Full text of "Gamma-ray astrophysics; a symposium held at NASA Goddard Space Flight Center, April 30 to May 2, 1973"

See other formats

NASA SP-339

GAMMA-RAY ASTROPHYSICS

A symposium held at

GODDARD SPACE-FLIGHT CENTER

Green belt, !Maryland

April 30-May 2, 1973

u. s. ^•

NATIONAL AERONAUTICS AND SPACE ADMINISTRATION

NASA SP-339

GAMMA-RAY ASTROPHYSICS

A symposium held at NASA Goddard Space Flight Center

April 30 to May 2, 1973, sponsored by the

National Aeronautics and Space Administration and the

American Physical Society

Edited by

Floyd W. Stecker and Jacob I. Trombka
Goddard Space Flight Center

Prepared by NASA Goddard Space Flight Center

Scientific and Technical Information Office 1973

NATIONAL AERONAUTICS AND SPACE ADMINISTRATION

Washington, D.C.

The requirement for the use of the International System of Units
(SI) has been waived for this document under the authority of
NPD 2220.4, paragraph 5.d.

For sale by the Superintendent of Documents,

U. S. Government Printing Office, Washington, D. C. 20402

Library of Congress Catalog Card No. 73-600345

FOREWORD

Opening Remarks

Dr. John F. Clark, Director

Goddard Space Flight Center

Significant advances have been made recently both in experimental and
theoretical investigations in 7-ray astrophysics. It is thus most appropriate
at this time to gather together to discuss the meaning of these results.
This, then, is the first of what we hope will be many fruitful international
symposia devoted exclusively to this subject.

Because of the controversial and still unsettled nature of some aspects of
this young and exciting subject, it is fitting that we will devote fully one-
half of our time to free and open discussion. We also will try to utilize the
observational and theoretical data presented at this conference to help guide
us in charting our future investigative efforts in 7-ray astrophysics. To this
end, we will culminate the Symposium with a panel discussion on the future
of this field.

The three days of the Symposium will be spent considering the observational
data from about 0.3 MeV to a few hundred GeV and theoretical models of
production mechanisms that may give rise to both galactic and extragalactic
7-rays. We also hope to measure a large interaction cross section between
the theorists and experimentalists gathered here.

We feel that since Goddard Space Flight Center has been heavily involved in
both the theoretical and experimental aspects of 7-ray astrophysics, it is
fitting that such a symposium be held here. We thank the Division of
Cosmic Physics of the American Physical Society for cosponsoring this
meeting.

We further would like to thank the distinguished members of the interna-
tional scientific community who have taken the time to come here and
actively participate in this Symposium.

April 11, 1973

PREFACE

The first international symposium and workshop devoted to gamma-ray
astrophysics was held at Goddard Space Flight Center, Greenbelt, Maryland
April 30 to May 2, 1973. The Symposium was cosponsored by NASA and the
Division of Cosmic Physics of the American Physical Society. Significant
advances have been made recently both in theoretical and experimental inves-
tigations in the field so that 7-ray astrophysics is coming into its own as a
separate branch of astrophysics. This led Prof. Kenneth Greisen, of Cornell,
who was one of the session chairmen, to make the remark that, this Symposium
marks a "birthday of 7-ray astronomy."

Our philosophy in organizing the Symposium was to devote equal time to both
theory and observation. To this end, the organizational work was shared by a
theoretician and an experimentalist.

The Symposium was divided into morning sessions of invited papers which
surveyed all aspects of present work on 7-ray astrophysics and related X-ray
astrophysics, and afternoon workshop-discussion sessions where brief reports
were contributed and discussions of controversial aspects of the field were held.
The final afternoon session consisted of a review of the Symposium (contained
here in the introduction) and a panel discussion on future directions for
research in the field.

