppler strips and delay rings decomposes
the entire planet into a set of "resolution cells" symmetric about the apparent
equator. Each cell is specified by its set (t, Af) of delay-resolution coordinates.
There is an inherent North-South ambiguity (since one strip intersects a ring
twice), but this may be resolved either in time through rotation of the surface,
or spatially by using a narrow beamwidth (as has been done for the Moon). The
present radar capability does not permit narrow beamwidth operation at the dis-
tance of Mars. Although it is possible theoretically to produce reflectivity maps
of the entire disk, only areas close to the subradar points have been studied to
date, on Mars. A set of "range -gated" spectra (spectra separated according to
delay rings) taken by Goldstein et al. (1970) is shown in Fig, 2.
Observation Techniques
Usually continuous waves (CW) are employed to measure the radar cross-
section of a planet. '■' The actual measurement may be perfornned either directly,
by using a radiometer (measuring total power), or indirectly, by integration of
the Doppler spectrograms.
The total radar cross- section represents the echo power reflected from
the entire disk; that is, it must total both quasi- specular and diffuse scattering
portions of the echo, and may be expressed: u - (^g + o"d • At short (centimeter)
wavelengths, the quasi- specular returns have a higher (peak) intensity, but total
power is greater in the diffuse component of the cross- section. At longer (meter)
^Pulses may be employed to measure <r, but the pulse length (duration t) must
be as long as the radar depth of the planet (2R/c), so that a full hemisphere is
illuminated at one time to obtain the total cros s - section. The radar depth for
Mars is 22.6 msec.
Sec. 3.3, page 6 C. M. Michaux, JPL November 15, 1971
JPL 606-1
Radar Properties
a) AERIA
b) CANDOR
I ,^ — -^ -
L
-320
■320
-160
160
320
-160 160
DOPPLER SHIFT (Hz)
320
DOPPLER SHIFT (Hz)
Fig. 2, Samples of range-gated frequency power spectra set
(Goldstein et al. , 1970),
wavelengths, it appears likely* that the opposite is true. Good measurements
of total cross -section can only be obtained with a high signal-to-noise ratio.
Early measurements were dominated by the quasi-specular component.
A radar observing run consists of a succession of alternate transmis-
sions and receptions, usually continuing as long as the planet is conveniently
above the horizon. Transmission is made in a circular polarization mode in
order to avoid the problems of Faraday rotation (of a linearly polarized wave
traversing our ionosphere). Reception of the echo waveform from Mars
because of the weak signal-to-noise, has, so far, only been made in the sense
of circular polarization, 180° opposite in phase to that transmitted. This 180°
return corresponds to specular reflections from a smooth surface. Only the
so-called "polarized" component of the echo has been received from Mars
When radar capability improves, it will be possible to receive the same 'sense'
as transmitted; that is, the weaker so-called "depolarized (or cross-polarized)
component, " which is the depolarization product of surface roughness.
Antenna aiming is done automatically by computer. The central frequency
of the receiver is also automatically adjusted to compensate for Doppler shift
introduced by the relative velocity of the observing station and the target (Mars).
-If Mars behaves like the Moon, which appears likely as discussed later.
November 15, 1971
C. M. Michaux, JPL
Se<
3 . 3, page 7
Radar Properties
JPL 606-1
Because of the fast rotation rate of Mars itself, the subradar point
travels some 100° in longitude during one night of observation, maintaining
nearly the same circle of latitude. Because the terrestrial rotation is some
37 minutes faster than Mars, several weeks are necessary in order to cover a
full 360° in longitude over Mars from a single station. (There is considerable
overlap of sectors visible in successive nights. ) After several weeks, the sub-
radar point also will be circling at a somewhat different latitude, the excursion
in latitude being given by the change in areocentric declination (Dj^ in American
Ephemeris) of the Earth. Figure 3 shows these excursions for the observing
periods of the experiments since 1963.
Early Radar Observations (1963 and 1965)
The first successful radar observations of Mars were made by two groups,
during the unfavorable opposition of early 1963: In the United States, by a JPL
group (Goldstein and Gilmore, 1963), at 12.5-cm wavelength; and in the Soviet
Union, by a Radioengineering and Electronics Institute group (Kotelnikov at al. ,
1963),' at 43-cm wavelength. Both groups reported a cross - section varying
directly with the surface oresented: <t varied from 2 to 13 percent as the sub-
radar point circled the plknet at 13°-14°N latitude. Kotelnikov et al. obtained
an average of 7 percent, similar to that of the Moon. Their results are consid-
ered unreliable, however, because of very low signal-to-noise ratio (~1. 5-2.5).
Goldstein and Gilmore, covering a full rotation of Mars, produced the first radar
brightness map along the 13 °N parallel- -which visually is practically all bright
NOTE OPPOSITION DATES ASE
INDICATED 8Y THE SIGN T
I 1 i
i"iyio30l 10 ^H~io~20 3ol 10 20 30 10 20 3ol 10 20 30l 10 20 30' 10 20 30 10 20 30i 10 20 30 10 2(N 1" 20 30
Jan.1 Feb.l MarJ Apr.l May 1 June! July 1 Aug-l Sept.l Oct . 1 Nov.1 Dec.l
DATE
Fig. 3. Excursions in Martian latitude of the subradar point
for the 1963 to 1971 apparitions of Mars.
Sec. 3.3, page
C. M. Michaux, JPL
November 15, 1971
JPL 606-1 Radar Properties
area-- and they noted that the visually dark Syrtis Major appeared bright to
radar. They also found, from their average Doppler spectrogram, that Mars
is somewhat smoother than Venus.
After 1963, the Russians appear to have discontinued radar observations
of Mars. From that time through 1970, reports have come only from investiga-
tors in the United States. At the opposition of March 1965, Goldstein (1965),
with improved radar equipment, secured 36 Doppler spectrograms, each an
average of many runs, from successive longitude intervals near the 21 °N lati-
tude circle. Some areas, such as Trivium Charontis-Cerberus , showed high,
narrow peaks, indicating very smooth, strong reflecting areas. Others (e. g. ,
Nodus Laocoontis) had a wider, lower peak, suggesting a rougher, but still
strong, reflecting area. Surprisingly, the northern tip of Syrtis Major showed
no strong echo. The large Amazonis desert was a poor reflector, but the poor-
est reflecting areas were the dark features Ascraeus Lacus and Albis Lacus.
On the whole, Mars was found to be significantly smoother than Venus at
X12.5 cm.
At MIT, Evans et al. (1965) derived a cross- section average of 14 per-
cent by summing the echoes obtained over four nights of observations at \23 cm.
Investigations which may be considered the forerunner of more modern
observations were conducted by Dyce, et al. (1967), at Arecibo Observatory in
1965 (AIO). Using the 1000-foot antenna at X70 cm, they measured an average
radar cross-section of 7 percent, varying with rotation between 3 to 13 percent,
as the subradar point latitude reinained within 1° of 22 °N. These results are
very similar to the first results of 1963 at U^N-M^N latitude,
Dyce et al. (1967) illuminated Mars with a stream of short (4 and 19 msec)
phase-coherent pulses, and analyzed the echoes by the delay-Doppler method.
The resultant cross-section measurements confirmed that dark areas crossing
the central meridian gave strong echoes, but the correlation was noted to be
imperfect. The "significant differences in precise alignment" were attributed
to large-scale slopes passing the subradar region. Their plot of the radar
cross-section versus longitude at the 22 °N parallel was compared against a
simplified strip map (see Fig, 4). An example of this comparison indicates
that Syrtis Major corresponds to a local minimum between two high peaks,
while Trivium Charontis lies slightly off an isolated high "peak. "
Secondarily, the delay-Doppler analysis for Mars definitely established
the presence of two components of scattering. These were (1) a prominent,
narrow-band, which was somewhat variable with the longitude (the "central'
peak" of the spectrum). This was attributed to the quasi-specular reflections
from tlie central region of the disk; and (2) a low-level, relatively broad-band,
less variable component which corresponded to the large, more distant sur-
rounding regions extending almost halfway to the limbs. Dyce et al, presumed
the latter component to be due to diffuse scattering.
By refining their analysis through measurements of bandwidths and shifts
of the central peaks, they calculated that the large-scale "smooth" slopes
responsible for the asymmetry of the spectral peaks varied from about 10° to
less tha n 3 ° .
November 15, 1971 C. M. Michaux, JPL Sec. 3.3, page 9
Radar Properties
JPL 606- 1
MARE
ACIOALIUM
TRIVIUM
CHARONTIS
SYRTIS
MAJOR
^ 15
I I I M
NORTH
lATITUDE
+ 10"
+ 20*
+ 30°
MARS 1965
I I I
I I 1
I I I I I I I
I I 1
100 200
EAST LONGITUDE (DEGREES)
300
360
Fig. 4. The variation of the radar cross-section near the 22°N latitude
as a function of the longitude of the central meridian of the visible
disk as obtained at 70-cm wavelength (Dyce et al. , 1967).
Recent Radar Observations (1967 and 1969)
Carpenter (1967) transmitted X12.5 cm CW at Goldstone, using an 85-foot
antenna for transmission and the Z 10-foot antenna for receiving the signal
returns. Many Doppler-broadened spectrograms of Mars were obtained during
April and May 1967, scanning the 21 °N latitude area of Mars; see Fig. 5(a).
The maximum bandwidth of 3 . 7 kHz excluded only half of the total spectrum
(approximately 30° in longitude). Utilizing the 60 composite (average) spectra
obtained at every 5° longitude aroimci /lavs, Carpenter measured the radar
cross -section variations, and attemptf-d 'o r-s^iniate surface roughness from the
half-power widths. His results are ph. ' .rd :a l<'ig. "-(b) as a f^-actional cross-
section versus areocentric longitude, with specific Martian features also noted.
The mean ct/ttR is 0.063, but remarkable variations rar.L'i;ig fmni '""i^
to 0.123 were found along the explored area of 21 °>; pjrallel. Cai /-ntti fodiiti
no clear relationship be'tween the visual and radar -derived a;)peararice of Mars.
Pettengill et al. (1969), in 1967, and Rogers et al. (1970), in 1969, in a con-
joined effort at MIT's Lincoln L,aboratory, explored Mars at \3.8 cm by using
unmodulated CWfor cross -section and Doppler spectrograms determination. These
Sec . 3.3, page 10
C. M. Michaux, JPL
November 15, 1971
JPL 606-1
Radar Properties
z-S
(a)
HARD
OUTLINE
--- VAGUE
OUTLINE
-«■ ALSORT
\ " ' J :'^^
. V- ELVSIUM 1 -T V-T
' TRIV1UM * ' "*
CHARONTIS ■" LACOONTIS 4 NEPENTHES
l^w
"yt^^'
y ,' NIX OLVMPICA "IT^
7m<
320 340 20 40
WEST LONGITUDE, d»o
(b)
Fig. 5. Radar cross-section variation with longitude near the 21 °N parallel
of latitude, as obtained at 12.5-cm wavelength
(Carpenter, 1967).
November 15, 1971
CM. Michaux, JPL
Sec . 3.3, page 1 1
Radar Properties JPL 606- 1
were spectrally analyzed by the "one-bit autocorrelation technique" of Goldstein
(1961). The sampling frequency in 1969 provided a small window only 4.8-kHz
wide, permitting 120-Hz resolution (as compared to 1 kHz in 1967). This proved
sufficient to resolve the narrow central peak in the quasi- specular portion of the
echo. This frequency resolution corresponded to 0.5° in longitude on Mars near
the subradar point; however, because of the rapid rotation of Mars during the
10-minute integration period (required to improve the signal-to-noise ratio), the
frequency resolution was degraded to ~2° in longitude. The spectrum obtained
was best-fitted by a theoretical spectrum based on Hagfors' exponential scatter-
ing law. This fit required adjustment of three parameters: slope parameter C
(rms slope =l/ \rC), frequency offset Af from ephemeris-predicted "zero fre-
quency shift" (effect of a tilted subradar region), and radar cross- section o-g .
The cross- sections obtained are plotted in Fig. 6. Additional cross-sections
were obtained as a "byproduct" of topographic studies using phase-coded CW.
These are presented in Fig, 7 for four latitudes from 3° to 22° N.
The small radar cross-section values obtained (0.02-0.04 ttR^) pertain
only to the central region of the disk around the subradar point, because of the
narrow-band windows used in both the CW (4.8-kHz) and phase-coded (1-kHz)
measurements, compared to the 23,5-kHz limb-to-limb Doppler spread. Thus,
a substantial portion of the total echo power from the full disk, or total radar
cross-section, was excluded. The measured relative radar cross- section is
almost exclusively due to the quasi-specular portion of the echo.
According to Pettengill et al. (1969), the quasi-specular portion of the
echo, at X3.8 cm, contributes from less than 0.006 to about 0.05 ttR^ to the total
cross-section of Mars, while the diffuse portion contributes a "background"
cross-section about 0.1 ttR^, largely independent of longitude. The total for
the entire disk averages 0.11 vR^ , or gp^ = 0.11, at X.3.8 cm, according to
them.
Using this value (cr = 0.11 ttR^) as a basis, these investigators recalcu-
lated the average total cross -section values, which previous observers would
have obtained, taking account of the limitations due to too narrow spectral
acceptance of the receiver (CW window), or poor delay resolution from too short
pulses. Table 2 lists these results for comparison. The new values generally
are in fair agreement with the older ones actually reported. Exceptions are
those of Evans et al. (1965) and especially of Kotelnikov et al. (1963).
In contrast to topography, the fractional radar cross- section, or the
reflectivity p if g = 1 , exhibits abrupt variations with latitude as well as
longitude. A? present, it is unknown whether the variations are due to local
surface roughness, electromagnetic properties of the material (dielectric con-
stant), or surface geometry.
If conductivity and permeability effects are neglected, an effective dielec-
tric constant k may be derived from 'the fractional radar cross - section (t/ttR ,
or reflectivity p^ (assuming g = 1), through the relation
\
L
M
Sec. 3.3, page 12 C. M. Michaux, JPL November 15, 1971
JPL 606-1
Radar Properties
10
0.08 -
c
o
o 06
o
2
O
P 0.04
O
LiJ
<J1
if)
CO
g 02
o
o
o O
o
O O o°
ho o o
^o ° ° o
o o
o o
o o
° o °o
o oo
%
o
o
on
O O ^po
o o
J 1 I I I I I I
<9o
J L
J L
40 80 120 160 200 240
MARTIAN LONGITUDE (west)
280
320
360
Fig. 6. Relative radar cross-section variation with longitude as inferred
from CW echo-Doppler spectrograms at 3. 8 cm wavelength
(bandwidth 4. 8 kHz) (Lincoln Laboratory, 1970).
1 1 1 1 1 1 ■ 1 1 1
• *l * *
'-W, :Ji^.J.-f I.'/''v.V-
•
L.T . .1- ^'
••• . '■ . .-••■ i
•
•
•1. J
L«T . S-
1 1 1 1 1 1 1 1 1 J 1
MARTIAN LONGITUDE ldtgl(M'«t»I
Fig. 7. Relative radar cross-section variation with longitude near four
latitudes, as inferred from phase-coded (ranging) measurements at
3. 8 cm wavelength (bandwidth 1 kHz) (Lincoln Laboratory, 1970).
November 15, 1971
C. M. Michaux, JPL
Sec. 3.3, page 13
Radar Properties
JPL 606-1
Table 2. Comparison of radar cross -sections obtained by various observers
(Pettengill et al., 1969).
Observer
Goldstein and GiUmore (1963), JPL
Koti-1'nlkov et al. (1963), USSR
Goldstein (1965); Sagan et al.
(1967), JPL
!:vans et al. (1965), .\aT/LL
Dyce et al. (1967), AIO
Pettingill et al. (1969), KOT/LL
(C^^' observations only)
Date
1963
1963
1965
1965
1965
19 67
Martian
latitude,
deg
13°N
14°N
21°N
zrN
22 °N
19 °N
Radar
wavelength,
cm
12.5
-40
12.5
23
70
3.8
Normalized^
reccnving
bandwidth
0.053
0.0016
0.48
0.125
0.20
1.00
0.04
Observed radar
cross section
max
(ra2)
0.07
0.15
0.16
0.13
0.14
0.05
mm
(Tra-i)
0.01
0.03
0.04
0.03
0.09
0.02
avE
(ira^)
0.032
0.07
0.087
0.14
0.07
0.11
0.031
Calculated^'
average
c ross
section
,Va2)
0.037
0.0014
0.097
0.067^
0.072^
0.11
0.0^4
^Normalized to the aopropriate limb-to-limb bandwidth.
''Calculati-d from the indicated normalized bandwidth of the receiver on the basis of the average cross
section and (typical) scattering law observed at 3.8 cm.
^A factor associated with resolution in delay has also been taken into account in these calculations.
Unfortunately, the 1969 measurements are not included in this table.
Using p = 0. 0915 (Pettengill et al. , 1969) leads to k = 3. 5, an intermediate
value between that of the Moon (~3) and of Venus (4. 5 at decimeter wavelengths).
This implies that the material is loosely compacted, like sand.
At JPL in 1969, Goldstein et al. ( 1970) performed delay- Doppler analysis
on the cchos from subradar areas. They used coded CW at a wavelength of
12. 5 cm, because they were primarily interested in topography. They looked
at only the central 640 Hz of the 7600-Hz limb-to-limb bandwidth, but had 10-Hz
resolution within the frequency range studied. The cross sections shown in
Fig. 8 are therefore relative cross - sections for local areas.
3. 3. 3 ANGULAR SCATTERING AND ROUGHNESS
Funda mental Concepts
Information on the average surface roughness of the planetary targets is
contained in the angular backs catter mg law specified by the a_yerage echo power
P(<?!)) versus angle of incidence (or reflection, since i = ^). P{^) , also caUed
the "angular power spectrum, " is usually derived from measurements of P(t),
the echo power versus delay function, through the uhiquee})— t correspondence,
discussed on pages 5 and 6, when short pulses (<1 msec) are transmitted. The
function P(4)) actually averages the echo power at a given incidence angle, as
contributed by the delay annulus corresponding to that angle. Therefore, P((}>)
represents a backscattering law which implicitly assumes a uniformly rough
surface over the spherical target.
Sec . 3.3, page 14
C. M. Michaux, JPL
November 15, 1971
JPL 606-1
Radar Properties
(E
Z
o
U
tlJ
tf>
v>
tn
O
CE
o
0.00
WEST LONGITUDE (DEGREES)
Note: The digits refer to latitudes according to the code:
3 = 3 °N, 4=4°N, ... 0=10°N, 1 = 11 °N, and 2=12 °N.
Fig. 8. Relative radar cross- section variation with longitude near
several latitudes 3 °-12 °N (Goldstein et al. , 1970).
An alternative, but more involved, procedure for obtaining P(4>) is by
derivation from the average echo power frequency spectrum P(Af), usually
written P(f ) (a "Doppler spectrogram"), through one of two mathematical trans-
formations available. These are (1) the Fourier- Bessel transformation as used
by Hagfors, Nanni, and Stone (1968), and (2) the inverse Abel transformation,
as used by Carpenter (1964). The latter transformation has also been employed
for Mars by Pettengill et al. (1969), at X 3. 8 cm.
Radar echoes P(t) or P(4>) fromi the terrestrial planets are composed of
two parts: (1) the "quasi- specular " portion--a strong central highlight from the
subradar region, which is attributed to near-normal reflections from large
smooth surface elements; and (2) the "diffuse scattering" portion- -a much
weaker, or subdued "background" from the surrounding concentric regions, as
far as the limbs, which is attributed to scattering from a large number of
smaller- scale elements (sometimes called the "diffuse scatterer s ").
There are two kinds of "surface roughnesses" associated with scattering.
One kind (diffuse scattering) is due to the small surface elements, such as
rocks strewn on the surface or buried in shallow depth. The larger of these
elements gives some normal backscatter, while the very small, and especially
the angular elements, can also give returns by diffractive scatter and multiple
November 15, 1971
C. M. Michaux, JPL
Sec . 3.3, page 15
Radar Properties JPL 606- 1
reflections. The stronger echoes (quasi- specular ) are from the frontal area
and large surface elements (those normal to the line of transmission).
Surface roughness is usually measured statistically by the rms slope
pertaining to a certain horizontal scale size. Slopes on a planet can be derived
from the radar data by means of several procedures, each yielding an rms
slope on scales larger than the wavelength. The three procedures which were
applied to Mars radar data are as follows:
Backscattering Model Fitting (for very small-scale slopes)
The central, quasi- specular portion of the echo may be approximated by
a theoretical backscattering law of either gaussian of exponential form, as was
done for the Moon. From experience, it was found that the following law, suc-
cessfully developed by Hagfors (1964) for the Moon, also provides a good fit for
Mars :
4 Z -3/2
P(0) oc (cos (^ + C sin 4>)
where C is a parameter related to the rms surface slope. For large C (as in
the case of Mars) the rms slope ^^l/x/C . This law assumes a uniform exponen-
tial distribution of surface slopes over the target (uniform roughness), and
applies to a horizontal scale size of about 1 to 10 times the probing wavelength
X. (For further elaboration, consult the Lunar Scientific Model, JPL Document
900-Z78, Radar section.) The particular parameter C value, providing the
best fitting curve (by least squares of residuals) to the planet's angular power
spectrum P(0), yields the rms small-scale slope. No applicable theoretical
model is available for the diffuse scattering portion of the echo.
Power Spectrum Frequency Offset Measuremient (for intermediate slopes)
The power spectra obtained from analysis of CW returns may show an
asymmetry of frequency displacemient of the central peak from the expected
(ephemeris-predicted) "zero-frequency shift" position (Af = 0). This frequency
offset may be interpreted as the effect of an overall general tilt of the subradar
region. This offset is then equal to
,, 2f' dr(t) 2f' , dh
c dt c dL
where f is the center frequency, c is the velocity of light, and r(t) is the dis-
tance of the subradar point from the center of mass of the target (Mars). The
quantity r is a function of time t because of the planet's (apparent) rotation at
rate w' . The apparent rotation differs somewhat from the true rotation of Mars,
because of the relative motion of Mars and the radar equipment. The quantity
L is longitude, while h is height (which also varies with time, of course).
The obtainable tilt or E-W slope dh/dL refers to a much larger horizon-
tal scale (60-120 km, if L = l°-2°) than the slope inferred from backscattering
modeling parameter C. With an improved signal-to-noise ratio, finer resolu-
tion can be obtained with this technique.
Sec. 3.3, page l6 C. M. Michaux, JPL November 15, 19V1
Radar Properties JPL 606-1
Height Profile Differentiation (for large-scale slopes)
The height- versus- longitude profiles obtained from ranging measurements
permit, by simple differentiation procedure (i.e. , taking derivative of topo-
graphy), derivation of E-W slopes on a still larger scale (e.g., 300-600 km
from L = 5°-10°). If a contour map of heights is available for several latitude
belts, the slopes in other directions may be similarly derived.
Experimental Results
Pettengill_et al. (1969), have provided a complete 0°-90° angular back-
scattering curve P(«^) for Mars at their operating wavelength of 3. 8 crn. They
derived P(<^) from a representative echo power frequency spectrumi P(f) by
using the Bessel transformation, see Fig. 9 (a and b).
Figure 9(b) compares the scattering behavior of Mars with that of the
Moon, Venus, and Mercury, utilizing their P(<^) curves at a wavelength of 3. 8 cm.
The fast drop of the Mars curve shows that Mars is definitely smoother than
Venus, which, in turn, is smoother than either Mercury or the Moon.
Using all available data (MIT at 3. 8 cm, JPL at 12. 5 cm, and AIO at
70 cmi), Zachs and Fung (1969) derived the P(<A) curves at three wavelengths
from the planetary averaged P(f) curves, which they also computed by using
Abel and Bessel mathematical transformations as a crosscheck. Their results
are shown in Fig. 10, where, for convenience, the processed P(f) curves are
also shown. They note that the 70-cm data pertained only to a small sec_tor
(208°-221°) of Martian longitude. Therefore, the derived 70-cm curve P(0) is
regional, and cannot readily be compared to the averaged planetary curves
obtained at 12. 5 and 3. 8 cm. The 12. 5- and 3. 8-cm curves do, however, show
variation with wavelength, invalidating the statement, by Dyce et al. (1967) and
Pettengill et al. (1969), that the Martian scattering law does not vary with wave-
length. These two curves exhibit scattering similar to that observed on the
Moon. Indeed, it appears that the fraction of echo power in the diffuse portion
P(4>) increases as the wavelength decreases. The 12. 5- and 3. 8-cm curves are
characteristic of planetary smoothness, while the 70-cm curve is either
anomalous or indicates regional roughness significantly different from that
obtained for the other curves.
Interpretation
The surface slopes for the north equatorial belt, derived by the MIT,
Lincoln Laboratory group in 1969, utilizing the three procedures outlined in the
earlier paragraphs on Fundamental Concepts, are as follows:
1) On a small scale of about 1-10 \, corresponding to parameter
C = 300, giving best fit using Hagfors' backscattering model,
an rmis slope of 3.3° was derived.
2) On the large scale of about 120 kms, corresponding to the ~2 °
longitude resolution of the CW measurements, an rms slope of
0. 5° was derived.
November 15, 1971 C. M. Michaux, JPL Sec. 3.3, page 17
Radar Properties
JPL 606-1
(a)
\
1 1
1 1 1 1 1 1
1 1 1 1
P(f) X= 3.8 cm
1
\
MARS
:
V
-
V
.
-
X
234.2°
1 1
AVERAGE
1 1 1 1 1
,1,1
5
10
FREQUENCY (kHz)
(by Bessel
transform)
o
<
-1
lij
a:
15 20 25
<t> <deg)
ANGLE OF INCIDENCE OR REFLECTION
Fig. 9. Frequency power spectrum for Mars converted by Bessel
transformation for comparison with the Moon, Mercury, and Venus
Top (a): Pettengill et al (1969). Bottom (b): Evans (1969).
Sec. 3.3, page 1<
C. M. Michaux, JPL
November 15, 1971
JPL 606-1
Radar Properties
\^ X-70cnn
f/f (normalized frequency)
(by Abel or
P(<^)
-25-
Bessel transformation)
t
MARS
• MOON
X=3.6cm
CD
I -10
o
o.
E -15
1-
^ -20h X = 12.5cm
X=3.8cm
>^ X=23cm
^V ****** ^
^2^=68cm
X=70cm
1
0° 10° 20° 30° 40° 50°
Angle of incidence
60°
-0
Fig. 10. Average doppler spectrograms (top) and angular backscattering curves
(bottom) for Mars at several wavelengths (Zachs and Fung, 1969).
November 15, 1971
C. M. Michaux, JPL
Sec. 3. 3, page 19
Radar Properties JPL 606-1
3) On the still larger scale of about 180 kms, corresponding to the
~3 ° longitude resolution of the ranging measurements, an rms slope
of 0. 3° was derived.
3.3.4 TOPOGRAPHY
Fundamental Concepts
The fundamental method used to measure Martian topography is the
round-trip delay technique. By very accurate timing of the strong central echo,
using either narrow pulses or phase-coded CW, it has been possible to obtain
elevation differences. Typically, the measurement made refers to the mean
return from a small spherical front cap a few areocentric degrees in size.
In such topographic ranging measurements, it is necessary to know the
accurate relative orbits (or ephemerides) of both planets. It is also necessary
to accurately determine the center, the rotation rate, radius, and the degree
of flattening (shape). Since the (equatorial) radius and orbit of Mars are not
accurately known, ranging measurements are also used to obtain (by solution of
simultaneous equations) better determination of these quantities.
The topography may be checked at "closure points, " v/hen available.
Closure points are defined as a pair of altitude (or delay) observations''' made
at very close or coinciding locations on the planet. However, these must be
made at different times, separated by approximately one or more synodic
rotations of the two planets (i.e. , ~40 days in the case of Mars). Any important
discrepancy noted between these two altitudes leads to a correction of the rela-
tive orbits (or ephemerides) and serves as a check on the topography.
The zero level of topography, as used on the illustrations and (for one
latitude strip around the planet), is obtained arbitrarily as the average of heights.
When many such strips have become available for each accessible latitude over
the years 1965 to 1975, a grand average of heights will furnish a better zero
level.
Experimental Results
An early attempt was made in 1965 by Dyce et al. (1967) at Arecibo to
derive topographic information for Mars. They were only able to set an upper
limit of about 15 km for the greatest elevation difference, at about 17 °N latitude.
The MIT efforts of Pettengill et al. (1969) in 1967 and Rogers et al. (1970)
in 1969 used 3.8 cm phase-coded CW for their round-trip delay topographic
studies. The 1969 work involved an elaborate least- squares fit to a theoretical
delay profile, based on Hagfors' (1964) backscatter ing law, found adequate from
previous experimentation. The fit involved adjustment of three parameters:
slope parameter C, radar cross - section (t, and delay to the subradar point.
=■' The time-delay "residuals" (or difference between observed and theoretically
predicted Earth-Mars displays) are easily converted into altitudes by multiply-
ing by - 1/2 the velocity of light or by -0. 15 km per (xsec.
Sec. 3.3, page 20 C. M. Michaux, JPL November 15, 1971
JPL 606-1 Radar Properties
Reduction of the entire set of delay residual data for 1967 and 1969 produced
the relative altitude- versus -longitude plots for four separate 3°-4° wide north
equatorial latitude belts, centered on 3°, 7°, 11°, and 22 °N, as showm in
Fig. 11. The spatial resolution for each plotted point on the surface is about
5 ° in longitude or latitude, or 300 km. The error bars are considered conserva-
tive estimates; the measurement repeatability was found to be within 200 meters
rms .
There is a striking similarity between the four topographic profiles.
The topographic variation in the north equatorial belt, although large in magni-
tude, changes only slowly, contrary to the variation in radar reflectivity. The
profile at 22 °N differs from the 11° and 7 °N profile s, mainly because of the
broad (~30° wide), high (~5 kmi) Elysium bright area. This tends to confirm
the statemients often made by Capen (1966) and Binder (1969), that Elysium is
a high plateau. Near 40° longitude, the Chryse • Xanthe desert is a broad dep-
ression extending northward and apparently including part of Niliacus Lacus, a
dark patch just south of Mare Acidalium. At 285° W longitude, the dark area
Syrtis Major exhibits a rather abrupt increase in elevation (~5 km) over about
ten degrees toward the westward bordering bright area, Aeria, The 7°N
profile shows a minor peak at 250° longitude, in the Aethiopis desert. This peak
does not appear in the other profiles. Contrary to Binder's (1969) conjecture,
no indication of depressions were detected for the Phison and Euphrates canals
located in the Arabia desert, near 330° longitude.
The profiles do show narrower structure, but because of the 5° spatial
resolution and measurement accuracy, it is doubtful that they actually represent
local variations in elevation.
The highest region found in all three profiles was in Tharsis, at about
100° longitude. The maximum elevation difference obtained was ~12 km.
During the 1969 opposition, the JPL team of Goldstein et al. (1970) used
phase-coded 12. 5-cm CW radar for delay studies of Martian topography. They
looked at five different range gates simultaneously, each gate spaced 750 m
apart in range. The output of each range gate was sampled every 1/640 second,
the autocorrelation function calculated by a computer, and then Fourier-
transformed to yield power spectra, 640-Hz wide (and resolution of 10 Hz),
which finally were averaged over the 9 minutes round-trip flight time (integra-
tion time), and displayed on an x-y plotter. This real-time display permitted
observers to follow the large changes in topography (some 12-km variation,
which is several times greater than the 4.5-km "field of view" of the set of five
range gates), by properly adjusting the delay. The computer could only handle
five range gates simultaneously. The 640-Hz width is the largest "effective"
Doppler spread in the subradar region, while the limb-to-limb bandwidth is
7600 Hz, at \12.5 cm. The range (to the subradar point or center of each range
zone) was estimated by using cross-correlation techniques and comparing each
power spectrum obtained with a set of (30) theoretical power spectra, computed
for a planet of the same radius, apparent rotation rate, and with a surface
roughness assumed to give a similar exponential scattering law at very small
angles ('^<2°). The limiited range resolution (0.3 km) and Doppler resolution
(10 Hz) was also compensated for in this computation. The best fit gave the
best estimiate of range.
November 15, 1971 C. M. Michaux, JPL Sec. 3.3, page 21
Radar Properties
JPL 606-1
I
o
LAT
8oc^
i \ L
90 160 270 360
MARTIAN LONGITUDE (deg)(West)
Fig. 11. Topography variation with longitude in four latitude
steps from 3° to 22°N (Lincoln Laboratory, 1970).
Sec. 3.3, page Z2
C. M. Michaux, JPL
November 15, 1971
JPL 606-1
Radar Properties
The JI-^L results (Goldstein et al. , 1970) of relative altitude- versus -
longitude were given in one plot (see Fig. IZ), with subradar data points varying
in latitude from 3 ° to 12 °N. Each point represents an average measurement
over a rectangular area ~90 km wide in a N-S direction by ~220 km in the E-W
direction. They are very similar to tht: MIT results given in previous para-
graphs. Tharsis was found to have the highest elevation, with Aeria next highest,
and neighboring Syrtis Major dropping by 4 km "nearly linearly across its full
width from Aeria to Moeris Lacus" (its eastward border). Mare Cimmerium ' s
northern tip (~3 ° latitude) is lower. The lowest region found was in the large
Amaz(jnis desert, just west of Tharsis. The maximum elevation difference
o bta i n e d wa s ~1 1 km .
3.3.5 TOPOGRAPHY - CROSS-SECTION - ROUGHNESS CORRELATION
The MIT (Lincoln Laboratory, 1970) results of 1967 and 1969 on topo-
graphy, crrjss - section, and roughness were compared with de Vaucouleurs'
(1967) luminance results from his photometric map, which averaged 1941 and
1958 data, in an attempt to establish any possible correlation. The mathemati-
cal techniques of correlation (cross -correlation functions) were systematically
applied to these data by the MIT group (see Pettengill et al. , 196 9).
The results of the correlation analysis were as follows:
1) Topography is not correlated with cross - section cjr luminance,
and is probably not correlated with roughness.
2)
Cross- section is anticorrelated with both luminance and roughness.
+ 50
180 140 100 60 20 340 300
LONGITUDE (DEGREES) (WEST)
260
220
180
Note: The digits refer to latitudes according to the code;
3=3 °N, 4=4°N, ... 0=10°N, 1=11 °N, and 2=12 °N.
Fig. 12. Topography variation with longitude in
latitudes 3 °- 12 "N (Goldstein et al. , 1970).
November 15, 1971
C. M. Michaux, JPL
Sec. 3.3, page 23
Radar Properties JPL 606-1
3) High cross- sections and low luminances (visually dark areas)
tend to be associated with west-rising slopes. This conclusion,
however, may be influenced by the fact that the only large cross -
section regions mapped, Syrtis Major and Trivium Charontis,
are both west-rising slopes.
These results apparently support the view that dark areas produce a
strong quasi- specular echo, and suggest that high cross -section regions are,
on the whole, smoother than their surroundings. Also, results indicate that
dark areas tend to lie on the eastern slopes of highlands, which is consistent
with the wind-blown model defining dark areas as bare rock (exhibiting strong
radar reflectivity), and the bright areas as dust- covered surfaces (exhibiting
low radar reflectivity).
Sec. 3.3, page 24 C. M. Michaux, JPL November 15, 1971
JPL 606-1 Radar Properties
BIBLIOGRAPHY
The American ephemeris and nautical almanac, 1963, 1965, 1967, 1969, 1971:
Wash. , D. C. , U. S. Government Printing Office.
Binder, A B. , 1969, Topography and surface features of Mars: Icarus, v. 11,
no. 1, p. 24-35, July.
Capen, C. F. , 1966, The Mars 1 964 -1 965 apparition: JPL-TR 32 -990 (187 p. ),
December 15.
Carpenter, R. L. , 1964, Study of Venus by CW radar: Astron. J. , v. 69, no. 1,
p. 2-11, February.
Carpenter, R. L. , 1967, Radar observations of Mars: p. 157-160, in JPL-SPS
37-48, Vol. Ill (Supporting research and advanced development for the
period October 1 to November 30, 1967), December 31.
de Vaucouleurs, G. , 1967, A low^-resolution photometric map of Mars: Icarus,
V. 7, p. 310-349.
Dyce, R. B. , Pettengill, G. H. , and Sanchez, A. D. , 1967, Radar observations of
Mars and Jupiter at 70 cm: Astron, J. , v. 72, no. 6, p. 771-777, August.
Dyce, R. B, , 1965, Recent Arecibo observations of Mars and Jupiter: Radio
Sci. -J. Res. NBS, v, 69D, p. 1628-1629.
Evans, J. V. , 196 9, Radar studies of planetary surfaces: in Ann. Rev. Astron. &
Astrophys. , v. 7, p. 201-248.
Evans, J. V. and Hagfors, T. , 1968, Radar astronomy: New York, McGraw-
Hill Book Co,
Evans, J. V. , Brockelman, R. A. , Henry, J. C. , Hyde, G. M. , Kraft, L. G. ,
Reid, W. A. , and Smith, W, W. , 1965, Radio echo observations of Venus
and Mercury at 23 cm wavelength: Astronom. J. , v. 7 0, no. 7,
p. 486-501, September.
Goldstein, R. M. , 1961, Amplitude modulated system: p. 40-44, in Chapter 4,
Radar exploration of Venus: Goldstone Observatory Report for March-
May 1961: Victor, W. K. , Stevens, R. , and Golomb, S, W. , Editors :
JPL-TR 32-132 (103 p. ), August 1.
Goldstein, R. M. , 1965, Mars: radar observations: Science, v. 150,
p. 1715-1717.
Goldstein, R. M. and Gillmore, W. F. , 1963, Radar observations of Mars:
Science, v. 141, p. 1171-1172.
Goldstein, R. M. , Melbourne, W. G, , Morris, G. A. , Downs, G. S. , and O'Handley,
D. A. , 1970: Preliminary radar results of Mars: Radio Sci. , v. 5, no. 2,
p. 475-478, February.
November 15, 1971 C. M. Michaux, JPL Sec. 3. 3, page 25
Radar Properties JPL 606-1
Hagfors, T. , 1964, Backscattering from an undulating surface with applications
to radar returns from the Moon: J. Geophys. Res. , v. 69, p. 3779-3784.
Hagfors, T. , Nanni, B, , and Stone, K. , 1968, Aperture synthesis in radar
astronomy and some applications to lunar and planetary studies: Radio
Sci, , (New Series), v. 3, no. 5, p. 491-509, May.
Kotel'nikov, V. A. , et al. , 1964, Radar studies of the planet Mars in the Soviet
Union: Sov. Phys. -Dokl. , v. 8, no. 8, p. 760-763, February. Translation
of 1963 article.
Lincoln Laboratory (MIT), 1970, Radar studies of Mars: Final Report (79 p. ),
January 15.
Pettengill, G. H. , 1965, A review of radar studies of planetary surfaces:
Radio Sci. -J. Res. NBS, v. 69D, p. 1617-1623.
Pettengill, G, H. , Counselman, Rainville, L. P. , and Shapiro, I. L , 1969, Radar
measurements of Martian topography: Astron. J. , v. 74, no. 3,
p. 461-482, April.
Rea, D. G. , Hetherington, N. , and Mifflin, R. , 1964, The analysis of radar
echoes from the Moon: J. Geophys. Res. , v. 69, p. 5217-5223.
Rogers, A. E. E. , Ash, M, E, , Counselman, C. C. , and Shapiro, L L , 1970, Radar
mieasuremients of the surface topography and roughness of Mars: Radio
Sci., v. 5, no. 2, p. 465-473, February.
Sagan, C. , Pollack, J. B. , and Goldstein, R. M. , 1967, Radar doppler spectro-
scopy of Mars. I: elevation differences between bright and dark areas:
Astronom. J. , v. 72, no. 1, p. 20-34, February.
Skolnik, M, L , 1962, Introduction to radar systemis: McGraw-Hill Book Co. ,
(648 p. ).
Westman, H. P. , 1956, Reference data for radio engineers; Editor: New York,
International Telephone and Telegraph Corp.
Zachs, A, and Fung, A. K. , 1969, Radar observations of Mars: Space Sci. Rev. ,
V. 10, no. 3, p. 442-454, Decemiber.
Sec. 3.3, page 26 C. M. Michaux, JPL November 15, 1971
JPL 606-1 Radar Properties
APPENDIX
MARS RADAR OBSERVATIONS IN 1971: TOPOGRAPHY AND
RADAR CROSS-SECTIONS
During the very favorable opposition of 1971, two groups of investigators,
at MIT and at JPL, using the same techniques, have made extensive measure-
ments of Martian topography and radar cross -sections in the Southern Hemis-
phere (equatorial belt 14° to 18°S), utilizing improved radar and computer
processing capabilities. Both groups were able to resolve some of the larger
craters, and also some steep scarps, crater rims, etc. Observations were
begun in June 1971, and were continuing into November, as of this writing.
Only preliminary results of their work are available and are presented here.
MIT Observations (Pettengill et al. , 1971)
The MIT group at Haystack operated at 3. 8 cm, transmitting phase-coded
CW. Lateral surface resolution was 1.3° in latitude and 0.8° in longitude.
Range resolution was of 0. 90 km. Repeatability reached 75 meters at best
when reflectivity was high. Results given concerned the topography mainly.
Figs. A- 1(a), (b), and (c) show the altitude variations around Mars at latitudes
from 14. 5° to 16. 8°S. They are very similar to those obtained by Goldstein and
co-workers (in 1971). Many craters also were resolved, as well as rims and
scarps. In particular, there is a sudden drop of 4. 5 km in Pyrrhae Regio near
45 °W (lowest depth encountered). Then, an extremely radar bright area (high
radar cross - section) near I5°W in Deucalionis Regio. Three abrupt transi-
tions (A, B, C) near 345 °W obviously correspond to craters (one B with depth
1 km). Between 310° and 275 °W in lapygia they noted an apparently very large
(2000 km) basin some 2 km deep, with another nearly concentric crater 500 km
across and 0.5 km deep. There is a nearly level plain from 258° to 243 °W with
a few small craters. At 230° W, a large crater 340 km across, 1 km deep (D)
with a well developed lip is noted. It corresponds to a crater seen in a Mariner
1969 FE photograph. From 198° to 183 °W there is a large depression (E and F)
with a very irregular bottom, in Zephyria. At 167 °W, a small, 2 km deep
feature (G) was obtained in Titanum Sinus. At 149 °W, there is a crater (H).
Then, the remarkable "twin peaks" at 122° and 100°W in Phoenicis Lacus, the
latter displaying extremely high radar brightness (strongest at 108°W). These
features seem to continue in the South Hemisphere from those seen in the North
Hemisphere, and, may thus form a double ridge of major importance. Between
125° and 121 °W no echoes could be obtained. Either the region is exceedingly
rough (at wavelength scale) or of exceedingly low radar reflectivity or both.
From 98° to 80 °W there is a very regular downslope from the peak with strong
reflectivity, which is apparently free of craters.
One surprising general result was the lack of correlation between areas
of high radar reflectivity and visually dark features (as found in the north
equatorial belt).
December 1, 197J C. M. Michaux, JPL Sec. 3.3, Appendix, page 1
Radar Properties
JPL 606-1
7
6
5
4
? '
~- 2
f
.S? 1
a>
• Of
>
•I
'^ -2\-
-3
-4
-5
-6
-1 — I — I — I — 1 — r
-I 1 . 1 r-
-i — . — I — ■ — I — 1 — r
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il '1.
■ w
,■>'•■'
^ V'lr'"'^-"'
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Lat: -16.5' t''
*0.T
,a_._i I 1 I 1 I 1 1 j_
60" 50' ' 40" 30° 20° 10* 0° 350° 340° 330" 320° 310° 300*
West longitude
Fig. A-l(a). Topography variation with longitude near 16.5 °S latitude
(Pettengill et al. , 1971) - (MIT at \3.8 cm).
8h
7
6
5
4
? '
-^ 2
-i — , — I — I — I — < — 1 — ' — \ — i — I — ' — r — ' I '~~\ ' I ' I ^
lapygia MareTyrrhenum Hesperia Mare Cimmerium Zephyria _
-1 -
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-3 -
-4 -
-5 -
-6 -
. 1
V.
Lat: -15.8' *"'!
II
'I '■ '. «■, H
_, I . I . I . I . 1 I I L- 1 J 1 . 1 1 1 > L
300* 290* 280° 270° 260* 250° 240° 230° 220° 210° 200° 190° 180°
West longitude
Fig, A-l(b). Topography variation with longitude near 16° S latitude
(Pettengill et al. , 1971) - (MIT at X3.8 cm).
Sec. 3. 3, Appendix, page 2 C. M. Michaux, JPL
December 1, 1971
JPL 606-1
Radar Properties
8
' 1
I '
I ' I '
1 '
.t ' I ' I •
t
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7
-
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-6
. 1
,
1 .
111.
1 1
I.I.I."
180* 170' 160* 150* 140* 130* 120' 110* 100' 90° 80' 70' 60*
West longitude
Fig, A- 1(c). Topography variation with longitude near 1 5. °S latitude
(Pettengill et al. , 1971) - (MIT at \3.8 cm).
December 1, 1971
C. M. Michaux, JPL Sec. 3. 3, Appendix, page 3
Radar Properties JPL 606-1
JPL Observations (Downs at al. , 1971)
The JPL group at Goldstone operated at 12,5 cm with a 300 kW trans-
mitter and a 20°K system temperature. Using delay-Doppler mapping on the
subradar region, they recorded every 30 seconds (integration time) an array
of 32 delay (rings) by 64 Doppler (strips) resolution cells. The delay rings gave
a resolution of 0.45 km in range, while the Doppler strip resolution was 9.4 km.
The lateral surface resolution was ~100 km, corresponding to a front-cap diam-
eter -1.6°. Repeatability of altitudes obtained reached 40 meters, whenever
the points were very close and the echo strong.
Results were of two sorts:
1) Topography. Large craters and small craters were resolved with
occasional resolution of their rims. Figure 2(a) shows the altitude
profile obtained around Mars at latitudes from 13.8° to 14.6 °S. It
is similar to that obtained in northern equatorial belts. There is
a difference in altitude of some 13 km between lowest (120°W) and
highest points (85 °W), Slopes are of the order of 1°. The fine
structure is due to craters with some large craters (at 230° and
188°W) identifiable on Mariner 1969 FE pictures.
2) Rada r cross - section . Wide fluctuations are noted as before, due
either to variations in dielectric constant and/or surface rough-
ness. Crater floors are bright but their crater walls are not.
Figure 2(b) shows the fractional radar cross -section around Mars
at the same latitudes as Fig. 2(a).
BIBLIOGRAPHY
Downs, G.S., Goldstein, R.M. , Green, R.R. , and Morris, G. A. , 1971, Mars
radar observations; a preliminary report: Science, v. 174, no. 4016,
p. 1324-1327, December 24.
Pettengill, G.H. , Rogers, A. E.E. , Shapiro, I.I. , 1971, Martian craters and a
scarp as seen by radar: Science, v. 174, no. 4016, p, 1321-1324,
December 24.
Sec. 3.3, Appendix, page 4 C. M, Michaux, JPL December 1, 1971
JPL 606-1
Radar Properties
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C. M. Michaux, JPL Sec. 3. 3, Appendix, page 5
JPL 606-1 Chemical and Physical Properties
3.4 CHEMICAL AND PHYSICAL PROPERTIES
INTRODUCTION
So far, very little is known about the chemical, physical, mineralogical
or petrological nature of the Martian surface. The observational information
available to date has been gained solely by means of remote -sensing methods:
spectrophometry from UV to IR (and filter colorimetry), polarimetry (mostly
in the visible), infrared and microwave radioraetry, and radar probing. Of
these methods used from Earth, only UV and IR spectrometry, filter colorime-
try (in the visible), and infrared radiometry have been implemented on board
spacecraft (Mariners 6 and 7), and used in conjunction with the television
imagery.
Beyond the limited observational information which provides some guide-
lines, present "knowledge" is speculative. Speculations about the nature of the
Martian surface are based on limited knowledge and linked to speculations about
types of geological processes which are and have been operative for some time
at or near the surface. Thus, considerable uncertainty prevails and probably
will continue to prevail until landers return more exact and detailed information
from the surface itself.
Presented here is the factual and interpreted information obtained from
remote -sensing and also some speculations derived from them.
This section contains two separate portions: one dealing with the physical
and chemical properties of the actual ground or soil, the other with the chemical
properties of the polar cap deposit.
PHYSICAL PROPERTIES OF THE GROUND SURFACE MATERIAL
The top layer of Martian surface material in bright and dark areas is
expected to be granular, much of it fine-grained and not particularly cohesive.
Atmospheric gases (CO2, CO, H2O) must be adsorbed in this porous soil,
which appears to be stirred at least sometimes during the Martian year.
Large eroded blocks of rocks (boulders) are also expected near the older
craters, but partly or completely buried in the fine material and the impact
rubble which no doubt forms most of the epilith ( regolith) of Mars. Smaller
angular blocks are probably present also, near the more recent small craters.
The Mariners did not have sufficient resolution to provide direct photographic
information on the small-scale appearance of the epilith. In any case, erosion
has been much more active than on the Moon, as can be seen by comparing the
Martian and lunar craters (see Section 3. 6, Photographic Atlas), and the fine
material has been distributed widely over the whole planet, mostly by aeolian
transport and deposition. '
The bearing strength of the granular dry Martian "soil" is probably lower
than that of lunar soil, which is on the order of 2 - 5 X 105 dynes cm"^.
Cohesion at the top is lower because there is no high vacuum nor particle
sintering effects; at little depth, it must rapidly increase from packing and
December 1, 1971 C. Michaux, JPL Sec. 3.4, page 1
Chemical and Physical Properties
JPL 606-1
perhaps some humidity. Bearing strength varies with grain size distribution
as well as mean grain size and packing, all of which are surely governed by
the aeolian transport and deposition processes. It is thought that one of the
important differences between bright and dark areas lies in the mean grain
sizes. See subsection on Granularity, where other physical properties, such
as density and dielectric constant estimates are also given. Table 1 summa-
rizes the best estimates of the miost important physical parameters of the
Martian top material.
Table 1. Estimates of average physical properties of Martian surface material
Bright Areas
Dark Areas
A*
-Bolometric Albedo (assumed)
(dimensionless)
0. 25
0. 15
-1
Y
-Thermial Inertia
-1 ,^ r^\V2'
V = (KpC) '
0.004
0. 006
/I -2 -1/2 , -K '
(cal cm sec ' deg )
P
-Density
(g cm" )
(1.2)
(2)
C
-Specific Heat Capacity (assumed)
0. 15
0. 15
(cal g' deg' )
K
-Thermal Conductivity
/I -1 -1 ^ -i\
(cal cm sec deg )
5. 10-5
12. 10-5
^Ao
Dielectric Constant
(dimensionless)
(2.5)
(3.5)
d
Mean Particle Diameter
(miicrons)
50
200
Granularity
The bright areas of Mars are most probably constituted of fine, granular
material of high porosity. Many lines of evidence converge to support this
statement. First, the visual observations of yellow clouds indicate they almost
always arise over bright areas; second, the colorimetric observations of both
these clouds and bright areas are sin-iilar; third, the polarimetric observations
show a negative branch at small phase angles; fourth, the spectrophotometric
observations of both the yellow clouds and bright areas are similar, with a
characteristic decline in albedo (reflectivity) from the red to the blue; fifth, the
analyses of yellow cloud fallout times are in accord with expected particle sizes;
sixth, the infrared radiometric observations of bright areas yield thermal
inertias characteristic of granular, porous material; and finally, the radar
observations sho-w low dielectric constants for bright areas, indicative of
pulverized material to a depth of at least a meter. To this series of evidence
Sec. 3. 4, page 2
C. Michaux, JPL
December 1, 1971
JPL 606-1 Chemical and Physical Properties
secured from Earth, may be added the ones returned by the Mariner 6 and 7
spacecraft, which have been discussed under the heading Dark and Light Areas
in Section 3.5 on Morphology and Processes.
The Earth-based observations agree rather well in consistently indicating
a mean particle diameter of about 50 \i for the bright areas, which means fine
dust is covering them. Such dust is certainly capable of being transported by
winds (to produce the patterns over crater floors in the Meridiani Sinus region,
as revealed by Mariners 6 and 7), and even, provided winds are strong enough,
of being lifted up to form yellow clouds, as Ryan (1969) concluded from a study of
the dynamics of dust-devils, or Sagan et al. (1971) from Martian wind regimes.
For the dark areas, the evidence of granularity is less certain. Photom-
etry, polarimetry, and infrared radiometry indicate somewhat larger grains --
about 200 \i in average--for the dark areas. On the other hand, radar shows
higher dielectric constants, which independently suggest either compacted
granular material or possibly solid rock, which, of course, could exist just
below the granular layer.
Density
Tentative estimates of densities (and porosities) can be obtained from the
radar observations, if one is willing to make many uncertain assumptions.
hence the speculative nature of such estimates. Indeed, there is a formula
(Rayleigh's) which relates density to the dielectric constant of pulverized and
nonpulverized material. If one assumes a pow^dery limonite layer for the
bright areas, then the 'observed' dielectric constant can yield the density.
Pollack and Sagan (1970) thus obtained a density of 1 . 2 g cm"3 for Amazonis
(using Goldstein [1965] radar observations giving a dielectric constant of 2.6).
For the dark areas, they assumed compacted limonitic material (a questionable
assumption), and from the higher dielectric constant (3. 5-4. 5) observed by
radar over Syrtis Major, they derived a density of 1. 5-2, which they considered
probably a lower limit.
CHEMICAL PROPERTIES OF THE GROUND: GENERAL ASPECTS
The coinposition of the Martian surface may be quite different from that
of the Earth (and perhaps more akin to that of the Moon), for two basic reasons:
(1) the starting materials of this small, apparently incompletely differentiated
planet (see Section 2 on Interior) may have been quite different from those of
Earth, and (2) aqueous erosion (flowing water) and weathering have probably
not been a dominant process (see Section 3. 5 on Morphology and Processes) in
the last few (perhaps four) billion years of Mars' history. On Earth, aqueous
erosion and weathering have produced an immense variety of sedimentary rocks
differing in type and composition, and often later these rocks have been trans-
formed by metamorphism; i. e. , action of pressure and temperature at depths.
It is not too surprising, therefore, that recent evidence from our most informa-
tive remote -sensing method from Earth (spectrophotometry) indicates that the
surface rocks or soil of Mars as a whole appear to be predominantly igneous,
with iron-bearing basalt the fundamental type of rock. Very slow chemical
weathering (possibly oxidation by traces of O2 or O3 catalyzed by minor H2O
available) has most likely stained their surfaces with ferric oxides of various
degrees of hydration. This will be treated in detail under "Composition Inferred
December 1, 1971 C. Michaux, JPL Sec. 3.4, page 3
Chemical and Physical Properties JPL 606-1
From Reflectance Spectrophotometry. " Other products of weathering, such as
carbonates (whose formation may require catalysis by liquid H2O), might be
expected. From the IR spectra, however, carbonates seem very low in
abundance (as none of the strong carbonate bands in the 1. 7-2. 5 )jl region has
been detected so far). Products of reaction with a perhaps N2-rich former
atmosphere, namely nitrates, have likewise not been detected (but their
spectral lands are very weak anyway).
There should be an abundance of products of asteroidal or meteoritic
cratering and ablation (by the atmosphere): metals and oxides, glass spherules,
shocked minerals, etc. , besides the common meteoritic silicates in the
probably thick rubble -and-dust epilith (regolith) layer of Mars.
K th'? crust has been geologically active with mild volcanism (outgassing,
hydrothermal action, . . ) then sublimates (e.g. , sulfur) and hypogene minerals
(e.g. , sulfide ores) may occur in favored localities in the upper crust. If
magmatic intrusions have taken place, contact metamorphic minerals would
then also be present. If the crust has been active with tectonism, such as
faulting, thrusting, . . , then one could expect cataclastic metamorphism locally;
and, with large orogenic uplifts, regional dynamothermal metamorphism in the
deeper parts of the crust. Of course, these possibilities are mere speculations .
Likewise, and even more speculative is the case: if, early in the history
of the planet, there was once a wet age on Mars (ocean or dense watery atmo-
sphere), then deep sedimentary layers should be present in parts of the crust,
and b< . ause of later history (impacting, etc. ), near surface rocks may show
strong compositional variations. One might still find evaporite minerals (salts,
such as Na, Mg chlorides, sulfates, . . ) in remnants of ancient marine beds.
3.4. 1 Co-np sitio;. Inferred From Reflectance Spectrophotometry
So far, the only remote -sensing method effective in providing definite
clues to the composition of the Martian surface has been reflectance spectro-
photoi-netry. The reflectance spectra of rocks and minerals in the 0. 2 to 2. 6 (i
range exhibit two types of absorption band: (1) broad bands resulting from
electronic transitions, and (2) narrow bands due to vibrational modes.
An example of the first type is the ferrous Fe ion band, occurring at
about lfJ.(0.9tol.l|jL)in the mafic silicates, while exaraples of the second
type are the 1.4 and 1 . 9 p. sharp bands due to the hydroxyl OH" stretching
vibration, as occurring in quartz, feldspars, and other light-colored minerals ,
and in some hydrous minerals. Such bands in a reflectance spectrum, either
individually or in combination, can be used to identify minerals or to narrow
do-zm their possible choices. The position of the band also depends on the
crystal field symmetry.
In recent years, the reflectance spectra of many rocks and minerals have
been obtained; for example, by Hovis and Callahan, 1966; Adams and Felice,
1967; Hunt and Salisbury, 1970; Greenman et al. , 1967; Adams, 1968; and
Ross et al. , 1969. Some of the results, other than those mentioned above, are
given in the listing that follows (see Adams, 1968).
Sec. 3.4, page 4 C. Michaux, JPL December 1, 1971
JPL 606-1 Chemical and Physical Properties
1) Most of the common iron-bearing silicates (olivines, orthopyroxenes ,
clinopyroxenes) have characteristic spectra, with a major band
between 0.9 and 1 . 1 [j. due to Fe + + (usually in sixfold coordination,
and another band between 1 . 8 and 1 . 9 fj. due to Fe + + , residing
"probably in a highly disordered octahedral site" (Bancroft and Burns
1967; White and Keester, 1967).
Z) Ferrous oxides (magnetite, ilmenite) are opaque and show no
structure .
3) Ferric oxides (hematite, goethite, limonite*) exhibit an absorption
band at 0. 85 to 0. 89 y. due to ferric ion Fe + + + in sixfold coordination.
Goethite and limonite have no OH" band.
4) Iron-bearing carbonates have both a broad band near 1 p. and narrow
bands between 1.7 and 2.6 |jl, which are quite characteristic.
Mixing of different minerals in a rock tends to average the spectra; for
example, a basalt combining orthopyroxene ( 0. 9 h- band) and olivine ( 1 02-1 05 fa
band) would show a band only at 0.95 ^l. Lack of structure is ambiguous and may
be due to very small particle size (<10 ^) or glass (as found in lunar samples)
The particle size effect, however, may act the other way (i.e. , small size may
enhance structure), depending on the opacity or transparency of the mineral
(See Salisbury and Hunt; 1968, 1969. ) Thus, the identification or narrowing
down of possibilities is often a difficult, if not impossible, matter. It is also
complicated by the presence in actual reflectance spectra of additional bands
due to Mars atmospheric constituents, such as CO2 and H2O, as well as similar
(telluric) bands due to the Earth's atmosphere, which may overlap or block out
the characteristic bands of the minerals to be identified. Fortunately, these
atmospheric bands are well known for both Mars and Earth (see Section 5. 1).
3. 4. 2 Reflectance Spectra of the Bright and Dark Areas (Earth-based results)=;":=
The spectral reflectivity curves or the spectral geometric albedo curves
for Martian bright areas and dark areas in general are now well established,
with an accuracy of a few percent, in the 0.3 to 1.1 [j. spectral region. Refer to
the review articles by McCord and Adams (1969), and especially to McCord et al.
(1971), which gives the most recent results. These results were obtained at the
time of the I969 opposition by McCord and Westphal (1971), at Cerro Tololo
Inter -American Observatory (Chile), by means of a narrow-band spectro-
photometer using 52 interference filters spanning from 0.3 to 2.5 [i. Figure 1
summarizes the best available data for dark and bright areas. Figure 5 of Se
tion 3.2 represents the McCord and Westphal (1971) curves (0.3 - 1 . 1 fx) for
seven areas. Seasonal changes in dark areas were also investigated in the
0.3-1. 1 |j. region, by McCord and Adams (1969) for Syrtis Major (see Fig. 2).
)ec -
)r
*Limonite is a mixture of hematite, goethite, water of hydration and many
impurities.
^♦Regretfully, no Mariner I969 IRS final results on the soil composition of the
scanned Martian surface have been available (January 1972). The presence of
silicates has however been repeatedly mentioned by Pimentel and Herr.
December I, 1971 C. Michaux, JPL Sec. 3.4, page 5
Chemical and Physical Properties
JPL 606-1
0.5
0.4
O
o
0.3
0.2
0.1
ARABIA
(MAY 1969)
•*_ • • • •• .
r.'Xf>0°r
OOo°°oO °0o
SYRTIS MAJOR
(MAY 1969)
°oOo°
• BINDER AND CRUIKSHANK 1963
ESIPOV AND MOROZ 1963
gj McCORD AND
4 MOROZ 1964
SINTON 1967
+ TULL 1966
W^
_L
X
0.5
1.0
1.5
2.0 2.5
WAVELENGTH {^)
3.5
4.0
Fig. 1. Spectral geometric albedos of a typical bright area (Arabia) and
dark area (Syrtis Major) for the 0.3 to 2.5 [i spectral region according
to the data of McCord and Westphal (1971) and other investigators.
McCord at al. (1971)
30
o
-o
a;
TO 20
a>
E
o
— INTEGRAL DISK (MOSTLY BRIGHT REGIONS)
DARK REGIONS
. SPRING
o MIDSUMMER
« LATE SUMMER
010
o 9) o ooo o o
_L
_L
6 .8
Wavelength (n)
1.0
1.2
Fig. 2. Seasonal changes in dark areas (McCord and Adams, 1969).
Sec. 3.4, page 6 C. Michaux, JPL December 1, 1971
JPL 606-1 Chemical and Physical Properties
The characteristic absorption features and slopes seen in the 0.3 to 1.1 \j. curves
are amenable to interpretation.
The extension of these curves beyond 1 . 1 |j. and up to almost 4 fji has been
attempted by various investigators as seen in Fig. 1, but the results are not in
good agreement. Recently, however, the portion up to Z.5 \i seems to have been
established with fair reliability and accuracy by the work of McCord and
Westphal {19V1) mentioned above. The authors caution that the data points in
this extended region show more scatter, implying that the newly revealed
absorption features need confirmation before proper interpretation can be car-
ried out. The full interpretation of the 0.3 to 2.5 (jl reflection spectrum or
spectra already promises improved deciphering of the composition and miner-
alology of the surface materials.
From 2.5 to about 4 \i, the continuation has been very much more
uncertain, being based only on the older measurements from Moroz (1964)
and Sinton's results (1967), which were marred by calibration difficulties;
very recently, however, Beer et al. (1971) published their very high resolution
2.5-3.2 \i spectrumi. This region also may reveal some most interesting
absorption features. However, there is insufficient laboratory data on rocks
and minerals in the mid -infrared region to afford safe interpretation of the
observed Martian spectral features. Beyond 4 [x, thermal emission of the
planet starts to take over, predominating increasingly over reflection.
Description and Interpretation of the Martian Spectra (see Figs. 1 and 3)
The reflection curves of bright and dark areas of Mars show an overall
similarity, with the same steep slope rising from 0.4 to 0.7 \j., indicating
generally a very reddish material indeed; Syrtis Major and other dark areas
are not green or grey, as somLetimes previously reported, but reddish spec-
trally.* Arabia and other bright areas are higher in albedo and redder than
dark areas.** Although the reflection curves of light and dark areas are simi-
lar, they differ in their near infrared absorption features, as well as described
below. The differences are significant enough to indicate a probable "different
mineralogical composition, and not simply a particle size difference, " to quote
McCord and "Westphal, 1971, whose curves serve as a basis for the interpreta-
tions given below.
*"Color" here means relative spectral color and refers specifically to the
spectral flux distribution, and not to the customary physiological sensation
of color (to human eyes).
**The stability of their relative color, which earlier conclusions asserted
(McCord and Adams, 1969), now seems questionable: relative to Syrtis
Major, which remained constant, Arabia was found to be significantly redder
and brighter in March than in May 1969 (McCord and Westphal, 1971). This
unexpected observation is in conflict with that of the traditional wave of
darkening and remains to be explained. However, more data at different
seasons and planetary phase angles are needed to define these effects.
December 1, 1971 C. Michaux, JPL Sec. 3.4, page 7
Chemical and Physical Properties
JPL 606-1
WALIASTONITE
CaSiO,
ENSTATITE
MgSiOj
DIOPSIDE
(Co, Mg) SiOj
HEDENBERGITE
(Co, Fe++) SiO,
FERROSIUTE
F.-H-SiO,
There are two main series of pyroxenes:
1) Calcium-rich pyroxenes, or clinopyroxene series (because
they are all monoclinic), with end members: diopside and
hedenbergite. The intermediate members are grouped
under the name augite.
Z) Calcium-poor pyroxenes, or orthopyroxene series (mainly
orthorhombic ), with end members: enstatite and ferrosilite.
Intermediate members are bronzite and hypersthene. The
orthopyroxenes are less abundant terrestrially than the
c lino pyroxenes.
Fig. 3. Ternary diagram (Ca, Mg, Fe++) Si03 showing the
compositional variations of pyroxenes-
Sec. 3. 4, page
C. Michaux, JPL
December 1, 1971
JPL 606-1 Chemical and Physical Properties
The 0. 3 to 1.1 fjL portion of the curves exhibit the following features:
1) A strong absorption in the blue-UV (0. 3-0. 4 fi) which is stronger for
the bright areas. This usually characterizes ferric oxides --if only
common geological materials are considered. The bright areas
would then be enriched in ferric Fe +++ oxides.
2) A possible band at ~0. 85-0. 87 M, only in bright areas. Its intensity
is very weak: 2 to 3 percent at most according to Sinton ( 1967) while
Younkin (1966, 1969) claims it is not present at all. There appears
to be a suggestion of it in the curve for Moab. Such a band is due to
the ferric Fe+++ ion and appears in the spectrum of limonite (Sagan
et al. , 1965; Draper et al. , 1964); however this band is not unique to
hmonite, although it is commonly referred to as the "limonite band. "
If this band truly were present in the bright areas spectra, the state-
ment of Fe+++ enrichment in such areas would be reinforced.
3) A weak absorption feature at 0. 80-0. 85 (i in the Syrtis Major curve,
even v/eaker in the Arabia curve. This feature is difficult to distinl
guish from feature 2), because data here obtained with one filter is
not well resolved. No specific interpretation is given, although it
appears to belong to the darker colored ferromagnesian silicate
minerals, such as augite (clinopyroxenes). See liand location in
Adams (1968). Interestingly, McCord and Wcstphal mention that
this absorption feature becomes stronger as the dark area becomes
darker.
4) A narrow absorption feature between 0. 9 and 1.0^, centered at
0.95 t^for the Arabia, Moab, Neith Regio bright area curves only.
This feature appears compatible with the presence of the lighter -
colored ferromagnesian silicate minerals, such as hypersthene and/
or bronzite (orthopyroxenes ). * In fact, McCord and Westphal note
that this feature becomes less strong as the albedo decreases.
This narrow feature was unreported previously because spectral
resolution was insufficient and because older data usually pertained
to an average of many bright areas.
5) A broad absorption depression from 0. 90 to I . ] - 1 . 2 h- present only
m the dark area curves, namely for Syrtis Major, Mare Acidalium,
lapygia, and Meridiani Sinus. This important feature is character-
istic of ferrous Fe++ ion in sixfold coordination. (The exact position
of the band depends on the crystal structure. ) It is probably originat-
ing here m a ferromagnesian silicate mineral (since typically Fe++
resides in them), perhaps olivine and clinopyroxene or a mixture (as
these dark minerals often occur together). McCord and Westphal
mention that the broad feature becomes stronger as the dark area
becorriL'S darker.
*Confirmahon of such an identification would be provided by a second band at
-1.8 ^^(see Adams, 1968). The spectrum of McCord and Westphal (1971)
unfortunately has a gap between 1.8-1.95 ja (telluric water vapor variations
being responsible).
December I, 1971 C. Michaux, JPL Sec. 3.4, page 9
Chemical and Physical Properties ^^^
Earlier McCord and Adams (1969) actually reported that this broad
feature is stronger in dark areas, probably not distinguishing it from
the then unknown narrow feature at 0.95 h- in bright areas, which they
apparently saw as a weak absorption at about I. p..
From the above analysis, one may conclude that although bright and dark
areas are basically constituted by similar rock-forming minerals namely
ferromagnesian silicates, they differ in that bright areas are richer m
orthopyroxenes, hypersthene for example, and ferric oxides, while dark ^ ^
areas are richer in clinopyroxenes and olivme, augite for example. These
Tnterpretations are, of course, only possibilities . Confirmation or rejection
should be provided when the presence or absence of other characteristic
features in the spectrum beyond 1. 1 fa is more firmly established.
In the 1 1 to 2. 5 u portion of the spectra, not yet firmly established,
^^ertainly several absorption features exist. We shall mention only ( 1) a general
depress on between about 1. 3 and 1. 7 ^or Syrtis Major and (2) a general
depression between about 1. 3 and 2. 2 (. for Arabia. Both curves show a posi-
tive feature at about 1. 8 h- greater in Syrtis Major. The sharp depression at
Z u in both curves is caused by the atmospheric COz absorption. No interpre-
tations have yet been offered for tiie features of this spectral region.
W a ter of Hydration
Undoubtedly from the above, water of hydration is present in the Martian
soil or rock minerals. Additional evidence has been reported such as the follow-
ing: Sinton (1967), with his birefringent interferometer, obtained 2-4 |i spectra
of Mars (see Fig. 1) which showed a very strong broad band at about 3 |j., the
shape of which followed that of a prominent H2O band. For bright areas, the
minimum was at 3.0 ix, and for dark areas 3. 1 |j.. Later, this general observa-
tion was confirmed by the high-resolution 2.4 - 4 ^x spectra of Mars (integrated
light) taken with a Connes' type interferometer (Beer et al. , 1971). The best
interpretation given is that this broad band is probably due to water of hydration
or crystallization of hydrated minerals at the surface of the planet. The amount
of this chemically bound water could not be estimated.
The shift from 3. ^ to 3. 1 to. noted by Sinton (1967) "may indicate composi-
tional differences for the hydrated minerals" of the bright versus dark areas.
Laboratory Sim ulation Experiments
Laboratory spectral reflectivity cu.-ve. from 0.3 to 1.1 \x for the Martian
bright and dark areas were modeled bv fi'ls^ns and McCord (1969), using only
geochemically expectable materials. They acuiev-.d a close fit only with an
oxidized basalt.* Starting with the same fundamental basalt material, they
l-Specifically, a fresh, dense olivine basalt from Little Lake, Calilornia, svas
attacked by dilute nitric acid, which dissolved and oxidized only the magnetite
grains; upon drying, the solution precipitated a rust-colored stain on the
unattacked mineral grains of the basalt. This stain was identified as
limonite.
„ , 1 r, r i^/fif-hanv TPL December 1, 1971
Sec. 3.4, page 10 <^« Jviicnaux, ji-j^
JPL 606-1 Chemical and Physical Properties
achieved a reasonable fit for both area types (see Fig. 4). The only difference
was that bright areas required the more oxidized and finer grained material
{<50 [J. mean particle size). It was inferred that both area types were consti-
tuted of fundamentally ferromagnesian silicate rock material.
le
To simulate the seasonal darkening of Syrtis Major, variations of the
oxidation state and/or mean particle size were not adequate. Only two darken-
ing mechanisms were .found to be satisfactory: (1) partial covering of the sur-
face by a very dark grey or black material, possibly growth of grey vegetation
or black microorganisms, the decay of which provides return to the original
state; (2) addition of moisture to the surface, for example by condensation of
atmospheric water vapor below 0°C. Drying restored the original aspect;
therefore this mechanism was reversible.
Such modeling represents only a first-order attempt. Clearly much more
work is needed. For example, the 0. 95 |jl feature in bright areas was not
matched. The limonite, or ferric oxides, coating or staining the silicate matrix
are strong pigments, only a little of which is necessary to impart a reddish
(orange or ochre) coloration to the Martian surface silicate particles. This
linionite stain model, proposed several years ago by Van Tassel and Salisbury
(1964), and Binder and Cruikshank (1965), is more plausible than the unrealistic
but fashionable model of the fifties and early sixties. These unrealistic models
proposed that powdered limonite or ferric oxides were the major surface con-
stituent rather than silicates, as on Earth (Dollfus, 1957; Sharonov, 1961). Such
outdated models were based on very meagre, strictly physical information
derived from polarimetry and wide-band photometry/colorimetry. * These
methods can give only particle size, albedo, and color, and cannot solve the
compositional problem, as is possible via narrow-band spectrophotometry.
Exotic Interpretations
An ingenious compositional model of decidedly speculative caliber was
proposed by Plummer and Carson (1969). It is based on the hypothetical carbon
suboxide C3O2, supposedly formed photochemically from atmospheric CO7 and
CO according to the reaction; COz + 2CO ^ C3O2 + OZ- This compound
readily polymerizes into heavier, rather hygroscopic molecules (C307)n, which
settle onto the surface. These molecules exhibit a range of color from pale
yellow through orange, reddish brown, and violet, to nearly black, depending
upon temperature and UV radiation. Plummer and Carson found the reflection
spectrum of the dull yellow polymer to quite nicely match the reflection spec-
trum of Mars between 0. 2 and 1. ^. Therefore, they proposed the polymer
as responsible for the colors of Mars, instead of the traditional limonite or
ferric oxides. Major objections are that there is no certainty that such a
polymer can be a stable geochemical material, under prolonged UV irradiation
-Cutts (1971) concluded from his broad-band colorimetric measurements (his
.u °. °r:'' ''^,*'''"y" P^°^'^ °^ ^^^^' g^^^"' ^"d red pictures taken by Mariner 7
that it may be possible to match the reflectivity characteristics of both light
and dark areas with oxidized basalt, but if this is the case there must be a
change in particle size as well as composition. "
December 1, 1971 C. Michaux, JPL Sec. 3.4, page II
Chemical and Physical Properties
JPL 606-1
T3
0)
^
.-H
nJ
U
• rH
U
4->
<u
s
o ^
QJ O
bJOvXJ
^ <>
rt -'
M
tiTJ'
0) '-I
a o
tn U
., U
Si ^
-i-»
. ^3
f-i C
'^
4d "i
^
tn tn
X
^ d
h-
ni rt
(')
in T3
z
i<
OJ
_J
'd tn
LU
<U (Ti
>
N (D
<
=5
^ -^
o t^
-^ 4^
43 XI
•H
w ni
-p
C -'-'
«^^
^■^
• M f-i
fH ^
11 -^
a S
X .;^
dj +j
c l;i
•4J '^
(%) oaaaiv Diyiawoao
Sec. 3. 4, page 12
C. Michaux, JPL
December 1, 1971
JPL 606-1 Chemical and Physical Properties
on Mars surface. In addition, the traditional pattern of dark and bright areas
would necessarily have been created and maintained by an atmospheric supply
of polymer. And more fatally, beyond 1.0 (jl the polymer spectrum diverges
from the Martian one. ''•-
Despite the preceding objections, the polymer hypothesis is still appeal-
ing to some chemical experimenters. Perls (1971), for example, critically
examined and enlarged the C3O2 polymer hypothesis in an attempt to explain a
number of Martian phenomena, both on the surface (seasonal changes) and in
the atmosphere (large yellow clouds, blue clearings, seasonal variation of
water vapor, etc. ). Assuming the C3O2 compound is geochemically stable over
long periods of time. Perls recognizes still the difficulties encountered by his
'working' hypothesis, and has enumerated many atmospheric constraints.
Furthermore, he proposes additional laboratory experimentation as a means of
verifying this hypothesis.
Neither the gaseous monomer C3O2 nor the polymer (C302)n have been
detected in any spectra of Mars, whether taken by Earth-based telescopes at the
highest resolutions (IR interferometry of Beer et al. , 1971) or obtained by
Mariner I969 (IR spectrometry of Horn et al. , 1972). These investigators
placed a very low upper limit on the abundance of C3O2 gas in the Martian
atmosphere (200 X 10-4 and 32 X lO'^, respectively). Their results should
seriously restrict the possibility of widespread distribution of the polymer on
the Martian surface.
Distribution of Martian Surface Mater ials; Preliminary Results
Binder and Jones (1972) performed a Mars mapping program of wideband
spectrophotometry, using 10 channels over the 0.6-2. 3 |jl region, to obtain fur-
ther data on the distribution of surface materials. (See Section 3. 2 on Photom-
etry. ) The resulting spectral albedo curves obtained for 150 well-distributefl
regions (at spatial resolution of 300-500 km) fell into two distinct and uniform
color/albedo groups, corresponding to the bright areas and dark areas as sur-
face units. They interpreted this distinction as indicating that "the surface of
Mars consists of basically only two types of materials, " characterizing two
lithologic provinces, bright and dark areas. Using laboratory comparison data
(0. 6-2. 3 n) they concluded that the general color of the two unit groups is due to
a lirnonite stain coating the basic rock materials (basalts, andesites) or soil
particles. This conclusion was consistent with prior findings by other investi-
gators. Binder and Jones further indicated that the small color difference (in
the 1 (J. region) indicates that the dark areas are "richer in olivine and/or
pyroxene" than the bright areas. However, in their testing they found that "soil
particle grain size plays virtually no role in the albedo for limonite -stained
materials. " The difference between the two basic lithologies would be small,
such as basalt versus andesite or peridotite.
*Cutts (1971) noted from his Mariner 7 tricolor -reflectivity plots that "the light
and dark areas cannot be explained as simply different polymers of carbon
suboxide"; i.e., various increases in its degree of polymerization.
December 1, IQ?! C. Michaux, JPL Sec. 3.4, page 13
Chemical and Physical Properties JPL 606-1
Stability o f Goethite on Mars
From the reflectance spectra, ferric oxides appear to be normal
constituents of the bright areas top material, with amounts in the order of 1%
or less. These oxides may be present in the mineralogical form known as
goethite FeO(OH), or rather its impure variety, limonite. However, goethite
or limonite should be subject to dehydration under the very low partial pressures
of water vapor, and the question arises as to the stability of goethite on Mars.
Therefore, it was necessary to investigate the system:
2 FeO(OH) ^^ Fe^O^ + H^O vapor
goethite hematite
and perform exact thermodynamic calculations to determine the equilibrium
water vapor pressure-versus-temperature curve in presence of these two
minerals.
Preliminary estimates were made by Adamcik (1963), who concluded that
goethite was sufficiently stable and that the system could well act as a buffer to
the Martian H2O vapor in the atmosphere. However, his estimates were
severely criticized by Fish (1966) and O'Connor (1968) on the basis of new cal-
culations proving goethite too unstable.
Recently the problem was treated anew more exactly and more comprehen-
sively by Pollack et al. (1970a and 1970b). Not only did they produce improved
thermodynamic calculations of the p versus T curve, but they also investigated
the kinetics of the hydration/dehydration process and applied their results to
conditions prevailing in the subsurface soil layers of Mars. Their vapor pres-
sure equilibrium curve is shown in Fig. 5. The dehydration rate of goethite was
found to be very slow (the time t^ for dehydration to 1/e of the original weight
for powdered goethite with an average -50 [i particle size, typical of bright
areas, was 67 hours at 2Z5 °K). Application to Mars confirmed that goethite
would be unstable at the surface itself under the heat of daytime, but would be
stable deeper within the topsoil layers where the heat wave is damped out and
the temperature remains constant and low (about 200 °K). (See Section 3. 1 on
Thermal Properties. ) If the top layers are stirred relatively frequently by
winds and cyclones, i. e. "dust storms, " dust devils, or the like, then the
surface and subsurface can be considered coupled, with the result that goethite
can appear on the surface, where it only slowly and partially dehydrates before
the next storm buries it again, permitting it to rehydrate. To the authors, this
plausible Martian hydration/dehydration cycle for goethite (and/or possibly
other hydrated minerals) appears to be balanced because the dehydration time
is much greater than the 'characteristic' time for vertical mixing by winds or
cyclonic storms. And, provided goethite or limonite is sufficiently abundant on
Mars surface, it may even assume the role of long-term regulator or buffer
(over millions of years) of the atmospheric water vapor content -because this
goethite "would contain much more water than is present in the atmosphere and
possibly in the polar caps. " (Pollack et al. 1970b. )
Sec. 3.4, page 14 C. Michaux, JPL December 1, 1971
JPL 606-1
Chemical and Physical Properties
At a giveii temperature T when the actual partial pressure of water
vapor PH2O lies below the equilibrium curve, hematite is stable and
goethite decomposes to hematite and water vapor; and conversely when
PH2O is above the curve.
Fig. 5. Equilibrium vapor pressure curve for the goethite-hematite system.
Pollack et al. (1970b)
3.4. 3 Adsorption of Volatiles: CO2 and HO
The highly particulate nature of the Martian epilith, at least in the bright
areas, suggests that adsorption of atmospheric gases may occur to a very sig-
nificant extent. This phenomenon could be a major mechanism for storing key
volatiles such as CO2 and H2O especially at nighttime or when the temperature
is low. Conversely, desorption would occur during daytime nr when tempera-
ture rises, releasing volatiles from the topmost few cm of epilith. Such an
adsorption-desorption cycle may not only have meteorologicdi consequences but
should be of vital importance to the metabolism of possible mic roorL;anism.s
residing in the soil. The few meters of epilith below could thus store more or
less permanently a significant amount of water, in order for the epilith to be in
physical equilibrium with the atmosphere.
Decemiber 1, 1971
C. Michaux, JPL
Sec. 3. 4, page 15
Chemical and Physical Properties JPL 606-1
The extent of the adsorption will depend generally upon the temperature,
the relative pressure of the gas, and the amount, state of pulverization (or
porosity) and adsorptive properties of the material. Thus, the larger the
internal surface area (porosity), the lower the temperature, and the greater the
relative pressure of the gas at that temperature, the more gas is adsorbed.
CO2 Adsorption in a B right Area: Experiments and Calculations
Davis (1969) made a conservative estimate of the amount of CO2 adsorbed
in a typical bright area under reasonable assumptions of temperature, partial
pressure of COz and adsorptive properties of the Martian dust assuming a par-
ticle size of diameter 2.8 \i (value given by Koval and Morozhenko, 1962).
Using Brunauer et al. (1938) measurements of CO2 adsorption on silica gel, he
calculated that the dust adsorbs and desorbs (upon solar warming) 4.4 X IQ-o
mole of CO2 per cc of bulk volume (particles + voids) of particulate material.
This is a significant amount for a tenuous atmosphere.
Fanale and Cannon (1971), experimenting with vacuum- pulverized basalt
of mean particle diameter 22 |jl to simulate the Martian epilith material, found
a very high adsorptive capacity: 6.75 X 10-5 mole/cc of material under 6.5 mb
CO2 and at 196°K (-77°C). They produced adsorption isotherms for this
material under a CO2 pressure range of 1-25 mb and at three different temper-
atures (-77% 0°, and 29°C). Applying this data to Mars, and assuming an
epilith depth of 10 meters and bulk density 1.5 g cm-3, they found that about
1500 cc STP (3 g) of CO2 can be adsorbed per cm^ of a hypothetical 15 meter
epilith column at 196 °K, or about 1000 cc STP (2 g) of CO2 at 210 °K.
H2O Adsorption: Experiments
Fanale and Cannon (1971), experimenting with the same pulverized basalt,
predict an even greater adsorption of water vapor than that of CO2 on Mars. An
adsorption of 8.5 mg of H2O vapor per gram, at 29°C, was noted when the
partial pressure of H2O vapor was almost at the saturation value (99%). They
calculated that many monolayers of H2O were then adsorbed. The number of
H2O monolayers absorbed, according to them, is determined primarily by
the relative pressure of H2O (ratio of ambient to saturation pressure; 1. e. ,
the relative humidity) and only very moderately by temperature. * Extrapolating
to a Martian atmosphere containing 25 precipitable micron H2O and with a base
at ~200°K (where, they remark, would prevail a near-saturation humidity), a
similar amount of H2O adsorption could be expected at the proper season.
The relative humidity being very sensitive to temperature at constant
partial pressure H2O, only a slight warming (10° or 20») of the subsurface
(above the saturation temperature) would cause rapid desorption of most of the
H2O. The released desorbed H2O vapor could then condense in the colder lower
atmospheric layer as ice crystals, fog. The authors propose that atleast some
of the (morning) whitenings seen in the Martian tropics may be explained
by such a diurnal desorption of H2O vapor from the soil.
=:<There is some evidence that this is the case, at least for some adsorbers,
down to about -30°C. Data on H2O adsorption far below the freezing point is
exceedingly sparse, however, and that for adsorption on free silicate surfaces
under such conditions is virtually nonexistent (Fanale, 1971b).
Sec. 3.4, page 16 C. Michaux, JPL December 1, 1971
JPL 606-1 Chemical and Physical Properties
3.4.4 Martian Permafrost: Speculations
The possibility that the Martian subsurface holds appreciable quantities of
H^O permafrost (permanently frozen ground) is very attractive, especially to
biologists. The idea has been advanced by a number of investigators, for
example: Lederberg and Sagan (1962), Strughold (1965), Salisbury (1966),
and Katterfeld and Frolov (1968). The biologists are contemplating a possible
source of soil moisture indispensable for the survival of a hypothetical
Martian microbiology or microflora. The geologists are interested in explain-
ing some of the morphology (for example: chaotic terrain) photographed by the
Mariners. (Note: No typical terrestrial permafrost forms were seen nor could
be seen since the resolution was insufficient by at least an order of magnitude.
A resolution of one to ten meters would be necessary. ) In any case, if perma-
frost beds or layers truly were present extensively on the planet, our concepts
about the development of its surface and atmosphere would need serious
revision. Hence, the appropriateness of examining speculations about the
possible occurrence of Martian permafrost.
Formation of permafrost presupposes three conditions are met:
1) A porous soil or subsurface material
2) A supply of water (atmospheric or subterranean)
3) Subfreezing temperatures nrost of the time
On Mars, conditions 1) and 3) are certainly fully met. Condition 2) is
partially met. There is definitely atmospheric water vapor, in very small
variable amounts, but is the amount sufficient? Ground (juvenile) water is not
known yet to occur. From terrestrial analogy, juvenile water usually occurs
in areas where the crust is active (outgassing, volcanism, . . . ); the amount of
released H2O would be adequate but the occurence most likely localized.
The presence of permafrost on Earth produces a number of very sinall
characteristic topographic features at the surface (in arctic regions), which are
due to two main processes (Wade and DeWys, 1968):
1) Frost-heaving
2) Ice /sand wedging
The typical forms produced include patterns of polygonal ridges and troughs,
conical and irregular mounds, etc.
Wade and DeWys (1968) hypothesized that such forms could well develop
on Mars, at almost any latitude, provided there is an adequate supply of water.
They rely primarily on the availability of subterranean water (juvenile water
from outgassing of the interior), rather than on atmospheric (meteoric) water.
This juvenile water would come up from the warm depths below and freeze upon
reaching the base of the permafrost layer, thus adding to its thickness gradu-
ally. The patterned ground develops from thermal contraction and expansion of
the layer, with subsequent filling of the resulting cracks with ice or sand, etc.
December 1, 1971 C. Michaux, JPL Sec. 3.4, page 17
Chemical and Physical Properties JPL 6 06-1
]'e rniaf rost Fi-orn Atmospheric Water: Calculations
Leighton and Murray (1966) calculated the behavior of Martian water vapor
assuming only an average abundance of ~10~^ g cm"'^ or 10 \i of precipitable
H^O in the atmosphere. They showed that under the known regime of insolation
(with resulting average annual temperatures below freezing at all latitudes) much
permafrost H2O (ice) could be locked in the subsurface of the polar and even
temperate regions (down to 40° -50° latitude). Only atmospheric (meteoric)
water vapor was used. The mechanism postulated was the following: if the
water vapor can penetrate into the porous soil to a depth of at least a few
meters, it may reach a region where the temperature is perpetually below the
1 90 ° K condensation temperature which is necessary to condense the H2O vapor
Jit the partial pressure of 3.7 X 10"^ mb, corresponding to the 1 X 10-3 g cm"^
abundance (Fig. 6). The water would tend to migrate poleward condensing as
permafrost. After millions of years the trapping layers would be saturated with
ice and extend downward tens of m.eters, so that 'possibly several hundred
grams of H^O per cm3 would be present in the pores of the soil. " They esti-
mated that "the top of the permafrost layer should be only a few centimeters
below the actual surface except near the (lateral) boundaries' (40° -50° latitude),
and that at a given latitude it should be shallower in the Northern than in the
Southern Hemisphere. See Fig. 7. (Note: this depth difference would define
the amount of water potentially transferable between hemispheres during the
50,000-year effective precessional cycle. They did not calculate this amount.)