The formal program for the Symposium was as follows:

Monday, April 30, Morning

Chairman: Dr. George F. Pieper

Goddard Space Flight Center

Dr. John F. Clark, Director,
Goddard Space Flight Center
Welcome

Lawrence E. Peterson, University of California at San Diego,
and Jacob I. Trombka, Goddard Space Flight Center, on
Low-Energy Gamma-Ray Observations (with emphasis on
results from Apollo-15, -16, and -17 and discussion of
induced activity in crystal detectors)

W GAMMA-RA Y ASTR OPHYSICS

Floyd W. Stecker, Goddard Space Flight Center, on
Mechanisms for Production of Diffuse Gamma-Ray
Continuum Radiation

Donald D. Clayton, Rice University, on

Prospects for Nuclear-Gamma-Ray-Line Astronomy

Monday, April 30, Afternoon

Chairman: Dr. James I. Vette

Goddard Space Flight Center

David J. Forrest, University of New Hampshire, on
Observations of Gamma-Ray Emission in Solar Flares

Reuven Ramaty, Goddard Space Flight Center, on
Theory of Gamma-Ray Emission in Solar Flares

Vitaly L. Ginzburg, Academy of Sciences USSR, P.N.
Lebedev Physical Institute, Moscow (remarks authorized
by Prof. Ginzburg which were presented in his absence), on
Gamma-Ray Astronomy and Cosmic-Ray Origin Theory

Workshop Session Discussion of:

1. Experimental techniques and errors involved in
7-ray measurements (spallation, and so forth)

2. Gamma-ray production mechanisms and theoretical
production rates

3. Gamma-ray astronomy and cosmic-ray origin theory

4. Topics related to morning session

Tuesday, May 1 , Morning

Chairman: Dr. Maurice M. Shapiro

Naval Research Laboratory

William Kraushaar, University of Wisconsin, on
Diffuse Soft X-Ray Observations

Donald Kniffen, Goddard Space Flight Center and
Gerald Share, Naval Research Laboratory, on
10-100 MeV Gamma-Ray Observations

Daniel Schwartz, American Science and Engineering, on
Diffuse 1 keV - 1 MeV X-Ray Observations

Ramanath Cowsik, University of California at Berkeley, on
Theory of the Diffuse X-Ray Background

INTRODUCTION vii

Tuesday, May 1 , Afternoon

Chairman: Dr. George W. Clark

Massachusetts Institute of Technology

Workshop Session Discussion of:

1 . Theory of 7-ray sources

2. Interpretation of SAS-2 results and related
experimental results

3. Cosmological implications of 7-ray
measurements

4. Solar and galactic 7-ray line emission

Wednesday, May 2, Morning

Chairman: Dr. Kenneth Greisen
Cornell University

Roland Omnes, Laboratory of Theoretical and High-
Energy Physics, Orsay, France, and Evry Schatzman
and Jean-Loup Puget, Paris Observatory, France, on
Baryon-Symmetric Cosmology and Gamma-Ray
Observations

Gary Steigman, Yale University, on
Antimatter in the Universe?

Wednesday, May 2, Morning

Giovanni G. Fazio, Harvard Observatory and
Smithsonian Astrophysical Observatory, on
Observations of Ultra-high Energy Gamma Rays

Arnold Wolfendale, University of Durham, England, on
Theory of Ultra-high Energy Gamma Rays

Wednesday, May 2, Afternoon

Floyd W. Stecker

Goddard Space Flight Center

Concluding Remarks, Summary

SPECIAL SESSION: Panel Discussion on Future Directions
in Gamma-Ray Astronomy

Jacob Trombka, Chairman, Goddard Space Flight Center

Evry Schatzman, Paris Observatory

Giovanni Fazio, Harvard and Smithsonian Astrophysical Observatories

viii GAMMA-RA Y ASTR OPHYSICS

Carl Fichtel, Goddard Space Flight Center

Albert Metzger, Jet Propulsion Laboratory

Kenneth Greisen, Cornell University

Glenn Frye, Case Institute

Floyd W. Stecker, Goddard Space Flight Center

Important new results on the diffuse 7-ray background as obtained by Apollo
were presented by L. Peterson (University of California at San Diego) and
J. Trombka (NASA/GSFC) and results obtained by SAS-2 were presented by
D. Kniffen (NASA/GSFC), who also reported observations of the galactic
plane. The results from SAS-2 confirm some important qualitative results
first obtained by OSO-3 that the galaxy is an intense source of 7-radiation
above 100 MeV and it stands out above the extragalactic background in this
energy range. The spectrum is harder above 100 MeV than the 7-radiation
seen at high galactic latitudes, which is presumably extragalactic. The SAS-2
results also indicate that the extragalactic (high-galactic latitude) background
spectrum is quite steep above 40 MeV (roughly ^E "3).

Results from balloon observations by groups at the Max Planck Institute and
the U. S. Naval Research Laboratory, reported by G. Share (NRL), are con-
sistent with the Apollo and SAS-2 results, which present a continuous obser-
vational spectrum from 300 keV up to 135 MeV. These data suggest a bulge
in the 7-ray spectrum above 1 MeV, in spite of background corrections which
are of most importance below 4 MeV as discussed by J. Fishman and C. Dyer.
This bulge has been interpreted as a new component of 7-radiation at energies
above 1 MeV. This argument is even more important if the X-rays below 1 MeV
are thermally produced and are falling off exponentially in energy above
100 keV as was suggested by D. Schwartz (American Science and Engineering)
and R. Cowsik (University of California at Berkeley). Problems with the
thermal interpretation were discussed by W. Kraushaar (University of
Wisconsin).

The interpretation of the 1 MeV to 100 MeV bulge in the 7-ray spectra was
discussed by F. Stecker (NASA/GSFC) who gave a review on 7-ray production
mechanisms. He concluded that the excess is most likely caused by matter-
antimatter annihilation. The Apollo and SAS-2 observational data present an
excellent fit to the predicted annihilation spectrum up to 135 MeV. The
matter-antimatter-symmetric cosmology was discussed by R. Omnes (Labora-
tory of Theoretical and High-Energy Physics, Orsay) and E. Schatzman and
J. Puget (Paris Observatory, Meudon).

The exciting aspects of the matter-antimatter cosmology reported on by
R. Omnes, E. Schatzman, J. Puget, and F. Stecker indicate that, in addition
to implying baryon symmetry on a universal scale, it can explain such diverse
phenomena as: the cosmic 7-ray background spectrum; the ratio of photons

INTRODUCTION ix

to nucleons in the universe of ~109; annihilation as the energy source for
generation of large-scale turbulence leading to galaxy formation; and the
consequent observed sizes, mean densities, and rotational velocities of
galaxies. G. Steigman (Yale) discussed the observational restrictions on
matter-antimatter cosmological models.

The galactic 7-ray flux in the 100-MeV range seen by SAS-2 and OSO-3
indicates an increase in the direction of the galactic center. The most likely
implication is that there is a cosmic-ray gradient toward the galactic center,
as was pointed out in remarks by A. Wolfendale (University of Durham) and
in a communication by V. Ginzburg (Lebedev Institute, Moscow).

Measurements reported on the Crab Nebula and Pulsar by various groups were
discussed by J. Share and talked about by K. Greisen, G. Fazio, and G. Frye.
They indicate that the 7-ray spectrum from the Crab Nebula goes all the way
up to the highest energies yet observed and that there are time variations at
about 1012 eV; this also tells us something about the magnetic field strength
in the Crab Nebula. At ultra-high energies, A. Wolfendale (Durham) discussed
the possibility of observing a flux of 7-radiation in the 1019- and 1020-eV
energy range.

If cosmic rays in this energy range are universal and the cascading process
which he discussed occurs, then we may very well be able to observe the
resultant 7-rays. The situation is a little more pessimistic if there is an extra-
galactic magnetic field of average strength above 10"10 G, because synchrotron
losses would then cut off the cascade process. So, by looking for these 7-rays
in air showers, we may be able to learn something about ultra-high energy
cosmic rays of metagalactic origin.

Wolfendale also discussed joint work with A. Strong and J. Wdowczyk on the
possible electromagnetic cascading at lower energies in the early big bang
universe to possibly explain the 7-ray background in the 1- to 100-MeV range
or, alternatively, to use the 7-ray data to rule out cosmic-ray production at
early epochs on the scale suggested by Hillas. This model requires a rather
low intergalactic gas density at present of ~10"9 atoms/cm3.