Possible Frost-Heaving Caused by Atmospheric Water
Otterman and Bronner (1966) had suggested that frost-heaving phenomena
(producing microroughness of the surface because of textural changes associated
with freezing) could account for the dark areas of Mars. The niechanism sug-
gested was: atmospheric H2O (vapor) adsorbed soil H2O (water) permafrost
frozen H2O (ice).
Anderson et al. (1967) examined this possibility, using experimental data
on the adsorption characteristics of sodium montmorillonite for water vapor
and liquid (from Mooney et al. , 1952), and concluded that Otterman and
Bronner's hypothesis was "highly improbable" on the scale of dark areas.
Exception, however, was made for the unlikely (imrealistic) case where Martian
soils contain great quantities of strongly deliquescent salts, since these would
attract and retain liquid water in the soil interstices. (Note: Such salts not
only lower the freezing temperature of the soil water solution, but also can
"greatly increase the water vapor sorption capacity of any soil material" —
Anderson et al. )
Anderson et al. (1967), in their article, did not consider the possibility of
frost-heaving caused through ascent and freezing of juvenile water from the
interior.
3.4. 5 Liquid Water (at or Near Surface)
Pure liquid w^iter nornrally cannot exist for very long on the Martian sur-
face. This follows from the classical phase diagram of water (Fig. 8) and the
fact that the average surface pressure on Mars is about 5.5 inb; i.e., below the
Sec. 3.4, page 18 C. Michaux, JPL December 1, 1971
JPL 606-1
Chemical and Physical Properties
30 60
LATITUDE (deg)
90
Fig. 6. Mean annual temperature as a function of latitude, with indication of
condensation temperatures of water vapor for three atmospheric
abundances (Leighton and Murray, 1966).
30 60
LATITUDE (deg)
90
Fig. 7. Depth of top surface of H2O permafrost, as a function of latitude.
The difference in depth between the two hemispheres defines an amount
of water that should be exchanged between the hemispheres or during
the 5 X 104-year precessional period (Leighton and Murray, 1966).
December 1, 1971
C. Michaux, JPL
Sec. 3. 4, page 19
Chemical and Physical Properties
JPL 606-1
U8»K|=-125*C)
EDGE OF SPRING
POLAR CAP ON MARS
TEMPERATURE, °C
DIURNAL MAXIMUM
ON MARS 310 'K
(NEAR PERIHELIONl
Fig. 8. Phase relationships of COz and H2O.
(modified after Wade and De Wys, 1968)
triple point pressure of water (p^ = 6.105 mb). In low areas where surface
pressure is as high as 8 mb (or 10 mb), liquid water would survive for a longer
period (at a temperature just above freezing or O'C) but would still be evapor-
ating very rapidly. (Note that the rapid evaporation may create a protective
layer of ice in places, retarding the evaporation. )
Concentrated solutions of hygroscopic salts (FeCl3, CaCl2, NaCl, . . . ),
by significantly depressing the freezing point (by 10° to 50°C) as well as the
equilibrium vapor pressure (to below 1 mb), could certainly exist for much
longer periods. Such pools could form locally at hydrothermal locations, but
would of course undergo the diurnal and seasonal freezing and thawing cycles.
The availability of surficial liquid water could occur from the insolation
melting of permafrosts which are composed of eutectics of such salts, provided
the permafrost "table" is near the surface (a few centimeters below), but it is
more likely that it would originate from subterranean sources with subsurface
temperatures kept above freezing by internal (e. g. , radioactive) heating.
Surface H2O frost (condensed from the atmosphere) is not expected to
contribute liquid (melt) water under insolation, because calculations (Ingersoll,
1970, for example) show that sublimation will take place faster than melting
under the known Martian surface conditions.
Sec. 3. 4, page 20
C. Michaux, JPL
December 1, 1971
JPL 606-1 Chemical and Physical Properties
Subsurface H2O frost in soil interstices and capillaries can produce a
liquid phase (from buildup of partial H2O vapor above p^), provided the diffusion
rate is low, but then this would contribute hardly any liquid water to the surface.
See, for example, Sagan et al. , (1967) or Smoluchowski, (1968).
Direct condensation from the atmospheric H2O vapor into liquid droplets
on surface "active" spots is possible but would produce only tiny (submicron-
sized) droplets, as shown by Mukherjee (1968).
CHEMICAL PROPERTIES OF THE POLAR CAP DEPOSIT
The long-standing question of the chemical nature of the Martian polar
caps is finally reaching settlement after the 1969 flyby of Mariner 7. Of the
two acceptable theories as to the bulk of its constitution, H2O (Lowell, 1906,
reproposed by Miyamoto and Hattori, 1968) or CO2 (Wallace, 1907, revived by
Leighton and Murray, 1966 and Leovy, 1966), the latter is undoubtedly the cor-
rect one. The former, however, also has its place, simultaneously, but to a
much more restricted degree. From the evidence listed below, it can be stated
today that the Martian polar caps are composed mainly of solid carbon dioxide —
the major constituent —with solid water also, as a minor constituent. Since
1963 the following observational and theoretical evidence has been rapidly
accumulating toward this view.
1) Kaplan et al. (1964): Spectroscopic determination of 4 mb of CO2 and
14 p, of precipitable H2O vapor over Mars.
2) Kliore et al. (1965): Mariner 4 radio occultation determination of an
approximately 5 mb total surface pressure and inference of at least
5 0% CO2 atmosphere.
3) Leighton and Murray (1966): Theoretical calculations indicating that
polar temperatures fall low enough to precipitate CO2, and that the
polar caps probably are predominantly CO2 plus minor amounts of
H2O.
4) Leovy (1966) and Leovy and Mintz (1969): Calculations of the mini-
mum atmospheric temperatures required to prevent CO2 condensation
and simulation of the Martian atmospheric circulation respectively
led to essentially the same conclusion.
5) Morrison (1968): Re-analysis of the Sinton and Strong (I960) IR
radiometric temperatures, permitting extrapolation for the extended
cap's edge (60° latitude) to a CO2 condensing temperature (145° K).
6) Kieffer (1968): Laboratory IR spectra of CO2-H2O frosts indicating
that very small amounts of H2O can mask the characteristics of the
solid CO2 spectruin, thus rending illusory the earlier conclusions of
Kuiper (1952) and Moroz (1964) in favor of an H2O cap. *
='=Polar cap spectra obtained from Earth are limited to those of Kuiper (1952)
and Moroz (1964) who concluded that the caps consisted of H2O.
December 1, 1971 C. Michaux, JPL Sec. 3.4, page 21
Chemical and Physical Properties JPL 606-1
7) Schorn et al. (1969): Spectroscopic observations of greater than
average amounts of precipitable H2O vapor over the late spring/early
summer rapidly shrinking polar cap (35 micron precipitable H2O).
8) Sharp et al. (1971): Mariner 7 TV observations of the morphology of
the expanded South Polar Cap revealing interior frost cover thick-
nesses in the order of meters or more.
9) Neugebauer et al. (1971): Mariner 7 IRR determination of the very
low (148°K) and constant daytime temperature of the sublimating
expanded South Polar Cap.
10) Herr and Pimentel (1969, 1971): Mariner 7 IRS identification of spec-
tral absorption features characteristic of solid CO2, plus the possi-
bility of H2O and impurities for the South Polar Cap composition.
3.4.6 South Polar Cap: Mariner 7 IRS Results
On August 5, 1969, when the South Polar Cap was starting to shrink
during early Martian spring, the IRS channel Z of Mariner 7 recorded some
19 spectra. These spectra between 1.88 and 6.00 fi. were obtained over the cap
fromi 61° to 80°S latitude with an areal resolution of about IZO km square.
Although final analysis of the data has not been completed, it has been con-
cluded that the dominant absorption features are "obviously attributable to
solid CO2" (Herr and Pimentel, 1969). Typical polar cap spectra is shown in
Fig. 9, with comparison laboratory spectrum. Besides the common solid CO2
features, such as the 2.0 \x (4900 cm"!) intense sharp band, two distinct new
features, labelled X and Y, were noted near 3.31 |j. (3020 cm'^) and 3.03 \x
(3300 cm-1) respectively. The maximum intensity of these bands occurred at
about 68 °S, 19''W, not far from the cap's edge (61 °S). These bands (X and Y),
which happen to match some bands of methane and ammonia, appear to be
previously unreported spectral features of solid CO2 in thicknesses of several
millimeters, according to the laboratory experiments of Herr and Pimentel
(1969). These two absorptions correspond to "forbidden" transitions of the
CO2 molecule (which violate the spectroscopic selection rules). Pimentel and
Herr tentatively attributed these transitions to lattice imperfections in the
deposit's microstructure. It was noted that these features disappear further
south, nearer the pole (~80°S).
3.4.7 Near-Infrared Reflection Spectra of CO2-H2O Frosts
For the purpose of interpreting the pre-Mariner 1969 spectra (Kuiper,
1952; Moroz, 1964) of the Martian polar caps, Kieffer (1968, 1970a, b) performed
laboratory measurements on the spectral reflectance from 0.8 to 3.2 |i of the
frosts likely to occur on Mars: pure CO2, pure H2O, mixed CO2-H2O, and H2O
on CO2 frosts in a wide range of grain sizes. His main results are summarized
below. (For the spectra themselves, see the original papers. )
1) CO2 Frosts: Uniformly high reflectance (>90% at \ below 2.5 |j.) from
0.8 to 3.2 \i except for two major absorption bands centered at 2.0 and
2.7 |x (quite broad: 2.6-2.8 fx usually). These two bands, which are
the triplet and doviblet observed in the gas phase, deepen in absorp-
tion as grain size increases, while new, weaker features unique to
Sec. 3.4, page 22 C. Michaux, JPL December 1, 1971
JPL 606-1
Chemical and Physical Properties
1
lifOJt rOLAt C*f: 49'S, 3\3'[
ON POLAR CAP: dfi'S. i4\n
N[A« SOUTH POlf: >0*S. je'f
/"-v.
vv-'
NOTE: A portion of the 3 jjl region is recorded twice on either side of the
spectrometer spike.
Fig. 9. Near-infrared spectra of the South Polar Cap by Mariner 7 (left)
and comparison laboratory spectriam of solid CO2 at 7 7 "K (right)
(Herr and Pimentel, 1969).
Deccrnljcr I, 1071
C, Michaux, JPL
3cc. 3. 4, page 23
Chemical and Physical Properties JPL 606-1
the solid phase start appearing (such as a strong one at Z.62 \x, also
2.85 10., and weak ones at 1.87, 2.12, 2.28, 2.34, 2.90, and 2.03 \x).
The 2.7 fj. band is saturated (1% reflectance) for frosts of textural
scale 50 |Jl grain size. The 2.9 to 3.2 fi. continuum is very sensitive
to H^O impurity.
2) H2O Frosts: Broad absorptions centered at 1.56, 2.04, and 3.0 \x for
fine frosts especially, and saturation adsorption from 2.9 to 3.2 |i for
all frosts. The reflectance showed great variability with growth con-
ditions. Compared to CO2 frost, the H2O frost spectra have lower
near-IR reflectance, generally decreasing from 0.8 to 3.2 |j. (and
especially with thin, fine H2O frosts).
3) Mixed CO2-H2O Frosts: Small amounts of H2O have a strong effect
on CO2 spectra, an effect which increases with grain size. When
H2O is >10% by mass, it becomes difficult to identify CO2, because
the H2O features predominate. The presence of CO2 is revealed
only by the high reflectance at 2.5 fji for fine frosts, and only by minor
absorption detail at 2.0 and 2.7 [x for coarse frosts.
4) H2O on CO2 Frosts: The effect is even more drastic, since a surface
layer of only a few nag cm"^ of H2O will mask an underlying thick CO2
deposit (about 7 mg cm~^ will suffice; with 1 mig cm"^, H2O already
subdues the CO2 features).
3.4.8 Possible O3 Adsorption by the Polar Cap: Mariner 7 UVS Results
All 45 of the Mariner 7 UVS spectra^of the spring South Polar Cap showed
a broad absorption band centered at 2550 A with a shape corresponding well to
that expected from ozone O3 (Hartley continuum band between 2000 and 3000 A);
see Fig. 10 from Barth and Hord (1971). Ozone is a normal product occurring
in the photochemistry of CO2. This ozone could be either in the atmosphere,
preferentially over the cap, or adsorbed (or trapped) in the solid CO2 of the cap
itself as si;ggested by Broida. In the first case, the amount of atmospheric
ozone was calculated by Barth and Hord to be 1 X 10"^ cm-atm or 3 X 10^°
molecules cm"'^ of O3 (or a mixing ratio of ~10"' for 03/C02)« In the second
case, the amount of adsorbed ozone could be estimated on the basis of the lab-
oratory experiments of Broida et al. (1970), but no figure is available.
The polar cap C O2 miight also adsorb or trap other gases (such as O2,
CH4, NH3, . . . ), if present in Mars atmosphere, and serve as a "sink" for
these gases (Barth and Hord, 1971).
3.4, 9 Speculations on the Composition and Structure of the Polar Caps
The brilliant white tnaterial constituting the Martian polar caps is now
generally believed to be the icy condensation products of the abundant gaseous
carbon dioxide and the scarce water vapor unquestionably present in the
Martian atmosphere. The saturation vapor pressure curves of solid CO2 and
H2O (Fig, 8) show that, under the prevailing Martian mean surface pressure
of 6.5 mb (Kliore et al. , 1969), and a H2O/CO2 mass ratio of lO-"^, the con-
densation temperatures are very low: for CO2 it is 148 °K, while for H2O it is
Sec, 3,4, page 24 C. Michaux, JPL December 1, 1971
JPL 606-1
Chemical and Physical Properties
5-
4-
■- 3-
a:
2-
— \ —
2000
1 1 1 1 1 —
2500
Wavelength (A)
3000
Fig. 10. Ratio of reflectance of polar cap to reflectance of a desert region.
The minimum at about 2550A suggests absorption by ozone.
Barth and Hord (1971)
~210°K. (The condensation temperature T^ varies in the same sense as the
surface partial pressure. ) Thus, normally solid H2O should first condense out
as the temperature goes down and form snow or ice fog in the atmosphere or
frost on the cold ground. These products should be finely divided since the
atmospheric HtO abundance is very low. Then, at a temperature of 153 °K (if
pC02 is still 6.5 mb), according to the Miller and Smythe (1970) experiments
(see Fig. 11), the finely divided H2O ice combines with gaseous CO2 rather
rapidly (several hours) to form the carbon dioxide clathrate hydrate C02»6H20,
which also has the appearance of ice. (All the H2O ice becomes converted into
C02'6H20 hydrate. ) Finally, at the slightly lower temperature of 148 °K, gas-
eous CO2 condenses out into solid CO2 in the presence of the hydrate, so that
both coexist. The hydrate is stable down to about 121 °K, below which teinpera-
ture it dissociates into H2O ice and CO2 ice. Thus, if the temperature in the
Martian polar regions remains above l2l°K, the polar cap formed ''can consist
of water ice, water ice + CO2 hydrate, or CO2 hydrate + solid CO2, but not
water ice + solid CO2" (Miller and Smythe, 1970). If the surface pressure is
different (than 6.5 mb), one can find the exact combination or condensation tem-
peratures (for CO2 hydrate and solid CO2) from Fig, 11, which shows the dis-
sociation pressure curve of CO2 hydrate and also the vapor pressure curve of
solid CO2.
the
The above sequence of condensation products applies particularly to the
formation of the polar cap in wintertime. It appears that the composition of t,..
formed cap (at end of winter) may well be layered: the lower layers being rich
in CO2 hydrate and representing the "fall" cap, while the upper layers would be
rich in solid CO2 and represent the "winter" cap. The formed cap, however,
will have a bulk composition of solid CO2 (and a small fraction of CO2 hydrate)
since the atmospheric H2O vapor ratio is always small. The composition of the
shrinking cap in late spring-early summer is expected to be quite different from
December 1, 1971
C. Michaux, JPL
Sec. 3. 4, page 25
Chemical and Physical Properties
JPL 606-1
400
1
i
1 \ .
1
200
-
Hydrate +COj(s)
/Hydrate + /
/
100
—
^ CO,(g) /
""
80
_
/
—
X /
^
^ 60
V^/
/
—
XI
c /
/^
E
cr/
/
■
X /
^
£ 40
3
—
i/ /
/
"~
vt
a)
20
/ //
Ice + CO,(g)
10
8
:/
/ ^/
-
6
A
I
I 1
!
150
160 170
Temperature (°K)
180
190
Fig. 11. Phase diagram of carbon dioxide hydrate. Experimental
dissociation presstire measurements, O; vapor pressure of solid
CO^. Miller and Smythe (1970)
that of the winter cap for the same physical reasons. If one considers the more
rapid sublimation of solid CO2 than H2O ice — (resulting from the dissocation, *
under the higher temperatures, of the CO2 • 6H2O hydrate into H2O ice and CO2
gas) —this cap should with the advancing season more and more rapidly enrich
itself in its H2O ice proportion so that before its final late summer stage the
remnant cap should exhibit a bulk composition closer to H2O ice (with a small
fraction of CO2 hydrate remaining). The Mariner 1971 orbital mission might
determine this for the remnant south cap. (The Mariner 7 in 1969 viewed an
early spring south cap. ) In some places however, the summer shrinking south
cap may be covered with dust layers protecting the rapid CO2 sublimation,
since major dust storms usually occur in southern late spring and early
sumnier (see Section 4, 2 on Seasonal Activity).
*This dissociation, taking place at 153°K under 6.5 mb, is the regular hydrate
dissociation with temperature rise. The dissociation at temperatures below
121 °K is an anomaly —known only for this hydrate —and caused by the fact that
its dissociation pressure is greater than the vapor pressure of solid CO2 below
121 °K (Miller and Smythe, 1970).
Sec. 3. 4, page 26
C. Michaux, JPL
Deceml^er 1, 1971
JPL 606-1 Chemical and Physical Properties
BIBLIOGRAPHY
Adamcik, J. A. , 1963, The water vapor content of the Martian atmosphere as a
problem of chemical equilibrium: Planet. Space Sci. , v. 1 1, no. 4,
p. 355-359, April.
Adams, J. B. and Filice, A. L. , 1967, Spectral reflectance 0.4 to 2.0 microns of
silicate rock powders : J. Geophys. Res. , v. 72, no. 22, p. 5705-5715,
November 15.
Adams, J. B. , 1968, Lunar and Martian surfaces: petrologic significance of
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p. 1453-1455, March 29-
Adams, J. B. and McCord, T. B. , 1969, Mars: interpretation of spectral reflec-
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Anderson, D.M. , Gaffney, E.S, , and Low, P.P., 1967, Frost phenomena on
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Barker, E.S,, Schorn, R,A. , V\roszczyk, A, , Tull, R.G. , and Little, S.J, , 1970,
Mars: detection of atnnosphe ric water vapor during the southern hemi-
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1310, December 18.
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experiment results: topography and polar cap: U. of Colorado, Boulder,
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Beer, R. , Norton, R.H. , and Martonchik, J. V. , 1971, Astronomical infrared
spectroscopy with a connes-type interferometer: II - Mars, 2500-
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Binder, A. B, and Cruikshank, D. P, , 1966, Lithological and mine ralogical inves-
tigation of the surface of Mars: Icarus, v,5, no. 5, p. 521-525, September.
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Sec. 3.4, page 28 C. Michaux, JPL December 1, 1971
JPL 606-1 Chemical and Physical Properties
Herr, K. C. and Pimentel, G. C. , "Infrared Spectroscopy Chap. 6, p. 83-96 of
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(94 p.).
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v. 75, no. 3, p. 501-509, January 20.
Kliore, A. , Cain, D.L. , Levy, G. S. , Eshleman, V. R. , Fjeldbo, G. , and
Drake, F.D., 1965, Occulation experiment: results of the first direct
measurement of Mars' atmosphere and ionosphere: Science, v. 149,
no. 3689, p. 1243-1248, September 10.
Koval,I.K. and Morozhenko, A. V. , 1962, Several properties of the yellow haze
observed on Mars during 1956: Astron. Zh. , v. 39, no. 1, p. 65-72,
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December 2.
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JPL 606-1 Chemical and Physical Properties
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December 1, 1971 C. Michaux, JPL Sec. 3.4, page 31
Chemical and Physical Properties JPL 606-1
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Sec. 3.4, page 32 C. Michaux, JPL December 1, 1971
JPL 606-1 Morphology and Processes
3.5 MORPHOLOGY AND PROCESSES
INTRODUCTION
Prior to our era of spacecraft and radar exploration of planetary
surfaces, only speculations (based on poor-resolution telescopic observations
and using lunar and terrestrial analogies) could be evolved as to the morphology
of the surface of Mars. In 1965, Mariner 4 revealed the presence of many
craters, as on the Moon, but the Martian craters exhibited significant erosion.
Mariners 6 and 7, in 1969, while confirming the wide extent of cratered terrain,
discovered two new types of essentially uncratered terrain: the chaotic terrain
of Pyrrhae Regio, and the featureless terrain of Hellas, which was found to be
a basin. Mariners 6 and 7 also investigated the relationships between dark and
bright areas, established the nature and estimated the thickness of the South
Polar Cap, and apparently dispelled the lingering myths about the Martian
canals. Radar observations had separately established the topography of the
northern equatorial belt, and determined surface properties such as dielectric
constant and average roughness.
With the flood of new information, speculations on a different level about
the geology of Mars have replaced the older, rudimentary ideas. The present
speculation will, in turn, be influenced by new data obtained from the current
(197 1) orbital missions.
Processes at work on the Martian surface and beneath it (subsurface),
responsible for its morphology (other than atmospheric and impact modifica-
tions), still remain unknown or are very poorly understood. Some of these
processes have apparently operated continuously or episodically, over very
long times to shape the present distinct Martian surface. Other processes
appear relatively more recent in the planet's history.
The following review attempts to present only the factual information
about Martian morphology proper, as gathered primarily through examination
of the near encounter (NE) pictures of Mariners 4, 6, and 7, which were taken
several thousand kilometers away from Mars. The tableau presented here is
very incomplete because of the limited surface coverage of the planet (less
than 20%). Only some 80 total NE pictures were taken. A selection of
Mariner 6 and 7 pictures is presented in Section 3. 6, Photographic Atlas. The
far encounter (FE) pictures, although obtained at coarser surface resolution,
have provided a global survey which has permitted new maps to be drawn. As
for the probable processes at work, this review is necessarily speculative in
character, based on the present state of knowledge.
It must be mentioned that relatively few geological or morphological
studies of Mars have been published since the Mariner 6 and 7 encounters. At
present, there is little doubt that most geologists attracted by this subject are
eagerly awaiting the more extensive information expected from a successful
completion of the Mariner 9 orbital mission (and also from the Mars 2 and 3
Russian missions), before undertaking a comprehensive study of the Martian
features and processes.
January 1, 1972 C. Michaux, JPL Sec. 3.5, page 1
Morphology and Processes JPL 606-1
3.5.1 TOPOGRAPHY
The topography of substantial portions of the Martian surface has been
established with a fair degree of reliability by Earth-based radar and by the
infrared and ultraviolet spectrometers on board Mariners 6 and 7. The radar
coverage, geometrically restricted to latitudes less than 25°, was mostly in
the Northern Hemisphere, while the often-overlapping swaths of the Mariner
IRS and UVS extended prongs deep into the Southern Hemisphere. Some
locations were sensed or scanned by all three methods, which permitted a
comparison of results. It was found that the data do not always agree; in fact,
some of the discrepancies are rather notable (4 km). In such cases, preference
has been given to the radar results, which should provide correct relative
altitudes (along latitudinal scans around the planet) and are not affected by
undetermined systennatic errors caused by atmospheric (aerosols, etc.) and
photometric (contrast) effects, as are the IRS and UVS results. The radar
method is explained in Section 3. 3, Radar Properties. The IRS and UVS
methods are briefly discussed here, with indications of the best resolutions
attained. An attempt is then made to summarize the topographic information
obtained by these three methods. Indications on the 1971 radar results are also
included. Additional Earth-based IRS measurements were also made by two
teams (Belton and Hunten, 1971; and Wells, 1971 a), but the resolution
is necessarily much coarser (order of 1000 km), and many sources of error
(seeing, calibration, etc. ) have apparently affected the results, which do not
generally agree with those presented here.
Spectroscopic Methods
Infrared (Mariner IRS)
The infrared spectroscopic method measures the absorption of a
prominent CO2 band (2 \x in Mariner IRS, or 1.05 \i in Earth-based IRS) and
derives, through laboratory curves of growth, the abundance of the CO2 in the
atmospheric path sampled. From the CO2 abundance a corresponding surface
partial pressure of CO2 (at the base of the atmospheric column) is obtained,
which is finally converted to topographic height above a reference level, upon
adoption of an atmospheric model. (For the Mariner IRS data, a 100% CO2
atmosphere, isothermal at 200°K, was chosen by Herr et al. , (1970.)*
The zero altitude chosen was that corresponding to po = 6.0 mb. The
best horizontal resolution (on ground) was 130 X 8 km. The resolution corre-
sponds to the instrument's field of view (2.0° X 0.23°) and on the slant range to
the planet. The absolute accuracy claimed in heights measured was ±1 km.
Ultraviolet (Mariner UVS)
The ultraviolet spectroscopic method measured the UV reflectance
originating primarily from atmospheric scattering (molecules and possible
small aerosols), which is taken as directly proportional to the total number
>:'The Mariner IRS results of Herr and Pimentel are being revised, with
adoption of a more appropriate or realistic atmospheric model. No results
have been published yet.
Sec. 3.5, page 2 C. Michaux, JPL January 1, 1972
JPL 606-1 Morphology and Processes
of scatterers in the atmospheric column sensed. This number is in turn con-
verted to topographic height by means of an atmospheric model, as in the IRS
method. Basic assumptions are (1) homogeneous scattering atmosphere, and
(2) uniform ground reflectance.
The zero altitude chosen was that corresponding to Pq = 6,105 mb (at the
triple point pressure of H2O). The best horizontal resolution was 140 X 14 km,
but 300 X 30 km at the linnbs. The relative accuracy clainned was 1%,
Summary of Present Topographic Information
The below-listed topographic highlights were obtained via all three
remote-sensing methods. The information, categorized by data source, is
related to the classical surface features (dark areas and bright areas as shown
in Fig. 7), starting at 0° longitude and working eastward across the surface of
Mars without regard to latitude. Comparison of the detailed surface pressure
and derived altitude results from Mariner 6 and 7 UVS and IRS experiments is
shown in Figs. 1, 2, and 3, from Hord (1971).
1) From the Mariner I969 UVS and IRS, and the 1971 Radar Results
Deucalionis Regio, in its western half, bordering Pandorae Fretum,
is equally a high region (2-4 km) to both UVS and IRS; but this
disagrees markedly with the 1971 radar information, indicating it
is some 4 km lower. The 1971 radar traverses show the whole of
Deucalionis Regio to gently slope from 1 to -3 km along its 2500-km
length, parallel to the equator.
Pandorae Fretum and from northeast Noachis desert to
Hellespontus form a high plateau about 3 -km high, stretching some
2000 km (UVS and IRS).
Hellas is a large basin with its central region as low as 3 km below
mean or zero level (UVS and IRS). Its western edge in Yaonis
Regio and Yaonis Fretum is a steep slope (looking toward Hellas,
with a 4-km altitude difference over only 300 km (UVS).
lapygia is at mean elevation (0 km) as is Trinacria to the
east of it (radar 1971).
Northern Aurorae Sinus near Juventae Fons shows a greater dis-
agreement of ~5 km between UVS (-3 km) and IRS (2 km). In
Juventae Fons and northwest of it, however, the UVS and IRS
agree, with altitudes of 2 to 3 km.
Argyre (north of Mare Australe) appears to be rolling terrain at
medium elevation (UVS and IRS), and apparently slopes down
toward Ogygis Regio (UVS).
January 1, 1972 C. Michaux, JPL Sec. 3.5, page 3
Morphology and Processes
JPL 606-1
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Fig. 1. Mariner 6 UVS and IRS surface pressures and derived altitudes.
(Hord, 1971)
Sec. 3. 5, page 4
C. Michaux, JPL
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JPL 606-1 Morphology and Processes
Margaritifer Sinus at its southwestern side, and Eos, are close to
mean level (0 to 1 km) according to both UVS and IRS, but are
definitely low (-3 to -4 km) according to the 1971 radar. This is
again a major disagreement of 4 km.
Tip of Margaritifer Sinus and Aram channel are at mean level
(0 km), according to both UVS and IRS; likewise, Thymiamata (UVS
and IRS). But all this disagrees with the 1969 radar data showing a
-Z and -3 km elevation.
2) From the 1967-1969 Radar Results
Contiguous deserts, Eden, Moab, and Arabia, appear to be at
mean level. But Aeria and the western half of Syrtis Major are high
ground (3 km, even 4 km), while the eastern half of Syrtis Major is
a steep slope (3 km over 200 km), which, in fact, continues through
Isidis Regio and especially in Moeris Lacus forms a hole 3 km
deep. Amenthes is at mean level, but most of Aethiopis is some-
what lower,
Elysium appears to be an isolated mountain, at least from its
southern approaches or flank.
Complex Cerberus -Trivium Charontis is below mean level (-1 to
-2 km), and southeast of it Mesogaea is a deep hole reaching -5 km.
Amazonis likewise is low, but not as much (-3 to -4 km).
Nix Olympica, mostly low ground, sits on a fair slope climbing east-
ward from -4 to 1 km. Hougerius Eacus (east of Nix Olympica) is
above Nix Olympica on this same large slope which reaches great
heights of 6 km, where it meets Ascraeus Lacus to form one of the
highest domes or uplifts on Mars, with summits at 7 or 8 km. Then
there is an eastward downslope across Tractus Albus to Candor,
which is apparently at zero level.
Lunae Palus has a high (2 km), and east of it are found markedly
lower areas stretching as far as Thymiamata, Especially low is
the part of Xanthe near Niliacus Lacus (-5 km) and the whole of
Chryse (-3 km), Oxia Palus is at about -2 km.
3) From the 1971 Radar Results
Southern Aurorae Sinus (near Capri Cornu) has a deep gorge
plunging to -6 km. Southern Margaritifer Sinus is quite low (-3 to
-4 km). Then there is a steady rise eastward over Deucalionis
Regio from -3 to 1 km in eastern Sabaeus Sinus. From Sabaeus
Sinus to Mare Tyrrhenum, there is a long stretch of mean-level
rolling plain (0 to -1 km). lapygia and Trinacria are near zero
level; then Aeolis and Zephyria have lows reaching -3 km.
Zephyria is quite low (-1 to -4 km), with depressions near
Laestrygonum Sinus, apparently. Memnonia is generally low, but
slopes up from -3 to 3 km going east. Then eastward this slope
continues to rise steadily up to the highest summits (the "twin
January 1, 1972 c. Michaux. JPL Sec. 3.5, page 7
Morphology and Processes JPL 606-1
peaks"), which may be part of the high uplifts previously
mentioned. Phoenicis Lacus and the similarly shaped marking west
of it (unnamed) are apparently at 6 and 7 km. Then eastward is a
very gradual downslope to Sinai (Thaumasia) and Coprates, which
are still high (2 km) however.
Radio occultation results of the three Mariner flybys have yielded values
of "surface' pressure at six different locations on Mars (see Section 5, 2,
Table 1). The areas of northern Mare Acidalium, Electris, and Boreosyrtis,
Hougeria are low regions. The tip of Meridiani Sinus is at average altitude,
while southern Hellespontus is an elevated region. Cross-checking correlation
with the above summary is possible only for the latter two or possibly three
locations. Good agreement is noted.
Interpretation
For a small planet, Mars has a pronounced topography with altitude
differences of at least 12 km. There are high elevations, plateaus, valleys,
basins, ridges, gorges, etc.. besides the many craters and ring-like struc-
tures. This implies that the structural (tectonic) development of Mars was
rather complex and extensive, producing orogenetic uplifts and perhaps
intrusive and extrusive magmatic activity (Binder, 1971). This supports the
new concept of a differentiated Mars (see Section 2, Interior). It is interesting
to note that some investigators (Leonardi, 1966; Katterfeld and Hedervari,
1969) of Mariner 4 pictures pointed out some striking resemblances in groups
of craters and cirques on Mars, to eruptive and collapsed structures on Earth
and presum^ably on the Moon.
3.5.2 NEW MARS MAPS
Several new maps and charts of Mars have been prepared since 1969.
These maps and information concerning their source data and method of
preparation are presented in the following sequence:
1) Mariner Mars 1969 Chart (NASA), including both polar regions.
2) International Planetary Patrol Photographic Maps of Mars 1969
and 1971 (Lowell Observatory).
3) Mariner Mars 1971 Planning Chart and Planning Charts of South
Polar Regions (G. de Vaucouleurs).
4) Mariner Mars 1969 Photomap (J. Cutts, C.I. T. ).
5) Mariner Mars 1969 Regional Maps (C. Cross).
a) Meridiani Sinus.
b) South Polar Region.
Sec. 3.5, page 8 C. Michaux, JPL January 1, 1972
JPL 606-1 Morphology and Processes
Mariner Mars 1969 Chart (NASA)
One of the most significant applications of the Mariner 1969 TV
photography was the production of a much more reliable map of Mars, based
for the first time on fixed topographic features (i.e., craters), and constructed
according to modern photogrammetric principles. Accurate positioning of the
Martian features was possible through a new areodetic control net established
by Davies and Berg (1971) in cooperation with the Aeronautical Chart and
Information Center (A.C.I. C). This net consisted of 112 control points, or
"clearly identifiable marks, " on the surface of Mars. These control points
were selected principally from crater centers, and when these were not avail-
able, the centers of particularly dark or light spots, or tips of markings, were
used. Measurements of the positions of these control points on geometrically
corrected FE and NE TV pictures permitted derivation of their Martian
coordinates by analytical triangulation, knowing the position of the spacecraft
and assuming a certain equatorial radius 3393-4 km) and polar flattening
(21 km). The orientation of the coordinate system was determined by the direc-
tion of the rotational axis.
This new NASA Mars Chart 1969, Fig. 4, issued by the A. C.I. C. in
August 1970, is a Mercator projection at the scale of 1:25,000,000 at the
equator. In addition, there are two stereographic projections of the polar
regions (60° to 90" latitude). The appearance of the main chart is quite
striking when compared to earlier Mars charts, such as the MEC-2 produced
in 1967 by the A. C.I. C. Besides the usual albedo mapping of dark and bright
areas, it presents two large inserts corresponding to the equatorial and polar
swaths covered, where topographic relief is qualitatively incorporated by air-
brush renditions. Such interpretative renditions are of course subjective,
depending upon the experience and judgment of the mapping staff guided only by
visual inspection of the photographs. For example, the drawings in niany
places have tended to accentuate the relief.
The chart also includes the Mariner 4 (1965) swath from the equator
to the south polar region, mostly across Mare Sirenum and Phaetontis.
The albedo mapping is that obtained from the July-August 1969
encounter period. The traditional permanent dark and light areas are easily
recognizable and some secular changes are noticeable to those familiar with
Mars, but no seasonal changes were yet manifested at the time.
The basic nomenclature, or list of areographical names, is that
adopted by the 1958 lAU Congress. However, additional well-known names are
used where needed.
The longitude system is the standard westward -counting to 360 degree
system long used by areographers and telescopic observers of Mars.
January 1, 1972 C. Michaux, JPL Sec. 3.5, page 9
Morphology and Processes
JPL 606-1
u
en
^1
<
Sec. 3. 5, page 10
C. Michaux, JPL
January 1, 1972
JPL 606-1 Morphology and Processes
International Planetary Patrol Photographic Maps of Mars 1969 and 1971
(Lowell Observatory)
Mars 1969
This new albedo map of the Mars 1969 apparition was prepared at
Lowell Observatory's Planetary Research Center, from thousands of photo-
graphs taken with identical photographic equipment (camera systems and focal
lengths, filters). The photos were obtained under the control of the Inter-
national Planetary Patrol Program, by a network of six observatories located
around the Earth. The participating observatories were Lowell Observatory,
Flagstaff, Arizona (Coordinating Center); Maunea Kea Observatory in Hawaii;
Cerro Tololo Inter-American Observatory, Chile; Mount Stromlo Observatory,
Australia; Republic Observatory, Johannesburg, South Africa; and the
Magdelena Peak Station of New Mexico State University, New Mexico, The
photographic period of a month and a half transpired between May 19 and
July 19, 1969 (opposition date. May 31). Although filters in all colors were
used, only the red filter* photographs were selected for the mapping to obtain
maximum contrast of Mars surface features. Areographic positions were
obtained on the images by superimposing the appropriate orthographic graticule
of latitude -longitude lines, through a specially built projection image reader.
Mean positional error was estimated to be less than a degree for latitudes
between 40°N and SCS, The new Mercator map is shown in Fig. 5 and contains
the names of 191 Martian features properly inscribed so as to form a key map.
The list of names adopted comprises 113 names from the Standard lAU 1958
List, plus a selection of traditional names most widely used today. Both the
key map and the list of names, which included the approximate 1969 location
coordinates, are shown in Table 1. The objectivity of the cartographic pro-
cedures employed by the authors (Inge et al. , 1971) of this Earth-based
telescopic map undoubtedly provided a high level of reliability. In their own
words, "Our cartography was carried out completely "de novo, " without the use
of earlier maps or earlier notions concerning the nature of various markings.
Care was taken not to over -interpret the photographic evidence. "
Mars 1971
A similar photographic map for the 1971 apparition (Fig. 6) was pre-
pared, using the same techniques. The same observatories participated (in
supplying the red filter photos), with the exception of Magdalena Peak Observa-
tory, and the addition of Perth Observatory, Western Australia, and Kavalur
Observatory, Indian Institute of Astrophysics, India. Resolution was a little
better since Mars was closer in 1971, especially for the Southern Hemiisphere
turned toward Earth.
Mariner Mars 1971 Planning Char ts
An albedo map of Mars as it is expected to appear about January 1,
1972 -- the time of the Mariner 1971 orbital mission -- was prepared by
G. de Vaucouleurs (1971) with the aid of J. Roth, artist. At this time, Mars
*Red filter passband 300 A wide, centered on X6200 A.
January 1, 1972 C. Michaux, JPL Sec. 3.5, page 11
Morphology and Processes
JPL 606-1
AMAZONIS \ otZc *'=•*"" iu~«
MESOSAEA
THARSIS
MEMNONIA f
PANCHAIA
WOKWTfi
/ I
\j' I CaRENIA «
■i PKOPOHTM
JT,.*.. I * %% *■ T ^ ELYSIUM • ' ,
AUSOHIA " ^i^ ^Vi^
•OMALis *%^ *«<UM M.
fH»*l
T AliSONtA
" AUSrtAUS E«IDANIA
MARE
OIA
A U S T R A
oeveis
^ HH-LHTONTICA
OCtANIOUM M. °-
aECTRIS
CHRONIUM M.
L B
_L.
-i_
Fig. 5. 1969 Mars Patrol Photographic Map (Lowell Observatory)
(Inge et al. , 1971).
Sec. 3.5, page IZ
C. Michaux, JPL
January 1, 197 2
JPL 606-1
Morphology and Processes
Table 1. List of names used on the International Planetary Patrol
Photographic Map of Mars 1969 (Lowell Observatory)*
1. Achillis Fvns, 53°, +23°
2. AcliilUs Fom, 30°, +37°
3. Acklalium Marc, 28°, +48°
4. Acidaliui Fans, 63°, +58°
5. Acolis, 212°, -10°
C. Aciia, 310°, +-15°
7. Acthcria, 240°, +40°
a. Acthiopis, 235°, -|-10°
9. Agathodaemon, 65°, —14°
10. Albor, 208°, +18°
11. Alcyonius Nodiu, 268°, +35°
12. Atcyonius, 260°, +50°
13. Amazonis, 160°, +20°
14. ^mfiroj/o^ 85°, —38°
15. Amcnthcs, 251°, +3°
16. Amphitrites Mare, 322°, -58°
17. Anian, 228°, +48°
IP Anligones Fans, 295°, +20'
1". Aonius Sinus, 105°, —47°
20. Arabia, 320°, +28°
21. Aram. 12°, -5°
22. Araxcs, 117°, -24°
23. Arcadia, 115°, +42°
24. Arethusa Lacus, 337°, +58°
25. Argus, 10°, 0°
26. Argyrc I, 35°, -48°
27. Arnon, 337°, +.50°
28. Ascraeus f.acus, 100°, +20°
29. Asciiris Lacus, 95°, +53°
30. Astaboras, 305°, +26°
31. Astusapes,29S\ +30°
32. Athos, 153°, +48°
33. Atlantis, 173°, -30°
34. Aurorac Sinus, 50°, —13°
35. Aiisonia Australis, 250°, —40°
36. Ausonia Borealis, 275°, —23°
37. Australc Mare, 90°, —65°
38. Azania, 185°, +30°
39. Baltia, 40°, +63°
40. Bathys, 92°, —38°
41. Biblis Fans, 132°, +10°
42. Bidis, 182°, +45°
43. Boreum Mare, 95°, +65°
44. Bosporus Gemrtiatus,
63°, -43°
45. Catlirrhoes Sinus, 3°, +50°
46. Candor, 75°, +5°
47. Casius, 275°, +43°
48. Castorius Lacus, 150°, +55°
49. Cebrenia, 215°, +48°
50. Cecropia Mare, 305°, +67''
51. Ceraunius, 96°, +42°
52. Cerberus, 212°, +9°
53. C>iaos,215°,+35'
54. Chronium Mare, 215°, —60°
55. Chryse, 32°. +8°
56. Chrysokeras, 100°, —52"
57. Cimmerium Mare, 210°, -25°
58. Claritas, 102*, -30°
59. Coloe Palus, 304», +43'
60. Copals Palus, 288°, +58°
61. Crocea, 293°, 0°
62. Cyclopia, 218°, 0°
63. Cydonia, 345°, +50°
64. Daedalia, 120°, -34°
65. Deltoton Sinus, 304°, —5°
66. Deucalionis Regie, 345°, -18°
67. Dcutcronilus, 358°, +35°
68. Dia, 88°, -60°
69. Diacria, 170°, +47°
70. Dioscuria, 315°, +54°
71. £(icn, 350°, +28°
72. Edora, 345°, -4°
73. Elcctris, 190°, -52°
74. Eleus, 168°, +40°
75. Elysium, 215°, +25°
76. £oj, 37°, -15°
77. Erebus, 182°, +20°
78. Eridania, 218°, -45°
79. Erythraeum Mare, 30°, —30°
80. Eunostos, 225°, +15°
81. Eurotas, 125°, +58°
82. Euxinus Lacus, 155°, +43°
83. Fastigium Aryn, 358°, 0°
84. Ganges, m°, +5°
85. Gchon, 358°, +15°
86. Geryon, 75°, —22°
87. Gamer Sinus, 230°, —2°
88. Hades, 192°, +33°
89. Hadriacum Mare, 270°, —40°
90. Hammonis Cornu, 316°, —13°
91. Hellas, 295°, -50°
92. Hellespontica Dcprcssio,
358°, -58°
93. Hellespontus, 330°, —47°
94. Hesperia, 240°, —20°
95. Hiddckel, 347°, +18°
96. Hyblaeus, 228°, +30°
97. Hydrae Pons, 48°, —3°
98. lani Fretum, 10°, —10°
99. lapygia, 295°, -15°
100. Icaria, 124°, -45°
101. Idaeus Fans, 53°, +35°
102. Isidis Regio, 275°, +20°
103. Ismenius Lacus, 335°, +42°
104. Jamuna, 44°, +10°
105. Jaxartes, 22°, +65°
106. Juventae Pons, 62°,— 4°
107. Labotas, 345°, 0°
108. Laocoontis Nodus, 246°, +15°
109. Lemuria, 230°, +70°
110. Libya, 275°, 0°
111. Lunae Lacus, 71°, +15°
112. Mareotis Lacus, 96°, +32°
113. Margaritifer Sinus, 20°, — 10°
114. Memnonia, 142°, —20°
115. Meridian! Sinus, 0°, — 5»
116. Meroe Insula, 290°, +30°
117. Mesogaea, 168°, -2°
118. Midas, 165°, +56°
119. Moab, 338°, +-10°
120. Moeris Lacus, 278°, +8°
121. Nectar, 60°, —28°
122. Neith Regio, 275°, +30°
123. Nepenthes, 268», +8°
124. Neudrus.A'.—M'
125. Niliacus Lacus, 32*, +-27"
126. Nilokeras, 55', +28'
127. Nilosyrtis, 280«, +30«
128. Nilus, 82°, +25°
129. Nix Cydonia, 3°, +40°
130. Nix Lux. 110°, -7°
131
132
133
134,
135.
136.
137.
138.
139.
140.
141.
142.
143.
144.
145.
146.
147.
148.
149.
150.
151.
152.
153.
154.
155.
156.
157.
158.
159.
160.
161.
162.
163.
164.
165.
106.
167.
168.
169.
170.
171.
172.
173
174
175
176
177,
178.
179.
180.
181.
1'82.
183.
184.
185.
186.
187.
188.
189.
190.
191.
Nix Olympica, 132°, +21°
Nix Tanaica, 55°, +52°
Xoachis, 355°, —40°
Nubis Lacus, 264°, +24°
Nymphaeum, 300°, +10°
Oceanidum Mare, 35°, —60°
Ogygis Regio, 60°, —53°
Ophir, 65°, -10°
Ortygia, 350°, +65°
Oxia, 20°, +20°
Oxia Palus, 17°, +-8°
Oxus, 12°, +20°
Panchaia, 205°, +62°
Pandorac Fretum, 345°, —25°
Phaethontis, 150°, —50°
Phison, 308°, +35°
Phlegethon, 125°, +35°
Phlegra, 190°, +45°
Phoenicis Lacus. 110°, —15°
Pierius, 310°, +59°
Pontica Depressio, 85°, —47°
Propontis L 180°, +40°
Propontis H, 179°, +58°
Protei Regio, 50°, —22°
Protonilus, 320°, +42°
Pyriphlegethon, 140°, +20°
Pyrrhae Regio, 30°, —22°
Sabaeus Sinus, 335°, —12°
Scandia, 150°, +66°
Scythes, 75°, +64°
. Serpentis Mare, 320°, —28°
, Sigeus Portus. 335°, -8°
Sinai, 65°, -23°
Sirenura Mare, 140°, —40°
Sitacus, 338°, +17°
.Sithonius Lacus, 230°, +58°
Solis Lacus, 85°, —30°
Stymphalius Lacus, 205°, +54°
Styx, 200°, +25°
Syria, 90°, -20°
Syrtis Major, 290°, +10°
Syrlis Minor, 260°, —8°
Tanais, 50°, +55°
Tempe. 75°, +40°
Tempcs, 63°, +47°
Tharsis, 103°. +8°
Thaumasia, 75°, —35°
Thoana Palus, 256°, +35°
Thotii, 263°, +15°
Thymiamata, 6°, +10°
Tithonius Lacus, 80°, —5°
Trilonis Sinus, 240°, —10°
Trivium Charontis,
198°, +14°
Typhon, 322°, —4°
Tyrrhenum Mare, 270°, —13°
Umbra, 290°, +49°
Utopia, 265°, +56°
Vulcani Pelagus, 25°, —40°
Xanthe. 50°, +15°
Yaonis Regio, 318°, -43°
Zephyria, 190°, 0°
* Italics indicate names not used on
January 1, 1972
the International Astronomical Union's 1958 map of Mars
C. Michaux, JPL
Sec. 3. 5, page 13
Morphology and Processes
JPL 606-1
Fig. 6. 1971 Mars Patrol Photographic Map (Lowell Observatory)
Sec. 3.5, page 14
C. Michaux, JPL
January 1, 197 2
JPL 606-1 Morphology and Processes
will be experiencing midsummer in its Southern Hemisphere. The heliocentric
orbital longitude of Mars will be r| = 45 °, or the areocentric longitude of the
Sun: Lg = 320°. This "planning chart" includes only the relatively stable
features as they appeared in 1941 and 1958, at nearly the same r\ for Mars,
plus updating for the most recently observed secular changes (1969 observa-
tions by both Mariner FE sequences, and by Earth observatories). The 1971
perihelic opposition is to provide further updating data for later versions of the
basic chart. *
The basic cartographic information- -the areodetic net employed - -was
produced by G. de Vaucouleurs (1965, 1969) for his Mars Map Project, begun
in 1958 at Harvard Observatory and completed at the University of Texas in
1969. The basic data is derived from hundreds of both visual observations,
since the time of Schiaparelli (1877), and photographic observations, begun
in 1924. This should ensure the reliability of the control stations (79
altogether). Preliminary comparisons with the NASA Mars Chart 1970 indicate
agreement within ±2°, in both coordinates.
The resolution limit on this Mercator chart is on the order of 50 to
150 km (i.e., 1 " to 3 " areographic), depending upon the degree and gradient of
brightness contrast and quality of the basic photographic data. No attempt was
made to include the topographic details (craters, etc, ) seen in the NE frames
of Mariners 6 and 7 or the FE frame fine structure observed in the Tharsis
region.
The nomenclature is basically that of Antoniadi (1930), with some
revisions and additions recommended by the lAU in 1958, or required by recent
surface changes.
The planning chart is shown in Fig. 7. The list of names identified on
the planning chart is shown in Table 2.
Charts for the South Polar Region and Cap
To supplement their 1971 Planning Chart, the authors later issued
(November 1971) two orthographic projection charts of the South Polar Region
(Fig. 8), showing the cap at two different stages of regression, as follows:
One chart of this region shows nearly all bare ground, with only a very small
(7° across) eccentric residual cap close to the pole, as it appears in Southern
midsummer when Lg = 320°. The other chart represents the same region
overlaid by the expanded polar cap of midspring when Lg = 220°, and actively
regressing; it displays the well-known (recurring every Martian year) pattern
of rifts ( "rima"), brilliant patches ("mons"), and dark patches ("depres sio").
Observations used for the Lg = 320° chart dated from 1941 and 1958,
especially, with others since 1877. Observations used for the Lg = 220° chart
were made at the perihelic oppositions of 1909, 1924, and 1971.
=:=An updated version containing 1971 opposition information up to
September 17, 1971 has already appeared.
January 1, 1972 C. Michaux, JPL Sec. 3.5, page 15
Morphology and Processes
JPL 606-1
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Sec. 3. 5, page 16
C. Michaux, JPL
January 1, 197 2
JPL 606-1
Morphology and Processes
Table 2. List of names used on the Mariner Mars 1971 Planning Chart
(de Vaucouleurs, 1971).
1 Acidalium, Marc
30
+48
2 Achillis Pons
30
+38
3 Acolis
210
-05
4 Aeria
310
+ 15
5 Acthciia
240
+ 35
6 Actliiopis
235
+05
7 Alba
106
+45
8 Ainazonis
150
+03
9 Anuiulus
255
+ 15
10 Aoiiius Sinus
115
-50
11 Arabia
325
+20
12 Aram
13
-05
13 Arcadia
120
+45
14 Argyre I
30
-46
15 Argyre 11
72
-65
16 Ascraeus Lacus
100
+ 18
17 Atlantis
168
-30
18 Aurorac Sintis
50
-14
19 Ausonia
250
-40
20 Australc, Marc
25
-60
21 Bosporus
75
-40
22 Candor
70
23 Capri Cornii
50
-20
24 Ciaralis Pons
155
-42
25 Casius
265
+42
26 Castoriiis Laciis
155
+53
27 Ccbrcnia
215
+45
28 Cerauniiis
95
+25
29 Cerberus
208
+ 10
30 Chcrsoncsus
260
-53
31 Chronium, Marc
180
-60
32 Chrysc
35
+10
33 Cbrysokeras
98
-55
34 Cimmeriiim, Marc
210
-30
35 Claritas
102
-32
36 Coloe Pahzs
299
+ 44
37 Copais Pahis
275
+ 5(i
38 Copratcs
65
-15
39 Crocea
285
-05
40 Cyclopum Sinus
220
-08
41 Cydonia
355
+45
42 Daedal ia
118
-27
43 Dcltoton Sinus
305
-07
44 Dcucalionis Regio
345
-17
45 Deuteronilus
357
+35
46 Dia
85
-60
47 Diacria
163
+48
48 Dioscuria
318
+48
49 Eden
350
+20
50 Edom
51 Elcctris
52 Elysium
53 Eos
54 Eridania
55 Euxinus Lacus
56 Erythraciun, Mare
57 Ganges
58 Gehon
59 Gomcr Sinus
60 Gorgonum Sinus
61 Hadriacum, Mare
62 Hellas
63 Hellespontica
Depressio
64 Hellespontus
65 Herculis Pons
66 Hcspcria
67 Hougcria
68 Hoiigcriiis Lacus
69 lapygia
70 Icaria
71 Idacus Eons
72 Isidis Regio
73 Ismenius Lacus
74 Juvcntae Pons
75 I.aestr\goniim Sinus
76 Libya
77 Lunae Palus
78 Margaritifer Sinus
79 Melas Lacus
80 Memnonia
81 Meridian! Sinus
82 Meroe
83 Mesogaca
84 Moab
85 Moeris Lacus
86 Nectar
87 Ncith Regio
88 Nepenthes
89 Nereidum Fretura
90 Niliacus Lacus
91 Nilokeras
92 Nilosyrtis
93 Nix Olympica
94 Noachis
95 Noctis Lacus
96 Nodus Gordii
97 Oenotria
98 Ogygis Regio
99 Ophir
345
-03
100
Oxia
18
+20
180
-48
101
Oxia Palus
17
4 10
215
-23
102
Palinuri Frtitini
145
-60
37
-15
103
Pandorae Frettnii
345
-25
218
-45
104
Pavonis Lacus
114
157
+43
105
Pbaetliontis
140
-48
30
-33
106
Phkgra
190
+30
60
+05
107
Plioenicis I.aciis
108
-15
357
+ 15
108
Pronietliei Sinus
260
-02
225
-05
109
Propontis I
182
+43
149
-30
110
Propontis 11
177
+55
278
-35
111
Protonilus
318
+42
294
-47
112
Pyrrbac Regie
20
-25
113
Rasena
192
-26
345
-62
114
Sabaeus Sinus
330
-10
323
-40
115
Scamander
197
-48
180
+50
116
Scrpentis, Mare
320
-25
240
-20
117
Simois
160
-18
144
+25
118
Sinai
75
-20
130
+20
119
Sir<'nuin, Mare
155
-32
298
-15
120
Sirentim Sinus
130
-35
123
-40
121
Solis Lacus
85
—27
53
+30
122
Styx
202
+28
275
+20
123
Syria
98
-20
333
+40
124
Syrtis Major
290
+ 12
62
-05
125 Syrtis Minor
260
-10
198
-20
126
Tenipe
68
+45
272
-01
127
Thaumasia
82
-38
65
+20
128
Tliarsis
105
-03
23
-10
129
Tbyle 1
150
-67
73
-13
130
Thyle 11
225
-67
148
-20
131
Thymiamata
5
+ 15
358
-05
132
Tiphys Fretum
220
—57
290
+32
133
Titanum Sinus
168
-20
170
134
TiiluiMiiis Lacus
83
-113
340
+20
135
Iiaitus Albiis
270
+08
(Australisi
95
67
-28
136
Iraclus Albtis
272
+35
(Borealis)
75
+28
265
+ 15
137
I'rinacria
275
-25
55
-45
138
Tritonis Sinus
245
-06
32
+32
139
Trivium Charoniis
200
+ 15
58
+34
140
Fvrrlunum, Mare
255
—OO
280
+43
141
I 'ml)ia
285
-^50
138
+20
142
Itopia
245
+52
350
-45
143
Xantlic
52
+ 12
94
-10
144
Vaonis Frelimi
310
-35
130
-05
145
Vaonis Regio
315
-33
298
-02
146 Zea Lacus
290
-47
63
-42
147 ;
^ephyria
182
-10
65
-10
January 1, 197 2
C. Michaux, JPL
Sec. 3. 5, page 17
Morphology and Processes
JPL 606-1
O -
O
DC
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a.
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Sec. 3. 5, page IS
C. Michaux, JPL
January 1, 197Z
JPL 606-1 Morphology and Processes
Mariner Mars 1969 (J. Cutts)
A photomap of Mars, Fig, 9, assembled from Mercator projection
sections of the Mariner 6 and 7 FE pictures and obtained by a special computer
program (GEOM), transforming a quadrilateral into a rectangle, was produced
in the JPL Image Processing Laboratory, using the technique described by
Rindfleisch et al. (1971).* The photomap accurately provides, with no
subjective interpretations, the appearances and positions of the Martian surface
features at the time of photography (end of July-early August 1969), for lati-
tudes between 50°N and 70''S. The areodetic positions were established from
the Davies and Berg (1971) control net for Mars. When first assembling the
Mercator sections, a problem was encountered due to a sharp discontinuity in
brightness between edges of segments, caused by differences in solar illumina-
tion and viewing angle. By using an appropriate photometric function
(Minnaert's model, see Section 3.2 on Photometric Properties), it was possible
to obtain matching of the sections. This was done by selecting parameter k in
Minnaert's photometric function B/Bq = (cos i)k (cos e)^-l, as k = 0.6.
The photomap of Mars is illustrated in Fig. 9, including a diagram of
the GEOM transformation scheme used by Cutts et al. (1971).
Mariner Mars 1969 Meridiani Sinus Region Map (C. Cross)
Cross (1971a) has produced a map of the Meridiani Sinus region (see
Fig. 10), which incorporates both the albedo and the topography information
from the Mariner 6 and 7 TV pictures. He used the fine detail revealed by
the "Max-D" versions, as well as the true -intensity contrasts provided by the
"photometric" versions of the NE pictures. The overall albedos were adjusted
to match those of the FE pictures. The result was a "compromise" map
blending the two aspects. The terrain features were rendered by an artist's
charcoal-and-stump technique. The positional accuracy of the feature locations
and orientations were assured by transferring these from the NASA Mariner
Mars 1969 Chart, using the available control points of the areodetic net of
Davies and Berg (1971).
The northern and southern border regions of Fig. 10 are drawn only
from the FE pictures and indicate a near-absence of craters. Solar elevation
is near zenithal for the western edge of the chart, where there are no
shadows, and about 75° from zenith toward the eastern edge (near 350° W),
where shadows become pronounced.
Of particular interest are the dark highland promontory at the northern
tip of Margaritifer Sinus (near 17 "W), the area of chaotic terrain (centered at
35 °W) with channels extending into both Margaritifer Sinus and Chryse, and the
very large craters (300 and 225 km) in Edom and Deucalionis Regio.
*A similar computer photomap of Mars, but at much higher resolution
(obtained by using a much finer grid system) is presently under preparation
by R. B. Leighton of C. I. T.
January 1, 1972 C. Michaux, JPL Sec. 3.5, page 19
Morphology and Processes
JPL 606-1
NCRT H
30° N ■
30 'S
■ ' s
I
1
>T^ '10"
30' 0° 3 30' 300- 2 .'0°
WFST LONGir,0E.
'JO'
Fig. 9. Mariner Mars 1969 Photomap (J. Cutts).
Mariner Mars 1 969 South Polar Region Map (C. Cross)
Cross (197 1b) also prepared a stereographic projection chart of the
portion of the South Polar Region and Cap photographed by Mariner 7, Fig, 11.
Again he combined albedo and topography, using the same charcoal rendering
technique. Positioning of features was based upon the areodetic control points
for the South Polar Region used in making the NASA Mariner Mars 1969 polar
chart.
Sec, 3, 5, page ZO
C, Michaux, JPL
January 1, 1972
JPL 606-1
Morphology and Processes
35 30 25 20
NORTH
15 10 5
3^5 350 345 34Q 335
MERCATOR PROJECTION
Scalp 1:25,000,000 at the equator
500 1000
KILOMETERS
Fig. 10. Mariner Mars 1969 Meridiani Sinus Region Map (C. Cross).
3.5.3 TYPES OF TERRAINS
Cratered Terrain
Mariner 4 Photography*
In 1965, Mariner 4 discovered heavily cratered terrain on Mars. The
most striking morphological features seen on the photographs were these
numerous craters of smooth and subdued appearance, ranging in size from
-See Fig. 24 (Section 3. 5) from the first edition (1968) of this document, which
shows Mariner 4 pictures 7-12. The complete set of 22 pictures was presented
and fully documented in the final report by Leighton et al. , 1967.
January 1, 197 2
C. Michaux, JPL
Sec. 3. 5, page 21
Morphology and Processes
JPL 606-1
SlTREOOR.M'Mir PROIECIION
Sidle 1:^5, 1)00, 'i(Kl
50°
50°
70°
80'
KILOMETERS
Fig. 11. Mariner Mars 1969 South Polar Region Map (C. Cross).
the limit of available resolution (3 km) up to at least 180 km. Over 600 definite
and possible craters were counted in Mariner 4 frames 4N3-l6.=;= The craters
were classified - -rather subjectively- -according to three categories, based on
degree of preservation: well-preserved, of intermediate preservation, and
poorly preserved. Table 3 gives a summary of the compilation by the principal
investigators (Leighton et al. , 1967), for successive picture pairs. The table
includes the diameters of the smallest and largest craters, as well as the
number of craters displaying central peaks and polygonal outlines. It is seen
that the percentage of well-preserved craters is not high (14%), and that they
are more abundant in pictures 4N7-14, located, in most part, across the
^Picture designation conforms to that used for Mariners 6 and 7 explained in
Section 3.6, Photographic Atlas.
Sec. 3. 5, page 2Z
C. Michaux, JPL
January 1, 197 2
JPL 606-1
Morphology and Processes
Table 3. Summary compilation of Martian crater data from Mariner 4
pictures 4N3-16 (Leighton et al. , 1967).
Pictures
Total
number
Ccntrol
peaks
Polygonal
outline
Smollett
die., km
Largest
dia., km
Well
preserved
Intermediote
preservation
Poorly
preserved
3, 4
88
1
6
80
5
37
46
5,6
92
6
2
5
64
9
40
43
7,8
115
7
10
4
180
18
45
52
9, 10
132
11
10
4
123
19
64
49
11,12
97
S
9
3
175
20
42
35
13, 14
67
4
4
5
350(?)
12
19
36
15. 16
45
2
3
7
1U
3
22
20
636
39
38
Smolleit
3 km
Lorgest
350(?| km
86
269
281
Mare Sirenum-Mare Cimnierium complex. Only two craters of 3 -km size were
counted, while a very large one (350 km) seemed present (its rim appears on
4N13). About 6% of the craters displayed polygonal outlines or central peaks.
Few display both characteristics. The numbers given are considered some-
what uncertain, however, since crater recognition and classification on the
Mariner 4 pictures is difficult, mainly due to the limited intensity discrimin-
ability of the imaging system and the poor photographic rendition under the
unfavorable high-Sun illumination conditions, which created little if any
shadow in the best pictures (4N7-12). Recognition of craters in the 3-4 km
size range was especially difficult.
Mariner 4 Crater Statistics and Analyses
Leighton et al. (1967) Analysis . The cumulative size -frequency
distribution of craters recognized in 4N7-12 was plotted, with subdivision into
"definite, " "probable, " and "possible" categories, reflecting degree of rim
preservation (see Fig. 12). This plot shows for the 20-60 km sizes, or above,
a curve slope approximating that for the lunar upland craters and an absolute
abundance only slightly lower. For sizes under 20 km, the Mars curve flattens
out to lower frequencies. This bend, indicating a deficiency in small crater
abundance, may be considered real (as did Leighton et al. , 1967), and inter-
preted as due to crater modification processes (erosion and deposition), which
become more effective on smaller craters; or it may be considered suspect
(as did Chapman et al. , 1969) and attributed in large part to observational
incompleteness in the crater counts. In any case, the Martian craters,
whether all counted or not, generally appear much more eroded than those of
the lunar uplands.
January 1, 1972
C. Michaux, JPL
Sec. 3. 5, page 23
Morphology and Processes
JPL 606-1
2000
lOoo
500
1 TA — r
ALL CRATERS \
E
10
O
200
- 100
A
LUNAR UPLANDS
50
20
10
WELL-PRESERVED
CRATERS ONLY
10 20 50
DIAMETER, km
100 200 500
Fig. 12. Cumulative size-frequency distribution of craters recognized
in Mariner 4 pictures 4N7-12 (Leighton et al. , 1967).
Other observations made by the investigators from picture examination
were as follows:
1) Central peaks and polygonal outlines were associated mostly with
the relatively well-preserved craters.
2) Well-preserved craters are more abundant in the small and
medium sizes, while many of the larger craters are faint.
3) Polygonalization was more abundant in intermediate -size craters
(15-45 km).
4) Irregular circular outlines characterize a number of craters,
particularly the largest ones; and features resembling slump blocks
are found at the base of their inner wall slopes.
5) Some of the larger craters are girdled by rough terrain suggesting
ejecta sheets of rubble.
6) Smaller crater clustering in random fashion was noted.
Sec. 3. 5, page 24
C. Michaux, JPL
January 1, 1972
JPL 606-1
Morphology and Processes
7) One or two broad domes were observed in crater floors.
8) Albedo differences between crater floors and surrounding terrain
were observed, but without consistent pattern. Floors were either
lighter or darker.
Chapman et al. (1969) Analysis . Two years after the final report of
Leighton et al. (1967), Chapman et al. (1969) issued a more complete statistical
analysis. The Chapman et al. conclusions are essentially similar, and will not
be elaborated upon here. Their classification system consisted of four classes
of crater degradation (Classes 1, 2, 3 and 4 represented fresh, less fresh,
eroded, and 'ghosts' in essence, respectively), and a quality identification
(Qualities A, B, and C denoted definite, probable, or uncertain crater,
respectively). Some of the tabular results are shown in Table 4.
Mariner 6 and 7 Photography
The three Mariners have shown that cratered terrain is the most
common type of terrain encountered on Mars. Cratered terrain seen in the
Mariner 6 and 7 pictures were defined by Murray et al. (1971) as terrain in
which craters are the dominant landforms recognizable at the available
resolution (0.3 km at best).
The far-encounter (FE) views as well as the near-encounter (NE) views
of the 1969 flybys have revealed that Mars, like the Moon, has an abundance of
craters of all sizes, from the limit of resolution (0.3 km) up to several hundred
kilometers across, and that there are roughly two distinct morphological types,
which are described below, summarizing from Murray et al. (1971).
The first type, large flat-bottomed craters, ranges in diameter from
about 15 km to several hundred km. They are highly modified from their
presumed initial appearance as impact craters. Rims are missing or greatly
subdued, central peaks rare and, if present, also quite levelled, rays are
absent; the ejecta blankets or swarms of secondary craters commonly associated
with large impact craters are degraded generally beyond recognition. Some,
termed vestigal or 'ghost' craters, have only faintly visible wall relief. Thus,
Table 4. Crater percentages by class at several diameter intervals
for Mars (Chapman et al. , 1969).
January 1, 1972
Region
Class
Diameter Interval (km)
5-10
10-15
15-20
20-30
30-60
>60
Mars:
pictures 4N7-14,
Quality A
and B
1
2
38
58
11
26
14
24
4
9
4
8
14
3
2
43
29
46
43
43
4
2
20
33
41
54
43
C. Michaux, JPL
Sec. 3. 5, page 25
J n JPL 606-1
Morphology and Processes
one may distinguish two extreme states of preservation, or two ^ub Ypes of
flat-bottomed craters (see. for example, 6NI6 for comparison). Although
various degrees of preservation can be seen within one crater popula ion. no
completely satisfying transition has yet emerged; perhaps this reflects complex
episodic modification processes.
The second type, small bowl-shaped craters, is the majority of those
with diameters below 10-15 km. Of fresh-looking appearance, th^^ 7/^";bl^
the small lunar primary impact craters with their associated impact features,
at lea't as far as the limit of resolution permits to see; for example one cannot
expect to see slump-blocks or secondary crater swarms. Apparently of
remarkable uniformity, their morphology shows little degradation; hence, they
must be the product of recent impacts.
In comparison with the lunar upland craters, the Martianflat-bottomed
craters are less numerous and more highly degraded, and their mtercrater
areas are smoother. No large fresh-looking crater, such as Tycho. was seen;
no rays and secondary crater swarms are present.
The local irregularities of the large old crater walls are usually pre-
served desp te the smLthing of elevated rims and ejecta sheets This indicates
;ronounccd\orizontal and regional redistribution of the material, rather than
a local one from impact fragmentation and slumping as on the Moon.
Neither dark maria nor lava-flooded plains were detected, but this does
not rule out their presence. Detection may be difficult if they have been thickly
covered by dust.
Features other than craters are present in this terrain, such as sinuous
channels and ridges, as well as some short linear subparallel markings. No
sinuous "rilles. ''flow fronts, and partially flooded craters, which are so
characteristic of lunar maria. are seen.
The small bowl-shaped craters seem to be morphologically similar to
those of the Moon.
Correlations were sought by the CIT team (Murray et ^1-, 197 1). but no
significant correlations were found between cratered torram and latitude, or
Tafk/brtght areas, or topography. However, the analysis was only general.
Mariner 6 and 7 Crater Statistics and Analyses
Complete crater statistics covering the entire set of pictures returned
by the tw spacecraft in 1969 are not yet available. Only three studies of size-
fJequency distributions of craters can be presented here, with their tentative
conclusions, and with no attempt to unify them.
resul ti
Sec. 3.5. page 26 C. Michaux. JPL January 1, 1972
JPL 606-1
Morphology and Processes
■10''
CVJ
ID
o
a:
UJ
0. 1000
tr
UJ
»-
UJ
<
Q 100
z
<
I
ijj
I-
<
LlI
cr
o
10
CD
Z
SN-ie.aOAND 22
AVERAGED
6 N 17, 19 AND 21 AVERAGED
SBC
- LFB
±
1
.1
100
^ 10
CRATER DIAMETER D (KM)
Fig. 13. Cumulative size-frequency distribution of craters
in Deucalionis Regio (Murray et al. , 1971).
are in the traditional "lunar" form; i. e. , logarithmic plots in terms of
cumulative numbers of craters larger than a certain diameter per unit area
(km per 10^ km^).
Examination of Fig. 14 disclosed that (1) there are no major areographic
variations in the large (flat-bottom) craters density in Deucalionis Regio (data
is from A-frames), and (2) there are apparently large areographic variations
in the small bowl-shaped craters (data is from both A- and B-frames), but
these are considered by the authors as not more than "possible minor" varia-
tions, because they feel that their counts are unrepresentative of the small
crater population in the 5-15 km range, due to inadequate resolution of the
A-frames and insufficient areal coverage of the B-frames.
Comparison with lunar crater curves was also made (see Fig. 15). The
general conclusion was that the Martian craters have a similar overall distribu-
tion in form, except that (1) there is under saturation for small and large
Martian craters, and (2) the distribution for the small craters diverges sig-
nificantly from that of the lunar uplands as size decreases, which again may be
due to unrepresentative counts.
From the uniformity of appearance of small bowl-shaped craters,
Murray et al. (1971) believe that the more recent crater modification history
on Mars has probably been episodes of crater removal coupled with continuous
rate of formation.
January I, 1972
C. Michaux, JPL
Sec. 3. 5, page 27
Morphology and Processes
JPL 606-1
Fig.
10 100
CRATER DIAMETEn D (KM)
14. Plots of crater abundances (similar to those in Fig 13) for individual
wide-angle (A) and narrow-angle (B) frames (Murray et al. . 1971).
J XT- ^ l^Q■7^\ Analv<?is Woronow and King (1971) made
,„„... ZTlTuet^Z li::i'^S^"iaia.e surroundings in Deucaiionis
Regio, Thymiamata, and Edom.
Their analysis of size-frequency distributions of crater s on both
A-frames (6N11, 13. 19. and 7N25) and B -frames (6N10. 12. 18. 20, and mb.
22) obtained the following results:
1) The size-frequency curves for the four A-frames closely coincide
(Fig 16). Such a distribution may be representative m general
of the Martian crater population observed on the large scale
recorded by the A-cameras.
2)
The curves for the B-frames do not cluster as well as A-frames
and diverge most at diameters well above the B-frame resolution
Wt (pYg 17). The observed discrepancies appear to be due to
aTomMnltion of two reasons: (1) insufficient number o^ c-t
on any one image to provide a significant sample ^^^^ ^^P^^^,^^ °!,^^
deftni^g the populations, and (2) true differences m the populations
of the smaller craters in the different image areas.
Sec. 3.5, page 28
C. Michaux, JPL
January 1, 197 2
JPL 606-1
Morphology and Processes
10-
o
cc
LlI
a.
o
f^ -1000
<
<
tr
cr
LlJ
m
3
-100
•10
MARS COMPARED WITH
LUNAR MARIA
.RANGERSm
M^ffF TRANQUIL-
LI TAT IS
RANGERSn
MARE
COGNITUM
MARS COMPARED WITH
LUNAR UPLANDS
SOUTH POLAR
REGION
10 -100 1
CRATER DIAMETER D (KM)
TSIOLKOVSKY REGION -
lO'O
Fig. 15. The Deucalionis Regio crater abundances of Fig. 13
compared with those of the lunar maria (left) and the
uplands (right) (Murray et al. , 1971).
3) When the bright and dark area portions, which cover Meridiani
Sinus and immediate surroundings in Deucalionis Regio,
Thymiamata, and Edom (frames 6N11, 13, and 19), were plotted
separately, a significant displacement of the two curves was
found in the 20-50 km diameter range. The divergence is due to
a greater percentage of large craters on the bright area terrain
and a greater percentage of small craters (less than 15 km) on the
dark area terrain.
4) A greater total crater density is apparent in the dark area.
The authors' interpretation of item 3) and 4) results is as follows:
A longer impact exposure age for the bright areas (to account for its
greater density of larger craters).
A possible greater production of endogenetic craters in the dark areas
(to account for its greater total crater density).
January 1, 197Z
C. Michaux, JPL
Sec. 3. 5, page Z9
Morphology and Processes
JPL 606-1
999
99 5
99
98
95
90
I 80
U
t 70
a.
60
1 50
I 40
=• 30
o
20
10
5
2
1
05
01
20 40 60 80 100
Crater diameter (dm)
Fig. 16. Cumulative size-frequency probability distributions of craters
found in wide-angle frames (Woronow and King, 1971 and 1972).
0^5-
0.1-
6NI0,6NI2,eNI8
6N20, 7N6, 7NE2
»9.e-
99.9-
CRtTCII DKMETEK I
k:?
Fig. 17. Cumulative size-frequency probability distribution of craters
found in six narrow-angle frames (Woronow and King, 1971).
NOTE: The "frequency distributions were plotted with probability ordinates,
so that random samples from like populations or the same population
would graph identically, provided that the sample size is sufficient"
(Woronow and King, 1972). "This eliminates the necessity of
normalizing the crater count to a unit area" (Woronow and King, 1971).
Sec. 3. 5, page 30
C. Michaux, JPL
January 1, 197 2
JPL 606-1 Morphology and Processes
More rapid erosion and filling of crater s - -particularly the small sizes --
in the bright areas (to account for the greater percentage of smaller
craters and greater total crater density in the dark areas).
McGill and Wise (1971) Analysis . McGill and Wise (197 1) obtained more
extensive statistics aimed toward a study of the regional variations in average
density of craters of various sizes and their degree of topographic degradation,
for an area about 4.6 million km^ lying within the four regions Meridiani Sinus,
Margaritifer Sinus, Deucalionis Regio, and Hellespontus -Noachis . They used
frames 6N10, 12, 13, 16, 17, 18, 19, 20, 21, 22, and 7N24, 25, 26. Each
crater was classified according to its size and a "degradation number, '' which
is the sum of score values from 1 (fresh) to 4 (highly degraded) assigned to
their rim, their inner wall, and their floor, as systematized in Table 5. Thus
the degradation number or class may range from 3 (sharpest craters) to 12
(barely visible 'ghosts'). The Mariner imagery resolution permitted classifica-
tion by this degradation number of craters in only two size ranges: ''small
craters" (diameters 1-8 km) on B-frames and "large craters" (diameters
greater than 16 km) on A-frames.
Results:
1) Density of Craters
The size frequency distribution plots (Fig. 18) resemble those
of other authors, but they reveal that significant regional
differences exist only for small craters (1-8 km). The fall -off
toward the limit of resolution is greatest for Hellespontus -Noachis,
then Meridiani Sinus, and least for Deucalionis Regio. Since this
order is the same as that of their average degradation numbers,
(see 2) Degree of Degradation), it suggests that the fall-off is due
to the fact that difficulty in recognizing small craters increases
with degree of degradation.
2) Degree of Degradation of Craters
The comparison of average degradation numbers for small and
large craters in each of the four regions (see Fig. 19) shows that
there are significant regional differences only for the small crater
category. Thus, Deucalionis Regio appears to have fresher small
craters on the average than has Meridiani Sinus.
Furthermiore, the plots (not shown here) of degradation nunibc r
versus crater density, for the various sizes, indicate that
disparity in degradation exists only in the small crater sizes,
becoining increasingly apparent as size decreases, and that
marked distinctions exist between the regions. Thus, Deucalionis
Regio has a great abundance of fresh 1-2 km craters, while
Margaritifer Sinus has many moderately degraded small craters.
January 1, 197 2 C. Michaux, JPT Sec. 3.5, page 31
Morphology and Processes
JPL 606-1
Table 5. Classification of Martian craters by degradation number
(McGill and Wise, 197 1).
Total degradation number for a crater is the sum of values for rim,
inner wall, and floor.
Point
Value
3
4
Rim
Sharp, strong
relief
Moderately
strong relief
Barely visible
Completely
absent
Inner "Wall
High and
steep
Moderately
strong relief
Barely visible
Completely
absent
Floor
Cup-shaped (small craters)
or with well-defined central
peak (large craters)
Part of floor, flat and
featureless
Mostly flat and featureless
Entirely flat and featureless
The contrasts in degradation-density curves of the small (1-8 km)
craters in between the four regions are summarized in Fig.^ 20.
Of the three possible interpretations given by McGill and Wise,
their favored was a uniform degradation of small craters formed
locally in episodes (or bursts). All four curves show an increase
in crater density toward the degraded end of the plot, where, they
suggest, a steady-state distribution is reached (assuming that the
degradation number is some function of age). In Fig. 21, they
illustrate the hypothetical chronological sequence of degradation-
density curves for small craters.
For the large craters, McGill and Wise find a wide range of
degradation numbers, but with the density progressively increasing
with degradation. This argues against a twofold division into
"degraded" craters predating and "not degraded" craters post-
dating some global catastrophic event. They caution that this may
not apply to "the very largest and oldest craters, which are
inadequately sampled at present and may represent a distinctly
different, primordial catastrophic period. " They find the apparent
absence of fresh, large craters (Copernicus -type) "expectable if
their predicted rate of formation is very small compared to the rate
at which topographic details are modified by Martian surface
processes" and that the expected abundance of such craters
estimated by Murray et al. (197 1) was "based on an extrapolation
of the apparently anomalous high density of small fresh craters m
Deucalionis Regio, "
Sec. 3. 5, page 32
C. Michaux, JPL
January 1, 197 2
JPL 606-1
Morphology and Processes
O
O-
CvJ P.
CO O
q: ~
UJ
I-
go
O
CC
LU
Q_
O
t/) O
a: ~
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q:
o
u.
O
q:
LU
GO
o-
LUNAR TERRAE
MODIFIED FROM
HARTMANN, 1967)
.25
HELLESPCNTUS - NOACH I S
MERIDIAN! SINUS
MARGARITIFER SINUS
DEUCALIONIS REGIO
LUNAR MARIA X^
(MODIFIED FROM
HARTMANN, 1967)
.5
2
4
DIAMETER IN
-I 1 1 —
8 16 32
KILOMETERS
64 128
Fig. 18. Size -frequency distribution of Martian craters in four regions
(McGiU and Wise, 1971).
The almost continuous sequence of degradation found by McGill and Wise
for small Martian craters was not accepted by some geologists (for exan^iple,
Murray, Soderbloom) on the grounds that small degraded craters are too
difficult to see and to classify reliably, especially when seen at large slanting
ranges at the resolution of the B -frames of Mariners 6 and 7. Also, the
subjectivity of the new classification scheme yielding a degradation number for
each crater has been looked upon suspiciously. Nevertheless, the authors
claim that useful, "reasonably reliable data' has been obtained by their scheme.
January 1, 197Z
C. Michaux, JPL
Sec. 3.5, page 33
Morphology and Processes
JPL 606-1
1-8 KM DIAMETER CRATERS
cr
CD '^'
g .
o
<
"^ 7-
^
GRAND AVERAGE,
232 ' 1-8 KM
CRATERS
LjJ 6-
O
<
(T
O
LlI 5-
CM
^
Z
OJ
(M
CM
00
CD
CM
CD
^11- +16 KM DIAMETER CRATERS
UJ GRAND AVERAGE, 557
^ /+I6 KM CRATERS
9-
<
(r
o
UJ
o 6-
<
ir
UJ
4-
(T>
^D
Fig. 19. Average degradation numbera for smrJl craters (above) and large
craters (below) in four Martian regions (McGill and Wise, 1971).
Sec. 3.5, page 34
C. Michaux, JPL
January 1, 197 2
JPL 606-1
Morphology and Processes
1 r
9 10 II
- ^''^'^^ DEGRADATION NUMBER ^DEGRADED)
Fig. ZO. Summary plots contrasting distribution of small (1-8 km diameter)
craters among degradation classes in four Martian regions
(McGill and Wise, 1971).
Crater Modification Processes
From analysis of the Mariner 4, 6, and 7 pictures, only a few crater
modification processes are thought to be operative under the present conditions
existing on Mars, which lacks substantial amounts of water. There are three
important exogene processes to consider, as follows:
1) Meteoroidal impact fragmentation and obliteration must be similar
to that on the Moon, except for the slight retarding effect of the thin
Martian atmosphere. Dycus (1969), however, calculated that the
5-mb atmosphere has no appreciable effect upon the impact velocity
of objects creating craters larger than about 10 meters. Direct-hit
impacts can partially or completely obliterate the existing craters,
depending upon the size, velocity and course of the impacting body.
Significant erosion or damage by a nearby impact still can occur,
either through the filling in of the existing crater by ejection of
debris or indirectly by mass wasting (downslope mass movement)
through the propagating shock wave along the surface.
January 1, 1972
C. Michaux, JPL
Sec. 3.5, page 35
Morphology and Processes
JPL 606-1
FRESH — ► DEGRADED
(I) STEADY-STATE
DISTRIBUTION
i
(2) EXCESS FRESH
CRATERS AFTER
EPISODE OF CRATER
FORMING ACTIVITY
I
(3) EXCESS CRATERS
PARTIALLY
DEGRADED
I
(4) EXCESS CRATERS
EXTREMELY
DEGRADED
i
(5) excess craters
obliterated:
return to
steady- state
DEGRADATION NUMBER
Stages (1) through (5) represent a chronological sequence with unknown time
scale.
Fig. Zl. Model explaining differences in degradation - density curves
for small craters from four Martian regions
(McGill and Wise, 197 1).
Sec. 3. 5, page 36
C. Michaux, JPL
January 1, 197Z
JPL 606-1 Morphology and Processes
One certain result is that much impact rubble is produced , and the
Mars surface , with its many craters, must be thickly mantled by
such debris. Chapman et al. (1969) attempted quantitative treat-
ment of the impact erosion process, sometimes called pelting, and
found that it seems to account for much of the damage sustained by
the Martian craters .
2) Aeolian erosion and deposition can be very effective on a dusty/
sandy surface, such as that of Mars, provided the winds are suffi-
sufficiently strong and frequent. The aeolian mechanism is through
transport of dust or sand- sized particles, either by suspension,
saltation, or traction, depending upon the strength of the wind. It
appears from the studies of Gierasch and Sagan (1971 ) that local
topographic (or slope) winds, induced thermally by sharp relief on
Mars,, can alone reach considerable velocities of the order of 50
msec (180 knnhr" ) and even greater (100 msec ) above the sur-
face boundary layer; and, if superimposed on the global seasonal
thermal wind (which itself can reach 50 msec , according to the
Leovy and Mintz 1969 calculations), such winds appear to be an
adequate driving force, since it was estimated that wind velocities
about 80 msec were required to raise particles 400 fx in dia-
meter in low areas under a surface pressure of p = 10 mib (Sagan
and Pollack ,1967,1969). This recent optimistic view has stemnned
from the new knowledge that Martian topography is highly variable.
Similar calculations were performed by Arvidson (1972) for surface
pressures of 6 to 7. 6 mb, and the results are given here: see Fig.
22 for the threshold drag velocities and Fig. 23 for the corresponding
lowest threshold wind velocities (at 1 m. ). Settling velocities for
spherical particles have also been estimated by Arvidson (1972), see
Fig. 24.
Another line of studies undertaken by Ryan (1969) and centered on
explaining the yellow clouds on Mars as resulting from a dust-devil
generation mechanism concludes that the high threshold velocities
required for the suspension of dust under surface pressures p^ = 5
to 10 mb are likely to be attained in the larger of these vortices.
Crater obliteration on Mars by aeolian dust deposition (especially
of small craters) has recently been discussed by Hartmann (1971b),
3) Mass-wasting, namely the downslope movement of loosened material
( rock, debris, soil) under the influence of gravity is another crater
modifying or levelling process which must be operative on Mars.
Shearing stresses due to gravity exist in every slope, constantly
causing a very slow, continuous downslope movement known as
'creep'. When the shearing force reaches the failure point, a sud-
den rapid collapse of mass called 'landslide' can occur. Mass-was-
ting is promoted by thermal expansions and contractions, freeze
and thaw, seismic events, vibrations from meteoroidal impacts,
loading at top of slope (by aeolian deposition , for examiple), also
weathering (decreasing the strength of the material).
January 1, 1972 C. Michaux, JPL Sec. 3.5, page 37
Morphology and Processes
JPL 606-1
4.0
>-
(J
O
_i
UJ
>
<
Q
O
X
3.0
2.0
T?
:v.nC^>:-;::"'>>v*."- •
:>:S:V;:;;v:<.v:.-.W"::-;
'."'.'.■•.'■I'i
MARS
(p.
■jii':.'i-f^:':i:'--^'>^>'^^:- ;'^.v'.''";-:.''
7.6 mb) V^/JS:i$gV^^JSS^'
^
.004
.01 .10
PARTICLE DIAMETER (cm)
1.0
Fig. 22. Threshold drag velocities plotted over a range of particle sizes
for Mars and Earth. Earth curve after Bagnold (1941). Particle
density for Martian curves = 3.0 gm cm"^ (Arvidson, 1972).
Thermal creep was advanced by Sharp (1968) as probably quite
effective in modifying Martian craters, because of the large diurnal
fluctuations of temperature. It is difficult yet, however, to assess
its importance on Mars, since the effectiveness of this process
depends on many parameters (slope angles, structure of the rubble
layers, coefficients of friction, etc.) in addition to temperature
fluctuations.
Processes such as thermal fracturing are considered ineffective on the
Moon by Ryan (1962) while freeze and thaw proposed by Wade and DeWys (1968)
requires a substantial permafrost layer, thus being still speculative for Mars.
Other possible crater- modifying processes, not evident in the cratered
terrain photographed by Mariner 6 and 7, would include endogene processes:
volcanism under various forms (lava flows and flooding of crater floors, ash
deposits, etc.) and tectonism (crustal readjustment through faulting), orogeny
(mountain building), and geothermal activity. Thus for example. Pike (1971)
drawing from certain similarities between lunar and Martian craters above
10-20 km indicated the probable importance of: (1) slow or gradual tectonic
adjustment, and (2) localized magmatism in large Martian craters. Green
(1971) claims evidence of volcanism in some large craters observed by
Mariner 4.
Sec. 3.5, page 38 C. Michaux, JPL January 1, 1972
JPL 606-1
Morphology and Processes
ru
1 —
-1 1—
' 1
■
.. (p^ = 6 mb)
MARS
60
*>■:■:.•■'■.■■.■:'
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OLD
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-2
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—
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20
^
^
^
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10
ft
—
EARTH
-
,
1 1
1 1
.
.001
.01
.03
SURFACE ROUGHNESS (cm)
Fig. 23. Lowest threshold wind velocities for Mars and Earth (1 meter
above surface). Maximum surface roughness for terrestrial
deserts = 0.03 cm (Ryan, 1964). Lower bound was
set arbitrarily (Arvidson, 1972).
Age of the Large Craters
The large and well-eroded, flat-bottomed craters, visible in pictures
obtained by all three Mariners, are undoubtedly very ancient structures,
apparently several aeons old. They may be "primeval" craters dating from the
end of the planet's major accretional phase, or they can be "later" craters.
Differences of opinion concerning these craters existed between the experts
in 1965. After Leighton et al. (1965) announced that the craters must be
2-5 billion years old, a series of criticisms appeared (Anders, 1965; Baldwin,
1965; Witting et al. ,1965) with much lower estimates, around 1 billion years
old. Subsequently, Opik (1965), Hartmann (1966), Binder (1969), and Chapman
et al. (1969) estimated that the craters must be at least 4 billion years old,
and this estimate became generally accepted. After the Mariner 6 and 7 fly-
bys, Murray et al. (1971) declared that these ancient structures must be the
survivors of the accretional phase some 4.5 billion years ago, while Hartmann
(1971a) finds as a result of his analysis, that they cannot go that far back and
apparently are slightly later craters surviving from an early intense erosional
phase on Mars. Three lines of evidence are given by Hartmann: (1) under-
saturation compared to lunar upland craters, (2) possible preservation of the
January 1, 1972
C. Michaux, JPL
Sec. 3. 5, page 39
Morphology and Processes
JPL 606-1
!