There was much discussion of the 470-keV feature, which has been observed
by R. Haymes' group (Rice) in the galactic center region. Three very interesting
theoretical explanations of the 470-keV feature were presented. D. Clayton
(Rice University), suggested that it may be caused by lithium. R. Ramaty
(NASA) suggested that this feature could be attributed to red-shifted positron
annihilation produced at the surface of neutron stars. M. Leventhal gave a
very interesting interpretation that this feature may be due to positronium,
and that the positronium spectrum has been altered by the finite energy reso-
lution of the detector, so that the edge at 51 1 keV appears as a bump at
-470 keV.

x GAMMA-RA Y ASTR OPHYSICS

D. Clayton gave a review of astrophysical processes which should be important
for the production of 7-ray line spectra. D. Forrest (New Hampshire) dis-
cussed the observations of 7-ray line emission in solar flares and R. Ramaty
discussed the theoretical interpretation of these observations. In order to
observe discrete line emission experimentally with 7-ray detectors, higher-
energy resolutions are required. The use of solid-state detectors capable of
such high -energy resolutions aboard satellites was considered by G. Nakano
and W. Imhoff (Lockheed). Experimental results were also presented.
A. Metzger (JPL) described the experimental detectors being planned for
flight aboard HEAO utilizing solid-state techniques.

In planning the Symposium program, we purposely mixed the theoretical and
observational papers in order to maximize the interactions between theoretical
and observational workers in the field. However, for a more logical organization
of the Symposium proceedings, we have divided this book into four sections,
one on observations, one on theoretical papers, one on cosmological implica-
tions, and the last section consisting of an edited transcript of the panel dis-
cussion on the future of the field.

An examination of the transcripts of the discussion indicated that heavy
editing was required in order to make sense out of some of the discussion.
In addition, many of the speakers incorporated points made in the discussion
into their final manuscripts. Thus, discussion material was eliminated which
was deemed to be either incoherent or redundant after speakers were given a
chance to revise any unclear material. The remaining discussion material is
appended to pertinent chapters.

When this Symposium was organized, it was planned to cover all aspects of
the field, so that the proceedings would be a comprehensive up-to-date
reference. However, shortly after the Symposium was held, the exciting
discovery of cosmic 7-ray bursts was reported by the Los Alamos group in
the Astrophysical Journal. The editors therefore felt it important to attempt
to include some discussion of this topic as a special addition to the proceedings.
We have therefore added two special short papers on this subject prepared by
people at Goddard Space Flight Center because of time limitations. These
are an observational paper by Cline et al., on energy spectra of cosmic 7-ray
bursts and a theoretical paper by Stecker and Frost on the stellar superflare
origin hypothesis of these bursts. We realize that other and important work
on this topic such as that by S. Colgate and the Los Alamos group should also
be included in any well-rounded discussion, but, unfortunately this was not
possible here because of our publication schedule.

F. W. Stecker

J. I. Trombka

Goddard Space Flight Center

August 1973

ACKNOWLEDGMENTS

We would like first of all to thank the many authors who have contributed
solid, thoughtful manuscripts to these proceedings and who gave excellent
presentations at the Symposium. We particularly appreciate the time and
effort made by the authors to prepare this material under a severely short
time schedule so that these proceedings would be available to the astro-
physics community in a time comparable to that for publication of a
journal article.

We would also like to thank Dr. Maurice M. Shapiro and Dr. Frank B.
McDonald for coordinating this Symposium with the Division of Cosmic
Physics of the American Physical Society. We would like to thank Dr. John
F. Clark, Dr. George F. Pieper and Dr. James I. Vette, Dr. Theodore G.
Northrop and Dr. Aaron Temkin for their support which enabled us to hold
the Symposium at the Goddard Space Flight Center.

We also thank the attendees, many of whom came from great distances,
for their contributions in making the Symposium a success.

Most special thanks go to Mrs. Sandra J. Walter, for her untiring work in
handling almost all of the administrative details which are so numerous in
an undertaking of this type, and to Barbara Welsh and Elizabeth R. Miller
for help in preparation of these proceedings.