J
^
o
z
til
M
10
10
-3
y * MARS
.0001
.001
.01
.10
PARTICLE DIAMETER (em)
Fig. 24. Settling velocities over a range of particle sizes for Mars and
Earth. Earth curve after Ryan (1964). Martian curve is an average
of the curves for surface pressures p = 6 and 7. 6 mb. (Arvidson, 1972).
fossil asteroid mass distribution, and (3) extrapolation backward a few aeons
of the asteroidal impact rate can account for most of the craters. He
hypothesizes that a severe erosional phase wiped out all the original accretional
craters, and that those we currently observe are the product of a lesser, but
still intense, bombardment period which extended for some 10^ years after
accretion. One very pertinent observation is that the intercrater areas are
quite smooth, reflecting the mentioned under saturation of craters. An early
short-lived, dense atmosphere on Mars has been proposed by several authors:
on lithological grounds by Binder and Cruikshank (1966), and on geochemical
grounds by Fanale (1971).
Sec. 3. 5, page 40
C. Michaux, JPL
January 1, 1972
JPL 606-1 Morphology and Processes
Chaotic Terrain
The Mariner 6 TV photography revealed a rough, uncratered type of
terrain which Sharp et al. (1971a) described as "chaotic" because of its
appearance as an irregular jumble of topographic forms. "Chaotic terrain
consists of a rough, irregular complex of short ridges, knobs, and irregularly
shaped troughs and depressions" at the kilometer scale (Sharp et al. , 1971a).
It is best seen in the high-resolution (B) frames 6N6, 8, and 14. In frame 6N6,
a northeasterly grain is noticeable. Chaotic terrain essentially lacks recogniz-
able craters; only three small, faint, marginal ones were "tentatively
identified. " It is difficult to recognize small craters (< 5 km) amidst the jumble
of features. The contact line between chaotic and cratered terrain is irregular
and not always well defined. The 1;ransition from cratered to chaotic terrain
often is sharp and marked by abrupt structural patterns: apparently arcuate
scarps, blocks and ridges with intervening depressions on the chaotic side,
and apparently huge cracks on the cratered side of the contact line (see 6N8
and 14), which suggests a definitely lower elevation of chaotic terrain. Con-
firmation of this lower elevation was in fact given by the Lincoln Laboratory
(1970) radar results, as well as by the Mariner IRS C02-pressure results
(Herr et al. , 1970) over these equatorial areas.
The albedo of chaotic terrain often contrasts with adjoining cratered
terrain; usually, but not necessarily, it appears brighter.
Distribution
Subsequent to the discovery of some 12,500 km^ of characteristic
chaotic terrain in the B -frames previously mentioned, a much greater area--
about 1.5 million km^--of possible chaotic terrain was delineated in A-frames
6N5, 7, and 9, as illustrated in Fig. 25. This interpretive map was constructed
by extending the chaotic -cratered contact line, traced from the B -frames onto
the A-frames, "on the basis of structural patterns, crater distribution, bright-
ness contrast, and characteristic regional trends" (Sharp et al. , 1971a). The
distribution of postulated chaotic terrain appears highly irregular, like lace-
work, with odd-shaped, sometimes disconnected, patches. Areographically,
this distribution is within equatorial latitudes 15 "N to 15 "S and is centered in
Pyrrhae Regio (a mixed dark and bright area), with extensions into the dark
areas Aurorae Sinus and Margaritifer Sinus, as well as into bright area Chryse.
It is quite possible, however, that chaotic terrain is present elsewhere on the
Martian surface in areas not photographed by the Mariners. In fact, the
Mariner 4 frame 4N2 shows "irregular streaky patterns, " hinting at possible
chaotic terrain near 25 °N, in Amazonis, northeast of Trivium Charontis.
Relative Age
The distinctively fresh and sharp topographic features of the chaotic
terrain, especially at its contact with adjoining cratered terrain, the lack of
craters, and the low elevation, all suggest that it "formed at the expense of
cratered terrain" (Sharp et al. , 1971a) and is relatively young. At least, it is
younger than the large (>15 km) flat-bottomed craters, since these are
definitely erased. It is not clear whether the chaotic terrain is as young as the
small bowl-shaped craters seen on adjacent cratered terrain.
January 1, 1972 C. Michaux, JPL Sec. 3.5, page 41
Morphology and Processes
JPL 606-1
Postulated Areas of
Chaotic Terrain
O Crater
'''' Obscure Crater
1° at Equator = 59 km
20*N
10'N
Fig. Z5. Interpretive map of chaotic-terrain distribution constructed
from Mariner 1969 photos (Sharp at al. , 1971a).
Origin and Possible Processes
The following speculations on the mode of formation of chaotic terrain
have been made by Sharp et al. (1971a). If chaotic terrain is younger than
cratered terrain, then it must be the result of either some recent events or
some continuous processes erasing most craters. No analogs can be found on
either the Moon or the Earth. The irregular, hummocky terrain seen on the
Moon in the Montes Apennines, Mare Vaporum, and Cassini quadrangles are
ejecta sheets formed by asteroidal impact and/or volcanic explosions, and not
depressed areas. On Earth, large-scale collapse usually results from
volcanism, but the calderas and volcano-tectonic formations developed do not
have the geometric form or distribution pattern found in Martian chaotic
terrain. The unusual features of the latter are similar to those found in
terrestrial collapse, slumps, and slide areas; however, their scale is much
larger. It is difficult to explain such large-scale mass movements on Mars
in a relatively stable, recent period. It w^ould be necessary to have enough
large impacts and fair regional slopes of the bedrock to permit movement of
loose impact rubble induced from vibrations of these large impacts. This is
Sec. 3.5, page 42
C. Michaux, JPL
January 1, 1972
JPL 606-1 Morphology and Processes
unlikely. Other possible explanations, such as collapse from decay of
segregated bodies of ice (developed earlier in Mars' history), or even aeolian
deflation, are also difficult to support. Processes of internal origin, such as
volcanism accompanied by defluidization and/or tectonic deformation, might be
closer to the correct explanation. This, however, would presume that these
processes are an expression of a recently begun maturing stage in Mars'
geothermal evolution (Sharp et al. , 1971a). Recent models of the Mars interior
(Anderson, 197Z) do favor a melted core and partial differentiation (see
Section Z, Interior).
Featureless Terrain
Another type of uncratered terrain, called "featureless, " was discovered
in the Mariner 7 TV photographs 7N27-30 over the southern bright area Hellas.
Featureless terrain was defined by the CIT geologists (Sharp et al. , 1971a) as
terrain which appears to lack any kind of recognizable topographic features at
the available l/2-km resolution and over sizable expanses.* It has, therefore,
the appearance of smoothness in the pictures mentioned, but with improved
resolution this may not be the case. The irregular, diffuse variations in
shading seen may be due to gentle surface undulation under a low Sun or albedo
differences of the ground. The possibility of ground clouds or haze obscuring
the surface was considered but was discarded after careful scrutiny of both
FE and NE sets of pictures.
So far, featureless terrain has been found only in Hellas. The A-frames
covered about 65% (1.6 million km^) of Hellas, and revealed only a mesa-like
knob some 300 km east from its Hellespontus border (Thorman and Goles,
1971), three small flat-floored craters near this border (Sharp et al. , 1971a),
see 7N27, and possibly two ghost craters toward the center (see improved
versions of 7N29). Only two B-frames, confirming the lack of features (7N28
and 30), were taken over Hellas, but it was inferred that probably most of
Hellas is featureless.
The change in morphology from featureless Hellas to heavily cratered
Hellespontus proceeds rapidly through a transition zone some 150-350 km wide.
This zone displays, in addition to the flat-bottomed type crater found in
Hellespontus, a ''series of discontinuous overlapping scarps and narrow ridges,
individually 20-90 km long" and facing Hellas (Sharp et al. , 1971a). There
appears to be no small bowl-shaped craters. This zone is darker than
Hellespontus proper and forms a sharp albedo contrast with bright Hellas. The
contact line with Hellas is quite irregular and poorly defined topographically,
except near 45 °S (latitude of center of Hellas), where "there is an abrupt
change from cratered to featureless. . .marked by an irregular east-facing
escarpment" (Thorman and Goles, 1971). The abundance of large and small
flat-bottomed craters in the transition zone appears similar to that of
Hellespontus. The statement by Sharp et al. (1971a) that "nowhere do the
scarps or ridges crosscut flat-bottomed craters; rather, the craters appear to
interrupt or deform the ridges" was contradicted by a number of example
cases pointed out by Thorman and Goles (1971). (See their diagram. Fig. 26,
♦ An area at least 10,000 km^, to avoid confusion with the smooth intercrater
areas lying usually within cratered terrain.
January 1, 1972 C. Michaux, JPL Sec. 3.5, page 43
Morphology and Processes
JPL 606-1
7N27
(Fiducial Marks Indicated)
DIAGRAM OF THE HELLESPONTUS TO
HELLAS TRANSITION ZONE
(YAONIS FRETUM) AS VIEWED IN
7N27
• CIRCLED NUMBERS REFER TO CRATERS
• PLAIN LINES INDICATE RIDGES
• HACHURED LINES INDICATE SCARPS
• LEHERS REFER TO GRABENS AND
HORST5
Fig. 26. Diagram of the Hellespontus to Hellas transition zone (Yaonis
Fretum) as viewed in 7N27 (Thorman and Goles, 1971).
of an enlargement of the zone viewed in 7N27. ) In fact, these authors find
evidence of more cases against than for Sharp's statement. They also found
that not all scarps face Hellas; they indicate on the same diagram, the
presence of at least two "grabens" (indicated as A and B), one of which forms
an embayment of Hellas into the transition zone. The same authors also report
the presence of two north -trending "horsts" (indicated C and D) in the transition
zone.
The morphology of the transition zone, with its ridges parallel to scarps
facing Hellas, and the appearance of the contact line, reminiscent of creep,
suggest a general slope downward toward Hellas. This confirms other data
from Mariner 7 (see Topography subsection) indicating that Hellas is a large
depression some 2 km below the mean Martian level and some 5 km below
Hellespontus which is a high region. The circularity of Hellas certainly is
suggestive of a basin.
Origin and Age of the Hellas Basin
Hellas, considered as a large circular depression, appears to be
extremely old, dating perhaps as far back as the planet's accretionary phase.
Three possible modes of origin were considered by Sharp et al. (1971a): impact.
Sec, 3. 5, page 44
C. Michaux, JPL
January 1, 1972
JPL 606-1 Morphology and Processes
volcanic explosion, and subsidence. A combination of early accretionary
impact and later isostatic subsidence "over a dense mass within the crust or
near crustal interior (O'Leary et al. , 1969)" to quote Sharp, seems more
promising. Sharp et al. and Thorman and Goles are in agreement on this view.
The scarps and ridges of the transition zone appear to be related to,
and may provide some clues concerning the formation of Hellas. For the time
being, however, the question of their age relative to the craters of the zone is
an unsettled matter: Sharp et al. (1971a) consider the scarps and ridges to be
older than the old flat craters, arguing that they are distorted by them, while
Thorman and Goles (1971) feel they are generally younger. In the latter case,
the scarps and ridges would be accounted for by the late subsidence of Hellas,
rather than by the early impact itself, as suggested by Sharp et al.
Origin and Age of the Featureless Floor of Hellas
The featureless floor of Hellas is thought to be the product of either a
recent continuing surface process "capable of obliterating craters as rapidly
as they are formed" (Sharp et al. , 1971a) or a recent episodic event that swept
the surface clean. Continuing surface processes include aeolian burial (by
transport and deposition of dust), * unusually active creep, and basal surges
associated with impacts (a less likely mechanism). Episodic events include
volcanic extrusions (ash or tuff), fluidization of the rubble layer by volcanic
gases, activation of creep and/or other mechanisms by geothermal warming,
decay of frozen ground, and atmospheric cometary explosion. See Sharp et al. ,
1971a. Objections may be found for each of the processes listed, so that the
origin of the featureless terrain remains cloaked in mystery.
3. 5.4 SOUTH POLAR CAP
The Mariner 1969 photography provided the first close look at a Martian
polar cap: the fully developed, cloud-free, early springtime South Polar Cap.
The many FE pictures by both spacecraft, taken a half-million kilometers
away, first revealed that the very irregular edge was due in part to the presence
of many large craters, some exceeding 100 km in diameter. In addition, the
cap's interior displayed an irregular mottling with features up to 300 km, which
presumably represent real differences in the reflectivity of its surface. It was
also noted that the fuzziness and limb darkening toward the morning terminator
(on the west) may be due to a special photometric function of its surface (Leovy
et al. , 1971). Figure 27 shows the successive hour -interval views taken of the
Polar Cap by Mariner 6. Marginal irregularities can easily be followed through
their rotation with the planet.
Ten Mariner 7 NE frames, 7N10-20, taken at slant ranges of 5000-
6000 km and oblique angles of 40° -48° from vertical, revealed in spectacular
detail the structure of the frost cover from its edge, near 60 °S, to the South
Pole proper. It became clear that a rather well-cratered surface underlies
the white, gleaming frost almost everywhere except near the South Pole, at
least in the sector photographed (from 270° to 60°W). The frost cover
*An attractive process, in view of the denser air layer over Hellas.
January 1, 1972 C. Michaux, JPL Sec. 3.5, page 45
Morphology and Processes
JPL 606-1
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C, Michaux, JPL
January 1, 197Z
TPL 606-1 Morphology and Processes
appeared substantially thicker, although more variable, than previously
thought. Also noted were a number of strange -looking "snowforms or features
which presently defy explanation.
Mariner 7 Photography and Observations
Morphology
In studying the pictures of the South Polar Cap, the analyzing team of
CIT geologists (Sharp et al. , 1971b) has distinguished the following three con-
centric zones: (1) the edge or margin, a narrow area subdivided into three
subzones; extra, outer, and inner marginal subzones. (2) the interior, and
(3) the central region.
Summarizing the findings of the CIT team, these zones are reviewed
below.
1) Edge or Marginal Zone
a) The extra-marginal subzone (7N11, 13) is the bare ground
area (not frost-covered) just outside the frost patches of the
cap at the time of photography, but still within the polar
region limits or area annually covered by the extended polar
cap. Here craters appear to be similar to and as abundant as
those in other heavily cratered Martian terrain, and apparently
are unmodified by the cap's annual waxing and waning. There
is a lack of craters in the western third of 7N11, possibly due
to cloud obscuration.
b) The outer -marginal subzone (7N11-13, 15) is the immediate,
sparsely frost -covered edge of the cap itself, or in the words
of Sharp et al. (1971b), the zone "characterized by preservation
of disconnected frost patches. " This is to be viewed in the true-
intensity mosaic (without AGC). The Max-D mosaic (with AGC)
shows an unfortunate black band artifact here which is not to be
interpreted as the classical dark polar collar seen from Earth.
(When the AGC effect is removed, there is no suggestion of it. )
But this cannot be used to disprove the existence of the dark
polar collar, as the time of photography was too early in the
Martian spring for its appearance. In fact, the polar collar was
first seen in late September 1969, according to Capen (1970).
In this zone, the crater floors are frost-covered either
completely, appearing as solid white ellipses, or partially,
appearing as irregular white patches. Some of the southward-
facing inner walls, sheltered from the sun, display lingering
frost, appearing as white crescents, concave south. This frost
accentuates crater visibility and ground topography in general.
Also noted (7N11, right center) are "some scattered, irregular,
dense white patches" (Sharp et al. , 197 1b), which may be mis-
taken for clouds (Leovy et al. , 1971) but are^ore likely frost
accumulations due to topographic irregularities.
January 1, 1972 C. Michaux, JPL Sec. 3.5, page 47
Morphology and Processes jp^ £,o6 1
c) The inner marginal subzone (7N11, 13, 15) is a more deeply
and irregularly frost-covered zone of ragged appearance. The
frost cover is "essentially continuous except for islands of
frost-free, or nearly frost-free, ground" which register dark
to black in the Max-D pictures (AGC effect). Many craters are
expressed "as dark-line ellipses owing to obliquity of views or
as solid dark crescents concave north, " with some displaying a
dark central dot, possibly a peak (Sharp et al. , 1971b). The
visibility of abundant craters again owes much to frost
accentuation.
2) Polar Cap Interior
The interior (7N13-20) is the zone largest in area, and is shown as
essentially continuous frost cover with an abundance of craters and
various other topographic features, some of which were unexpected.
Most craters display bright rims and darker floors, although the
floors probably are not frost-free, because they do not appear black
like the bare ground at the cap margin. The surrounding (inter -
crater) surfaces are intermediate in brightness. Outlines of many
large craters appear much more irregular, even ragged, compared
to unfrosted ones, probably because of irregular frost accumula-
tions on their rims. A few medium-size craters display peripher-
ally "a radial pattern of short ridges and furrows" (7N19, upper
center). Some craters have central white dots suggestive of central
peaks. Craters over 15 km appear to be flat-floored, and some of
30-50 km size have distinct rims. Many smaller craters are bowl-
shaped, with rims. Abundance is comparable to other well-cratered
Martian terrain.
Unusual noncrater features are seen: irregular depressions with
angular outlines appearing often in frost-filled craters; they were
named "etch-pits. " They have lighter rims and darker floors like
the craters. Associated with them may be found "etch-furrows, "
or elongations of similar character.
Also found are "a series of short, parallel linear features aligned
in a WNW direction" and seemingly connecting tiny nodes: they were
called "beaded lineations. " Other lineations elsewhere have less
regular alignments. The cap in fact appears to be "a grooved,
fluted, and scoured surface" (Sharp et al. , 1971b). Smaller
subdued features of positive relief also are present.
Interesting larger, irregular features of positive relief were
noticed in frame 7N19. These are: one crater about 15 km across,
surrounded by features resembling (again quoting Sharp et al. ,
(1971b) "a pile of volcanic flows extruded from a central vent, "
then "an area of irregular hummocky terrain looking much like a
lunar ejecta sheet, " and "a belt featuring a number of short
irregular ridges" of northeasterly trend.
Finally, some large (20-120 km across) irregularly shaped white
patches and bands of high luminance were seen (7N17, 19) which,
if they are not clouds (Leovy et al. , 1971), remain unexplained. '
Sec. 3.5, page 48 C. Michaux, JPL January!, 1972
JPL 606-1
Morphology and Processes
3) Central Polar Region
The central region (7N17, 19, 20) is the region seen near the
South Pole and presumably encircling it without necessarily being
centered on it. Actually, only a sector of 165 degrees longitude
(between 235° and 40 °W) was photographed. Its boundary, with cap
interior, east of the prime meridian (0°W) is well-defined, between
80-75''S and displays an abrupt crater -scalloped arc somewhat like
the edge of a lunar mare. The boundary west of ° is not as well
defined. Craters are rare and barely visible in this central region,
with one notable exception: the forefoot of a large crater pair called
the "giant's footstep" crossing the boundary (see 7N19, 20). The
most unusual series of sinuous linear features, appearing to spread
out like waves from the pole, probably best characterizes the
central polar region. Called "quasi-linear markings" by Sharp
et al. , these enigmatic lineations are conspicuous between 230° and
40 °W, with lengths up to 300 km and widths of 10 km (see Fig. 28).
Their separation varies, and gnarls are present, as well as seg-
mentation. Faint indications of them swinging around the pole
were detected in highly processed pictures.
o
Marginal ^ o
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Fig. 28. Sketch of South Polar Cap: interior and central region,
morphological features appearing in 7N17 (Sharp et al. , 1971b).
January 1, 1972
C. Michaux, JPL
Sec. 3. 5, page 49
Morphology and Processes JPL 606-1
Processes
Marginal Zone . Processes responsible for the lingering of frost in
crater bottoms of the outer marginal zone are probably (1) protection from
solar radiation and winds, (2) differential CO2 vapor -pressure effect with
altitude, rather than temperature, as discussed by Sagan and Pollack (I966).
Similarly, in the inner marginal zone, the frost-free, or nearly frost-free,
rims of craters and other high-standing features (central peaks, for example)
are thus denuded for converse reasons.
Note: Gentle slopes (5° or less) for the individual topographic features
of the marginal zone were inferred from frost wastage (denudation) observa-
tions, assuming only solar radiation (insolation) is responsible (see Sharp
et al. , 1971b).
Polar Cap Interior . The process responsible for the unusual brightness
of highstanding features (crater rims, central peaks, or isolated knobs)
relative to the surrounding surface, especially crater bottoms and etch pits,
may be the textural differentiation of the CO2 frost cover experiencing a more
intense history of sublimation and condensation under the local meteorological
conditions (Sharp et al. , 1971). Also, adiabatic compression (heating) and
expansion (cooling) of the CO2 Martian air blowing over features of strong relief
may result in sublimation of low-lying C02-frost and condensation over high
areas (see Leighton et al. , 1971).
The etch features (pits and furrows), which are depressions carved into
the frost cover, with irregular, angular shapes, are quite unique to the cap
interior and seem to be the result of either (1) differential ablation caused by
local thin accumulations of dark debris (dust) blown in at one time--which would
explain their lesser reflectivity- -or (2) wind erosion 'which can scour and
pluck in unusual fashion" Again to quote Sharp et al. (197 1b): "the sharpness
and angularity of etch features suggest undermining of a frost layer udth a
slabby structure. "
The beaded lineations and other smaller feature's seem to have been
shaped by wind motion, that is by scour and drifting of the frost, "Fine detail
on larger crater rims resembles scalloping and fluting by wind scour. Much
of the pole cap surface looks swept because of faint lineations and elongation
of minor features, possibly produced through scour and drifting by wind"
(Sharp et al. , 1971b).
Central Polar Region . The rennarkable long, quasi -linear markings
definitely are ground features and not clouds, but their origin is mysterious.
From the photography, it is not clear whether they are troughs, ridges, or
scarps. They do resemble our longitudinal dunes, but are on a much larger
scale and more irregular. If they are ridges, they may also represent
accumulations of snow or ice, and/or rock material, suggesting large
Sec. 3.5, page 50 C. Michaux, JPL January 1, 1972
JPL 606-1 Morphology and Processes
moraines. If they are considered like scarps, then they conceivably could be
"the edges of platy layers of remnant ice" (Sharp et al. , 19? lb). These
authors speculate that great masses of perennial ice may have accumulated
under a more favorable phase of the effective precessional cycle of 50, 000
years,* and that this has occurred repeatedly for the south polar region. Per-
haps the unusual brightness of the central polar region, the paucity and faint-
ness of craters, and the quasi-linear markings are the result of a cumulative
effect of alternate large and small remnant summer central caps over many
niillions of years.
Thickness of Frost Cover
The frost cover is thin in the marginal zone, but becomes much thicker
in the cap interior with increasing latitude. However, the thickness is
obviously quite variable in the interior, locally because some large craters are
partly buried while elsewhere many small ones (1 km across) do show their
presence. Such variations probably are due to drifting and piling up of the
frost by wind. The thickness of these accumulations appears to be "on the
order of tens of meters" (Sharp et al. , 1971b). On the other hand, the etch
pits of large size (such as the "elephant's footprint" in 7N14) are negative
relief features suggesting thicknesses "of at least tens -of-meters. " They may,
however, not be due to differential ablation.
These thicknesses inferred from Mariner 7 photography of the South
Polar Cap appear to be roughly in agreement with the theoretically estimated
average thicknesses (about a meter) obtained by Leighton and Murray (1966),
and Cross (1971b).**
Permanence of Frost or Ice
Sharp et al. (1971b) speculate from quantitative estimates of the total
mass of solid CO2 possibly involved in the Martian polar caps through either
the annual or precessional cycles, that most likely there are sizable local
masses of old "dirty" perennial CO2 deposits which in summer escape detection
urider an insulating blanket of dust at least a few centimeters thick. The etch-
pits and quasi-linear markings may be indications of such deposits. The
Mariner 1971 mission might find evidence of their permanence.
*The~50, 000 year effective precessional cycle is the result of the combined
effects of the rate of advance of the perihelion (or line of nodes) and the rate
of regression of the north pole of rotation (or line of equinoxes or solstices) -
see Section 1 on Orbital and Physical Data.
**Cross (1971b) made approximate calculations of the seasonal variations in
size and average thickness of both polar caps. He found that the Spring
Southern Cap is not only larger but one -fifth thicker than the Northern one
(disregarding the perennial CO2 cap deposits).
January 1, 1972 C. Michaux, JPL Sec. 3.5, page 51
Morphology and Processes JPL 606-1
The total mass of solid CO2 estimated to be present now is either
1) Between ~10 and ~100 g cm~^ (average over planet). In this case,
a perennial CO2 frost cap normally exists at one pole, alternating
from S to N through the 50,000-year precessional cycle. Local
perennial deposits may occur at both poles as a result of wind drift-
ing and/or dust blanketing.
or:
2) Greater than ~100 g cm"^. In this case both poles harbor perennial
CO2 caps, and large masses of CO2 might exist under dust blankets.
3. 5. 5 DARK AND BRIGHT AREAS (MERIDIANI SINUS REGION):
BOUNDARIES AND MARKINGS
Near their closest approach in 1969, both Mariners photographically
sampled one of the prominent, stable dark areas of Mars: Meridiani Sinus
(well-known by its forked shape) and the appending arm of Sabaeus Sinus.
Mariner 6 picture quality and resolution were especially high (see 6N11, 13,
7N4, 6, and 8). Thus it was possible for the GIT geologists (Cutts et al. , 1971)
to make a special analysis of this dark area and its immediately surrounding
light areas. Their results are presented below.
Morphologically, this complex dark area, Meridiani Sinus -Sabaeus
Sinus, displayed around its periphery three types of boundaries:
1) Linear boundary, characterized by dark streaks within the light
area and nearly parallel to the overall boundary trend. This is
seen between Sabaeus Sinus and adjacent Moab light area.
2) Transverse boundary, characterized by projections perpendicular
to, or at steep angle to, the overall boundary trend with some dark
"outliers" in the light area. This is seen on the eastern boundary
of Meridiani Sinus next to light area Edom, which is a large crater.
3) Diffuse boundary, characterized by a gradual transition in albedo.
This is seen on the western boundary of Meridiani Sinus next to
light area Thymiamata, and continues along the southern boundary
next to light area Deucalionis Regio.
The character of the boundary may be controlled by the local topography.
Along crater walls (e.g., Edom) and scarps, the boundary is sharp with light
material on the lower side; the dark area then shows uniform albedo, while the
light area usually varies in albedo. In regions lacking relief the boundary is
diffuse, as near Deucalionis Regio. However, between Sabaeus Sinus and
Deucalionis Regio there are many short linear depressions nearly parallel to the
main diffuse boundary trend and which disappear gradually into the light area.
Cratered terrain extends over Meridiani Sinus and surrounding light
areas, with no apparent change in density and morphology of the large flat-
bottomed craters, except in the northern light areas, where craters are fewer
and more subdued.
Sec. 3.5, page 52 C. Michaux, JPL January 1, 1972
JPL 606-1 Morphology and Processes
Some craters in Meridian! Sinus display unusual albedo markings:
crescents of high albedo are on the northern part of crater floors and southern
slopes of crater walls and rims. Since this is contrary to the effects of solar
illumination, it was suggested by Cutts et al. (1971) that this may indicate
aeolian transport of "light material" in or out of the craters.
3.5.6 CANALS AND LINEAMENTS
Canals
The three Mariners have not confirmed the presence of the classical
system of canals. Instead, only a few linear formations and irregularly
elongated dark patches could be detected. The Mariner 4 TV track, which, it
was argued at the time, did not cross many prominent canals, apparently
revealed some dark bands and linear (tectonic) structures in frames 4N1-3.
These markings were "recognized" by Katterfeld and Hettervari (1969) in
particular as being portions of the large canals Erebus and Orcus. (Others,
less distinct, in 4N1 were interpreted as belonging to the canals Hades and
Boreas.) But this was not clear to most other geologists. The Mariner 6 and
7 TV NE tracks did cross prominent equatorial canals, but nothing very con-
vincing has been seen. (Note: Many so-called "lineaments" were detected,
but, as defined below, the terra lineaments refers to much smaller and narrower
markings which cannot be identified as canals. ) The FE pictures which
covered the planet at higher than telescopic resolutions showed only very few
dark elongated markings which could approximate in form the canals presumed
to exist at their locations. Such were the wide dark peninsula Coprates and the
elongated island Cerberus, both called "canals" on pre-Mariner maps.
Besides these two important examples (which correspond to wide canals seen
at the telescope), there were reports by Leighton et al. (1971b) of quasi -linear
alignments of dark-floored craters (at the position of the canal Gehon, for
example) and also mention by Cutts et al. (1971) of much more subtle boundary
lines between differently toned bright areas (north of Meridiani Sinus), which
appear to be contrast effects. It is possible that morphological features, other
dark patches and craters will later be found associated with other canals, but
this remains to be seen.
In summary, the present evidence supporting the existence of "canals"
is therefore unsubstantial. So far, the canals cannot be identified as a distinct
physiographic unit on Mars surface. (Note: It was argued that the season on
Mars at the time of the 1969 flybys was improper for the canals to show up at
greatest contrast since the canals supposedly follow the wave of darkening,
and are darkest in late spring-summer. This argument is expected to be
answered during the Mariner 9 surveillance of the planet's variable surface
features. )
Lineaments
It is now known from Mariner photographs that there exist on the surface
of Mars structural linear features interpreted as faults, grabens, horsts,
fractures, ridges, rilles, valleys, crater chains, etc. Some of these "linea-
ments" may be considered to be lines of weakness in the Martian crust and
January 1, 1972 C. Michaux, JPL Sec. 3.5, page 53
Morphology and Processes JPL 606-1
serve to indicate zones of past, and possibly present, tectonic activity (crustal
deformation processes). Many Martian lineaments have already been mapped
on the Mariner NE photographs. By plotting direction and frequency of the
lineaments in polar coordinates (as radius vector), one can produce an
azimuthal frequency diagram, more commonly called "rose diagram, " for a
certain region, or a group of regions, and eventually for the whole planet.
Figures 29 and 30 illustrate examples of regional rose diagrams: one was
issued by Katterfeld (1969) from mapping of lineaments on Mariner 4 frames
4N3-15; the other is by Binder (1971) from mapping on a number of Mariner 4, 6,
and 7 frames. These diagrams often exhibit clearly the prevailing trends or
directions of regional tectonism or crustal instability. For example, Katter-
feld's diagram shows strong prevalent NW-SE and NE-SW directions and weak
N-S and E-W directions. Generalization to the whole planet at present, how-
ever (as done by Wells, 1969),* appears premature until near-global NE photo-
graphic coverage by spacecraft has been achieved.
Oases
The oases, usually recorded on early maps as round or oval dark spots
at the intersection of canals in bright areas, were seen by Mariners 6 and 7 to
have irregular shapes (as observed in more recent telescopic observations) and
to resolve into fine structure, where "circular and annular markings may
correspond to large individual craters" (Cutts et al. , 1971). See for example
Oxia Palus in 7N5. Possibly a smaller oasis is formed by a single large crater.
Hartmann(197 lb) hypothesized that the darkness of the oases, considered as
impact craters, can be accounted for (in his model of aeolian crater obliteration)
by "dark rocky ejecta of bedrock deposited on top of a thin veneer of desert
material. "
*Wells (1969) inferred the existence over the Martian globe of these two
overlapping grid systems from Katterfeld's regional diagram and a more
global diagram based on the orientation of canals. (1969 References: see
Wells, 1971 b)
Sec. 3.5, page 54 C. Michaux, JPL January 1, 1972
JPL 606-1
Morphology and Processes
315*
Fig. 29. Rose diagram showing the azimuthal distribution of 86i
lineaments mapped from the Mariner 4 photographs 4N3-15
(Katterfeld (1969) modified by Wells, 1971 b).
January 1, 1972
C. Michaux, JPL
Sec. 3. 5, page 55
Morphology and Processes
JPL 606-1
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C. Michaux, JPL
January 1, 1972
JPL 606-1 Morphology and Processes
BIBLIOGRAPHY
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JPL 606-1 Morphology and Processes
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Hartmann, W. K. , 1966, Martian cratering: Icarus, v. 5, no. 6, p. 565-576,
November.
Herr.K. C, Horn, D. , McAfee, J. M. , and Pimentel, G. C. , 1970, Martian
topography from the Mariner 6 and 7 infrared spectra: Astron. J. , v. 75,
no. 8, p. 883-894, October.
Hord, C. W, , 1971, Mariner 6 and 7 ultraviolet spectrometer experiment:
photometry and topography of Mars: Boulder, Colorado, U. of Colo. ,
Report from the Laboratory for Atmospheric and Space Physics (78 p. ),
October. (To appear in Icarus, April 1972. )
Inge, J. L. , Capen, C. F. , Martin, L. J. , Faure, B. Q. , and Baum, W. A. , 1971,
A new map of Mars from planetary patrol photographs: Sky and Telescope,
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Katterfeld, G. N. , and Hedervari, P. , 1968, Ring-shaped and linear structures
on Mars: Astron. Zh. , v. 45, no. 5, p. 1091-1100, September-October;
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and 7 radio occultation results on the atmosphere of Mars: COSPAR
Space Research XI (Xlll-th Plenary Meeting, Leningrad, USSR);
Kondratyev, K. Y. , Rycroft, M. J. , and Sagan, C. , Editors : Akademie-
Verlag, Berlin, p. 165-175.
Leighton, R. B. , Horowitz, N. H. , Murray, B. C. , Sharp, R. P., Herriman, A. H. ,
Young, A. T. , Smith, B. A., Davies, M. E. , and Leovy, C. B. , 1971,
Mariner 6 and 7 television pictures: preliminary analysis, p. 259-294 in
Planetary atmospheres, (I. A. U. Symposium No. 40, held in Marfa, Texas,
October 26-31, 1969; Sagan, C. , Owen, T. C. , and Smith, H. J. , Editors :
D. Reidel Publishing Company, Dordrecht, Holland, (408 p.)-
Leighton, R. B. , and Murray, B. C. , 1966, Behavior of carbon dioxide and other
volatiles on Mars: Science, v. 153, no. 3732, p. 136-144, July 8.
Leighton, R. B. , Murray, B. C. , Sharp, R. P. , Allen, J. D. , and Sloan, R. K. ,
1967, Mariner IV pictures of Mars: Pasadena, Calif., Jet Propulsion
Laboratory, Tech. Rep. 32-884 (178 p. ), December 15. (Mariner Mars
1964 Project Report: Television Experiment, Part 1 -Investigators '
Report. )
Leighton, R. B. , Murray, B. C. , Sharp, R. P. , Allen, J. D. , and Sloan, R. K. ,
1965, Mariner IV photography of Mars: initial results: Science, v. 149,
no. 3684, p. 627-630, August 6.
Leonard!, "P. , 1966, Osservazioni geomorfologiche sui crateri lunari e
Marziani: Atti Accademia Nazionale dei Lincei, CI. sci. fis. , mat.e
natur., Rendiconti, v. 40, no. 5, p. 763-769, May,
January 1, 1972 C. Michaux, JPL Sec. 3.5, page 59
Morphology and Processes JPL 606-1
Leovy, C. , and Mintz, Y. , 1969, Numerical simulation of the atmospheric cir-
culation and climate of Mars: J. Atmos. Sci. , V, 26, no. 6, p. 1167-1190.
Leovy.C.B., Smith, B. A., Young, A. T. , and Leighton, R. B. , 1971, Mariner
Mars 1969: atmospheric results: J. Geophys. Res. , v. 76, no. 2, p. 297-312,
January 10.
Lincoln Laboratory (MIT), 1970, Radar studies of Mars: Final Report (79 p. ):
Lexington, Mass., Lincoln Laboratory, MIT, January 15.
McGill, G. E. , and Wise, D. U. , 1971, Regional variations in degradation and
density of Martian craters: Preprint of paper presented at the JPL
Symposium on the Geology of Mars: April 26-27. (Also submitted to
J. Geophys. Res. ) Revised Jan. 15, 1972.
Murray, B.C., Soderbloom, L. A, , Sharp, R. P. , and Cutts, J. A. , 1971, The
surface of Mars 1: cratered terrains: J. Geophys. Res. , v. 76, no. 2,
p. 313-330, January 10.
O'Leary, B. T. , and Rea, D. G. , 1967, Mars: influence of topography on forma-
tion of temporary bright patches : Science, V. 155, no. 3760, p. 317-319,
January 20.
O'Leary, B. T. , Campbell, M. J. , and Sagan, C. , 1969, Lunar and planetary
mass concentrations: Science, v. 165, no. 3894, p. 651-657, August 15.
Opik, E. J. , 1969, Mars: the changing picture: Irish Astron, J. , v. 9, no. 4,
p. 136-148, December.
Opik, E. J. , 1965, Mariner IV and craters on Mars : Irish Astron. J. , v. 7,
no, 2 and 3, p. 92-104, June -September.
Pang, K. , and Hord, C. W. , 1971, Mariner 7 ultraviolet spectrometer
experiment: photometric function and roughness of Mars' polar cap
surface: Icarus, v. 15, no. 3, p. 410-453, December.
Pike, R. J. , 1971, Genetic implications of the shapes of Martian and lunar
craters: Icarus, v. 15, no. 3, p. 384-395, December.
Pollack, J. B. , and Sagan, C. , 1970, Studies of the surface of Mars (very early
in the era of spacecraft reconnaissance): Radio Sci. , v. 5, no. 2,
p. 443-464, February.
Ryan, J. A. , 1969, Study of dust devils as related to the Martian yellow clouds:
McDonnell Douglas Corp., DAC -63098 (Volume 1 -Final Report, 117 p.),
January.
Ryan, J. A. , 1962, The case against thermal fracturing at the lunar surface:
J. Geophys. Res. , v. 67, no. 6, p. 2549-2558, June.
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no. 5208, p. 791-794, August 23.
Sec. 3.5, page 60 C. Michaux, JPL January 1, 1972
JPL 606-1 Morphology and Processes
Sagan, C. , and Pollack, J. B. , 1968, Elevation differences on Mars : J. Geophys.
Res. V.73, no. 4, p. 1373-1387, February 15.
Sagan, C. , and Pollack, J. B. , 1967, A windblown dust model of Martian surface
features and seasonal changes: Cambridge, Mass., Smithsonian
Astrophys. Observ. , Special Report no. Z55 (44 p. ), November 8.
Sagan, C, Veverka, J. , and Gierasch, P. , 1971, Observational consequences of
Martian wind regimes: Icarus, v. 15, no. 2, p. 253-278, October.
Sharp, R, P, , 1968, Surface processes modifying Martian craters : Icarus, v. 8,
no. 3, p. 472-480, May.
Sharp, R. P., Soderbloom, L. A. , Murray, B. C. , and Cutts, J. A. , 1971a, The
surface of Mars 2: uncratered terrains: J. Geophys. Res. , v. 76, no, 2,
p. 331-342, January 10.
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1971b, The surface of Mars 4: South polar cap: J. Geophys. Res. , v. 76,
no. 2, p. 357-368, January 10.
Thorman, C. H. , and Goles, G. G. , 1971, The relative age of the transition zone
between Hellas and the Martian cratered terrain: Preprint of a paper
presented at the JPL Symposium on the Geology of Mars, April 26-27.
Wade, F. A. , and DeWys, J. N. , 1968, Permafrost features on the Martian
surface: Icarus, v. 9, no. 1, p. 175-185, July.
Wells, R. A. , 1971 a, Analysis of large-scale Martian topography variations, I:
data preparation from Earth-based radar. Earth-based CO2 spectroscopy,
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Sciences Laboratory (63 p. ), June 23.
Wells, R. A., 1971 b, Martian surface harmonics and continental drift: Phys .
Earth Planet. Int. , v. 4, no. 3, p. 273-285, April.
Witting, J. , Narin, F. , and Stone, C. A. , 1965, Mars : age of its craters :
Science, v. 149, no. 3691, p. 1496-1498, September 24,
Woronow, A., and King, E, A. , Jr., 1972, Size frequency distribution of
Martian craters and relative age of light and dark areas: Science, v. 175,
no. 4023, p. 755-757, February 18.
Woronow, A. , and King, E. A. , Jr., 1971, A crater size frequency study of
Mariner 6 and 7 imagery: Preprint of a paper presented at the JPL
Symposium on the Geology of Mars, April 26-27.
January 1, 1972 C, Michaux, JPL Sec. 3.5, page 61
JPL 606-1
A mosaic of the Mariner TV r^i^f,.
model oi M.„. pTcZl c^ughMheS":?"; ' ""? t ^^°'°^'^^ o! a hand-pain.ed
observed it from a position about 8 oon „•', °\^^S'. o! the planet as Mariner ly^"""'"
range of 10, 500 miles. As the s J«rr,7^ ^'°" ""^ P'"^ »' Mars' orbit at a slant
C,P'°f'"-^ toward tt lo:,ra"\"„1'irp!? r/t. "" ''r"' ''= '^levVsron scan"'
lo^t^plft^e'-dtte^r^"^'- ^'■' ■"' •^'« ^'«-"(^0,"rld=zT:e'^: t'arn'in^'"
July 1, 1968
'"^X'TFfiAMg ;
photograph of a hand -painted
planet as Mariner IV
ine of Mars' orbit at a slant
planet, the television scan
-ed across the terminator
and 22) were taken in
Picture 7
Picture 7 is wholly within the south-
eastern part of the bright region
Zephyria, near Mare Sirenum.
Location of picture center:
Latitude .230
Longitude .174°
Dimensions of area:
East-west 290 km (180 mi)
North -south ... .290 km (180 mi)
Spacecraft distances:
^^i^^d^ 13,192 km
5, , (8, 179 mi)
Mant range 13,582 km
(8,421 mi)
Time and lighting:
^°^^^i"^^ about noon
Phase angle 59.6° to 60 5°
Zemth angle ... .Sun is 29° north
of the zenith
Filter ....
green
°^^^^aP Lower right corner
overlaps picture 8
Time of exposure . . . 00:25:45 GMT,
July 15, 1965
Picture
Picture 8 contains pa
and the dark Mare Si:
border between these
trends across the mic
picture .
Location of picture ce
Latitude
Longitude
Dimensions of area:
East-west 2S
North -south .... 27
Spacecraft distances-
Altitude
Slant range
Time and lighting:
Local time
Phase angle f
Zenith angle . . . . Sui
Filter
Overlap Uppe
overl
Time of exposure . . . 0(
Ir OLD
ti«-'T FRAME 1
■''-*^"i;!!t'.'~''' '-»"'**'■
-■■.J- '
Picture 8
Picture 8 contains part of Zephyria
and the dark Mare Sirenum. The
border between these two features
trends across the middle of the
picture .
Location of picture center:
Latitude -17°
Longitude 173°
Dimensions of area:
East-west 290 km (180 mi)
North-south .... 275 km (170 mi)
Spacecraft distances:
Altitude 13, 056 km
(8,095 mi)
Slant range 13, 373 km
(8,291 mi)
Time and lighting:
Local time about noon
Phase angle 59. 5° to 60. 5°
Zenith angle .... Sun is 32° north
of the zenith
Filter orange
Overlap Upper left corner
overlaps picture 7
Time of exposure . . . 00.26:33 GMT,
July 15, 1965
Picture 9
Picture 9 is largely in the dark area
Mare Sirenum, bordering on the
light area Atlantis in the southwest
corner of the picture.
Location of picture center:
Latitude -24°
Longitude 169°
Dimensions of area:
East-west 275 km (170 mi)
North-south . . . .260 km (161 mi)
Spacecraft distances:
Altitude 12, 790 km
(7,930 mi)
Slant range 13, 004 km
(8,062 mi)
Tinne and lighting:
Local time about noon
Phase angle 59.4° to 60.4°
Zenith angle .... Sun is 39° from
the zenith
Filter orange
Overlap Lower right corner
overlaps picture 10
Time of exposure . . . 00:28:09 GMT,
July 15, 1965
Picture 10 is
Mare Sirenuj
light area At
on Mare Sire
corner of the
Location of p
Latitude. ,
Longitude
Dimensions c
East-west
North -sou
Spacecraft di
Altitude . .
Slant rang
Time and ligl
Local tim<
Phase ang!
Zenith ang
Filter . ,
Overlap.
Time of expo:
» O.'.DOUT FRAME 3
J. de Wys, JPL
■..r
«^-:^^ii^^'-
e darkarea
I on the
southwest
,-24'
169°
m (170 mi)
m (161 mi)
12,790 km
(7,930 mi)
13,004 km
(8,062 mi)
ibout noon
" to 60.4°
39° from
the zenith
. . orange
,'ht corner
picture 10
:09 GMT,
15, 1965
Picture 10
Picture 10 is partly in the darkarea
Mare Sirenum, and largely in the
light area Atlantis, which borders
on Mare Sirenum in the northeast
corner of the picture.
Location of picture center:
Latitude -27°
Longitude 2570
Dimensions of area:
East-west 275 km (170 mi)
North-south .... 260 km (I6I mi)
Spacecraft distances:
Altitude 12,660 km
(7,849 mi)
Slant range 12, 843 km
(7,963 mi)
Time and lighting:
Local time about noon
Phase angle 59- 4° to 60. 3°
Zenith angle Sun is 42° from
the zenith
filter g,^^^
°^^^lap Upper left corner
overlaps picture 9
Time of exposure . . . 00:28:57 GMT,
July 15, 1965
Picture 11
Picture 11 probably is principally
within the dark Mare Sirenum, but
may be in or near the lighter area
Atlantis between Mare Sirenum and
Mare Cimmerium.
Location of picture center:
Latitude -33°
Longitude 162°
Dimensions of area:
East-west 275 km (170 mi)
North-south . . . . 240 km (149 mi)
Spacecraft distances:
^titude 12,407 km
(7.692 mi)
Slant range 12, 564 km
(7,790 mi)
Time and lighting:
Local time about 12:40
Phase angle 59 . 3° to 60 . 3°
Zenith angle .... Sun is 49° from
the zenith
Filter g,,^^
Overlap Lower right corner
overlaps picture 12
Time of exposure . . . 00:30:33 GMT,
July 15, 1965
fOLDOUT FRAAIg 4
Morphology and Processes
r«^ ■■•••■
Picture 12
;ip6Llly Picture 12 probably lies across the
m, but poorly defined border between the
r area dark Mare Sirenum and the lighter
lum and area Atlantis.
-33'
,162'
(170 mi)
(149 mi)
,407 km
692 mi)
, 564 km
790 mi)
at 12:40
to 60.3°
9° from
e zenith
. . green
corner
:ture 12
3 GMT,
5, 1965
Location of picture center:
Latitude -36°
Longitude 160°
Dimensions of area:
East-west 275 km (170 mi)
North-south .... 240 km (149 mi)
Spacecrait distances:
Altitude 12, 284 km
(7,616 mi)
Slant range 12, 446 km
(7,717 mi)
Time and lighting:
Local time about 12:40
Phase angle 59. 3° to 60. 2°
Zenith angle .... Sun is 52° from
the zenith
Filter orange
Overlap Upper left corner
overlaps picture 11
Time of exposure . . . 00:31:21 GMT,
July 15, 1965
Fig. 24. Mariner IV pictures 7 to
12, and the locations of Martian
regions photographed in pictures
1 to 19. North is at the top.
'3!-D0UT FRAME 5
Sec. 3.5, page 29
JPL 606-1 Mariner 1969 Photographic Atlas of Mars
3. 6 MARINER 1969 PHOTOGRAPHIC ATLAS OF MARS
INTRODUCTION
The successful flybys of Mariners 6 and 7 past Mars, on July 31 and
August 5, 1969, respectively, have provided more than 200 complete television
pictures of the planet which have significantly increased our knowledge of its
surface and atmosphere. The volume and quality of the data returned greatly
surpassed that obtained from the initial flyby of Mariner 4 in 1965. Not only
was the whole planet photographed during several rotations, but many high-
resolution closeup photos were obtained over its Southern Hemisphere which
had just entered its spring season. It is the purpose of this section to present
the pictorial highlights of the 1969 Television Imaging Experiment results. The
selected pictures presented here, with relevant photographic and areographic
information, are often referred to and discussed in other sections of this docu-
ment. The complete set of pictures will be published in book form as a
NASA SP-263, under the title "The Mariner 6 and 7 Pictures of Mars" (Collins,
1971), and may more properly be called a "photographic atlas. "
TELEVISION EXPERIMENT DESIGN
The television observations period for each spacecraft during the Mars
encounter was divided according to two periods:* (1) a two to three day far
encounter (FE) period terminating several hours before actual closest approach
of the planet, and (2) a half-hour near encounter (NE) period centered on
closest approach.
The scan platform of both spacecraft carried two cameras of different
focal lengths, providing two resolutions differing by a factor of ten. The short
focal length or wide-angle camera (camera A), for the lower resolution, was
equipped with a rotating set of four filters sequenced red, green, blue, and
green, on a filter wheel. The long focal length or narrow-angle camera
(camera B), for the higher resolution, had only one blue cutoff (or yellow haze)
filter. The characteristics of the camera optics are contained in Table 1 and
the transmission curves of the filters utilized are shown in Figs. 1, 2, and 3.
A full description of the optics and filters may be found in Montgomery and
Adams (1970), and Danielson and Montgomery (1971).
The FE pictures were taken with camera B to obtain highest resolution
of the disk. The NE pictures were taken alternatively with cameras A and B
so that a (high resolution) B franne was nested within the overlapping region of
two successive (low resolution) A frames. The camera A frames were taken
successively through the alternate filter sequence indicated above, starting with
blue for Frame No. 1. All B frames were taken through the unique yellow haze
filter.
*With Mariner 7, an additional, originally unscheduled, period of observations
took place between the FE and NE periods. During this late far encounter
period (LFE) an additional 88 valuable pictures were secured through the two
r ^ rvi «^ "r a G _
cameras,
June 1, 1971 C. Michaux, JPL Sec. 3.6, page 1
Mariner 1969 Photographic Atlas of Mars
JPL 606-1
Table 1. Characteristics of the Mariner 6 and 7 Camera Optics.
Type
Focal Length
Field of View
Aperture
Shutter
Filters
(effective
wavelength)
Exposure Time
(fast, slow)
Camera A
(wide -angle)
Zeiss Planar f/Z
multielement lens
stopped down to f/5.6
5 2 mm
11 " X 14"
10 mm
4 -position rotary-
Red (5730 A)
Green (5260 A)
Blue (4690 A)
90 and 180 msec
Camera B
(narrow- angle)
Modified Schmidt-
Cassegrain telescope f/2. 4
508 mm
1.1° X 1.4°
200 mm
2-blade, right-left
Blue cut-off (5600 A)
(Schott GC-14 glass)
6 and 12 msec
100
<
Damklson and Mo.ntcomery
400 440
480 520 560
WAVELENGTH, nm
600
640
680
Fig. 1. Spectral transmission, Mariner 6 camera A filters. Spectral
response characteristics of each filter incorporated in the
A camera flown on Mariner 6.
Sec. 3. 6, page 2
C. Michaux, JPI.
June 1, 1971
JPL 606-1
Mariner 1969 Photographic Atlas of Mars
<
a:
00
1 I ■ T ■ - r
1
1
80
BLUE —
7 r\L
/—RED
<
60
40
/ '/ \\ k^
/ '/ i /
V /
,^ GREEN 1
/—GREEN II
20
J
J,
400 440 480 520 560
WAVELENGTH, nm
600
640 680
Fig. 2, Spectral transmission. Mariner 7 camera A filters. Spectral
response characteristics of each filter incorporated in
the A camera flown on Mariner 7.
Danielson and Montgomery (1971)
300
400
500
WAVELENGTH, nm
600
700
Fig. 3. Transmission characteristics as a function of wavelength
for the camera B filter.
June 1, 1971
Danielson and Montgomery (1971
Sec. 3. 6, page 3
Mariner 1969 Photographic Atlas of Mars
JPL 606-1
The designation of each frame is identified by a code which cites
spacecraft number (6 or 7), sequence letter (F or N) and frame sequence
number. For example, 6N3 is decoded as: 6 = Mariner 6, N = near
encounter, 3 = frame number 3. All odd-numbered, near encounter frames
were taken by camera A (low resolution) and all even-numbered frames by
camera B (high resolution).
CAMERA SYSTEM
The Television Camera System was composed of a vidicon image tube,
complex analog-to-digital conversion assembly, and a magnetic tape recorder
system for recording both analog and digital video data prior to transmission
to Earth, This system was designed to improve the visual contrast of the
Martian surface, * as well as maximize the transmission of both TV data and
other scientific data obtained by the Mariners, The onboard video processing
is described in detail by Danielson and Montgomery (1971). Essentially, each
camera scanned a raster of 945 picture elements (or pixels) per line and
704 lines per frame, so that a complete picture consisted of 665,280 pixels.
Each pixel's brightness was digitized (encoded) into an 8-bit binary word,
providing 256 ( = 2^) possible levels of brightness or shades of gray to be
recorded. The onboard encoding was accomplished, however, in three
versions:
1) A composite analog video (CAV) picture, consisting of a fully sam-
pled raster, with each pixel encoded to 6 significant bits. Before
encoding, the analog signal from the vidicon was modified so as to
emphasize low-contrast detail; this was done through two amplifiers;
an "automatic gain control" (AGC) which enhanced the visibility of
the small-scale detail, and a "cuber'' which enhanced the local con-
trast (by a factor of ~3), After encoding, the data was recorded on
tape.
2) An "every seventh" spacing digital video (DV) picture, consisting of
sampling every seventh pixel along each scan line, with encoding to
8 significant bits, **
3) An "every-twenty-eighth" (ETE) spacing digital picture, from
sampling every 28th pixel along each scan line, with encoding to
6 significant bits. ■'■'■''
The encoding scheme used in versions 2) and 3) permitted transmission
of both the nonvideo science data (forming a "data bar" in the middle of the DV
pictures) and information on the nonlinearity of the AGC and cubing amplifiers.
IMAGE PROCESSING
The image processing performed at the JPL Image Processing
Laboratory was designed to give two final products or versions of the images:
(1) maximum discriminability version ("Max-D"), and (2) photometric version.
*Mariner 4 revealed very low contrast factor.
=;=':=With the two most significant bits truncated.
Sec. 3. 6, page 4
C. Michaux, JPL
February 15, 197 2
JPL 606-1 Mariner 1969 Photographic Atlas of Mars
The Max-D version strongly emphasizes the fine-scale brightness
variations and, therefore, local abrupt contrasts of small surface features,
while sacrificing true contrasts between large surface areas. This effect is
due to the AGC-cuber combination, which acts as a photometric high-pass
filter, eliminating the slow, low-frequency variations of brightness (see Dunne
et al. , 1971). Only the NE pictures were processed in the Max-D version.
The photometric version gives a large-scale true brightness (or
'photometric') representation of the surface albedo differences. This version
is useful for study of the classical dark areas/bright areas dichotomy or
albedo patterns of large areas of the planet. All the FE pictures were
processed in this version only.
To obtain the Max-D versions, the following processing was performed
on the CAV picture streamis:
1) Rectification or correction to remove the various distortions due to
the vidicon system. These distortions include system noises of
several kinds: periodic noise producing a basket-weave pattern,
long-line or streak noise, and isolated (black and white) spike
noise.
2) Enhancement cjf frequency response to compensate for modulation
transfer function fall -off.
3) Ge(jmetric distortion correction to remove optical and especially
electronic imaging aberrations. (Note: The geometric distortion
here is not the projection distortion due to oblique viewing. )
To obtain photomietric versions, the processing necessitated first per-
forming a reconstruction task of a "full analog video" (FAV) picture (that
is, a picture as could be obtained at the vidicon before the AGC and cubcr).
This task consisted of two parts:
1) Restoration of the 'most significant bits' truncated before
transmission, in the DV and ETE picture data streams.
2) Reconstruction proper of a full picture, from recombination of the
three picture data streams (CAV, DV, and ETE).
This included a correction to remove the nonlinear photometric effects
of the AGC-cuber combination.
Then, the same rectification processing, as done for the Max-D
version, was performed; that is, correction of distortions due to the vidicon
system (and in addition the removal of the 'residual image'), and correction
of geometric distortions.
Finally, a photometric decalibration was performed to remove shading
and light-sensitivity variations due to the vidicon (nonlinear and spatially
varying sensitivity properties). A discussion of the JPL imiage processing
techniques applied to the Mariner 6 and 7 pictures has been given by
Rindfleisch et al. , (1971).
February 15, 1972 C. Michaux, JPL Sec. 3.6, page 5
Mariner 1969 Photographic Atlas of Mars JPL 606-1
3. 6. 1 FAR ENCOUNTER
Introduction
All the far encounter (FE) pictures were taken with the narrow-angle
camera (B) to obtain highest possible resolution. Mariner 6 took a total of ^
49 FE pictures in two series, extending over almost two days and two rotations
of Mars. The Mariner 7 camera was activated a day earlier and produced
a total of 91 FE pictures in three series, covering three complete rotations of
the planet.
Only the most significant pictures are presented here, starting with the
initial photo at 1.7 million km and concluding with the closest FE photo taken at
130,000 km from Mars. The terrestrial date and the longitude of the central
meridian on the Martian disk are indicated below each photograph. "When the
full disk is seen, a fixed size is used for the image of the disk. Later, as the
spacecraft's TV field of view is reduced to a portion of the disk, we present the
picture as it was received. The first series of pictures obtained by Mariner 7
does not show much more than was previously known from Earth telescopic
observations, because of the low resolution. As the spacecraft approached
closer to the surface of Mars, as shown in the second series of Mariner 7 or
first series of Mariner 6 photos, the improving resolution revealed many of
the classical surface feature aspects not previously observable. As the
spacecraft approached 300,000 km from the planet a wealth of detail began to
appear.
In presenting this brief photographic reconnaissance of Mars by the two
spacecraft of 1969, it will be seen that many previous opinions and speculations
concerning the true nature of the surface of Mars, and associated atmospheric
phenomena, have had to be revised. Only the highlights of the FE sequence are
shown and annotated in the photos contained herein.
The two spacecraft approached Mars slightly south of its equatorial
plane: 6 degrees offset for Mariner 6, and 4 degrees for Mariner 7. Con-
sequently, the cameras viewed the entire Martian globe except for a small
zone centered on the north pole (see Fig. 4). The Martian surface features,
therefore, 'drifted' on flattened elliptical paths across the disk.
The classical westward increasing system of longitudes is used
(0 to 360 degrees) in denoting position on Mars. The present zero meridian
passes about 3 degrees west of Fastigium Aryn (the old origin of longitudes,
located inside the fork of Meridiani Sinus).
The trajectory of the spacecraft was such that Mars was photographed
in the FE sequence with a noticeable solar phase angle of about 22 degrees.
The dark crescent beyond the morning terminator appeared on the left side of
all the pictures, which are oriented with the Martian north pole at the top.
Since the direction of rotation for Mars is the same as for the Earth, the
surface features will appear to drift from west to east, that is from left to
right, as the Mariners approached the planet.
The planetary photography data pertaining to each of the 140 FE
pictures taken are listed in Table 2, Far Encounter Photoreference Data.
Sec. 3.6, page 6 C. Michaux, JPL June 1, 1971
JPL 606-1
Mariner 1969 Photographic Atlas of Mars
MORNING
TERMINATOR
NORTH POLE OF ROTATION
EAST
SOLAR RAYS
SOUTH POIAR CAP
CENTRAL MERIDIAN (CM.) OF THE DISK
(Note in fh!» figure CM. = 320")
Fig. 4. The globe of Mars.
June 1, 1971
C. Michaux, JPL
Sec. 3. 6, page 7
Mariner 1969 Photographic Atlas of Mars
JPI. 606-1
Table 2. Far encounter photoreference data.
Picture
Number
Shutter
Range,
km
Phase
Angle,
deg
Center (
f Disk
Resolution
km /Pixel
Date
l-ongitude,
deg
Latitude,
deg
Mariner 6
6F1
July 29
Bottom
1,244,221
21
100°W
-6
34.0
6F2
Bottoni
1,228,388
21
109
-6
33.4
6F3
Bottom
1,212,556
21
117
-6
32.9
6F4
Bottom
1,196,724
21
126
-6
32.5
6F5
Top
1,180,281
21
136
-6
32.1
6F6
Top
1,164,448
21
145
-6
31.7
6F7
Top
1,148,615
21
153
-6
31.3
6F8
Top
1,132,782
21
162
-6
30.8
6F9
Bottom
1,116,342
21
172
-6
30.4
6F10
Bottom
1,100,507
21
181
-6
30.0
6F11
Bottom
1,083,790
21
190
-6
29.5
6F12
Top
1,068,232
21
199
-6
29.0
6F13
Top
1,052,399
21
208
-6
28.6
6F14
Top
1,036,568
21
217
-6
28.2
6F15
Top
1,020,733
21
226
-6
27.8
6F16
Bottom
1,004,291
21
235
-6
27.3
6F17
Bottom
988,457
21
244
-6
26.8
6F18
Bottom
972,624
21
253
-6
26.4
6F19
Top
956,181
21
262
-6
26.0
6F20
Top
940,347
21
27 1
-6
25.6
6F21
Top
924,514
21
280
-6
25.2
6F22
Top
908,679
2!
289
-6
24.7
6F23
Bottom
892,236
21
298
-6
24.3
6F24
Bottom
876,402
21
307
-6
23.8
6F25
Bottom
860,568
21
316
-6
23.4
6F26
Bottom
844,733
21
325
-6
23.0
6F27
Top
828,289
21
334
-6
22.6
6F28
Top
812,455
21
34 3
-6
22.1
6F29
Top
796,619
21
352
-6
21.7
6F30
Bottom
780,175
21
I
-6
21.2
6F31
July 29
Bottom
764,339
21
10
-6
20.8
6F32
July 30
Bottom
748,454
21
:9
-6
20.4
6F33
Bottom
732,667
21
28
-6
19.9
6F34
Bottorri
568,196
21
120
-6
15.4
6F35
Bottom
540,171
21
136
-6
14.7
6F36
Top
512,753
21
151
-6
13.9
6F37
Top
484,723
21
167
-6
13.1
6F38 ,
Bottom
457,301
21
182
-6
12.4
6F39
Bottom
429,267
21
198
-6
U. 6
6F40
Bottom
401,840
21
213
-6
10.9
6F41
Bottom
401,231
21
214
-6
10.9
6F42
Bottom
376, 849
21
227
-6
10.2
6F43
Bottom
352,463
21
241
-6
9.5
6F44
Bottom
328,075
21
255
-6
8.8
6F45
Top
304,293
21
268
-6
8.2
6F46
Top
279,898
21
282
-6
7.5
6F47
Top
255,498
21
295
-7
6.9
6F48
Top
231,092
21
!09
-7
6.2
6F49
July 30
Top
206,680
21
B22
-7
5.5
Mariner 7
7F-1
Aug 2
Top
1,844,034
22
50
-4
50.3
7F2
*
Bottom
1,720,371
22
121
-4
46.9
7F3
Top
1,697,072
22
134
-4
46.3
7F4
Bottom
1,685,722
22
141
-4
46.0
7F5
Top
1,674,371
22
147
-4
45.7
7F6
Bottom
1,663,021
22
154
-4
45.3
7F7
Top
1,651,670
22
160
-4
45.0
7F8
Bottom
1,640,320
22
167
-4
44.7
7F9
Bottom
1,628,371
22
173
-4
44.4
7F10
Top
1,617,023
22
180
-4
44.1
7F11
Bottom
1,605,67 1
22
186
-4
43.8
7F12
Top
1,594,321
22
193
-4
43.5
Sec . 3.6, pa
ge
C. Michaux, JPL
June 1, 1971
JPL 606-1
Mariner 1969 Photographic Atlas of Mars
Table 2. Far encounter photoreference data (continued)
Pic turf
Nunibf !
Date
Kan^n,
km
Phase
Angle,
deg
Center of Disk
l-^ongitude,
deg
Latitude,
deg
Resolution
km /Pixel
7F1^
7F14
7F15
7F16
7Fn
7F18
7F19
7FZ0
7F21
7F22
7F23
7F24
7F25
7F26
7F27
7F28
7F29
7F30
7F31
7F32
7F33
7F35
7F36
7F37
7F38
7F39
7F40
7F41
7F42
7F43
7F44
7F45
7F46
7F47
7F48
7F49
7F50
7F51
7F52
7F53
7 F54
7F5 5
7F56
7F57
7F58
7F5 9
7F60
7F61
7F62
7F63
7F64
7F65
7F66
7F67
7F69
7F70
7F7)
7F7 2
7F73
71 /4
7F75
', r,-6
rF77
7F78
7F79
Aug 2
Aug 2
Aug 3
i
Mariner 7 (continued)
Aug 3
Aug 4
Bottoin
1,582,971
Top
1,571,621
Top
1,559,673
Bo 1 torn
1,548,323
rop
1,536,973
Bottom
1,525,623
T,,p
1,514,167
Rotloni
1,502,923
Bottom
1.490,977
Top
1,479,626
Bottom
1,468,279
Top
1,453,533
Bollom
1,445,577
Top
1,433,630
Top
1,422,281
Bottom
1,410,931
Top
1,399,581
Bottom
1,388,233
Top
1,376,882
Top
1,364,937
Bottom
1,353,586
Bottom
1, 199,474
Top
1,184,540
Top
1,169,009
Bottom
1,154,076
Bottom
1,138,545
Top
1,123,613
Top
1,108,081
Top
1,092,551
Bottom
1,077,617
Bottom
1,062,086
Top
1,047,152
Top
1,031,621
Bottom
1,016,687
Bottom
1,001,156
Bottom
985,625
Top
970,691
Top
955,160
Bottom
940,225
Bottom
924,693
Top
909,759
Top
894,227
Top
878,695
Bottom
863,760
Bottom
848,227
Top
833,292
Top
817,760
Bottom
803,248
Bottom
787,291
Top
772,355
Top
756,821
Top
741,287
Bottom
726,351
Bottom
710,816
Bottom
535,132
Bottom
514,811
Top
495,087
Top
474,764
Bottom
455,037
Top
435,309
Top
414,982
Bottom
395,251
Bottom
374,920
Top
355,185
Bottom
335,449
22
22
22
22
22
22
22
22
22
22
22
22
22
22
22
22
22
22
22
22
22
i. i
23
23
23
23
23
23
23
23
23
23
23
23
23
23
23
23
23
23
23
23
23
23
23
23
23
23
23
23
23
23
23
23
23
23
23
23
23
23
2 3
23
I99-W
-4
206
-4
213
-4
219
-4
226
-4
233
-4
239
-4
246
-4
252
-4
259
-4
265
-4
27 2
-4
27 8
-4
285
-4
292
-4
298
-4
305
-4
311
-4
318
-4
325
-4
331
-4
60
-4
68
-4
77
-4
86
-4
95
-4
103
-4
112
-4
121
-4
13
-4
139
-4
147
-4
156
-4
165
-4
174
-4
183
-4
191
-4
200
-4
209
-4
218
-4
226
-4
235
-4
244
-4
253
-4
26J
-4
27
-4
279
-4
287
-4
296
-4
305
-4
314
-4
323
-4
331
-4
340
-4
81
-4
93
-4
104
-4
116
-4
127
-4
138
-4
150
-4
161
-4
173
-4
184
-4
195
-5
43.2
42.9
42.5
42.2
41.9
41.6
41.3
41.0
40.7
40.4
40.1
39.8
39.4
39.1
38.8
38.5
38.2
37.9
37.6
37.3
37.0
32.8
32.4
31.9
31.5
31.1
30.7
30.3
29.9
29.5
29.0
28.6
28.2
27.8
27.3
26.9
26.5
26.1
25.7
25.3
24.8
24.4
24.0
23.6
23.2
22.8
22.4
22.0
21.5
21.1
20.7
20.3
19.9
19.4
14.6
14.0
13.5
12.9
12.4
11.9
11.3
10.8
10.2
9.7
9.1
June 1, 1971
C. Michaux, JPL
Sec. 3. 6, page 9
Mariner 1969 Photographic Atlas of Mars
JPL 606-1
Table 2, Far encounter photoreference data (continued).
Center o
f Disk
l-icture
Nunibi;r
Date
Shutter
Range,
km
Phase
Angle,
deg
Resolution
km/Pixel
Longitude,
deg
Latitude,
deg
Mariner
7 (continued)
7J-80
Aug 4
Bottom
315,112
23
207 " W
-5
8.6
7F81
Top
295,370
23
218
-5
8.0
7F82
Bottom
275,625
22
230
-5
7.5
7F83
Bottom
255,279
22
241
-5
6.9
7F84
Top
235,527
22
252
-5
6.4
7F85
Top
215,173
22
264
-5
5.8
7F86
bottom
195,412
22
27 5
-5
5.3
7F87
Top
175,645
22
286
-5
4.8
7 FSB
Top
155.273
22
298
-6
4.2
7F89
fop
150,478
22
300
-6
4.0
7F90
1
Bottom
145,084
22
304
-6
3.9
7F91
1
Bottom
140,288
22
306
-6
3.7
7F92
1
Top
134,892
22
309
-6
3.6
7F93
Aug 4
Top
130,095
22
312
-6
3.5
Sec. 3. 6, page 10
C. Michaux, JPL
June 1, 1971
JPL 606-1 Mariner 1969 Photographic Atlas of Mars
Mariner 7 — First Series
The first series of pictures, obtained by Mariner 7 (7F1-33), were
taken from distances of 1.7 million km to 1.35 million km.. These lower
resolution pictures show the face of Mars from Solis Lacus (7F2) across the
great Martian desert (Amazonis, Tharsis, Arcadia, Tempe), to the island
Trivium Charontis (7F16), to Syrtis Major (7F28) and Sabaeus Sinus (7F33).
Not much detail is revealed, but the same areas reappear later with much
greater resolution. The edge of the South Polar Cap was noted to display a
large irregularity.
June 1, 1971 C. Michaux, JPL Sec. 3.6, page 11
Mariner 1969 Photographic Atlas of Mars
JPL 606-1
Sunrise or Morning
Terminator
MARE SIRENUM
AUGUST 2, 1969
TEMPE
MARE ACIDALIUM
Late Afternoon Limb
SOUS UCUS
AONIUS SINUS
SOUTH POLAR CAP
CM. 121°
1 .72 MKm
Fig. 5. Far Encounter Frame 7F2.
NORTH POl^R MISTS
NODUS
LAOCOONTIS
(NODUS LACUS)
MARE TYRRHENUM
AUGUST 2, 1969
AUSONIA
TRIVIUMCHARONTIS
AMAZONIS
MARE CIMMERIUM
ELECTRIS
ERIDANIA
SOUTH POLAR CAP
Sec. 3. 6, page IZ
CM.: 219°
Fig. 6. Far Encounter Frame 7F16.
C. Michaux, JPL
1.55 Mkm
June 1, 1971
JPL 606-1
Mariner 1969 Photographic Atlas of Mars
NORTH POLAR HOOD
DELTOTON SINUS
ISIDIS REGIO
SYRTIS MAJOR
lAPYGIA
MARE TYRRHENUM
SABAEUS SINUS
HELLESPONTUS
HELLAS
SOUTH POLAR CAP
AUGUST 2, 1969
CM.: 298°
1.41 Mkm
Fig. 7. Far Encounter Frame 7F28.
ISMENIUS LACUS
SABAEUS SINUS
MERIDIANI SINUS
SYRTIS MAJOR
lAPYGIA
HELLAS
DEUCALIONIS REGIO
NOACHIS
MARE SERPENTIS
SOUTH POLAR CAP
AUGUST 2, 1969
June 1, 1971
CM.: 331°
Fig. 8. Far Encounter Frame 7F33.
C, Michaux, JPL
1.35 Mkm
Sec. 3. 6, page 1 3
Mariner 1969 Photographic Atlas of Mars JPL 606-1
Mariner 6 — First Series and Mariner 7 — Second Series
The first series of pictures taken by Mariner 6 (6F1-33) and the second
series obtained by Mariner 7 (7F35-67), both viewing the same area, start
again at Solis Lacus and go all the way to Sabaeus Sinus, even to its termina-
tion into Meridiani Sinus. These pictures were taken from distances of
l.Z million km to approximately 700,000 km, at which distance the resolution
attained becom^es distinctly superior to that obtainable from Earth.
Sec. 3.6, page 14 C. Michaux, JPL June 1, 1971
JPL 606-1
Mariner 1969 Photographic Atlas of Mars
MARE ACIDALIUM
ARCADIA
AMAZONIS
PHOENICIS LACUS
MARE SIRENUM
AONIUS SINUS
CAP EDGE IRREGULARITIES
AUGUST 3, 1969
JUVENTAE PONS
AURORAE SINUS
COPRATES
TIT HON I US LACUS
(GROUP OF 'OASES')
SOUS LACUS
SOUTH POLAR CAP
CM. 103°
1.12 Mkm
The complex area of Solis Lacus and Tharsis is seen with much
fine detail in 7F40. The features known as Coprates Canal, Juventae
Fons, Tithonius Lacus are clearly recognizable. The Phoenicis Lacus
"oasis" to the west is less distinct. Notice in the south the large dark
areas Mare Sirenum and Aonius Sinus hugging the brilliant South Polar
Cap, and the two cap edge irregularities.
June 1, 1971
Fig. 9. Far Encounter Frame 7F40.
C. Michaux, JPL Sec. 3.6, page 15
Mariner 1969 Photographic Atlas of Mars
JPL 606-1
NIX OLYMPIC A
MARE SIRENUM
AUGUST 3, 1969
CLOUD-LIKE
MOnLINGS
PHOENICI5 LACU5
TITHONIUS LACUS
(GROUP OF 'OASES')
SOLIS LACUS
SOUTH POIARCAP
AONIUS SINUS
CAP EDGE IRREGULARITIES
CM. 139°
1.06 MKm
Northwest of the Solis Lacus area, the vast Amazonis -Thar sis
desert appears in 7F44, where many paler mottlings are conspicuous.
Notice a circular marking, which will be identified later as Nix
Olympica.
Fig. 10. Far Encounter Frame 7F44.
Sec. 3.6, page l6 C. Michaux, JPL
June 1, 1971
JPL 606-1
Mariner 1969 Photographic Atlas of Mars
PROPONTIS
NIX OLYMPICA
TRIVIUM CHARONTIS
CERBERUS
MARE CIMMERIUM
AUGUST 3, 1969
ERIDANIA
ELECTRI5
CLOUD-LIKE
MOTTLINGS
SOUTH POLAR CAP
CM. 174°
MARE SIRENUM
1.00 MKiT
Trivium Charontis and Propontis, two large northern islands,
appear in 7F48, while the large southern dark areas Mare Sirenum and
Mare Cinamerium stretch across the disk uninterruptedly, revealing an
irregular border facing Amazonis. In the afternoon limb at the equator,
a bright cloud-like mottling appears.
June 1, 1971
Fig. 11. Far Encounter Frame 7F4{
C. Michaux, JPL
Sec. 3. 6, page 17
Mariner 1969 Photographic Atlas of Mars
JPL 606-1
NORTH POUR HOOD
PROPONTIS
CERBERUS
TRIVIUM CHARONTIS
CLOUD-LIKE
MOTTLINGS
AETHIOPIS
MARE CIMMERIUM
DISCRETE
BRIGHTENING
AUSONIA
ERIDANIA
SOUTH POLAR CAP
AUGUST 3, 1969
CM. 209°
a 94 MKm
In 7F5Z, the north polar hood is quite conspicuous. It is also clear
in the Mariner 6 pictures. Elysium appears bright above Trivium
Charontis. Propontis stands out. Mare Cimmeriuni reveals a large
protrusion, unmistakably the Cyclopia (Angustus Sinus) canal. The sharp
outline of the South Cap has one large irregularity.
Fig. 12. Far Encounter Frame 7F52.
Sec. 3. 6, page li
C. Michaux, JPL
June 1, 1971
JPL 606-1
Mariner 1969 Photographic Atlas of Mars
SYRT IS MAJOR
MARE TYRRHENUM
NODUS LAOCOONTIS
(NODUS LACUS)
CERBERUS
AUSONIA
HELLAS
SOUTH POLAR CAP
AUGUSTS, 1969
CM. 270°
0.83 MKm
Syrtis Major reappears in the morning sun at the left in 7F59. In
the center of the disk is the prominent dark Mare Tyrrhenum, with a
large elongated island, often referred to as Nodus Laocoontis, off its
irregular border. In the south, the bright patches of Ausonia merge into
Hellas, which appears circular in other photographs.
June 1, 1971
Fig. 13. Far Encounter Frame 7F59.
C. Michaux, JPL Sec. 3.6, page 19
Mariner 1969 Photographic Atlas of Mars
JPL 606-1
NORTH POLAR HOOD
MERIDIANI SINUS
SYRTIS MAJOR
lAPYGIA
MARE SERPENTIS
SABAEUS SINUS
DEUCALIONIS REGIO
PANDORAE FRETUM
HELLAS
YAONIS FRETUM
HELLESPONTUS
NOACHIS
SOUTH POLAR CAP
AUGUST 4, 1969
CM. 340°
0.71 MKrr
7F67, taken at 700,000 km, shows (much better than Earth-based
photographs) the "tree" formed by Hellespontus -Mare Serpentis as the
trunk of the two boughs with lapygia-Syrtis Major on the right and Sabaeus
Sinus -Meridian! Sinus on the left. Below lapygia, Hellas is quite con-
spicuous, and below Meridiani Sinus, Deucalionis Regio and Pandorae
Fretum are quite distinct. Above Sabaeus Sinus the three deserts Aeria,
Arabia, and Moab are quite bright. None of the classical canals are
visible.
Fig. 14. Far Encounter Franie 7F67.
Sec . 3.6, page 20
C. Michaux, JPL
June 1, 1971
JPL 606-1
Mariner 1969 Photographic Atlas of Mars
NORTH POLAR HOOD
ISMENIUS LACUS
MARE ACIDALIUM
NILIACUS LACUS
OXIA PALUS
MARGARITIFER SINUS
AURORAE SINUS
MARE ERYTHRAEUM
JULY 30, 1969
MERIDIANI SINUS
SABAEUS SINUS
PANDORAE FRETUM
HELLAS
CM. 19«
0.75 MKm
6F32 shows for the first time the face of Mars which was omitted by
the gap between the first and second series of Mariner 7. In the center,
bright Aram-Thymiamata separates the prominent Meridiani Sinus from
Margaritifer Sinus along the equator. Below Pandorae Fretum the desert
Noachis appears diffuse, and below Margaritifer Sinus stretches Mare
Erythraeum. The triangular Oxia Palus oasis is quite visible off the tip
of Margaritifer Sinus, and in the north the large Mare Acidalium shows a
sharp southern outline. Notice in all of the FE pictures the limb dark-
ening of the South Polar Cap.
June 1, 1971
Fig. 15. Far Encounter Frame 6F32.
C. Michaux, JPL Sec. 3.6, page 21
Mariner 1969 Photographic Atlas of Mars JPL 606-1
Mariner 7 — Third Series and Mariner 6 —Second Series
The final series of Mariner 7 Far Encounter photos, consisting of
Z5 pictures (7F69-93), was taken after an interruption of 6-1/2 hours and
only the first eight are of the full disk. The distances were from 535,000 km
(7F69) to 130,000 km (7F93). Mariner 6 took 16 pictures (6F34-49) in its
second and last series. The distances were from 570,000 km (6F34) to
ZOO, 000 km (6F49), with only the first six pictures being of the full disk.
All these pictures are taken at distinctly better resolution than can be
obtained from Earth.
Sec. 3.6, page 22 C. Michaux, JPL June 1, 1971
JPL 606-1
Mariner 1969 Photographic Atlas of Mars
TRACTUS ALBUS
THARSIS
PHOENICIS LACUS
NILOKERAS
MARE ACIDALIUM
LUNAE PALUS
TITHONIUS LACUS
(6 'OASES')
JUVENTAE PONS
AURORAE SINUS
COPRATES
NECTAR
MARE ERYTHRAEUM
AONIUS SINUS
SOUS LACUS
THAUMASIA
SOUTH POLAR CAP
AUGUST 4, 1969
CM. 93°
asi MKm
7F70 shows the Solis Lacus face of Mars for the third time, but
now with a great wealth of detail heretofore unseen. The six or more
components of the Tithonius Lacus multiple oasis are entirely resolved.
Phoenicis Lacus is well resolved as the right 'twin' arrowhead. Juventae
Fons remains a single large black dot, while Coprates (Agathodaemon)
is seen to be a continuous unusually large elongated peninsula rather than
a canal. Solis Lacus is linked to Mare Erythraeum by a twisted channel
swollen into a large knot recognized as Nectar. It is well known that both
Solis Lacus and Nectar have varied much in size and shape in the past.
In the Tharsis area a number of circular bright streaks are visible around
dark centers. In the north the southwestern double promontory, Nilokeras,
of Mare Acidalium is very conspicuous. In the desert Chryse just off
Aurorae Sinus a number of smaller black dots and/or streaks are notice-
able. A cleft is visible in the edge of the South Polar Cap.
June 1, 1971
Fig. 16. Far Encounter Frame 7F70.
C. Michaux, JPL Sec. 3. 6, page 23
Mariner 1969 Photographic Atlas of Mars
JPL 606-1
NORTH POLAR HOOD
NIXOLYMPICA
PHOENICIS LACUS
MARE SIRENUM
JULY 30, 1969
DISCRETE BRIGHT FEATURES
(CLOUDS OR FROSTS)
MARE ACIDALIUM
TITHONIUS LACUS
JUVENTAE FONS
AURORAE SINUS
COPRATES
SOUS LACUS
AONIUS SINUS
CM. 120°
0.57 MKn
6F34 has much less contrast but shows the north polar hood as
rather extensive and irregular, with at least two brighter cloud-like
tongues, which will rotate nearly unchanged with the planet.
Fig. 17. Far Encounter Frame 6F34.
Sec. 3.6, page 24 C. Michaux, JPL
June 1, 197 1
JPL 606-1
Mariner 1969 Photographic Atlas of Mars
NORTH POLAR HOOD
DISCRETE BRIGKT FEATURES
F AND G (LEOVY)
NIXOLYMPICA
AMAZONIS
MARE SIRENUM
ELEMENTS 0,P,Q (LEOVY)
OF THE FORMING W-CLOUD.
THEY ARE ASSOCIATED WITH
LARGE MULTI-RINGED
STRUCTURES.
TITHONIUS LACUS
ELEMENT R (LEOVY)
OF FORMING
W-CLOUO
PHOENICIS LACUS
SOLIS LACUS
OGYGIS REGIO
AONIUS SINUS
SOUTH POLAR CAP
AUGUST 4, 1969
CM. 138°
0.44 MKm
7F74 reveals annular Nix Olympica as a white ring with a white
central dot. Its appearance is suggestive of a very large crater with
frost-covered rims and a central peak. Other large multiringed struc-
tures are present to the southeast of Nix Olympica. With these are
associated diffuse cloud masses which are enlarging and brightening
during the afternoon to form a W- pattern. Cloud elements O, P, Q, R,
labelled by Leovy are indicated here.
as
June 1, 1971
Fig. 18. Far Encounter Frame 7F74.
C. Michaux, JPL Sec. 3.6, page 25
Mariner 1969 Photographic Atlas of Mars
JPL 606-1
NORTH POLAR HOOD
DISCRETE BRIGHT FEATURE
(CLOUD OR FROST)
SMALL LOBES
MARE SIRENUM
AUGUST 4, 1969
NIXOLYMPICA
PHOENICIS LACUS
SMALL ROUND
WHITE SPOT
SOUTH POLAR CAP
CM. 161°
0.40 MKm
7F76 shows again the region of Nix Olympica. The Amazonis-
Tharsis desert, southeast and west of it, is strangely streaked, ringed
and dotted into a complex of subtle pale marksings. The shores of
Mare Sirenum are sharp and irregular with protrusions in the form of
lobes and even a finger. It is interesting to mention that the Mariner 4
track was located across Mare Sirenum as indicated in Fig. 20.
Fig. 19. Far Encounter Frame 7F76.
Sec. 3.6, page Z6 C. Michaux, JPL
June 1, 1971
JPL 606-1
Mariner 1969 Photographic Atlas of Mars
'ii-i' ,?.
EEB
ana
Fig. ZO. Far Encounter Frame 7F76 modified to show the positions of
Mariner 4 pictures 4N7 through 4N14.
June 1, 1971
C. Michaux, JPL
Sec. 3. 6, page 27
Mariner 1969 Photographic Atlas of Mars
JPL 606-1
NORTH POLAR HOOD
TRIVIUM CHARONTIS
CERBERUS
MARE CIMMERIUM
AUGUST 4, 1969
NIXOLYMPICA
W-CLOUD*
(PART OF)
CM. 184°
MARE SIRENUM
OR 'THARSIS'
CLOUD (G.deV.)
a36 MKm
7F78 shows on the eastern limb the seasonal "W cloud" over
Tharsis which is a well-known afternoon recurrence at that time of the
year. Whether the phenomenon is actually a cloud or ground frost is
not yet known. The development of this bright feature is seen from
7F77 to 7F79, or 6F38 and 6F39.
Fig. Zl. Far PZncounter Framie 7F78.
Sec. 3. 6, page 28
C. Michaux, JPL
June 1, 1971
JPL 606-1
Mariner 1969 Photographic Atlas of Mars
PROPONTIS
EUXINU5
LACUS
TRIVIUM CHARONTIS
CERBERUS
MARE CIMMERIUM
MARE CHROMIUM
JULY 30,1969
NIXOLYMPICA
W-CLOUD
('THARSIS' CLOUD)
ELEMENTS O, P,R
MARE SIRENUM
LARGE SPIKE
OF CAP EDGE
CM. 182°
0.46 MKn
6F38 also shows the development of a portion of the W-cloud, but
on another late afternoon five days earlier. The North Polar haze is
quite extensive as in preceding frames. Below it the two elongated large
dark forms entering at the western end of Amazonis are Propontis and
Trivium Charontis -Cerberus. The South Polar Cap exhibits a fairly-
prominent spike on its morning side; this makes a sizable encroachment
into Mare Chronium bordering it, below Mare Cimmerium.
June 1, 1971
Fig. 22. Far Encounter Frame 6F38.
C. Michaux, JPL Sec. 3.6, page 29
Mariner 1969 Photographic Atlas of Mars
JPL 606-1
NORTH POLAR HOOD
BRIGHT SPECKS
PROPONTIS I
EUXINUS LACUS
TRIVIUM
CHARONTIS
CERBERUS
MARE CIMMERIUM
MARE CHROMIUM
W-CLOUD
(ELEMENT O)
SOUTH POLAR CAP
AUGUST 4, 1969
CM. 207°
0.32 MKm
7F80, taken at 315,000 km, shows the interesting fish-like structure
of the Trivium Charontis -Cerberus complex, a famous variable feature.
Bright Elysium above appears strinkingly pentagonal and here speckled
with two brighter spots also seen in the Mariner 6 pictures. Prominent
Mare Cimmerium displays a multitude of finger -like extensions into
Aeolis .
Fig. Z3. Far Encounter Frame 7F80.
Sec. 3. 6, page 30
C. Michaux, JPL
June 1, 1971
JPL 606-1
Mariner 1969 Photographic Atlas of Mars
NODUS
LAOCOONTIS
(NODUS LACUS)
MARE
TYRRHENUM
TRIVIUM
CHARONTIS
*' ERIDANIA ■■,.<^.f:i
ANGUSTUS
SINUS
(G.de V.)
LARGE
CRATER
MARE
CIMMERIUM
AUGUST 4, 1969
CM. 241°
0.26 MKm
7F8 3 shows more of the ragged shores of Mare Cimmerium; the
large finger Cyclopia (Angustus Sinus) seems to be a chain of craters.
The interior of the large Mare Cimmerium appears riddled with similar
circular features. The brightness of Elysium is decreasing with the
afternoon hours. The surrounding deserts Aeolis, Aethiopis, and
Aetheria are in view. On the boundary of Aethiopis the large island
known as Nodus Laocoontis has diffuse form.
June 1, 1971
Fig. 24. Far Encounter Frame 7F83.
C. Michaux, JPL Sec. 3.6, page 31
Mariner 1969 Photographic Atlas of Mars
JPL 606-1
SYRTIS MAJO
MARE
SERPENTIS
ISIDIS REGIO
MOERIS LACUS
JULY 30, 1969
HELLAS
CM. 282°
LIBYA
MARE
TYRRHENUM
0.28 MKn
6F46 exhibits the entire Syrtis Major-Hellas region. Note the
circularity and northern brightness of Hellas, probably due to haze. The
feature known as Zea Lacus is not visible in any of the FE pictures.
Fig. Z5. Far Encounter Frame 6F46.
Sec. 3.6, page 32 C. Michaux, JPL
June 1, 1971
JPL 606-1
Mariner 1969 Photographic Atlas of Mars
AERIA
HE LIAS
7F91, taken at 140,000 km from Mars, gives a closeup view of the
area linking Syrtis Major to northern Hellas and Sabaeus Sinus. In
l3^Pygia» ^ very large crater is seen, while numerous smaller craters
are unmistakably recognizable in the Deltoton Sinus promontory and
across lapygia. Craters are even visible in the neighboring light areas
Aeria on the left and Libya on the right, but they are more subdued in
contrast.
June 1, 1971
Fig. 26. Far Encounter Frame 7F91.
C. Michaux, JPL Sec. 3.6, page 33
Mariner 1969 Photographic Atlas of Mars
JPL 606-1
MARINER 7 PHOTOGRAPHED
MARS' MOON PHOBOS
OVER THE DESERT AERIA
ON 4 AUG. 1969 23:35:34
INSET IS X4 MAGNIFICATION
1 PIXEL
^SCAN
Plane
Equatorial
f PHOBOS 1
17.5
Icm
[^r,. J
-IP Lao
« — 22.5 km— «.
Limb profTle of Phobos in frame 7F91 . The size of a
single television picture element (pixel) and vidicon
scanning direction is indicated. The orbital plane of
Phobos (not shown) lies nearly in the Martian equa-
torial plane. Geometric corrections and coherent
noise removal hove changed the shape from that seen
in the raw version (Smith, 1970).
An enlarged portion of 7F91 (16 times) showed the image of Phobos
as an oval-shaped, dark object (see upper portion of picture). Only four
other FE frames have positively shown Phobos.
Fig. 27. Far Encounter Frame 7F91 (magnified portion showing Phobos).
Sec. 3.6, page 34 C. Michaux, JPL June 1, 1971
JPI. 606-1
Mariner 1969 Photographic Atlas of Mars
DELTOTON
SINUS
SABAEUS
SINUS
jAUGUST4, 1969
SYRTIS
MAJOR
LIBYA
MARE
TYRRHENUM
0.13 MKm
7F93, taken at 130,000 km, is the last and closest FE frame and is
centered on lapygia, showing clearly the large crater over 300 km across,
mentioned earlier, as well as other craters surrounding it.
June 1, 1971
Fig. 28. Far Encounter Frame 7F93.
C. Michaux, JPL Sec. 3.6, page 35
Mariner 1969 Photographic Atlas of Mars
JPL 606-1
DELTOTON SINUS
SABAEUS SINUS
MARE SERPENTIS
YAONIS FRETUM
HELLESPONTUS
JULY 30, 1969
j^ lAPYGIA
HELLAS
CM. 322°
SYRTIS
MAJOR
MARE
TYRRHENUM
THIN HAZE
LAYER
MARE
HADRIACUM
0.21 MKm
6F49, taken at ZOO, 000 km from Mars, is the last of the Mariner 6
FE pictures. It shows in large scale Sabaeus Sinus, Mare Serpentis, and
lapygia with its enormous crater, and portions of the deserts Noachis and
Hellas where thin haze layers seem to be present. Note the irregular
northern edge of Sabaeus Sinus. Mare Tyrrhenum and Mare Hadriacum on
the right limb are rather diffuse, perhaps because of the same veil of
haze.
Fig. Z9. Far Encounter Frame 6F49.
Sec. 3.6, page 36 C. Michaux, JPE
June 1, 1971
JPL 606-1 Mariner 1969 Photographic Atlas of Mars
3.6.2 NEAR ENCOUNTER
The near-encounter (NE) period, lasting about 20 minutes, was centered
on closest approach of the planet and provided 59 pictures, some as close as
3500 km, with a resolution attaining 0.3 km. The photographic coverage by the
two spacecraft encompassed over 10% of the Martian surface, dependent on slant
range (and perhaps 20% if limbs are included), mostly of the Southern Hemis-
phere. Mariner 6 took 26 NE pictures covering an equatorial swath (from
to 25 °S latitude), and five days later Mariner 7 took 33 NE pictures covering
both a mid-latitude and polar swath in the Southern Hemisphere. Two resolu-
tions were used alternatively: the low resolution pictures (3 km at best) were
taken with the wide-angle camera "A, " and the high resolution pictures (0.3 km
at best) with the narrow-angle camera "B. " The approximate picture locations
are illustrated on painted globes of Mars, one for each spacecraft, as seen in
Figs. 30 and 31.
The following pages present a selection of the most significant NE pic-
tures, each illustrating a particular aspect of the Martian surface or atmos-
phere. The classification is somewhat arbitrary. The caption comments,
however, place the emphasis on the main aspect as indicated by the six group-
ings which are as follows:
1) Cratered Terrain
2) Chaotic Terrain
3) Featureless Terrain
4) Atmospheric Haze (Limb Pictures)
5) Dark and Light Areas (Meridiani Sinus Complex)
6) South Polar Cap
Each of the above groupings is introduced by at least one mosaic pro-
viding a panoramic spacecraft view of the regions encompassed. Odd-numbered
pictures were taken by camera A and even-numbered pictures by camera B.
The Near Encounter Photoreference Data, Table 3, provides specific data for
each of the NE pictures taken.
June 1, 1971 C. Michaux, JPL Sec. 3.6, page 37
Mariner 1969 Photographic Atlas uf Mars
JPL 606-1
Fig. 30. Mariner b picture locations on a painted globe of Mars.
Sec . 3.6, page 3i
C. Michaux, JPL
June 1, 1971
JPL 606-1
Mariner 1969 Photographic Atlas of Mars
Fig. 3 1. Mariner 7 picture locations on a painted globe of Mars,
June 1, 1971
C. Michaux, JPL
Sec. 3. 6, page 39
Mariner I969 Photographic Atlas of Mars
JPL 606-1
Table 3. Near encounter photoreference data.
At Center of Picture
Picture
Shutter
or Pi Iter
Slant
Range,
Phase
Angle,
Sola r
Elevation
Angle,
deg
Viewing
Angle,
Longi -
tude.
Lati-
tude,
Picture Dimen.
on Surface, km
iTVAlillVt- I
km
deg
deg
deg W
deg
Horiz
Vert
Ma riner
6
6N1
Blue
6621
44
88
45
53
-6
1292
6N2
Top
7389
52
71
70
67
4
551
143
6N3
Green
6598
44
83
57
56
-2
1324
6N4
Bottom
6159
52
85
51
50
-5
238
118
6N5
Red
5699
45
79
42
43
-8
2283
1127
6N6
Top
5355
52
73
37
37
-10
162
103
6N7
Green
5030
51
66
30
31
-13
1492
990
6N8
Bottom
4778
52
61
25
26
-14
128
92
6N9
Blue
4930
51
49
41
14
1511
1228
6N10
Top
4727
52
44
39
10
-1
123
111
6N11
Green
4541
51
39
37
4
-3
1177
1124
6N12
Bottom
4428
52
35
38
-3
109
107
6N13
Red
4331
52
29
40
354
-4
1090
1130
6N14
Top
4903
80
71
62
37
-13
253
94
6N15
Green
4404
73
60
50
25
-16
2202
865
6N16
Bottom
4105
80
52
42
18
-17
123
79
6N17
Blue
3865
80
44
34
10
-18
127
756
6N:8
Top
3746
80
38
31
4
-16
105
73
6NI9
Green
3617
80
31
25
357
-17
983
718
6N20
Bottom
3546
80
26
21
351
-16
91
69
6N21
Red
3501
80
20
17
345
-16
889
697
6N22
Top
3498
80
14
15
340
-15
86
69
6N23
Green
3522
80
8
15
334
-14
880
706
6N24
Bottom
3584
80
3
18
328
-13
88
71
Mariner
7
7N1
Blue
8725
37
79
44
9
-5
1767
7N2
Top
9766
43
58
75
13
20
201
7N3
Green
9118
44
67
67
59
12
2144
7N4
Bottom
8492
44
75
59
5
4
401
162
7N5
Red
7995
44
82
52
3
-2
1598
7N6
Top
7552
44
88
46
359
-7
269
144
7N7
Green
7136
44
85
41
356
-12
3143
1406
7N8
Bottom
6774
44
80
36
353
-17
203
130
7N9
Blue
6443
44
74
31
350
-21
1989
127 2
7N10
Top
6693
35
39
48
33
-54
168
190
7N11
Green
6381
35
38
45
27
-57
1650
2025
7N12
Bottom
6095
35
36
43
21
-61
149
159
7N13
Red
5886
35
34
43
17
-65
1512
1690
7N14
Top
5662
35
31
42
9
-68
142
143
7N15
Green
5495
35
28
43
1
-71
1484
1513
7N16
Bottom
5318
35
25
43
349
-74
14 1
132
7N17
Blue
5195
35
21
45
334
-77
1549
1394
7N18
Top
5069
35
18
46
314
-77
148
124
7N19
Green
5013
35
13
49
291
-78
1837
1356
7N20
Bottom
4971
35
8
53
269
-75
17 2
124
7N21
Red
5337
80
76
66
354
-21
1066
7N22
Top
4818
80
66
56
346
-28
210
92
7N23
Green
44 31
80
57
47
339
-34
1951
868
7N24
Bottom
4154
80
49
40
331
-38
132
80
7N25
Blue
3938
80
43
34
3 24
-42
1218
773
7N26
Top
3778
80
36
29
316
-44
104
73
7N27
Green
3656
80
31
24
308
-46
989
723
7N28
Bottom
3679
80
24
28
I'^S
-41
94
77
7N29
Red
3633
80
19
26
291
-41
925
779
7N30
Top
3636
80
13
27
284
-39
89
78
7N31
Green
3660
80
8
28
277
-38
915
811
Sec. 3. 6, page 40
C. Michaux, JPL
June 1, 1971
JPL 606-1 Mariner 1969 Photographic Atlas of Mars
CRATERED TERRAIN
Mosaic of Seven Camera A Frames 6N9 Through 6N23
This series sweeps eastward (left to right) across the heavily cratered
Martian equatorial regions from dark areas Margaritifer Sinus and oasis Oxia
Palus through dark area complex Meridiani Sinus -Sabaeus Sinus with adjacent
bright area Deucalionis Regio.
June 1, 1971 C. Michaux, JPL Sec. 3.6, page 41
Mariner 1969 Photographic Atlas of Mars
JPL 606-1
rd
T3
GJ
a
nj
O
Z
O
Xi
be
O
CO
bC
Sec . 3.6, page 4Z
C. Michaux, JPL
June 1, 1971
JPL 606-1
Mariner 1969 Photographic Atlas of Mars
6N11 — Western side of prominent dark area Meridiani Sinus.
Diffuse boundary or transition into bright area Thymiamata at left.
Cratered terrain extends east-west across both areas.
June 1, 1971
Fig. 33. Near Encounter Frame 6N11.
C. Michaux, JPL Sec. 3.6, page 43
Mariner 1969 Photographic Atlas of Mars
JPL 606-1
6N13 — Meridian! Sinus: its forked northern portion is viewed here
in the Max-D version, processed to reveal details of the local topography.
Note the asymmetric shading of craters and the irregular northeastern
boundary of Meridiani Sinus. Part of 300-km crater Edom appears at
far right.
Fig. 34(a). Near Encounter Frame 6N13 (Max-D version).
Sec. 3.6, page 44 C. Michaux, JPL June 1, 1971
JPL 606-1
Mariner 1969 Photographic Atlas of Mars
.if,M:T.
6N13 — Meridiani Sinus: In this photometric version, the large
scale albedo variations are visible.
Fig. 34(b). Near Encounter Frame 6N13 (photometric version).
June 1, 1971 C. Michaux, JPL Sec. 3.6, page 45
Mariner 1969 Photographic Atlas of Mars
JPL 606-1
6N17 — Cratered terrain consisting mostly of flat-bottomed craters.
The location of this frame contains parts of Margaritifer Sinus at left and
the bright area Thymiamata at right. The large scale albedo variations
have been supressed in this Max-D version.
Fig. 35. Near Encounter Frame 6N17.
Sec. 3.6, page 46 C. Michaux, JPL
June 1, 1971
JPL 606-1
Mariner 1969 Photographic Atlas of Mars
■^' Li;
6N18 — Closeup view of a typical large flat-bottomed crater about
30 km across. The large crater shows some terracing, but no central
peak. Location is south of Meridiani Sinus near longitude zero. Note
the large grooved structure at bottom left.
June 1, 1971
Fig. 36. Near Encounter Frame 6N18.
C. Michaux, JPL Sec. 3.6, page 47
Mariner 1969 Photographic Atlas of Mars
JPL 606-1
1
1
•,. ■• • 'd
6N19 — A multitude of flat-bottomed craters in the bright area
Deucalionis Regio are clearly visible in the lower two-thirds of the
picture. Dark area Meridiani Sinus appears in the upper third of the
frame. Distinction between light/dark area does not show in this Max-D
processed version of the picture.
The craters have sizes ranging from a few km to about one hundred
km across. Under the favorable low solar illumination angle, some
large faint 'ghost' craters are discernible.
Fig. 37. Near Encounter Frame 6N19.
Sec. 3. 6, page 4i
C. Michaux, JPL
June 1, 1971
JPL 606-1
Mariner 1969 Photographic Atlas of Mars
6N20 — Small bowl-shaped craters of various sizes. A double -
ring, concentric structured 'ghost' crater is visible on the left. Faintly
visible on the right is a low, irregularly sinuous ridge, oriented roughly
north to south.
June 1, 1971
Fig. 38. Near Encounter Frame 6N20.
C. Michaux, JPL Sec. 3.6, page 49
Mariner 1969 Photographic Atlas of Mars
JPL 606-1
mj: ■
6N21 — This picture shows a continuation of the cratered terrain
of frames 6N19 and 6N17, eastward over Deucalionis Regio, in the lower
two-thirds of the picture, and Sabaeus Sinus in the upper third. Large
flat-bottomed craters show up distinctly in the low sun. One crater is
about 250 km across.
Approximately 250 flat-bottomed craters with diameters over 7 km
were counted in the three frames 6N17, 19, and 21 for a size -frequency
distribution analysis.
The boundary between light Deucalionis Regio and dark Sabaeus
Sinus (barely visible in this Max-D picture) has a "diffuse'' character,
as in frame 6N19.
Fig. 39. Near Encounter Frame 6N21.
Sec. 3.6, page 50 C. Michaux, JPL
June 1, 1971
JPL 606-1
Mariner 1969 Photographic Atlas of Mars
6N22 — Closeup view showing a 'smaller' flat-bottomed crater upon
the rim and the irregular wall of the large 250-km crater seen in frame
6N21.
A number of very small bowl-shaped craters are also visible.
From this frame and frame 6N20 a count was made of 104 craters with
diameters over 0.7 km.
Fig. 40. Near Encounter Frame 6N22.
June 1, 197 1
C. Michaux, JPL
Sec. 3. 6, page 5 1
Mariner 1969 Photographic Atlas of Mars JPL 606-1
CHAOTIC TERRAIN
Mosaic of Four Camera A Frames 6N 1 Through 6N7
The mosaic shows an overall view of the chaotic terrains, photographed
by Mariner 6 in the equatorial regions of Mars. These terrains are irregularly
shaped and irregularly distributed in the mixed light and dark area Pyrrhae
Regio, with prongs extending into dark areas Aurorae Sinus to the west,
Margaritifer Sinus to the northeast, and into light area Chryse to the north.
Sec. 3.6, page 52 C. Michaux, JPL June 1, 1971
JPL 606-1
Mariner 1969 Photographic Atlas of Mars
Fig. 41. Mosaic 6N1 through 6N7 (chaotic terrain).
June 1, 1971
C. Michaux, JPL
Sec. 3. 6, page 53
Mariner 1969 Photographic Atlas of Mars
JPL 606-1
6N3 — This picture is another limb view encompassing more terri-
tory eastward (left) into dark area Aurorae Sinus, where some chaotic
terrain patches already appear amidst cratered terrain. Xanthe desert
is shown toward the limb.
The heavy dark band is the residual image of the limb from a
previous frame.
Fig. 42. Near Encounter Frame 6N3.
Sec. 3. 6, page 54 C. Michaux, JPL
June 1, 1971
JPL 606-1
Mariner 1969 Photographic Atlas of Mars
6N5 — Many irregular patches of chaotic terrain covering a sizable
area are visible in this oblique view of Aurorae Sinus, with part of
Pyrrhae Regio visible at the top right.
This region, in which chaotic and cratered terrains are intricately
intermixed, is difficult to analyze geologically.
June 1, 1971
Fig. 43. Near Encounter Frame 6N5.
C. Michaux, JPL Sec. 3.6, page 55
Mariner 1969 Photographic Atlas of Mars
JPL 606-1
290° 300°
JSO* 330" 340»
EAST LONGITUDE
3S0»
-■lO*
300*
310'
320* 330-
EAST LONGITUDE
360-
Fig. 44. Distributions of light and dark markings (top) and chaotic terrain
(bottom) in equatorial region photographs 6N5, 7, and 9
(after Cutts at al. , 1971).
Sec. 3. 6, page 56
C. Michaux, JPL
June 1, 1971
JPL 606-1
Mariner 1969 Photographic Atlas of Mars
h-a ..
\NI^'
4^
6N6 — This picture is a high- re solution view of an area of chaotic
terrain with display of its characteristic pattern of ridges and troughs.
June 1, 1971
Fig. 45. Near Encounter Frame 6N6.
C. Michaux, JPL
Sec. 3. 6, page 57
Mariner 1969 Photographic Atlas of Mars
JPL 606-1
'W^''^'
^
• .'^ -t
6N7 — This shows the irregular but well-defined patches of chaotic
terrain in Aurorae Sinus at left and Pyrrhae Regio at right. Note the
abrupt vertical scarp shadow line at the top left. The remainder of the
terrain is fairly well cratered with old, flat-bottomed craters.
For high-resolution frames of this area see 6N6 and 6N14, and
6N8, where the details of chaotic structure are strikingly revealed.
Fig. 46. Near Encounter Frame 6N7.
Sec. 3.6, page 58 C. Michaux, JPL
June 1, 1971
JPL 606-1
Mariner 1969 Photographic Atlas of Mars
6N8 — A high-resolution view of the transition from cratered ter-
rain to chaotic terrain. The transition to chaotic "displays all the charac-
teristics of a slump zone on Earth, although scale is unusually large, "
Sharp et al. (1971b). Note the abrupt scarps at bottom right.
The cratered terrain has a rather large, old, flat-bottomed crater
and many fresh-looking small bowl-shaped craters.
June 1, 1971
Fig. 47. Near Encounter Frame 6N8.
C. Michaux, JPL Sec. 3.6, page 59
Mariner 1969 Photographic Atlas of Mars
JPL 606-1
6N14 — High-resolution view of the "Chaos Valley" barely visible
in frame 6N7. To quote Prof. Sharp (it) "looks like a feature which has
extended itself headward and sidewise into cratered terrain" (Sharp
et al. , 1971b).
Fig. 48. Near Encounter Frame 6N14.
Sec. 3.6, page 60 C. Michaux, JPL
June 1, 1971
JPL 606-1
Mariner 1969 Photographic Atlas of Ma
rs
6N15 — This picture shows clearly defined brighter patches in the
distance at left which are probably chaotic terrain across otherwise well-
cratered terrain. Details of these patches cannot be seen because of
highly oblique viewing.
This frame extends over much territory, including the base of
Margaritifer Sinus, Pyrrhae Regio, and Aurorae Sinus at left. As many
as five high-resolution B -frames were taken within it: 6N4, 6, 8, 14,
and 16. Three of these (6N6, 8, and 14) unmistakably show some chaotic
terrain.
June 1, 1971
Fig. 49. Near Encounter Frame 6N15.
C. Michaux, JPL Sec. 3. 6, page 61
Mariner 1969 Photographic Atlas of Mars JPL 606-1
FEATURELESS TERRAIN
Mosaic of Frames 7N21 Through 7N31
This mosaic depicts six A-frames spanning from the west limb, over
the Meridiani Sinus dark band (7N21, left), bright strip Deucalionis Regio,
Pandorae Fretum, a variable patch, bright desert Noachis, dark well-
cratered Hellespontus, and across featureless circular desert Hellas into
evening terminator (7N31 at right). Maximum discriminability versions (Max-D)
were used. Max-D brings out topographic detail at the expense of albedo
contrast.
Sec. 3.6, page 62 C. Michaux, JPL June 1, 1971
JPL 606-1
Mariner 1969 Photographic Atlas of Mars
u
(U
+->
Ul
w
0)
>—(
i)
u
(1)
cm
o
u
o
•r-l
June 1, 1971
C. Michaux, JPL
Sec. 3. 6, page 63
Mariner 1969 Photographic Atlas of Mars
JPL 606-1
7N25 — This shows the heavily cratered dark area Ilelle spontus,
bordered to the east by featureless Hellas, just discernible at the top
right-hand corner. In the wide transition zone into Hellas, the craters
not only become rarer, flatter, and fainter, but there exists a scries of
overlapping narrow linear ridges and scarps facing Hellas and parallel
to its border. These features suggest definite crustal creep downward
toward Hellas with a succession of abrupt drops in elevation.
Width of transition zone varies from 150 to 325 km.
ridges and scarps range from 20 to 90 km.
Length of
Fig. 51. Near Encounter Frame 7N25.
Sec. 3. 6, page 64
C. Michaux, JPL
June 1, 1971
JPL 606-1
Mariner 1969 Photographic Atlas of Mars
7N26 — Closeup picture of the conspicuous ridges and scarps in
the transition zone from Hellespontus to Hellas. Length of ridges is
about 40 km, while scarps, at bottom left, are shorter.
This local area of the transition zone is unusually poor in small
craters.
June 1, 1971
Fig. 52. Near Encounter Frame 7N26.
C. Michaux, JPL Sec. 3.6, page 65
Mariner 1969 Photographic Atlas of Mars
JPL 606-1
7N27 — This shows the interior of Hellas is a monotonously smooth
surface with only three small flat- floored craters visible, on its western
edge.
The transition zone into Hellespontus on the west is irregularly
defined along its Hellas edge, by an abrupt contrast fronn light (in Hellas)
to dark in the transition zone. The contrast does not sho-w up here in
these Max-D versions used to enhance local topographic details. The
scarps and ridges on the western, or Hellespontus, edge of the zone do
stand out.
Flat-floored craters are present in the transition zone, but they
gradually disappear as the zone slopes downward into Hellas.
Fig. 53, Near Encounter Frame 7N27,
Sec. 3.6, page 66 C. Michaux, JPL
June 1, 1971
JPL 606-1
Mariner 1969 Photographic Atlas of Mars
7N28 — This high-resolution frame confirms the featureless ter-
rain within Hellas. Resolution is ahout one-half km.
June 1, 1971
Fig, 54. Near Encounter Framie 7N2J
C. Michaux, JPL
Sec . 3.6, page 67
Mariner 1969 Photographic Atlas of Mars
JPL 606-1
ye-
-^f¥ '
7NZ9 — This is another picture of the featureless central parts
of Hellas.
Fig. 55. Near Encounter Frame 7NZ9.
Sec. 3.6, page 68 C. Michaux, JPL
June 1, 1971
JPL 606-1
Mariner 1969 Photographic Atlas of Mars
ATMOSPHERIC HAZE (Limb Pictures)
Mosaic of Frames 7N1 and 7N3
This mosaic covers the highly oblique northwestern view of complex
region, composed of dark fork Meridiani Sinus, bright channel Thymiamata,
and dark oasis Oxia Palus. Due to incomplete erasure by the vidicon system,
only the upper, or northern, portions of these areas appear in this mosaic.
The bright desert Chryse is beyond, on the limb. See frame 7N2 for a magni-
fication of the limb portion.
7N2
7N3
Fig. 56. Mosaic 7N1 and 7N3 (atmospheric haze).
June 1, 1971
C. Michaux, JPL
Sec. 3. 6, page 69
Mariner 1969 Photographic Atlas of Mars
JPL 606-1
7N1 — This is a northwest limb, blue-filter frame over Chryse
desert (on limb), revealing a thin haze layer distinctly detached from the
limb. Altitude of the base of the haze layer was estimated at about 5 km
above the surface of the desert. Geometrical thickness of the haze is
about 5 km (Leovy et al. , 1971).
Photometric comparison of the brightness of this layer in three
colors, using this blue-filter frame and 7N3 (green) and 7N5 (red), shows
it to be essentially white, with an absence of blue coloring anywhere in
the layer.
Fig. 57, Near Encounter Framie 7N1.
Sec. 3. 6, page 70 C. Michaux, JPL
June 1, 1971
JPL 606-1
Mariner 1969 Photographic Atlas of Mars
7N2 — High-resolution view of the detached thin haze layer
appearing in frames 7N1 and 7N3. Under even higher magnifications and
special processing, the layer reveals its multilayered structure of at
least three layers (Leovy et al. , 1971).
June 1, 1971
Fig. 58, Near Encounter Frame 7N2.
C. Michaux, JPL
Sec. 3. 6, page 7 1
Mariner 1969 Photographic Atlas of Mars
JPL 606-1
DARK AND LIGHT AREAS (Meridiani Sinus Complex)
M osaic of Three A Frames 7N5, 7N7, and 7N9 (Photometric Version)
This mosaic shows an overall oblique view of the equatorial region of
prominent dark area Meridiani Sinus, which forms the enlarged hand-like
end of the long dark arm Sabaeus Sinus, alongside the Martian equator.
(Limb orientation is about NNW, toward Mare Acidalium, ) Seen here at
8000-km slant range, Meridiani Sinus appears well cratered, not only within,
but across its unusually varied boundary.
The light area Edom, which appears on maps as a round notch at the
corner of Meridiani Sinus and Sabaeus Sinus, turns out to be a very large,
bright-floored, flat-bottomed crater. See top right portion of frame 7N9.
Crater wall defines a sharp boundary between light and dark areas, suggesting
topographic control of the albedo, with the light area at lower elevation (Cutts
et al. , 1971),
7N9
7N8
Fig. 59. Mosaic 7N5, 7N7, and 7N9 (dark and light areas).
Sec. 3.6, page 7 2 C. Michaux, JPL June 1, 1971
JPL 606-1
Mariner 1969 Photographic Atlas of Mars
7N5 — A limb picture. A sharp detached haze layer appears in
other versions. Prominent dark area Meridian! Sinus is well cratered.
The craters display remarkable albedo patterns with light markings on
their south-facing inner slopes. Note the isolated, dark-floored craters
in neighboring bright areas also. The crater complex projecting from
top left, not far from the limb, is the "oasis" Oxia Palus.
June 1, 1971
Fig, 60, Near Encounter Frame 7N5.
C. Michaux, JPL Sec. 3.6, page 73
Mariner 1969 Photographic Atlas of Mars
JPL 606-1
7N6 — This shows a high-resolution view of the craters in the
middle of Meridiani Sinus. Note their as/mnietric shading (not due to
local highnoon sun), an aeolian deposition process of the light material
inside dark-floored craters is suggested here.
Fig. 61. Near Encounter Frame 7N6.
Sec. 3.6, page 74 C. Michaux, JPL
June 1, 1971
JPL 606-1
Mariner 1969 Photographic Atlas of Mars
7N7 — Meridian! Sinus is shown slightly more to the east than in
frame 7N5. The dark arc is a residual image. West of the crater Edom
(just appearing at right) there is a bright area encroaching erratically into
the main body of Meridiani Sinus. Note that its boimdary is very sharp
as well as highly irregular, in that "small dark outliers" are found on
the bright side, which has variable albedo (see Cutts et al. , 1971).
June 1, 1971
Fig. 62, Near Encounter Frame 7N7.
C. Michaux, JPL Sec. 3.6, page 75
Mariner 1969 Photographic Atlas of Mars
JPL 606-1
SOUTH POLAR CAP
Mosaic of Frames 7N11 Through 7N19
This mosaic is presented in the two computer-processed versions. The
photometric version (vv^ithout AGC), Fig. 63 shows the actual brightness of the
South Polar Cap's white expanse against the bare ground surrounding it. Cra-
ters are visible along its edge and in its interior, where floors appear darker,
although frost-filled. The grid lines of latitude and longitude permit location
of the South Pole and show that the cap extends northward beyond 60 °S.
Fig. 63. Mosaic 7N11 through 7N19 (photometric version).
Sec. 3.6, page 76 C. Michaux, JPL June 1, 1971
JPL 606-1
Mariner 1969 Photographic Atlas of Mars
The maximum discriminability (Max-D) version mosaic (with AGC),
Fig. 64 reveals, in a spectacular way, an intricate wealth of detail across the
cap. Besides craters, there is evidence of a surprising variety of forms, due
to both frost and underlying ground.
All Max-D versions of the polar cap edge show an unfortunate black band
artifact which is not to be interpreted as the traditional dark polar collar.
When the AGC effect is removed, as in the photometric version, there is no
suggestion of it.
Fig. 64. Mosaic 7N11 through 7N19 (Max-D version).
June 1, 1971 C. Michaux, JPL Sec. 3.6, page 77
Mariner 1969 Photographic Atlas of Mars
JPL 606-1
■ih:-
7N11 — The extra-marginal edge of the South Polar Cap is best
displayed on the west side of this wide-angle frame. The craters are
apparently as abundant and of similar shape and size, to those elsewhere
on Mars, where cratered terrain is prevalent. Note the presence of
the black band artifact, which is absent in the photometric version
(Fig. 65b).
Fig. 65(a). Near Encounter Frame 7N11 (Max-D version).
Sec. 3.6, page 78 C. Michaux, JPL June 1, 1971
JPL 606- 1
Mariner 1969 Photographic Atlas of Mars
Fig. 65(b). Near Encounter Frame 7N11 (photometric version).
June 1, 1971
C. Michaux, JPL
Sec. 3.6, page 79
Mariner 1969 Photographic Atlas of Mars
JPL 606-1
:.#■
7N12 — This shows a closeup of a portion of the outer-marginal
zone of the South Polar Cap, located in the dark band artifact mentioned
earlier. Craters within this zone appear grotesquely distorted as a
result of the unusual combination of low solar illumination and frost
enhancement.
The brilliant patches are frost accumulations on portions of
crater bottoms and on south-facing slopes which are more protected
from solar rays.
Fig. 66. Near Encounter Frame 7N12.
Sec. 3.6, page 80 C. Michaux, JPL
June 1, 1971
JPL 606-1
Mariner 1969 Photographic Atlas of Mars
1 "^
.-m^'
7N13 — The edge of the polar cap is shown again in this frame,
as in frame 7N11, and also much more of the interior of the cap, where
the frost layer becomes thick rapidly. Outlines of frost-filled craters
are quite recognizable, but new unfamiliar forms make their appearance.
See frame 7N14 closeup.
Note again along the edge of the cap how craters have their south-
facing slopes covered with frost, while their north-facing slopes may be
bare, due to solar exposure.
Fig. 67(a). Near Encounter Frame 7N13 (Max-D version).
June 1, 1971 C. Michaux, JPL Sec. 3.6, page 81
Mariner 1969 Photographic Atlas of Mars
JPL 606-1
Fig. 67(b). Near Encounter Frame 7N13 (photometric version).
Sec . 3.6, page 82
C. Michaux, JPL
June 1, 1971
JPL 606-1
Mariner 1969 Photographic Atlas of Mars
7N14 — This polar cap interior view shows a surprising etch-pit
pattern developed within a frost-filled crater, humorously named the
"elephant's footprint."
Frost cover in this narrow-angle frame otherwise appears remark-
ably even, except for the display of some small subdued features of posi-
tive relief, which are characteristic of the polar cap interior, as are the
etch-pit depressions.
June 1, 1971
Fig, 68. Near Encounter Frame 7N14.
C. Michaux, JPL Sec. 3.6, page 83
Mariner 1969 Photographic Atlas of Mars
JPL 606-1
.1.
■ 3k
7N15 — This picture shows the portion of the polar cap interior
cut across by the zero meridian. The marginal zone of the cap, top
left, borders Depressio Hellespontica, well known as one of the areas
where the "wave of darkening" starts in southern spring. The interior
of the cap with continuous frost cover displays three types of features:
craters of various sizes, irregular depressions called 'etch pits', and
parallel features called 'beaded lineations'.
The etch pits are in the center. Note their lighter rims and som-
ber crater-like floors of irregular, angular outline. The beaded linea-
tions, to the south, are short alignments in WNW direction connecting
tiny nodes. Such features may be the result of wind-drifted and scoured
snow.
Fig. 69(a). Near Encounter Frame 7N15 (Max-D version).
Sec. 3.6, page 84 C. Michaux, JPL June 1, 1971
JPL 606-1
Mariner 1969 Photographic Atlas of Mars
Fig. 69(b). Near Encounter Frame 7N15 (photometric version).
June I, 1971
C. Michaux, JPL
Sec. 3. 6, page 85
Mariner 1969 Photographic Atlas of Mars
JPL 606-1
7N16 — This closeup picture shows the rather uniform frost coating
of the cap interior. Local relief is very well evidenced by low solar
illumination. Regions of positive relief are present, very few small
craters are seen, and there are some shallow, depressed regions with
irregular boundaries.
Fig. 70. Near Encounter Frame 7Nl6.
Sec. 3.6, page 86 C. Michaux, JPL
June 1, 1971
JPL 606-1
Mariner 1969 Photographic Atlas of Mars
7N17 — Here is another view of the polar cap interior. The same
general appearance, with unusually bright crater rims and somber floors,
and etch pits are shown, as in frame 7N15. Also displayed in the interior
are isolated, irregularly shaped patches and bands of unusual brightness
(at top center and right), which must be the result of local topography
and meteorology.
The central polar region appears in the lower right part of the
frame, with eye-catching, long sinuous lineations, roughly concentric
on the South Pole and spreading eastward. These may be up to 300-km
long and 10-km wide. Note the sharp buckle of one. These enigmatic
lineations are probably surface features and not clouds, because of their
sharp boundaries, shape, shading, and distinct shadows. They may jae
troughs, ridges, or scarps, composed wholly of snow or ice combined
or not with rock debris (see Sharp et al. , 1971a).
The central polar region is also characterized by paucity and
faintness of craters and the unusual smoothness of its snow or ice cover,
which is probably a rather thick layer at least tens of meters deep.
Fig. 71(a). Near Encounter Frame 7N17 (Max-D version).
June 1, 1971 C. Michaux, JPL Sec. 3.6, page 87
Mariner 1969 Photographic Atlas of Mars
JPL 606-1
Fig. 71(b). Near lilncounter Frame 7N17 (photometric version).
Sec. 3. 6, page
C. Michaux, JPL
June 1, 1971
JPL 606-1
Mariner 1969 Photographic Atlas of Mars
7N18 — This closeup shows three large (--50 km diameter) crater
forms under the thick frost cover, taken well within the cap interior not
far from the central region. The striations in the frost accumulation
suggest the effect of persistent winds blowing predominantly in one direc-
tion. The presence of the three small craters, probably bowl-shaped,
indicate that frost thickness cannot be too great.
June 1, 1971
Fig. 72. Near Encounter Frame 7N18.
C. Michaux, JPL Sec. 3. 6, page 89
Mariner 1969 Photographic Atlas of Mars
JPL 606-1
7NI9 — Again the polar cap interior and central regions are shown,
as in frame 7N17. Sunset terminator is at the right. The linear and
blotchy features are the same as in frame 7N17.
A pair of rather large craters formiing the "giant's footprint" can
be found in the center right among the many varied craters outlined by
frost. Crater abundance of the cap interior is high in the northeast
corner, as high in fact as in the most heavily cratered terrains photo-
graphed. A few craters in the left part of the frame have "small white
central dots suggestive of central peaks or an unusual accumulation of
frost. ''
The central polar region is best seen in this franne and frame
7N17. "Its boundary is most clearly defined in the central and eastern
part of frame 7N19 as an abrupt, irregular, crater- scalloped arc with
some vertical relief. " (Quotations from Sharp et al. , 1971a).
Fig. 73(a). Near Encounter Frame 7N19 (Max-D version).
Sec. 3.6, page 90 C. Michaux, JPL June 1, 1971
JPL 606-1
Mariner 1969 Photographic Atlas of Mars
'*#
Fig. 73(b). Near Encounter Frame 7NI9 (photometric version).
June 1, 1971
C. Michaux, JPL
Sec. 3.6, page 91
Mariner 1969 Photographic Atlas of Mars
JPL 606-1
7N20 — This shows a high-resolution oblique view of the crater
pair, forming the "giant's footprint," The irregularly ridged arch of
the footstep is part of the boundary separating the central polar region
(south) from the polar cap interior (north).
The Sun was low, only about 8° above the horizon, bringing out
all details in this truly unusual view over the Martian South Cap.
Fig. 74, Near Encounter Frame 7N20.
Sec. 3.6, page 92 C. Michaux, JPL
June 1, 1971
JPL 606-1 Mariner 1969 Photographic Atlas of Mar;
BIBLIOGRAPHY
Collins, S. A. , 1971, The Mariner 6 and 7 pictures of Mars, NASA SP-263,
December.
Cutts.J.A. , Soderbloom, L. A. , Sharp, R. P. , Smith, B. A. , and Murray, B, C. ,
1971, The surface of Mars, 3: light and dark markings, J. Geophys. Res. j
V. 76, no. 2, p. 343-356, January 10.
Danielson, G. E. , Jr., and Montgomery, D. R. , 1971, Calibration of the Mariner
Mars 1969 television cameras, J. Geophys. Res. , v. 76, no. 2, p. 418-431,
January 10.
Dunne, J. A., Stromberg, W. D. , Ruiz, R. M. , Collins, S. A. , and Thorpe, T. E. ,
1971, Maximum discriminability versions of the near -encounter Mariner
pictures, J. Geophys. Res. , v. 76, no. 2, p. 438-472, January 10.
Leovy, C. B. , Smith, B. A., Young, A. T., and Leighton, R. B. , 1971, Mariner
Mars 1969 atmospheric results, J. Geophys. Res . , v. 76, no. 2,
p. 297-312, January 10.
Montgomery, D. R. , and Adams, L. A. , 1970, Optics and the Marine r imaging
instrument, Appl. Optics, v. 9, p. 277.
Rindfleisch, T, C. , Dunne, J. A., Frieden, H. J. , Stromberg, W. D. , and
Ruiz, R. M. , 1971, Digital processing of the Mariner 6 and 7 pictures,
J. Geophys. Res. , v. 76, no. 2, p. 394-417, January 10.
Sharp, R. P., Murray, B.C., Leighton, R. B, , Soderbloom, L, A. , and
Cutts, J. A. , 1971a, The surface of Mars. 4: south polar cap,
J. Geophys. Res. , v. 76, no. 2, p. 357-368, January 10.
Sharp, R. P., Soderbloom, L. A. , Murray, B.C., and Cutts, J. A. , 1971b,
The surface of Mars. 2: uncratered terrains, J. Geophys . Res. , v. 76,
no. 2, p. 331-342, January 10.
June 1, 1971 C. Michaux, JPL Sec. 3.6, page 93
JPL 606-1 Observational Phenomena
SECTION 4 CONTENTS
4. OBSERVATIONAL PHENOMENA
Introduction j
4. 1 Clouds and Hazes
Introduction
4. 1. 1 Data Summary
The Violet Layer
Blue Clearing
White Clouds
Yellow Clouds
4. 1. 2 Discussion 2
Early Observations 2
The Violet Layer, Blue Clouds, and Blue Clearing 2
White Clouds £,
Yellow Clouds g
Gray Clouds and Bright Spots 14
Bibliography 26
Figures
1. 1971 Yellow storm, days 1-7; contours overlaid on
pre-storm map 10
2. 1971 Yellow storm, days 8-17; contours overlaid on
pre-storm map 11
Tables
1. Table of observations of Martian blue clearings 4
2. Martian clouds and hazes 1966-68 7
3. Table of major Martian "dust storms" 12
4. Table of bright flares and spots observed on Mars 15
4. 2 Seasonal Activity
Data Summary 1
The Polar Caps and Hoods 1
Seasonal Behavior of Clouds and Whitenings 1
The Wave of Darkening 1
Seasonal Behavior of Surface Features 2
Introduction 2
Polar Caps and Hoods 3
Polar Caps 3
March 1, 1972 Sec. 4, Contents, pagei
Observational Phenomena JPL 606-1
4. 2 (cont'd)
Polar Hoods 6
The Dark Polar Fringe 6
Boundaries of North and South Polar Caps 7
South Cap Regression 9
North Cap Regression 11
Seasonal Behavior of Clouds and Whitenings 13
White Clouds and Hazes 17
Seasonal and Recurrent White Clouds 21
Great Yellow Clouds 23
Whitening Areas 24
The Wave of Darkening 25
Seasonal Behavior of Surface Features 29
General Comments on Martian Colors 29
Bibliography 43
Appendix — Martian Seasonal Dates A-1
Figures
1. Measured width of the south polar cap of Mars for
various oppositions from 1798 to 1924 4
2. Seasonal evolution of the Martian polar caps and hoods
according to polarimetric and photometric measurements 5
3. Seasonal boundaries of the South Polar Cap as viewed
from the south 8
4. Seasonal boundaries of the North Polar Cap as viewed
from the north 9
5. Regression curve of the South Polar Hood and Cap,
mean 1905-1965 10
6. Mariner Mars 1971 Planning Charts of the South
Polar Regions 12
7. Regression curve of the North Polar Cap and Hood,
mean 1905-1965 13
8. Color map of the Martian surface in northern fall-
winter and southern spring-summier, with white and
yellow cloud activity during these seasons *
9. Color map of the Martian surface in northern spring
and southern fall, with wave of darkening, frost, and
white and yellow cloud activity during these seasons -''
10. Color map of the Martian surface in northern summer
and southern winter, with wave of darkening and white
and yellow cloud activity during these seasons *
11. Map and table of Martian place-names and their locations *
12. North Polar Cap micrometric regression curve, 1968-69 14
13. North Polar Cap photographic regression curve, 1966-67 14
14. North Polar Cap photographic regression curve, 1964-65 15
15. North Polar Cap photographic regression curve, 1962-63 15
16. Seasonal evolution of the North Polar Cap (Ls =60°) 17
17. Seasonal evolution of the North Polar Cap (Ls =90°) 18
18. Seasonal evolution of the North Polar Cap (Ls = 140°) 19
19. Seasonal evolution of the North Polar Cap (Ls = 160°) 20
'.^Material is physically located behind new Section 4. 2, page 45, dated
February 1, 1972.
Sec. 4, Contents, page ii March 1, 1972
JPL 606-1 Observational Phenomena
4. 2 (cont'd)
20. Schematic base map of Mars 1967 showing whitening
areas detected in 1962-1968 26
21. The seasonal waves of darkening of the dark areas of Mars
according to the photometric data of Focas 28
22. Mars, seasonal variation of polarization 29
Tables
1. North cap regression rates 1962-1969 16
2. Mean areographic longitude of centroid of white island
remnants of summer North Polar Cap 16
3A. Seasonal meteorological activity, 1966-1970 22
3B. Relative occurrence of seasonal cloud activity, 1966-1968 .... 22
4. Table of seasonal changes in northern dark areas of Mars .... 30
5. Table of seasonal changes in equatorial dark areas of Mars ... 32
6. Table of seasonal changes in northern light areas of Mars .... 34
7. Table of seasonal changes in equatorial light areas of Mars ... 36
8. Table of seasonal changes in south polar region of Mars 39
9. Table of seasonal changes in the north polar region of Mars ... 41
March 1, 1972 Sec. 4, Contents, page iii
JPL 606-1 Observational Phenomena
OBSERVATIONAL PHENOMENA
INTRODUCTION
This section of the Mars Scientific Model includes discussion of many of
those aspects of Mars which were originally discovered by direct visual obser-
vation. Section 4. 1 encompasses discussion of such persistent, but often
changing, phenomena as the clouds and hazes, while Section 4. 2 is devoted to
phenomena which change seasonally, such as the polar caps and the so-called
"wave of darkening. " There is a seasonal character to cloud activity which is
also discussed in Section 4. 2. It is hoped that a third section can be added at
some future date to provide a discussion of the many secular changes in the
Martian surface noted during the past 100 years.
More than brief mention of these phenomena is important in order to
convey a proper appreciation of Mars as a dynamic, changing planet. This
perspective is often lost in the presentation of quantitative data and simplified
theoretical models. Yet, the very transitory nature of these phenoraena has
made quantitative study from Earth so difficult as to be practically nonexistent.
It has resulted in fuzzy, subjective descriptions such as "violet layer, " "blue
clearing, " "wave of darkening, " etc. , which have no root in quantitative
physical study. The following sections are important primarily as a prelimi-
nary taxonomy based upon remote observation. The observations are of
sufficient difficulty that some of them will prove partially fallacious, while
attempts at physical explanation are at best tentative hypotheses. Nevertheless,
this tentative data is the product of 100 years of labor and does offer an
extended time base within which the new physical studies can be better under-
stood. Therefore, the following information should be treated with reservation
but not simply rejected.
December 15, 1971 R. Newburn, JPL Sec. 4, page 1
JPL 606-1 Clouds and Hazes
4. 1 CLOUDS AND HAZES
INTRODUCTION
Two types of Martian clouds are always distinguished by ground -based
observers, white clouds, and yellow clouds. Many observers also recognize
a distinct class of blue clouds, seen only in blue or violet light and distinct in
properties and behavior from the white clouds. White clouds in turn exhibit
several distinct types of behavior, and some of the so-called "clouds" may not
be clouds at all but rather surface frost. Attempts at physical understanding
of Martian cloud morphology are primitive at present, and many traditional
"explanations" are subject to question by the Mariner series of Mars probes.
Such is the class of phenomena discussed in this section.
4.1.1 DATA SUMMARY
The Violet Layer
The violet layer, sometimes called the blue haze or violet haze, is a high
altitude atmospheric layer of submicron-sized particles hypothesized to explain
the lack of contrast observed on Mars in blue and violet light. In fact, much
of this lack of contrast is an intrinsic property of the surface material, and
only remains a part of contemporary Martian theory because of the apparent
need to explain the "blue clearing" phenomenon.
Blue Clearing
Generally, Martian surface detail can only be observed at wavelengths of
4500 A or longer.
However, at times (with a transition period of about a day), the atmo-
sphere of Mars becomes clear enough for a period of one or more days for
surface detail to be seen clearly at wavelengths of 4250 A or less. This
"blue clearing" of the Martian atmosphere is attributed to a temporary dissi-
pation of the violet layer (by those who believe in the violet layer).
White Clouds
White clouds are those which can be photographed at short wavelengths,
but disappear in red or infrared photos. Sometimes an added distinction is
made for certain "blue clouds" based upon behavior and polarization properties.
White clouds are generally thought to be a fog of submicron ice particles
similar to terrestrial cirrus clouds. White clouds are of all shapes and sizes,
ranging from terminator haze lasting only a few hours to dense, 2000-km giants
lasting days or weeks.
Yellow Clouds
Yellow clouds are those readily photographed in yellow or red light, but
which are not seen in blue light. They may be small, dense, orange or yellow
objects lasting only one or a few days, or they may start large and grow larger
until they become a yellow veil covering most of Mars and lasting a month or
December 15, 1971 R. Newburn, JPL Sec. 4.1, page 1
Clouds and Hazes JPL 606-1
more. Yellow clouds are believed to be clouds of dust particles, perhaps
2-50 fx in diameter, raised by high surface winds.
4.1.2 DISCUSSION
Early Observations
Yellowish clouds and veils were reported by visual observers of Mars at
least as early as 1809 (by Flaugergues), and white clouds were reported by
Secchi in 1858 (Maggini, 1939). Such observations were made with refracting
telescopes, color-corrected for visual observing. Under such conditions,
color descriptions are quite subjective, but the differences in appearance which
resulted in a distinction between yellow and white clouds were real enough.
In 1871, the dry photographic plate was developed, raaking astronomical
photography a practical possibility rather than a stunt, as it had been during
the wet plate era. The "natural" spectral sensitivity of a photographic emul-
sion extends from the ultraviolet to about 5200 A, there being some dependence
upon the exact ratio of halides used in the emulsion and the physical processing
during manufacture. Although it was discovered by Vogel in 1873 that the range
of spectral sensitivity could be extended with dyes, reasonably fast red plates
were not commerically available until 1906 (Mees, 1954). Most early photog-
raphy of Mars was done with "natural" blue plates exposed through a refracting
telescope, color-corrected in the visual (yellow) region of the spectrum. The
resulting blurred images were rather discouraging.
Lowell Observatory began an experimental Martian photographic program
in 1901 with "encouraging" results (Lowell, 1905). Four plates taken in 1903
and 99 plates from 1905 of useful scientific quality still exist (Baum et al. ,
1967). In 1907, Lowell described their process of extending plate sensitivity to
6800 A and the use of a filter to isolate the yellow-green region, for which
their 24-inch refractor was color-corrected. Many excellent photographic
images were obtained in 1907, largely with an 18-inch refractor taken to Chile,
where Mars came much nearer to the zenith. Some of these photographs are
reproduced in Slipher's 1962 book. Even more important was the development,
for use in 1909. of filters for separately isolating blue and red light (Slipher,
1962). The blue images "surprisingly" showed no detail whatsoever.
The Violet Layer, Blue Clouds, and Blue Clearing
The Lowell Observatory's featureless, blue photographic plates of 1909
resulted in the belief of a "violet layer" or "blue haze" of sufficient opacity
to make surface features undetectable, in the blue and violet wavelengths.
Actually, most of the loss of contrast in the blue is an intrinsic property of
the surface (see Sections 3. 2 and 3. 4). The question of a tenuous atmospheric
layer cannot be dismissed so simply, however. Whereas surface detail
normally becomes quite difficult to see at 4500 A and below (McCord and
Westphal, 1971; Slipher, 1962), at times, for periods usually lasting no more
than a few days, surface detail becomes quite obvious at 45 00 A, and Slipher
(1962) found good contrast even at 4250 A during such periods. This is the
so-called phenomenon of "blue clearing" discovered by Slipher (1937), the
Sec. 4,1, page 2 R. Newburn, JPL December 15, 1971
JPL 606-1 Clouds and Hazes
name indicating belief that dissipation of an atmospheric layer was responsible
for the improved visibility of surface features. A "clearing" may be
hemisphere-wide, or more limited in extent, and occurs within a period of
about one day.
Spacecraft observations have not yet contributed greatly to a solution of
the "blue clearing" problem. Mariners 4, 6, and 7 all carried imaging sys-
tems with passbands too far to the red. The Mariner 6 and 7 ultraviolet
spectrometer data suggest there is about three times as much Martian atmo-
spheric scattering in the ultraviolet as would be expected from a pure molec-
ular atmosphere (Barth and Hord, 1971). However, the thin haze layers
detected at the Martian limb by Mariners 6 and 7 are whitish, not blue, and
seem incapable of causing the contrast change attributed to the "blue haze"
(Leovy et al. , 1971).
Van Blerkom (1971) has carried out a series of calculations on the
contrast changes that can be produced by forward scattering or isotropic
scattering hazes of various optical depths. For values of optical depth and
limb darkening which seem appropriate to Mars, as seen by Mariner 4 and
given by Young (1969), it is possible to reduce contrast by means of isotropic
scatterers at the sub-Earth point from 15 percent to 10-12 percent for varying
sun angles. Forward scattering hazes are even less efficient in reducing
contrast. An optical depth of 0. 3 is required to reduce contrast to 2-3 percent,
and this is some six times the value observed by Mariner 4. These observa-
tions alone may be sufficient to conclude that an aerosol layer is not respon-
sible for the blue clearing phenomona, but Mars is a dynamic planet, and high
quality photometry spread over some weeks or months is really needed to
understand it.
Attempts have been made to correlate visual and photographic obser-
vations of blue clearings with areographic longitude, Martian season, and
time from opposition. These results are given in Slipher (1962), but are
not reproduced here because of considerable correlation of each with observa-
tional selection (ease of making the measurement). Some data on blue clear-
ings are tabulated in Table 1.
If blue clearing does occur, and it is not (or not totally) an atmospheric
effect, then some change in the surface itself must be responsible. The
photometry of Boyce and Thompson (1971) suggests the intensity of the Martian
bright areas is strongly phase dependent while that of the dark areas is not.
Such a change in contrast with planetary phase and rotation could easily be
responsible for many, if not all, of the "blue clearings. " Pollack and Sagan
(1967) have preferred to go the whole way and suggest there is no real clearing,
the apparent clearing being due entirely to excellent seeing at times when the
Martian atmosphere is cloud free, and when there is an area of intrinsic high
contrast near the Martian central meridian. It is certainly true that a blue
clearing cannot be detected unless there is a surface region of some contrast
beneath. It is equally true that good seeing is required, since bad seeing can
destroy visibility of all surface detail even in the red or infrared.
December 15, 1971 R, Newburn, JPL Sec. 4.1, page 3
Clouds and Hazes
JPL 606-1
Table 1. Table of observations of Martian blue clearings.
Dates of clearing
May 26-Jun 1, 1890
Nov 2-3, 1926
May 20-21, 1937
Jul 18-25, 1939
Oct 10, 1941
Nov 22, 1941
Jun 13-14, 1954
Jun 27-Jul 2, 1954
Aug 7 and 11, 1956
Aug 23-Sep 3, 1956
Oct 26, 1956
Sep 3, 1958^
Oct 13-15, 1958
Nov 6-10, 1958
Nov 31, I960;
Jan 17 and 27, 1961*=
Sep 26-28, 1964;
Oct 3-4, 1964;
Dec 30, 1964-Jan 1, 1965;
Mar 8-9, 1965^
Mar 6, 1967;
Jan 14, 1968^
Opposition
date
May 27, 1890
Nov 4, 1926
May 19, 1937
Jul 23, 1939
Oct 10, 1941
Oct 10, 1941
Jun 24, 1954
Jun 24, 1954
Sep 11, 1956
Sep 11, 1956
Sep 11, 1956
Nov 17, 1958
Nov 17, 1958
Nov 17, 1958
Dec 30, I960
Mar 9, 1965
Apr 15, 1967
Source
de Vaucouleurs, 1954
de Vaucouleurs, 1954
de Vaucouleurs, 1954
de Vaucouleurs, 1954
de Vaucouleurs, 1954
Slipher, 1962
Slipher, 1962
Pettit and Richardson,
1955
Slipher, 1962
de Vaucouleurs, 1957;
Slipher, 1962
Slipher, 1962
Richardson and Roques,
1959
Slipher, 1962
Slipher, 1962
Smith, 1961
Capen, 1966
Capen, 1970
^This clearing occurred during the beginning of the great dust storm.
Not all of the planet was covered during this period, and those dark
areas visible in yellow light were almost equally visible in blue light.
Note that this clearing occurred 74 days before opposition, clearly
indicating the apparent association with opposition to be an effect of
observational selection.
'^A number of other dates showed lesser clearing than these three.
A number of other dates showed lesser clearing than these four groups.
^Moderate clearing was noted on 30 days between Dec 66 and Jan 68.
Sec. 4. 1, page 4
R. Newburn, JPL
December 15, 1971
JPL 606-1 Clouds and Hazes
Nevertheless, there seem to be periods of excellent seeing, when surface
detail is easily visible in the yellow and red, and when no hint of a surface
feature can be seen in blue-violet light (Capen*).
Normal blue images of Mars, taken in good seeing, often have an irregu-
larity, showing structure that is not correlated with surface features (Humason,
1961). Often bright areas on blue photographs correlate with white clouds seen
visually. In other cases, white clouds are visible only as a thickening or
brightening in the blue and violet images and are sometimes called blue or violet
clouds. Dollfus (1961a) prefers to reserve the term "blue clouds" for clouds
visible only in photographs using "deep blue filters" and located near the morn-
ing and evening limbs. Such blue clouds differ radically in polarization proper-
ties from white clouds and typically seem to be produced by 3 ^jl droplets (Dollfus,
1961a). Sometimes a blue image shows a distinct planet-wide banded structure
(Slipher, 1962; Humason, 1961). In summary, blue images of Mars are far
from being uniform and featureless at all times or from showing structure cor-
related only with surface features as would be anticipated for a pure, molecular
atmosphere. At least there are certainly some atmospheric features of con-
siderable opacity.
The observer of Mars "sees" the result of the combination of several
phenomena occurring simultaneously. The reflectivity of the Martian surface
decreases sharply from 6000 A to 3500 A (due possibly to the presence of ferric
oxides), and contrast between light and dark areas is greatly reduced. Contrast
between bright and dark areas may change with phase angle and angles of inci-
dence and emission. Somewhere near 3000 A, the light scattered by atmo-
spheric gas molecules becomes equal to that reflected from the surface, but
molecular scattering decreases toward longer wavelengths according to the
\-4 Rayleigh law dependence. Added to these is light scattered from aerosols,
which are certainly present in the Martian atmosphere, even though they may
not constitute a violet layer or blue haze which can be dissipated, plus any
changes in ground reflection. Even under conditions of perfect terrestrial
seeing, or from a space probe, all surface detail except the polar caps and pos-
sible frost patches (which have a high constant reflectivity through the visible
part of the spectrum) should disappear somewhere between 3000 A and 4000 A
without any added atmospheric scattering,
"Explanation" of a phenomenon whose very existence is debated is always
controversial. Dollfus and Focas (1966) find a photometrically determined sur-
face pressure for Mars of 30 mb. The actual surface pressure is known to be
about one-fifth of that amount. The difference, if not an instrumental artifact,
must be caused by particulate scattering. Martian polarization, as measured by
Gehrels and Teska (1962), is not consistent with pure Rayleigh scattering but
might be compatible with atmospheric scattering by submicron particles and/or
ground effects, in addition to some Rayleigh scattering. Kuiper (1964) has com-
pared these results with theoretical scattering calculations for mixtures of
variable size submicron ice particles with indefinite results. He notes that
"pumping liquid nitrogen into an open vessel placed in very dry air" results in a
blue cloud of submicron ice crystals (frozen by the evaporating nitrogen) having
the sort of extinction properties required of the violet layer. When the air is
not dry, a white cloud results.
^Private communication,
December 15, 1971 R. Newburn, JPL Sec. 4.1, page 5
Clouds and Hazes JPL 606-1
The most likely composition of particles in Mars' lower atmosphere
would seem to be frozen H2O or dust. Wilson (1958) compared a series of low
dispersion (100 A/mm) blue spectra taken during blue clearings with an identi-
cal set taken under normal conditions. He found periodic maxima and minima
in the ratio of light reflected from Mars that were exactly supplementary to
light transmitted through terrestrial noctilucent clouds. Noctilucent clouds are
generally thought to be composed of submicron- size meteoritic dust or ice-
coated particles of meteoritic dust. It is also interesting to note that the inten-
sity of light, scattered from particles comparable in size to its wavelength,
characteristically goes through a number of maxima and minima (Mie
scattering). Furthermore, Leovy et al. (1971) have given reasonable arguments
that the white "thin hazes" they observe can not be frozen CO2 (although neither
are they a violet layer). Above an elevation of about 20 km, frozen COo is
possible even in the equatorial regions, however, and CO-, lies frozen on the
ground in the pole caps, so CO2 particles are certainly present on Mars.
In sulnmary, there is good evidence for aerosols in the Martian atmo-
sphere, and there seems to be fair (though not perfectly quantitative) evidence
for occasional changes in the apparent contrast of Martian surface features in
blue light. Spaceprobe evidence is insufficient because of the limited time
Mariners 4, 6, and 7 spent near Mars and because quantitative photometry in
the wavelength regions from 3000 A to 4500 A was not included in its investiga-
tions. It remains unclear whether there is any association between the aero-
sols and the apparent contrast changes.
White Clouds
Originally white clouds were defined as those which seemed white or
bluish when seen visually. With the advent of filter photography, it became
apparent that white and blue clouds were easily photographed at short visible
wavelengths but disappeared in the red and infrared, a simple operational dis-
tinction. Differentiation between white and blue clouds is not so simple, nor is
there general agreement that such a distinction exists, as both can be photo-
graphed in the blue. Dollfus' distinction, given on page 5, was based upon loca-
tion and polarimetry. Polarimetry of dense white clouds suggests they are a
fog of ice particles, similar to terrestrial cirrus clouds (Dollfus, 1961a).
Some white clouds are quite large, as much as 2000 km across, and
remain visible for days or weeks. They may form or dissolve at their edges,
may move at relatively high speeds, and have a tendency to appear above certain
specific regions of the planet (Dollfus, 1961b). They may appear as bright
prominences at the limb of Mars. Under excellent seeing conditions, large
clouds may exhibit considerable fine structure. Large clouds also may be sur-
rounded by an even larger area of thin haze, detectable only with a polarimeter.
As previously noted, bands covering the full apparent diameter of Mars are
sometimes visible on blue images of the planet.
Other white clouds are small, bright, and generally remain fixed in an
isolated location (Dollfus, 1961b). These may be surrounded by a large, fainter
cloud structure. The polar caps are usually said to be covered in Martian
winter with a hood of clouds similar in character to these small bright clouds
(see Section 4.2).
Sec. 4. 1, page 6 R. Newburn, JPL December 15, 1971
JPL 606-1
Clouds and Hazes
Clouds or hazes are often seen near the morning terminator of Mars
but usually disappear in a few hours. Sometimes they are also seen to form
near the evening terminator. Both morning and evening hazes are seen most
frequently in the Martian spring (Dollfus, 1961b).
C. F. and V. W. Capen (1971) have considered the following mor-
phological classes of clouds: polar hood, polar haze, planetary system cloud
banding, limb brightening (nonrotating haze), diurnal cloud, recurrent cloud
(diurnal orographic), seasonal cloud (stable topographic), and white area
(frost or fog). This classification of clouds is primarily geometric, with only
minor reference to possible causes or physical mechanisms, and at the current
state of knowledge, it is more valuable than a speculative attempt at physical
discrimination. Table 2 from C. F. and V. W. Capen (1971) is an indication
of the fraction of observing nights on which the listed phenomena were observed
somewhere on the planet during 1966-1968. Such a list is necessarily biased
by the superior observations possible during the particular season (northern
midsummer), when Mars was nearest Earth, and is hurt statistically by the
single observing station, but it gains from uniformity, and does give some
useful measure of just how common these phenomena may be.
Table 2. Martian clouds and hazes 1966-68
after C. F. and V. W. Capen (1971)
Classification
% observed
of nights
viewed
Arctic hood or haze (latitudes +65° to +90°)
76. 3
Northern hemisphere clouds and haze (latitudes 0°
to +65°)
98.
Southern hemisphere clouds and haze (latitudes 0°
to 90°)
63.5
Antarctic hood or haze (latitude -65° to -9 0°)
39. 3
Morning cloud and haze
53. 1
Afternoon cloud and haze
87. 2
Cloud band
16.
Recurrent cloud
35. 5
Terminator projection
10.
White area (frost/fog)
68. 3
December 15, 1971
R. Newburn, JPL
Sec. 4. 1, page 7
Clouds and Hazes JPL 606-1
The famous "W" cloud of Mars seems to be a peculiar recurring white
i'loud. It was first observed and photographed by Slipher in 1907 and has be;.n
seen during several additional apparitions of Mars since then (Slipher, 1962),
It was particularly prominent and received wide public notice in 1954. The
W-shaped cloud always appears in the same place, the Tharsis region near
Lacus Phoenicis, and "the main stems of the cloud pattern appear to coincide
with the main (so-called) canals in the area" (Slipher, 1962). The "W" cloud
was observed by Mariners 6 and 7 during the far encounter sequences (see
Leovy et al. , 1971; also 7F74 in Section 3. 6). It must be noted that the cloud
appears as a "W" in astronomical orientation. With north "at the top" it is an
"M" cloud.
Although most "moving" clouds are of the yellow variety (see paragrapT^f
following), some white clouds appear to show motion, both in the sense of
growth and in apparent motion from day to day (Martin and Baum, 1969). This
does not necessarily imply real motion. If the white clouds are the result of
condensables, the apparent ixiotion may only represent a propagating change in
atmospheric conditions.
At least some white "clouds" are almost certainly made up of CO2 and/or
H2O ice particles. The conditions and mechanisms by which such clouds, fogs,
or frost can form are considered in some detail in Section 3. 4. The composi-
tion and formation dissolution of the polar caps are also discussed in Sec-
tion 3.4. The polar hoods and their relationship to the polar caps are discus-
sed in Section 4. 2.
Yellow Clouds
Yellow clouds usually appear yellow or orange when observed visually.
They are easily photographed in yellow or red light, but cannot be seen on
blue photographs. Their disappearance in the blue may be caused by lack of
contrast and by their being lower in the atmosphere than white and blue clouds,
essentially the same reasons for the disappearance of surface detail in blue
photographs. Even in red light, yellow clouds may be difficult to see because
of the lack of contrast when above a Martian bright area, althoiagh when first
forniing they are often very bright. The existence of extended elements of a
yellovv' clond can sometimes be detected polarimetrically when they cannot be
seen visually.
It is almost universally accepted at the present time that yellow clouds
nre dust clouds. -i^Polarimetric and thermal studies of the Martian surface
indicate the virtually certain existence of finely divided material (see Sec
tion 3. 1) The color of the clouds ctjrrelates with what wotild be expected for
disperses =iurface material. Polarimctric studies of the c^vids themscl\""-
!-As final prepa"^ati j.: v/as under way on this docunient. Mariner " -v ■ -' '••-
orbit around Mars 'luring * he biggest "dust storm" ever observc-J or ^"'--^
planet. Early ne-vs releases from the Mariner 9 project seem to indicate
that considerable evidence has been obtained confirming that yellow clouds
are made up of airborne surface particulates (dust), but quantitative
scientific data is not available at this writing.
Sec. 4.1, page 8 R. Newburn, JPL December 15, 1971
JPL 606- 1 Clouds and Hazes
do not disagree with the dust hypothesis (Dollfus, 1961a). Yet, this evidence
is more circiimstantial than direct, for example Perls (1971) offered an
alternative carbon suboxide yellov/ cloud hypothesis.
Yellow clouds have a large range of sizes and shades of color. The
most obvious are the great yellow "storms, " which may grow in size until
they nearly cover the planet. The great dust storm of 1956 began with local
activity on August 20 (Earth date) and greatly intensified on August 28 with
a "brilliant, orange-colored cloud" over Noachis, according to Slipher (1962).
Kuiper (1957) first saw a group of "five or six bright yellow clouds" over
Mare Sirenum on August 29- The differences in descriptions of the cloiul
development are probably due, at least in part, to the fact that Slipher was
observing in Bloemfontein, South Africa, and Kuiper in west Texas; they
were thus seeing somewhat different parts of the planet. However, both agree
that by September 2, almost the entire planet was covered by a yellow veil
which partially obscured most surface detail. Actually, a part of the surface
was still unobscured as seen from Australia on September 3 (de Vaucouleurs.
1957). There were local areas of greater opacity which completely obscured
the surface. This activity reached a maximum on September 7, according
to Slipher (1962), and continued to some extent until at least September 22.
A third description of the 1956 storm, the greatest ever obsor^ *^d on Mars
until that of 1971, is given by Dollfus (1961b).
The yellow storm of 1971 first appeared, on photographs taken Septem-
ber 22, as a bright bar, some 2400 km long, lying across Hellespontus and
Noachis in a northeast-southwest direction (Capen and Martin, 1971). Thcio
was no evidence of "activity" on photographs taken the previous day. JJurii
the succeeding 16 days, the westward moving front of the storm traveled
completely around the planet to link with the more slowly expanding eastern
front, at a longitude of about 240° (Capen and Martin, 1971). Figures 1 and 2
from the Planetary Research Center of Lowell Observatory show the day-to-day
progress of the storm during this period. By the twentieth day, October 12,
the entire planet was obscured, showing no surface detail on red images.
After 60 days, the surface of Mars was still obscured, making the 1971 storm
the greatest observed in duration, as well as area covered. Information on
other major yellow storms is given in Table 3.
Small yellow clouds have been observed at most favorable apparitions
of Mars, but they are not common. Some yellow clouds appear to move and
have been followed for as long as 16 days (Gifford, 1964). In a study of old
observations, Gifford (1964) found several yellow clouds which see^ned to move
with velocities in excess of 100 km/hr. The inean rate of advance of the
1971 storm appeared to be about 40 km/hr for a 16 -day period (Capen ind
Martin, 1971). A detailed, uniform study of 95 discrete clouds o.'' pH lypes by
the Planetary Research Center revealed 43 clouds which appeared to n^ove
(Martin and Baum, 1969). Thirty-five of these, those for v.-hich observatic f.
spanned more than 72 hours, showed a mean velocity of oi:'-' '', .( ]-:'rJ\\v cTi^
a nnaximum velocity of 15 km/hr (Martin and Bauiri, 196").
It is clear that the expansion of the cloud in a great storm could be
caused by an expansion of the disturbance, as well as by motion of the actual
cloud material. A small yellow cloud presumably cannot disappear too rapidly,
December 15, 1971 R. Newburn, JPL Sec. 4.1, page 9
Clouds and Hazes
JPL 606-1
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Clouds and Hazes
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Clouds and Hazes
JPL 606-1
Table 3. Table of major Martian "dust storms,'
Date
September - December 1877
July 1909
August 1909
August 15 -October 26, 1909
September - November 1911
November 3 -
December 23, 1911
August -
mid October 1924
Early December 1924-
mid February 1925
August - November 1926
Late May - July 1939
August - September 1939
November 12-28, 1941
August - November 1941
August 19 -
September 22, 1956
August 19 -November 1956
September - November 1958
September 22, 1971 -
January 1972
Description
Schiaparelli observed vast clouds totally obscuring
the "equatorial continent" between Syrtis Major and
Ganges. On around the equatorial zone north of
Mare Sirenum and Mare Cimmerium everything
was veiled but not totally obscured.
A large part of the visible disk was covered by a
yellow veil (grayish in some areas).
Considered here to be pre-storm veiling.
Major yellow storm — nearly planet-wide —
observed and sketched by Antoniadi (also by
G, Fournier):
By August 12, Mars had turned lemon yellow,
where one could hardly distinguish markings
normally as dark as Mare Tyrrhenum, Syrtis
Major, and Sinus Sabaeus. Only Mare Sirenum
maintained in August its usual intensity.
Major yellow storm seen on Lowell Obs. plates.
Major yellow storm seen on Lowell Obs. plates.
An orange yellow veil covered the entire region
south of Syrtis Major and Sinus Sabaeus as far
west as Thaumasia, and persisted for weeks.
Also observed and sketched by Antoniadi.
Major yellow storm seen on Lowell and Lick Obs.
plates.
Major yellow storm observed and sketched by
Antoniadi, and measured polarimetrically by Lyot.
Polarization began to decrease about Dec. 1, and
by mid-month surface detail was gone, hidden by
a thick yellowish veil. Both came back to normal
during February.
Major yellow storm seen on Lowell Obs. plates.
Total obscuration of the Utopia - Umbra region by
a very persistent large yellow cloud.
Late July, a general abnormal paleness of the
disc was noted. Solis Lacus - Bosporos very pale.
Questionable major storm. Little data on Lowell
Obs. plates.
A large cloud grew over Libya on Nov. 12 and
began moving south. On Nov. 15 Slipher saw it as
a vast system of several clouds more than 1000 km
across. It was last seen over Phaethontis on
Nov. 28. de Vaucouleurs describes it as yellowish
or pinkish except on Nov. 13, when it appeared
bluish-white.
Major yellow storm seen on Lowell Obs. plates.
Details are given in the main body of the text.
Major yellow storm seen on Lowell Obs. plates.
Major yellow storm seen on Lowell Obs. plates.
Major yellow storm - planet-wide. Details of
developing stages are given in the main body of
the text.
Source
Maggini (1939), p. 323-325
Maggini (1939), p. 318-320
Antoniadi (1930), p. 40-41
Capen and Martin (1972a)
Capen and Martin (1972a)
Maggini (1939), p. 321-322
Antoniadi (1930), p. 41-42
Capen and Martin (1972a)
Antoniadi (1930), p. 42-43
de Vaucouleurs (1954),
Plate Vin
Capen and Martin (1972a)
de Vaucouleurs (1954),
p. 88 (see ref. given)
Capen and Martin (1972a)
de Vaucouleurs (1954),
p. 339-341; Slipher (1962),
p. 108-109
Capen and Martin (1972a)
Slipher (1962): Kuiper (1957);
de Vaucouleurs (1957)
Capen and Martin (1972a)
Capen and Martin (1972a)
Capen and Martin (1972a, b)
Capen and Martin (1971)
Refs: -Antoniadi, E.-M., 1930, La plan^te Mars, 1659-1929; Hermann et Cie., Paris 239 p. ....j,,.
-Capen, C. F. and Martin. L. J. , 1972a, Photographic survey of Martian yellow storms, to be published.
(Abstract in Bull. Am. Astronom. Soc, v. 4, no. 3, pt. II) ^ , t- ,
-Capen, C. F. and Martin, L. J., 1972b, Mars' great storm of 1^71: to appear in Sky and Telescope,
V. 43, no. 5, May.
Sec, 4. 1, page 12
R. Newburn, C. Michaux, JPL
February 15, 1972
JPL 606-1 Clouds and Hazes
however, as some time is required for the particles to settle out of the
atmosphere. Therefore, apparent motion of such a discrete object is likely
to be real motion. Martin and Baum (1969) suggest the high apparent veloc-
ities noted by Gifford (1964) may be caused by errors in thf> oriqinal oKt^r-
vations reported in the old literature. There is no obvious objective way to
handle such nonphotographic data.
The Planetary Research Center work indicated that clouds pr>rerally
tend to avoid the darker areas on Mars and are particularly prevalent in
the zone from 20°N latitude to the equator (Martin and Baum, 1969). These
workers also found the majority of clouds to miove east and Vv-e^-t, espet >:'lly
to the east, rather than to the north or south. Their charts and those of
others show a strong preference for cloud activity in certain discrete geo-
graphic locations. Capen and Martin (1971) note the strong preference for
the formation of great yellow clouds in the Hellas-Noachis region (1909, 191 1,
1924, 1939, 1956, and 1971), for example.
It is difficult to believe that volcanic activity or meteoritic impacts
could be responsible for the great yellow storms of 1956 and 1971. Therefore,
it is generally assumed that the dust particles are raised by suitably high
winds. (It must be recognized that miore than one mechanism actually could
be active. ) Very high wind velocities are required to initiate movement of
particles of any size, large particles (-1/2 mm) being easier for wind to move
than small ones (Ryan, 1964). Even particles of optimum size require near-
surface winds (at 2-m elevation) of 450-500 km hr"! to miove them in a 5-mb
atmosphere (Ryan, 1969).
Neubauer (1966) showed that large dust-devils offered an attractive
mechanism for moving dust aloft. Extensive studies by Ryan (1969) indicated
that velocities adequate to entrain dust should be available in the form of
"vortex systems of the dust devil type. " These systems originate in a shallow,
superadiabatic layer near the surface. Conditions on Mars are such that they
should originate rather more frequently and with greater diameter (and lifting
ability) than on Earth (Ryan, 1969). Once large particles begin to move, they
collide with small particles, knocking them into the air. Dust-devils have a
good upward component of velocity (Neubauer, 1966). Larger-scale winds may
also have a vertical component, perhaps as great as 6 m sec"^ (Gierasch and
Goody, 1968). There is good reason to believe that dust can be raised from
the surface, at least to the tropopause, to form yellow clouds.
The altitudes of yellow clouds have often been estimated, but the results
vary wildly. They are "mostly low level objects, lying generally betv/een 3
and 5 miles above the surface, " according to de Vaucouleurs (1954). Slipher
(1962) states that "most of the examples best observed and most susceptible
of accurate measurement have been found at heights of 1 8 to 20 miles. " Alti-
ttide "measurements" have generally been made on terminator clouds (clouds
that remain illuminated on the dark side of the terminator because of their
height above the ground) by assuming the apparent horizontal distance between
cloud and terminator is a direct function of altitude. However, statistical
errors are quite large for the small angles involved (<0"1), and there are
numerous possibilities for systematic error in such measurenients.
December 15, 1971 R. Newburn, JPL Sec. 4,1, page 13
Clouds and Hazes JPL 606-1
Small yellow clouds usually last for no more than two to four days. The
particles involved must be at least 1-2 |a, in diameter (larger than the wave-
lengths of yellow and red light in which they are easily photographed). They
very probably never rise higher than the tropopause. In fact, particles must
be at least 20 |j. in diameter and rise no more than ~5 km, if they are to settle
out of the atmosphere in four days. Therefore, yellow clouds must include
many particles 20-50 |i in diameter. The yellow veils lasting a month or more
can easily be accounted for by including smaller particles or by maintaining
the lifting mechanism.
In summary, it is certainly possible, based upon all direct and indirect
observational evidence, that the yellow clouds are dust clouds. It is, in fact,
probable that the yellow clouds are dust clouds, but proof of this theory must
await the detailed evidence from Mariner 9 or future studies.
Gray Clouds and Bright Spots
As is quite obvious, Martian meteorological knowledge is largely
qualitative, and at times geological changes in the surface even may be confused
with atmospheric phenomena. Identical events may be described quite differ-
ently by different observers or differing events described similarly. It is this
lack of precision that makes it difficult to evaluate reports of potentially great
importance, such as those that follow. These reports are a few among many.
Ley (1963) has summarized reports of strange dark-gray clouds, made
by several of the best known Japanese observers of Mars. Four of these
clouds were seen during 1950 and 1952, in an area about 500 mi across,
extending from Mare Sirenum across Electris to Eridania. White clouds often
appear somewhat gray, but the experienced Japanese observers felt those
reported were unique both in color and in the great height to which they appeared
to rise. They have chosen to attribute the gray clouds to volcanic activity.
Such a hypothesis is possible, of course, but can hardly be supported from the
observations alone.
Over the years, there have been a number of reports of bright spots
or flares on Mars, typically lasting five or ten minutes. Some of these reports
have been made by reputable, experienced observers. At least two Martian
sites have reportedly exhibited repeated flares, again resulting in hypotheses
of volcanic activity. A brief list of flare observations and characteristics is
given in Table 4. Whether these flares are truly the result of volcanic activity,
even if volcanos exist, or whether they might be specular reflection from a
temporary patch of surface ice or some other surface phenomena, is pure
speculation.
Sec. 4. 1, page 14 R. Newburn, JPL December 15, 1971
JPL 606-1
Clouds and Hazes
Table 4. Table of bright flares and spots observed on Mars
(from Katterfeld, 1966).
InslA-ument
Und
Observatory
Location of
flare or spot
Charatteristics and duration
Jvin 4,
193 7
Dec 8,
1951
Jul 1,
1954
Jul 24.
1954
Nov 5.
1958
Nov 6,
19 58
Nov 10,
1958
Nov Zl,
1958
Nov 21,
1958
Sitzuo
Mayeda
Tsuneo
Saheki
8-in. reflector
Tsuneo
Saheki
Clark
McClelland
Sadao
Mu rayama
Sigeji
Tanabe
Sanenobu
F\ikui
Tsuneo
Saheki
Ilsiro
Tasaka
8-in. reflector,
planetarium at
Osaka, Japan
8-in. reflector,
planetarium at
Osaka, Japan
13-in. refractor,
Allegheny Obs.
of Pittsburgh
University, Penn.
20-cm refractor,
National Science
Museum in
Tokyo, Japan
8-cm reflector,
Sitsuoke, Japan
25-cm reflector,
Kobe, Japan
32. 5-cm
reflector,
V a Ic a y a n 1 a , Japan
Close to
Sithoniua Lacus
+55* lat. ,
240" long.
Western portion
of equatorial
Tithonius Lacus
Edoni
Promontorium
{at the equator )
Edom
Promontorium
South of Tanais
Plateau,
southwestern
edge of
Aphrodite Mare
(Acidalia);
+ 35" lat, ,
42 ° long.
Southern edge of
Tithonius Lacus
Northeastern
part of
Soils Lacus
Northern edge
of Hellas and
Edom
Promontorium
Northern edge
of Hellas and
Edom
Promontorium
Considerably brighter than the polar
cap and the white clouds. Flickered
like a star, and after 5 min it was
hidden from view (possibly due to
rotation of the planet). (See Saheki,
1955. )
Brighter than the north polar cap.
Flickering light and stellar bright-
ness of the 6th niagnitude for
5 min. It then began to be extin-
guished and changed into a grayish
cloud having a diameter of more
than 300 km. The entire phenom-
enon lasted about 40 min. (See
Saheki, 1955. )
In 1 sec the color changed from
a whitish-yellow to a bright, pure
white, and then changed to yellowish-
white. Duration of the flare was
5 sec. (See Saheki, 1955. )
Flare was visible for about 58 sec.
In the opinion of the observer, it
was caused by a volcanic eruption.
Small but very briglU ypot, while in
color. Lasted about 5 min. From
Jul 23 to Aug 3, 1954, a similar
white spot was observed at the sarrie
location by Tsuneo Saheki in the
form of a very l>right, small cloud.
Brightness same as for the polar
cap for 4 min.
Brightness same as for the polar
cap for 5 min. Diameter of spot
(according to a figure) aliout
25 km.
Two bright spots. Visibility belov
5 according to the standard .scale.
The same spots as al)o\i', hut
visibility 6 + 7. Yellowish-white cloud
over northern part of Ilellas. Both
flares lasted about 5 min, togolher
with phases of inirea.se and decrease
in ijrightness - - 1 5 niiti. After several
minutes, the flares i-e.ipijcn red.
December 15, 1971
R. Newburn, JPL
Sec. 4. 1, page 15
Clouds and Hazes JPL 606-1
BIBLIOGRAPHY
Barth, C.A. and Hord, C.W. , 1971, Mariner ultraviolet spectrometer; topog-
raphy and polar cap: Science, v. 173, p. 197-201.
Baum, W.A. , et al. , 1967, Mars cloud survey report no. 1: Flagstaff, Ariz. ,
Lowell Obs. , Planet. Res. Center.
Boyce, P.B. and Thompson, D.T. , 1971, A new look at the Martian "violet haze"
problem I. Syrtis Major- Arabia, 1969: in press.
Capen, C.F. , 1966, The Mars 1964- 1965 apparition: Pasadena, Calif. , Jet
Propulsion Laboratory, Tech. Rep. 32-990.
Capen, C. F. , 1970, Martian blue-clearing during 1967 apparition: Icarus, v. 12,
p. 118-127.
Capen, C. F. , 1971, Martian yellow clouds-past and future: Sky and Telescope,
V.41, no. 2, p. 2-4.
Capen, C.F. and Capen, V.W. , 1971, Martian meteorological phenomena:
to be published.
Capen, C.F. and Martin, L.J. , 1971, The developing stages of the Martian yel-
low storm of 1971: Lowell Obs. Bull. No. 157, v. HI, p. 211-216.
de Vaucouleurs, G. , 1954, Physics of the planet Mars: London, Faber and
Faber.
de Vaucouleurs, G. , 1957, Photographic observations in 1956 of the blue clearing
on Mars: Pub. Astron.Soc. Pacific, v. 69, p. 530-532.
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Bedford, Mass. , Air Force Cambridge Research Laboratories, Contract
AF-61(052)-508, final report.
Dollfus, A. , 1961a, Polarization studies of the planets, chapter 9 in Planets
and satellites, v. Ill of The solar system; Kuiper, G.P. and Middlehurst,
B.M. , Editors : Chicago, U. of Chicago Press.
Dollfus, A., 1961b, Visual and photographic studies of the planets at the Pic du
Middi, chapter 15 in Planets and satellites, v. Ill of The solar system;
Kuiper, G.P. and Middlehurst, B. M. , Editors : Chicago, U. of Chicago
Press.
Evans, D.C., 1965, Ultraviolet reflectivity of Mar s: Science, v. 149, p. 969-972.
Gehrels, T. and Teska, T.M. , 1962, The wavelength dependence of polariza-
tion: Comm. Lunar Planet. Lab., v.l, no. 22, p. 167-177.
Gierasch, P. and Goody, R.M. , 1968, A study of the thermal and dynamical
structvire of the Martian lower atmosphere: Planet. Space Sci. , v. 16,
no. 5, p. 615-646.
Sec. 4.1, page 16 R. Newburn, JPL December 15, 1971
JPL, 606-1 Clouds and Hazes
Gifford, F. , 1964, A study of Martian yellow clouds that display movement:
Mon. Weather Rev. , v. 92, p. 435-440.
Humason, M. L. , 1961, Photographs of the planets with the 200-inch telescope,
chapter 16 in Planets and satellites, v. Ill of The solar system;
Kuiper, G.P. and Middlehurst, B.M. , Editors : Chicago, U. of Chicago
Press.
Katterfeld, G. N. , 1966, Volcanic activity on Mars: Wash. , D.C. , National
Aeronautics and Space Administration, Tech, Trans. F-410, Translation
of, 1965, Vulkanicheskaya aktivnost' na Marse: Priroda, no, 8,
p. 103-109.
Kuiper, G.P. , 1957, Visual observations of Mars, 1956: Astrophys.J, , v,125,
p. 307-317.
Kuiper, G. P. , 1964, Infrared spectra of stars and planets, IV, The spectrum
of Mars, 1-2.5 microns, and the structure of its atmosphere: Comm.
Lunar Planet, Lab, , v,2, no. 31, p. 79-112,
Leovy, C.B, , Snnith, B,A. , Young, A, T,, and Leighton, R,B. , 1971, Mariner
Mars 1969: atmospheric results: J.Geophys,Res. , v,76, p. 297-312.
Ley, W, , 1963, Watchers of the skies: New York, Viking Press,
Lowell, P. , 1905, The canals of Mars — photographed: Lowell Obs,Bull, , v,l,
no. 21, p, 134-135.
Lowell, P. , 1907, On a new means of sharpening celestial photographic images;
and applied with success to Mars: Lowell Obs,Bull. , v.l, no, 31,
p,183-185.
Maggini, M. , 1939, II pianeta Marte: Milan, Italy, Ulrico Hoepli, Editore-
Libraio della Real Casa (400 p. ).
Martin, L.J. and Baum, W.A. ,. 1969, A study of cloud motions on Mars, Lowell
Obs. , Planetary Research Center, Final Report, Part B, August,
Mees, C.E.K. , 1954, The theory of the photographic process. Rev. Edition:
New York, Macmillian Co,
McCord, T,B. and Westphal, J.A. , 1971, Mars; narrow-band photometry, from
0,3 to 2,5 microns, of surface regions during the 1969 apparition:
Astrophys.J., v. 168, p. 141-153,
Neubauer, F,M, , 1966, Thermal convection in the Martian atmosphere:
J.Geophys.Res. , v,71, p. 2419-2426.
Perls, T,A, , 1971, Carbon suboxide on Mars; a working hypothesis: Icarus,
V.14, p. 252-264,
Pettit, E, and Richardson, R.S, , 1955, Observations of Mars made at
Mt. Wilson in 1954: Pub. Astron.Soc. Pacific, v. 67, p, 62-73.
December 15, 1971 R. Newburn, JPL Sec. 4.1, page 17
Clouds and Hazes JPL 606-1
Pollack, J.B. and Sagan, C. , 1967, An analysis of Martian photometry and
polarimetry: Wash, , D.C. , Smithsonian Inst. Astrophys.Obs, ,
Spec. Rep. Z58.
Richardson, R.S. and Roques, P.E, , 1959, An example of the blue clearing
observed 74 days before opposition: Pub. Astron.Soc. Pacific, v,71,
p.321-323.
Ryan, J. A. , 1964, Notes on the Martian yellow clouds: J.Geophys.Res. , v,69,
p. 3759-3770.
Ryan, J. A. , 1969, Study of dust devils as related to the Martian yellow clouds:
McDonnell Douglas Astronautics Co. , Rpt, DAC-63098, January.
Saheki, T. , 1955, Martian phenomena suggesting volcanic activity: Sky and
Telescope, v. 14, no. 4, p. 144-146.
Slipher.E. C. , 1937, An outstanding atmospheric phenomenon on Mars:
Pub. Astron.Soc. Pacific, v,49, p. 137-140,
Slipher, E. C. , 1962, Mars, the photographic story: Cambridge, Mass. , Sky
Pub. Corp, , and Flagstaff, Ariz, , Northland Press.
Smith, B.A. , 1961, Blue clearing during the 1960-61 Mars apparition: Pub.
Astron.Soc. Pacific, v. 73, p, 456-459.
Van Blerkom, D.J. , 1971, The effect of haze on the visibility of Martian surface
features: Icarus, v. 14, p. 235-244.
Wilson, A.G. , 1958, Spectroscopic observations of the blue haze in the atmo-
sphere of Mars: Santa Monica, Calif. , RAND Corp. , Rep. P-1509.
Young, A. T. , 1969, High-resolution photometry of a thin planetary atmosphere:
Icarus, v.ll, p. 1-23.
Sec. 4.1, page 18 R. Newburn, JPL December 15, 1971
JPL 606-1
Seasonal Activity
4.2 SEASONAL ACTIVITY
DATA SUMMARY
Th e Polar Caps and Hoods
The Martian polar caps appear to be deposits of some solid substance,
most probably CO2 (and small amounts of H2O), condensing during the fall and
winter in each hemisphere and then subliming during the spring and summer.
The polar hoods are white clouds which hide the polar regions when photographed
in blue or violet light during the fall and winter. Typical time behavior of these
features is shown below. Lines indicate typical periods when these features
are visible. Lengths of seasons are given in Earth days. (The symbols Ls and
r| are defined in the Appendix. )
(n=85°)
Ls = 0'
Vernal*
Equinox
(11=175°)
Ls=90°
Summer =!