FWS
JIT

CONTENTS

Page

FOREWORD iii

PREFACE v

ACKNOWLEDGMENTS xi

SECTION 1 -OBSERVATIONAL DATA 1

Chapter I 3

A. Diffuse Cosmic X-Rays Below 1 keV

William L. Kraushaar 3

Chapter II 15

A. The X-Ray Emissivity of the Universe: 2 to 200 keV

Daniel Schwartz and Herbert Gursky 15

B. Atmospheric Corrections to Balloon X-Ray Observations

H. Horstman 37

Chapter III 41

A. The Measurement and Interpretation of the Cosmic
Gamma-Ray Spectrum Between 0.3 and 27 MeV as
Obtained During the Apollo Mission

L. E. Peterson, J. I. Trombka, A. E. Metzger, J. R. Arnold,

J. I. Matteson, and R. C Reedy 41

B. Induced Radioactivity Contributions to Diffuse
Gamma-Ray Measurements

G. J. Fishman 61

C. Preliminary Results from the First Satellite of a
High-Resolution Germanium Gamma-Ray Spectrometer:
Description of Instrument, Some Activation Lines
Encountered, and Studies of the Diffuse Spectra

G. H. Nakano, W. L. Imhof, J. B. 'Reagan, and

R. G. Johnson 71

xiv GAMMA-RAY ASTROPHYSICS

Page

D. Preliminary Results from the First Satellite of a
High-Resolution Germanium Gamma- Ray Spectrometer:
Backgrounds from Electron Bremsstrahlung and from
Electron-Positron Annihilation

W. L. Imhof, G. H. Nakano, R. G. Johnson, and

J. B. Reagan 77

E. Further Considerations of Spallation Effects

Give Dyer 83

F. HEAO Gamma- Ray Astronomy Experiments

A. Metzger 97

Chapter IV . . . 103

A. Recent Observations of Cosmic Gamma-Rays from
10 MeV to 1 GeV

Gerald H. Share 103

B. Report on Gamma-Ray Astronomy Results Obtained
in Europe Since the IAU Symposium No. 55

K.Pinkau 133

C. Preliminary Results on SAS-2 Observations of
> 30 MeV Gamma Radiation

D. A. Kniffen, C. E. Fichtel, andR C Hartman .... 139

Chapter V 153

A. Observations of High-Energy Gamma Rays

G. G. Fazio 153

Chapter VI 165

A. Observations of Gamma- Ray Emission in Solar Flares
D. J. Forrest, E. L. Chupp, A. N. Suri, and
CReppin 165

Chapter VII 175

A. Energy Spectra of Cosmic Gamma-Ray Bursts
T. L. Cline, U. D. Desai, R W. Klebesadel, and
I. B. Strong 175

CONTENTS xv

Page

SECTION 2-THEORY 183

Chapter VIII 185

A. The Astrophysics of the Diffuse Background of
X-Rays and Gamma-Rays
Ramanath Cowsik 185

Chapter IX 211

A Mechanisms for Production of the Diffuse Gamma-Ray
Continuum Radiation
F. W. Stecker 211

Chapter X 249

A. Gamma- Ray Astronomy and Cosmic- Ray Origin Theory

V. L Ginzburg 249

B. Galactic Gamma Rays: Models Involving Variable
Cosmic- Ray Density

A W. Strong, J. Wdowczyk, and A. W. Wolfendale . ... 259

Chapter XI 263

A. Prospects for Nuclear-Gamma- Ray Astronomy

Donald D. Gay ton 263

B. Positronium Formation Red Shift of the 51 1-keV
Annihilation Line

M Leventhal 291

C. Nuclear Gamma-Rays from Solar Flares

R. Ramaty 297

Chapter XII 315

A. Ultra-High Energy Gamma Rays

A W. Strong, J. Wdowczyk, and A. W. Wolfendale. ... 315

Chapter XIII 329

A. A Comparison of the Recently Observed Soft Gamma-Ray
Bursts with Solar Bursts and the Stellar Superflare Hypothesis
F. W. Stecker and K. J. Frost 329

xvi GAMMA-RA Y ASTR OPHYSICS

Page

SECTION 3-COSMOLOGY 333

Chapter XIV 335

A. Matter- Antimatter Cosmology

R Omnes 335

B. The Deuterium Puzzle in the Symmetric Universe

B. Leroy, J. P. Nicolle, and E. Schatzman 351

C. Antimatter in the Universe?

Gary Steigman 361

Chapter XV 367

A. Gamma-Ray Background Spectrum and Annihilation
Rate in the Baryon-Symmetric Big-Bang Cosmology

/. L. Puget 367

B. Distortion of the Microwave Blackbody Background
Radiation Implied by the Baryon-Symmetric Cosmology
of Omnes and the Galaxy Formation Theory of
Stecker and Puget