Solstice
(^=265°)
Ls=180»
Autumnal''
Equinox
('1=355°)
Ls=270°
Winter*
Solstice
Northern
Hemi-
sphere
Length
Southern
Hemi-
sphere
Spring
Summer
Fall
Winter
-i
1
--
199 days
183 days
147 days
158 days
Fall
Winter
Spring
Summer
('1=85°)
Ls=360°
or 0°
Vernal*
Equinox
Cap
Hood
Hood
Cap
Seasonal Behavior of Clouds and Whitenings
See Figs. 8 through 10.
The Wave of Darkening
The wave of darkening is "a progressive albedo decline of the Martian
dark areas starting in local springtime from the edge of the vaporizing polar
ice cap and moving towards and across the equator" (Sagan and Haughey, 1966).
Its reality as a wave recently has been seriously questioned.
*Strictly for the Northern Hemisphere. Adoption of the same convention for
designating the equinoctial and solstitial orbital points as used for Earth in
astronomy.
February 1, 1972
C. Michaux, R, Newburn, JPL
Sec. 4. 2, page 1
Seasonal Activity JPL 606-1
Seasonal Behavior of Surface Features
Besides intensity (albedo) changes, dark areas show change in color,
shape, size, and internal appearance, while light areas also change in color
and structure. These changes are best shown in detailed descriptions of the
changes occurring in individual areas (Tables 4-7) and in a series of color
maps (Figs. 8-10).
INTRODUCTION
The first known drawings of Mars showing features which can definitely
be identified are Christiaan Huygens' maps of November 28, 1659, clearly
showing Syrtis Major, and of August 13, 1672, showing the south polar cap
(Ley, 1963). J. D. Cassini saw both caps in I666. The first astronomer to
realize that both the polar "white spots" and the equatorial dark areas changed
in appearance from opposition to opposition was Giacomo Filippo Maraldi, who
observed every opposition of Mars from 1672 until at least 1719 (Ley, 1963).
For more than 250 years, then, it has been realized that Mars is a changing
world, and famous planetary astronomers such as Herschel, Schroeter, Beer
and Von Madler, Secchi, Lockyer, Kaiser, Dawes, and Proctor studied this
planet, which appeared more Earthlike than any other. The modern period
of good maps and really useful records of surface changes began with
Schiaparelli's work during the excellent opposition of 1877.
Mars undergoes many sorts of change. There are the apparent changes
as seen from Earth caused by the rotations of Mars and Earth and by the
tremendous variation in the distance separating the two planets. Because the
axes of Earth and Mars point in different directions, the sub-Earth point on
Mars changes through almost 50° of areographic latitude, causing a great
change in perspective. During this time there are also more subtle changes in
appearance caused by changes in the photometric coordinate (see Section 3. 2).
There are changes caused by the appearance and disappearance of various
meteorological phenomena such as are discussed in Section 4. 1. Some of these
phenomena are not completely random, but appear to be somewhat a function of
season. Changes in polar caps and in the appearance and photometric
properties of Martian dark areas are obvious seasonal changes and the primary
subject of this section.
Since 1969, the many changes occurring on Mars, both in its atmosphere
and on its surface, are monitored photographically, approximately hourly,
through each apparition, by an International Planetary Patrol (IPP) network of
observatories distributed around the world. Each patrol station is furnished
with identical 35 -mm cameras and filters (red, green, and blue), and returns
its (fourteen-exposures) filmstrips to the coordinating Planetary Research
Center at Lowell Observatory for development, editing, mounting, cataloguing,
copying, etc. An almost continuous and homogeneous coverage of the planet
during its apparition has thus become available. This data permits detailed
studies of the various seasonal and secular changes on Mars. Many thousands
of photographs of Mars have already been obtained under this program, more
than doubling the pre-program Lowell Observatory collection. A description
of this program has been given by Baum et al. (1970).
Sec. 4.2, page 2 C. Michaux, R. Newburn, JPL February 1, 1972
JPL 606-1 Seasonal Activity
The rare flyby missions, such as Mariner 1964 and Mariner 1969,
obviously cannot monitor the planet's complex seasonal evolution in just a few
days period. They can only photograph the planet at a particular seasonal time.
The encounter time of Mariner 1969 with Mars took place during its Northern
Hemisphere's early autumn, and Southern Hemisphere's early spring (L^ ~ 200°).
Orbital missions, on the other hand, such as Mariner 1971, are designed
to make a detailed investigation and lengthy surveillance of the planet's fixed
and variable features and study time -varying phenomena of both the surface and
atmosphere. Specific variable phenomena to be studied closely by Mariner 9
include: (1) wave of darkening, (2) polar caps phenomena, (3) non-polar white
clouds and whitenings, (4) yellow clouds and dust storms, and (5) hazes (for a
description of the Mariner 1971 TV Experiment, see Mazursky et al. , 1970).
The seasonal activity is to be studied over a (terrestrial) year or more,
depending upon the spacecraft operating lifetime. It is hoped, therefore, that
at least half a Martian year of seasonal activity will be recorded by its TV
cameras and other instruments.
POLAR CAPS AND HOODS
Polar Caps
The first recognition (in 1784) that the variation in the "white polar
spots" of Mars is seasonal, and the suggestion that the spots are true polar
caps of snow and ice, were due to William Herschel (Ley, 1963; Slipher, 1962).
Measurements of the waxing and waning of the caps, begun by Herschel, have
shown little change over a period approaching 200 years. Figure 1, compiled
by Slipher (1962), shows the regression of the south cap as observed through
many seasons.
There is considerable difference in the north and south polar caps of
Mars caused by the asymmetry in Martian seasons. The south cap is formed
during the long, 382 (terrestrial) days of southern fall and winter when Mars is
near aphelion, and it covers more than 70 areocentric degrees at greatest
extent (extending below latitude -5 5°). The north cap is formed during the
short, warmer 305 (terrestrial) days of northern fall and winter when Mars is
nea"r perihelion, and it usually measures only about 53° at maximum extent
(Slipher, 1962).
Quantitative photometric and polarimetric data on the seasonal evolution
of both polar caps and hoods have been given by Focas (1961 and 1962). See
Fig. 2.
The polar caps (and hoods) of Mars are difficult to observe at the
telescope from Earth because they are viewed from the side, with highly
oblique viewing angles, especially near the poles. Because of the orbital
geometry and inclination of the rotational axis with respect to Earth, it is
presently the larger springtime South Cap which is closer and more favorably
tilted for viewing during perihelic oppositions, while the smaller springtime
February 1, 1972 C. Michaux, R. Newburn, JPL Sec. 4.2, page 3
Seasonal Activity
JPL 606-1
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87 67 47 27
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DAYS COUNTED FROM THE SUMMER SOLSTICE
Fig. 1. Measured width of the south polar cap of Mars for various
oppositions from 1798 to 1924. This figure embraces the Martian
season from before the vernal equinox to 94 days after the summer
solstice. Other measures derived from drawings in 1781, 1783,
1815, 1830, 1845, and 1862 were checked with those shown here,
but no notable deviations were found other than accidental errors
attributable to optical limits of the observer's telescope. The
plotted measures shown in the figure agree very w^ell indeed, and
the deviations in the measures by the same observer are of about
the same order as those that occur bet'ween different observers.
This study revealed no evidence of any irregularity in the melting
of the south cap at any of these oppositions during this long period
of observation. (Slipher, 1962)
Sec. 4. 2, page 4
R. Newburn, C. Michaux, JPL
February 1, 1972
JPL 606-1
Seasonal Activity
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February 1, 1972
C. Michaiix, JPL
Sec. 4. 2, page 5
Seasonal Activity JPL 606-1
North Cap is best seen during aphelic oppositions. Thus, understandably, more
observations have been made of the South Cap in the past (see the literature
from Schiaparelli to de Vaucouleurs). Some observers (for example
C. Capen), however, have specialized in the North Cap, which is more difficult
to study. There is a compensation, however, for observers situated in our
Northern Hemisphere, for at aphelic opposition times Mars is then at much
higher declination, usually resulting in better seeing and longer observing
periods.
Polar Hoods
During the fall and winter in either hemisphere of Mars, the polar areas
are covered by a very large hood of clouds, often somewhat dull bluish-white
in color and with an irregular diffuse edge. The early fall north polar hood was
photographed by Mariners 6 and 7 in the FE series (see Section 3. 6). The
clouds usually appear large and bright in violet and blue photographs, although
varying considerably in extent from day to day. Yellow and red pictures may
show a much smaller bright region, where perhaps the forming cap is visible
through the hood, or they may show almost nothing. Sometimes in late fall-
early winter the hood becomes tenuous in places or may pull back partially,
apparently revealing a portion of the deposited cap. By mid-winter the hood
appears to be very thick and stable. Around vernal equinox (spring) time, the
polar hood breaks up and finally dissipates, unveiling the much smaller true
polar cap which is brilliant white and sharp. (The cap appears sharp in red
light and brilliant in green and blue light. ) Throughout spring (and summer) the
cap remains well-defined, appearing practically constant in size from night to
night, but very gradually regressing. A beautiful series of photographs
illustrating these points has been published by Slipher (1962).
The southern polar hood reaches, in average, a latitude of -40°, spanning
then an equivalent breadth of 100" areocentric. Its outline and seasonal
behavior are quite irregular. The northern hood is not as large, but is
definitely more regular. Usually extending down to slightly below +60°
(equivalent breadth 60° areocentric), it may reach +50° latitude as in 1967
(Capen and Capen, 1971, 1972), or even exceptionally lower to +40° as in I960
(Slipher, 1962).
Photometric and polarimetric data on the seasonal evolution of the polar
hoods are shown in Fig. 2.
The Dark Polar Fringe*
The so-called dark polar fringe, also known as the polar band, or collar,
surrounding the retreating polar caps, is said to have been seen first by Beer
and Von Madler as early as 1830 (Slipher, 1962). The fringe is not seen at the
telescope when a polar cap is at its miaximum extent nor is one visible during
the slow, final demise of a cap (de Vaucouleurs, 1954). During the time of
retreat, however, a zone most often described as 'bluish' develops contiguous
to the cap, reaching its greatest width at the time of maximum rate of vaporiza-
tion (de Vaucouleurs, 1954).
^Recently, Pang & Hord (Icarus, Dec. 1971) interpreted the collar as a photometric
effect due to oblique viewing of glazed CO^-ice layers bared by the retreating cap.
Sec. 4.2, page 6 C. Michaux, R. Newburn, JPL February 1, 1972
JPL 606-1 Seasonal Activity
The dark fringe appears to many observers to be more than just an effect
of contrast between the brilliant cap and its relatively dull surroundings.
Quantitative studies by de Vaucouleurs and by Dollfus identified a contrast
effect, yet seemed to confirm the reality of the fringe (de Vaucouleurs, 1954),
The behavior of the fringe may be identical to that of other areas of the planet
that take part in the wave of darkening.
The dark fringe has sometimes been called the melt zone, implying it to
be an area where liquid water exists for a brief time, wetting the ground before
evaporating. This is most unlikely since sublimation of surface ice will occur
rather than melting on Mars, as shown by Ingersoll(1970). Furthermore,
theoretical polar temperatures will not exceed 0" C.
The existence of the dark fringe has been held in doubt, especially in
recent years, by a number of theoreticians who consider it to be just another
optical-physiological illusion (a contrast effect) produced in our ground-based
telescopic observations. The Mariner 7 flyby over the South Polar Cap could
not resolve this controversial question, since the flyby took place too early in
Southern Spring. The polar dark fringe was first seen two months later at the
telescope (late September 1969). The black border appearing in the NE
Mariner 7 pictures around the brilliant South Cap is only the product of the TV
camera system's "automatic gain control" used to obtain "maximum discrimi-
nability" versions of the pictures, and does not appear in the photometric
versions (see Section 3.6, Photographic Atlas). Thus, the polar dark fringe
is still a questionable Martian phenomenon.
Boundaries of North and South Polar Caps
Fischbacher, Martin, and Baum (1969) made a statistical study of the
boundaries of both caps during their regressional phase on more than 3,000
photographs (yellow and red), obtained between 1905 and 1965. Measurements
of boundary latitude at each of 36 meridians, separated by 10° in areographic
longitude, were determined by superposition of the appropriate coordinate grid
or graticule on each photograph. The Martian year was divided into 36
"seasonal" intervals, each spanning 10° in areocentric longitude of the Sun (Ls).
The mean latitude (combining all apparitions 1905-1965) at each selected
meridian and for each 10" seasonal interval was calculated by computer, and
the results tabulated. Figures 3 and 4 are the final diagrams showing the
successive mean boundaries for each regressing cap.
Conclusions from this statistical study are as follows:
The South Cap recedes more regularly than the North Cap. During the
entire southern spring (Ls: 180-270°), the South Cap boundary was found to be
"unusually well-defined and exceedingly repeatable in its behavior from one
Martian year to another. " Only small but significant differences were detected
for individual meridians. Comparison of the 36 meridian curves with the
average curve for all meridians, again showed a more regular behavior for the
South Cap.
February 1, 1972 C. Michaux, JPL Sec. 4.2, page 7
Seasonal Activity
JPL 606-1
270
to Ls.
Fig. 3. Seasonal boundaries of the South Polar Cap as viewed from the south
(Fischbacher et al. , 1969).
Although the two caps differ substantially in their spring -summer decay
patterns, each cap follows a rather well-defined curve, repeatable from year-
to-year. However, there are slight variations from year to year, as will be
seen later, especially for the North Cap.
Sec. 4. 2, page
C. Michaux, JPL
February 1, 1972
JPL 606-1
Seasonal Activity
270
ers refer fo L5 .
Fig. 4. Seasonal boundaries of the North Polar Cap as viewed from the north
(Fischbacher et al. , 1969).
South Cap Regression
The regular regression of the South Cap is illustrated schematicallv bv
the regression curves of Antoniadi (1930), Slipher (1962, Fig. 1) and
Fischbacher et al. (I969; Fig. 5) which are essentially in good agreement.
The disintegration of the sublimating South Cap is well documented histori-
cally up to modern times and follows the familiar pattern described in the follow-
ing: In early southern spring (Ls = 200-215") the fully expanded South Cap
February 1, 1 972
C. Michaux, JPL
Sec. 4. 2, page 9
Seasonal Activity
JPL 606-1
o
o
uj
O
O
O
I
<
u
2
o
-&■
UJ
Q
90
80
70
AG
50
40
30
SOUTH POIAR HOOD
SOUTH POLAR CAP
• • •
• • •
O
SOUTHERN AUTUMN!
WINTER
SOUTHERN SPRING
SUMMER
60
120 180 240
SEASONAL DATE L,
300
360
NOTE: South Me at oil
oreographic longitudes.
Fig. 5, Regression curve of the South Polar Hood and Cap, mean 1905-1965
(Fischbacher et al., 1969).
spans about 60° areocentric, and begins slowly shrinking. By late spring
(Ls = 240°) its diameter has decreased to about 40° and its regression rate
has started to speed up considerably. Rifts or dark fissures (called Rima,
Depressio, Fretum, . . ) begin to make their appearance, signaling the rapid
breakup phase of the cap, with gradual isolation of large bright promontories or
islands (called Mons) along its edge. Thus, the bright Thyles Mons appears
south of Phaethontis, separated from the main cap body by the often large
Depressio Parva (at longitude 150°-210°W) and the bright double-lobed Argenteus
Mons, south of Argyre I, separated by the narrow rift Rima Angusta (at
30° -50° W). The large bright Novus Mons, otherwise known as the "Mountains
of Mitchel, ■' covering the underlying surface feature Novissima Thyle, becomes
well separated by the large Rima Australis (at 300° -340° W). Later, by early
Summer (Ls = 280°), this Novus Mons becomes completely detached from the
main cap by both Depressio Magna extending eastward the enlarged Rima
Australis, and Rima Brevis. By then, the cap is going through its most rapid
regressional phase; the rifts are very dark, and the promontories and islands
have become very brilliant. (In 1971, de Vaucouleurs found Novus Mons to be
"almost sparkling. ") The controversial "dark polar fringe" is then best seen
at the telescope. Other rifts are also formed; in particular, Ulyxis Fretum
and Rima Brevis which separate bright Thyles Collis, south of Thyle II (at
220° -270° W), on two sides.
Sec. 4. 2, page 10
C. Michaux, JPL
February 1, 1972
JPL 606-1 Seasonal Activity
Around the disintegrating irregular cap, there are a number of darkened
areas: Depressio Hellespontica (near longitude 330° W), Promethei Sinus
(near 260° W), Noti Sinus (near 200° W), Palinuri Fretum (near 150° W),
Tiphys Fretum (near 220° W), and Depressiones Aoniae (near 120° W). All
these areas, as well as the rifts and promontories or islands mentioned, are
easily located on de Vaucouleurs MM'71 Planning Charts of the South Polar
Regions (see Fig. 6), where both mid-spring and mid-summer aspects of the
cap are represented. When mid-summer comes (Ls = 320°), only a roughly
triangular remnant cap, Hypernotius Mons, of about 7° maximum extent is left.
According to de Vaucouleurs (1972): "notwithstanding the frequent published
statements to the contrary, the residual polar cap never evaporates completely
before haze (hood) begins to form over the south polar regions at the end of
summer. " The mean position of the centroid of the small eccentric residual
cap is at 23°W, 84. 7°S (de Vaucouleurs 1972, from a dozen best determinations
at 1877 to 1941 perihelic oppositions).
North Cap Regression
A mean regression curve (period 1905-1965) for the somewhat irregular
North Cap was issued by Fischbacher et al. 1969 (see Fig. 7). This curve, it
should be noted, is not in agreement with the classical curve given previously
by Antoniadi (1930) covering mostly an earlier period (1856-1929). It is to be
considered more reliable than Antoniadi's.
Regression curves for the North Cap were obtained by Capen and Capen
(1970, 1972) for the four successive apparitions in the period 1962-1969, when
it was advantageously tilted towards Earth. See Figs. 12 through 15. These
curves, which are in fair agreement with Fischbacher 's mean curve, show a
generally similar overall behavior in the retreat, but with irregularities.
The general evolution, as depicted by Capen and Capen (1970, 1972, is as
follows: a slow regression in the first part of spring (1° per 10-20 days when
Ls = 0-45°), followed by a rapid regression in the second part of spring
(Ls = 45°-90°), with a temporary halt in late spring (caused by arctic hazes),
resumption to maximunn rate near summer solstice (1° per 3 days), then con-
siderable slowdown by mid-summer (1° per 20-30 days when Ls ~ 135°), and
finally a virtually static cap by late summer (Ls ~ 160°) with a remnant
diameter of about 6°. The evolution of the late 1967 cap was abnormal: no halt
in late spring, slower rate around solstice and in early summer, finally a large
remnant cap of 10°. Evidently, the arctic climate was colder than normal in
1967. Table 1 compares the rates at various seasonal times for the four
apparitions. The halt in the retreat takes place when Mars is close to aphelion
passage ( r] = 155°, or Ls = 70°). It lasts only a few days, but the North Cap
may even increase in diameter by 1 or 2°, as was observed in 1963 and also in
I960 (Miyamoto and Hattori, 1968). The usual remnant cap of 6° diameter agrees
with the average given by Slipher (1962). The centroid of the residual cap is
only about 1° away from the North Pole.
The areal pattern displayed by the regressing North Cap is repeatable in
norinal years. The rapidly disintegrating summer cap (near maximum rate)
produces three prominent white areas at the same locations, which linger
around the main cap edge as bright projections, varying in size and brightness,
and eventually become detached as islands, as the cap further retreats to less
February 1, 1972 C. Michaux, JPL Sec. 4.2, page 11
Seasonal Activity
JPL 606-1
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Sec. 4. 2, page 12
C. Michaux, JPL
February 1, 1972
JPL 606-1
Seasonal Activity
o
o
2
o
Q
3
90
80
70
60
50
40
30
NORTH POLAR CAP
NORTH POLAR HOOD
• • •
• • •
NORTHERN SPRING
SUMMER
ALTTUMN
WINTER
(50
120
]80
SEASONAL DATE L
240
300
360
Fig. 7.
NOTE: North Me at all
creographic longitudes.
Regression curve of the North Polar Cap and Hood, mean 1905-1965
(Fischbacher et al. , 1969).
Tn^r\l° ^'V^*f "^"^"^^^)- J^^^^ ^^Jor remnant white islands are named
after the underlying areas they occupy: lerne (-l-^l' W\ i o^ • ^^^'^
(-200- W) and Cecropia (W- W). ^TH.lnZj ,ilLoVs\fZ-\\:i°Jr''
worth Polar Cap as prepared by Capen (1972) especially for this review.
SEASONAL BEHAVIOR OF CLOUDS AND WHITENINGS
mete
Outside those of the polar hoods, very few seasonal patterns of
orologxcal activity have been established so far. Three addmonal
^n rr^iJt^elra^rr ^'-^'^^'^^^ ^^^^ity which have been gradually defined
1)
2)
3)
Whitenings (surface frost or fog)
Seasonal and recurrent white clouds (orographic afternoon cloud <
and localized persistent clouds) t,ioua.
Great yellow clouds (major dust storms)
February 1, 1972
C. Michaux, JPL
>ec, 4. 2, page 13
Seasonal Activity
JPL 606-1
60° 70°
80°
9(7> 100° 110° 1 20° 130° 140° ISP* 160^
+48°
+51°
+54°
+57°
+60°
+63°
+66°
+69°
+72°
+75°
+78°
+81°
+84°
-87°
+90°
U
O
MARTIAN DATE
Fig. 12. North Polar Cap micrometric regression curve. 1968-69
(Capen and Capen, 197Z).
X
»—
Q
5^
a.
<
O
a.
X
I—
o
z
80'
90° 100° 110° 120° 130° 140° 150° 160°
20 30
MAR
10 20
APR
Fig 13. North Polar Cap photographic regression curve.
(Capen and Capen, 1970).
a:
t—
Z
UJ
U
o
1966-67
Sec. 4. 2, page 14
C. Michaux, JPL
February 1, 1972
JPL 606-1
Seasonal Activity
X
a
u
o
a:
o
z
84»
78°
72'
66'
60«
54'
48"
42"
36"
30"
24'
IS"
12"
6*
0*
cr
10" 20"
30«
40-
50"
60«
70'
80'
1
1 1
1
1
1
1
1
1
- ~
•
•s
i»l
-
>(
\
k
-
N
K»
90°
20
MAR
30
I
10 20
APR
30
10 20
MAY
30
MARTIAN DATE
10 20
JUN
100" 110° 120° 130° 140° 150° 160°
—\ 1 \
1 — r
30
30
+48°
+51°
+54°
+57°
UJ
+60"
a
3
+63°
t—
»—
+66°
<
+69°
y
+72°
7
UJ
+75°
U
O
+78°
Ui
<
+81"
+84°
+87°
+90°
10 20 30
AUG
Fig. 14. North Polar Cap photographic regression curve, 1964-65
(Capen and Capen, 1970).
0° 10° 20° 30° 40° 50° 60° 70° 80° 90°
X
t—
Q
U
a:
<
o
a.
X
I—
a:
o
z
Fig. 15.
10 20 30
MAR
10 . 20
APR
10 20
MAY
MARTIAN DATE
30
10 20
JUNE
<
O
North Polar Cap photographic regression curve, 1962-63
(Capen and Capen, 1970).
February 1, 1972
C. Michaux, JPL
Sec. 4. 2, page 15
Seasonal Activity JPL 606-1
Table 1. North cap regression rates 1962-1969 (areocentric degrees
per Martian days) (Capen and Cave, 1971)
Apparition Early Spring Max. Rate Early Summer Mid -Summer
1962-63 1°/I0d l°/3d l"/8d
1964-65 l"/23d l°/3d l"/6d l°/30d
1966-67 1°/I5d l°/4d 1°/I0d l°/30d
1968-69 l°/20d i=/2d l"/7d l°/20d
Table 2. Mean areographic longitude of centroid of white island remnants
of summer North Polar Cap (1879-1968). (Capen and Capen, 1970)
Observer
lerne
Lemuria
Cecropia
121 °
208°
310°
122°
206°
311 °
122°
208°
309°
136°
213°
278°
142°
227°
292°
140°
196°
290°
G. Schiaparelli, 1879-88
P. Lowell, 1901-5
E. M. Antoniadi, 1903-29
M, Maggini, 1918-35
A, Dollfus, 1946-52
C. and V. Capen, 1962-68
Our knowledge of Martian seasonal meteorological activity is still very
limited for the basic reason that apparitions have generally been covered only
for a few months centered on opposition. A reconstruction from a full 15-17
year cycle of oppositions therefore yields a rather uncertain general picture
over a Martian year. (This is of course true also for surface changes,^ but
it is much more so for the highly variable and often elusive meteorological
changes.) Another important fact, especially pertinent to meteorological
phenomena (since these are best detected and identified with the use of color
filters), is that a systematic colorimetric procedure using a variety of well-_
defined filters did not come into widespread use in planetary observing practice
until recently. (An exception has been the work done at Lowell Observatory
by Slipher who pioneered filter photography of planets about 1905. )
The cloud overlays of Figs. 8 through 10 give areographic locations of the
center of many clouds - white (O) or yellow (O) - reported in the different
seasons.* Since the compilation is by no means exhaustive nor homogeneous
in coverage (through the last 100 years of observational record searched), these
overlays cannot be considered as very representative of the clouds' seasonal
frequencies. Thus, the northern autumn-winter overlay indicates more clouds,
but this is probably biased by the fact that more observations have been made
at the m.ore favorable (near -perihelic) oppositions.
-Diversified sources were used for the cloud overlays, the most important
were: Slipher (in Annals of Lowell Observatory), Antoniadi (1930), Focas
and Dollfus (many reports since 1948), Wells (1966), Capen (many reports
since 1954), de Vaucoaleurs (many reports since 1937), Miyamoto, etc.
Sec. 4.2, page 16 C Michaux, JPL February 1, 1972
JPL 606-1
Seasonal Activity
White Clouds and Hazes
To supplement the dearth of 'continuous' information over one Martian
year, Capen and Capen (1972) undertook an unusually lengthy filter photographic
and visual patrol of Mars for a period of 20 terrestrial months, centered on
the April 15, 1967 opposition. Their I966-I968 Mars apparition covered close
to one full Martian year, from Lg = 5 ° to 326" (only the second part of northern
winter was excepted). They placed particular emphasis on recording the
Northern Mid-Spring (L = 60°)
The map was produced from 1963 observations
made with the JPL Table Mountain Observatory
16-inch Cass, reflector.
Fig. 16. Seasonal evolution of the North Polar Cap (Capen, 1972).
February 1, 1 972
C. Michaux, JPL
Sec. 4. 2, page 17
Seasonal Activity
JPL 606-1
Nor
thern Summer Solstice (L„ = 90°)
The map was produced from 1969 observations.
Fig. 17. Seasonal evolution of the North Polar Cap (Capen, 1972).
Sec. 4. 2, page H
C. Michaux, JPL
February 1, 1972
JPL 606-1
Seasonal Activity
_*
180
Northern Mid- Summer (L = 140°)
The map was produced from 1967 observations.
Fig. 18. Seasonal evolution of the North Polar Cap (Capen, 1972).
February 1, 1972
C, Michaux, JPL
Sec. 4. 2, page 19
Seasonal Activity
JPL 606-1
Northern Late -Summer (L„ = 160°)
The map was produced from observations made
with the JPL Table Mountain Observatory Z4-inch
reflector and the McDonald Observatory 82-inch
reflector. The three North Polar Cap residual
white areas are labeled; I Lemuria (Olympia),
II lerne, and III Cecropia.
Fig. 19. Seasonal evolution of the North Polar Cap (Capen, 1972).
Sec. 4. 2, page 20
C. Michaux, JPL
February 1, 1972
JPL 606-1 Seasonal Activity-
meteorological activity. Table 3 gives their seasonal comparison statistics*
for white clouds and hazes. For nonpolar regions, the results are as follows:
1) The Northern Hemisphere definitely had more clouds and hazes
throughout the year than the Southern Hemisphere, and particularly
in Northern autumn and spring.
Z) The afternoon limb was considerably more covered by clouds and
hazes than the morning limb especially during Northern autumn,
winter and even summer.
3) The recurring (mostly afternoon) clouds, however, started appearing
in Northern late spring and were frequent throughout summer while
decreasing in early autumn.
4) Cloud bands (mostly equatorial) appeared only during Northern
spring and summer.
It should be noted that for this apparition the Northern Hemisphere was best
observed from Earth.
Seasonal and Recurrent White Clouds
There are two types of seasonal white clouds known which recur over the
same areas of Mars from year to year: the "orographic white clouds, " which
recur diurnally in the afternoon at the same spot, and the "localized white
clouds, " which persist for days, shifting in positions around specific areas
>'< >I<
1) Orographic White Clouds (afternoon clouds)
These recurring white clouds form in the early afternoon over
small areas of the surface, and produce patterns, such as the
famed "W" in the Tharsis region, during spring and summer
apparently in both hemispheres. Usually, but not always***, they
are located in the tropics. They grow larger and brighter by
sunset, and as a rule do not reappear at the morning terminator
(one exception has been recorded however, by Capen). They are
described in Section 4. 1 as blue -white clouds, although they some-
times photograph well in yellow-green. The formation of these
clouds appears definitely linked to the topography and special mete-
orological conditions in spring or summer (see Wells, 1967). The
"W" cloud seen by Mariner 1969 was not solid CO2, according
*Percentages given refer to percentages of observing nights (from total),
when the listed meteorological phenomena occurred.
**Capen, in his reports, has called the first 'recurrent or recurring afternoon
clouds" and the second "seasonal clouds in the classical sense, "
***Capen found them restricted to the zone -10° to +40°.
February 1, 1972 C. Michaux, JPL Sec. 4.2, page 21
Seasonal Activity
JPL 606-1
Table 3A. Seasonal meteorological activity, 1966-1970^
(Capen and Capen, 1972)
Spring
Lg 0»-90''
(%)
Summer
Lg 90°-180''
(%)
Autumn
Lg 180''-270°
(%)
Winter
Lg 270°-360°
(%)
Arctic
30.6
78.7
98.0
100.0
Antarctic
72.1
100.0
13.6
10.0*
Northern hemisphere
84.2
90.4
92.8
97.2
Southern hemisphere
50.9
87.5
30.1
80.6*
Morning
77.2
60.3
19.3
33.3
Afternoon
75.4
100.0
74.7
72.2
Cloud band
20.5
30.3
1.2
0.0*
Recurrent cloud
(orogenic)
40.8
57.4
18.0
1.0*
'Data for the 1969-1970 apparitio
more complete coverage.
n were incorporated by the au
thors for
*These percentages are probably
the unpredictable nature and inte
more variable
nsity of yellov^
during this season due to
1 storms.
Table 3B. Relative occurrence of seasonal cloud activity, 1966-1968
(between hemispheres and morning and afternoon).
(Capen and Capen, 1972)
Spring
Lg 0°-90°
(%)
Sunimer
Lg 90°-180°
(%)
Autumn
Lg 180°-270°
(%)
Winter
Lg 270°-360°
(%)
Northern hemisphere
70.8
57.9
79-0
65.3
Southern hemisphere
29.2
42.1
21.0
34.7
Morning
48.2
40.7
19.2
32.6
Afternoon
51.8
59.3
80.8
6r.4
Sec. 4. 2, page 22
C. Michaux, JPL
February 1, 1972
JPL 606-1 Seasonal Activity
to the analysis by Leovy, who went on to speculate that it could be
the condensation product of H2O vapor percolating up from the
ground warmed by the noon insolation maximum. (A quantitative
discussion is found in Leovy et al. , 1 97 1 . )
Besides this conspicuous "W cloud" repeatedly observed (see
Slipher, 1962, for example) in southern spring-summer, a number
of other examples of orographic clouds were reported by Capen
(1966, 1970) who particularly monitored the Northern Hemisphere.
Areas conducive to orographic cloud formation included: Elysium,
Nix Olympica, Ascraeus Lacus, and an area at +40° in Arcadia.
Other such regions are Hellas in the Southern Hemisphere, Edom
and Eden near the equator, and the Candor-Tharsis region where
Capen again detected the "W" formation in northern spring-summer
(thus the "W" group seems to be bi-seasonal).
2) Localized White Clouds
Recurring regionally and seasonally, these dense, whitish clouds of
limited extent, persist for days but with displacement within the
region. White at first (the clouds are best seen however through
blue and blue -green filters) their color changes somewhat, becoming
prominent in green, and sometimes visible in yellow. This led
Capen (1972) to suspect they might be combinations of white and
yellow clouds.
One remarkable example of such localized white clouds was
retraced by Capen (1972) as far back as 1911: the Libya -Crocea-
Oenotria cloud, which recurs every Martian year at about the
time of northern summer solstice and lasts through the first part
of summer (Lg: 88° to 140°), circulating around Syrtis Major
(sometimes even reaching its northern tip). The prominently
dark-blue Syrtis Major changes to light blue, and may even
partially disappear, perhaps the result of an important seasonal
atmospheric change.
Great Yellow Clouds (major 'dust storms ')=!=
While smaller yellow clouds, as uncommon as they usually are, have been
reported in almost any season (see overlays of Figs. 8-10), all of the reported
giant or "great" yellow clouds, obscuring large parts of Mars, have occurred
only during southern late spring and summer (according to preliminary
results of searches of the observational records; see for instance Capen (1971),
or Capen and Martin (1972). These major storms probably recur every
Martian year at about that time. They seem to start suddenly when Mars is
near its perihelion ( q = 335°, or Ls = 250°), as for example, in 1956,
August 19 (Lg = 248°), in 1971, September 22 (Lg = 259°). Once started, these
storms may actively develop, sometimes expanding planet-wide (as in 1971) and
may persist for weeks or even a few months, through the first part of summer,
and finally decay and subside (as fine yellow veils) in an equally long period.
*See Table 3, of Sec. 4.1, p. 12
February!, 1972 C. Michaux, JPL Sec. 4.2, page 23
Seasonal Activity JPL 606-1
These major yellow clouds or storms may be classified in a 'recurrent'
yellow cloud category, regionally as well as seasonally, because the records
show their initial phase appears to always be associated with the same
(Southern Hemisphere) regions: Hellas-Hellespontus-Noachis and vicinity in
particular. (See Capen and Martin, 1971, and 1972.)
The sudden occurrence of major yellow clouds or "dust" storms near the
perihelion time and in certain preferred regions located in the zone of
maximum insolation* on Mars, strongly points to the seasonal buildup of
particularly unstable atmospheric conditions (leading to the onset of strong
advection and convection with associated turbulence, vortices, etc). This
buildup occurs over thermally and perhaps hygrometrically preconditioned
ground, which itself is probably already favored by its topography (location
and relief) as well as its physical-chemical properties**. It appears that dust
is raised suddenly by some triggering meteorological mechanism at a precise
time, and is maintained aloft by a persistent, dynamically exceptional circula-
tion which rapidly enlarges the cloud outburst. Analysis of the 1971 Mariner 9
data may not clarify the exact triggering mechanism, since the global dust
storm had already been in progress for 6 weeks before the arrival of the
spacecraft. (L =290°, Nov. 10, 1971, fir at pre-orbital TV picture s) .
Whitening Areas
Many areas of Mars exhibit a conspicuous whitening which appears to be
seasonal, although there may be a diurnal effect superimposed. The whitening
(as seen at the telescope) may last all day, or may not, depending upon
seasonal time. In the latter case, the white patch is seen to be brightest in the
early morning, as if formed during nighttime, and gradually fades away with
the heating rays of the Sun. Most whitenings occur during one season, lasting
sometimes two successive seasons and a few even three seasons (Hellas).
According to Capen and Capen (1970), white areas are usually more prevalent
at a time when the polar cap of the same hemisphere is rapidly regressing,
that is around summer solstice time (a notable exception is Hellas). Some areas
whiten during two separate seasons: for example, Elysium, Nix Olympica,
Tharsis, etc. Capen (1972), who has devoted much time since 1962 in observing
and recording whitening areas, noted that the second whitening seems to
correlate with the rapid regression of the other polar cap; only a systematic
study could confirm this. So far, only sketchy or incomplete information is
available on the seasonal activity of white areas. The survey by Capen and
Capen (1970) for the 1967 apparition yielded the following percentages of
observing nights when whitenings occurred: spring 24%, summer 66%,
♦ The thermal equator is then at about -24° latitude.
♦♦Comparison of the 1956 and 1971 cases by Capen and Martin 1971, disclosed
that the behavior of the atmosphere prior to their formation was "some-
what alike" (little white cloud activity, and high contrast of surface
features) .
Sec. 4.2, page 24 C Michaux, JPL February 1, 1972
JPL 606-1 Seasonal Activity
autumn 10%, (winter: insufficient data). The observations, both visual and
photographic were made with filters ranging from blue -green to near infrared.
A list of all whitening areas detected and charted by Capen and Capen
(1970) during their 1962 to 1968 observations is contained on Fig. 20 which
locates many of these on Capen's 1967 base map of Mars. The whitenings are
seen to generally affect the light* areas of Mars' surface and are not restricted
by latitude in either hemisphere (while the "afternoon clouds" seem mostly
restricted to tropical latitudes). Thus, they are widely distributed over Mars,
but affect only preferred areas. In these areas, the topographical, thermal,
and meteorological conditions become, at certain times of the year, favorable
to condensation of a common volatile. The whitening formed is apparently
either an actual surface deposit (frost) or a near-surface dense fog (ground
fog), or both.** The condensing volatile is presumably H^O vapor, but this has
yet to be demonstrated. Condensation of CO2 is less likely because of the
extremely low temperatures required (below -148 °K). However, one may
envisage condensation of both H2O and CO2, or only CO2 at the higher latitudes.
THE WAVE OF DARKENING
The wave of darkening has been described as "a progressive albedo
decline of the Martian dark areas (but not the bright areas) starting in local
springtime from the edge of the vaporizing polar ice cap, and moving toward
and across the equator" (Sagan and Haughey, 1966). There is quantitative
evidence that dark areas darken during the Martian spring, reaching maximum
darkness after the summer solstice. Whether the darkening occurs as a
"wave" from the pole has been argued. A statistical analysis by Pollack,
Greenberg, and Sagan (1967) showed that while there are areas which "violate
the concept of an invariable wave, " there is "a very significant correlation of
latitude with time of maximum darkening. "
The waves start alternately from the two polar caps at half-year intervals,
cross the equator, and fade at about 22° latitude in the opposite hemisphere
from which they began (Focas, 1962). The rate of propagation is variable but
averages about 35 km per day. The time from beginning of darkening to
maximum darkening is 0.30 to 0.35 Martian years in the circumpolar and
temperate areas, 0.30 years at the boundary of the equatorial zone, and 0.15
years in the equatorial area (Focas, 1962). The total duration of darkening -
minimum to maximum and fading back to minimum —is 0.67 Martian years in
the circumpolar areas, for the wave proceeding from the North Cap, and 0.55
years for the north wave, at its southern limit. The wave proceeding from the
:=The white areas as shown in Fig. 20, which was primarily derived from
visual observations, generally do not encroach over bordering dark areas.
^to
**de Vaucouleurs (1971) cautions that not all early morning white patches are
frost deposits on the surface and that "at least some are due to high-altitude
icy haze (or in some cases dust layers), the elevation of which can be
estimated when the cloud is completely or, more often, partly detached
from the terminator. " He cites some rare examples of clouds projecting at
the sunrise terminator.
February 1, 1972 R. Newburn, C. Michaux, JPL Sec. 4.2, page 25
Seasonal Activity
JPL 606-1
o^co^ooco^towo»^o^oo^o«l
i^ ' A ' A ' i ' iid'ito ' ito' A>
Fig. 20. Schematic base map of Mars 196?''' showing whitening areas
detected in 1962-1968 by Capen and Capen (1970).
Note: The list of whitening areas observed is given below with approximate
coordinates of their centroids:
Abalos (030, +85), Aeolis (212, -10), Aeria(310, +18), Aethiopis (245, +08),
Albor (205, +22), Amazonis (145, +22; 170, +22), Aram (015, 00), Arcadia (Alba)
(115, +50), Argyrel(030, -45), Ausonia(250, -45), Azania(185, +36), Baltia
(040, +65), Candor (075, +08; 078, +20), Cebrenia (2 15, +42), Cecropia (290, +75),
Chryse (038, +02), Claritas ( 100, -30), Crocea (290, -02), Daedalia ( 122, -12),
Deucalionis Regio (345, -11), Dia (090, -60), Dioscuria (3 10, +52), Edom
(350, -02), Electris (180, -45), Elysium (215, +30), Eos (040, -10), Eridania
(215, -45), Hammonis Cornu (315, -10), Hellas (295, -45), Hellespontus
(328, -42), lerne (140, +80), Isidis Regio (280, +20), Lemuria (Olympia)
(185, +80), Lemuria (235, +75), Libya (275, 00), Memnonia ( 165, -20; 155, -05),
Meroe Insula (295, +35), Neith Regio (275, +35), Nix Cydonia (015, +45), Nix Lux
(112, -08), Nix Olympica (135, +26), Nix Tanaica (048, +50), Noachis(345, -40),
Nymphaeum (305, +08), Ogygis Regio (065, -48), Ophir (068, -08), Ortygia
(015, +55), Oxia(018, +30), Panchaia (250, +62), Phaethontis ( 1 35, -45), Phlegra
(192, +48), Propontis Quadrangle ( 170, +50), Scandia(l60, +68), Scandia-Boreum
Mare (110, +80), Sinai (062, -25), Syria (090, -15), Tempe (065, +47), Tharsis
(100, +02), Thaumasia (075, -30; 090, -40), Utopia (265, +52), Uchronia (260, +60),
Xanthe (050, +15), Zephyria ( 190, -08).
"Ref.: Capen, C.F., A 1967 photovisual chart of Mars: J, ALPO, v. 22,
nos. 7-8, August.
Sec. 4. 2, page 26
C. Michaux. JPL
February 1, 1972
JPL 606-1 Seasonal Activity
South Cap lasts 0.50 Martian years in the circumpolar area and 0,40 years at
its northern limit (Focas, 1962). '■-
The average "intensity" of dark areas on Mars increases from pole to
equator. The additional intensity resulting from the wave of darkening
decreases from poles to equator. This is balanced by the fact that two -waves
affect the equatorial regions. The behavior of the wave of darkening is shown
diagrammatically in Fig. Zl taken from Focas (1961), which indicates dark
areas affected. An average behavior is also shown on a separate overlay for
Figs. 9 and 10.
The seasonal variation of polarization (differences between dark and
bright areas) is likewise shown in Fig. 22, taken from Dollfus (1961). It
indicates that the variations take place principally during spring and summer
in each hemisphere and are maximized at the end of spring.
The photometric and polarimetric curves of Focas and Dollfus, which
exhibit a very definite seasonal relationship to the sublimating-regression of
the spring-summer caps, quantitatively depict the wave of darkening phenomenon
presumed to exist on Mars.
The cause of the wave of darkening was commonly attributed in the past
to water vapor, released from the vaporizing polar cap, that somehow
interacted with the surface material or with "vegetation" to cause the darkening.
Other explanations, which now (since about 1965) have gained in favor, generally
involve seasonal transport of dust on and off of dark areas. See for example the
Sagan and Pollack (1967) "windblown dust" theory. Both classes of explanation
of the wave of darkening are only hypotheses. Pollack, Greenberg, and Sagan's
(1967) statistical analysis is unable to differentiate between the two hypotheses.
At present, the wave of darkening remains one of Mars' greatest enigmas.
The reality of an actual "wave" of darkening sweeping down from regress-
ing pole to equator has been seriously questioned in very recent years (since
about 1969) by B. Smith of NMSU, and by W. Baum and C. Capen of Lowell
Observatory, for example. They claim that extensive examination of photo-
graphic records (such as Lowell Observatory's collection of Mars plates) reveal
no convincing evidence for the existence of a darkening wave.** However, it
appears that the study was only cursory (a systematic analysis requiring much
more work), and therefore it is too early to consider this view here.
*Focas (1962) pointed out that under the highest telescopic resolutions, the
dark areas resolve into a mottled pattern of small dark spots or nodes of the
order of 100-km across, on a dusky background; and that it is these nodes
that do the darkening.
**Interestingly, some evidence of the inverse phenomenon, a seasonal brighten-
ing of (adjacent) bright areas, was obtained recently through photoelectric
scanning of the disk (Boyce and Thompson, 1971). This would confirm what
Capen earlier (1967-1969) had detected through densitometric measurements
of photographic plates.
February 1, 1972 R. Newburn, C. Michaux, JPL Sec. 4.2, page 27
Seasonal Activity
JPL 606-1
n - 1 S0° 2W° S20°I!°W° 120° 20 0°
' I — I — ' — ' — ' — ' — r^ — I — I — I
\-'hiver'\'print.'^ e/e-^au/.-\
South Cop->> ^ .^.kii^ti^j .„r
Polar hood '\\_^
nuOffe polaire M
Depress. 1
Hellesp. ;
M. Australe ~^^
M, Chronlum
Hellespontus
AoniuB S.
calofe en fusion
Phrbd R.
Thaumasla
M, HadrUcum
M. Tyrrhen.
M. Sirerum
Soils L.
M, Serpentli
Pandor. Fr.
M. Erythr.
M. Clmmer.
Aurorae 5.
laplgla
Margar. S.
S. Sabaeut
S. Merldlanl
TIthoniui L.
Equator
SjTTtls major
-50
■W
-30°
-20° -
y^i
Lunae palui
Trlv, char. +20'
Ncpentbea
Tlioth
NUUcus L
Ismen. L.
M. Acldallum
PropontU
Doreosyrtli
North Cap -
nua^e polaire
Fblar hood
^ Sublimating cop
caloite en fusion
Aotomn Winter Spring Sorprper
-• ou/ -p }!iyer^prin/. •+» efe H
Jl'280° 0" 80° ISO' :w°
Fig. 21. The seasonal waves of darkening of the dark areas of Mars according
to the photometric data of Focas. The waves are seen to stem alternatively
from the two poles and fade out as they reach about 22° in the opposite
hemisphere (Focas, 1961).
Sec. 4. 2, page 21
C. Michaux, JPL
February 1, 1972
JPL 606-1
Seasonal Activity
%
+ 4
+ 2
-2
-4
winter
spring
8U(nnf>er
fall
winter
s
Quthern
• a
.: ^•^
•
• •
Hemispher
• ^^
s«
^/^
^v •
y/9
-
• <►
-
—
-
-
—
m
• •
A KInrlkiArn
—
• ^x^^
(
* Henisphere
—
Nv • y
•
—
•
winter
1
1 Spring
•umnter
1
fall
1
winter
1
1
1
'OO
+4
+2
-2
-4
90
180 270
90 7) 180*
Fig. 22. Mars, seasonal variation of polarization. Polarization differences
between dark markings and orange deserts, for phase angle a = 25°,
plotted against heliocentric longitude (Dollfus, 1961).
SEASONAL BEHAVIOR OF SURFACE FEATURES
The wave of darkening is a statistical phenomenon differing in its effect
from area to area, and is in no sense an absolute description of even the
behavior of intensity of all dark areas. Besides such albedo changes, there are
color changes and changes in shape, size, and internal appearance of the
various dark areas. Bright areas also show seasonal changes in color and
sometimes in structure. These changes in individual northern equatorial and
polar areas have been summarized in Tables 4 through 7* prepared by C. Capen
for this document. Tables 8 and 9 (Capen, 1972) give the seasonal changes for
the south and north polar regions, respectively.
General Comments on Martian Colors
Winter colors of dark areas tend to be very subdued, often grayish or
brownish hues. Lack of contrast sometimes causes features to disappear.
Blue -greens, yellow-greens, and blacks appear common in late spring. The
'canals, ' which react to the wave of darkening like any dark area, become
prominent in spring. Summer is a period of deepening color with changes to
purples, browns, and grays, which fade as summer progresses. Fall is a
drab period very like winter.
-i=The regional maps are frora the Mars 1969 I. P. P- map prodviced at Lowell
Observatory.
February 1 , 1 972
R. Newburn, C. Michaux, JPL
Sec. 4. 2, page 29
Seasonal Activity
JPL 606-1
Table 4. Table of seasonal changes in northern dark areas of Mars
(Capen, 1972).
Area and location
(center of area)
Martian
st'asoii
Typp of change
Color
Shape aijd size
Map
Mare Acidalium
+ 50° latitude,
35° longitude
Spring
Dark gray ,
blue -gray,
with oasis
g ray-green;
gray in general
Boundaries weak
1 1
g ORTYGIA
<
BALTIA 5
ARETHUS
\
::,, * ACIDALIUM / o
. . CALLIRtHOES O '^
M. V jC) <
^ ISMENIU!
CTOONt*
/?■ NILIACUS EDEN
.O .<=„,.. L.
' ^ OXIA $ < if ^^s
AORORAE % .J^^efe. SASAfU
Summer
Black-green
central area;
large, gray-
green oases;
changes to
gray-green at
end of summer
Large , swollen in
size
Fall
Lightens to
blue-gray,
losing contrast;
still dark,
variated
Borders fade
Winter
Grays and
browns
Center fades to
match edges
Niliacus Lacus
+ 32* latitude,
32° longitude
Spring
Dark gray
unmapped oasis;
late: dark gray
Internad changes in
shape;
borders and center
darken
ACIOALIUS
,r,c. ACIDALIUM
EMPE C.Z.-.
'p ^'''^
^ NILIACUS EDEI
^.CHKLrS L. ^
LUNAE \^ OXIA 0- i ^
L -.^ 5 ''L
Summer
Dark gray-
Swollen in size;
locks larger
Fall
Dark, gray
Winter
Gray
Borders weak
Mare Boreum
(Baltia-Boreum)
+65" latitude,
85° longitude
Spring
Dark gray;
blue tint and
black shades as
cap melts;
black blue
1 1 1
°'* B O R E U M M.
\
ACIOAilUS
ORIUS '"«°"' '■
ASCURIS
L.
s \
z
ARCADIA <
S TEMPE
^V MARtOTIS
>1- NIK .c^o.c. re
Summer
Black-blue
Fall
Less contrast;
dark gray
111 -defined borders
Winter
Dark gray
Bene-ith cap
Sec. 4. 2, page 30
C. Michaux, JPL
February 1, 1972
JPL 606-1
Seasonal Activity-
Table 4. Table of seasonal changes in northern dark areas of Mars
(Capen, 1972) (continued).
Area and location
(center of area)
Propontis Connplex
+ 50° latitude,
170* longitude
Nodus Laocoontis-
Alcyoniua
+ 30' latitude,
255* longitude
Utopia-Boreosyrtia
+ 50" latitude.
270' longitude
Martian
season
Spring
Summer
Fall
Type of change
Color
Darker gray
Gray black;
loses contrast in
late sumnier
Winter
Spring
Skimmer
Fall
Dark gray to
black
Mid-gray shade
Shape and size
Highly complex
changes in shape
Shape: oases
swollen, canals
dark and seasonal
canal structure;
Size: larger oases
Fading shape
Very active;
A: gray-green
to black-green;
B: small,
black-green
oasis;
C: intense,
saturated green;
D: not seen
C: black-green,
darker
Spring
Summer
Fall
Winter
C: fades,
gray-green;
others gray
Inactive;
A and C seen;
medium gray
Increase in
contrast;
dark gray 61
hue
General: dark
blue -gray;.
mid-Bummer:
intense blue-
black with
variations,
dark browns,
olive drabs,
ochre
Fading con-
trast;
dark gray
Shape: simple
parallelogram with
corner oases;
Size: small weak
oases, not
connected
Well defined
border, secular
changing position
toward S. E.
'ANCHAIA
C: most color -
saturated (green)
on Mars, 1964-65
Changing shape
Internal structure
changes in size
At tinnes yellow
cloud changes its
shape and '
appearance
Dark gray
Vo-<cv°'^'
AU50NIA
BOREALIS
"^-k %^--^
%
CECROPIA
LEK
February 1, 1972
C. Michaux, JPL
Sec. 4. 2, page 31
Seasonal Activity
JPL 606-1
Table 5. Table of seasonal changes in equatorial dark areas of Mars
(Capen, 1972).
Area and Location
(center of area)
Meridiani Sinus
0* latitude,
0* longitude
Margaritifer Sinus
-10' latitude,
25* longitude
Aurorae Sinus
-12' latitude,
60* longitude
Trlvium Charontis-
Cerberus [
+ 15* latitude,
200'* longitude
Martian
season
Spring
Summer
Fall
Winter
Spring
Summer
Fall
Winter
Spring
Summer
Fall-
Winter
Spring
Type of change
Color
Dark gray shade
Black shade
Dark gray to
black
Sometimes
black'blue
Early : dark
gray shade;
late : blue -gray
hue + dark gray
Dark gray-
Changes to
dark gray +
mottled brown
Dark gray +
mottled dark
brown
Early : black
shade;
late: black-blue
Black-blue
Black and dark
gray shade
Shape and size
Shape: Caret*
extension toward
Argus;
Size: Fastigium
Aryn fills in (dark-
ens) between two
canal Carets (does
not become black)
W. Caret dims;
no change in size
Size: Fastigium
Aryn becomes light
Map
Internal canal
changes In size
111 -defined
Fall
Winter
Stays dark
Medium dark
gray and brown
tints;
olive drab
seasonal change
takes on intense
dark gray
Stays dark
Dark gr4y and
brown
Medium dark
gray and brown
tints
Size: below nornnal
contrast. 1964-65
i
pandorae fr.
►»►■'
TITH0NIU5 J^'f"
SYRIA
A^KORAE
CLARITAS
/ SINAI
//.
^^'■"
.s>^^'
.^
.<^
\
rnLc^snA
PROPONTIS
/
ELYSIUM
1
'^
( . SYtTB
/
OP TRIVIUM
^ CHAtONTIS
tlV
.IM''
V
o> c
o.
■icy
o'^''
ZEPHYRIA"
*Term used to designate root-like darkening toward canal.
Sec. 4. 2, page 32
C. Michaux, JPL
February 1, 1972
JPL 606-1
Seasonal Activity
Table 5. Table of seasonal changes in equatorial dark areas of Mars
(Capen, 1972) (continued).
Area and location
(center of area)
Martian
season
Type of change
Color
Shape and size
Map
Mare Cimmerium
-30" latitude,
200" longitude
(equatorial and
southern area)
Spring
Medium dark
gray with dark
brown and
purple
Summer
Nondescript
medium dark
gray with dark
brown and
nnedium gray
streaks
Well-defined.
Heaperia becomes
light.
Fall
Dark purple-gray
Winter
Changes to dark
brown with
purple -gray
streaks
Edges: N. E. Gomer
Sinus has been
showing secular
change.
ELECTRIS
CHKONIUM M.
Mare Tyrrhenum-
lapygia
-15" latitude,
270" longitude
Spring
Summer
Fall
Winter
Purple and dark
brown
Carets increase in
contrast;
no change in shape
ELYS
No change
Fading purple
and mottled
brown;
also grays
Carets decrease in
contrast
Medium gray with
mottled brown and
yellow-green;
most prominent
and colorful
feature on Mars
Yellow-cloud
affects lapygia
Syrtis Major
+ 1 1" latitude,
290" longitude
Spring
Blue-gray
late : bright,
saturated
blue-green
Summer
Dark blue-green;
late ; returns to
blue -gray
North tip secular
change
Fall
Dark blue-gray
Winter
Dark gray-blue
Seasonal shift W
Sabaeus Sinus-
Mare Serpentis
-10* latitude,
330* longitude
Sp ring
Dark gray
Size: very stable,
high' contrast area
Summer
Black
Possible S border shift
Fall
Dark gray shade
Stable
Winter
Black and
dark gray
Possible So.
border shift.
Yellow cloud
affects Serpentis
February 1, 1972
C. Michaux, JPL
Sec. 4. 2, page 33
Seasonal Activity
JPL 606-1
Table 6. Table of seasonal changes in northern light areas of Mars
(Capen, 1972).
Area and
location
(cente r of
area)
Cydonia -
Ortygia
+ 50* latitude,
0" longitude
Tempe
+40° latitude,
70° longitude
Arcadia,
Scandia
+45" latitude,
135° longitude
Propontis
+ 50* latitude
180" longitude
Type of change
Normal ochre desert
color south of Noveni
Viae in Nix Cydonia;
exhibits diurnal sea-
sonal whitening
Large ochre region;
filling-in by Tempes
canal (secular);
Nix Tanaica exhibits
diurnal seasonal
v/hitenine
Map
ORTYGIA
BALTIA
ARETHUSA
L.
jHUIS
ACIDAUUM
NfUACUS
L
CALLIMHOES O oc
jv^''
EDEN
B O R E U M M.
\
BALTIA
Arcadia;
large ochre area,
over which recurrent
afternoon clouds form;
yellow dust clouds
have been observed
Scandia;
dark ochre;
some seasonal
darkening;
sometinnes filled in,
appears similar to
nnaria
Propontis complex
seasonal darkeninu.
Secular motion
E-W and SW
ASCURIS
L.
\
DIA
<
of
o
MAREOTIS
ASCRAEUS
L
TEMPE
LUNAE
L.
.vVacmm
THARSIS
%.
ACIDA
M.
NIUAC
L.
1
SCANDIA
1
BORE
PROPONTIS
" »«oaj CASTORIUS
1
DIACRIA f
BJROTAS
ASCU
L.
•r
EUXINUS
ARCADIA ?
'X
AMAZONIS
\
NtX
OLYMflCA
■litis
O
MAREC
ASCRAEI
L.
THARSIS
Sec. 4. 2, page 34
C. Michaux, JPL
February 1, 1972
JPL 606-1
Seasonal Activity
Table 6.
Table of seasonal changes in northern light areas of Mars
(Capen, 1972) (continued).
Area and
location
(center of
Type of change
Map
area)
Phlegra,
Phlegra:
1 1
Cebrenia-
variable color hues
LEMURIA
Aetheria
ranging from light to
dark ochre, and
+45° latitude.
yellowish;
PANCHAIA
210° longitude
exhibits seasonal
_
darkening due to
WTHONIUS PROPONTi;
filling-in
PIA I II
^ STYMPMAUUS
Cebrenia-Aetheria;
/
normal ochre region
^ 5 CEBRENIA s,
T -^ PHLEGRA §
r^ PROPONTIS
«« IHOAN* -S- „^oi 1
ELYSIUM ' '^ ,
Dioscuria
Normal ochre area;
some seasonal
1 I
+ 50' latitude.
whitening
CECROPIA
320° longitude
raius COPAIS
'• UTOPIA
DIOSCURIA 4?
UMBRA . . ^
MOTONILUS '' r, ' ^jt
Panchaia-
Normally dark ochre
Cecropia
region;
darkening during
+65° latitude,
spring;
180-340°
light ochre in late
longitude
sunnnrier, fall;
covered by polar cap
Lemuria
in winter;
+ 70° latitude, .
within Cecropia
240° longitude
( + 65° latitude.
275-315° longitude)
a large, white, oval
frost patch has been
observed from time
to time
Lemuria dar^k when
first uncovered by
cap
LEMURIA
ORTYSIA
CECROPIA
PANCHAIA
ARETH
I.
USA "Mius COfAB
$nH<»,us PROPONTIS
r.
JTOPIA L 1"
February 1, 1972
C. Michaux, JPL
Sec 4.2, page 35
Seasonal Activity
JPL 606-1
Table 7. Table of seasonal changes in equatorial light areas of Mars
(Capen, 1972).
Area and
location
(center of
area)
Type of change
Map
Arairi ,
Chry se ,
Xanthe
+ 10° latitude,
30° longitude
Aram:
light ochre;
seasonal whitening
Chryse:
normal desert ochre
hue. Secular
darkening 1969
Xanthe:
known 'to whiten in
some areas
Candor -
Tharsis
+ 10" latitude,
90° longitude
Light ochre hue;
clouds are generated
in this, region;
entire region becomes
seasonally whitened
LUNAE
L.
THARSIS
''%
Syria-Sinai -
Thaumaria etc.
-30° latitude,
90° longitude
Soils Lacus
-28* latitude
90° longitude
Daedalia
-15° latitude,
125° longitude
Variegated colored
area with hues rang-
ing from dark orange,
light ochre, to yellow;
seasonal local-area
whitening;
sometimes covered
by yellow clouds
Soils L. shifts
position and shape
seasonally and
secularly
llNONIA
V
r, ^ TITH0NIU5 '^«
* SYRIA /
J SINAI
,J(iri&!^?tj ■ NEC
AR(
Normal ochre color;
some seasonal
darkening.
MESOGAEA \ „,.
PMOtNICIS
MEMNONIA J>
CLARITAS
tema
L \^^^^^'''"'
y
Sec, 4. 2, page 36
C. Michaux, JPL
February 1, 1972
JPL 606-1
Seasonal Activity-
Table 7. Table of seasonal changes in equatorial light areas of Mars
(Capen, 1972) (continued).
Area aiid
location
(ccntf r of
area
Meninonia -
Mesogaea
area
-10° latitude,
160° longitude
Amazonis
+20° latitude.
60° longitude
Zephyria,
Aeolis
-10° latitude,
200° longitude
Elysium,
Aethiopis
+20" latitude,
235' longitude
Typf of change
Normal ochre color;
seasonal w'hitening
i<nown to occur on
soutliern border of
Memnonia
Very bland, feature-
less area (Tempe-
Arcadia similar);
recurrent cloud
activity;
Nix Olympica in
Amazonia becomes
seasonally whitened
(historically known);
yellow clouds have
been observed in this
region. A few well
defined dark circular
oases seen at times.
Zephyria:
normal ochre color
seasonal whitening
Aeolis:
light region;
whitening at times
in summer and
winter
Elysiunn;
normal light ochre
(Albor area gray-
white);
becomes dazzling
white during late
spring with rapid
regression of north
polar cap;
diurnal retreat of
whitening leaves pink
tint
Aethiopis:
normal ochre color
dark, transient
and secular features
Map
MESOSAEA
THARSIS
I
Nnr
I
MMNUS
ARCADIA
X
f AMAZONIS \
i .
Imjesogaea
V
ASCR^
L.
THARSIS
»*^
ELYSIUM * "^
|\\ <^.
»#^
'iUM
"lONTIS
ZEPHYRIA--
MESOSAI
AMAZ'
February 1, 1972
C. Michaux, JPL
Sec. 4. 2, page 37
Seasonal Activity
JPL 606-1
Table 7. Table of seasonal changes in equatorial light areas of Mars
(Capen, 1972) (continued).
Area and
location
(center of
area)
Hespt'ria
-20° latitude,
235° longitude
Type of change
Libya-Crocea,
Isidis Regio-
Neith Regio
area
+ 20° latitude,
280° longitude
Dark ochre narrow'
area;
shows some seasonal
£illing-in (darkening)
and lightening during
winter
Aeria,
Arabia-Eden,
Edom
+ 20° latitude,
330° longitude
Libya-Crocea:
ochre region;
displays seasonal
morning whitening,
blue limb haze, and
summer cloud.
Isidis Regio-Neith
Regio;
light ochre desert
regions;
much morning
seasonal whitening
and limb haze.
Aeria:
yellowish hue with
enclosed Nymphaevim
being gray-white
shade;
similar characteris-
tics to Elysium-Albor
area;
may become pink
around Nymphaeum
Arabia-Eden;
normal ochre color;
morning limb frost
has been observed in
Eden area
hi^lit desert area;
seasonal whitening
Map
^O-^ ZE
ERIDANIA
DIOSCURIA
UTOPIA
UMBRA
nOTONILUS
o
■•A
\
' ^ if \ ^\ I -^^ -
/
uiuav^uniA
UMBRA
ISMENIUS PROION11U5 "S*
I. ■
^^^rJAPYGIA ^
Sec. 4. 2, page 38
C. Michaux, JPL
February 1, 1972
JPL 606-1
Seasonal Activity
The color base-maps (Figs. 8 through 10) are based upon filter
photography, color photography, and visual studies. Color saturation has been
increased to aid reproduction. The maps represent a useful description of
seasonal changes on Mars, but again the colors should be accepted with due
reservation.
An overlay gi^'ing names and locations of prominent Martian features is
included and identified as 'Place-Name of Surface Features' and can be used
with any of the color base-maps (Figs. 8 through 10). A list of names and
coordinates of the features appears as Fig. 11.
Table 8. Table of seasonal changes in south polar region of Mars
(Capen, 1972).
Area
Argenteus Mons
-70° lat.
40° long.
M. Oceanidum
-60° lat.
35° long.
M. Australe
-65° lat.
90° long.
Season
Winter
(N. Hemi.
sphere)
Spring
Summer
Fall
Winter
Spring
Summer
Fall
Winter
Spring
Summer
Fall
Characteristic
Dark gray and ocher.
Medium contrast until covered by
winter haze hood.
Covered by South Polar Cap.
Seasonally bright white with
South Polar Cap projection.
Later becomes dark.
Dark gray.
Very low contrast. Hazy.
Covered by winter -type hazes and
South Polar Cap.
When first uncovered by South Cap
it becomes dark.
Dark gray to black. Affected by
yellow clouds.
Dark gray until covered by xsinter-
type haze.
Covered by South Polar Cap.
Darkens as South Polar Cap
retreats.
February 1, 1972
R. Newburn, C. Michaux, JPL
Sec. 4. 2, page 39
Seasonal Activity
JPL 606-1
Table 8. Table of seasonal changes in south polar region of Mars
(Capen, 1972) (continued).
Area
Argyre II
-66° lat.
70° lone
Thyles Mons
-70° lat.
155° long.
Thyle I and II
-65° lat.
160° long.
220° long.
M. Chronium
-60° lat.
215° long.
Thyle Collis
-70° lat.
230° long.
Promethei S.
Cher sonesus
Euripus I
-63° lat.
260° long.
Season
Winter
Spring
Summer
Fall
Winter
Summer
Fall
Winter
Spring
Summer
Fall
Winter
Spring
Summer
Fall
Winter
Summer
Fall
Winter
Spring
Summer
Fall
Characteristic
Light ocher, seasonal whitening.
Light area until covered by polar
hazes.
Covered by South Polar Cap.
When first uncovered by cap, a
dark ocher. Later a light ocher.
A light ocher area.
Covered by South Polar Cap.
Exhibits a bright polar cap
projection late fall.
Light ocher area.
Light and ill -defined.
Covered by South Polar Cap.
Dark ocher hue.
Varigated dark gray and browns.
Dark and ill -defined.
Covered by South Polar Cap.
Darkens and expands as cap
retreats.
Light ocher area.
Covered by South Polar Cap.
Retains a bright cap projection.
Dark gray.
Ill -defined.
Covered by South Polar Cap.
Appears dark when uncovered by
South Polar Cap.
Sec. 4. 2, page 40
C. Michaux, JPL
February 1, 1972
JPL 606-1
Seasonal Activity
Table 8. Table of seasonal changes in south polar region of Mars
(Capen, 1972) (continued).
Area
Season
Characteristic
Novissima Thyle
-70° lat.
325° long.
Winter
Summer
Fall
A dark brown and gray area.
Covered by South Polar Cap.
A bright cap remnant known as
Mountains of Mitchel.
Table 9. Table of seasonal changes in the north polar region of Mars
(Capen, 1972).
Area
Season
Characteristic
M. Baltia
+ 63° lat,
40° long.
Hyperboreus L.
+80° lat.
55° long.
M. Boreum
+ 65° lat.
95° long.
Scandia
+ 66° lat.
150° long.
Winter
Spring
Summer
Fall
Winter
Spring
Summer
Fall
Winter
Spring
Summer
Fall
Winter
Spring
Summer
Fall
Covered by North Polar Cap.
Dark gray when uncovered by cap.
Dark ocher.
Light ocher.
Covered by North Polar Cap.
Covered by North Polar Cap.
Becomes very dark gray and brown
and expands as cap retreats.
Very dark and partly haze
covered.
Covered by North Polar Cap.
Appears dark gray when
uncovered by cap.
Remains dark.
Medium gray and ill-defined.
Covered by North Polar Cap.
Dark ocher and brown when not
covered by cap.
Medium to light ocher. Contains
a Polar Cap white remnant spot.
Light and ill -defined.
February 1, 1972
C. Michaxox, JPL
Sec. 4. 2, page 41
Seasonal Activity
JPL 606-1
Table 9. Table of seasonal changes in the north polar region of Mars
(Capen, 1972) (continued).
Area
Panchaia
Lemuria
+ 65° lat
205° long,
Olympia
+ 80° lat.
210° long.
Utopia
Uchronia
+ 60° lat.
260° long.
M. Cecropia
+ 67° lat.
305° long.
Ortygia
+ 65° lat.
350° long.
Season
Winter
Spring
Summer
Fall
Summer
Fall
Winter
Spring
Summer
Fall
Winter
Spring
Sumnier
Fall
Winter
Spring
Summer
Fall
Characteristic
Covered by North Polar Cap.
Dark ocher and brown when not
covered by cap.
Light ocher.
Light ocher and ill -defined.
White remnant area of
North Polar Cap.
Light and hazy.
Covered by North Polar Cap
Utopia dark gray when uncovered.
Dark gray.
Light gray and varigated browns.
Covered by North Polar Cap.
Dark ocher when uncovered.
Dark to light ocher. Contains a
Polar Cap white remnant spot.
Light and ill-defined.
Covered by North Polar Cap.
Dark ocher when uncovered.
Medium to light ocher.
Light and ill -defined.
Sec. 4. 2, page 42
C. Michaux', JPL
February 1, 1972
JPL 606-1 Seasonal Activity
BIBLIOGRAPHY
Antoniadi, E. -M. , 1930, La planete Mars 1659-1929: Hermann et Cie, Paris,
239 p.
Ashbrook, J. , 1958, The new lAU nomenclature for Mars: Sky and Telescope,
V. 18, no. 1, p. 23-25, November.
Baum, W. A. , Millis, R. L. , Jones, S.E., and Martin, L. J. , 1970, The inter-
national planetary patrol program: Icarus, v. 12, no. 3, p.435-439i May.
Boyce, P. B. , and Thompson, D. T. , 1971, A new look at the Martian 'violet
haze' Problem. I. Syrtis Major - Arabia, 1969, Preprint of Paper from
the Planetary Research Center, Lowell Observatory, Flagstaff, Arizona
(submitted to Icarus, 1971).
Capen, C. F. , 1972, (Pasadena, Calif., Jet Propulsion Laboratory) private
communications to C. Michaux, January-February.
Capen, C. F. , 1971, Martian yellow clouds -- past and future: Sky and
Telescope, v. 41, no. 2, p. 117-120, February.
Capen, C. F. , 1970, Observational patrol of Mars in support of Mariners VI and
VII: Pasadena, Calif. , Jet Propulsion Laboratory, Tech. Rep. 32-1492,
12 p. June 15.
Capen, C.F., 1966, The Mars 1964- 1965 apparition: Pasadena, Calif., Jet
Propulsion Laboratory, Tech. Rep. 32-990, 187 p. , December 15.
Capen, C. F. , and Capen, V. W. , 1972, Meteorological phenomiena in Physical
observations of Mars 1966-1967-1968: Pasadena, Calif., Jet Propulsion
Laboratory, in preparation.
Capen, C. F. , and Capen, V. W. , 1970 Martian north polar cap, 1962-68: Icarus,
V. 13, no. 1, p. 100-108, July.
Capen, C. F. , and Cave, T. R. , 1971, Mars 1969 -- The north polar region --
ALPO Report II: The Strolling Astronomer (J. ALPO), v. 23, nos.3-4,
p. 67-75, August; andnos.5-6, p. 79-85 November.
Capen, C.F., and Martin, L. J., 197Z, Photographic survey of Martian yellow
storms: Abstract to appear in B uU. Am . Astronom . Soc . v. 4, no. 3, pt. II.
Capen, C.F., and Martin, L. J., 1971, The developing stages of the Martian
yellow storm of 1971: Lowell Observatory Bulletin No. 157, (or:
Bulletins V. 7, no. 20, p. 2 1 1 - 21 6) November 30 .
de Vaucouleurs , G. , 1972, Telescopic observations of Mars in 1971 -- III: Sky
and Telescope, v. 43, no.l, p. 20-21, January .
February 1, 1972 C. Michaux, JPL Sec. 4.2, page 43
Seasonal Activity JPL 606-1
de Vaucouleurs, G. , 1971, Cloud activity on Mars near the equinox: comparison
of the 1937 and 1969 oppositions in Planetary atmospheres (International
Astronomical Union Symposium. No. 40, held in Marfa, Texas,
October 26-31, 1969): Sagan, C. , Owen, T. C. , and Smith, H, J. , Editors :
D.Reidel Publishing Co., Dordrecht, Holland, and Springer -Verlag
New York Inc., New York, 408 p. , see p. 3 10-3 19-
de Vaucouleurs, G. , 1965, Charting the Martian surface: Sky and Telescope,
V. 30, no. 4, p. 196-201, October.
de Vaucouleurs, G. , 1962, Precision mapping of Mars: La Physique des
Planete's, CoUoque International Universite de Liege.
de Vaucouleurs, G. , 196la, Sources of areographic coordinates 1909-1954:
Harvard College Obs., Sci. Rep. No. 2, ARDC Contract AF19(604)-746l,
AFCRL 257.
de Vaucouleurs, G. , 1961b, Areographic coordinates for 1958: Harvard College
Obs., Sci. Rep. No. 4, ARDC Contract AF19(604)-746l, AFCRL 818.
de Vaucouleurs, G. , 1954, Physics of the planet Mars: London, Faber and
Faber, 365 p.
Dollfus.A., 1961, Polarization studies of planets, chapter 9, p. 343-399 in
Planets and satellites, v. Ill: The solar system; Kuiper, G. P. ,
and Middlehurst, B. M. , Editors : 1961, U. Chicago Press, Chicago,
601 p.
Fischbacher,G. E. , Martin, L. J. , and Baum, W. A. , 1969, Mars polar cap
boundaries: Flagstaff, Ariz., Planetary Research Center, Lowell
Observatory, JPL Contract 95 1547, May.
Focas, J. H. , 1962, Seasonal evolution of the fine structure of the dark areas of
Mars: J. Planet. Space Sci. , v. 9, no. 5, p. 371-381, July.
Focas, J. H. , 1961, Etude photome'trique et polarimetrique des phenomenes
saisonniers de la planete Mars: Annales d'Astrophysique, v. 24, no. 4,
p. 309-325, July-August.
Ingersoll, A. P. , 1970, Mars: occurrence of liquid water : Science, v. 168,
no. 3934, p. 972-973, May 22.
Leovy, C.B., Smith, B. A. , Young, A. T., and Leighton, R. B. , 1971, Mariner
Mars 1969: atmospheric results: J. Geophys. Res. , v. 76, no. 2,
p. 297-312, January 10.
Ley, W., 1963, Watchers of the skies: New York, Viking Press, 528 p.
Sec. 4.2, page 44 C Michaux, JPL February 1, 1972
JPL 606-1 Seasonal Activity
Mazursky, H. , Batson, R. , Borgeson, W. , Carr, M. , McCauley, J. , Milton, D. ,
Wildey, R. , Wilhelms, D. , Murray, B. , Horowitz, N. , Leighton, R. ,
Sharp, R. , Thompson, W. , Briggs, G. , Chandeysson, P. , Shipley, E. ,
Sagan, C. , Pollack, J. , Lederberg, J. , Levinthal, E. , Hartmann, W. ,
McCordjT. , Smith, B. , Davies, M. , de Vaucouleurs, G. , and Leovy, C. ,
1970, Television experiment for Mariner Mars 1971: Icarus, v. IZ,
no. 1, p. 10-45, January.
Pollack, J. B. , Greenberg, E. H. , and Sagan, C. , 1967, A statistical analysis
of the Martian wave of darkening and related phenomena: J. Planet.
Space. Sci., v. 15, no. 5, p. 817-824, May.
Sagan, C. , and Haughey, J. W. , 1966, Launch opportunities and seasonal
activity on Mars, chapter I6 (p. 283-291) in Biology and the exploration
of Mars; Pittendrigh, C. S. , Vishniac, W. , and Pearman, J. P. T. ,
Editors ; Wash., D. C. , Natl. Acad. Sci. Nat. Res. Council, Pub. 1296,
516 p.
Sagan, C. , and Pollack, J. B. , 1967, A windblown dust model of Martian
surface features and seasonal changes: Cambridge, Mass. ,
Smithsonian Astrophysical Observatory, Special Report No. 255,
November 8, 44 p.
Slipher, E. C. , 1962, Mars, the photographic story: Cambridge, Mass. , Sky
Publishing Corp., and Flagstaff, Ariz. , Northland Press, 168 p.
Wells, R. A., 1967, Some aspects of Martian clouds and their relationship
to the topography in Moon and planets (proceedings of the 7th Inter-
national Space Science Symposium, COSPAR, Vienna, May 10-19,
1966); Dollfus, A. , Editor: North-Holland Pub. Co. , Amsterdam,
336 p. , see p. 262-273.
Wells, R. A., 1966, An analysis of Martian clouds and their topographical
relationships: E. S. R. O. Scientific Note ESRO SN-54, 59 p. May.
February 1, 1972 C. Michaux, JPL Sec. 4.2, page 45
JPL 606-1 Seasonal Activity
APPENDIX
MARTIAN SEASONAL DATES
Seasonal date is indicated by the value of Ls> areocentric longitude of the
Sun, measured in the orbital plane of the planet from its vernal equinox. Thus,
Ls = 0° corresponds to the beginning of spring in the Northern Hemisphere, and
LS = 90° to the beginning of summer, etc. (See Fig. 8 in Section 1 on Orbital
and Physical Data. )
It is sometimes indicated by r\, the heliocentric orbital longitude of the planet,
measured from Earth's vernal equinox, or First Point of Aries T. The rela-
tion connecting r\ to Lg, following the adoption of the Martian equator of de
Vaucouleurs (1964) by the American Ephemeris and Nautical Almanac in 1968,
is now very nearly: t^ = Ls + 85° (the 85° "constant" deviates very slowly with
time, according to the precession of the equinoxes).
February 1, 1972 C. Michaux, JPL Sec. 4.2, Appendix, page 1
V
^..
K^
JPL 606-1 Seasonal Activity
Fig. 8. Color map of the Martian surface in
northern fall -winter and southern spring -
Slimmer, with white and yellow cloud activity
during these seasons. The color map was
compiled from observational data by C. F.
Capen (color contrasts of features were nec-
essarily increased to aid reproduction). The
Mercator format was obtained from the
International Astronomical Union (Ashbrook,
1958) and from de Vaucouleurs' map (1965)
■nb-'^ and areographic coordinates (196la, 196lb,
1962).
C. Capen, April 1, 1967
J. de Wys, JPL Sec. 4.2, page 19
r
(
r\ £
«iiO
CO
Q
k
Si
!
I
DC
2
CO
Q
3
O
I
CO
S
to
•c
K
5
to
It
I
«t
o
CO
5
o
5
-» O «
JPL 606-1 Seasonal Activity
c.
(^
Fig. 9. Color map of the Martian surface in
northern spring and southern fall, with wave
of darkening, frost, and white and yellow
cloud activity during these seasons. The
color map was compiled from observational
data by C. F. Capen (color contrasts of
features were necessarily increased to aid
reproduction). The Mercator format was
obtained from the International Astronomical
Union (Ashbrook, 1958) and from de Vaucou-
leur s' map (1965) and areographic coordinates
(1961a, 1961b, 1962).
C. Capen, April 1, 1967
J. de Wys, JPL, Sec. 4.2, page 21
- »»^
i ^
I
I
I
I
a"
HI
o
I
I
-1
O
o
O
CO
i
5
-1
1 .1 t :
1 1 s S
C 9)
« ^ a ^
ti « ii
q:
Co
a:
Ui
1
5
ft
o
to
ft
ft
q:
O
ft
% ^
l«
If
•4
-^
1 ^
5
•j
IB
ef
^-
ft
8)
ft"
11^
f
JPL 606-1 Seasonal Activity
n
O
Fig. 10. Color map of the Martian surface in
northern summer and southern winter, with
wave of darkening and white cind yellow cloud
activity during these seasons. The color
map was compiled from observational data
by C. F. Capen (color contrasts of features
were necessarily increased to aid reproduc-
tion). The Mercator format was obtained
from the International Astronomical Union
(Ashbrook, 1958) and fromi de Vaucouleurs'
map (1965) and areographic coordinates
(1961a, 1961b, 1962).
C. Capen, April 1, 196?
J. de Wys, JPL Sec. 4.2, page 23
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JPL 606-1 Atmosphere
SECTION 5 CONTENTS
ATMOSPHERE
Introduction 5
5. 1 Atmospheric Composition
Data Summary 1
Discussion 1
5.1.1 Observed Constituents 1
Carbon Dioxide 1
Carbon Monoxide 4
Isotopes of Carbon and Oxygen 4
Water Vapor 4
Upper Atmospheric Constituents 4
Ozone 8
5. 1. 2 Constituents Sometimes Assumed Present 8
Argon 8
Molecular Nitrogen 8
Molecular Oxygen 8
Oxides of Nitrogen 9
Sinton Bands 9
"Reduced Gases" 10
5. 1.3 Upper Limits on Possible Atmospheric Constituents 10
Bibliography 11
Figures
1. Water vapor abundances (northern hemisphere) 6
2. Water vapor abundances (southern hemisphere) 7
5. 2 Surface Pressure
Data Summary 1
Discussion 1
5. 2. 1 Historical 1
5. 2. 2 Spectroscopic Surface Pressures 2
Theory and Techniques 2
Observational Results 5
5. 2. 3 Occultation Surface Pressures 6
Theory and Techniques 6
Observational Results 7
5. 2. 4 Conclusions About Mean Surface Pressure 8
Bibliography 10
Tables
1. Occultation surface pressures 8
March 1, 1972 Sec. 5, Contents, page i
Atmosphere JPL 606-1
5.3 Lower Atinosphcr
e
Data Smiiniarx' i
Discussion i
Layers of the Lower Atmosphere i
Physics of the Lower Atmosphere 2
Troposphere 2
Stratosphere and Mesosphere 4
Contemporary Models of the Lower Atmosphere 5
Typc-s of Models 5
Convective 5
Radiative £,
Convective -Radiative ^
Lower Atmosphere Models I, II, and III ■. g
Conclusions o
Bibliography 21
Figure s
1. Table of ground air temperatures for Mars referred to northern
seasons j^O
2. Lower Atmosphere Model I for ground air temperatures 180,
190, 200, and 210'K H
3. Lower Atmosphere Model I for ground air temperatures 220,
230, 240, and 250°K 12
4. Lower Atmosphere Model I for ground air temperatures 270
and 290= K I3
5. Lower Atmosphere Model II for ground air temperatures 180,
190, 200, and 210°K 14
6. Lower Atmosphere Model II for ground air temperatures 220
230, 240, and 250°K 15
7. Lower Atmosphere Model II for ground air temperatures 270
and 290^ K 1^
8. Lower Atmosphere Model III for ground air temperatures 180
190, 200, and 210°K 17
9. Lower Atmosphere Model III for ground air temperatures 220,
230, 240, and 250°K 18
10. Lower Atmosphere Model III for ground air temperatures 270
and 290'^ K ig
11. Table of contemporary models for lower atmosphere of Mars .... 20
5.4 Upper Atmosphere
Data Summary 1
Discussion 1
Layers of tlie Upper Atmosphere 1
Physics of the Upper Atmospliere 2
Photodis sociation Region 2
Ionosphere 3
Ionization Processes 4
Thermal Processes 5
Sec. 5, Contents, page ii March 1, 1972
JPL 606-1 Atmosphere
5.4 (cont'd)
Contemporary Models of the Upper Atmosphere 8
Preliminary E-Model 8
F^ -Model 9
F2 -Model 11
Conclusions 13
Bibliography 20
Figure s
1. Upper Atmosphere Fj -Model: table of calculated neutral number
densities and temperatures vs. altitude ■ 14
Z. Upper Atmosphere F^ -Model: ion and electron density vs.
altitude ' 14
3. Upper Atmosphere F^ -Model: temperature vs. altitude 15
4. Upper Atmosphere F^ -Model: neutral density vs. altitude 16
5. Upper Atmosphere F^-Model: profiles of photoionization rate of
neutral constituents 16
6. Upper Atmosphere F2-Model: table of calculated neutral and
electronic number densities and temperature vs. altitude 17
7. Upper Atmosphere F2 -Model: number density of electronic and
neutral constituents vs. altitude 17
8. Upper Atmosphere F2-Model: tennperature vs. altitude 17
9. Table of significant reactions in the Martian ionosphere for a
pure CO2 lower atmosphere 18
10. Table of incident solar flux densities at Mars and absorption
cross sections for selected wavelength regions 18
1 1 . Table of models for upper atmosphere of Mars based on
Mariner IV results 19
March 1, 1972 Sec. 5, Contents, page iii
JPL 606-1 Atmosphere
5. ATMOSPHERE
INTRODUCTION
That Mars has an atmosphere was well appreciated by astronomers of the
19th century, who saw a disk with fuzzy edges, brilliant polar caps that came
and went, and even what appeared to be clouds or haze that at times obscured
the surface. By the turn of the century the great Princeton astronomer, Charles
A. Young, had authored a standard textbook which included interpretations of
planetary atmospheres. He made the following statement which presents the
best opinion available at the time, although by modern standards the reasoning
is not exactly flawless:
"This (Mars) atmiosphere, however, contrary to opinions
formerly held, is probably much less dense than that of
the Earth, the low density being indicated by the infre-
quency of clouds and of other atmospheric phenomena
familiar to us upon the Earth, to say nothing of the fact
that, since the planet's superficial gravity is less than
two fifths of the force of gravity on the Earth, a dense
atmosphere would be impossible.""''
The atmosphere of a planet is an immensely complex thing, an order of
magnitude more complex than an ordinary hot stellar atmosphere where only
atoms need be considered. (Some molecules become important in the coolest
stars.) Not only does a planetary atmosphere consist largely of molecules , but
even matter in liquid or solid state may be present. Furthermore, a planetary
ati-nosphere is in a perpetually nonisotropic radiation field, as insolation varies
with planetary rotation and revolution. We have really begun to understand the
Earth's atmosphere only as it has been penetrated by aircraft, balloons, and
rockets. I£ we deinand perfect accuracy, a "simple" Z4-hr weather prediction
for a terrestrial city is still beyond our practical capabilities, due to lack of
sufficient radiosonde data to properly delineate boundary conditions for the
problem and to lack of electronic computers with sufficient speed to solve the
equations involved even at the rate at which the physical phenomena are
occurring .
Fortunately the real engineering needs of our space program are not so
great as our scientific curiosity. In the pages that follow an attempt is made to
define those gross atmospheric parameters most needed for a successful entry
into the Martian atmosphere and a landing (intact) on the Martian surface.
These data will be refined as additional observational and theoretical results
become available. Meanwhile, an attempt is made to assess realistically the
probable errors in the results presented.
Young, C. A., 190Z, p . 36 1 -37 1 in Manual of astronomy: Mar s: Boston, Mas s . ,
Ginn and Co .
April 1, 1967 R. Newburn, JPL Sec. 5, page 5
JPL 606-1
Atmospheric Composition
5. 1 ATMOSPHERIC COMPOSITION
DATA SUMMARY
The species C, O, H, H2O, O3, CO2, and CO have been observed in the
Martian upper atmosphere. The isotopes Cl3o2l6, Cl2ol6ol8, Cl2ol6ol7,
Cl3ol6, and Cl2ol8 have been detected in addition to the normal forms of CO2
and CO. No other atmospheric species have been observationally confirmed to
date. However, there are many observational upper limits and theoretical
amounts of other species (see text). The abundances of those observed are
identified below.
Carbon dioxide, CO.
Carbon monoxide, CO
72 ±5 m atm (average over planet). This
amount is at least slightly variable with
season.
12.5 ,' . cm atmi.
-0.4
Water vapor, H2O
Ozone, 0_
DISCUSSION^
5.1.1 Observed Constituents
0-50 ji precipitable; variable in both time
and place (see text and Figs. 1 and 2).
10 |Ji atm. This could be ozone trapped in
the polar caps, rather than free in the
atmosphere (see text).
Carbon Dioxide
Carbon dioxide was the first (and for many years the only) gas that had
been detected spectroscopically in the Martian atmosphere. Kuiper (1952) dis-
covered the relatively strong, Zv\ + 2^2 + ^3* (1.57 (o.) and v\ + 4v2 + ^3 (I.6O fi),
bands on one of his pioneering spectrometer tracings of Mars made October 7,
1947, He added the discovery of the 2^1 + 1^3 (1.96 fx), <jy + 2v2 + v3 (2.01 p.), and
4^2 ■'■ ^o. (2.06 |ji) bands when Mars was near opposition in February 1948. These
bands are sufficiently strong to be seen in both Martian and terrestrial
atmospheres.
Kuiper's initial attempt to give an abundance took no account of pressure
differences between Martian and terrestrial atmospheres and, as he himself
Ta brief discussion of spectroscopic theory and techniques is given in Section 5. 2.
'^The "old" band designations for CO2 are used throughout this document. The
newer nomenclature is perhaps a bit more descriptive in a quantum mechani-
cal sense, but at least four different versions of it have been used since 1965,
causing considerable confusion to the nonspecialist. Furthermore, none of
the new systems is universally accepted as yet.
April 15, 1971
R. Newburn, JPL
Sec. 5. 1, page 1
Atmospheric Composition JPL 606-1
recognized, was quite inadequate. The separation of the superimposed
absorptions was handled more properly by Grandjean and Goody (1955), who
used Kuiper's data to derive a pressure-concentration product P| C = 1.6 ±0.5
X 102 mb2 (Pg n surface pressure in mb, C^ = volume mixing ratio). A con-
centration cannot be derived from a strong band alone, unless very high resolu-
tion of individual rotational lines has been achieved. Using the surface pressure
figures, available in 195 5, Grandjean and Goody found the CO2 concentration to
be about 2% by volume. Using modern values for the pressure, the concentra-
tion calculated exceeds 100%. Considering the relatively poor resolution in
these pioneering infrared spectra, the lack of good laboratory line strengths,
and the fact that the theory used made no allowance for such refinements as
temperature differences or the proper individual damping constants for each
atmosphere, a factor-of-six error is not surprising.
In 1963, while searching for (and finding) water vapor in the Martian
atmosphere, Spinrad, Munch, and Kaplan discovered on the same plate a num-
ber of faint spectral lines in the 8700 A. region which are not present in solar
spectra (Kaplan et al. , 1964). These proved to be lines of the very weak 5v,
band of CO2, a band so weak that it should not have appeared at all unless the
CO2 abundance of Mars were much higher than previously thought by astrono-
mers. Such proved to be the case, and Kaplan, Miinch, and Spinrad derived a
CO2 abundance for Mars of 50 ±20 m atm>:' assuming a 200° K Martian tempera-
ture and based upon that one plate (Kaplan et al. , 1964). A detailed calibration
by Owen of the same plate corrected the Kaplan, Miinch, and Spinrad air -mass
function and allowed for doubling back of the band, thereby resulting in a re-
vised abundance value of 46 ±20 m atm (Owen, 1964). Much of the large prob-
able error arose from the uncertainty in the measurement of the one plate of
Mars upon which the numerical values derived from all of the elaborate theory
depended. This particular plate initiated the modern era in spectroscopic
investigation of the atmosphere of Mars.
During the Mars apparitions of 1964-1965 and 1967, there were intensive
efforts by many workers to determine CO2 abundance (and surface pressure).
Many additional photographic plates were taken of the 513 band at the highest
available resolution (Spinrad et al. , 1966; Owen, 1966; Barker, 1967). The
primary source of error here was the limited resolution available for so weak
a band and the problem of continuum location. Photoelectric scans of the 2 u-t +
3v3 (1.04 \i) band (Belton et al. , 1968) and the v^ + 2^^ + Sv, (1,05 [i) band
(Belton and Hunten, 1966; Belton et al. , 1968; Giver et al. , 1968), which are
stronger and somewhat easier to measure accurately, were carried out. These
bands, being stronger, are no longer on the linear part of the curve of growth
(see Section 5. 2), however, and are affected by pressure. Also, there is still
a problem of continuum location, which tends to become even worse at longer
wavelengths. Carleton et al. (1969) used a PEPSIOS Fabry-Perot interferom-
eter to scan two lines of the v^ + Zv-^ + 3 v^ band. This work was observationally
of high accuracy, but errors in the reduction process appear to have resulted in
a value approximately 10% too large.
=l=See Appendix A, Units Used for Atmospheric Abundances.
Sec. 5. 1, page 2 R, Newburn, JPL April 15, 1971
JPL 606-1 Atmospheric Composition
While visiting at the Jet Propulsion Laboratory during 1964, J. Connes
and P. Connes (1966) perfected an interferometric spectrometer capable of
high spectral resolution of Mars in the near infrared (1.2-2.5 fx). During the
April 1967 Martian opposition, the Connes were able to obtain spectra at a
resolution in excess of 0.08 cm"^ (Connes et al. , 1969). The resolution of the
strong COo bands in this region was sufficiently high to attempt determination
of both abundance and pressure from the strong bands. Preliminary results
were given by Gray Young (196 9), and an analysis fitting 15 CO2 bands, line-by-
line, using Voigt profiles (sec Section 5. 2) has now been completed (Gray Young,
1971a). This last result, 72.1 ±0.5 m atm of CO^, must be considered the best
planetwide spectroscopic analysis to date. Because this value is sufficiently
more accurate as a planetwide mean than the preceding results, none of the
previous results will be discussed further. Perhaps the largest remaining
source of uncertainty, a potential systematic error not reflected in the quoted
error, is the effective path through the atnnosphere. Gray Young (1971a) used
r| =3.5 (see Section 5. 2), as a reasonable value for an intermediate degree of
planetary limb darkening. The extreme possibilities add no niore than roughly
±3 m atm to the uncertainty. Pressure broadening by argon is only about half
as effective as CO-, self-broadening. The presence of a small amount of argon
would result in a slightly smaller Lorentz half width (see Section 5. 2) for Gray
Young's calculations and would require a slightly larger amount of CO2 to match
the measured line equivalent widths (see Section 5. 2). It is always difficult to
allow properly for systematic error, but it does not appear overly conservative
to suggest 95% confidence that the limits were no more than ±5 m atm froin
72 m atm during April 1967.
One reason that considerable confidence can be placed in the interfer-
ometric result is that it represents a inean over an entire hemisphere and, in
fact, niorc than a hemisphere, since Mars rotates during the long integration
time required to obtain a good interferogram. A spectrograph slit typically
will project a footprint on Mars covering only abovit 1% of the hemisphere,
although planetary rotation will considerably increase the averaged area. A
spacecraft radio occultation experiment literally provides only two points, one
at ingress and one at egress, and radar studies have indicated Mars to have
large elevation differences (see Section 3. 3). In fact, several attempts have
been made both from the ground (Wells, 1969; Belton and Hunten, 1969; Wells,
1971) and from spacecraft measurements (Herr et al. , 1970) to determine
topography from the varying COo abundance.
In 1966, Leovy (1966a, 1966b) and Leighton and Murray (1966) provided
evidence that the polar caps of Mars could consist primarily of frozen CO2.
The Mariner 7 infrared radiometer has confirmed that the temperature of the
south polar cap is about 150 °K, a temperature low enough to permit condensa-
tion of CO2 (Neugebauer et al. , 1969). This indicates the atmospheric CO2
abundance must be, at least, slightly variable. The imaging experiment on
Mariner 7 gave evidence that in some areas the south polar cap thickness was
"at least tens of meters" rather than millimeters or centimeters (Sharp et al. ,
197 1). However, there is no positive method to estimate the bulk density or
fractional CO2 content of the polar cap material. The CO2 abundance quoted
here, 72 m atm, appears proper for April 1967, which corresponds to mid-
summer in the northern hemisphere of Mars. The abundance could vary
significantly during the Martian year.
April 15, 1971 R. Newburn, JPL Sec. 5. 1, page 3
Atmospheric Composition JPL 606-1
Carbon Monoxide
The discovery of CO on Mars was first reported by Kaplan, Connes, and
Connes (1969) from interferometric spectra taken during the 1967 opposition.
Some 50 lines of the 2-0 (2.35 fj.) band were identified, as well as a few unblended
lines in the 3-0 (1.57 fa) band. The abundance reported was 21 cm atm in the
total optical path (or 5.6 cm atm in a vertical column, using the suggested
'I = 3.75). A line-by-line fitting of Voigt profiles to the same data by Gray
Young (1971b) indicated 47 +^| cm atm in the total path. The difference is
caused by a small change in continuum placement, a subjective choice, which
results in a 25% difference in the measured equivalent widths of the lines.
Unfortunately, on the appropriate part of the curve of growth, a 25% difference
in equivalent widths is reflected as a -200% difference in abundance. Accepting
Gray Young's result as the more thorough, and applying the H = 3.75 suggested
by Kaplan, Connes, and Connes (1969) as appropriate for lines in (that)
"location on the curve of growth, " a CO abundance of 12.5 "^§'5 cm atm is
suggested as the best available figure. This gives a COrCC)"^ mixing ratio in
the bulk lower atmosphere of 1.7 X 10"-^.
Isotopes of Carbon and Oxygen
Kuiper (1964) first reported Cl^O^^O^^ and C^^O^^ to be present on Mars,
but his spectral resolution was too poor to give accurate abundances. Kaplan,
Connes, and Connes (1969) found lines of Cl3ol6 g^-^^^^ cl2ol8 in their work on
the normal isotope Cl2ol6. Gray Young (1971a) reported on C'^^o'^^O'^^,
Cl2ol6ol7^ and Cl^oi" in her comprehensive work on Martian CO^. No
departure from terrestrial isotopic ratios was found for C^-^ or O . Gray
Young's (1971a) O abundance appears large, but the band is very weak, and
she does not feel that the difference from the terrestrial ratio is necessarily
significant.
Water Vapor
The question of the existence of water vapor in the atmosphere of Mars
is a classic one in planetary spectroscopy. Pioneer studies were carried out
at Lowell and Lick Observatories, shortly after the turn of the century, and at
times yielded apparently positive results, the uncertainty being due to inade-
quate equipment and technique. Adams and his co-workers at Mt. Wilson gave
considerable attention to the problemi between 1925 and 1943, at times with
apparently positive results, but always with large probable errors. A review
of these early efforts is given by de Vaucouleurs (1954).
The 1962-1963 apparition of Mars resulted in two independent reports of
successful H2O detection on Mars. Spinrad, Munch, and Kaplan, working at
the 100-in. Mt. Wilson reflector, obtained a spectrogram of the Zv\ + v -p + vo
(8200-A) water band at a dispersion of 5.6 A/mm. It showed 11 lines witn
satellites at the doppler displacement appropriate to Mars, most of them
•i^Old calibration. New laboratory line strengths and temperature corrections
indicfite these values are too large by a factor of about two (Farmer, 1971a).
(Sec Appendix A for definition of units. )
Sec. 5. 1, page 4 R. Newburn, JPL April 15, 1971
JPL 606-1 Atmospheric Composition
apparently free of blends with other terrestrial or solar lines (Spinrad et al. ,
1963). Detailed analysis indicated an abundance of 14 ±7 [i'!= of precipitable
water. The analysis was based upon about one-third of the length of the image
of the lines, namely, that part covering the polar region (Kaplan et al. , 1964).
Working with a 50-cm (20-in. ) reflector at the scientific station on the 3600-m
(12, 000-ft) Jungfraujoch in Switzerland, Dollfus used a Lyot filter and half-wave
plate to alternately isolate the 1.4 |jl band and two adjacent bands. He then sub-
tracted the terrestrial component as determined from measurements of the
Moon and other objects. His result was an average over the planet of 200 |j. of
precipitable water (Dollfus, 1963). With recalibration, the value from the same
observations was later reduced to 45 |ji (Dollfus, 1965). However, the spectro-
gram, of Spinrad, Miinch, and Kaplan marked the effective discovery of water
on Mars, since it was the first widely accepted evidence of water, being based
upon specific identification of spectral features isolated from telluric inter-
ference. That spectrogram also resulted in a major program to improve
knowledge of water on Mars.
An intensive observing program was carried out by Schorn et al. , during
the 1964-1965 apparition of Mars, in an attempt to refine the previous result
(Schorn et al. , 1967). Over a 9-month period, 19 well-exposed spectographic
plates were obtained at McDonald and Lick Observatories, at dispersions of
4.09 and 4.14 A/mm, respectively. Those taken from September through mid-
November 1964 showed no water vapor; those taken during late December and
January indicated about 15 [}.'■' of precipitable water in the northern hemisphere
only. The season was late spring in the northern hemisphere of Mars. Further
measurements were impossible until May, due to insufficient doppler shift
(small relative radial velocity) during the opposition period, which occurred in
March. During May and June, 10 to 25 |j,* of water were detected in both hemi-
spheres; apparently more detected in the southern hemisphere than the northern.
The season was early summer in the northern hemisphere of Mars.
Equipnnent construction and bad weather combined to prevent further
water vapor observations during the 1967 Mars apparition. In 1969, good dry
observing conditions combined with new equipment, allo-wing 2 A/mm disper-
sion, yielded a large number of plates and the best results yet obtained. Even
photographic reproductions of the spectrograms clearly show the Martian water
lines. Schorn, Farmer, and Little (1969) found 26 ±5 jj. of water in the northern
hennisphere of Mars and less than one-third of that amount in the southern
hemisphere, during February and March of 1969. In a limited observing pro-
gram, Owen and Mason (1969) found similar results, when corrected by
Farmer's (1971a) line strengths. TuU (1970, 1971) described the latitude
variation of water vapor in more detail, finding a peak abundance of about 48 \i
at a latitude 30''-40''N, in late March 1969. Little (1971) presented extensive
results for 110° < Lg < 145° during the 1969 opposition. Barker et al, (1970),
have extended these studies through March 1970, finding water vapor increasing
from undetectable (<20 \i) in August 1969, to 45-50 \i, over a wide latitude range,
in March 1970. Figure 1 is adopted from Farmer (1971b), who uses most of
the observations noted above to show the abundance of water vapor on Mars as
*See footnote on page 4.
April 15, 1971 R. Newburn, JPL Sec. 5. 1, page 5
Atmospheric Composition
JPL 606-1
80
60
40
20
o
t
-20
/
f ■
SPRING
SEASONS IN THE NORTHERN HEMISPHERE
SUMMER FALL
WINTER
NORTH CAP
>,
t ■
^.. ^
f
J)
t ^
^' .
*
.•».
K ••
iHl
'
* :• n
■^%' .
0:0 ®
^
35
i45!
;30i
l^\ ^^
'01
iiol
1201
i20t
ilOf
t-i
90 180 270
AREOCENTRIC LONGITUDE OF THE SUN (Lj)
Fig. 1, Water vapor abundances (in precipitable miicrons). This is a plot
of abundance as a function of season (the abscissa) and latitude
(the ordinate). It does NOT show diurnal or longitudinal
variation, functions which are still es sentially unknown.
Sec. 5.1, page 6
R. Newburn, JPL
April 15, 1971
JPL 606-1
Atmospheric Composition
a function of location and season. All abundances shown on Fig. 1 have been
derived from the original measured equivalent widths, using a commion miethod
of reduction to the vertical water content. Therefore, they should contain mini-
mum relative errors, resulting from the choice of differing averaging factors,
intrinsic strengths, and analytical techniques. The recent data derived by
Barker et al. (1970), and Barker (1971) have not been added to Fig. 1, because
their areographic resolution was poor (Mars had an angular diameter less than
5 arc seconds by March 1970). These data are shown separately in Fig, 2.
There is some evidence for symmetry, however, with water vapor rising to a
peak in the temperate latitudes of each hemisphere during the midsummer
season. A highly informative account of the history of the search for water
vapor on Mars was produced by Schorn (1971).
Upper Atmospheric Constituents
On July 31 and August 5, 1969, Mariners 6 and 7 flew pajst Mars carrying
ultraviolet spectrometers sensitive to wavelengths from 1100 A to 4300 A, The
instruments have been described by Pearce et al. (1971). Besides CO^ and
CO, the spectrometers found COj, CO"*", C, O, and H (Barth et al. , 1971),
None of these is really surprising as an upper atmosphere constituent, as they
are obvious dissociation and ionization products of the three molecules known
to be present in the lower atmosphere. The Lyman a resonance line of atomic
hydrogen was seen to an altitude of more than 20, 000 km above the Martian
surface (Barth et al. , 1971), which was unexpected. Hydrogen is not easily
retained by a planet as small as Mars, and this would imply a considerable
continuous loss of water (assuming that water is the source of the hydrogen,
rather than captured solar wind protons or some other mechanism).
360°
Fig. 2. Water vapor abundances (in precipitable microns)
These data from Barker (1971) are for the southern
hemisphere Spring-Summer season, values ol L
(areocentric longitude of the Sun) not studied ^
until 1970.
April 15, 1971
R. Newburn, JPL
Sec. 5. 1, page 7
Atmospheric Composition JPL 606-1
Ozone
Theoretically, there should be at least a very small amount of ozone in
the atmosphere of Mars, Belton and Hunten (1969b) calculated an abundance of
about 2 X 10l6 cm-2 (7,4 ^j. atm), assuming their oxygen detection (see molecu-
lar oxygen) to be valid. In fact, that value would seem, at best, to be an upper
limit on the ozone abundance, since the tentative O2 identification seems to have
proven incorrect. Barth and Hord (1971) report the existence of an absorption
feature in Mariner data over the polar regions, a feature having all the charac-
teristics of the Hartley absorption band of O3. It could be explained either by
10 (JL atm (3 X IOI6 molecules) of O3 in the atmosphere or by O3 trapped in the
solid CO2 of the polar cap. It is intended that the Mariner 1971 flights make a
new study of ozone on Mars (Hord, Barth, and Pearce, 1970),
5,1.2 Constituents Sometimes Assumed Present
Argon
Argon has no spectral lines in regions of the spectrum observable from
the surface of the Earth, yet there has been a general assumption that some
argon should be present in the atmosphere of Mars. On Earth, argon has
resulted mainly from decay of potassium 40. If Mars has undergone the same
process of differentiation and surface concentration as the Earth, then the
argon abundance in the Martian atmosphere would be expected to be proportional
to its surface area, relative to the Earth. The surface area of Mars is roughly
28% that of the Earth. Since the Earth has about 74,5 m atm of argon, Mars
might be expected to have about 28 m atm. In fact, the abundance of argon is
almost certainly less than 10 m atm. The total atmospheric pressure (see
Section 5. 2) is so near the partial pressure of CO2 that the presence of a large
amount of argon is untenable.
Molecular Nitrogen
At one time, molecular nitrogen was thought to be the major constituent
of the Martian atmosphere. Like argon, N2 has no detectable absorption
features visible from the ground in its spectrum, so this was simply a guess by
terrestrial analogy. As a relatively accurate surface pressure and CO2 abun-
dance became known, it was realized that there was no room left for any large
amount of nitrogen. Nitrogen has a number of electronic bands in the ultra-
violet, but none of these v/ere detected by Mariners 6 and 7 (Barth et al, , 1971).
Dalgarno and McElroy (1970) have thereby set an absolute upper limit of 5%
nitrogen in the Martian atmosphere and a probable upper limit of 0.5%, assum-
ing the eddy transport coefficient is the same in the Martian atmosphere as in
that of the Earth. Thus, nitrogen must be considered a minor constituent.
Molecular Oxygen
The search for molecular oxygen on Mars is second only to that for water
vapor, both historically and in expended effort. The searches of Adams, Dun-
ham, and St. John at Mt. Wilson, during the period 1926-1934, have been
summarized by Dunham (1952). The results of the searches were negative. A
Sec, 5. 1, page 8 R. Newburn, JPL April 15, 1971
JPL 606-1 Atmospheric Composition
search by Kaplan, Munch, and Spinrad was also negative and resulted in an
upper limit of 70 cm atm (Kaplan et al, , 1964). Then, Belton and Hunten (1968)
reported a tentative identification of 20 cm atm of O^ on Mars. A more thorough
search by Margolis, Schorn, and Gray Young (1971) has set an upper limit of
15 cm atm (assuming r\ = 3), however, and molecular oxygen must still be
considered an undetected minor constituent of the Martian atmosphere. Since
there is atomic oxygen in the upper atmosphere of Mars, there should be a
small amount of molecular oxygen of photochemical origin as well, but theoreti-
cal calculations of such abundance are based on too few data to be considered
reliable.
Oxides of Nitrogen
In 1960, Kiess, Karrer, and Kiess presented "A New Interpretation of
Martian Phenomena, " a claim that many observational results of long standing
were due to the presence of the various oxides of nitrogen on Mars. That
paper contained no new observational results, consisting of a rediscussion of
previously observed results. Virtually every statement in the paper has since
proved untenable. The paper is mentioned because it unfortunately resulted
in a concern about nitrogen oxides which took a long time to dispel.
Sinton (1961), Kaplan (1961), and Huang (1961) immediately objected to
Kiess, Karrer, and Kiess' work, Sinton and Kaplan on observational grounds,
and Huang on theoretical grounds. In an attempt at rebuttal, Kiess, Karrer,
and Kiess (1963) showed a section of microphotometer tracings of Mars and the
Sun. This proved chiefly that their Martian spectrogram exhibited a very poor
signal-to-noise ratio. The 1963 paper also contained a statement that 1 to 2 mm
atm was sufficient NO2 to account for many of the observed effects. Meanwhile,
a new observational study by Spinrad set an upper limit of 1 mm atm for the
NO^ abundance in the Martian atmosjihere (Spinrad, 1963). This was later
reduced to an upper limit of 8 [i atm by Marshall (1964). Another detailed
observational study by O'Leary indicated, very conservatively, that the upper
limit of NO2 abundance was no more than 0.1 mm atm (O'Leary, 1965). Still
another study by Owen verified Marshall's results (Owen, 1966).
Kuiper (1964) has made spectrographic searches for N2O (<800 \x atm) and
NO (<20 cm atm). Sagan et al. (1965) set theoretical limits on NO and HNOt,
based on observed NO2 limits. They also set limits on N2O4, based on the
well-known relationship between monomer (NO2) and dimer (N2O4). Beer,
Norton, and Martonchik (1971) have set an observational upper limit on N2O4
of 500 \x atm, using an interferometric spectrometer. Thus, there is a great
deal of evidence that oxides of nitrogen are, at most, insignificant components
of the Martian atmosphere.
Sinton Bands
In 1957, Sinton reported the presence of three bands at 3.43 fji, 3.56 fi,
and 3.67 fo. (later revised to 3,45 |ji, 3.58 fi, and 3.69 |x) in a spectrometer tracing
of Mars, made with the 200-in. Palomar reflector. These were generally
April 15, 1971 R. Newburn, JPL Sec. 5. 1, page 9
Atmospheric Composition
JPL 606-1
attributed to some compound with a C-H bond, although such interpretations
were far from completely satisfactory (Rea et al. , 1963). Later, rather
positive identification with telluric HDD was made of the 3,58 |i and 3.69 \i fea-
tures (Rea et al. , 1965). Recent high resolution studies with an interferomet-
ric spectrometer (Beer, Norton, and Martonchik, 1971) have shown the 3.45 fi
band, the weakest of the three, to be spurious. (At least it did not appear in
1969.)
"Reduced Gases"
Connes, Connes, and Kaplan (1966) reported the discovery of unidentified
absorption bands in the near infrared, which they suggested were in part, prob-
ably caused by "reduced gases, including substituted methanes. " There has
been no confirmation of the existence of the lines, let alone positive identifica-
tion, and the observations are now generally considered invalid.
5.1.3 Upper Limits on Possible Atmospheric Constituents
It has been possible to place upper limits on the abundances of a number
of relatively simple molecules which are conceivable components of the Martian
atmosphere. The limits in the following table have been reported by Kuiper
(195Z and 1964), who surveyed the 1-Z.5 \i region with a standard spectrometer,
and Beer et al. (1971), who surveyed the 2.8-4.0 \i region with an interfero-
metric spectrometer.
Molecule
Carbon suboxide
Ammonia
Methane
Ethane
Ethylene
Acetylene
Hydrogen sulfide
Carbonyl sulfide
Formaldehyde
Formic acid
Hydrogen chloride
C3O2
NHo
CH,
C2H4
H^S
COS
HCOH
HCOOH
HCl
Abundance
Upper Limit
200 (j.-atm
1 mm-atm
1 mm-atm
1 mm-atm
30 mm-atm
20 mm-atm
30 mm-atm
1 . 5 mm-atm
50 p.-atm
70 fjL-atm
1 1 ^j.-atm
Reference
Beer et al. , 1971
Kuiper, 1964
Kuiper, 1964
Kuiper, 1952
Kuiper, 1952
Beer et al. , 1971
Beer et al. , 1971
Beer et al. , 1971
Beer et al. , 1971
Beer et al. , 1971
Beer et al. , 1971
Sec. 5. 1, page 10
R. Newburn, JPL
April 15, 1971
JPL 606- 1 Atmospheric Composition
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JPL TM33-266.
Sagan, C. , Hanst, P. L. , and Young, A. T. , 1965, Nitrogen oxides on Mars:
Planet. Space Sci. , v. 13, p. 73-88.
Schorn, R. A. , Spinrad, H. , Moore, R. C. , Smith, H. J. , and Giver, L. P. ,
1967, High dispersion spectroscopic observations of Mars, 11. The
water-vapor variations: Astrophys.!., v. 147, p. 743-752.
Schorn, R. A. , Farmer, C.B., and Little, S. J. , 1969, High-dispersion spectro-
scopic studies of Mars. III. Preliminary results of 1968-1969 water-
vapor studies: Icarus 11, p. 283-288.
Schorn, R. A. , 1971, The spectroscopic search for water on Mars -- a history:
in lAU Symposium No. 40: Planetary atmospheres; Sagan, C. , Owen,
T.C., and Smith, H. J, , Editors : D. Reidel Publ. Co. , Dordrecht-Holland.
Sharp, R. P. , et al. , 1971, The surface of Mars, 4 South Polar Cap
J. Geophys . Res. , v. 76, p. 357-368.
Sinton, \V. M. , 1957, Spectroscopic evidence for vegetation on Mars: Astrophys.
J. , V. 126, p. 231-239.
Sec. 5. 1, page 14 R. Newburn, JPL April 15, 1971
JPL 606-1 Atmospheric Composition
Sinton, W. M. , 1961, An upper limit to the concentration of NO2 and N2O4 in
the Martian atmosphere: Pub, Astron. Soc. Pacific, v. 73, p, 125-128.
Spinrad, H. , 1963, The NO2 content of the Martian atmosphere: Pub. Astron.
Soc. Pacific, V. 75, p. 190-191.
Spinrad, H. , Munch, G. , and Kaplan, L. D. , 1963, The detection of water vapor
on Mars: Astrophys. J. , v. 137, p. 1319-1321.
Spinrad, H. , Schorn, R. A. , Moore, R. , Giver, L. P. , and Smith, H. J. , 1966,
High dispersion spectroscopic observations of Mars, I. The CO2 content
and surface pressure: Astrophys . J. , v. 146, p. 331-338.
Tull,R.G., 1970, High-dispersion spectroscopic observations of Mars, IV.
The latitude distribution of atmospheric water vapor: Icarus 13,
p. 43-57.
Tull, R. G. , 1971, The latitude variation of water vapor on Mars: in lAU
Symposium No. 40: Planetary atmospheres; Sagan, C, , Owen, T.C.,
and Smith, H. J. , Editors : D. Reidel Publ, Co., Dordrecht-Holland.'
Wells, R. A., 1969, Martian topography: large-scale variations : Science,
V. 166, p. 862-865.
Wells, R. A., 1971, Martian topography from range-gated radar, ground-
based CO2, and Mariners 6 and 7 CO2 measurements: Bull.Amer.
Astron. Soc, v. 3, p. 277.
April 15, 1971 R. Newburn, JPL Sec, 5. 1, page 15
JPL 606-1
Atmosphe ric Composi tion
APPENDIX
UNITS USEJ_J FOR ATMOSPHERIC ABUNDANCES
A number of somewhat specialized units are in common use by
atmospheric physicists to describe abundances. A b ri ef glossa ry of sornr of
these terms is given below, followed by derivations of the numerical relation-
ships between some of thein.
Amagat
Meter-atm
Meter -aniagat
Millibar (mb)
The amagat is a dimensionless unit of density normalized
to STP conditions (1 atm & 0°C). One amagat implies a
number density equal to Eoschmidt's number.
A meter-atm of gas is that abundance which would
occupy a path length of 1 m at 1 atm pressure. The
temperature is often assumed (as here) to be O'C, but
this is not a part of the latest spectroscopic definition
wherein the temperature must be specified. Units of
cm-atm, km-atm, etc., have obvious analogous
definitions.
A nieter -amagat of gas is that abundance v.'hich would
occupy a path length of 1 m at a density of 1 amagat.
The millibar is a unit of pressure, of course, equal to
e s ^i r c
Tor r
10^ newtons m"2 or 10^ dynes crn"'^. Abundanc
quoted as partial pressures in many Martian papers,
commonly using millibars as the pressure unit.
The torr is also a unit of pressure, being that pressure
sufficient to raise a column of mercury at 0°C by one
millimeter in a one standard gravity field. Abundances
are sometimes quoted as partial pressures in torr.
The atmosphere is still another pressure unit, derived
from standard Earth atmospheric pressure, used to give
abundances as partial pressures.
On occasion, an abundance will be given in units of the
mass of a given molecule above some unit area, g cm-'^
being most commonly used.
Water vapor abundance is often given in terms of the
thickness of the layer which would form if the vapor were
"precipitated out. '' Sometimes the unit is gi\-en as
'precipi table microns" or any other convenient unit of
length.
Other naeasures of abundance include volume percent (or volume mixing
ratio), mass percent, and nuinber density, all of which have obvious meaning.
Atmosphere (atm)
Grams per square
centimeter
Microns (H2O)
April 15, 1971
R. Newburn, JPL
Sec. 5. 1, Appendix, page 1
Atmospheric Composition JPL 606-1
The numerical relationship among the pressure units is as follows:
1 mb = 1.000 X 10^ newtons m'^
1 torr = 1.333 X 10^ newtons m'^
1 atm = 1.013 X 10^ newtons m-2
The meter -amagat can be related to partial pressure units as follows:
, , .. .^. ^ F mg P = pressure
by definition P = "X " "a"
^ ^ F = force
, . , m m^ A = area
obviously -r- - -TT-
m = mass
Path length in m-amagat units g = local acceleration of gravity
refers to STP conditions. ^ = path length
Therefore, under standard con-
ditions of temperature and V = volume
pressure ^ _ molecular mass (phys. scale)
ui N = Avogadro's number (phys. scale)
- n o °
N o
o
n = Loschmidt's number
o
n 2
P = — ^ u^g = 4.4601 X 10" fig £ ^ in mKs units
N ^ ^ ^^ m-amagat
o
For Mars specifically, g = 3.7 1 ms" and:
_2
P (newtons m ) = 0.1655 u^ ,
^ ' ^ m-amagat
or
P (mb) = 1.655 X 10"^ \xi .
^ ' '^ m-amagat
The mass per unit area M is equal to local pressure divided by local
surface gravity. Its expression in terms of m-amagats can be taken directly
froni above
M (kg m"^) = 4.4601 X 10'^ [jl £
^ ° ' m-amagat
Sec; 5.1, Appendix, page 2 R. Newburn, JPL April 15, 1971
JPT> 606-1 Atmospheric Composition
In various partial pressure units the relationship becomes
^, / -2^ in^ P (mb)
M (g cm ) = 10 ^^ —J
g (cm s )
or
^ , ,, -Z> P (newtons m )
M (kg m ) = ^ 32
g (m s )
Precipitable water units W can be converted fromi microns to mass per
unit area simply by multiplying by the density of water (1 g cm"^). Thus,
1 micron of water equals lO""^ g cm"^ or 10"-^ kg m-2. Converting to equivalent
path units
W (microns) = 8.028 X 10^ 2
^ m-amagat
_2
In partial pressure units (using g = 371 cm s )
4
W (microns) = Z.695 X 10 P (mb)
April 15, 1971 R, Newburn, JPL Sec 5.1, Appendix, page 3
JPL 6O6-I Surface Pressure
5.2 SURFACE PRESSUllE
DATA SUMMARY
The mean surface pressure of Mars is 5.Z ±1.0 nib, exclusive of seasonal
variations. It is considered unlikely that seasonal variations can exceed
±1.5 mb, and they are probably much smaller. Pressures at the extreme high
and low elevations may reach 2,3 mb (+10 km) and 11.5 mb (-10 km), respec-
tively, which is of greater significance.
DISCUSSION
5.2.1 Historical (Photonu;tric and Polai-imetric Surface Pressures)
The history of the study of Martian surface pressure is an excellent
example of the way science often progresses; tin apparent convergence toward
an answer, followed by one or more quantum jumps to completely different
answers as systematic errors are isolated, and then a new convergence
(hopefully toward the final answer).
Within the past half century, Donald II. Menzel (1926), using pliotometric
data, established an upper limit of 50 torr (66 mb) for Martian surface j^ressure,
He noted that if the atmosphere contributed only 2% of the planetary albedo, the
surface pressure would be 20 torr (26 nib). In 1929, Lyot (1964), reporting
upon polarimetric observations, suggested that the surface pressure must be
less than 18 mm (24 mb). These pioneering efforts were followed by many
other photometric and polarimetric studies, and de Vaucouleurs (1954) con-
cluded, in summary, that the surface pressure must be 64 ±3 torr (85 ±4 ml:)).
All the "classical" attempts to determine the Martian surface pressure
involved unverifiable assumptions. Each successive worker pointi;d out the
"unwarrantable assumptions" of his jjr edeces sor s ;:uid |jroceeded to malve a new
set of his own. It is difficult to fault the work of Menzel or Lyot, who very
clearly stated their assumptions and qualified their conclusions, but many later
workers were distinctly overconfident in their assignment of probable errors.
A detailed critique of these surface pressure measurements was presented by
Chamberlain and Hunten (1965). Their criticisms are summarized in the follow-
ing paragraphs.
In even the best polarimetric work done to d.ite, it was assumed
1) that surface polarization variation across the disk is wavelength
independent.
2) that the phase-angle and zenith-distance vary independently witli
surface brightness.
3) that the atmosphere is a pure Rayleigh scatterer (least- reliable
assumption).
July 30, 1971 R. Newburn, JPL Sec. 5.2, page 1
S LI r f ai c I" P r t; s s u r e JP L 606-1
Fui'thcr, the conversion from intensity to surface pressure involves the absolute
planetary surface brightness and atmospheric composition. Finally, the obser-
va.tions themselves are difficult and involve some error, even with the best of
cc[uipnient and observers. According to Chamberlain and Hunten, the total
error, excepting the composition effect and non-Rayleigh component effect,
could be at least ±50°^;..
The composition effect enters in the conversion of a derived atmospheric
scattering intt-nsity to the number of molecules or pressure. A oure CO2 atmo-
spht^re has lialf £is many n^iolecules and about 75"'i of the surface pressure of an
N2 atmosphere with an identical scattering intensity.
If there is a non-Rayleigh component to any atmosphere, its effect
dejjends upon the size and type of particles involved. If they arc very small
particles, such as typical haze or fog-type particles (a few wavelengths of
light or less), the gas pressure will bt; overe stiniated by either polarimetric
or phototnetr ic technique's, which assume ]Dure Rayleigh scattering for Lht;
atmosphere. Large; jaarticles, such as ice crystals, would add a component
of negative polarization at small phase angles which would conversely cause
an underestimate of surface pressure. Such effects could be quite gross, caus-
ing errors of several hundred percent. In fact, a considerable amount of
submicron material would simply invalidate the whole approach (photc^metry
has shown that there is at least some non-Rayleigh comjDonent).
The photometric approach usually contains its own assumptions, such as
ignoring illumination of the ground by the atmosphere, illumination of the atmo-
sphere by the ground, and absorption. Even if these (assumptions are allowed,
it is impossible to correct for the non-Rayleigh scattering component known to
exist.
Photometry and polarimetry remain valid methods of studying the surface
of Mars, and perhaps with accurate knowledge of the atmosphere gained from
other methods, they may provide useful results on the non-Rayleigh component
of the Martian atmosphere. However, photometry and polarimetry, by them-
selves, are not cajoable of establishing a useful surface pressure for the i^lanet.
5.2.Z Spectroscopic Surface Pressures
Theory and Techniques
Astronomical molecular spectroscopy is an active and complex field of
research which cannot be easily sumniarizcul in a few pages. The following
paragraphs are intended to acquaint readers without an astrophysical back-
ground, with some of the basic concepts and terminology of this field.
An astronomical spectrogram is a photograph of a light source which has
been instr umcntally dispersed (spread out) in wavelength, usually by means of
a diffraction grating or, occasionally, a prism. In planetary spectroscopy the
source is usually sunlight reflected from a planet, and the instrument is usually
one of several types of large spectrographs, although recently Michelson inter-
ferometers have been used to produce spectra in so-called Fourier spectroscopy
(see ]Dage 5).
Sec, 5.Z, page 2 R. Ncwburn, JPL July 30, 1971
J PL 606- ]
Surfactj Pressure
In an astronomical spectrogram of the planets, th ""reflected continuun-i of
light will normally exhibit certain, relatively discrete, p rtions weakened rela-
tive to adjacfuit parts of the spectrogram. These wea] sned wavelenghts, or ab-
sorption features, correspond to the natural absorbing frequencies of atoms or
molecules between the source (Sun) and the detector (spectrometer). In the
spectra of Mars, the absorbers may be on the Sun itself (Fraunhofer lines), in
the atmosphere of Mars, or in the Earth's atmosphere. The major Martian
absorber, C02> is not present on the Sun; therefore, solar interference with
absorption lines at that frequency would be random and due to something other
than C02- Also, by means of the Doppler effect (identical frequencies, on
different bodies, shift slightly with respect to each other if the bodies have any
relative radial velocity), absorptions at the same wavelength can be separated
by making observations at a time of large radial velocity. This effect is used
particularly to separate lines of the same gas appearing on more than one body,
such as water vapor on Earth and Mars.
If a mic rodensitomcter tracing is made of a single absorption line on a
photographic plate and the image densities are converted to intensities, the
results may be somewhat bell-shaped and usually will be fairly symmetrical,
as shown below. Similar displays can be made on a strip chart recorder,
CONTINUUM
ZERO INTENSITY
if photoelectric detection is used, scanning through the spectrum on a step by
step basis. Assuming the continuum intensity to be essentially constant and
unabsorbcd over the interval of integration, the total absorption A in the line
(the shaded area) can then be expressed as
A
= I
'1 -
and I v^ the continuum intensity.
where v is the frequency, I^ the intensity at
In practice, the integration (measurement) is not carried to infinity but rather
is terminated when the difference ly - ly cannot be separated frotri the noise
(unless a mathematical model for I,, is used). The effect of finite instrument
resolution (instrument profile) is a complication vvhieh v/ill not be -;isc issed
here, except to note that it causes line shapes to spread, but usually le-ives
total absorption unaltered (although it may cause nearby lines ^o make that
absorption difficult or inapossible to measure).
'In actual practice, mic rodcnsitomete r or spectrometer data maybe transferred
onto magnetic tapes for computer jDrocessing.
July 30, 1971
R. Newburn, JPL
Sec. 5.2, pi
S u 1- f a c e P 1- e s s u r e JP L 6 6 - 1
Total absorption is generally expressed in terms of equivalent width W,
the width of an equivalent rectangular line of zero intensity everywhere within
the line. In other words, A = I,, W, and
a
dth.
ne
the
ce
The relationship between equivalent width W of a given spectral line and
the product of the number of molecules of absorbing gas (in the columin creatin
the spectral line) times the intrinsic strength of the line is called the curve of
growth . When that product is small (the weak line region) the relationship is
linear one, and doubling the number of miolecules doubles the equivalent widtl
When the product becomes larger and the line begins to saturate, there is a
transition to a square root law (the strong line region), requiring four tinnes as
many molecules to double the equivalent width. In the transition and strong li
regions, the gas density also enters the relationship. In a real atmosphere, 1
number of molecules cannot be changed, but molecular absorption bands or
lines of different intrinsic strengths can be utilized, determining an abundan
with lines of a very weak band and density-abundance product with a strong
band, thus allowing derivation of atmospheric pressure. The Curtis -Godson
approximation allows ready comparison between a real atmosphere (with tem-
perature and pressure decreasing with height) and a fictitious homogeneous
atiTiosphere (for details, see Chap. 6, Goody, 1964). In fact, the pressure
determined spectroscopically, assuming a homogeneous atmosphere, is very
nearly half the surface pressure in the real atmosphere. This is the technique
that was used in the first and miost of the successive spectroscopic pressure
determinations (Kaplan, Munch, and Spinrad, 1964).
The detailed shape of a spectral line is usually the result of combined
effects of Doppler (thermal) broadening and pressure (collisional) broadening,
referred to as a Voigt profile. The integral describing the area under such a
profile is a function of two parameters, usually identified as the Voigt function
(also referred to as the Hjerting function). Numerical tabulations have been
produced by Hummer (1965), among others. The relative contributions of
Doppler and pressure broadening to a Voigt profile are a function of tempera-
ture, pressure, and line strength.
Rotational temperature can be derived by comparing the relative inten-
sities of lines in a given rotation-vibration band (e. g. , Gray, 1969). Then for
an assumed abundance and pressure, Voigt profiles can be used to calculate
precise equivalent widths for each line in many different bands. These calcu-
lated widths are then compared with measured equivalent widths, and the pro-
cess is iterated until the residuals reach a miinimunn. Using this process, it
is not necessary that bands of any particular strength be used, but a greater
number of bands and a wider range of strengths provide a more accurate
solution.
Sec. 5.2, page 4 R. Newburn, JPL July 30, 1971
JPL 606-1 Surface Pressure
Observing techniques have improved considerably in recent years.
Initial work on the Martian surface pressure was carried out with a standard
high dispersion coude spectrograph, using photographic detection. Such instru-
ments have been described in detail by Bowen (1962). Photoelectric detection
has often been used on the 1.04, 1.05, and 1.06 \i. bands to improve detective
quantum efficiency and permit acquisition of spectra of adequate resolution in
a reasonable observing period. All conventional spectrographs necessarily
have an entrance slit which accepts only a small fraction of the light froir. a
planetary image. In the infrared bands where the flux is low, the lirnited
amount of light available results in poor resolution. The develnpnient of the
P'ourier spectrometer by P. Connes (Connes and Connes, 1^66) has made it
possible to accept the total flux from a planet and, further, to record virtually
all wavelengths simultaneously (to multiplex), even in infrared regions where
photographic plates are unavailable. This significantly improves the resolution
of the infrared bands and provides far more accurate data. The latest spectro-
scopic results are based largely upon Fourier spectroscopy. An excellent
elementary discussion of Fourier Spectroscopy has been written by Beer (1968).
Observational Results
Following the initial 1963 application of spectroscopy to the determination
of the surface pressure of Mars by Kaplan, Munch, and Spinrad (1964), a n-iajor
effort was conducted by many observers to improve the accuracy of the results,
partly because an accurate surface pressure was critical for the design of
spacecraft to be landed on the Martian surface. These results appeared in a
long series of papers in the "Astrophysical Journal, " "Icarus, " and the "Journal
of Quantitative Spectroscopy and Radiative Transfer" between 1966 and 1969.
Some details of this work have been given in Section 5. 1, under Carbon Dioxide.
The definitive spectroscopic surface pressure is the result of Gray Young
(1971), using Fourier spectra taken by the Connes during April 1967, which
exhibited resolution in excess of 0.08 cm"l between 1.2 and 2.5 \i (Connes et al. ,
1969). This resolution was sufficient to allow Gray Young to use the technique
of iterative fitting of Voigt profiles described in previous paragraphs. Gray
Young was able to fit 15 separate bands varying in intrinsic strength by more
than three orders of magnitude, and most bands having 50 or more unblended
lines, with a result of 5.16 ±0.64 mb. This quoted error does not include several
potential sources of a small amount of systematic error. There is a small
dependence upon composition, and the result given assumes a pure CO2 atmo-
sphere. Pressure broadening by argon is only about half as effective as CO2
self-broadening. The presence of a smiall amount of argon would result in a
slightly smaller half-width for Gray Young's calculations and would require a
relatively larger amount of CO^ to match the measured line equivalent widths.
This would change the pressure slightly. Gray Young used an effective atmo-
spheric path of 3.5, which appears reasonable for an intermediate degree of
linib darkening. A longer path would imply a smaller pressure and vice versa,
but at most, the potential error is only a few tenths of a millibar.
July 30, 1971 R. Newburn, JPL Sec. 5.2, page 5
Surface Pressure jpr /(-,/ ,
One reason that considerable confidence can be placed in the
intorferometric result is that it represents a mean over an entire hemisphere,
and in fact, more than a hemisphere, as Mars rotates during the long integra-
tion time required to obtain a good interferogram, A spectrograph slit typically
will project a footprint on Mars covering only about 1"/,, of the hemisphere,
although planetary rotation will considerably increase the averaged area.
Gray Young's value can be considered a good mean surface pressure for Mars
during April 1967 . This seasonal qualification is necessary, since there is
clearly a change in the total amount of the Martian atmosphere as the polar
CcLps increase and d'^creasc in size.
Belton and Hunten (1969), and Wells (1969), attempted to derive local
elevations from photoelectric observations of variations in the strength of the
1.05 (J. CO^ band across the Martian disk. Actual pressures were not reported.
The elevation trends were generally in accordance with radar and spacecraft
results, but the absolute values were not reliable.
Herr et al. (1970), used the Mariner VI and VII infrared spectromicter
data on the 2 p. CO2 bands to make 114 pressure (and altitude) measurements
which averaged 5.3 ±0.3 mb. Individual values varied from 3.7 to 8.1 mb and
the correlation with other topographic data is good. The large number and
localized nature of these measurements makes them extremely valuable topo-
graphically and generally useful for mean surface pressure.
5.2.3 Occultation Surface Pressures
Theory and Techniques
When an electromagnetic wave passes through a planetary atmosphere, its
amplitude, phase, and direction are changed. The ability to measure and inter-
pret these changes depends greatly on the frequency involved. Unfortunately,
there has never been an accurate observation of the optical occultation of a
relatively bright star by Mars; although such observations have been made for
Venus, Jupiter, Neptune, and lo. A successful occultation experiment at micro-
wave frequencies (2300 MHz) was first carried out by Kliore, Cain, Levy,
Eshleman, Fjeldbo, and Drake using the communication system of the Mariner IV
spacecraft (Kliore et al. , 1965).
In a modern phase-coherent communication system with the frequency
reference on Earth, extremely precise measurements of phase changes can be
made, although direction and precise power loss are much more difficult to
determine. The integrated refractive index along the path followed by a signal
transmitted from a spacecraft through a planetary atmosphere, is a direct
function of the measured phase change (Kliore and Tito, 1967). The actual
refractive index profile, as a function of altitude, can be determined if the atmo-
sphere is assumed spherically symmetric, and the indc^x is small enough to
ignore bending of the ray (Fjeldbo and Eshleman, 1965). As the atmosphere
penetrated by the signal becomes dcnst-r, the refractive index increases, bend-
ing of the ray increases, and knowledge of the spacecraft trajectory is necessary
to correct for the effective increase in atmospheric path (Fjeldbo and Eshleman,
1965). If the atmosphere is nons phe rical , the refractive- index profile can still
Sec. 5,2, page 6 R. Newburn, JPL July 30, 1971
TPL 606-1 Surface Pressure
be determined if the shape of the planet's atmosphere is known, as well as the
trajectory of the transmitting spacecraft with respect to that atmosphere.
The refractive index profile is a function of the electron density and the
neutral particle density. If the ionosphere is separated in altitude from the
bulk of the neutral atmosphere, accuracy is improved. In an ideal experiment,
at least two, frequencies will be transmitted, since the refractive index of the
neutral atmosphere is virtually independent of frequency, while that of the
ionosphere varies inversely as the square of the frequency (Fjeldbo and
Eshleman, 1965).
Assuming ionospheric effects have been removed, the refractive index
profile of the neutral atmosphere is available to attempt fits with density pro-
files or the equivalent. (In most literature, reference is made to refractivity,
which is just the refractive index minus one. ) Assuming a pure CO2 atmosphere
(mass), density can be calculated directly from refractivity, and pressure and
temperature are derived from the perfect gas law and the equation of hydro-
static equilibrium. This has proven quite satisfactory for Mars.
For more complex multi-component atmospheres, the data reduction
cannot result in unique answers without additional assumptions or knowledge.
Typically, ratios of temperature to mean molecular weight may be derivable for
at least part of such an atmosphere. The more that is known about an atmo-
sphere from other studies, the more useful the occultation technique becomes
in precisely defining an atmosphere, since the quantities actually measured are
measured with great precision.
Observational Results
It is now known that the Martian atmosphere is almost entirely CO^ with
at most a few percent of argon or nitrogen (see Section 5. 1). As a result, occul-
tation experiment refractivities can be converted into atmospheric variables of
state with great accuracy. Mariners IV, VI, and VII have each provided two
surface pressures, obtained during ingress and egress of the occulted signal.
For Mariner IV, the ingress data is somewhat the better, since the uplink fre-
quency reference transmitted from Earth remained until the signal was lost.
The egress data had only the internal spacecraft frequency standard for refer-
ence until phase lock with the Earth equipment was reacquired, which occurred
about 7.5 seconds after "reappearance, " On Mariners VI and VII, independent
local oscillators furnished frequency reference.
The pressure results obtained by the three Mariner flights are contained
in Table 1. The primary references for Mariner IV are Kliore, Cain, and Levy
(1967) and Fjeldbo and Eshleman (1968). Kliore, Fjeldbo, and Seidel (1970)
produced detailed results on Mariners VI and VII, following a preliminary
report by these authors and Rasool (1969). These six data differ presumably,
because of physical differences in local elevation, since occultation results are
referenced to a point on the planetary surface.
July 30, 1971 R. Newburn, JPL Sec. 5.2, page 7
Surface Pressure
JPL 606-1
Table 1, Occultation surface nrcssurc!
'■->' :"•. il\
A:,rir...r IV
M,i.-inrr VJ
Marin.- r VII
:r.r..,, 1 K...-.,
;v-..-.'SS
]•:.--. s
IiiL' r,->^
I--I^l-. .- -■
"7T ,^' .',,^7
\1 i!-,'
A., 1-1 .lilKIl
\!-jM.li ini
t'.ii r- 1 .s\- rt i -
.s ■('!■:
ll'■l^■^l)u^.tn-i
A-:M/.i'ni s t rv' A r-'/.t-i i i
■V- )■:
, — :.-._: —
-, , _ , , ,
T " }■:
.1 i 1 • 1-:
I ■.■ It L , ■ (■.!.-■)
■ • r -■ S
■ ■ r-
■- S " \
: '■ • 1 : : ■■ . )
:- t 1 .
■' , ■ i ■ 1 . ■-•
'..'11
I . J i 1 -
.■'■ '. 1 "■ :• I
[ill-, -1, ;'-.. -
li:l-.- •; 1 . ! I'l.'l
AUL- S, 1 ii|,'i
v^.r" t;.
'^.^-''Vr
5.Z.4 Conclusions About Mean Surface Pressure
All photometric and polarimetric surface pressures arc considered
worthless, as noted earlier. Most of the individual spectrographic studies
made from 1966 through 1969 xvcre of good quality, but cannot be compared with
the latest results. The Fourier spectrometer offers the best Earth-based
results for mean surface pressure, since it has both high spectral resolution
m the region of the strong (infrared) CO2 bands, and large scale spatial cover-
age. The spacecraft spectrometers on Mariners VI and VII produced excellent
results on localized spots, but these results were based upon relatively low
spectral resolution and are not as accurate for the mean as those obtained with
the Fourier spectrometer. The occultation results are individually the most
accurate of all, but they refer to six highly localized areas on Mars. It is
suggested at this time that a mean surface pressure of 5.2 mb be used. The
total probable error, including systematic effects, should bo no more than
±1.0 mb, except for possible seasonal variations.
Most spectrographic measurements of CO2 on Mars have been made
during the Martian Spring and Summer, in the northern hemisphere, since
these have been the seasons on Mars during opposition with Earth in recent
years. Thc^ Gray Young measurement refers to midsummer, while the
Mariner IR Spectrometer results were taken in late-summer. The conventional
spectrographic measurements made from 1965 through 1969 used such a diverse
range of equipment and assumptions, that the somewhat higher surface pres-
sures derived cannot be considered significant. It would take a careful uniform
reanalysis of all the original data to even hope to derive a significant season-
pressure correlation.
CO-
Since the Martian atmosphere is now known to be composed largely of
'2, the early CO^ abundance figures appear directly applicable to pressure-
determinations. These are probably more reliable than the early pressures
themselves, because the early infrared strong band data was generally of very
low resolution. With only one exception, the early abundances vary by no more
Sec. 5. 2, page 8
R. Newburn, JPL
July 30, 197]
JPL 606-1 Surface Pressure
than ±Z5% from the best current value. Most of these studies average data
obtained across 2- to 6-month periods of observation, thereby inhibiting any
search for specific seasonal effect. Seasonal effects on the 5,2 mb mean sur-
face pressure should not be considered to exceed ±1,5 mb. In fact, there is no
clear evidence for any seasonal effect on pressure, except for the obvious
growth and shrinkage of the polar caps.
The pressure differences between high and low Martian topography at any
given time would therefore far exceed any seasonal variation in the mean pres-
sure, since a topographical difference of +10 to -10 km corresponds to 2.3 and
11.5 mb, respectively.
July 30, 1971 R. Newburn, JPL Sec. 5.2, page 9
Surface Pressure JPL 606-1
BIBLIOGRAPHY
Beer, R. , 1968, Remote sensing of planetary atmospheres by Fourier
spectroscopy: Physics Teacher, v. 6, (4).
Belton, M. J. S. , and Hunten, D. M. , 1969, Spectrographic detection of topo-
graphic features on Mars; Science, v. 166, p. 225-227.
Bowen, I.S., 1962, Spectrographs: Astronomical Techniques, Chapter 2,
Hiltner, W. A. , Editor ; U. of Chicago.
Chamberlain, J. W. , and Hunten, D. M. , 1965, The pressure and CO2 content
of the Martian atmosphere, Rev. Edition: Geophys., v. 3, p. 299-317.
Connes,J. , and Connes, P. , 1966, Near-infrared planetary spectra by Fourier
spectroscopy, I. Instruments and results; J. Opt, Soc. Am. , v. 56,
p. 896-910.
Connes, J., Connes, P. , and Maillard, J. P. , 1969, Atlas des Spectres dans le
Proche Infrarouge de Venus, Mars, Jupiter et Saturne: Editions du
Centre National de la Recherche Scientifique, Paris.
de Vaucouleurs, G. , 1954, Atmospheric pressure. Chapter IV (p. 99-127) in
Physics of the planet Mars, Part I; London, Faber and Faber.
Fjeldbo,G., and Eshleman, V. R. , 1965, The bistatic radar-occultation method
for the study of planetary atmospheres: J. Geophys. Res. , v. 70,
p. 3217-3225.
Goody, R. M. , 1964, Atmospheric radiation, I. Theoretical basis: London,
Oxford U. Press (Clarendon Press).
Gray, L. D. , 1969, Comparison of procedures used to analyze spectroscopic
observations: the 7820-A carbon dioxide band in the spectrum of Venus:
Icarus, V. 10, p. 90-97.
Gray Young, L. D. , 1971, Interpretation of high-resolution spectra of Mars-II
Calculations of CO2 abundance, rotational temperature and surface
pressure: J. Quant. Spectros . Radiat. Transfer, v. 11, p. 1075-1086.
Herr, K. C. , Horn, D. , McAfee, J. M. , and Pimentel, G. C. , 1970, Martian
topography from the Mariner 6 and 7 infrared spectra; Astron. J. ,
v. 75, p. 883-894.
Hummer, D.G., 1965, The Voigt function, an eight-significant-figure table and
generating procedure: Mem. Roy. Astron. Soc. , v. 70(1), p. 1-32.
Kaplan, L. D. , Munch, G. , and Spinrad, H. , 1964, An analysis of the spectrum
of Mars; Astrophys . J. , v. 139, p. 1-15.
Sec. 5.2, page 10 R. Newburn, JPL July 30, 1971
JPL 606-1 Surface Pressure
Kliore.A. J. , Fjeldbo, G. , and Seidel, B. L. , May 20-29, 1970, Summary of
Mariner 6 and 7 radio occultation results on the atmosphere of Mars:
COSPAR, Leningrad.
Kliore,A. J. , Fjeldbo, G., Seidel, B. L. , and Rasool, S. I. , 1969, Mariners 6
and 7: radio occultation measurements of the atmosphere of Mars:
Science, v. 166, p. 1393-1397.
Kliore.A., Cain, D. L, , and Levy, G. S. , 1967, Radio occultation measurements
of the Martian atmosphere over two regions by the Mariner IV space
probe, p. 226-239 JJ2 Moon and planets; Dollfus, A., Editor: Amsterdam
North-Holland Pub. Co.
Kliore,A. , Cain, D. L. , Levy, G. S. , Eshleman, V. R. , Fjeldbo, G. , and
Drake, F. D. , 1965, Occultation experiment: results of the first direct
measurement of Mars' atmosphere and ionosphere: Science, v. 149,
p. 1243-1248.
Kliore,A. , and Tito, D. A. , 1967, Radio occultation investigations of the
atmosphere of Mars: J. Spacecraft Rockets, v. 4, p. 578-582.
Lyot, B. , July 1964, Research on the polarization of light from planets and
from some terrestrial substances: NASA TT F-187, a translation from
Annales de I'Observatoire de Paris, Section de Meudon, VIII, v. 1, 1929.
Menzel, D. H. , 1926, The atmosphere of Mars: Astrophys J. , v. 63, p. 48-59.
Wells, R. A., 1969, Martian topography: large-scale variations : Science,
v. 166, p. 862-865.
July 30, 1971 R. Newburn, JPL Sec. 5.2, page 11
JPL 606-1 Lower Atmosphere
SECTION 5. 3
LOWER ATMOSPHERE
This section has not been revised, although it contains outdated informa-
tion written almost 5 years ago. However, the improved information on the
Martian atmosphere obtained over the past 5 years is so substantial that a major
research effort is required to evaluate and document this data properly.
The level of sophistication now possible in producing models from this
later data should provide extremely meaningful information in the near future.
The material contained in this section can be utilized as background
inforniation.
Until the later data has been documented, reference to the Viking 75
Project, Mars Engineering Model (M75- 1 25- 1 ) will provide the reader with
improved information.
March 1, 1972 Sec. 5.3, page
JPL 606-1 Lower Atmosphere
5.3 LOWER ATMOSPHERE
DATA SUMMARY
Nine models have been calculated for the lower atmosphere of Mars, each
model giving atmospheric profiles for 10 different ground air temperatures.
All use contemporary interpretations of the Mariner IV flight occupation data
plus the best results of Earth-based observations. ' ='= The three models listed
below are presented here, representing a reasonable range of ground air tem-
peratures (180 through Z90°K) for three combinations of surface pressure and
composition; other models are available from the author. At this time Model I
is recommended. For instructions on use of the models, see Fig. 1.
Su r f a c e
pressure Composition
Model I 10 mb
(Figs. Z through 4)
Model II 10 mb
(Figs . 5 through 7)
80% CO^, 10% Ar, 10% N2
60% CO2, 20% Ar, 20% N2
Model III 15 mb 60%CO2, 20% Ar, 20% No
(Figs. 8 through 10) ^
DISCUSSION
Layers of the Lower Atmosphere
Using terrestrial nomenclature for classifying various regions of the
atmosphere, the troposphere is the lowest region, where the source of heating
is conduction from the ground and absorption of infrared energy radiated by the
ground, and the principal transport of energy is by convection. It is a region
where the kinetic temperature therefore decreases at a rate approximating the
adiabatic lapse rate.
At some abrupt level in the lower atmosphere, the convective transport of
energy virtually ceases. This level at which radiative equilibrium becomes a
good first approximation is known as the tropopause. The height of the tropo-
pause is a function of latitude, season, time of day, and solar activity. Above
the tropopause the atmosphere is approximately in a state of radiative equilib-
rium and is called the stratosphere, a region where the temperature gradient
can be either positive or negative depending upon the local conditions of the
atmosphere (Goody, 1964). The numerical value of the temperature gradient
will depend upon latitude, season, time of day, and solar activity.
See page 21 for list of cross references.
September 11, 1967 E. Monash, JPL Sec. 5.3, page 1
Lower Atmosphere JPL 606-1
Above the stratosphere but below the thermosphere is a region called the
mesosphere, where the temperature in the terrestrial case decreases with
increasing height. The amount of infornnation available about the Earth's
mesosphere is very limited and has been obtained largely by the observation of
meteor trails through the region and from rocket data. The rocket data indi-
cate that very large wind speeds occur with magnitudes as great as 150 m sec"!
(Fleagle and Businger, 1963).
Physics of the Lower Atmosphere
The important physical processes which are believed to occur in the lower
atmosphere of Mars are based upon the theories of convective and radiative
equilibrium plus the results of the Mariner IV occultation experiment and vari-
ous Earth-based telescopic observations. These processes form the basis of
Lower Atmosphere Models I, II, and III presented in Figs. 2 through 10.
Troposphere ^
Mars is believed to have a normal troposphere, at least during daylight
hours, with convection the dominant process for transporting energy. With this
assumption, the temperature lapse rate is the dry-adiabatic lapse rate for the
given composition of the atmosphere. For a well mixed atmosphere (the result
of convection turbulence) the composition will be uniform with height; that is,
for the mass density p- of constituent i and the total mass density p = Sp-, we
have p:^^/p = constant. In an atmosphere of constant composition, the ^
dry-adiabatic lapse rate is
A = g /<c> (1)
&o' p
where g^ is the acceleration due to gravity and <Cp> is the mean value of the
specific heat at constant pressure. The height of the tropopause (top of the
troposphere) is below 10 km for Models I, II, and III; hence, g^ was treated as
a constant (376 cm sec'^). <c > is given by
fPi^Vi
<c > = (2)
P Sp.
i ^
Specific heat for N^ or Ar is nearly independent of temperature over the range
of temperatures in question, while the value for CO^ varies approximately as a
weak linear function of temperature. Values of Cp for Ar, N^, and CO^ were
taken from Hilsenrath et al. (I960). Thus, the lapse rate depends both on the
composition and the air temperature found in the troposphere.
The kinetic temperature of the troposphere decreases almost linearly
with height since A changes only very slowly.
Sec. 5.3, page Z E. Monash, JPL August 18, 1967
JPL 606-1 Lower Atmosphere
T = T„ - Ah (3)
where T^ is the ground air temperature and h the height above ground. Assum-
ing hydrostatic equilibrium and the ideal equation of state, the pressure profile
for the troposphere is
/T\f^goAA (4)
where P^ is the surface pressure in dyne cm-2, ^i is the mean mass per mole-
cule in gm (mean molecular mas s/Avogadro' s number), and k is the Boltzmann
constant. The total number of atoms and molecules in the troposphere is
P
(5)
kT
To o
the
3btain the partial concentrations [CO2], [Ar], and LN^], n is multiplied by
fractional abundance of the constituent; that is
[CO2] = xn (6)
[Ar] = yn (7)
[N2] = zn (8)
where x + y + z = 1 .
Anderson (1965) assumes that the altitude of the tropopause varies
linearly with the ground air temperature, and this assumption is adopted here.
\ = a(T^ + b) (9)
where h^ is the altitude of the tropopause and a and b are constants to be deter -
mmed. The constants a and b were fitted to the results of the Mariner IV data
and Model I (44% CO^, 10-mb surface pressure) of Prabhakara and Hogan
(1965). ^
Anderson interpreted the Mariner IV data (immersion) to indicate a very
shallow troposphere or none at all. From his interpretation he derived the
result that at T^ = 175°K we have h^ = . To determine the other constant in
Eq. (9), Model I of Prabhakara and Hogan was used, and the model gives the
result that at T^ = 230°K we have h^ = 3 km; hence, a = 3/55 and b = -175 for
Tq ^ 175°K. Models will be added later for T < 175°K based upon straight
radiative equilibrium calculations.
August 18, 1967 E. Monash, JPL Sec. 5.3, page 3
Lower Atmosphere JPL 606-1
Stratosphere and Mesosphere
In the Martian stratosphere, radiative transport is believed to be the
dominant mechanism of energy transfer. The results from the Mariner IV
experiment do not indicate that radiative transport is the only mechanism pres-
ent in the Martian stratosphere, but at the level of complexity justified by
present knowledge of Mars, this is the only process worth consideration here.
The results of the radiative equilibrium calculations of Prabhakara and
Hogan (1965) show temperature gradients which are all negative and range in
magnitude from 0. 9 to 1 . 2° K km- 1. Models I, II, and III (Figs. 2 through 10)
have negative temperature gradients above the tropopause, and the range in
magnitudes for the gradients is from 0.08 to 1.8°K km'l. The radiative tem-
perature gradients a in °K km"! are calculated from the expression
To - 1'75
a = -2 (10)
62.5
where Tq, as before, is ground air temperature. Equation (10) is derived from
Figure 2 of Anderson (1965). The temperature profile from h^ to 50 km is
T = Tj - Q;(h - h^) (11)
where Tj = T{^t) , that is, the kinetic temperature at the height of the tropo-
pause. From the equation of hydrostatic equilibrium and the ideal equation of
state, the pressure profile from h^ to 50 km is
1
T
where Pi = P(h(-) and g, is treated as a constant with the value of 370 cm sec'^
in this region of the atmosphere. The total density profile in the stratosphere
is given by Eq. (5), and the partial concentrations of [CO^], [Ar], and [N2J
are given by Eqs. (6), (7), and (8).
The Martian mesosphere, which is thought to be a layer approximately
40 km thick, in this model is a region where the temperature decreases linearly
with height to an altitude of 75 km; above 75 km to an altitude of 90 km the tem-
perature has the constant value of 155°K, which is consistent with the data pre-
sented by Gierasch and Goody (1967), and Prabhakara and Hogan (1965). The
mesospheric temperature profile used in Models I, II, and III is
Sec. 5.3, page 4 E. Monash, JPL August 18, 1967
JPL 606-1
Lower Atmosphere
jT2 - /3(h - 50) for 50 km < h < 75 km
't3 = 155°K for 75 km s h 5 90 km
(13)
where T^ = 175°K and /? = 0.8°K km-1. The value of T^ is fitted to the results
of Anderson (1965) at an altitude of 50 km, and the value of the temperature
gradient j3 is chosen to represent a mean value in the 50- to 75-km region.
From the equation of hydrostatic equilibrium and the ideal equation of state, the
pressure profile in the mesosphere is
^2|t;7") ^oi" 50 km < h < 75 km
T-
Po exp
^^g'
kT-
(h - 75)
(14)
for 75 km s h s 90 k
m
where P2 = P(h = 50 km) and g^ is treated as a constant with the value of
362 cm sec-2 in the region of 50 to 75 km. P3 is the pressure at an altitude of
75 km, and g-^ is treated as a constant with a value of 358 cm sec'^ in this part
of the atmosphere. The total number density n is calculated from Eq. (5).
The partial concentrations of the atoms and molecules are determined from
Eqs. (6), (7), and (8).
Contemporary Models of the Lower Atmosphere
Types of Models
Interpretation of the data from observational and theoretical studies has
produced three distinct types of models for the lower atmosphere of Mars
(Fig. 11). These can be classified by the dominant process which transports
the energy as convective, radiative, or convective -radiative . The "correct"
choice among them is not clear; hence, we have used a combination of the con-
vective and radiative models to produce a series of new models which are
classified convective -radiative to represent the "best" approximation to the
lower atmosphere of Mars.
Convective. A model which can be classified as convective for the lower
Martian atmosphere has been presented by Neubauer (1966). He gives a
detailed calculation for the development of thermal convection and then explores
the possibilities of the influence of this process in producing "dust-devils" on
Mars. In his treatment, thermal convection in the lower atmosphere of Mars
is approximated by Brunt's equation
8T
at
a
dZ
Ki-
(15)
September 11, 1967
E. Monash, JPL
Sec. 5.3, page 5
Lower Atmosphere JPL, 606-1
where T is the temperature, t is the time, Z is the altitude, K is the coefficient
of turbulent heat transfer, and T is the adiabatic lapse rate. K depends on the
time of day, the Richardson number, the season, solar activity, the latitude,
and altitude. Neubauer assumes K to be a function of Z only and assumes the
values of K for the Earth's atmosphere to be applicable to the Martian atmos-
phere. To extrapolate K from the Earth' s atmosphere to the Martian atmos-
phere without considering the seasonal, latitudinal, and diurnal variations of K
for the terrestrial case is a dubious exercise. The main result of Neubauer' s
work is a demonstration that dust-devil formation can occur more easily on
Mars than on Earth. ^
A by-product of Neubauer' s study is the diurnal variation of the surface
temperature of Mars for southern hemisphere summer solstice at midlatitudes ."
Figure 3 of Neubauer shows the diurnal variation of two temperature profiles,
one for the mean surface and one for 50 cm above the mean surface of Mars.
Two results are readily apparent: (1) the diurnal temperature wave becomes
damped very rapidly with altitude, and (2) the diurnal temperature waves for
the ground and for 50 cm above the ground are in phase with each other. The
first observation is disputed by the results of other authors (Leovy, 1966;
Goody and Belton, 1967). The validity of the second observation needs to be
established with more certainty.
le
Radiative. A model which can be classified as radiative is presented by
Gierasch and Goody (1967). They calculate a "simple" solution to the equation
describing the state of radiative equilibrium in the Martian atmosphere. From
their solution Gierasch and Goody discuss and evaluate the relative importances
of doppler broadening, the effect of water vapor in the Martian atmosphere,
solar heating, vibrational relaxation of CO^, and the development of a convective
troposphere. They do not present a complete model for the lower atmosphere
of Mars, only the temperature profile up to an altitude of 60 km. The tempera-
ture profile was calculated for a pure CO2 atmosphere with a surface pressure
of 4.9 mb and a surface gravity of 372 cm sec'^. A complete model (tempera-
ture, pressure, and number density profiles) can be calculated by using their
temperature profile together with the equation of hydrostatic equilibrium and
the ideal equation of state.
C onvective-Radiative . The models of Prabhakara and Hogan (1965) and
Leovy (1966) can be classified as convective -radiative models. Prabhakara and
Hogan present a detailed calculation of the thermal structure for the atmosphere
of Mars based on the absorption of solar photons in the ultraviolet and visible
regions of the spectrum by O2 and O3 and the absorption of infrared radiation
by C02.^ Prabhakara and Hogan use an iterative procedure to calculate the
atmospheric parameters of their model. When the surface pressure, surface
air temperature, and atmospheric composition are specified, a first guess for
the vertical temperature distribution will yield a pressure and number density
distribution through the use of the equation of hydrostatic equilibrium and the
ideal equation of state. From the vertical variation of the atmospheric constit-
uents, the opacity can be determined for any level of the atmosphere. With the
use of the radiative transfer equation plus the atmospheric parameters just
determined, a new temperature profile is generated. With this new temperature
distribution, they calculate new pressure and number density distributions and
then a new opacity. Then another temperature profile is determined with the
Sec. 5.3, page 6 E. Monash, JPL September 11, 1967
JPL 606- i Lower Atmosphere
use of the radiative transfer equation. This procedure is repeated until a
temperature distribution is obtained which satisfies a convergence criterion of
Prabhakara and Hogan, specifically, that at any level in the atmosphere the
temperature does not change more than 0. 1°K in two successive iterations.
The models of Prabhakara and Hogan (1965) predate the results of
Mariner IV and the recent Earth-based spectroscopic studies; hence, the mod-
els do not contain the most recent initial atmospheric parameters and should be
used with caution when they are employed to establish the boundary conditions
for upper atmospheric profiles.*
The model of Leovy (1966) was developed as a necessary preliminary to a
numerical study which attempts to simulate the atmospheric circulation on Mars
Leovy has gone into great detail to produce a model that will represent the
diurnal, seasonal, and latitudinal variation of ground and atmospheric tempera-
tures for an atmosphere with a surface pressure of 5 mb and composed entirely
of CO^.
The convective- radiative model of Leovy is derived for a two-layer
atmosphere; that is, the upper layer contains half the mass of the troposphere
and the mass of the stratosphere, and the lower layer contains the other half of
the mass of the troposphere. The two-layer approximation used by Leovy for
the lower atmosphere of Mars is not adequate to describe the vertical variation
of the atmospheric parameters. Leovy's profiles for the atmospheric param-
eters instead show the diurnal, latitudinal, and seasonal variations. From his
profiles for the temperature at an altitude of 2.89 km, Fig. 1 was derived by
extrapolation to the ground.
The atmospheric temperature near the Martian surface is given by
Tq = T(h=0) = T3 + Ah^ = T3 + 10.1 (16)
where T3 is the temperature at the reference altitude of 2.89 km, A is the tem-
perature lapse rate (3.5°K km"!), and h^ = 2.89 km. To obtain the ground air
temperatures for a given time of day, extrapolations of Leovy's results roughly
imply
Tq (sunrise) = T^ (noon) - AT^ (17)
Tq (sunset) = T^ (noon) - l/3 AT^ (18)
Tq (midnight) = Tq (sunrise) + l/3 AT^ (19)
where AT^ is the diurnal variation of the ground air temperature and the factor
1/3 is derived from Figures 3 and 6 of Leovy (1966). Use of Eqs . (16) through
(19) plus Figures 4 and 5 from Leovy allows calculation of the variation of the
ground air temperature with latitude and season as was done in Fig. 1.7
September 11, 1967 E. Monash, JPL Sec. 5.3, page 7
Lower Atmosphere JPL 606-1
Lower Atmosphere Models I, II, and III (Figs. 2 through 10)
These models are based on the general theory given on pages 2 through 5
and, in part, on the types of contemporary models discussed on pages 5 through
7. The models all use interpretations of data from the Mariner IV occultation
experiment and the best results of Earth-based observations. The assumptions
made in deriving Lower Atmosphere Models I, II, and III are listed below;
model parameters all have been adjusted to match the calculated and observa-
tional data of Mars given in items 1) through 4).
1) Near the base of the thermosphere the kinetic temperature matches
the calculated temperature of 155°K derived by Prabhakara and
Hogan (1965) based upon radiative equilibrium calculations.
2) The temperature profile is assumed linear with temperature
gradients which are consistent with the radiative equilibrium cal-
culations of Prabhakara and Hogan (1965) and Anderson (1965).
3) The altitude of the Martian tropopause, which is a function of the
ground air temperature, is derived from the ingress data of
Mariner IV and radiative equilibrium calculations.
4) The thermodynamic parameters (specifically pressure and density)
for the lower atmosphere of Mars are derived from the equation of
hydrostatic equilibrium.
5) The lower atmosphere of Mars is assumed to be in a state of
hydrostatic equilibrium above the tropopause.
6) Gas in the lower atmosphere of Mars is assumed to obey the ideal
equation of state.
7) Aerosol concentration of the atmosphere is assumed negligible.
8) Convective transport is assumed to be the dominant mechanism for
energy transport in the troposphere and negligible above; radiative
transport is assumed to dominate in the stratosphere.
9) The effect of circulation has not been included in the atmospheric
models .
10) Latitudinal, seasonal, and diurnal variations of the atmospheric
parameters are included as significant effects in the description
for the lower Martian atmosphere (see Fig. 1). ^
11) Storm activity is assumed negligible.
12) Solar activity has not been included in the models.
13) Formation of dry ice (condensed CO^) is assumed negligible.
14) The phenomenon of the "blue haze" has been ignored. ^
Sec. 5.3, page 8 E. Monash, JPL September 11, 1967
JPL 606-1 Lower Atraosphere
CONCLUSIONS
Lower Atmosphere Model I (Figs. 2 through 4), with surface pressure of
10 mb and composition of 80% CO2, 10% Ar, and 10% N2 , is recommended at
this time. Figure 1 gives detailed instructions for use of this model.
Contemporary models of the lower atmosphere of Mars which are pre-
sented in Fig. 11 are all post-Mariner IV, the model of Prabhakara and Hogan
(1965) thereby being excluded. From Fig. 11 we see that there is general con-
sistency among the various atmospheric parameters assumed for the theories.
The atmospheric composition assumed is primarily CO2 with trace amounts of
N2 and Ar. Ohring et al . (196?) prefer an atmospheric composition for the
lower Martian atmosphere of 74% CO^ and 26% N2 • The surface pressure used
in these models was 5 mb except for Ohring et al . (1967), who preferred 7 mb .
Values of the mean ground air temperatures are difficult quantities to establish
with any reasonable certainty. The values used in these models are believed to
represent the most reasonable estimates to date. The altitude of the base of the
mcsosphere is a highly uncertain quantity; two estimates are given in Fig. 11.
September 11, 1967 E. Monash, JPL Sec. 5.3, page 9
Lower Atmosphere
JPL 606-1
Latitude
To
Ground air temperature, °K
Ground air t
To
emperature, °K
Noon
Sunset
Midnight
Sunrise
Noon
Sunset
Midnight
Sunrise
+ 90
Summer
W
inte r
235
235
235
235
143
143
143
143
+ 70
260
247
235
220
143
143
143
143
+ 50
265
249
233
217
177
170
162
155
+ 30
265
248
232
215
225
212
198
185
+ 10
260
243
227
210
245
230
21 5
200
-10
245
23C
215
200
260
243
227
210
-30
225
212
198
185
265
248
232
21 5
-SO
177
170
162
155
265
249
233
217
-70
143
143
143
143
260
247
235
220
-90
143
143
143
143
235
235
235
235
+ 90
Fall
S
pring
143
143
143
143
143
143
143
143
+ 70
190
181
171
162
208
197
187
176
+ 50
235
221
206
193
235
220
205
190
+ 30
2 50
234
219
203
250
234
217
201
+ 10
257
241
224
208
260
243
226
209
-10
260
243
226
209
257
241
224
208
-30
2 50
234
217
201
250
234
219
203
-50
235
220
205
190
235
221
206
193
-70
208
197
187
176
190
181
171
162
-90
143
143
143
143
143
143
143
143
Note: Th
lat
e grou
ion and
id air te
should
Tiperatures
le used wit
are de ri\
1 caution.
;ed negl
ecting at
^Tosphe ric
circu-
Fig. 1. Table of ground air temperatures for Mars
referred to northern seasons. Data are the results of
calculations from Leovy (1966) and Neubauer (1966).
Leovy presents the calciilated thermal data in graphic
form, from which the table is derived.
To use Lower Atmosphere Models I, II, and III (Figs. 2 through
10), first refer to Fig. land read the appropriate ground air tem-
perature for the pertinent season, latitude, and time of day. Then
select the profile with the most nearly correct value for ground
air temperature from among the 10 profiles given for each model.
Data for Models I, II, and III are reproduced essentially in the
original computer printout format. The "E" or exponent notation
following each number in the figures indicates the power of ten by
which the number must be multiplied; e.g., . lOOE 06 = . 100 X
10^. A positive exponent is denoted by a blank after the E rather
than by a plus sign. All exponents are positive with the single
exception of those associated with zero altitude. Zero altitude is
always followed on the printout by a large negative exponent, a
characteristic of the computer.
Sec. 5.3, page 10
E. Monash, JPL
September 11, 1967
JPL 606-1
Lower Atmosphere
H
F^ei^ht above
Atmospheric parameters
Atmospheric parameters
T
P
N
T
P
N
mean si'.rface,
Kinetic
Total
Total
Kinetic
Total
To till
cm
tempe rature,
p res su re,
dyne cn-^."^
concent ration,
tempe rature.
pres sure.
concentration,
°K
c m ~ -^
°K
dyne cm
2
cm" 3
n. ooof:- ?8
Ground air temperature Tq: 180°K
Ground air temperature To: 200°K
0. 180000E 03
0. lOOOOOE 05
0.402442E 18
0.200000E 03
0. lOOOOOE
5
0. 3b21 98E 1 8
0. 1 OOE 06
0. 17851 IE 3
0. 900326E 04
0. 365351E 18
0. 194887E 03
0. 908307E
04
0. 337619E 18
0. 200f; 06
0. 1 78431 E 3
0. 81 0835E 04
0. 329183E 18
0. 192773E 03
0. 824143E
04
0. 309b94E 1 8
0. 300f: 06
0. 1 783S1 E 03
0. 730205E 04
0.296582E 18
0. 192373E 03
0. 74793bE
04
0.281642E 18
0.400F: 06
0. 1 78271 E 03
0.657562E 04
0.267ig7E 18
0. 191973E 03
0. 678639E
04
0.25b080E 18
0. 500E 06
0. 1 781 91 E 3
0. 5921 18E 04
0.240712E 18
0. 191 573E 03
0.61 5b3 7E
04
0.232792E 18
0.6 OOE 06
0. 17811 IE 03
0. 5331 61 E 04
0.216842E 18
0. 191173E 03
0. 558371E
04
0.211579E 18
0.700E 06
0. 1 78031 E 03
0.480053E 04
0.195330E 18
0. 190773E 03
0. 50b327E
04
0. 1 92261 E 18
0.«00E 06
0. 1779S1E 03
0.432214E 04
0. 175944E 18
0. 190373E 03
0.459041E
04
0.174b72E 18
0. 900E 06
0. 177871E 03
0. 389125E 04
0. 158474E 18
0. 189973E 03
0.416084E
04
0. 1 58b60E 1 8
0. lOOE 07
0. 177791E 03
0. 350314E 04
0.142733E 18
0. 189573E 03
0. 377070E
04
0. 1 4408bE 1 8
0. 1 50E 07
0. 177391E 03
0. 20701 IE 04
0.845352E 17
0. 187573E 03
0.229753E
04
0.887292E 17
0.200E 07
0. 1 76991 E 03
0. 122184E 04
0. 500079E 17
0. 185573E 03
0. 139249E
04
0. 543569E 1 7
0. 2 50E 07
0. 176S91E 3
0. 720306E 03
0.295477E 17
0. 183573E 03
0. 839399E
03
0.331235E 17
0. 3 OOE 07
0. 17bl91E 03
0.424129E 03
0. 174377E 17
0. 181573E 03
0. 5031'15E
03
0.200753E 17
0. 3S0E 07
0. 175791E 03
0.249435E 03
0. 102787E 17
0. 179573E 03
0. 2 9994b E
05
0. 120998E 17
0.400E 07
0. 175391E 03
0. I46518E 03
0. 605145E 16
0. 177573E 03
0. 177759E
03
0. 7251 5bE lb
0.45OE 07
0. 174991E 03
0. 859603E 02
0.355842E 16
0. 175573E 03
0. 104724E
03
0.4 3208 3E lb
0. 500E 07
0. 174591E 03
0. 503704E 02
0.208991E 16
0. 173573E 03
O.bl 3241 E
02
0.2 5 5',i3 3E lb
0. 550E 07
0. 1710001-; 03
0. 296967E 02
0. 12 5802 E 16
0. 171000E 03
0. 36 1 547E
02
0. 1 531bOE lb
0.600E 7
0, 10 7000 E 3
0. 172906E 02
0.750015E 15
0. lb7000E 03
0. 21 0507E
02
0. 91 51 1 7E 15
0.6 50E 07
0. 1 63000E 03
0. 99361 3E 01
0.44157hE 15
0. 163000E 03
0. 120969E
02
0. 5 57b04E 1 5
0, 700E 07
0. 1 59000E Oi
0. 5631 78E 01
0. 256581 E 15
0. 1 59000E 03
0. b85650E
01
. 3 1 2 3 7 9 E 15
0.7 50E 07
0. 1 S5000E 03
0. 314t)24E 01
0. 147040E 1 5
0. 1 55000E 03
0. 383044E
01
0. 1 7901 bE 1 5
0,800E 07
0. 1 55000E 03
0. 175598E 01
0.820661E 14
0. 1 55000E 03
0. 213784E
01
0.999125E 14
0.850E 07
0. 1 T5000E 03
0. 980045E 00
0.458027E 14
0. 155000E 03
0. 1 19317E
01
0. 5 576 31 E 1 4
0. 900E 07
0. 1 55000E 03
0. 546982E 00
0.255634E 14
0. 1 55000E 03
0.b65932E
00
0. 31 1225E 14
0.000E-3K
Ground air temperature Tq: 190°K
Ground air temperature T,,: 210°K
0. 190000E 03
0. 1 OOOOOE 05
0.381261E 18
0.210000E 03
0. 1 OOOOOE
05
. 3 4 4 9 5 1 E 18
0. lOOE 06
0. 1SS720E 0?
0. 903950E 04
0. 352584E 18
0. 204950E 03
0. 912541E
04
0. 3225 38E 18
0.200E 06
0. 185480E 5
0. KI7378E 04
0. 319230E 18
0.200309E 03
0. 830'>53E
04
0. 30050bE 1 8
0. 300E 06
0. 185240E 03
0. 7 39002 E 04
0.288993E 18
0. 199749E 03
0. 75b849E
04
0. 274474E 1 8
0.400E 06
0. 1 85000E 03
0. 668053f; 04
0. 261 587E 1 8
0. 199189E 03
0. b891 72E
04
0.2 50b34E 18
0. 500E 06
0. 1 84760 E 3
0.f)Oi8 3 7E 04
0.236749E 18
0. 1986 29E 03
0. 627382E
04
0. 22880bE 1 8
. 6 E 06
0. 184520E 03
0. 545721E 04
0.214242E 18
0. 1980o9E 03
0. 570980E
04
0.208825E 18
0.7 OOE 06
0. 1 842K0E 03
0.4931 34E 04
0.193849E IH
0. 197509E 03
0. 519510E
04
0. 1 90 5 39 E 1 8
0.800E 06
0. 184040E 03
0. 445556E 04
0. 175375E 18
0. 196949E 03
0. 472553E
04
. 1 7 3 8 1 E 18
0. 900E 06
0. I 83800E 03
0.40251 5E 04
0. 1 58640E 1 8
0. 19b389E 03
0, 42972 5E
04
0. 1 58508E 1 8
0. 1 OOE 07
0. 183S60E 3
0, 363583E 04
0. 143484E 1 8
0.195829E 03
0. 390b72E
04
. 1 4 4 5 1 5 E 18
0. 1 50E 07
0. 182360E 03
0. 218194E 04
0.866742E 17
0. 193029E 03
0.241bl8E
04
0.90b742E 17
0. 200E 07
0. 181 160E 03
0. 130502E 04
0.521833E 17
0. 190229E 03
0. 148387E
04
0. 5b50b4E 1 7
0.250E 07
0. 179960E 03
0. 777872E 03
0.313119E 17
0. 1 87429E 03
0. 904 74 3 E
3
0. 349b7bE 1 7
0. 300E 07
0. I78760E 03
0.462057E 03
0. 187242E 17
0. 184b29E 03
0. 547544E
03
. 2 1 4 8 3 1 E 17
0. 3 50E 07
0. I77S60E 03
0. 273502E 03
0. 11 1 58 IE 17
0. 181829E 03
0. 32883bE
3
. 1 3 1 7 1-: 17
0.400E 07
0. 1 76360E 03
. 1 6 1 3 1 7 E 3
0.662608E 16
0. 179029E 03
. 1 9 5 ',) 3 1 i-;
03
. 7 9 2 7 ',) 1-; 1 f J
0.4S0E 07
0. 17S160E 03
0. 948056E 02
0. 392081E 16
0.176229E 03
0. 1 1 579 5E
03
47 ^ij7Ay 1 ,,
0. 500E 07
0. 173960E 5
0. 5551 38E 02
0.231169E Ih
0. 173429E 03
0. b78585E
02
. 2834 3 9E lb
0. 5S0E 07
0. 1 71000E 03
0. 327292E 02
0.138648E lb
0. 171000E 03
0. 400072E
02
0. 1 b94 80E 1 b
0.6 OOE 07
0. 16 7000 E 3
0. I90562E 02
0. 826602E 1 5
0. lb7000E 03
0. 2 32 9 58E
02
0. 1 01 041 E lb
0.650E 07
0. lb3000E 03
0. 1 09507E 02
0.486667E 15
0. 1 63000E 03
0. 1 3 3859E
02
0. 5 948 88E 1 5
0.700E 07
0. 1 S9000E 03
0.620686P: 01
0.282782E 15
0. 1 59OO0E 03
0. 758709E
01
. 3 4 5 b b 4 E 15
0.7S0E 07
0. 1 5SOO0E 03
0. 346751 E 01
0. 16205 5E 15
0. 1 55000E 03
. 4 2 3 8 5 ') E
01
0. 1 98091 E 1 5
0.800E 07
0. 1 S5000E 03
0. 19 3529E 01
0.9044blE 14
0. 1 55000E 03
0. 23t,5b4E
01
0.11 0559I-; 1 5
0. 8S0E 07
0. 1 S5000E 3
0. 108012E 01
0.504797E 14
0. 1 55000E 03
0. 132031E
01
0.bl7049f: 14
0. 900E 07
0. 1 55000E 03
0. 602837E 00
0.281737E 14
0. 155000E 03
0. 73b890E
00
0. 344387E 14
Z. Lower Atmosphere Model I for ground air temperatures 180, 190,
and 210°K. Surface pressure 10 mb (0. 10 X 10^ dyne cm'^); atmos-
Fig.
ZOO,
pheric abundance 80% CO2, 10% Ar, 10% N^ by volume;^ mean molecula]
mass of atmospheric constituents 0.697119 X 10-2Z
September 11, 1967
E. Monash, JPL
Sec. 5.3, page 1 1
Lower Atmosphere
JPL 606-1
H
Height above
Atmospheric parameters
Atmospheric parameters
T
P
N
T
F
N
mean surface,
Kinetic
Total
Total
Kinetic
Total
Total
cm
ten^pe rature,
pressure ,
concentration,
tempe rature ,
pres sure.
concent rat i ori ,
°K
dyne cm'2
c m ~ 3
°K
dyne cm-2
cr.]~ ''
0. OOOE-'iS
G round air
temperature T
-y. 220°K
Ground ai r
tempe rature T
,: 240 = K
0.220000E 03
0. lOOOOOE
05
0. 329271E 18
0. 240000E 03
0. 1 OOOOOE OS
0.3O1832I': IH
0. lOOE 06
0. 215012E 03
0. 916401 E
04
0.308744E 18
0. 235131E 03
0. 9231 82E 04
0. 28441 5J-; 18
0.200E 06
0. 2I0024E 03
0.838072E
04
0.28rj060E 18
0. 2302b3E 03
0.850841E 04
. 2 b 7 b 7 1 E IS
0. 300E Oh
0. 207364E 03
0.7b5619E
04
0.267458E 18
0. 225394E 03
0.782803E 04
0.251 ShbE 18
0.400E 06
0. 20b644E 03
0.699S40E
04
0.245225E 18
0.222266E 03
0.7194S0E 04
0.234479E 18
0. 500E 06
0.205924E 03
0. 6 3 8962 E
04
0.224773E 18
0.22122bE 03
O.bbl 512f: 04
0. 21b 544 E 18
0.600E 06
0.205204E 03
0. S8344SE
04
0.20S963E 18
0.22018bp: 03
0.bO7b3OE 04
0. 1 \»990bE 1 8
0.700E 06
0. 2044K4E 03
0. S32S82E
04
0, 18H670E 18
0.21914bE 03
0. 5 58082 E 04
0. 184477E 18
0.800E 06
0. 20 5764E 03
0. 4K599bE
04
0. 17277SE 18
0. 21 810bE 03
0. SI 23b8E 04
0. 1701 7 5E IH
0.900F: 06
0.203044E 03
0.443342E
04
0. 158170E 18
0.2170bbE 03
0.470205E 04
0. 1 Sb'il 8E I 8
0. lOOE 07
0.202324E 03
0.404299E
04
0. 144754E 18
0. 21b02bE 03
0.431335E 04
0. 1 44b 39E 1 «
0. 150E 07
0. 198724E 03
0.2S3724E
04
0.924884E 17
0. 21082bE 03
0.278420F: 04
O.''5bb50f; 17
0.200E 07
0. 193I24E 03
0. 157878E
04
0.S86120E 17
0.20 5b2bE 3
0. 1 777b2E 04
0.b2b2 5 5E 17
0. 2S0E 07
0. 191S24E 03
0. 973742E
03
0.368296E 17
0. 200426E 03
0.1121 98f; 04
0.40SSIbiE 17
0. 300E 07
0. 1 87924E 03
0. S95092E
03
0.229392E 17
0. 1 95226E 03
0.b99044E 03
0.259b07i.: 17
0. 350E 07
0. 184324E 03
0. 360236E
03
0. 141 S73E 17
0. 19002bE 03
0.4307S7E 05
0. lb420'<E 1 7
0.400E 07
0. 1 80724E 03
0.215919E
03
0.86 5467 E 16
0. I 8482b E 3
0. 2blb63I-; 03
0. 1 02 5 5 SI-: 1 7
0.450E 07
0. 177124E 03
0. 128092E
03
0. 52 3 86 5 E lb
0. 1 79026E 03
0. 1 56702 E 3
0.t)31948E lb
0. 500E 07
0. 173524E 03
0.7S1786E
02
0. 31 3842E 16
0. 17442bE 05
0. 924407E 02
0. 383910E lb
0. 550E 07
0. 171000E 03
0.443229E
02
0. 1877b2E lb
0.171 OOOE 03
0. 545001 E 02
0.23087SE lb
0.600E 07
0. 167000E 03
0. 2S8065E
02
0.11 1941E 10
0. lb 7000 E 3
0. 317321 E 02
0. 1 57b441-: lb.
0.650E 07
0. lb3000E 03
0. 148298E
02
0.6S9060E 15
0. 163000E 03
0. 1 82 3 50 E 02
0.810 39(JE 15
0.700E 07
0. 1 5 9000 E OS
0. 8405 S4E
01
0. 382T52E 1 5
0. 1 59000E 05
0. 1 05 iSbi-: 02
0.47 8851-; 1 5
0. 750E 07
0. 155000E OJ
0. 4b95K2E
01
0. 2194bOE 1 5
0. 1 SSOOOE 3
0. S77405E 01
0. 2(j98S1 E 1 S
0. 800E 07
0. 1 S3000E 03
0. 2b2083E
01
0. 12248SE 15
0. 1 SSOOOE 03
0. 3222bl K 01
0. 1 50b091-; 1 5
0. 8 60E 07
0. 1 3S00OE 03
0. 1 4b273E
01
0. bK ibl 5E 1 4
0. 1 SSOOOE 03
0. 1 798bOE 01
0.840SS01-; 14
0. 900E 07
0. 133000E Oi
. 8 1 6 3 8 1 E
00
0. 381 S37F: 14
0. 1 SSOOOE 5
0. 1 003K3E 01
0.4b.J144E 14
O.000E-3H
Ci round ai i
temperature T
y. 2 30-K
Grounc! aii
temperature T
,: 2S0-K
0. 2 50000E 3
0. lOOOOOE
OS
0. 3144SSE 18
O.2S0OOOE 3
0. 1 000001-: OS
0.28<i7S81-: 18
0. 1 OOE Ob
0. 22 5072 E Oi
0.91 993 5E
04
0. 2 9b 081 E 1 8
0. 245189E 03
0. 92bl7bE 04
0.27 3b3 5E 18
0.200E Ob
0. 220I45E 3
0. 84471 9E
04
0.277958E 18
0. 240 578E 03
0. 85b499F: 04
0. 258 1 12I-: 18
0. 300E 06
. 2 1 5 2 1 7 E OS
0. 7 741 SbE
04
0.2b0572E 18
0. 2 55Sb7E 03
0. 7<'081 5E 04
0.243185!-: 18
0.400E 06
0.214i37E 03
0. 7096 S4E
04
0.239842E 18
0. 23075bE 03
0. 7289b3E 04
0.228839E 18
0. 500E 06
0.21 3437E 05
0. b502 94E
04
0.22068bE 18
0.229227E 03
0.b719b7E 04
. 2 1 2 3 5 3 E 18
0.600E 06
0, 212377E 03
O.S9Sb83E
04
0.202 990E 18
0.228027E 03
0.bl'i2 3 5E 04
0. 1 9b 71 8E 18
0. 700E 06
0. 21 lb97E Oi
0. 545461 E
04
0. 18b 049 E 1 8
0.22b827E 03
0. 570394E 04
0. 1 82 1 bl E 18
0. 800E 06
0. 2 1081 7E 3
0.499289E
04
0. 171 Sb3E 18
0.225b27E 03
0. S25I 77E 04
0. Ib8bl 3E 18
0. 900 F. 06
0. 209937E 03
0.45b8S7E
04
0. 157640E 18
0. 224427E 03
0. 4833 31E 04
'. 0. 1 Sb007E 1 8
0. 1 OOE 07
0. 2090S7E 03
0.417875E
04
0. 1 4479faE 18
0.223227E 03
0.444b22E 04
0. 144284E 18
0. 150E 07
0.204b 57 E 3
0. 26b010E
04
0. 941 5 58 E 17
0. 217227E 03
0. 290'101E 04
0. 970081 E 17
0.200E 07
0.200237E 03
0. lb7682E
04
O.bOoSblE 17
0.211227E 03
0. 188079E 04
0.b4S010E 17
0, 2 50E 07
0. 1 958S7E 03
0. 1046 22 E
04
0. 38r)954E 1 7
0. 205227E 03
0. 120082E 04
. 4 2 3 8 S 5 [■; 17
0. 300E 07
0. 1914S7E 05
0. b4S808E
03
0.244347E 17
. 1 9 't 2 2 7 E 3
0.7S(jt57E 03
0.275079E 17
, ; 5 E 07
0. 187057E 03
0. 3<1419bE
03
0. 1 52b56E 17
0. 193227E 03
0. 4()9',)45E 5
0. 17bl7't!-: 1 7
0.400E 07
0. lK2i)37I': 03
0. 23780 5E
03
0.943099E 16
0. 187227E 03
0.287Sb',)E 03
0. 1 1 1 2b 3 E 17
0. 450E 07
0. 1 7K2S7E 03
0. 1 41 701 E
03
0. 57 58 38 E 16
0. 181227E 03
0. 1 731 7b E 5
0.b922 14E lb
0. SOOE 07
0. 17 5KS7E 03
0.833499E
02
0. 34 7287 E lb
0. 175227E 03
0. 102S22E 03
0.423831K lb
0. SSOE 07
0.171 OOOE 03
0.491404E
02
0.208170E 16
0. 171000E 03
0.b04437E 02
0. 2SbOS4E 1 b
O.OOOE 07
0. 1I)7000E 5
0. 28b 1 1 SE
02
0. 124108E lb
0. lb7000E 03
0. 3S1927E 02
0. 1 52bSbF; 1 b
0.650E 07
0. 1()3000E 03
0. lb4417E
02
0.730694E 15
0. lb3000E 03
0.202237E 02
0.8<tH770F; IS
0.700E 07
0. 1 39000E 03
0. 931914E
01
0.424575E IS
0. 1 S9000E 03
0. 114b27E 02
0. 5222 37E 15
n. 760E 07
0. 1 S3000E 03
0. S20b21E
01
0.243314E IS
0. 1 SSOOOE 03
0.b4037(jE 01
0. 29^)281 E 1 S
0. 800 [■: 07
0. 1 SSOOOE 03
0. 2 90S69E
01
0. 1 5 5798E 15
0. 155000E 03
0. 35740l)E 01
0. lb703SE 1 S
. K T F-: 07
0. 1 SSOOOE 03
0. 162172E
01
. 7 5 7 9 1 5 E 14
0. 1 SSOOOE 3
0. 19947 5 E 01
0. <) 322S2E 14
0. 900E 07
0. 1 SSOOOE 3
0.9051 14E
00
0.423007E 14
0. 1 SSOOOE 03
. 1 1 1 3 3 1 1-: 01
0. 52 5 08I-: 1 4
-
Fig. 3. Lower Atmosphere Model I for ground air temperatures 220, 230,
240, and 250°K. Surface pressure 10 mb (0. 10 X 10^ dyne cm"^); atmos-
pheric abundance 80% CO^, 10% Ar, 10% N2 by volume; mean molecular
mass of atmospheric constituents 0.697 119 X 10-22 gj^.
Sec. 5.3, page 12
E. Monash, JPL
September 11, 1967
JPL 606-1
Lower Atmosphere
II
Ilci^lU a
>OVC
AliiK
)sph(^ i-ic paramo
■cry
J
\tnu
).spheric pa i
amel ery
r
P
N
r
F
ni(-an sur
face,
Kint'Uc
Total
Total
Kinetic
Total
T'otai
cm
trl npc ralu I'c .
pvvssyit'v ,
concent rat
on,
teiiijjo i-atiire.
pres sure
c tmcent rat
or-,.
- i«
C'jr(^vinfi ai
flync c m " *-
en-*
°K
d ai
(lynv cm"
'
c n 1 " ^
O.OOOK
r tompr ratu re T
„: 270°K
CirtJun
r tein|3e ratu re T
„: 2'U)'K
O.270000F:
03
0. 1 00000 f: 5
0.268295E
18
0. 290000E
3
0. 1 00000 f:
05
0.24 >7')2F
18
0. lOOK
Ofc
0.26S!00K
03
0.9 31 51 4E 04
0.2 54348 E
18
0. 28 540b E
Oi
0. '.nt,! i2F:
04
0. 2 57b02K
18
0. 200K
06
0, 260600E
03
O.Ht)6619E 04
0. 240896 K
18
0. 280M12E
Oi
0. 875405 F
04
().225824E
18
0. 500 f:
06
0. 2S590OE
03
0. 80 51 Kb E 04
0.227930E
18
0.27b2IKE
3
U.8I7712F:
04
0. 21444'tE
18
0.400F:
06
0. 251200F
03
0.7470K9F; 04
0.215441E
18
0.271b24F
03
0. 7b2949E
04
0. 205471 F
18
0. 500F,
06
0.246 500 F:
03
0.b92203E 04
0.203420E
18
0. 2670 iOE
3
, 7 1 1 1 2 f:
04
0. I'!28S2F
18
0. 600E
06
0.244402F
03
0.b4l275E 04
0. I 90071 E
18
0. 2b24i6E
Oi
0. bbl 800E
04
0. I 82b7 5F:
18
0.700E
06
0. 242KK2F
Oi
0. 59i934E 04
0. I77I4I E
18
0. 259845E
3
O.bl 58b2F
04
0. 1 71b90F;
18
0. HOOK
Ob
0. 241 56 2 K
3
0. 549824F 04
0. 1650I8E
18
0.258005E
03
0. 572985F;
04
0. lbOS7bE
1 8
0. 900K
Ot>
0. 239H42F
0!
0. 508741 E 04
0. 153656E
18
0. 256165E
03
. 5 3 2 8 I 8 F
04
0. I50b7iF
i 8
0. 1 oof:
07
0.2?Ki22E
Oi
0.470496E 04
0. 14301 1 E
18
0.254325E
05
0.495208E
04
0. 14I050F:
1 8
0. 1 5of:
07
0.2 i0722K
03
0. 31 5891 F 04
0. 991802E
17
0.24 51 2 5E
3
0. i40642F
04
0. 1 O0t>b7F;
18
0. 200 f:
07
0.225122F;
Oi
0. 209277E 04
0.679448E
17
0.235925E
3
0. 2 i0989F
04
0. 709240F:
17
0. 2S0E
07
0.21 5 522F:
OS
0. 1 366HIE 04
0.459404E
17
0.226725E
3
0. 1 542 i2E
04
0.492777E
17
0. soof;
07
0. 207922F
3
0. K791 32E 03
0. 306289E
17
0.217525E
03
0. 101 272 E
04
0. 3 57252E
17
0. 5 50F,
07
0.200322F
3
0. 556239E 03
0.201 USE
17
0.208325E
3
0.6 529'U E
5
0.227061 K
17
0. 400F
07
0. 192722F
3
0. 345764E 03
0. 129965E
17
0. I99125E
03
0.412778E
05
0. 1 501 b4F
17
0.450I-:
07
0. 1H5I22E
Oi
0.210857E 03
0.825101E
lb
0. 189925F;
Oi
0. 25 53 iOF
3
0. ',>738 5')E
lb
0. 500 f:
07
0. I77522E
03
0. I25948E 03
0. 513945E
16
0. 180725E
03
0. 1 5421 5E
03
0.bl8I 37F
It,
0. SWF.
07
0. 171000E
03
0.742549E 02
0. 314561E
16
0. 171000E
3
0.909203E
02
0. 3851 bOE
lb
0. 600 F
07
0. 167000E
Oi
0.432342E 02
0. 187537E
16
0. I67000E
3
0. 529374E
02
0. 22 'tb2 7E
lb
O.dSOK
07
0. lb3000E
03
0.248447E 02
0.110414E
16
0. 163000E
03
0. i04207F
02
0. I i5l 94 K
lb
0. TOOK
07
0. 1S9000E
03
0. 1408I9E 02
0.641 567E
15
0. 15 9000 E
03
0. I72424E
02
0. 78 5 5 5b E
1 5
0. 750E
07
0. 1 S5000E
03
0.786699K 01
0. 367666E
15
0. 155000E
03
0.9632b2E
01
0.4501«iE
1 5
0. «00F':
07
0. I 55OO0E
03
0.43g072E 01
0.205201E
15
0. 1 55000F:
03
0. 5 3761 5E
01
0. 251 25bE
1 5
0. 8tOE
07
0. 155000E
03
0.245055E 01
0. U4527E
15
0. 155000E
03
0. 300053E
01
0. 140231 E
1 5
0. 900F
07
0. 1 SSOOOE
03
0. 13b770E 01
0.639I97E
14
0. 1550OOE
03
0. 1 b746f>E
01
0. 782(.54F:
14
Fig. 4. Lower Atmosphere Model I for ground air temperatures 270 and
Z90°K. Surface pressure 10 mb (0. 10 X 10^ dyne cm"2); atmospheric
abundance 80% CO2, 10% Ar , 10% N2 by volume; mean molecular mass
of atmospheric constituents 0.697119 X lO"^^ gm .
September 1 1, 1967
E. Monash, JPL
Sec. 5.3, page 13
Lower Atmosphere
JPL 606-1
H
Height above
Atmospheric parame
ters
Atmospheric parame
er s
T
P
N
T
P
N
mean surface,
Kinetic
Total
Total
Kinetic
Total
Total
cm
temperature, |
pressure ,
concent ration, 1
tetiipe raturc, |
pressure.
concent ration.
°K
dyne cm"^
cm-'
= K
dyne cm^'^
cm-3
0. OOOE-38
Ground air temperature T
o: 180-K
G rounc
air
temperature Tot 200°K
0. 180000E
03
0. lOOOOOE
05
0. 402442E
18
0. 200000E
03
0. lOOOOOE
05
0. 362198E 18
0. lOOE 06
0. 178519E
03
0.904843E
04
0.3fa7169E
18
0. 194885E
03
0.912476E
04
0. 339171E 18
0.200E 06
0. 178439E
03
0. 818978E
04
0. 332475E
18
0. 192771E
03
0.831767E
04
0. 312563E 18
0. 300E 06
0. 178359E
03
0.741230E
04
0.301048E
18
0. 192371E
03
0.758351E
04
0. 285567E 1 8
0. 400E 06
0. 178279E
03
0.670832E
04
0.272578E
18
0. 191971 E
03
0.691281E
04
0.260853E 18
0. SOOE 06
0. 178199E
03
0.607092E
04
0.246789E
18
0. 191571E
03
0.630022E
04
0.238234E 18
0.600E 06
0. 1781 19E
03
0.549385E
04
0.223431E
18
0. 191 171E
03
0. 574080E
04
0.217534E 18
0.700E 06
0. 178039E
03
0.497140E
04
0.202274E
18
0. 190771E
03
0.523003E
04
0. 198595E 18
0.800E 06
0.177959E
03
0.449844E
04
0. 1831 13E
18
0. 190371E
03
0.476377E
04
0. 181271E 18
0.900E 06
0. I77879E
03
0.407029E
04
0. 165759E
18
0. 189971E
03
0.433823E
04
0. 165426E 18
0. lOOE 07
0. 177799E
03
0. 368272E
04
0. 150043E
18
0. 189571E
03
0. 394992E
04
0. 150936E 18
0. 1 50E 07
0. 177399E
03
0.223149E
04
0.911213E
17
0. 187571E
03
0.246417E
04
0.951662E 17
0. 200E 07
0. 176999E
03
0. 135060E
04
0. 552757E
17
0. 185571E
03
0. 152953E
04
0. 597069E 17
0. 250E 07
0. 176599E
03
0.816524E
03
0. 334933E
17
0. 183571E
03
0. 944494E
03
0. 372711E 17
0. 300E 07
0. 176199E
03
0.493074E
03
0. 202715E
17
0. 181571E
03
0. 580I59E
03
0. 231461E 17
0, 350E 07
0. 175799E
03
0.297413E
03
0. 122552E
17
0. 179571E
03
0. 354447E
03
0. 142986E 17
0.400E 07
0. 175399E
03
0. 179187E
03
0. 740043E
16
0. 177571E
03
0. 215357E
03
0.878543E 16
0.450E 07
0. 174999E
03
0. 107833E
03
0.446368E
16
0. 175571 E
03
0. 130111E
03
0. 536832E 16
0. 500E 07
0. 174599E
03
0.648174E
02
0. 268922E
16
0. 173571E
03
0.781560E
02
0. 326184E 16
0. S50E 07
0. 171000E
03
0.391880E
02
0. 166009E
16
0. 171000E
03
0.472524E
02
0.20O172E 16
0.600E 07
0. 167000E
03
0.234121E
02
0. 101555E
16
0. 167000E
03
0.282301E
02
0. 122454E 16
0.650E 07
0. 163000E
03
0. 138135E
02
0.613894E
15
0. 1630O0E
03
0. 166562E
02
0.740226E 15
0.700E 07
0. 159000E
03
0.804407E
01
0. 366484E
15
0. 159000E
03
0.969944E
01
0.441902E 15
0. 750E 07
0. 155000E
03
0.462023E
01
0,215927E
1 5
0. 155000E
03
0. 557101E
01
0. 260363E 1 5
0.800E 07
0. 155OO0E
03
0.265126E
01
0. 123907E
15
0. 155000E
03
0. 319685E
01
0. 149406E 1 5
0.850E 07
0. 155000E
03
0. 152139E
01
0.711 024E
14
0. 155000E
03
0. 183447E
01
0. 857345E 14 !
0.900E 07
0. 155000E
03
0.873028E
00
0.408012E
14
0. 15500OE
03
0. 1 05269E
01
0.491976E 14
0. OOOE-38
Groun
d ai
r temperature T
o: 190°K
Groun
d air temperature T
o: 210°K
0. 190000E
03
0. lOOOOOE
05
0. 381261E
18
0.210000E
03
0. lOOOOOE
05
0. 344951 E 18
0. lOOE 06
0. 185730E
03
0.908310E
04
0. 354266E
18
0.204935E
03
0. 916524E
04
0. 323969E 18
0.200E 06
0. 185490E
03
0.825272E
04
0.322295E
18
0.200280E
03
0.838302E
04
0. 303207E 18
0. 300E 06
0. 185250E
03
0.749733E
04
0.293174E
18
0. 199720E
03
0.766937E
04
0.278173E 18
0.400E 06
0. 185010E
03
0.681023E
04
0.266651E
18
0.199160E
03
0.701472E
04
0.255143E 18
0. 500E 06
0. 184770E
03
0.618533E
04
0.242498E
18
0. 198600E
03
0.641433E
04
0.233964E 18
0.600E 06
0. 184530E
03
0. 561707E
04
0.220506E
18
0. 198040E
03
0.586385E
04
0.214490E 18
0. 700E 06
0.184290E
03
0. 510037E
04
0.200483E
18
0. 197480E
03
0.535925E
04
0. 196588E 18
0.800E 06
0. 184050E
03
O.463062E
04
0. 182255E
18
0. 196920E
03
0.489683E
04
0. 1801 36E 1 8
0.900E 06
0. 183810E
03
0.420361E
04
0. 165665E
18
0. 196360E
03
0.447315E
04
0. 165020E 18
0. lOOE 07
0.183570E
03
0.381548E
04
0. 150565E
18
0. 195800E
03
0.408507E
04
0.1511 34E 18
0. 150E 07
0. 182370E
03
0.234618E
04
0. 931933E
17
0. 1930OOE
03
0.258479E
04
0.9701b3E 17
0. 200E 07
0. 181170E
03
0. 143806E
04
0. 575001E
17
0. 190200E
03
0. 162460E
04
0.618746E 17
0. 250E 07
0. 179970E
03
0.878581E
03
0.353638E
17
0. 187400E
03
0. 101409E
04
0.391997E 17
0. 300E 07
0. 178770E
03
0. 535001E
03
0. 216789E
17
0. 184600E
03
0.628528E
03
0.246643E 17
0. 350E 07
0. 177570E
03
0.324695E
03
0.132460E
17
0. 181800E
03
0. 386721E
03
0. 154092E 17
0.400E 07
0. 176370E
03
0.196393E
03
0.806639E
16
0. 179000E
03
0.236155E
03
0.955698E 16
0.450E 07
0.175170E
03
0. U8382E
03
0,489558E
16
0. 176200E
03
0. 143093E
03
0. 588288E 16
0. 500E 07
0. 173970E
03
0.711108E
02
0.29bl00E
16
0. 173400E
03
0.860115E
02
0. 359322E 16
0. 550E 07
0. 171000E
03
0.429929E
02
0. 182128E
16
0. 171000E
03
0.520017E
02
0.220291E 16
0.600E 07
0. 167000E
03
0.256853E
02
0, 11 I415E
16
0. 167000E
03
0.310674E
02
0. 1 34761E lb
0.650E 07
0. 163000E
03
0. 151 548E
02
0.673500E
15
0. 163000E
03
0. 183303E
02
0. 814626E 1 5
0. 700E 07
0. 159000E
03
0.882511E
01
0.402068E
15
0. 159000E
03
0. I06743E
02
0. 48631 7E 1 5
0.750E 07
0. 155000E
03
0. 506883E
01
0. 236893E
15
0. 155000E
03
0.613095E
01
0.286532E 15
0.800E 07
0. 1550OOE
03
0. 290868E
01
0. 135938E
1 5
0. 155000E
03
0. 351817E
01
0. 164422E 1 5
0. 850E 07
0. 155000E
03
0. 166911E
01
0.780062E
14
0. 15500OE
03
0.201885E
01
0.943516E 14
0. 900E 07
0. 155000E
03
0.957795E
00
0.447628E
14
0. 1550O0E
03
0. 1 15849E
01
0. 541424E 14
Fig. 5. Lower Atmosphere Model II for ground air temperatures 180, 190,
200, and 210°K. Surface pressure 10 mb (0. 10 x 10^ dyne cm'^); atmos-
pheric abundance 60% CO2, 20% Ar, 20% N2 by volume; mean molecular
mass of atmospheric constituents 0.663921 X 10"22 gm.
Sec. 5.3, page 14
E. Monash, JPL
September 11, 1967
JPL 606-1
Lower Atmosphere
H
Height above
Atmospheric parameters
Atmospheric parameters
T
P
N
T
P
N
mean surface.
Kinetic
Total
Total
Kinetic
Total
Total
cr7l
tenipe rature,
pressure
concentration,
tempe rature,
pre ssure
,
concentration,
°K
dyne cm"
}
cm" ^
°K
dyne cm"
2
cm" 5
0. 000F:-3H
G rour
d air temperature 1
o: 220°K
Groun
d air temperature T
o: 240°K
0.220000E
03
0. lOOOOOE
05
0.329271E 18
0.240000E
03
0. lOOOOOE
05
0. 301832E 18
0. lOOE 06
0.2149H4E
03
0.920214E
04
0. 310069E 18
0.235080E
03
0. 92t)695E
04
0.285560E 18
0. 200E 06
0.209969E
03
0.845134E
04
0.291573E 18
0.230160E
03
0. 857382E
04
0.269849E 18
0. 300E 06
0. 207296E
03
0.775381E
04
0.270957E 18
0.225240E
03
0. 791922E
04
0. 254691 E 18
0.400E 06
0.206576E
03
0.71 1490E
04
0.249497E 18
0.222083E
03
0.730716E
04
0.238347E 18
0. 500E 06
0.205856R
03
0.632669E
04
0.229670E 18
0.221043E
03
0.674323E
04
0.220987E 18
0.600E 06
0.20S136E
03
0.598529E
04
0.211358E 18
0.220003E
03
0.622047E
04
0.204819E 18
0. 700E 06
0. 204416E
03
0.S48713E
04
0. 194449E 18
0.218963E
03
0. 573604E
04
0. 189766E 18
0.800E 06
0. 203696E
03
0. 502889E
04
0. 178840E 18
0.217923E
03
0. 528730E
04
0. 175755E 18
0.900E 06
0. 20297bE
3
0.460750E
04
0.164436E 18
0.216883E
03
0.487176E
04
0. 162718E 18
0. lOOE 07
0.2022S6E
03
0.4220IOE
04
0. 151146E 18
0.21 5843E
03
0.44871 IE
04
0. 1 50593E 18
0. 15 OK 07
0. I 986S6E
03
0. 270740E
04
0. 987249E 17
0. 210643E
03
0.295632E
04
0. 101667E 18
0. 200E 07
0. 195056K
3
0. 172289E
04
0.639842E 17
0.205443E
03
0. 192755E
04
0.679659E 17
0. 250E 07
0. I914S6E
03
0. I087I9E
04
0.411349E 17
0.200243E
03
0. 124308E
04
0.449694E 17
0. 300E 07
0. I87K56E
03
0. 680072E
03
0.262244E 17
0. 195043E
03
0. 792462E
03
0.294323E 17
0. 3 50E 07
0. 1842S6E
03
0.421564E
03
0. 165736E 1 7
0. 189843E
03
0.499085E
03
0. 190439E 17
0.400E 07
0. I806S6E
03
0.2 58866 E
03
0. 103800E 17
0. 184643E
03
0. 310308E
03
0. 121741E 17
0.450E 07
0. 1770S6E
3
0. 157406E
03
0.644002E 16
0. 179443E
03
0. 190333E
03
0. 768360E lb
0. 500E 07
0. 173456E
03
0.947395E
02
0.395656E 16
0. 174243E
03
0. 115079E
03
0.478427E 16
0. 550E 07
0. 171000E
03
0. 572786E
02
0.242645E 16
0.171000E
03
0.b95754E
02
0.294738E 16
0. 600E 07
0. 167000E
03
0. 342200E
02
0. 148436E 16
0. 167000E
03
0.41 5666E
02
0. 180303E lb
0.6S0E 07
0. 163000E
03
0. 201904E
02
0.897291E 15
0. 163000E
03
0.245250E
02
0. 108993E lb
0. 700E 07
0. 159000E
03
0. 1 1 7575E
02
0.535666E 15
0. 159000E
03
0. 142817E
02
0.650b6oE 15
0. 750E 07
0, 1 SSOOOE
03
0.b75309E
01
0. 315607E 15
0. 155000E
03
0.820289E
01
0.383364E 15
0. 800E 07
0. 1 55000E
3
0. 3875I7E
01
0. 181107E 1 5
0. 155000E
03
0. 470712E
01
0. 219988E 1 5
0. 8S0E 07
0. 1 S5000E
03
0. 222372E
01
0. 103926E 15
0. 155000E
03
0.2701 12E
01
0. 126237E 1 5
0. 900E 07
0. 1 5 5000 H
03
0. 1 27605E
01
0. 596365E 14
0. 155000E
03
0. 1 55000E
01
0. 724396E 14
0.000E-3K
Groun
d ai
r tenipe rature T
o: 230-K
Groun
d air temperature T
o: 250°K
0.23O0OOE
03
0. lOOOOOE
05
0.314955E 18
0. 250000E
03
0. lOOOOOE
05
0.289758E 18
0. lOOE 06
0.22S033E
03
0.923592E
04
0.297311E 18
0.245126E
03
0. 929556E
04
0.274702E 18
0.200E 06
0.22006SE
03
0.K51 510E
04
0.280294E 18
0.240252E
03
0.862808E
04
0.2b0149E 18
0. 300E 06
0.215098E
3
0. 783598E
04
0.263897E 18
0.235379E
03
0.799631E
04
0.246093E 18
0.400E 06
0. 214218E
3
0. 721258E
04
0.243900E 18
0.230505E
03
0. 739901 E
04
0. 232525E 18
0. 500E 06
0. 21 33!KE
03
0. 663652E
04
0.225345E 18
0.228971E
03
0.684640E
04
0.216600E 18
O.bOOK 06
0. 21 24 58 K
03
0.610436E
04
0.208134E 18
0.227771E
03
0.633318E
04
0.201419E 18
0. 700E 06
0.211 578E
3
0. 561 2 92 E
04
0. 192174E 18
0.226571E
03
0. 585602E
04
0. 187230E 18
0. 800 t: 06
0.210698E
3
0. 51 5925E
04
0. 177379E 18
0.225371E
03
0. 541 2 56 E
04
0. 173973E 18
I 0. 900 E 06
0. 20981 KE
03
0. 474057E
04
0. 163668E 1 8
0.224171E
03
0. 500057E
04
0. 161591E 18
j 0. 1 OOF 07
0. 208938E
03
0.435432E
04
0. 150966E 18
0.222971E
03
0.461799E
04
0. 150031E 18
0. 1 50E 07
0. 204S38E
03
0.283142E
04
0. 100278E 18
0.216971E
03
0. 308161E
04
0. 102885E 18
0. 200E 07
0.2001 38E
03
0. 182400E
04
0.660194E 17
0.210971E
03
0.203318E
04
0.698118E 17
0. 2 50E 07
0. 195738E
03
0. 1 16358E
04
0.430625E 17
0.204971E
03
0. 132545E
04
0.468432E 17
0. 300E 07
0. 191 338E
03
0.734739E
03
0.278169E 17
0. 198971E
03
0.853157E
03
0. 310610E 17
0. 3 50E 07
0. 186938E
03
0.459011E
03
0. 177870E 17
0. 192971E
03
0. 541796E
03
0.203386E 17
0,400E 07
0. 182538E
03
0,283561E
03
0.1i2530E 17
0. 186971E
03
0.339168E
03
0. 131406E 17
0.450E 07
0. I7H1 38 K
03
0. 173127E
03
0.704020E 16
0. 180971E
03
0.209101E
03
0. 836998E 16
0. 500E 07
0. 173738E
03
0. I04406E
03
0.435318E 16
0. 174971E
03
0. 126828E
03
0. 525082E 16
0. 5S0E 07
0. 171000E
03
0.631227E
02
0.267403E 16
0. 171000E
03
0, 766792E
02
0. 324831E 16
0.600E 07
0. 167000E
03
0. 3771 15E
02
0. 163581 E 16
0. 167000E
03
0.458106E
02
0. 198713E 16
0.6S0E 07
0, 163000E
3
0,222504E
02
0.988841E 15
0. 163000E
03
0.270290E
02
0, 120121E 16
0. 700E 07
0. 159000E
03
0. 129571E
02
0. 590321E 1 5
0. 159000E
03
0. 157399E
02
0. 717I01E 15
0. 750E 07
0. 155000E
03
0. 74421 IE
01
0.347809E 15
0. 1 55000E
03
0.904042E
01
0.422506E 15
0. 800E 07
0. 1S5000E
03
0.427056E
01
0. 199586E 15
0. 155000E
03
0. 518772E
01
0.242450E 15
0. H50E 07
0. 155000E
03
0.245060E
01
0. 114530E 15
0. 1 55OO0E
03
0. 297691 E
01
0. 139126E 15
0. 900E 07
0. 1 55000E
03
0. 140625E
01
0.657213E 14
0. 155O00E
03
0. 17082bE
01
0. 798359E 14
Fig. 6. Lower Atmosphere Model II for ground air temperatures 220, 230,
240, and 250°K. Surface pressure 10 mb (0. 10 X 10^ dyne cm-2); atmos-'
pheric abundance 60% CO^, 20% Ar , 20% N2 by volume; mean molecular
mass of atmospheric constituents 0.663921 X 10-^2 gm _
September 11, 1967
E. Monash, JPL
Sec. 5.3, page 15
Lower Atmosphere
JPL 606-1
H
Height above
Atmospheric parame
ters
Atmospheric parame
ters
T
P
\
T
P
V
mean sur
ace,
Kinetic
Total
Total
Kinetic
Total
Total
cm
tempe rature ,
pressure.
concei^t ration,
tempe rature,
pres sure.
concenti'ation.
°K
dyne cm"
?
cm ~ ^
°K
dyne cm"
cm" ^
0. OOOE-
58
G roun
1 air temperature T
o: 270 = K
Groun
1 air temperature T
o: 290-K
0. 270000E
3
0. lOOOOOE
05
0.2b8295E
18
0.290000E
03
0. lOOOOOE
05
0.249792E 18
0. lOOE
06
0. 265216E
03
0. 934657E
04
0.25528bE
18
0.285303E
03
0. 939068E
04
0.238433E 18
0. 200E
06
0.260433E
03
0.872509E
04
0.242b89E
18
0. 280607E
03
0.880929E
04
0.227415E 18
0. 300E
06
0.2S5649E
03
0.81345faE
04
0.230497E
18
0.275910E
03
0.825499E
04
. 2 1 b 7 3 3 E 18
0.400E
06
0. 2S08fa6E
03
0.75739bE
04
0.218705E
18
0.271213E
03
0.772b94E
04
0.20b382E 18
0. 500E
06
0. 246082E
3
0. 704231 E
04
0.207 50bE
18
0. 2b6517E
03
0.722432E
04
0. 1 9b 3 58 E 18
O.bOOE
06
0.243969E
03
0.654712E
04
0, 1 94398E
18
0.2t)1820E
03
0.674632E
04
0. 1 8bb55E 18
0.700E
06
0.242449E
03
0.60a519E
04
0. 181815E
18
0.250201E
03
0.629851E
04
0. 1 7b02bE 1 8
0.800E
06
0.240929E
3
0. 563325E
04
0. lb9975E
18
0, 2 57 361 E
03
0. 58791 7E
04
0. 165481E 18
0.900E
06
0.239409E
OS
0. 524953E
04
0. 158839E
18
0. 255521E
03
0. 548504E
04
0. I55499E 18
0. lOOE
07
0.237889E
03
0.487233E
04
0. 148i68E
18
0. 2 5 3681 E
03
0. 51 147bE
04
0. 1 4b054E 1 8
0. 150E
07
0.230289E
Oi
0.3331b0E
04
0. 104799E
18
0. 244481 E
03
0. 357827E
04
0. 10b024E 18
0. 200E
07
0.222689E
03
0.224920E
04
0. 731b53E
17
0. 235281 E
03
0. 246927E
04
0.7b0254E 17
0. 250E
07
0.215089E
3
0. 149788E
04
0. 504470E
17
0. 226081 E
03
0. 1 b7895E
04
0. 537'^blE 17
0. 300E
07
0. 207489E
03
0. 98 5048E
03
0. 54i207E
17
0.21b881E
03
0. 112343E
04
0. 375234E 1 7
0,350E
07
0. 199889E
3
O.63510bE
03
0. 2}01b2E
17
0. 207681 E
03
0. 738739E
03
0.257b74E 17
0.400E
07
0. 192289E
03
0.403426E
03
0. I51980E
17
0. 198481 !•;
03
0, 4766 54 e;
03
0. 175157E 17
0.450E
07
0. 184689E
03
0. 251612E
03
0. 98b885E
lb
0. 1 89281 E
03
0. 301 193E
03
0. 1 1 52b9E 1 7
0. 500E
07
0. 177089E
05
0. 1 5 384 5E:
03
0. b293I3E
lb
0. 1 80081 K
03
0. 1 86025E:
03
0. 748 30 5 E lb
0.550E
07
0. 171000E
03
0.930129E
02
0.394024E
lb
0. 171000E
03
0. 1 124b9E
3
0.476444E lb
0.600E
07
0. 167000E
03
0. 555688E
02
0.24104IE
Id
0. lb7000E
03
0.b71923E
02
0.291460E lb
0.650E
07
0. 163000E
03
0.3278b5E
02
0. 14S708E
lb
0. lb3U00E
03
0. 3 9b44bE
02
0. 1 7bl 8bE lb
0.700E
07
0. 159000E
03
0. 190927E
02
0. 869852E
15
0. 1 5 90OOE
03
0.250863E
02
0. 1 05180E lb
0.750E
07
0. 155000E
03
0. 109661E
02
0. 512505E
15
0. 1550OOE
3
0. 132b00E
02
0.bl9707E 15
0.800E
07
0. 155000E
03
0.629277E
01
0.294094E
15
0, 155000E
03
0.760905E
01
0. 35 561 IE 15
0.850E
07
0. 155000E
03
0.36I103E
01
0. 168762E
15
0. 155000E
05
0.43b635E
01
0.2040b3E 15
0.900E
07
0. 155000E
03
0.207214K
01
0. 9b841 9E
14
0. 155000E
5
0.250557E
01
0. 1170',)9E 15
Fig. 7. Lower Atmosphere Model 11 for ground air temperatures 270 and
290°K. Surface pressure 10 mb (0. 10 X 10^ dyne cm-^); atmospheric
abundance 60% CO2 , 20% Ar, 20% N^ by volume; mean molecular mass
of atmospheric constituents 0.663921 X 10-^2 gm .
Sec. 5.3, page 16
E. Monash, JPL
September 11, 1967
JPL 606-1
Lower Atmosphere
.Atlr
nspho ric pa ram
• t e 1- .s
Atm
iisoiic ric na
ran-!
9 e r s
n
n
,
Kriylil .lOi.vr
1
1'
N
T
1'
K
11 ic.in ,su ri:i «•(■ ,
Kind ic
4'<it,ii
Tol.-U
f\i DftiC
Total
Tdt.il
^' M
('■Ml pc r.i 1 11 r-f ,
[) res.so r
riinccnt ration,
tfOipO f.itu l-C ,
p vv s sun
c once [-it ra
ion.
'K
dync cm
^
cm' '
"K
tiyne c-m
-^
cm" .
(J. 00 K- iK
Or'iiH^fl .1
1- t<-m]i<- loilurc 4
■<-,: 1.40 K
Ciriiunrl ,"ii
r t(>mpc raturc 1
o: 200" K
H
0. 1 KOOOOi-: (
0. 1 tOOOOI
: s
0.O03OO4K l«
0. 2000001.: i
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0. SOK7S7F
1 8
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0. I22H47I
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0.49«71 3K 18
0. 192771 E 03
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2
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Cr-iurul air t<_' r i ipe loitu rt' T
c,: 190"K
Cifounfl ,iir tcmpc r<-iturc T
,: 210'K
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0. i2SI8JF; 17
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0. '012 7') iF:
3
0. it.9',(bSE
1 7
0. 5S0K 07
(1. 1 77t7(1|.: i 0. 4.S704iI.:
i
0. I9K089E 17
0. 181 800].: i
1) . SK0082 i.:
i
0. .' i 1 1 i8|.-
1 7
o.4()or: 07
0. 1 7r, i7()I.: i
0. 2 04 S 90 1.:
i
0. I2099OE 17
0. 1 79000E i
0. iS42 i iF:
11 5 0. 14 i iSSF;
1 7
0. 4S01.: 7
0. 1 7SI 701.: i
. 1 7 7 =. 7 i !•:
i 0. 7 34 3 3 of: 1 1) 1
0. 1 7n200F: i '0.21 4b40F:
3 . 8 K 2 4 i ,' F'
1 b
0. TOOK 07
. 1 7 i 7 1.: i
0. 1 OOI, r, 1-:
Oi
0.4441 50 E It)
0.17 5400]-: i 0.1 2'i01 7f:
(1 i
. s i s - * 8 3 f:
1 1.
0. SSOK 07
0. 171 oooj.; i
0, n4-tK94F.:
2
n.273192E 16
0.171 OOOF: i 0. 780O2SF:
2
11. 3 30437]-:
1 b
o.^ooi^: 117
0. 1 1.70001.: i
0. !SS2H0i.:
2
U . 1 11 7 1 2 i E 16
0. 1 t,700(lF: i 0. 4i)r,01 IK
02
0. 211.' ] 4.' ].-
1 (J
(i.'iSoi-: 07
0. 1 1, iooof.: i
0.227i22I.:
02
0. 10102SE llJ
0. 1 n iooiiF: 1 0. 27.io-iS].:
02 0.122 1 -UF-
1 r.
0. 70()l-: 07
0. 1 SO0001-: Oi
0. 1 i2i771-:
02
0.oOil02E IS
0. 1 S'.iooof: ',
. 1 ij 1 1 -1 1-:
02
. 72 1 4 7 1 , f:
1 T
0.7tOI-: r)7
0. 1 T^OOOK i
0. 7lill i24|.:
01
0. 3SS3i9E IS
0. 1 ssooo].: Oi
0. 919n4iF:
01
0. 42','7'i7F;
1 S
. fs F. 7
0. 1 ^i^ooor-: o i
0. 4 in i02I.:
01
n.203907E IS
0. 1 ssooof: 03
0. S2772 SE
(11
0. 24i.t> i4F:
1 s
0. H-iOi,: 07
0. 1 stOooi.: i
0. 2S03i)Ol.:
01
0. 1 1 7009E 1 5
0. 1 ssooof: 3
(1. 3 028 2 8].:
0!
0. 141 S27E
1 s
,_
0. "OOP: 07
0. 1 stoooi-: i
0. \AU,<, >K
(11
0.671442E 14
0. 1 ssooof: i
0. 1 7 i7 74E
01
0.812 1 ioF:
14 .■
Fig. 8. Lower Atmosphere Model III for ground air temperatures 180, 190
200, and 210°K. Surface pressure 15 mb (0. 15 X 105 dyne cm'^); atm'os-
pheric abundance 60% CO^ , 20% Ar, 20% N2 by volume; mean molecular
mass of atmospheric constituents 0.663921 X 10"^^ gm.
September 1 1, 1967
E. Monash, JPL
Sec. 5.3, page 17
Lower Atmosphere
JPL, 606-1
H
Hrii^ht alxivi-
.Atit'o
s pin- ric pa ram e
e rs
.Mmoh
piu'ric ]jara
ri;et
.'!■ s
T
1'
N
T
1'
V
nu-an surface,
Kinetic
Tot.-il
Total
Kinetic
4 ot.,1
Tola!
Cltl
tcrTpc ratii re , j
p re s su re
concent rati
)n,
tempe r.itu rv .
p re s ,sn re
c once nt ra! i
. n,
"K
Hync cm"
?
cm- 5
"K
dyne c ti- -
-
cni- ■'
0. OOOE- W
C i r o 11 n
1 ail
tctr,]>c raUi
-c T
,: 220 'K
C'l r, )unH ."i i i
temp*- ratu r«- T
,: 2tO"K
0. 220000K
■;
0. 1 SOOOOE
f)S
0. 4'> 5't07E
18
0. 24 00 1-: 5
. 1 soofjO!-:
OS
0. 4S27-I81-:
1 8
0. lOOE 0()
0. 2I4')K4K
',
0. 1 580 52!':
OS
0. 46 SI 03E
18
0. 2 5S080I-: 5
0. 1 5 0004 E
OS
0. -128 540I-:
1 8
0. ZOOF: 06
0. 209ftiiO[';
5
0. 121)7701-:
OS
(1. 4 57 5 S ME
18
0. 2 501 nOl-: 5
. 1 2 H t , 7 I-:
5
fl. ■;o-t774i-;
1 '3
0. lOOE 06
0.207296f:
5
0. 1 1 t, 507E
OS
0. 4 0(,-r5SE
18
0.22S2401-: 5
0.11 878K1-:
OS
. 5 K 2 5 7 I-:
1 8
0.40nE 06
0. 206S7fiE
i
0. int.7241':
OS
0. 5 7-1 24 SE
18
0. 22208 5I-: 05
0. 1 OOn07E
OS
0. 557 520I-:
1 8
0. 500E Ot)
0. 2l)S,SSr,E
0>,
0. ■i7't00 5E
04
0. 544S0i)E
18
0.2 2 104 3 E 0''
0. 1 01 14'lE
OS
0. 5 51481 E
1 8
0.60 0E 6
0.20 SI 1r,E
i
0. K'>77'14K
04
0.51 7037 E
18
0. 22000 5E 5
0. '■! •• 3 071 E
04
0. 3 0722 9 E
1 8
0.700E 06
0.2044L6E
Oi
0. K2 5070E
04
0. 291674E
18
0. 21 H<'t) 5I-: 03
11. 8t,040t)E
04
0. 284t,48E
1 8
0. KOOE Ot)
0. 20?696E
(H
0. 7S4534E
04
0. 21)8261 E
18
0. 21 7' '2 5E 5
0.7') 50't41-:
04
0. 2()5n52E
1 H
0. 900E 06
0.202'l7i>!.:
3
0.6'M 12SE
04
0. 24tii,S4K
18
0. 2 1 tiKH 5E 5
0. 7 ill7i>4K
0-t
0. 244077I-:
1 H
0. lOOE 07
0. 2022 S(,E
3
0.05 5 1 1, E
04
0. 226720E
1 8
0.2 1 S84 5 F. 5
0. n7 30t,7E
0-f
. 2 2 5 H ■ t I-:
1 8
0. 150E 07
0, 1 986StjE
0!
0.4061 lOI':
04
0. 1 48087E
18
0. 2i On4 5E 3
0.443448K
04
0. 1 52 501 E
1 8
0. 200E 07
0. 19SnStjE
5
(1. 2SK4 5 }!■:
04
0. ')S-i7t) 3E
17
0. 20S44 3E 5
0. 2891 5 5E
04
0. 1 01 ''4'!i-:
1 8
0. 2S0E 07
0. 19l4Si,E
5
f). 1 i)5078f;
04
0. 6 1702 5 E
17
0. 20 024 5E 5
0. 1 8i,4t,2E
4
0. -.74 541 E
1 7
0. 500E 07
0. 1 H7KSf>E
05
0. 1 02 01 1 F.
04
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1 7
0. 1 'IS04 3E 5
0.11 8 8t,'ll-:
0-4
0.441 484 E
1 7
0. 150E 07
0. 1 K42St.h:
5
0. t) 52 54(>E
3
0. 24860SE
17
0. 1 8484 5E 5
0.748t,271-:
5
0. 2 8 St, -1 HE
1 7
0.400E 07
0. 1 KOt>st)f:
5
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U 5
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17
0. 1 8 4ti4 5E 5
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5
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! 7
0.4S0E 07
0. 1 770S6E
5
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3
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If,
0. 1 7944 51-: 05
. 2 H 5 S 1-;
5
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1 7
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0. 1 734S(,E
5
0. 1421 09E
3
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1 !,
. 1 7 4 2 4 5 1-: 5
0. 172',18E
(J 5
0.71 7, .-to 1-:
j ii
0. 5 50E 07
0. 171 OOOE
5
0. HS')|78K
02
0. 56 5 91)8 E
Itj
0. 1 7i 000 !•: 3
f). 1 04 5t, 5E
5
0. ■i42 1 071-;
1 1,
0.600E 07
0. lt)7000E
05
0. SI 5 5001-:
2
0. 222l)S4E
n.
0. 1 670001-: 5
f).t,2 54'lHI-:
02
(J. 2704 5 SE
1 1,
0.b50E 07
0. 1K3000E
!
0. 50285t)E
02
0. 1 54S94I-;
It,
0. 1 t) 500(11-; 5
. 5 1) 7 8 7 4 H:
(J 2
0. It) -otH'ii-:
1 n
0. 700E 07
0. 1 S9000E
3
(1. 1 7t. 5i)3E
02
0. HO 5500I-;
IS
0. 1 S'tOOOE 5
0. 2 1422 SE
02
(J . '(7 5 ■)■<■)[•;
1 T
0. 75GF'": 07
0. 1 S5000E
03
0. 1012')6E
02
0. 47 541 1 (•:
IS
0. 1 ssoooh: 5
0.12 504 51-;
02
0. T7-^04t)E
1 T
0. BOOK 07
0. 1 S5000E
03
0. S81 2 76E
01
0. 271t,ij| E
1 s
0. 1 SSOOOE 5
0. 7(Jlj0tj8I-:
01
0. 32 '1 t«2i-:
1 -)
0. K50E 07
0. 1 5S000E
03
0. 5 5 sssbf:
01
0. 1 SS889F;
1 s
0. 1 SSOOOE 5
0. 40S1 t,8E
Oi
0. 1 K't 5St,E
1 5
0. 900E 07
0. 1 SSOOOE
03
0. 1 91 408 E
01
0. 804S48I-;
14
0. 1 SSOOOE 5
0. 2 52S(JOE
01
0. 1 (jHt.S'lE
1 ~i
0. OOOE- iH
G roun
d ai
r tcnipe ratu rv T
„: 250"K
Cj ri)un<i ai
r te n ipe r.itvi I'e 1"
,,: 2-,0 K
0. 2 30000E
1
3
0. 1 SOOOOE
OS
0.472432E
18
0. 2 5 0000]-: (i5
(1. 1 SOOOOE
OS
0. 4 54', 581-;
1 H
0. lOOE 06
0, 22S0SSE
5
0. 1 5KS 5 9E
OS
0.44S96bI-:
1 H
(J. 24S12t)E 05
0. 1 5')4 5 5 1-;
OS
(j. 4 1 20 5 5 1-;
1 H
0.200E 06
0.2200i)SK
5
0. 1 2 7726 E
OS
0. 420442I-:
18
0. 2402S2I-: 5
0. 1 2''421 1-:
OS
0. 5'.0224l-;
1 H
0. 100E 06
0.21 S008E
3
0. 1 17 5401':
OS
0. 59S845E
18
0. 2 55 5701-: 5
0. 1 l'C'45l-:
OS
0. 5..-'l 3'tE
1 H
0.40 0E 06
0. 21421 KE
5
0. 1 081 h9E
OS
0. 5I)S8S0H;
18
0. 2 50-5OSE 05
0.11 o'iHSi-:
OS
0. 5487871-;
1 8
0. SOOE 06
0.21 i33«E
5
0. <t'>S477l-;
04
0. 5 5801 8E
1 H
0. 22H'i71 E 03
(j. 1 02n t(,E
5
0. 524 '001-:
1 8
0.600E 06
0, 2124SKE
3
0. '>\ St,S4E
04
0. 3 12201 E
1 8
0. 227771 E 05
0. ')-t'i't77E
04
0. 502 i 281-;
1 8
0. 700 E Ot.
0.211 S7SE
3
0. 841'I5'>1.:
04
0. 28H2i)2I-:
IK
0. 22t, ,71 !-; 5
0. 87840 51-:
4
0. 2808-t-tI-;
1 8
0. KOOE 06
0. 210t)9KE
5
0. 77 i887E
04
0. 20!)0(.';E
18
(9.225 571 I-: 5
0.811 .S8 -'.I-;
04
. 2 1 - ( t t n ■ ( E
1 H
0. '100 E 06
0. 209K1 HE
5
0. 71 1 08 S[.;
04
0. 24SS02I-:
1 8
0.224171 I-: 5
- 0. 7S00'8r,i-;
04
0. 242 581,1-:
1 H
0. lOOE 07
0. ZO«93HE
5
0. t>S 51 48 K
04
0. 221-449E
18
0. 222',I71 K 5
0. ),'t2t; 'tH I-;
0-t
0. 22-^')-leI-:
1 8
0. 1 SOE 07
0. 204SiSE
5
0.42471 iK
04
0. 1 S(I41 8 1-:
1 H
0.2 1 n')7 1 E >
(J. 4i,22-i 1 I-;
04
0. 1 5J 527I-:
i 8
O.ZOOE 07
0. 2001 38E
5
0.2 7 5f.00K
04
0. ')'t020 1 E
17
0. 210'!71 I-: (.15
(1. 504''77I-;
04
(1. 1 (i47 1 Hi-:
1 -^
0.2S0E 07
0. 19S7 3HE
1)3
0. 1 74 S 5 HE
04
. t, 4 S 9 5 8 I-:
1 7
0. 204 ,t71 I-: 5
0. 1 'i8H 1 7E
0-t
0. 7(J2n48i-:
1 "
0. ^OOE 07
0. 191 5 SHE
3
0. 1 1112 1 1 E
04
0.4I72S3E
17
0. 1 't8 171 1-: 5
0. 1 27't7-il-
(J4
0. 4t,=."i SI-:
1 "
0. ISOE 07
0. 18t)93KI-:
5
0.t,8KSll,l.-
5
0. 2l.'-80SE
17
0. I'll ill F 5
0. HI2).',)SE
5
0. 505(l7't|-:
1 ~
0.400E 07
0. 1H2S SBE
3
0.42S541 K
5
0. 1 i)87')tjl-:
17
0. 1 8(,'t71 F 5
IJ. S087S1 F
(.} 5
0. 1 '<71 1 OF.
1 7
0.450E 07
0. 1781 i8h:
(J 5
0. 2 S')6'l| |.
5
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17
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3
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1 7
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0. 17i75KE
5
0. 1 Si,6 09l-"
03
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1 1>
0. 1 74971 E 5
0. 1 '(024 5E
5
(J . 7 8 7 1., 2 5 1-:
1 '•
0. SSOE 07
0. 1 71000 E
03
0.94OH41 !■
02
0.401 I04E
It,
0. 171 OOOE 5
(1. 1 1 SO r 11-
-,
( (.). 4K7 2-171-:
1 ',
0.600E 07
0. 167000E
5
0. S6S672I-
2
0. 24S5721-:
M,
0. 10 70001-; 5
(1. r,871 S'(l-
02
0. 2')8()i,'ti.;
1 I;
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0. 163000J':
3
0. 5 5 57Sl,I-
02
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In
0. in 50001-: 5
0.4054 5 5 1-
2
(J. 1 HOI HI I-:
1 ti
0.700E 07
0. 1 soooot:
5
0. 1 •>4 5S7I-
2
0. H8S4H1 E
1 S
0. 1 5O0O0I-: 5
0.2 5t,0'iHI-
02
0. 1 9751, Si-
i t,
0. 7 SOE 07
0. 1 SSOOOE
5
0. 1 1 It, 521-
02
0. S2I71 5P:
IS
0. 1 SOOOOE 5
0. 1 5Tr.0t)|-
02
0... 5 '.7S'!|-
i 5
O.KOOE 07
0. 1 SSOOOE
n 5
0. t,40S841-
1
0. 2',!'! 5 78I-,
1 S
0. 1 '5S000]-; 5
0. 7 78 1 S i|-
1
0. 5n. 51,74 i-
1 -1
0. KSOt: 07
0. 1 SSOOOE
3
0. 567S'I| 1-
1
0. 171 7 941-
15
0. 1 S50(tOI-: 3
0. 44r,5 5n(-
01
0. 20Hn') (Ji-
0. 900E 07
0. 1 SSOOOE
5
0.2I0'I 571-
1
. ',» 8 S 8 1 9 h
1 4
0. 1 ssooof: 5
0. 2S(,2 5'(|-
01
ll. 1 1 '.7^-11-
Fig. 9. Lower Atmosphere Model III for ground air temperatures 220, 230,
240, and 250°K. Surface pressure 15 mb (0. 15 X 105 dyne cm'^); atmos-
pheric abundance 60% CO2, 20% Ar , 20% N2 by volume; mean molecular
mass of atmospheric constituents 0.663921 X 10"22 gm.
Sec. 5.3, page li
E. Monash, JPL
September 11, 1967
JPL 606-1
Lower Atmosphere
H
Height a
lOve
Atmospheric parame
ters
Atmospheric parameters
T
P
N
T
P
N
mean sur
face,
Kinetic
Total
Total
Kinetic
Total
Total
cm
tempe ratu
re,
pressure
concentrat
on.
tempe ratu
re,
pressure
concentration,
•K
dyne cm"
2
cm-3
°K
dyne cm"
2
cm" 3
0. OOOE
38
Groun
d air temperature T
o: 270°K
Ground air temperature T
o: 290-K
0,270000E
03
0. 150000E
05
0.402442E
18
0.290000E
03
0. 150000E
05
0.374688E 18
0. lOOE
06
0.265216E
03
0. 140198E
05
0.382930E
18
0.285303E
03
0.140860E
05
0.357650E 18
0.200E
06
0.260433E
03
0. 130876E
05
0. 364034E
18
0.280607E
03
0.132139E
05
0. 34U23E 18
0. 300E
06
0.255649E
03
0. 122018E
05
0.345746E
18
0.275910E
03
0. 123825E
05
0.325100E 18
0.400E
06
0.250866E
03
0. 113609E
05
0.328057E
18
0.271213E
03
0.U5904E
05
0.309573E 18
0. 500E
06
0.246082E
03
0.105635E
05
0.310958E
18
0.266517E
03
0. 108365E
05
0.294537E 18
0.600E
06
0.243969E
03
0.982068E
04
0.291597E
18
0.261820E
03
0. 101195E
05
0.279983E 18
0. 700E
06
0.242449E
03
0.912779E
04
0.272723E
18
0.259201E
03
0.944776E
04
0.264039E 18
0. 800E
06
0. 240929E
03
0.847988E
04
0.254963E
18
0.257361E
03
0.881875E
04
0.248222E 18
0. 900E
06
0. 239409E
03
0.787429E
04
0.238258E
18
0.255521E
03
0.822755E
04
0.233249E 18
0. lOOE
07
0. 237889E
03
0.730850E
04
0.222551E
18
0.253681E
03
0.767214E
04
0.219081E 18
0. 150E
07
0. 230289E
03
0.499739E
04
0. 157198E
18
0.244481E
03
0.536741E
04
0.159036E 18
0,200E
07
0.222689E
03
0.337379E
04
0. 109748E
18
0.235281E
03
0.370391E
04
0. 114038E 18
0.250E
07
0.21 5089E
03
0.224682E
04
0.756705E
17
0.226081E
03
0.251843E
04
0.806941E 17
0.300E
07
0.207489E
03
0.147457E
04
0. 514811E
17
0.216881E
03
0.168515E
04
0. 562850E 17
0.350E
07
0. i99889E
03
0.952660E
03
0.345243E
17
0.207681E
03
0. 1 1081 IE
04
0. 386511E 17
0.400E
07
0. 192289E
03
0.605138E
03
0.227969E
17
0. 198481E
03
0. 714951E
03
0.260935E 17
0.450E
07
0. I84689E
03
0.377418E
03
0. 148033E
17
0. 189281E
03
0.451789E
03
0. 172904E 17
0. 500E
07
0. 177089E
03
0.230767E
03
0.943970E
16
0. 180081E
03
0.279037E
03
0. 112246E 17
0. 550E
07
0, I71000E
03
0. 139519E
03
0. 591037E
16
0. 171000E
03
0. 168703E
03
0. 714666E 16
0.600E
07
0. 167000E
03
0.833532E
02
0.361 562E
16
0. 167000E
03
0. 100788E
03
0.437190E 16
0.650E
07
0. 163000E
03
0.491798E
02
0.218562E
16
0. 163000E
03
0. 594669E
02
0.264280E 16
0. 700E
07
0. isgoooE
03
0.286390E
02
0.130478E
16
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Fig. 10. Lower Atmosphere Model III for ground air temperatures 270 and
290°K. Surface pressure 15 mb (0. 15 X 105 dyne cm~^); atmospheric
abundance 60% CO2, 20% Ar, 20% N2 by volume; mean molecular mass of
atmospheric constituents 0.663921 X 10"^^ gn^-
September 11, 1967
E. Monash, JPL
Sec. 5.3, page 19
Lower Atmosphere
JPL 606-1
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E. Monash, JPL
September 11, 1967
JPL 606-1 L"^^''^" Atmosphere
BIBLIOGRAPHY
Anderson, A. D. , 1965, A model for the lower atmosphere of Mars based on
Mariner IV occiiltation data: Palo Alto, Calif., Lockheed Palo Alto
Research Lab., Tech . Memo . 6 -7 5 -65 -6Z .
Fleagle,R.G. , and Businger , J . A . , 1963, An introduction to atmospheric
physics (esp. p. 70): New York, Academic Press.
Gierasch,P., and Goody , R . M . , 1967, An approximate calculation of radiative
heating and radiative equilibrium in the Martian atmosphere: Cambridge,
Mass., Harvard U. , Preprint.
Goody, R.M., 1964, Atmospheric radiation, I. Theoretical basis (esp . p. 329):
London, Oxford U . Pres s (Clarendon Pre s s) .
Goody, R.M., and Belton, M . J . , 1967, Radiative relaxation times for Mar s: a
discussion of Martian atmospheric dynamics: Planet. Space Sci . , v. 15,
P.Z47-256.
Hilsenrath, J. , Hoge.H.J., Beckett, C . W . , Masi.J.F., Benedict, W . C . ,
Nuttall,R. L. , Fano.L., Touloukian, Y . S. , and WooUey , H. W . , I960,
Tables of thermodynamic and transport properties of air, carbon dioxide,
carbon monoxide, hydrogen, nitrogen, oxygen, and steam: New York,
Pergamon Press.
Leovy.C, 1966, Radiative -convective equilibrium calculations for a two-layer
Mars atmosphere: Santa Monica, Calif . , RAND Corp., Memo . RM-50 1 7 -
NASA.
Neubauer , P.M. , 1966, Thermal convection in the Martian atmosphere:
J.Geophys. Res. , v. 71, p.Z419-2426.
Ohring,G., House, F., Sherman, C, and Tang.W., 1967, Study of the Martian
atmospheric environmental requirements for spacecraft and entry
vehicles: Bedford, Mas s . , GCA Corp. , TR-67-12-G.
Prabhakara, C. , and Hogan, J . S. , Jr . , 1965, Ozone and carbon dioxide heating
in the Martian atmosphere: J . Atmos. Sci . , v. 22, p. 97-109.
March 1 1972 E. Monash, JPL Sec. 5.3, page 21
JPL 606-1 Upper Atmosphere
SECTION 5. 4
UPPER ATMOSPHERE
This section has not been revised, although it contains outdated informa-
tion written almost 5 years ago. Flowever, the improved information on the
Martian atmosphere obtained over the past 5 years is so substantial that a major
research effort is required to evaluate and document this data properly.
The level of sophistication now possible in producing models from this
later data should provide extremely meaningful information in the near futvire.
The material contained in this section can be utilized as background
information.
Until the later data has been documented, reference to the Viking 75
Project, Mars Engineering Model (M75- 125-1) will provide the reader with
improved information.
March 1, 1972 Sec. 5.4, page
JfL 606-1 Upper Atmosphere
5.4 UPPER ATMOSPHERE
DATA SUMMARY
Interpretation of the data from the Mariner IV occultation experiment has
led to development of three distinct types of models for the upper atmosphere of
Mars: a preliminary E-Model, an F^-Model at maximum solar flux (Figs. 1
through 5), and an F^ -Model at minimum solar flux (Figs. 6 through 8). i* In
addition to occultation data, each of the models uses the best results of Earth-
based observations, which are presented elsewhere in this document. ^
The E-, Fj-, and F2 -Models are distinguished by the type of Earth-analog
ionospheric layer which the ionized layer detected in the Martian atmosphere is
thought to represent; they thus differ primarily in the structure of the iono-
sphere as derived from the physical and chemical processes thought to occur
there. The "correct" choice among the models is not clear. The basic point of
disagreement is proper identification of the concentration peak of the electron
density profile measured at an altitude of approximately IZO km above Electris
at the time of immersion of Mariner IV' s S-band transmission on July 15, 1965.
DISCUSSION
Layers of the Upper Atmosphere
Using terrestrial nomenclature for classifying various regions of the
atmosphere, the thermosphere is the highest region of the atmosphere where
coUisional interactions are still important, a region where kinetic temperature
increases from its lowest value to a constant (which is quite high in the terres-
trial case). The thermosphere differs from lower atmospheric levels because
ionization and dissociation occur and diffusion becomes more important than
mixing. The ionization region, or ionosphere, and dissociation region lie
largely in the thermosphere. Both are produced primarily by interactions
(photoionization and photodissociation) of extreme ultraviolet solar photons
with atmospheric gases.
At some level in every atmosphere, called the critical level, the
horizontal mean free path of neutral particles becomes equal to the density
scale height. This critical level marks the base of the exosphere. In the exo-
sphere, particles move along ballistic trajectories to a first approximation. If
a magnetic field is present, it will fill a volume known as the magnetosphere ,
which generally extends beyond the exosphere and which must be considered in
terms of its effects on charged-particle trajectories. The exosphere or mag-
netosphere is the region whose conditions determine the probability of escape
of an atmospheric constituent . 3
See page 20 for list of cross references.
'uly3, 1967 E. Monash, JPL Sec. 5.4, page 1
Upper Atmosphere JPL 606-1
A planet without a permanent magnetic field may have a weak field
convected into it by the solar wind. This, it is thought, would add to the for-
mation of a collisionless bow shock in the solar wind ahead of and extending
around the planet. Solar wind flow in the zone behind the shock offers a further
factor affecting the exosphere. Attempts to account for this feature in relation
to the Martian atmosphere are still in early stages and are not included in the
atmospheric models presented in this section. (For additional reading about
upper atmospheres, see Craig, 1965, and Fleagle and Businger, 1963).
Physics of the Upper Atmosphere
The important photochemical and physical processes which are believed
to occur in the upper atmosphere of Mars are based upon the theories of photo-
chemical and diffusive equilibrium in a quasi -stationary (approximately steady-
state) atmosphere. These processes form the basis of the models presented
in this section. The only available direct measurements of the properties of
and processes occurring in the upper Martian atmosphere were provided by the
1965 Mariner IV occultation experiment.
Photodissociation Region
The photodissociative solar flux at a given height depends upon how much
attenuation of the flux has taken place above the given level. The primary con-
stituent that is photodissociated on Mars is CO2 . The significant reactions
occurring are
CO2 + hu ~* CO + O (1)
CO + O -♦ CO2 + hi/ (2)
O + O + M -♦ O2 + M (3)
O + O -» O2 + hi/ (4)
O2 + hi/ -♦ O + O (5)
O + CO + M -♦ CO2 + M (6)
where M is any third body in the three-body process, and hi/ is the energy of the
incident photon.
c
Absorption of the solar spectrum between 1800 and 2400 A by CO2 and O^
is negligible (Norton, 1964). Rayleigh (molecular) scattering of extreme ultra-
violet solar photons with wavelengths lying between 1800 and 2400 A is quite
small. The optical depth for Rayleigh scattering at X2100 at the surface of
Mars for a pure CO2 atmosphere is less than 0. 5 and decreases as the height
increases because there are fewer and fewer molecules to scatter the radiation
(Coulson and Lotman, 1962). This radiation is therefore able to penetrate to a
lower level, where it is available to be absorbed by ozone, if any ozone exists.
There will be some modification in detail if nitrogen exists in significant
quantities .
Sec. 5.4, page 2 E. Monash, JPL July 3, 1967
JPL 606-1
Upper Atmosphere
When atomic oxygen is released from photodissociation of CO2 , it is
available for the oxidation of CO. The atomic oxygen may react with molecular
oxygen through a third body to produce ozone.'* Ozone could then disappear by
absorption of the extreme ultraviolet solar radiation. In fact, there is good
evidence that there is little ozone present and that a great deal of ultraviolet
radiation penetrates to the Martian surface. ^
To establish the major chemical processes important in the photodissoci-
ation region, the reaction rates of Eqs . (1) through (6) must be known. Two
reactions produce CO2, one by two-body association (Eq. 2) and the other by
three-body association (Eq. 6). At the densities and temperatures of the Mar-
tian atmosphere, two-body association dominates (Smith and Beutler, 1967).
Both two-body (Eq. 4) and three-body (Eq. 3) association of O2 must be con-
sidered, as the rates are comparable (Smith and Beutler, 1967).
At 90 and 100 km the concentrations of CO, O^, and O in the F^ -Model
are given (Fig. 1) by photochemical equilibrium calculations based upon the
reactions given by Eqs. (1) through (6). These were modified between 100 and
110 km by estimating the effect of the transport terms in the continuity equation
(diffusion) to improve the values of the concentrations. From 110 km up to the
critical level (the base of the exosphere) at about 500 km, the various atmos-
pheric species are thought to be in diffusive equilibrium as given by
T
[x] = [x]
110km
110 km
T
exp
m /
' ~k J
^ dh
T
110 km
(7)
where [xJ^iQ km ^^*^ ^110 km ^-^^ ^^'^ concentration of any photodissociation
product and the temperature, respectively, at a height of 110 km, m is the
mean molecular mass, k is the Boltzmann constant, and g is the acceleration
due to gravity.
Equation (7) is based upon the assumptions of hydrostatic equilibrium and
the ideal equation of state and hence represents a diffusive distribution. Its
applicability to the upper atmosphere of Mars to describe the concentration of
various constituents as a function of altitude is valid even above the critical
level, provided the deviation from hydrostatic equilibrium is small. Chamber-
lain (1963), in his monograph on the physics of planetary exospheres, tabulates
the departure from hydrostatic equilibrium in terms of generally applicable
atmospheric parameters. In the F^ -Model of Smith and Beutler the critical
level occurs at approximately 500 km based upon the mean free path between
oxygen-oxygen collisions. Yet Chamberlain's work shows Eq. (7) to be applica-
ble to an altitude of nearly 30,000 km if magnetospheric phenomena are ignored,
Ionosphere
One method of formation of an ionospheric layer, the classic Chapman
layer, can be understood as follows. At large heights the density of any gas is
very low and the corresponding rate of formation of electron-ion pairs is low.
July 3, 1967
E. Monash, JPL
Sec. 5.4,
page
Upper Atmosphere JPL 606-1
At small heights the ultraviolet solar flux capable of ionizing a particular gas
may be reduced by its passage through the atmosphere; hence, the rate at this
height can also be low. At intermediate heights the formation rate of electron-
ion pairs will reach a maximum since the product of density and ionizing flux
reaches a maximum.
There are other ways to form a layer. In general, however, ionization of
various species occurs throughout a thermosphere, with rather broad maxima
being as likely as distinct layers. The Mariner IV occultation data indicated
the presence of an ionized layer in the Martian atmosphere at an altitude of
approximately 120 km (ingress).
Ionization Processes . The production rate of electrons and ions is in
competition with the loss rate. The primary reactions for the production of
electrons are
X + hi/ -> X"*" + e (8)
and I
XY + hu -» XY + e (9)
where X is the atom and X"*" the corresponding ion, XY is the molecule and XY
the corresponding ion. Production of ions is more complex than that of elec-
trons as a result of the various chemical processes that can occur. Figure 9
gives the reactions that may occur significantly in the ionosphere of Mars for a
pure COo lower atmosphere and shows that the primary constituents of the Mar-
tian ionosphere are CO"^, CO2. O"^, O2. and electrons. Smith and Beutler
(1967) derive the following equations for the concentrations of these constituents
based upon chemical equilibrium:
+ qgCco]
^^° ^ = ki[0] + (k3 +k5)[02] +k8[C02] + "sLe] ^^°^
^ q4[C02] +k8[CO + ]
^^°2] = (k2 +k7)[0] +k6[02] +Oc^le] ^^^^
^ qiCO] + (ki[CO + ] + k2[C0^])[0] + k3[CO+][02]
^° ^ " k4[02] +kg[C02] +ai[e] ^ '
+ qzC Oz] + (^4[0^] + ^5[CO^] + k8[CO^])[02] + k^[CO^][0]
[O2] - O^f^] ^ '
Sec. 5.4, page 4 E. Monash, JPL July 3, 1967
JPL 606-1 Upper Atmosphere
Charge neutrality requires that
[e] = [CO + ] + [COJ] + [0 + ] + [Oj] (14)
The q's, a's, and k' s in Eqs. (10) through (13) are rate coefficients for the var-
ious chemical processes appearing on the right side of Fig. 9-
The rate for photoionization q is not given explicitly in Fig. 9 but can be
determined from Fig. 10. Values of the photoionization rate at a given altitude
can be estimated by multiplying the extreme ultraviolet solar flux Jq(X) at the
top of the Martian atmosphere by the cross section for photoabsorption O'(X) and
by e-''"(X), where T(X) is the optical depth (at that altitude) at wavelength X, to
allow for absorption in higher atmospheric layers. The product jQ(X)a(X)e-'^(X)
is then integrated over X for each constituent of the atmosphere.
To discuss the rate coefficients Oc and k it is necessary to simultaneously
discuss the various chemical processes that can remove electrons and ions.
The simplest process for loss of electrons and ions is the inverse of Eqs. (8)
and (9), namely, radiative recombination.
X"*" + e -> X + hl^ (15)
XY"^ + e -♦ XY + hl^ (16)
A typical radiative recombination process which appears in Fig. 9 is
0+e-»0+hl^ (17)
with a rate coefficient Oi^ - 2 XIO"^'^ cm sec" .
Nonradiative processes can remove electrons and ions much faster. If a
molecular ion XY+ recombines with an electron, the energy of recombination
can dissociate the molecule into two atoms X and Y and will proceed with a far
greater rate constant than processes involving photon emission.
XY"^ + e -♦ X + Y (18)
Dissociative recombination is limited to the lower parts of the upper atmosphere
because under diffusive equilibrium these molecules will tend to be concentrated
lower in the atmosphere than lighter constituents such as atomic oxygen.
July 3, 1967 E. Monash, JPL Sec. 5.4, page 5
Upper Atmosphere JPL 606-1
The ionization balance of atomic species must take account of charge
transfer, that is,
X"^ + YZ -+ YZ"^ + X (19)
and ion-atom (or molecule) exchange
X"*" + YZ -♦ XY"^ + Z (20)
which with Eq. (18) offer a vehicle for efficient removal of X ions and elec-
trons. A typical charge transfer reaction appearing in Fig. 9 is
O"*" + O2 -♦ O2 + O (21)
wi
ith a rate coefficient k^ = 3. 5 X 10"^^ cm^ sec
In addition to the processes thus far mientioned, electron attachment can
become important in the lower parts of the ionosphere.
X + e+M-»X"+M (22)
where M is any third body. The negatively charged ion can combine with a
positively charged ion to produce two neutral particles.
X' + Y"^ -» X + Y (23)
The total recombination rate coefficient, which includes radiative and
dissociative recombination, must take account of dielectronic recombination.
Dielectronic recombination represents the capture of an electron by an ion in
a process which results in excitation of the resulting neutral into an upper
quantum state. In the process the captured electron gives up some of its extra
energy to one of the bound electrons. The result of dielectronic recombination
is thus an atom or molecule with a doubly-excited state which then normally
decays by ordinary photon emission.
Thermcd Processes. The vertical tennperature distribution and vertical
temperature gradient of the Martian ionosphere is calculated from a detailed
energy balance. The Martian ionosphere obtains its energy from the Sun and is
expected to be in quasi- stationary equilibrium with the local UV photon flux at
any given time of day. '
A certain fraction of the solar radiant energy is reflected or absorbed
directly by atmospheric gases or by clouds, ^ and the remainder is absorbed or
Sec. 5.4, page 6 E. Monash, JPL July 3, 1967
JPL 606-1 Upper Atmosphere
reflected by the Martian surface.* Of the energy absorbed by the Martian
surface a large fraction is then reradiated in the infrared, where it can be
absorbed by gases in the lower atmosphere. This energy trapping in the lower
atmosphere constitutes a greenhouse effect which results in an increase of the
surface temperature^ above the temperature for a corresponding atmosphere
that is optically thin (transparent) at all wavelengths.
The extreme ultraviolet solar photons that are absorbed by the Martian
atmospheric gases supply the energy required to maintain the thermosphere.
The rate of thermal energy input into the thermosphere for each constituent is
Ej(h) =ye.(X)CT.(X) N. J^(X) e"'^^'^) dA (24)
where Nj represents the number density of constituent j and € i(X) represents the
efficiency for conversion of ultraviolet solar radiation to heat via photochemical
processes by the j-th constituent; the other quantities have been defined previ-
ously. The total rate of heat input is the sum of E- for all constituents; hence,
E(h) = EE.(h) (25)
J
The thermosphere of Mars is cooled by the rotational, vibrational, and
electronic radiation of CO, CO2, and O, respectively. The total rate of energy
loss at each level of the thermosphere from the infrared radiation is
R(h) = E Ri(h) (26)
j
where Rj represents the infrared radiative term for each constituent j. The
only other mechanism considered significant in heat loss is thermal conduction.
The thermal conductive flux is AT^/^ dT/dh where (AT^/^) is a weighted mean
coefficient of thermal conductivity for CO2, O2, CO, and O.
The energy balance equation can now be written
where p is total mass density of the gas, Cy is specific heat at constant volume,
and 9T/9t is time rate of change of the kinetic temperature. Equation (27) is
solved for quasi-stationary equilibrium (9T/9t ci 0); hence,
AT^/^^ -/ (E-R)dh (28)
h
July 3, 1967 E. Monash, JPL Sec. 5.4, page 7
Upper Atmosphere JPL 606-1
Equation (28) determines the temperature gradient for the upper atmosphere of
Mars. The temperature profile is then found by numerical integration using a
previously derived lower atmosphere as the lower set of boundary conditions.
Since E(h) , through e-TiX), depends on the integrated atmosphere above h, an
iterative procedure is necessary.
Contemporary Models of the Upper Atmosphere
Contemporary models of the upper Martian atmosphere are based on the
general theory beginning on page 2 and ending above. They all use interpreta-
tions of data from the Mariner IV occultation experiment and the best results of
Earth-based observations. The models are classified by approximate analogy
to terrestrial ionospheric nomenclature as E-, F^-, or F2-(region) Models.
Ionospheric regions or layers are differentiated by the wavelength of the radia-
tion producing the region, by the principal ionic constituent, and by the domi-
nant loss mechanism of ions and electrons.
The E-region of the Earth's atmosphere is produced primarily by x-rays
from about 10 to about 100 A assisted by some ultraviolet radiation from 800 to
1026 A — particularly by emission lines such as Lyman /3 at 1025 A (Craig, 1965).
The most important ion produced is 02> and the loss mechanism is primarily
dissociative recombination. This region is a fair approximation to the classic
Chapman layer described briefly on page 3, a layer in which the peak electron
density is approximately coincident with the peak of the electron production.
The terrestrial F- region is produced by extreme ultraviolet radiation in
the wavelength range from 200 to 800 A with some contribution from longer
wavelengths. The principal ion present is 0+ (Craig, 1965). The loss mecha-
nism is ion-atom interchange up to a high level, where radiative recombination
becomes dominant. This is a height-dependent process which results in the
peak electron density being at a much higher altitude than the peak of production,
a so-called Bradbury layer. Actually, the F- region contains two peaks: a lower
layer called the Fi -layer near the production maximum, and the Bradbury layer
peak called the F2-peak (Craig, 1965).
Preliminary E-Model
In their preliminary paper, Chamberlain and McElroy (1966) interpret the
ionization layer above Electris to be an E-region. They assume a composition
of only 44% CO2, the remainder being N2 . Their model relies heavily on the
speculative assumption that mixing provides homogeneous composition of con-
stituent gases throughout the entire upper atmosphere. They assume that the
rate of photodissociation is so small that CO2 will be only partially photodis-
sociated, even at great heights, primarily because the rapid formation of O2
(from O resulting from the CO2 dissociation) shields the CO2. They further
assume a very large rate coefficient for dissociative recombination of CO2 to
prevent formation of an F;^ -region above the E-region (since of course none
was observed). A new determination of this coefficient by Weller and Biondi
(1967) makes it difficult to accept such an assumption.
These assumptions can res\alt in an ionosphere dominated by O2 (or
perhaps NO+) ^ in a Chapman layer created primarily by x-rays.' With appropri-
ate rate coefficients the layer is of the strength and at the height measured by
Sec. 5.4, page 8 E. Monash, JPL July 3, 1967
JPL 606-1 Upper Atmosphere
Mariner IV. Chamberlain and McElroy's assumptions seem questionable
according to proponents of other models; in addition, they do not present the
observed ionospheric scale height. However, Chamberlain and McElroy are
working on much more detailed calculations, and a final E-Model will be added
to this section when the results become available from these authors.
Fi -Model (Figs. 1 through 5)
This model is based upon the work of Smith and Beutler (1967), who have
interpreted the ionization layer observed over Electris to be an F-i -region. It
presents the largest reasonable values possible for density and kinetic temper-
atures above 90 km, without going to an E-type model, and utilizes maximum
ultraviolet solar flux and the best available rate coefficients for photochemical
and diffusive processes. (The model has been used to illustrate the discussion
on pages 2 through 8. ) More extreme models are possible if there are errors
in current theories and/or constants.
The calculated neutral number densities and temperatures versus altitude
for the Fj -Model are tabulated in Fig. 1. Figure 2 contains profiles of electron
den^sity and densities of the most abundant molecular and atomic ion species; the
CO2 abundance is very small and therefore is not shown. Figures 3 and 4
respectively give profiles of temperature versus altitude and of neutral density
versus altitude. Figure 5 shows the rate of photoionization of neutral atoms
and molecules above an altitude of 90 km. Since production of electron-ion
pairs by photoabsorption completely dominates other formation mechanisms,
the figure thus represents total ion production rate, which in turn is based upon
the rate of photoabsorption of extreme ultraviolet solar photons at sunspot max-
imum (Fig. 10) .
The Fj -Model results in three ionization peaks characterized by F-type
charge transfer loss processes. In the lower layers of the upper atmosphere
the most abundant molecular ion is O^- Peak concentration of O^ is 37,000
particles cm-3 at an altitude of I6O km. The dominant ion in the upper, or F^-
type, layer is O , which reaches a concentration of 8 . 8 X 105 particles cm-3 at
an altitude of 420 km. In fact, O'*' is the most abundant atomic ion anywhere in
the upper atmosphere of Mars.
The results of the Mariner IV occultation experiment as presented by
Kliore et al . (1967) indicate that only one of these layers, the lowest, was
detected; peak electron concentration (ingress) was 9.0 ±1.0 x 104 electrons
cm-3 a,t an altitude of 123 ±3 km. Smith and Beutler (1966) suggest that with low
solar activity and at large solar zenith angle, as encountered by Mariner IV, the
large upper peak (of the F2-type) at 420 km will disappear. They further irnply
that the two lower peaks at 120 and I6O km will coalesce under these conditions
into a single thin but rather dense O^ peak of the F^-type, creating the peak
observed by Mariner IV. In a brief paper that presents no detailed results,
Donahue (1966) also interprets the ionization layer over Electris to be an Fi -
region and comments that an F2-layer can be suppressed if loss of O"*" via the
reaction
O"^ + CO2 -» CO + O2 (29)
indeed has a rate coefficient of 1 . 2 x lO'^ cm^ sec"-^
July 3, 1967 E. Monash, JPL Sec. 5.4, page 9
Upper Atmosphere JPL 606-1
At an altitude of 90 km, which is roughly the lower boundary of the
thermosphere in the Smith and Beutler Fj -Model, the kinetic temperature,
total gas pressure, and concentration of CO2 are taken from lower atmosphere
Model I of Prabhakara and Hogan (1965). ' The concentration of CO^ for Model
I of Prabhakara and Hogan is changed from 44% to 100%. The assumptions used
to generate the Fj -Model of Smith and Beutler (1967) are listed below.
1) CO2 is assumed to be the only significant constituent of the lower
atmosphere. This is consistent with the results of Mariner IV data.
2) The atmosphere is assumed to be in a quasi -stationary equilibrium.
3) Advection, the horizontal transport of convective cells of matter, is
considered negligible. Inclusion of horizontal transport would
reduce the effectiveness of diffusion.
4) The solar extrenne ultraviolet spectrum used to represent the dis-
sociating and ionizing agent for the upper atmosphere is based on
data from maximum solar activity (Fig. 10). If the spectrum were
based on data from minimum solar activity, the flux contained in
Fig. 10 would be smaller by a factor of approximately two
(Ohring, 1967).
5) Mixing and eddy diffusion are assumed present in the lower atmos-
phere and represent the primary mechanism for the transport of
photodissociation products to this region. With this assumption
Smith and Beutler use the results of Prabhakara and Hogan (1965)
to obtain the concentration of CO2 and temperature at the base of
the photodissociation region.
6) The primary mechanisms that establish the adopted temperature
profile for the upper atmosphere of Mars are assumed to be heating
due to absorption of the extreme ultraviolet solar radiation and sub-
sequent cooling by radiation in the infrared and thermal conduction.
7) Both chemical and diffusive equilibrium are used to establish the
particle concentration in the upper atmosphere.
8) The equation of state for the upper atmosphere and ionosphere of
Mars is assumed ideal.
9) Gravity through the upper atmosphere is assumed to vary inversely
with the square of the areocentric distance (R^ + h) .
10) Accretion or escape of CO^ to the interplanetary medium is assumed
negligible. This implies no net flux or transport of CO^ other than
the removal by photodissociation.
11) The magnetic field of Mars is assumed negligible. A magnetic field
does not affect neutral constituents of an atmosphere. If the neutral
constituents are in diffusive equilibrium, they will maintain this
equilibrium with the injection of a magnetic field. On the other
Sec. 5.4, page 10 E. Monash, JPL July 3, 1967
JPL 606-1 Upper Atmosphere
hand, the molecular and atomic ions will be affected by lines of
force and ultimately will be able to diffuse only along the magnetic
lines of force. A magnetic field would lead to the probable develop-
ment of weak radiation belts high in the upper atmosphere. ^
12) Collisional dissociation and ionization by the solar wind is assumed
negligible. During maximum solar activity the solar wind and
extreme ultraviolet radiation will have comparable energy fluxes.
13) Aerosol concentration is considered negligible.
14) The nnodel is derived for the subsolar region of the atmosphere.
15) Latitudinal and seasonal variations in the atmosphere are neglected.
16) Charge neutrality is assumed; hence, no residual electric fields
exist.
Fp -Model (Figs. 6 through 8)
This model is based upon the work of Fjeldbo, Fjeldbo, and Eshleman
(1966b), who have interpreted the ionization layer above Electris to be an Fo-
region. It should be considered a limiting lower density model with the lowest
reasonable exospheric temperature; it uses minimum extreme ultraviolet solar
flux and the best available photochemical and diffusive constants. A more
extreme model is conceivable only if there are errors in the current theories
and/or constants.
The calculated neutral and electronic number densities and temperatures
versus altitude for the F2-Model are tabulated in Fig. 6. Figure 7 gives the
calculated concentrations of [O], [CO], [CO^J, and [e] for the upper atmos-
phere, the base of which starts at roughly 90 km — the altitude for minimum
temperature. Above this altitude [O] becomes the dominant constituent.
Figure 8 gives the temperature profile above an altitude of 90 km. Above
135 km the temperature becomes isothermal at a value of 80°K.
The F2-Model relies heavily on the assumption that loss of O"^ , which
Fjeldbo et al . assume to be the principal constituent of the ionosphere, is gov-
erned by Eq. (Z9) with a rate coefficient of 10~9 cm^ sec"l. The model does
not include additional nonradiative processes such as those found in Fig. 9-
At an altitude of 120 km the electron concentration and kinetic temperature
of the F^-Model agree with the results from the Mariner IV occultation experi-
ment (ingress), which also indicated that at the surface of Mars (Electris) the
pressure was 4.9 to 5.2 mb, air temperature was 175 to 180°K, and concentra-
tion of CO2 was 1 . 9 to 2. 1 X 10l7 particles cm "3 (Kliore et al. , 1967). The
model parameters are adjusted to match these data. Egress data over Mare
Acidalium indicated a surface pressure of 7 . 6 to 8 . 2 mb and a temperature of
235 to 240°K, but this information was not yet available when Fjeldbo et al.
produced their model. Since this is presented as a limiting model, the lack of
egress data is immaterial.'"
July 3, 1967 E. Monash, JPL Sec. 5.4, page 11
Upper Atmosphere JPL 606-1
The assumptions used to generate the F2-Model of Fjeldbo, Fjeldbo, and
Eshleman (1966b) are as follows:
1) The lower atmosphere is assumed to be pure CO2.
2) The photochemistry assumed to be significant in the upper atmos-
phere is photodissociation of CO^ and O2 by the extreme ultraviolet
solar radiation.
3) Three-body association is assumed to be the dominant raechanism
for formation of O^ and CO2.
4) Molecular diffusion is assumed to be the dominant physical process
that establishes the concentration of constituents.
5) "Hard" corpuscular radiation is assumed to be a negligible ionizing
agent .
6) The temperature profile represents an empirical fit to the data of
the Mariner IV occultation experiment. No energy balance was
assumed.
7) It is assumed that formation of dry ice (condensed CO2) can occur
near the base of the exosphere.
8) Aerosol concentration is assumed negligible.
9) The model is derived for local conditions over the occulting regions
on Mars .
10) Latitudinal and seasonal variations in the atmosphere are neglected.
11) The magnetic field of Mars is assumed negligible.
IZ) The equation of state for the upper atmosphere is assumed ideal.
13) Gravity through the upper atmosphere is assumed to vary inversely
with the square of the areocentric distance.
14) Charge neutrality is assumed; hence, no residual electric fields
exist.
Johnson (1965) has also interpreted the ionization layer above Electris as
an F2-region. The model of Johnson is an empirical model which uses the
plasma scale height observed by Mariner IV and assumes that O"*" is the domi-
nant ion in the upper atmosphere in order to derive an exospheric temperature
of 85° K. The temperature distribution is not obtained from a detailed energy
balance but is assumed to follow the vapor pressure curve for dry ice in the
phase equilibrium diagram." This highly speculative assumption is not based
upon calculated or observed results but represents the result of plausibility
arguments. Fjeldbo et al. (1966a) compare all post-Mariner IV models in a
good review of the subject.
Sec. 5.4, page 12 E. Monash, JPL July 3, 1967
JPL 606-1 Upper Atmosphere
CONCLUSIONS
Lacking detailed information, contemporary models tend to rely heavily
on terrestrial analogy even though the differences in atmospheric composition
between planets make such analogies of dubious value. For example, an Fi
ionization layer in the terrestrial atmosphere will generally appear more fre-
quently near midday and in summer than near sunset or sunrise or in winter
(Yonezawa, 1966). The reason for this is that the level at which the peak in the
electron production profile occurs in the terrestrial atmosphere is a function of
the solar zenith angle. Table VTI of Yonezawa demonstrates that the altitude at
which the peak in the electron production profile occurs increases as the solar
zenith angle increases. If the altitude of the electron production peak lies below
the altitude of the boundary that separates radiative and nonradiative removal
processes, then an Fj-region will appear. If the above conditions are not sat-
isfied, then an F^ -region will not form. The immersion data from Mariner IV
occurred in winter with a solar zenith angle of 67° above Electris (Kliore et al. ,
1967). 12 Therefore, it would seem inappropriate to call the ionization layer
observed over Electris an Fj -layer even though it has some of the properties
of such a layer .
All the models use photochemical and diffusive rate coefficients which are
not well determined. The result of using various values for the rate coefficients
that exist is to produce models that contain a great deal of uncertainty. With an
inappropriate identification of the ionization layer that was observed by Mariner
IV, and the ignoring of predicted layers not observed by the probe, plus the use
of different rate coefficients, plus some "hand-waving," many models can be
produced that will reproduce the Mariner IV data. Figure 11 outlines the
models that are based on the results of Mariner IV. From the figure it is seen
that there are great differences in atmospheric parameters among the various
models. The E- and Fj -Models show a high exospheric temperature while the
F^-Models show a low exospheric temperature. The effort to build a model
that correctly describes the Mariner IV occultation experiment has created
more questions than answers about the structure of the upper atmosphere of
Mars.
July 3, 1967 E. Monash, JPL Sec. 5.4, page 13
Upper Atmosphere
JPL 606-1
h
Height
above mean
surface.*
km
T
Kinetic
temperature
of atmos-
phe ric
constituents.
•K
Concentration.''
cm-3
H(OI
Density
scale height
for atomic
oxygen,
km
[C02]
[CO]
[02]
[O]
90
155
1.89'-3
1.16L^
i.^n
4.32>i
100
137
5. 92 '-2
1.82L^
1.34'-^
1.55'-^
110
132
1.47L2
1.3lL^
2.22'-0
1.26'-^
120
135
3.70Ll
5.28Li
3.585-
7.45I!
20.1
130
149
9.10'-0
2.08U
5.00^
4.20'-i
22.3
140
166
2.O4L?
8.97LO
9.3I
2.48'-i
24. 9
160
200
2.88-
2.08L"
2.4^
9.93'-»
30.4
200
327
1.13«
2.22^
1.6^
2.24'-»
50.9
240
415
1,37!
5.3li
-
9
8.94-
66.0
280
462
2.57^
1.77i
4.55^
75. 1
320
488
6. 10-
6.84-
-
2.57i
81.1
400
508
4.38^
1.261
9.67i
88. 1
480
514
3.70i
2.63*^
3.92-
92.8
S60
516
3.63^
6.00-
1.685-
97.2
640
520
-
9.25I
102. 1
720
520
-
5.22I
106. 1
800
520
3.00-
110.4
900
520
-
1.581
115.6
1000
520
-
8.42'-
121.3
"Smith and Be
utler (1967) <560
^m; Monash
>560 km including H(0)
and [0] ca
culated
from Eq. (7)
page 3.
The supersc
raised. For
ript in concentrati
example. 1 . 89L1
on values is
= 1.89 X IC
the power of ten to whic
13.
h the value
must be
Fig. 1. Upper Atmosphere
F]^ -Model: table of Ccdcu-
lated neutral number den-
sities and temperatures vs.
altitude. H(0) = kT/^JLg =
5.20 X IQl T/g, where g =
g^d +h/Ro)2. Tis in°K,
-2
g m cm sec
sec"^, and R
gQ = 375 cm
= 3380 km.
Fig. 2. Upper Atmosphere
Fj -Model: ion and elec-
tron density vs . altitude.
Note that below 200 km, the
curve for electron concen-
tration coincides with the
O2 concentration curve.
Above 240 km, the curve
for electron concentration
coincides with the O con-
centration curve.
10" 10
NUMBER DENSITY OF IONIZED CONSTITUENTS, cm
Sec. 5.4, page 14
E. Monash, JPL
July 3, 1967
JPL 606-1
Upper Atmosphere
610
570
530
490
450
410
370
330
290
250
210
170
130
90
1
/
/
^
y
^
^
1
225 275 325 375
KINETIC TEMPERATURE, »K
Fig. 3. Upper Atmosphere F^ -Model:
temperature vs. altitude. This figure is a
graph of the second column (Temperature)
of Fig. 1; the same precautions that exist
for Fig. 1 therefore apply.
July 3, 1967
E. Monash, JPL
Sec . 5.4, page 15
Upper Atmosphere
JPL 606-1
LOG (NUMBER DENSITY OF NEUTRAL CONSTITUENTS), cm
Fig. 4. Upper Atmosphere F^-Model: neutral density vs.
altitude. This figure is a graph of the third through sixth col-
umns (Concentration) of Fig. 1; the sanne precautions that exist
for Fig. 1 therefore apply.
10,000
PHOTOIONIZATION RATE COEFFICIENT OF NEUTRAL CONSTITUENTS, ion-eleclron pcirs,/cm ,' sec
Fig. 5. Upper Atmosphere Fj-Model: profiles of photoioniza-
tion rate of neutral constituents.
Sec. 5.4, page 16
E. Monash, JPL
July 3, 1967
JPL 606-1
Upper Atmosphere
HeiKht
T
Kinetic
Cnncent
atmn, '
H(0)
[density
ahi.ve mean
of dtm'is-
surface,
km
phe ric
constituents.
n.vynen
■K
rco^]
[O.
.CO]
-<-■-
'>li
sr,
7.-.S>J?
s.o.ii?
i.skU!
7. U
lorj
SI
l.z./i
1.2.1^0
1 .7« -
2.«ot
7.49
1 1
no
S, t>J-
!. In'i
1 . 7«"-
3. lot
8.8 7
li^n
70
1 , 1 ., -
1 . 00 -
2.00^-
H.sot
10.4
1 in
7S
i. 00-
3,54-^
1. =,4-
6,60l
11.7
1 40
BO
-
1.4ll
1 . 00-
4.781
IZ.O
1 SO
80
7, OH-
3.321
12. 1
u.o
«0
-
l. S4-
2.34I
12.2
170
SO
1 . Hn -
1 . 58l
12. 2
1»0
so
1 , 00-
-
1 . ool
12.3
190
,H0
5.03-
12.4
ZOO
«0
3.S,'^
-
12.4
'The supers
ript in conccn
ration va
lues is tt
e powe r
of ten tc
which the
value must
le raised. y„
exampl
. 7.9SL2
^ 7.95
. loio.
230
220
200
190
150
140
130
120
lie
100
Fig. 6. Upper Atmosphere F2-Model:
table of calculated neutral and elec-
tronic number densities and tempera-
ture vs. altitude. H(0) = kT/|jLg = 5 . 20
X lOl T/g, where g = §^7(1 + h/RQ)2.
T is in °K, g in cm sec"^. g = 375
cm sec"^, and R^ = 3380 km.
\
\
\
\
\
-M
^[']
\
^
\
\
\
\
>^
^^.^[^
oj \
y
[=°2] —
^^^
<\.
^
k
/
. \
^
'^
C""^s*^
'C 160
<
s
O 150
to
<
X
O
120
LOG|g (NUMBER DENSITY OF NEUTRAL CONSTITUENTS),
40 50 60 70 80 90
KINETIC TEMPERATURE, "K
Fig. 7. Upper Atmosphere F2 -Model: number
density of electronic and neutral constituents
vs. altitude.
Fig. 8. Upper Atmosphere
F2-Model: temperature vs.
altitude .
July 3, 1967
E. Monash, JPL
Sec. 5.4, page 17
Upper Atmosphere
JPL 606-1
Hf,u-tiMn
o t ■■
I' -• O 4 .-
r:o'
t O -• 0* t CO • 11. 18 ,.,
co\
• O - O* 1 CO, t 0. In ,
co^
• O, - O* ' CO, 4 u 7i
o^ *
1- -• O + hl>
O*" t
O, -• o' + O + 1 . ^M e.v.
^, ^
hf -• O^ + e
CO*
I O, -• ot ' CO * 1 . -1) e
CO,*
4 0,-0** CO, + 1.71
co;
4 O -• O* 4 CO t 1 . 30 e
°2 •
e ^ + O + b. 92 e.v.
CO t
hi/ ~* CO +- e
CO*
+ CO -• CO* + CO f 0.
CO*
* e -• C 4 O t 2.85 e.v.
CO,
+ hf -• CO, + e
o* ♦
CO, -• O* 4 CO * 1 . 20
CO*
4 (. -• CO t O + 8.23 .-. V
Rate c4jolticie
1 J -1
S, S X 10 cm s
,,-10 i -1
Fig. 9. Table of significant reac-
tions in the Martian ionosphere for
a pure CO2 lower atmosphere.
(Smith and Beutler, 1967)
Fig. 10. Table of incident
solar flux densities at Mars
and absorption cross sec-
tions for selected wave-
length regions. (Smith and
Beutler, 1967)
Intide.i
s<,lar
\Vavfler.,;th
fUix d^.
isity,
.c-l
Abs.jrpt.<,n cr<)
s s,<t.on.
Pholons
Er^s
(Oi)
<C02)
(O)
(CO)
1 'nQii - : 70{(
974 A 10*
10.95
U.25. 10-'"^
0, 0064 X 10"'"
1 70rj- 14,4,1)
"^58
4. 28
2.0
0. 0741
-
1 4,4jrj- 1 \:)'}
1 '1 i
2. 51
8.0
0. 04 52
-
1 iOO- 1 1 DO
ln4
2. 74
0, 5
1 . 88
1 1 00- i UIO
h. 44
0.131
0. 5
17. 5
IOiO-012
5. 33
0. 0942
5. 1 5
5. 1 5
912-885
2.48
0.0434
7. 28
7.28
2.69X 10-"*
885-71 2
1 . 77
0.0 54 3
1 D. 17
12.90
(.05
12.'«. X 10-'«
712-200
5. 88
0. 382
17. 00
12. 34
8. 80
12. 34
200-1 30
1 .06
0. 123
6.61
5. 71
3. 30
5. 71
1 30-90
0. 2 76
0. 0531
1 .67
1 .60
0. 87
1 .60
90-44
0.649
1 .85
0.84
0.65
0.42
0.66
Sec. 5.4, page 18
E. Monash, JPL
July 3, 1967
JPL 606-1
Upper Atmosphere
'ih.iM.n, ! ■■
[. ir>lhn, ,,mH
Fig . 11. Table of models for upper atmosphere of Mars based
on Mariner IV results.
March 1, 1972
E. Monash, JPL
Sec, 5. 4, page 19
Upper Atmosphere jPL 606-1
BIBLIOGRAPHY
Cl-ianiborlain, J . W. , 1963, Planetary coronae and atmospheric evaporation:
Planet, Space Sci. , v. 11, p. 901-960.
Chamberlain, J . W. , and McElroy , M . B . , 1966, Martian atmosphere: the
Mariner occultation experiment: Science, v. 152, p. 21-25.
Coulson,K. L. , and Lotman, M . , 1962, Molecular optical thicknes s of the
atmospheres of Mars and Venus: Philadelphia, Penn. , General Electric
Co., Space Sci. Rep. R62 SD 71.
Craig, R. A., 1965, The upper atmo sphere — meteorology and physics: New
York, Academic Press.
Donahue, T. M. , 1966, Upper atmosphere and ionosphere of Mar s: Science
V. 152, p. 763-764.
Fjeldbo,G., Fjeldbo, W. C. , and Eshleman, V. R . , 1966a, Atmosphere of Mar s:
Mariner IV models compared: Science, v. 153, p. 1518-1523.
Fjeldbo, G. , Fjeldbo, W, C. , and Eshleman, V. R. , 1966b, Models for the
atmosphere of Mars on the Mariner 4 occultation experiment:
J .Geophys.Res. , v. 71, p. 2307-2316.
Fleagle, R.G. , and Businger , J . A. , 1963, An introduction to atmospheric
physics: New York, Academic Pre s s.
Gross, S.H., McGovern, W.E. , and Rasool, S. I. , 1966, Mars: upper atmos-
phere: Science, v. 151, p. 1216-1221.
Johnson, F. S. , 1965, Atmosphere of Mar s: Science, v. 150, p. 1445-1448.
Kliore,A.J., Cain,D.L., and Levy , G. S. , 1967, Radio occultation measure-
ments of the Martian atmosphere over two regions by the Mariner IV
space probe, p . 226-239 in Moon and planets: Amsterdam, North -Holland
Pub. Co.
Norton, R . B. , 1964, A theoretical study of the Martian and Cytherian iono-
spheres: Washington, D. C. , National Aeronautics and Space Administra-
tion, NASA Tech. Note TND-2333.
Ohring,G., 1967, Study of the Martian atmospheric environmental requirements
for spacecraft and re-entry vehicles (esp. p. 21): Bedford, Mas s . , GCA
Corp., Spec. Rep. 67-147 .
Prabhakara, C. , and Hogan, J . S. , Jr . , 1965, Ozone and carbon dioxide heating
in the Martian atmosphere: J . Atmos . Sci . , v. 22, p. 97-109.
Smith, N., and Beutler , A. E . , 1966, A model Martian atmosphere and
ionosphere: Ann Arbor, Mich. , U . of Mich . Radio Astronomv Observatory
Rep. 66-3.
Sec. 5.4, page 20 E. Monash, JPL March 1, 1972
JPL 606-1 Upper Atmosphere
Smith, N., and Beutler, A.E. , 1967, A model Martian atmosphere and
ionosphere: Ann Arbor , Mich . , U. of Mich . Space Physics Research
Laboratory, Preprint.
Weller.C.S., and Biondi.M . A. , l'?67, Measurements of dissociative
recombination of cot ions with electrons: Phy s . Rev. Lett. , v. 19,
p. 59-61.
Yonezawa.T., 1966, Theory of formation of the ionosphere: Space Sci . Rev. ,
V.5, p. 3-56.
March 1, 1972 E. Monash, JPL Sec. 5.4, page 21
JPL 606-1 Cis-Martian Medium, Radiation
SECTION 6 CONTENTS
6. CIS-MARTIAN MEDIUM, RADIATION;
THE MAGNETIC, RADIATION, AND PARTICLE ENVIRONMENT OF MARS
Data Summary 3
Total Irradiation 3
Solar Spectral Irradiance 3
Solar Wind 3
Energetic Particles 3
Solar Interplanetary Magnetic Field 3
Magnetic Field at Mars 3
Martian Meteoroid Environment 4
Discussion 4
Solar Electromagnetic Radiation 4
Total Electromagnetic Irradiation 4
Solar Spectral Irradiance 5
Extreme Ultraviolet Radiation 5
X-rays 9
Radio Wave Radiation IZ
Absorption in the Martian Atmosphere IZ
6. 1 The Particle Environment 13
The Solar Wind 13
Solar Cosmic Rays 17
"Galactic" Cosmic Rays 18
6. 2 Magnetic Fields Zl
Solar Interplanetary Magnetic Field Zl
Martian Magnetosphere and Magnetic Moment 2 2
Surface Magnetic Fields ZZ
6. 3 The Meteoroid Environment Z3
Bibliography 27
Figures
1. Principal features of solar wind flow past Mars 16
2. Cosmic-ray-induced charged-particle flux at the surface
versus atmospheric mass for Mars and Earth 21
3. The miass-flux relationship 25
Tables
1. Solar spectral irradiance (from Thekaekara, 1970) transformed
to Martian mean distance 6
2. Solar XUV spectral irradiance (from Hinteregger, 1970)
transformed to Martian mean distance 10
3. Solar wind proton (H+) data for Earth and derived estimates
for Mars 14
4. Solar wind alpha-particle (He"*"^) data for Earth and derived
estimates for Mars 14
5. Solar wind electron data for Earth and derived estimates
for Mars 15
6. Maximum numiber of solar event protons (cosmic rays) reaching
the Martian surface, assuming an atmosphere consisting of
6 mb of CO2 18
March 1, 1972 Sec. 6, Contents, page i
JPL 606-1 Cis -Martian Medium, Radiation
6. CIS-MARTIAN MEDIUM, RADIATION:
THE MAGNETIC, RADIATION, AND PARTICLE ENVIRONMENT OE MARS
DATA SUMMARY (Sources are given in the Discussion.)
Total Irradiation
58Z.7 W m at mean solar distance
_2
7 09.0 VV m at perihelion
_2
487.5 W m at aphelion
Solar Spectral Irradiance
The distribution o£ the Sun's power with wavelength is shown for ultra-
violet, visible, c'-ind infrared light in Table 1, page 6. Extreme ultraviolet
irradiance is given in Table 2, page 10. Data on x-rays and radio emission
are contained in the body of the text.
Solar Wind
Detailed properties of the solar wind in the near-Mars and near -Earth
regions are listed in Tables 3, 4, and 5, pages 14 and 15, respectively.
I'^nergetic Particles
The maximum number of solar event protons reaching the Martian sur-
face are given in Table 6, page 18.
Solar Interplanetary Magnetic Field
Strength at 1.5 AU, the mean Strength is dependent upon solar
distance; of Mars froni the Sun activity and may fluctuate 1 to 2
orders of magnitude.
-5
Average Zy (1 y = 10 gauss)
Range to 25 7
Magnetic Field at Mars
The following upper limits are based upon the apparent absence of a
Martian shock wave or magnetosphere along the Mariner IV trajectory. The
magnetic moment of the Earth = M^ = 8.05 ±0.02 X 10^5 gauss cm^.
E
_4
Martian magnetic moment, 3X10 M
upper limiit
Surface magnetic field at 100 y
equator, upper limit
June 15, 1971 E. Haines, R. Newburn, JPL Sec. 6, page 1
Cis -Martian Medium, Radiation JPL 606-1
Piled-up interplanetary field -35^
at subsolar point (if internal
moment is zero)
Martian Meteoroid Environment
The terrestrial meteoroid mass-flux relationship is shown in Fig. 3,
page Z5. There are indications that the dust flux may be three to five times
laro-er at Mars. There is no data on larger bodies. Mars will encounter mete-
oroids with relative velocities between 10 and 58 km sec"-^. See the main text
for details.
DISCUSSION
Solar Electromagnetic Radiation
Total Electroniagnctic Irradiation
By definition, the solar constant is the total electromagnetic irradiation
per -unit-area normal to a solar radius vector, at a distance of one astronom-
ical unit (AU) from the Sun and outside the Earth's atmosphere. Thus, it is a
direct measure of the solar power output. At Mars' mean distance of 1.5237 AU
from the Sun, the total electromagnetic irradiation is only 0.4307 of the solar
constant, while at perihelion it is 0.5240 and at aphelion 0.3603.
For many years it was convenient to use Johnson's (1954) value of the
solar constant, 2.00 cal cm'^min"^ (1395 W m-2). Soviet balloon flights in
the 1960's to altitudes near 30 km, seemed to indicate an even larger value,
such as 2.016 cal cm'^^min"! from Kondratiev et al. (1967). Soviet analyses
of the probable value of the sokir constant reflect this work; for example,
2.03 ±0.15 cal cm'^^min"^ derived by Makarova and Kharitonov (1969).
More recently, Laue and Drummond (1968) have published a new value
for the solar constant which includes direct measurements made from an X-15
aircraft at an altitude of about 82 km. They found it necessary to reduce the
contribution for \ > 6070 A by about 7"'), compared to Johnson's work, and
derived a solar constant of 1.952 ±0.02 cal cm-^min"! (1361 W m"^). The
Mariner VI and VII spacecraft each carried a solar flux monitor to Mars in
1969. The data from these flights was quite consistent, giving a value of
1352.5 \\^ m"'^ (Plamondon, 1969). Plamondon feels a probable error of ±1%
was achieved.
An analysis of all solar constant measurements by the ^Standard's Sub-
committee of the Solar Radiation Committee of the lES suggests that 1353
±21 \V m"2 be adopted as the best value for the solar constant (Thekackara,
1970). However, there are still disagreements and uncertainties relative to
the various measurements, and a JPL committee is continuing work on the
j^roblem. A very comprehensive and useful discussion of work to 1967, in-
cluding extensive tables, has been prepared by Labs and Neckel (1968).
Sec. 6, page 2 R. Newburn, JPL June 15, 1971
JPL 606-1 Cis-Martian Medium, Radiation
Using 1353 W m-2 for the solar constant, the total electromagnetic
irradiation of Mars is
58Z.7 W m-2 at nnean solar distance
709-0 W m"2 at perihelion
487.5 W m-2 at aphelion
Solar Spectral Irradiance
The distribution with wavelength of the Sun's power can be approximated
roughly in visible light by that of a 5900°K blackbody, although the Sun's effec-
tive temperature (the temperature of a blackbody having the same total power
output as the Sun) is only about 5765 °K. The differences are caused by line
blanketing, the effect of the absorption in the many Fraunhofer lines of the solar
spectrum, and by the variation in photospheric depth (and therefore tempera-
ture) of the continuum observed at different wavelengths. The departure from
blackbody radiation is far more extreme in the ultraviolet than in the visible
wavelengths. Ultraviolet irradiance data is usually given in tabular form as a
function of wavelength due to these variations. Table 1 presents the data cited
by Thekaekara (1970) transformed to the mean solar distance of Mars.
Extreme ultraviolet, x-ray, and radio radiations from the Sun are direct
functions of the solar cycle and of sporadic solar events. Such radiation con-
stitutes an insignificant fraction (<2 X 10"^) of the total power output of the Sun,
but the very short wavelengths naturally have very high energies. Therefore,
they play a major role in various excitation and ionization phenomena. Varia-
tions in these extreme wavelengths are discussed in the following subsections.
Extreme Ultraviolet Radiation
The "extreme ultraviolet" is that part of the electromagnetic spectrum
beginning at approximately 50 A and extending up to about 1800 A, at which
point quartz begins to transmit and ordinary photographic emulsions become
usable. Above 1800 A there is strong and constant continuum radiation from
the Sun. At wavelengths shorter than 1600 A most solar power is in discrete
emission lines. No absorption lines are found below about 1700 A. The entire
extreme ultraviolet is dominated by the resonance line of atomic hydrogen
(Lyman o), at 1215.7 A, which emits a greater flux than all other emission
lines combined.
A basic problem in utilizing extreme ultraviolet solar data is the varia-
tion caused by the solar cycle, solar rotation, and solar flares. Although
rocket-borne spectrometers, etc. , have been used to study the Sun since
shortly after World War II, reliable quantitative results are not easy to obtain
even today, and it is difficult to separate true variations from calibration error,
particularly when comparing results of different equipment and experimenters.
In general, the variation for X. > 1300 A is probably small. Hinteregger (1970)
-suggests that the sum of various strong emission lines in the region \X 280-1300,
probably shows a relative variation with solar cycle of about 75% that of the
10.7 cm radio emission, a factor of about two. The variation of lines and
June 15, 1971 R, Newburn, JPL Sec. 6, page 3
Cis -Martian Medium, Radiation
JPL 606-1
Table 1. Solar spectral irradiance (from Thekaekara, 1970) transformed
to Martian mean distance.
X''
Pk"
Ax
<
X'-'
p^l)
Ax'
Dx"
0. 120
0.000004
0.00025843
0.00044
0.335
0.04657
1.937)7
3.323
0.140
0.000001
0.00031447
0.00053
0.340
0.04627
2.16926
3.721
0.150
0.000003
0.00033600
0.00057
0.345
0.04605
2.40005
4.117
0.160
O.OOOOIO
0.00040062
0.00068
0.350
0.04708
2.63288
4.517
0.170
0.000027
0.00058586
0.00100
0.355
0.04665
2.86723
4.919
0.180
0.000054
0.00099079
0.00169
0.360
0.04601
3.09888
5.316
0.190
0.000117
0.00184373
0.00316
0.365
0.04876
3.35381
5.723
0.200
0.000461
0.007317
0.0081
0.370
0.05088
3.58491
6.150
0.210
0.000986
0.011969
0.0205
0.220
0.00248
0.029286
0.0502
0.375
0.04984
3.83670
6.582
0.380
0.04825
4.08192
7.003
0.225
0.00280
0.042468
0.0728
0.385
0.04730
4.32077
7.413
0.230
0.00287
0.0566407
0.0971
0.390
0.04730
4.55726
7.819
0.235
0.00255
0.0702102
0.1204
0.395
0.05122
4.80358
8.241
0.240
0.00271
0.0833813
0.1430
0.400
0.06156
5.08553
8.725
0.245
0.00311
0.0979525
0.1680
0.405
0.07082
5.41646
9.293
0.250
0.00303
0.113321
0.1944
0.410
0.07543
5.78210
9.920
0.255
0.00448
0.132103
0.226
0.415
0.07642
6.16170
10.571
0.260
0.00560
0.157303
0.269
0.420
0.07526
6.54092
11.222
0.265
0.00797
0.191227
0.328
0.270
0.00999
0.236136
0.405
0.425
0.07293
6.91139
11.858
0.430
0.07060
7.27023
12.473
0.275
0.00879
0.283091
0.485
0.435
0.07164
7.62584
13.083
0.280
0.00956
0.328969
0.564
0.440
0.07797
7.99984
13.725
0.285
0.0136
0.386801
0.663
0.445
0.08280
8.40176
14.415
0.290
0.0208
0.472630
0.810
0.450
0.08641
8.82479
15.140
0.295
0.0252
0.587433
1.007
0.455
0.08861
9.26237
15.891
0.300
0.0221
0.705682
1.210
0.460
0.08900
9.70638
16,653
0.305
0.0260
0.825978
1.417
0.465
0.08822
10.1494
17.413
0.310
0.0297
0.965120
1.655
0.470
0.08758
10,5803
18.167
0.315
0.0329
1.12160
1.924
0.320
0.0358
1.29327
2.218
0.475
0.08805
1 1.0280
18.<?21
0.480
0.08934
11.4715
1Q.681
0.325
0.04200
1.48766
2.552
0.485
0.08512
1 1.9077
20.430
0.330
0.04562
1.70671
2.928
0.490
0.08400
12.3305
21.155
■'x
Wavc^onpth in r
nic rons .
1 ,
.Sol.ir spectra
irradiance av
<'raf;e(l over
small b.UK
Iwitlth centereci
at X . in watt ^
-2 - 1
cm (I
'A,
Area unrlcr Ih
(■ solar spell r
il irradi.mc
• curve in t
ho w.welenjjlh
ranee to X ,
n \\^ cit;
"n.
P(^ rci'iif a^i; of
ihi: sol.Lr i r r a
fliaiici' as so
iateil with
w.ivelon^ths sh
orter than X.
Sec. 6, page 4
R. Newburn, JPL
June 15, 1971
JPL 606-1
Cis -Martian Medium, Radiation
Table 1. Solar spectral irradiance (fronn Thekaekara, 1970) transformed
to Martian mean distance, (cont'd)
v
(1.5(H)
0.505
0.510
0.515
0.520
0.52 5
0.530
0.535
0.540
0.545
0.550
0.555
0.560
0.565
0.570
P-
0.0H443
n.O«366
0.08270
0.08107
0.07896
0.07896
0.07978
0.07935
0.07832
0.07681
0.07556
0.07431
0.07409
0.07302
0.07345
0.07375
0.575
0.07405
0.580
0.07388
0.585
0.07375
0.590
0.07323
0.595
0.07246
0.600
0.07177
0.605
0.07095
0.610
0.07043
0.620
0.06901
0.630
0.06763
0.640
0.650
0.660
0.670
0.680
0.690
0.700
0.06651
0.06509
0.06401
0.06272
0.06147
0.06040
0.05897
12.7516
13.1718
13.5877
13.9972
14.3973
14.7921
15.1889
15.5867
15.9809
16.3687
16.7496
17.1243
17.4953
17.8631
18.2293
18.5972
18.9668
19.3366
19.7056
20.0731
20.4373
20.7979
21.1547
21.5082
22.2054
22.8886
23.5593
24.2173
24.8629
25.4965
26.1175
26.7268
27.3269
n.
21.878
22.599
23.312
24.015
24.701
25.379
26.059
26.742
27.418
28.084
28.737
29.380
30.017
30.648
31.2 76
31.907
32.541
33.176
33.809
34.439
35.064
35.683
36.295
36.982
38.098
39.270
40.421
41.550
42.657
43.744
44.810
45.855
46.879
0.710
0.05790
0.720
0.05660
0.730
0.05557
0.740
0.05428
0.750
0.05320
0.800
0.04769
0.850
0.0426
0.900
0.0383
0.950
0.0360
1.000
0.0321
1.100
0.0255
1.200
0.0208
1.300
0.0171
1.400
0.0145
1.500
0.0124
1.600
0.0105
1.700
0.00870
1.800
0.00685
1.900
0.00543
2.000
0.00444
2.100
0.0039
2.200
0.0034
2.300
0.0029
2.400
0.0028
2.500
0.0023
2.600
0.0021
2.700
0.0019
2.800
0.00168
2.900
0.00151
3.000
0.00134
3.100
0.00112
3.200
0.000974
3.300
0.000827
27.9080
28.4805
29.0414
29.5907
30.1281
32.6503
34.9065
36.9279
38.7846
40.4872
43.3688
45.6864
47.5818
49.1585
50.5003
51.6441
52.6047
53.3823
53.9961
54.4894
54.9051
55.2691
55.5857
55.8700
56.1242
56.3439
56.5399
56.7165
56.8759
57.0180
57.1408
57.2455
57.3355
D.
47.882
48.864
49.826
50.769
51.691
56.818
59.889
63.358
66.543
69.464
74.409
78.385
81.637
84.342
86.645
88.607
90.255
91.589
92.642
93.489
94.202
94.826
95.370
95.858
96.294
96.671
97.007
97.3103
97.5838
97.8277
98.0383
98.2179
98.3724
VV^avc'lenyth in microns.
Solar spectral irradiance averaged over small Ijandwidth centered at \, in watts cm"'^ ^"'.
Area unrl(;r the solar spectral irradiance curve in the wavelength range to\, mW cm"^.
Percentage of thi^ solar irrarliancc associ.ited with wavelengths shorter than X.
June 15, 1971
R. Newburn, JPL
Sec. 6, page 5
Cis -Martian Medium, Radiation
JPL 606-1
Table 1. Solar spectral irradiance (from Thekaekara, 1970) transformed
to Martian mean distance, (cont'd)
\''
Pk"
^{
Dx"
\^
Px'^
Ax
Dx'
3.41)0
0.000715
57.4126
98.5047
9.0
0.0000164
58.234370
99.913939
3.500
0.000629
57.4798
98.6200
10.0
0.0000108
58.247939
99.937220
3.0OO
0.000582
57.5406
98.7238
11.0
0.00000732
58.256986
99.952742
3.700
0.000530
57.5961
98.8192
12.0
0.00000517
58.263232
99.963458
3.HO0
0.000478
57.6465
98.9056
13.0
0.0000037
58.267691
99.971 1 08
3.900
0.000444
57.6926
98.9847
14.0
0.0000024
58.270749
9v. 976356
4.00
0.00041
57.7353
99.0579
15.0
0.0000021
58.27298V
99.98019'^
4.100
0.00037
57.7745
99.1252
16.0
0.0000016
58.274863
99.983414
4.200
0.00034
57,8098
99.1861
17.0
0.0000013
58.275918
99.985964
4.300
0.00031
57.8421
99.2412
18.0
0.0000010
58.277534
99.987997
4.400
0.00028
57.8714
99.2915
19.0
0.00000086
58.278482
99.989623
4.500
0.00025
57.8981
99.3373
20.0
0.00000067
58.279257
99.990953
4.600
0.00023
57.9222
99.3787
25.0
0.000000263
58.281639
99.995836
4.700
0.00021
57.9438
99.4160
30.0
0.000000129
58.282617
99.996718
4.800
0.00019
57.9640
99.4504
35.0
0.0000000689
58.283112
99.997568
4.900
0.00018
57.9826
99.482195
40.0
0.000000040
58.283388
<:»9. 998037
5.000
0.0001650
57.999810
99.511500
50.0
0.000000016
58.283672
99.998525
6.000
0.00007539
58.119998
99.717708
60.0
0.0000000082
58.2837<^3
99.998736
7.000
0.0000426
58.179014
99.818965
80.0
0.000000003
58.283905
99.998928
H.OOO
0.0000258
58.213261
99.877723
100.0
0.000000001
58.283948
99.999002
1000.0
1 58.284534
100,0
'' V '.V,i\rolont;th in microns.
1 -2 -
\'\ Solar sjM'CtrHl irrarlianco avcirafrod ovor small baiicKviclth c-onter cd at \, in watts cm ^
' Av Area unrlcr the sclar spi'ctral irradiance curve in the wavelength range to \ , mW cm
]) ^ Pcree'ltafie of the solar irradi.mce associated with w.avelenpths shorter than \.
radio ennssion with solar rotation is more or le.ss comparable, again a factor
of about tvo (Hinteregger , 1970). In both cases, m^li-'idual -mi s sion line--- •^hov
changes ranging from very small to several times the average. Variation r.f
individual liijes in the W 30-280 region if uncertain, bn<- at ♦;'ie high ener-," ' >:'i
of this region Kreplm (1970) found a factor-of-ZO increasf- from ■:' ■ -i,- r,r '■ ■;-.
imum to sunspot maxiinum in a band from 44 A to bO A.
The data on spectral irradiance in Table 2, Part 1, were obtained from
the satellite OSO-III on March 11, 1967, a time of medium solar activity
(Hinteregger, 1970), and have been transformed to Martian mean distance, as
have Parts 2 and 3 of Table 2. The 10.7 cm radio flux at Earth on that date
Sec. 6, page 6
R, Newburn, JF'L
June 15, 1971
JPL 606-1 Cis -Martian Medium, Radiation
was 144 X 10-22 ^^ ^-Z Hz-1. The data in Table 2, Part 2, were obtained by
rocket at an altitude of 210 km on April 4, 1969, and were quoted by Hintoreggcr
(1970) from unpublished data of Heroux et al. The 10.7 cm flux at Earth was
177 X 10-22 w" nn-2 Hz-1, indicating somewhat greater solar activity but still
not "high" activity. The data in Table 2, Part 3, are from unpublished work by
Manson, quoted by Hinteregger (1970). They were taken on August 8, 1967,
when the 10.7 cm flux level at Earth was 143 X 10-22 w m-2 Hz'^, a medium
active Sun. These data collectively offer a reasonable picture of the average
extreme ultraviolet flux. A first approximation to the flux level under other
conditions can be obtained by multiplying the UV flux by the ratio of the 10.7 cm
radio flux level on Earth at the time in question to 150 X 10-22 \y rn-2 Hi^-l.
Unless the radio flux ratio varies more than 30% from unity, there is little
point in making the correction however, since there is at least that niuch uncer-
tainty in the data.
X-rays
Although a more precise definition could be given in terms of the origin
of the photon, the region of the electromagnetic spectrum between 0.1 A and
about 50 A is normally considered the x-ray region. The x-radiation from the
Sun consists of a continuum, which probably originates in recombination and
bremsstrahlung radiation of plasmas in active solar coronal regions, and super-
imposed emission lines of very highly ionized coronal atoms, many of them
stripped down to one electron.
There are large variations in solar x-ray flux, even during nonflare con-
ditions, and the^ variation increases toward shorter wavelengths. The flux at
1 AU in the 44 A-60 A band was 3 X 10-4 W m-2 in 1969, while at sunspot mini-
mum in 1964 it was only 1.5 X lO'^ W m-2 (Kreplin, 1970). In the 8 A-20 A
band, the flux was ~3 X 10-5 w m-2 in 1969 and more than 200 times less in
1964, undetectable in^the equipment then available (Kreplin, 1970). The 196^5
flux at 1 AU in the A-B A band was ~2 X 10-6 W m-2. As for visible radia-
tion, these flux figures can be transformed to Martian equivalents through mul-
tiplication by the inverse square of the distance in astronomical units, a factor
of 0.4307 for the mean distance of Mars from the Sun.
There is a sizable variation, of shorter period, as the Sun rotates and
brings regions of greater and lesser activity into view. During a 7-month
period of observations, from October 1967 thr^ough April 1968, using the satel-
lite QSO-4, there was a 7-to-l variation in 8 A-16^A flux, a 12-to-l variation
in 3^A-9 ^ flux, and a 120-to-l variation in 1 A-3 A flux (Pounds, 1970). In the
44 A-60 A band, the variation was only about 1.7 to 1 during March 1966
(Kreplin, 1970).
There appears to be a close relationship between solar Ho flares and
x-ray bursts or flares. Pounds (1970) suggests it is quite probable that the
correlation is 100%. During an import^ance 2B flare on April 11, 1967, Pounds
(1970) quotes data showing the 8 A- 12 A band flux increased by a factor of 20
in five minutes, and then decayed much more slowly (perhaps over several
hours). Whereas x-rays may be detectable only down to 1.5 A or 2 A during a
quiet Sun, x-rays with energies above 22 keV (<0.6 A) are common during
flares (Pounds, 1970).
June 15, 1971 R. Newburn, JPL Sec. 6, page 7
Cis -Martian Medium, Radiation
JPL 606-1
Table Z. Solar XUV spectral irradiance (from Hinteregger, 1970)
tr ansfor2-ned to Martian mean distance.
Pari 1
\ or
a ng('
(A)
1 ".Ot..O-l i04.0
1 30Z.Z
1260.7
1Z42.8
1238.8
1215.7
1206.5
1175 group
1 128.3
1122.5
1085 group
1037.6
1031.9
1 310-1027
1310-1027
1310-1027
1025.7
990 grou]:)
977.0
972.5
949.7
944.5
937.8
933.4
930.7
926.2
1027-91 1
1027-91 1
91 1 -890
904 group
890-860
860-830
Id,
if i fi c'.iti on
C) 1
O !
Si II
Si II
N V
N V
U Ly - o
Si III
C III
Si IV
Si IV
N II
O VI, C II
O VI
unresolved
excl. II Ly-o
integral
1 1 Ly - P
N III
C III
II Ly-Y
11 Ly--
S VI
II Ly - c
S VI
II Ly - t;
II Ly- r,
unresolved
int(.'gral
H conL
r. U
II cont.
11 cont.
I
0.00646
0.0035
0.00 31
0.0016
0.0020
0.0030
~ 2.2
0.025
0.016
0.0020
0.0016
0.00465
0.014
0.019
0.029
0.13
~2.3
0.029
0.00521
0.039
0.00706
0.0031
0.00082
0.0020
0.0012
0.0012
0.0012
0.01 1
0.099
0.030
0.0012
0.027
0.01 5
o
\ or Kange (A)
0.43
0.23
0.20
0.099
0.12
0.19
-129
1.6
0.95
0.12
0.095
0.25
0.73
0.99
1.6
7.54
-138
1.5
0.26
1.9
0.34
0.17
0.04 3
0.095
0.0 56
0.05i>
0.056
0.521
5.00
1.4
0.056
1.2
0.65
IdonI ifif ation
835 gro\ip
830-800
91 1-800
911 -800
800-770
790.2, 790.1
787.7
786.5
780.3
770.4
770-740
765.1
760
740-710
710-680
703 group
800-630
800-630
629.7
625.3
609.8
599.6
584.3
554 group
521.0
508 g rou])
504 -
499.3
465.2
630-460
6 30-460
460-370
368.1
O II, III
II cont.
unresolved
integral
II cont.
O IV
O IV
S V
Xe VIII
Xe VIII
1 1 cont.
X IV
O V
Jl cont.
II cont.
O HI
unresolved
integral
O V
Mg X
Mg X
O III
lie I
O I\^
Si XII
O III
He I cont.
Si XII
Ne VII
unr e sol\- ed
integral
integral
Mg IX
1
*
0.00530
0.00818
0.00099
0.086
0.0042
0.0028
0.0014
0.00086
0.001 3
0.0025
0.0022
0.0020
0.00082
0.0010
0.00052
0.0028
0.00482
0.028
0.012
0.0035
0.0069
0.0012
0.013
0.00474
0.0032
0.001 3
0.0086
0.00655
0.0029
0.00719
0.07^
0.01 3
0.013
0.22
0.34
0.04 5
3.6
0.17
0.11
0.05t)
O.OM
0.052
0.0'i't
0.082
0.078
0.034
0.03Q
0.01 7
o.O'ig
0.17
1.03
0.40
0.11
0.22
0.034
0.38
0.13
0.082
0.034
0.22
0.1b
0.0b<1
0.19
2.0
0.27
0.24
' I fl ux, er g cni sec
o
1, , -2-1
I flux, in units ol 10 photons cm sec
Sec. 6, page
R. Newburn, JPL
June 15, 197 1
JPL 606-1
Cis -Martian Medium, Radiation
Table 2. Solar XUV spectral ir radiance (from Hinteregger, 1970)
transformed to Martian mean distance, (cont'd)
Part 1 (cont'd)
\ Of i;,.n-r (A)
Iflcnl i fic.it ion
I ■•'
o
o
K or Uange (A)
Identification
I '-^
o
o
Ui4.8
0.0039
0.073
284.1
Fe XV
0.033
0.47
U.i(1.7
Fe XVI
0.0086
0.16
370-280
unresolved
0.0526
0.874
■: '0.4
Fc XVI
0.019
0.31
370-280
integral
0.28
4.44
•'. n -. . ,H
![,■ II
0.I5Z
2.3
Part 2
\i.r i;,,nyr (A)
Irli ntification
I^
)i. or Range (A)
Identification
I^-
^iSO-^ ■, 1
integral
0.099
176-153
integral
0.047
Z'U - j! (1 S
integral
0.056
205-153
integral
0.21 ]
IHU - in^
integ ral
0.1 55
153-100
integral
0.026
i05-17b
integral
0.164
Part 3
K CM- ]v ,1 HLJ (* (A )
Identification
I^
\or Range (A)
Identification
f-
1^8-1^(1
integral
0.0009
80-70
integral
0.011
1Z0-! l(i
integral
0.0017
66.3
Fe XVI
0.0009
](l -,.(,, KiS.^
Fe IX
0.00082
70-60
integral
0.014
1 1 0-100
integral
0.0051
50.5, 50.7, 55.3
Si X/IX
0.00 39
')4.0, Qf,.l
Fe X
0.0013
60-50
integ ral
0.01 3
1 r) - '->
integral
0.0090
44.1
Si XII
0.0013
80. S, Hf,.H
Fe XII/XI
0.0009
50-40
integral
0.0090
QO -80
i n t e g r ,a 1
0.013
80-40
integ ral
0.0465
1 .i - H
integral '
0.028
33.6
C VI
0.0 00 9
76.0
F<- XIII
0.0009
i I Oi>: , e r u rni si'c
\, 9 . 2. - ^
I liiix, ill units of 10 photons cm s(-c
flux, (■ r u < ■ 1 1
-2 - 1
see
June 15, 1971
R. Newburn, JPL
Sec. 6, page 9
Cis-Martian Medium, Radiation JPL 606-1
Radio Wave Radiation
At millimeter wavelengths, the Sun radiates very like a blackbody at a
temperature of about 5800°K. Near one centimeter, the brightness tempera-
ture begins to increase rapidly, however, and by the time a wavelength of three
meters is reached, the brightness temperature (of even the quiet Sun) has risen
to 106°K (Smith, 1967). The disturbed Sun may exhibit a brightness temper-
ature of 10^0°K at wavelengths of 3-10 meters (Castelli et al. , 1965), but the
flux involved is only about 10" 18 \y ni-2 Hz"!, at 1 AU. The detailed radio
frequency behavior of the Sun is quite complex, and is of interest primarily in
attempting to understand the Sun itself, as the small amounts of power and low
energies involved have little effect upon the Moon or planets.
It is useful to note that there is excellent correlation between the 10.7 cm
flux and the sunspot area on the solar disk. The correlation with XUV flux was
noted in previous paragraphs, and good correlation exists with the x-ray flux
(Pounds, 1970). Daily measurements of the 10.7 cm solar flux have been made
since the beginning of 1947, to serve as an indication of the level of solar activ-
ity (Castelli et al. , 1965).
Absorption in the Martian Atmosphere
The atmosphere of Mars is principally CO2, which has no significant
absorptions in the visible part of the spectrum. Ravleigh scattering will re-
move perhaps 3% of the incoming radiation at 6OOOA (Young, 1969). There
appear to be atmospheric aerosols generally present, but these are quite tenu-
ous (Leovy et al. , 1971) and in all probability create Little atmospheric opacity
or loss of contrast in the visible spectrum (Smith', Young, and Leovy, 1970;
Van Blerkom, 1971). There is also evidence of occasional major dust storms
which do cause great opacity (see Section 4. 1). Most of the time the solar spec-
tral irradiance falling on the Martian surface, at visible wavelengths, must be
virtually that which is incident upon the upper atmosphere.
In the infrared, CO2 has very strong rotation-vibration bands absorbing
at 15, 4.3, and 2.7 \x, plus additional weaker absorptions extending down to the
visible, but growing weaker with decreasing wavelength. The rare dust storm
would also create infrared opacity at wavelengths roughly equal to the particle
size or smaller.
In the ultraviolet, Barth and Hord (1971) have presented evidence for
scattering about three times as great as would be expected from pure Rayleigh
scattering. This is not necessarily inconsistent with the observations in the
visible, but does raise some question of disagreement unless rather special
assumptions are made about particle size or distribution. Ozone is the princi-
pal atmospheric absorber between 2000 A and 3000 A, and there is evidence for
perhaps 10 fi-atm of O3 in the Martian atmosphere (see Section 5. 1). Neverthe-
less, most sunlight X > 1800 k must still reach the Martian surface. While
there are significant absorptions by CO2 below 1800 A, there is relatively little
solar continuum to be absorbed at short wavelengths. Observational studies of
Martian aeronomy have been initiated by Barth et al. (1971) with their Mariner
6 and 7 experiment, but quantitative details of ultraviolet fluxes reaching the
Sec. 6, page 10 R. Newburn, JPL June 15, 1971
JPL 606-1 Cis-Martian Medium, Radiation
ground at X < 1800 A are not yet available. It can be stated that most x-rays
should be blocked from the Martian surface by the atmosphere, while perhaps
half of the y-rays reach the surface (see the following subsection, especially
Fig. 2).
6. 1 THE PARTICLE ENVIRONMENT
Particles arrive in the vicinity of Mars from a variety of sources. By
far the largest flux comes from the Sun, gently in the form of solar wind, or
explosively from solar flares and flare-related phenomena. By comparison,
the flux of cosmic rays is negligible, but the effects of these particles are sig-
nificant since one cosmic ray particle may contribute as much as 10 joules of
energy to the Martian atmosphere and surface, and create thousands of radio-
active nuclei. The nature and effects of the solar wind, solar flares, and cos-
nnic rays are discussed in this subsection.
The Solar Wind
The solar wind originates in the solar chromosphere, a sheath about
0.01 R thick, located outside the optically observed solar surface or photosphere.
The temperature of the chromosphere rises from 5 X IQS'k, at its boundary
with the photosphere, to about lO^'K at its outer limit, where it merges into the
corona and then continues to rise to a maximum of about 2 X lO^'K in the corona.
This nonequilibrium increase in temperature is believed due to acoustical or
magnetic noise created inside the photosphere, from turbulent hydrogen convec-
tion, and absorbed preferentially by the thinning gases of the chromosphere
(Brandt, 1970, Chap. 3). The acceleration of the hot chromospheric gases into
a supersonic plasma flow has been studied by Parker (1958). His theoretical
treatment uses the confining force of the Sun's gravity to form a "converging
throaty analogous to the throat of a rocket engine. Hydrodynamic acceleration
to sonic speeds in the throat leads to continued supersonic acceleration and
direct motion in the rarefied gas beyond the throat. Only the Parker model has
satisfactorily accounted for the high velocities and temperatures observed in
the solar wind.
Parameters generally describing the composition and state of the solar
wind are given in Tables 3 through 5. These are the results of spacecraft ex-
periments carried out in the area roughly between the orbits of Earth and Venus
and then extrapolated to Mars' distance from the Sun (although Mariner IV
carried an ion chamber which functioned part of the way to Mars and a solar
plasma experiment which operated throughout the mission).
The composition of the solar wind is known only in barest outline. The
earliest experiments on Mariner II (see Neugebauer and Snyder, 1962) revealed
alpha particles in the spectrum, along with protons. The ratio of alphas to
protons varies widely, from to about 0.25, averaging about 0.05 (Strong et al. ,
1970). The photospheric value is 0.1. The (average) preference for ejection of'
protons over alpha particles is understood in terms of diffusion in the corona
after the plasma has been thoroughly mixed in the chromosphere (Jokipii, 1966).
June 15, 1971 E, Haines, R. Newburn, JPL Sec. 6, page 11
Cis -Martian Medium, Radiation
JPL 606-1
Table 3. Solar wind proton (H ) data for Earth and derived estimates
for Mars (Neugebauer, 1971).
Solar wind proton
data for
Earth
Minimum
Average
Maximum
Mars
Minimum at aphelion
Average at aphelion
Average at perihelion
Maximum at perihelion
Velocity,
km sec ■ '
150
400
1000
150
400
400
1000
Density,
cm-3
1
6
200
0.2
2
3
100
Temperature,
°K
6 X lO""
1 X 10"
1 X lO"-
6 X 10-
1 X 10-
1 X 10'
1 X 10
nw
Directional
flux
1 X 10
3 X 10^
5 X 10^
4x10^
1X10°
9
2 X 10
3 X lo"^
nmv /2-'-
Energy density.
ergs cm- 3
3 X 10'
1 X 10
2 X 10
1 y 10
10
10
5 X IQ- -
7 X 10""^
1 X 1 0' '
mv"/ 2
Energy per
particle, t-v
1 • lo'-
8 X 10^
5 X 10^
1 ,-■ lo'
8 / lO'^
8 ■ IC^
5 V 10 '
:-Extreme values of nv and nmv /2 do
a generally inverse relation between
not correspond to the product of extreme values of n and v, due to
n and v.
+ + ^
Table 4. Solar wind alpha-particle (He ) data for Earth and derived
estimiates for Mars (Neugebauer, 1971).
Solar wind alpha-
particle data for
V
Velocity,
km sec - 1
n
Density,
cm - 3
T
Temperature,
°K
nv-'i
Directional
flux
nmv ! l-
Energy density,
ergs c m - ^
F!ne-rL'y per
pa f ti c 1 e , e\-
Earth
Minimum
150
6 X 10^
5
10"
A^■ e rage
400
0.3
4 X 10^
1 / lo'
1 X IQ- *
'-'•
1 -■■
Maximum
1000
2
2 xlO^
1 X 10*^
1 X 10'*^
2
10^
Mars
Minimum at aphelion
150
6 X 10^
5
10-
Axerage at aphelion
400
0.1
4 > 10^
4.10^
4. lO-'O
^.
1 "'
A\-erage at perihelion
400
0.1
4 ,10^
5 . 10^
Id"
Maximum at perihelion
1000
1
2 .- 10^
5 " lO'
_ Q
5 ■ 10 ■
1
10^
Extrcii'.c \-alucs of n\- an
1 n I ] 1 v '' / 2 do
nc>t corri.
s)5oncl tn the pr
oduct of extre
lie val ues oi n .ti
.'1 V, r
me I. •
,a penerally inverse rcla
ion ])et\^'Con
n and v.
Sec. 6, page IZ
M. Neugebauer, JPL
June 15, 1971
Jl^L 606-1
Cis -Martian Medium, Radiation
Table 5. Solar wind electron data for Earth and derived estimates
for Mars (Neugebauer, 197 1).
Sola
fltrcl r on
V
\'.-locit\V
km sec " 1
n
Density,
c m " ^
Tcmnerature,
nvjh/4
Omnidi rectional
flux/'
cm "^ sec - 1
3/2 n'r-"I
Energy densit\-,
<'rgs cni"-
panic If.-, 1 -
Mini
-] ■i\ui .
1 50
1
7 . 10^
4 :■ 10'
2 ,. 10-1'
r.
A\-er
1 ^ (■
400
6
1.5 .■ 10^
, .. 10«
2 . !0-l«
1 . in'
Mnxi
•T jurr;
lOon
200
2 X 10^
1 > 10>«
■ icr ' ■
1 ■ 1 cr
:i<in, fit
-ipheHon
150
0.2
7 ^ lO'*
1 ■: 1 '
't,. 10-1-
,
A., er
'.tie at
iphelirjn
400
2
1.5 ■ 10^
1 ,. 10«
7 . lO-'l
1 . 10^
'■-II<: Ht
)e rih t-li on
400
3
1.5 ■ 10^
2 . 10«
1 . 10-'^'
\:::xi
; ;'l: r , .1
:;»■ r-i hf-Won
1 non
100
2 .. ,0^
Q
5 ■ 10
a
S ■ 10
2 . 10"
'.)] y fcxer] away fron:i Sun.
The electrostatic analyzer aboard Vela III established the jDresence of
3lle+ + , sev^eral charge states of ^^O and possibly ^^C (Banie et al. , 1968). The
3nc/'^He ratio varied from 1.3 X 10-3 to less than 2 X 10-4, and the He/O ratio
from 25 to 80. The two spectra yielding these results were niade possible by
short periods of very low temperatures ( ;10'*K), which permitted peaks to Ijc
resolved.
The density of protons is derived from flux and bulk velocity measure-
ments made aboard spacecraft. Velocities are derived from ener gy-per-unit-
charge measurennents, made with electrostatic analyzers (either Faraday cu]ds
or hemispherical spectrometers). High bulk velocities are correlated with low
densities, implying a relatively constant flux. The average electron density
has been measured directly by means of radio propagation experiments (e. g. ,
Koehler, 1967). These measurements employ the differences in the phase lag
of two, different frequency, rf signals beamed to a distant spacecraft. There
have also been direct mieasurements of electron fluxes and velocity distributions
by electrostatic analyzer. The particle densities are consistent with bulk elec-
trical neutrality (no excess of electrons or protons plus ions).
The characteristics of the solar wind are subject to wide variations as the
ranges of different quantities has shown. The highest velocities and greatest
temperatures are attained during solar maxinnum, which is the period of peak
activity during the Sun's 11 -year cycle of activity. More sunspots appear during
this period, beginning early in the cycle in smaller numbers at high solar lati-
tudes and increasing in number and progressing toward the equator as the cycle
June 15, 1971
F^. Haines, R. Newbum, JPL
Sec. 6, page 13
Cis-Mai'tian Medium, Radiation
JPL 606-1
advances. Sector structure of the solar wind's magnetic field (see Magnetic
Fields) becomes more complicated during peak solar activity, and solar flares
and geomagnetic stornis occur frequently.
Insofar as Mars has no magnetic field (see Magnetic Fields), the solar
wind is deflected by the ionosphere and a bow wave is formed, not dissimilar
to that formed by interaction with the Earth's nnagnetic field, but much closer
to the planet (see Fig. 1). Observations of the Martian bow wave are nonexis-
tent, }:)ut a theoretical study by Spreiter, Summers, and Rizzi (1970), indicates
that it should be less than half a planetary radius above the surface normal to
the flow.
The details of the interaction of the solar wind and the Martian ionosphere
have not been studied, but Cloutier, McElroy, and Michel (1969) indicate the
possibility that the solar wind may provide a source of hydrogen for Mars.
9CVv A^AVE
NOSPHERE
lONOPAUSE
STREAMLINE
Fig. 1. Principal features of solar wind flow past Mars
(after Spreiter, Summers, and Rizzi, 1970).
Sec. 6, page 14
E. Haines, R. Newburn, JPL
June 15, 1971
JPL 606-1 Cis-Martian Medium, Radiation
Solar Cosmic Rays
Solar cosniic rays are solar particles which are distinguished from the
particles of the solar wind by their much greater energy, typically 100 Mev/
nucleon to over 1000 Mev/nucleon as opposed to a maximum of 0.02 Mev for
an Q particle and even less for protons and electrons in the solar wind. They
are generated by various sorts of disturbances in the solar photosphere, chro-
mosphere, and corona (McDonald, 1970). The most energetic and abundant
particles are associated with solar flare events when the particles are expelled
by what appears to be an explosion in the chromosphere. After one solar rota-
tion (Z7 days) the remnants of the flare again release particles, with reduced
energies, in the direction of the Earth. These are termed "recurrent flare"
particles. Even during relatively quiet periods, active centers in the chromo-
sphere release broad beams of energetic particles. These particles sometimes
pass 1 AU in a relatively undisturbed state and are called "active center
related"; alternatively they arrive in the turbulent plasma behind an interplane-
tary shock, in which case they are called "energetic storm particles. " Between
the active centers, the solar cosmic rays diminish in relative intensity, but a
small number of energetic particles are always present at 1 AU. It is not
known whether these represent some form of high energy emission from the
quiet Sun, or the time-averaged echo of many solar events trapped in the micro-
structure of the interplanetary medium.
The most energetic flare -as sociated-particles arrive at 1 AU several
hours after the visual observation of a solar eruption, perhaps 10 to 15 times
later than would l:ic expected for straight-line trajectories. The rise in flux is
rapid and may peak only a few hours after the onset and then slowly diminish
in a few days to the quiet Sun value. Peak fluxes vary from 10^ to 10^ times
the quiet Sun flux. Energy spectra are described by a power law;
dE
KE (1)
where J is the flux of particles (cm"^ sec"^ sterad"!), E is the energy-per-
unit-mass (MeV nucleon" 1) and K and y are parameters which may vary with
flare intensity or with time within a given flare. The exponent y ranges from
2.5 to 3.5 and tends toward the smaller value for the intense periods of flares,
(Smaller values of y make the spectrum flatter, giving more weight to higher
energy particles). Instantaneous fluxes reach as high as 10^ cm-2 sec"!
stcrad~l, but the time-averaged flux (averaged over the solar cycle) is very
much lower, probably on the order of 80 protons cm'^ sec"^ (4 sterad)"-^
(Shedlovsky, et al. , 1970).
The recurrent events represent only a small portion of the initial events,
Flux maxima are nearly symmetrical in tune, rising to approximately 2 to 10
times the quiet background, and lasting only a few hours. They contribute
little to the sustained flux at 1 AU.
June 15, 1971 E, Haines, R, Newburn, JPL Sec. 6, page 15
Cis -Martian Medium, Radiation
JPL 606-1
The particles associated with active centers display broad symmetrical
flux increases, which may rise a few hundred times above the quiet flux, and
last several days. These positively charged particle fluxes are distinguished
from other fluxes by their anticorrelation with energetic electron fluxes.
The energetic storm particles appear in the wake of a plasma shock wave
propagating out from a solar flare. They are believed to be trapped in the tur-
bulent magnetic flux behind the shock. Their fluxes may reach 100 times the
quiet flux level and diminish in a few days.
Solar cosmic rays are believed to represent essentially the solar compo-
sition. The ratio H/He derived from induced radioactivities in lunar materials
is 8:1, substantially in agreement with the photospheric value of 10:1.
The Martian surface dosages for solar flare protons are given in Table 6,
taken from calculations of Foelsche and Wilson (1970). These particles consti-
tute a significant natural radiation source on the Martian surface.
"Galactic" Cosmic Rays
"Galactic" cosmic rays are extremely energetic nuclear particles, of all
atomic numbers, and energetic electrons and gamma rays. Most of the rays
are nuclear, and most of the nuclear particles are protons. Origins of the rays
are unknown, and speculation about their source is quite varied. One theory
suggests that stochastic collisions between particles and fast moving turbulent
gas clouds will, on the average, accelerate the particles. Another theory
Table 6, Maximum number of solar event protons (cosmic rays)
reaching the Martian surface, assuming an atmosphere
consisting of 6 mb of CO^ (Foelsche and Wilson, 1970),
Solar event protons
(Cosmic rays)
Single event
Y e a r ly
Proton energy,
Mev
Fluence (>E),
protons/cm^
Maximum
Fluence (>E),
protons/cm2
Near Sunspot minimum
(1973-1976)
Fluence (--E), protons/cm^
Maximum
Most probable
>10
:20
MOO
:-200
■300
2.2 X lo"^
2.1 X lo"^
1.6 X 10^
1.1 X lo"^
7.2 X 10^
6.4 X lo'^
5.9 X lo"^
4.0 X lo"^
1.8 X 10^
9.6 X 10^
2.4 X 10^
2.2 X 10^
1.7 X lo"^
1.2 X 10^
7.7 X 10^
M.5 X 10^
1.4 / 10^
1.1 X 10^
M.7 X 10^
<5.0 X 10^
Sec. 6, page 16
E. Haines, R. Newburn, JPL
June 15, 1971
TPL 606 1 Cis -Martian Medium, Radiation
suggests that powerful supernova are frequent enough in our galaxy to supply
all the cosmic rays. Other theories include synchrotron acceleration m grow-
ing dipole magnetic fields, and the spiral acceleration of particles released by
pulsars in the pulsar's co-rotating magnetic field. It is almost certam that the
rays are formed and accelerated by a hierarchy of means, because they span
energies which exceed the means of any one mechanism. Disturbances on nor-
mal stars (see the preceding paragraphs on solar cosmic rays) may provide
the abundant, lowest energy particles. A variety of galactic disturbances may
give rise to, and accelerate, the intermediate energy particles. But inter -
galactic means are required to accelerate the high energy particles, whose
magnetic rigidities are too large to be contained by galactic magnetic fields.
(For reviews of cosmic ray physics, see Rossi. 1964, and especially Hayakawa,
1969. )
The flux of energetic galactic cosmic rays is very nearly constant and
isotropic, although there is some modulation by solar activity. The flux of
nuclear particles with energies greater than 100 MeV/nucleon is about 2 cm"^
sec-1 at 1 AU. The Mariner IV cosmic ray telescope indicated a 5% increase
in the integral proton flux and a 30% increase in a-particle flux at 1.5 AU
(O'Gallagher and Simpson, 1967). The electron flux is about 10% of the heavier
particle flux. The lowest energy portion receives episodic contributions from
solar flares whose most energetic components may reach several GeV/nucleon.
These are also distinctly anisotropic (see the preceding section). Other distur-
bances to the continuity of the flux occur as a result of the 11-year solar activ-
ity cycle Early experiments revealed a flux variation at sea level of about 6%;
higher flux at solar minimum and lower flux at solar maximum. These early
results were misleading about the total effect of solar maximum, because only
the most energetic particles were observed at sea level. Above the Earth s
atmosphere, the low energy cosmic ray flux is heavily attenuated during the
solar maximum, while the high energy flux is much less affected. Decreases
in cosmic ray fluxes are also observed in conjunction with magnetic storms.
Magnetic storms are believed to be shock waves from chromospheric explosions
propagating through the plasma medium. These decreases, called Forbush
decreases, affect cosmic particles in a much broader band, than do those
related to solar maxima. These phenomena are thought to reflect the increased
opacity of a turbulent magnetic field. The magnetic fields, related to the micro-
and meso-structure of the solar plasma, could be carried several AU into inter-
planetary space, which would provide a diffuse scattering shield against the low
energy galactic particles. Those particles with higher magnetic rigidity are
less affected. The shock waves apparently carry stronger and more turbulent
fields, thereby providing a more effective shield against those particles with
higher magnetic rigidity.
The differential energy spectrum at solar minimum rises between
100 MeV/nucleon and 1 GeV/nucleon, reaches a maximum between 1 and 2 GeV/
nucleon, and then falls monotonically to the highest energies observed. At
solar maximum it is this low energy part of the spectrum which is depressed,
pushing the maximum in the spectrum out to higher energies, -2-4 GeV/nucleon.
The energy spectrum below 1 GeV/nucleon may be modeled by
— - E^''^ (0.5 + E)"S E < 1 GeV/nucleon (2)
dE
June 15, 197 1 E. Haines, R. Newburn, JPL Sec. 6, page 17
Cis -Martian Medium, Radiation
JPL 606-1
"^"'"V' '' '^'?a[^of K'"''^ sec-l{4TTsterad)-lland F i s e^:nr..sed i n Ge V / nacl eon
(llayakawa. 1Q6Q). lietween 1 and 5 GcV the spectruni is bettor expressed by"
dE ~ ^"•'^' ^ ^-^ ■
1 GeV
iLic 1 eon
5 GeV
nucleon
(3)
,^'V' l''^^'"^^^^ indistinguishable frcn the K-2.5 i^vv above ^ GeV/nurb>on, which
hobls accurately to about 10' GeV. At 10? GeV the spectrum steenens some- "
what. Particles with en<>raies preater than 107 Qe V have reIative]^• shr>rt di[-
lusu,.i lifetimes in the lluctuatina magnetic fields rd the ualaxv IL i ■. ^p.M-ul'ated
tliat the steepening ol the spectrum at 10? GeV represents the loss of more ener-
I'r'" ''^'n^'^W'"^''"' '"' ^'''^ ^"'' favorable capture of inte r ualactic particles.
•^'^*''' ^'' ^'^'^ ^^'"■' spectrum again recovers its E-2.5 .^haoe out to the hi<diest
energy observed, -10^1 GeV (~1020 eV or -ID ir.nloc; f,.,,i, ' •"
, . vi^ V 1 LI cv, Ol lu joules, -i. truly macrosconic
eMT<' r siV : . ' '
Gomposition of the cosmic ray flux does not correspond to the solar abun-
rlance, or to that assumed t(, be the "cosmic" abundance. (For example see
Cameron, 1Q68.) Abundance investigators divide the elements into seve'rai
broad categories: light (L), Z . 3-5; medium (M), Z = 6-9; three categories
"''^^'-^^^y ^"3).. ^^ = 10-15; (FI-,), Z . 16-19; (Hi). Z ~ 20-23. Very heavy (VH)
'■^ "",; '"^^^ '''->'' P^'^M^. X ~ 24-28, and very, very heavy (VVII) to everything
"-^''^■^ '^ - 29. The o/p ratio, 1/8, is about the same in cosmic rays as in the
.solar and other cosmic abundances, while the cosmic rays are much richer in
^^ "'^"^^^^1^^'^' ^^^.- ^"^1 ^^- Tl^° ^t "^^clei, C, N, 0, and F, are qualitatively the
"';"^';, ^^^^ dominance of even-Z over odd-Z elements is already evident here
Lhe H, VII, and VVII elements are more abundant in cosmic rays, by almost
two orders of magnitude in the case of the VII elements. The Up group S CI
Ar, and K, is enriched relative to themther II elements. This subgroup forms'
a mmmuim m the solar and cosmic abundance curves which is filb^d in for cos-
'\'"^ ^'^y'^- ^^^ abundances for these groups, relative to the cosmic ray proton
<ibundanco, are: L, -1 X 10-3; M, 3 x 10-3; n, l y \q-3. ;,nH \TU X /in-4
and VH, 3 / 10'
Most of the ei
■nergy ol cosmic rays which impinges upon planetary matter,
such as Mars atmosphere and surface, is given up in nuclear cascade reactions
^^^ ^''^^/'' primary particles decreases with depth as nuclei are lost through
mternuclear collisions and as lower energy particles come to a stop. But the
overall tlux of particles rises to a nmximum as secondary narticles, protons
"^''^'■""^' ^'^^' niesons are prorjuced by collisions. The .itmosphere of ., olanet
thus absorbs the incident narticles, but ran multi,dy the net particle flu .!
Figure 2 shows char »ed -pa rticle flux at the planet surface versus
atmosplieric mass. The thin atmospher<; on M.a's i .creases the flux of char.-.a
particles expected at the surface, in comparison with that in space With a '"
i -vnh surtace pressure on Mars, the flux a' th - s-r;a(-e incr.v.^es by QO'''
Ilaines (1967) concludes that the fast neutron ihrn arising from ^osmic ray col-
lisions with atmospheric atoms d-es not ^^xceed SO cm-2 sec"!. This fju- is
probably insufficient to produce a measurable radiation field by neutron acti'-, -
tion ol the surtace (causing emission of elecG-ons and Y-rays in resultant beta-
decay orocesst's) but lov. b-v.d r adioac ti\d f y ,<m\ stable isot-aa' <hift^ -.- be
detectable. bhe flux drops off with depth to .1 fi'w -lercei
after traversing two meters of typical's ur face material.
>i its initial v^ilue
Sec. 6, page 18
H;
lines,
H. -Ncwburn, J PI.
June 15, 197 1
JPL 606-1
Cis -Martian Medium, Radiation
E 6
X
Z>
y
I—
<
O
>
X
o
X
o
Z
EARTH
ATMOSPHERIC MASS P , g cm
Fig. 2. Cosmic-ray-induced charged-particle flux at the surface versus
atmospheric mass for Mars and Earth. The dotted lines indicate values
for elevations 10 km above and below a mean elevation at 5.3 mb.
Based upon Haines, 1967.
6.2 MAGNETIC FIELDS
Solar Interplanetary Magnetic Field
The solar wind, being a nearly perfect conductor, carries with it the mag-
netic field lines which were trapped in the chromosphere plasma. These field
lines extend roughly perpendicular from the solar surface, but, as the Sun turns,
the lines which leave the surface radially arc into Archimedian spirals. Given
the average solar wind velocity (along with which the lines propagate), and the
27-day rate of solar rotation, the field lines at 1 AU are bent 45° from the
Sun-Earth direction; they may be positive at 135° and negative at 315°, or vice
versa. Particle velocity fluctuations are characteristically anisotropic about
the bulk velocity. Fluctuations parallel to the magnetic field are greater than
perpendicular fluctuations, leading to a temperature anisotropv of the solar
wind T II = 2Tj. ^'
June 15, 1971
E. Haines; R, J. Mackin, Jr.
R. Newburn, JPL
Sec. 6, page 19
Cis -Martian Medium, Radiation JPL 606-1
The spiral structure of the magnetic field is one manifestation of what is
termed the "macrostructure" of the solar wind. The structure is divided into
positive and negative regions (depending upon whether B is positive toward 135°
or toward 315°, respectively) called sectors. Typically 4 to 6 sectors (always
an even number) exist around the Sun. Smaller scale features in the velocity
and magnetic structure are called "mesostructure" and include the flux tubes
which conduct solar flare particles. Still smaller scale structure, the "micro-
structure" creates magnetic opacity which reduces the flux of interstellar cos-
mic rays. The average solar field strength at the orbit of Mars is thought to be
about Zy, but the instantaneous value may range up to 25 y . The larger field
values - and abrupt changes in direction of field lines - are associated with
hydromagnetic shock waves passing through the solar wind plasma.
Martian Magn etosphere and Magnetic Moment
A planet's magnetic field is expected to be bounded on the sunward side
by currents flowing in the solar wind plasma (q. v. ). The field is thus confined
to a roughly teardrop-shaped cavity known as the magnetosphere. The boundary
of the magnetosphere (magnetopause) is a surface along which, roughly speaking,
the planetary magnetic field pressure (B^/Stt) is balanced by the effective pres-
sure of the solar wind directed flow. "Radiation belts" of energetic electrons
and protons are expected to be contained within the magnetosphere.
The magnetosphere forms an obstacle to the solar wind flow. Because
this flow is "supersonic, " a shock wave (bow shock) is expected to be formed
in the flow ahead of the magnetosphere and to be readily detectable as a surface
of discontinuity in plasma and field parameters.
The Mariner IV instruments gave no indication of any effects associated
with a Martian magnetosphere, its radiation belts (Van Allen et al. , 1965), or
its associated bow shock (Smith et al. , 1965). This fact, interpreted in terms
of appropriately scaled terrestrial magnetospheric parameters, was used to
set a limit on the possible sunward projection of the bow shock (at 0.6 Rj^^ above
the surface) and thus to set an upper limit on the Martian magnetic moment.
Even if Mars has little or no intrinsic magnetic field, it may be expected
to have a bow shock, located about one-third planetary radius sunward from the
planet's surface. The effective "magnetosphere" is created by induced currents
in the Martian ionosphere that "pile up" interplanetary fields carried by the
solar wind flow. (Venus has been found to possess such a bow shock. Bridge
et al. , 1967.)
Surface Magnetic Fields
The Mariner IV data would permit an intrinsic surface (dipole) field as
large as 100 y and a magnetosphere extending to about 0.5 R^ above the surface
on the sunward side (Smith et al. , 1965). A diurnal field variation produced by
currents at the magnetopause would evidently be observable at the surface.
E. Haines; R. J. Mackin, Jr. ;
Sec. 6, page 20 R. Newburn, JPL June 15, 1971
JPL 606-1 Cis-Martian Medium, Radiation
For an intrinsic magnetic field less than about 35 y> the ionospheric
currents would replace the magnetopause with an "ionopause" (Bridge et al. ,
1967) and time-variable fringing fields, of the connpressed interplanetary fields
of about 35 Y. would be expected on the surface near the subsolar point, dropping
to interplanetary values or less past the terminator.
6. 3 thp: meteoroid environment
Most of what is known about the present-day meteoroid flux in the solar
system is derived from terrestrial and the near-1 AU spacecraft observations,
although Mariner 11 carried a "cosmic-dust experiment" to Venus and Mariner IV
carried one to Mars. The observations include eye-witness accounts of mete-
orite falls, photographic records of meteors in flight, radar observations of
ionization trails, photometry of the zodiacal light, detection of space -induced
rcidujactivity in particles from ocean sediments and polar ice, and spacecraft
penetration and acoustical measurements.
A meteoroid is a solid particle in space. A meteor is the electromagnetic
phenomenon produced when a meteoroid penetrates the atmosphere. A bolide is
a sound-producing meteoroid. The term meteorite is reserved for meteoroidal
(objects which have fallen to the Earth. Each meteoroid, no matter how small a
particle, maintains its own orbit about the Sun until it is deflected by a perturb-
ing gravitational field, degraded by radiation drag (Poynting-Robertson effect),
swept away by radiation pressure, or until the particle is destroyed either by
collision with another, or by sputtering or evaporation. New meteoroids con-
stantly replace those being destroyed. The source of meteoroids may be from
any one or a combination of the following: abrasion or collisions of asteroids,
fragmentation of comets, condensation of gaseous atoms and molecules, or the
capture of interstallar dust by the solar system.
Most meteoroids trace direct orbits about the Sun. Those with Earth-
crossing orbits have semimajor axes both larger and smaller than 1 AU, and
most have rather large eccentricities. The larger percentage of meteoroids
overtake the Earth, enter its atmosphere nearer the antapex and are thus more
frequent in the evening. Those with a semimajor axis less than 1 AU are over-
taken l^y the Earth and are slightly more frequent during the morning hours.
Retrograde meteoroids are less common and enter the leading hemisphere with
low frequency and great velocity. Statistics indicate there are two families of
meteoroids, designated sporadic and stream meteoroids. The sporadic mete-
oroids, whose orbits have elements which vary but are similar, appear ran-
domly in the Earth's sky. Stream meteoroids of any one stream have nearly
identical orbital elements. The sources of sporadic meteoroids are still sub-
ject to speculation, while the source of many meteoroid streams is clearly
related to short and long period comets.
Comprehensive reviews of the nature of meteoroids and their flux have
been written by Whipple (1967), Kaiser (1968), Singer (1969), Dohnanyi (1969),
and Whipple et al. (1969). More recent considerations based on spacecraft
experiments have been set forth by Gerloff and Berg (1970) and are concerned
primarily with the terrestrial and lunar environments. Extrapolation of this
E. Haines; R. J. Mackin, Jr.;
June 15, 1971 R. Newburn, JPL Sec. 6, page 21
Cis-Martian Medium, Radiation JPL 606-1
data to Mars is solely dependent upon Mariner IV results for reference. The
Mariner IV results indicated a five-fold increase in dust particle flux at 1.4 AU
(from 7.3 X 10"^ m"^ sec"l to 3.3 X lO"** m~^ sec"^), and then a decrease to
1.8 X lO"'* m"^ sec"-^ at encounter with Mars (Alexander, McCracken, and
Bohn, 1965). These results apply only to dust. Whether they have any rele-
vance to larger sporadic material is uncertain. In general, the meteoroid
streams intersected by Mars' orbit are different from those encountered by
Earth. There is no reason to expect the flux in the Mars intersected streams
to be grossly different from those intersecting the Earth's orbit. Velocities of
meteoroids further from the Sun will be somewhat lower. Sporadic meteoroids
and stream meteoroids obey different flux laws near Earth, These laws are
derived from the photographic, radar, and spacecraft experiments previously
described. Models of these mass-flux laws are presented in Fig. 3. Figure 3
shows the most important experimental points as well as the models developed
by Dohnanyi (1969) and Whipple et al. (1969). Stream meteoroids are seen to
have a much flatter mass -flux function, resulting in their being overwhelmed
by sporadic meteoroids at low mass. The mass cutoff at ~10"12 g represents
the removal of very small particles by radiation and solar wind pressure, and
the decreasing slope between 10"8 g and 10"12 g reflects the reduced lifetimes
of smaller particles due to radiation drag, collision, and sputtering. The
frame of reference is the geocentric frame. The flux presented in Fig. 3 is
that which a massless observer would experience while traveling in a near cir-
cular orbit at 1 AU. The effect of Mars gravity is to focus those meteoroids
having low areocentric velocities, thus increasing the flux somewhat. In addi-
tion, Mariner IV found a factor of 2 to 3 increase in flux for sporadic meteoroids.
An encyclopedic statement, even about geocentric meteoroid velocities,
is innpossible to make at this stage in our experimentation. It is difficult to
determine whether the difference in velocities with mass represents real varia-
tions in meteoroids, or differences in the methods of measurennent. Stream
meteoroid velocities can be rather accurately determined because many mete-
oroids of each stream are available for measurement. The eighteen major
streams encountered by Earth have geocentric velocities ranging from 16 km
sec"^ to 72 kni sec"^ and averaging about 40 km sec"^. The Earth's orbital
velocity plus the maximum rotational component is roughly 30.2 km sec"-'-.
The difference, 42 km sec"-^, is virtually the parabolic velocity (escape velocity)
at Earth's distance from the Sun. This indicates that many meteoroid orbits
are of large eccentricity; in fact, all but one of the eighteen major streams
encountered by Earth have eccentricity greater than 0.75 (Whipple et al. , 1964).
The mean orbital velocity plus the maximum rotational component of
Mars is roughly 24.3 km sec"^ and the parabolic velocity at its mean distance
is 34 km sec"^. Therefore, Mars can be expected to encounter streams with
relative velocities between about 10 km sec-1 and 58 km sec"-^.
The geocentric velocities of sporadic fireballs range typically from near
zero to about 30 km sec"^ and average about 15 km sec'l (McCrosky, 1968).
The fireball orbits tend to display low to moderate eccentricities and perihelia
slightly smaller than 1 AU. Radar observations of much smaller objects show
Sec. 6, page 22 E. Haines, R. Newburn, JPL June 15, 1971
JPL 606-1
Cis -Martian Medium, Radiation
O
O
2
-2
-4
-6
-8
-10
-12
-14
-16 -
-18
1 1 1 1 1 1 1
PIONEER SAND 9
(FOIL-GRID COINCIDENCE & TOF)
MARINER II AND IV (ACOUSTIC)
EXPLORER 16 AND 23
(PENETRATION)
PEGASUS
(PENETRATION)
MODEL
(SHADED;
MODEL 2
(DASHED)
SPORADIC
STREAM
RADAR METEORS
PHOTOGRAPHIC METEORS
-16 -14 -12 -10 -8 -6 -4
MASS (M) LOG^Q, g
Fig. 3. The mass-flux relationship. The flux is that observed
near 1 AU in the absence of gravitational focussing.
(Adapted from Berg and Gerloff, 1970a. )
semimajor axes ranging from 0.4 to >5 AU, with the number of objects increas-
ing with increasing eccentricity, and having more or less random inclination.
Further, geocentric velocities range from very low to extremely high (-'TO km
sec"-'^) with an average around 40 to 50 km sec"-'-. Most of the very small
objects measured by Pioneers 8 and 9 (Berg and Gerloff, 1970a and 1970b) had
low geocentric velocities and semimajor axes <1 AU. (One was retrograde,
semimajor axis less than 1 AU; one was direct, semimajor axis approximately
3 AU. ) This observation is interpreted by the authors to reflect radiation drag
on the very small particles, reducing their large, highly eccentric orbits to
small, less eccentric orbits.
Barring an encounter with an extra-solar system body (a hyperbolic mete-
oroid*), the velocities of sporadic meteoroids near Mars must also be in the
range of 10 to 58 km sec"^. However, the distribution of sporadic meteoroids
within that range is unknown.
*The existence of hyperbolic meteoroids is speculative and not comnnonly
accepted by all experts.
June 15, 1971
E. Haines, R. Newburn, JPL
Sec. 6, page 23
Cis -Martian Medium, Radiation JPL 606-1
Bulk density of photographic objects nriay be calculated from drag theory
and the photometric mass equation (McCrosky, 1968). No photographic object
has yielded a bulk density greater than 1,2 g cm"-^, with the average about 0.4 g/
cm"3. However, Pribram meteorite, whose photographic density was small,
had a measured bulk density of 3.5g cm"3. The uncertainty remains as to
whether our understanding of drag and ablation is in error, or whether that part
of Pribram which reached the ground represented only a small, dense fraction
of the whole meteoroid. The latter hypothesis is widely accepted, and the
density of most nneteoroidal material is considered to be < 1 g cm-3.
Sec. 6, page Z4 E. Haines, R. Newburn, JPL June 15, 1971
JPL 606-1 Cis-Martian Medium, Radiation
BIBLIOGRAPHY
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Brandt, J, C, 1970, Introduction to the solar wind: Freeman and Co, , San
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Bridge, H.S., Lazarus, A, J, , Snyder, C, W, , Smith, E, J. , Davis, L, , Jr. ,
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Dohnanyi, J, S. , 1969, On the origin and distribution of meteoroids: Bellcomm,
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Hames, E, L, , 1967, Estimate of the natural radiation on the surface of Mars:
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Kaiser, T,R., 1968, The interplanetary dust cloud, p. 323-342 in Physics and
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York, Springer Verlag,
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Koehler, R. L. , 1967, Interplanetary electron content measured between Earth
and Pioneer VI and VII spacecraft using radio propagation effects:
Stanford Electronics Laboratory, SU-SEL-67-051.
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^•■U.S GOVERNMENT PRINTING OFFICE l,i7i 7 i^-f-hH / |( i i 7 1- i
June 15, 1971 K. Haines, R. Newburn, JPL Sec. 6, page 27
JPL 606-1
Document Control List
DOCUMENT CONTROL LIST
These pages, together with the dividers and the two Mars maps (I, P.P.
Mars 1969 and MM'71 Mars Planning Chart in the pocket on the inside-front
cover), form the total content of the document as revised, publishing date
March 1, 1972.
Title page (i) 3-1-72
Frontispiece Caption (ii) 3-1-72
Frontispiece (iii/iv) 7-15-68
V thru ix 3-1-72
SECTION 1
i and ii 3-1-72
1 thru 32 10-15-71
Appendix A, 1 thru 11 4.1_67
Appendix B, 1 thru 9 10-15-71
SECTION 2
i a"fl ii 3-1-72
1 thru 21 11-15-71
22 thru 25 3-1-72
26and27 11-15-71
SECTION 3
i t'^ru x 3-1-72
3. 1
1 thru 25
2-15-72
3.2
i thru 24 10-1-71
Appendix A, 1 thru 19 10-1-71
Appendix B, 1 thru 4 10-1-71
3. 3
1 thru 26 11-15-71
Appendix, 1 thru 5 12-1-71
3.4
' tf^^" 32 12-1-71
3. 5
1 thru 61 1-1-72
"^9 7-1-68
3.6
i thru 93 6.1.71
SECTION 4
i thru iii 3-1-72
1 12-15-71
4. 1
i thru 11 12-15-71
12 2-15-72
13 thru 18 12-15-71
4.2
1 thru 45 2-1-72
Appendix, 1 2-1-72
=i'19 (+2 overlays and 1 color map) 4-1-67
*21 (+3 overlays and 1 color map) 4-1-67
*23 (+2 overlays and 1 color map) 4-1-67
*25 ( + 1 overlay) 4-i_67
SECTION 5
i thru iii 3-1-72
*5
4-1-67
5. 1
1 thru 15 4-15-71
Appendix, 1 thru 3 4-15-71
5. 2
1 thru 11 7-30-71
5.3
i 3-1-72
*1 thru 20 9-11-67
21
3-1-72
5.4
' 3-1-72
*1 thru 18 7-3-67
19 thru 21 3.1.72
SECTION 6
' 3-1-72
1 thru 27 6-15-71
APPENDIX
Document Control List 3.
1-72
* Denotes page numbers of material retained from original issue (July I968)
March 1, 1972
Appendix, page 1
CHANGE NOTICE
MARS SCIENTIFIC MODEL REVISION
The attached material replaces and/or supplements the information
contained in the Mars Scientific Model, Document No. JPL 606-1, published
July 15, 1968, The revised material reflects data obtained and/or derived
from Mariners 6 and 7, and other sources available through March, 1972.
INSTRUCTIONS FOR UPDATING
MAPS: Place the two new Mars maps (IPP Mars 19d9 and MM' 71 Mars Planning
Chart) into the map pocket located inside the front cover. Discard miaps
MEC-1 and MEC-2.
FRONT MATTER: Replace entirely with new issue of Front Matter except for
color Frontispiece (to remain as page iii),
SECTIONS 1 and 2: Replace entirely with new issues of Sections.
SECTION 3: Replace Contents pages with new issue of Contents.
SUBSECTIONS 3.1, 3.2, 3,3, 3.4: Replace entirely with new issues of Subsections.
SUBSECTION 3. 5: Replace entirely with new issue of Subsection except: Retain
p. 29-'' (Fig. 24) and insert as last page of Subsection.
SUBSECTION 3.6: Insert new Tab and new Subsection 3.6.
SECTION 4: Replace Contents pages with new issue of Contents.
SUBSECTION 4. 1: Replace entirely by new issue of Subsection.
SUBSECTION 4.2: Replace entirely with new issue of Subsection except: Retain
pp. 19''', 2 1'^% 23''- with related color maps and overlays, and p. 25"'-
(replacing overlay of Place-Names with new one) and insert as last pages
of Subsection.
SECTION 5: Replace Contents pp. 1 thru 3 with new issue of Contents (i thru iii),
and retain p. 5-'= as the last page.
SUBSECTIONS 5. 1 and 5.2: Replace entirely with nev/ issues of Subsections.
SUBSECTION 5,3: Retain entire Subsection except: Replace pp. 2 F'= and 23''=
with new p. 2 1 and insert new p. i in front of Subsection.
SUBSECTION 5.4: Retain entire Subsection except: Replace pp. 19* thru 22='=
with new pp. 19 thru 2 1 and insert new p. i in front of Subsection.
SECTION 6; Replace entirely with new issue of Section.
APPENDIX: Replace with new Appendix (Document Control List).
Verify contents of total updated Mars Scientific Model Document
by use of this new Documient Control List.
Denotes page numbers of material retained from original (1968) issue of
Mars Scientific Model Documient.
Change Notice, page 1
Live Music Archive
Librivox Free Audio
Metropolitan Museum
Cleveland Museum of Art
Internet Arcade
Console Living Room
Books to Borrow
Open Library
TV News
Understanding 9/11