F. W. Stecker and J. L Puget 381

SECTION 4-FUTURE DIRECTIONS IN GAMMA-RAY

ASTRONOMY 385

Chapter XVI

A. A Panel Discussion on the Future Direction of
Gamma- Ray Astronomy

Giovanni Fazio, Carl Fichtel, Glenn Fry e, Kenneth
Greisen, Albert Metzger, Evry Schatzman, Floyd
Stecker, and Jacob Trombka 387

INDEX- LIST OF AUTHORS 407

SECTION 1
OBSERVATIONAL DATA

Chapter I

A. DIFFUSE COSMIC X-RAYS BELOW 1 IceV

William L. Kraushaar*

University of Wisconsin

INTRODUCTION

The study of diffuse X-rays in the energy region below 1 keV has had a some-
what rocky past and has suffered from having attracted cosmological interest
early in its young life. Much of the available data and interpretation can be
found in recent review articles by Silk (1973, preprint), Felten (1972),
Field (1972), and Kato (1972). In this short review I cannot discuss all
the measurements or all the ideas that have been put forward. I will,
therefore, restrict my discussion to a description of those features of the
low-energy diffuse flux on which there is general observational agreement
and to some interpretive matters that I believe have been overlooked or at
least underemphasized. Also, most of the discussion will be restricted to
the energy region below 280 eV, the Carbon-K edge.

INTENSITY

The soft X-ray diffuse intensity is everywhere convincingly larger than would
be expected from an extrapolation of the high energy isotropic, unabsorbed,
and almost certainly, extragalactic power law spectrum. Data in support of
this conclusion are shown in Figures I.A-1 and I.A-2, taken from papers by the
Wisconsin (Bunner et al., 1971) and NRL (Davidsen et al., 1972) groups. The
solid curves in both figures are the predicted proportional counter response
given only the high-energy power law spectrum, with no interstellar absorption,
as an input spectrum. The prominent bumps in these curves result from the
X-ray transmission edges of the counter windows. The intensity ratio from
pole to plane is about 3 to 1, and, while there can be some argument about
a possible extragalactic contribution to the high latitude intensity, the plane
intensity must be of relatively local origin because the column density for unit
optical depth is only 2.5 X 1020 atoms/ cm2 or about 200 pc (with n =
0.4 atoms/ cm3) in these directions.

* Speaker.

OBSER VA TIONAL DA TA

PULSE HEIGHT IN KEV.

Figure I.A-1. Proportional counter pulse-height spectra near the galactic
plane and at a high galactic latitude (Bunner et al., 1971 ).

SPATIAL STRUCTURE

The soft X-ray intensity shows three broad classes of spatial structure.

First, there is the gross tendency for the intensity to be small in the galactic plane
and enhanced by perhaps a factor of 3 at high northern galactic latitudes. This is
shown in Figures I.A-3 and I.A4, surveys of the NRL (Davidsen et al., 1972) and
Wisconsin (Bunner et al., 1972; Williamson, F. W., 1973; Sanders, W., 1973)
groups. The polar enhancement is more obvious in the north than in the south,
although there are some isolated line scans that make the case for apparent
enhancement in the south more convincing (Bunner et al., 1969, 1971 ; Garmire
and Riegler, 1972).

DIFFUSE COSMIC X-RA YS BELOW I keV

POLE
FLON WINDOW

PLANE
TEFLON WINDOW

t+/i

HI \

ENERGY ( KEV )

Figure I.A-2. Pulse-height data taken with Kimfol and Teflon counter
windows (Davidsen etal., 1972).

Secondly, the soft X-ray intensity is by no means just a simple function of galac-
tic latitude nor is it correlated, except in the grossest sense, with the column
density of interstellar hydrogen gas. There are large high intensity spatial
features. None of these features except the North Polar Spur appear to correlate
well with other astrophysical phenomena. Figure I.A-5, taken from part of the
Wisconsin survey (Bunner et al., 1972) shows soft X-ray counting rate versus
time along the scan path plotted together with estimated expected transmission.
The bands on the time axes coincide with the North Polar Radio Spur and
approximately, it is seen, with regions of enhanced X-ray intensity. Notice that
there is little if any detailed correlation of X-ray intensity with gas transmission.
This, together with the observed large intensity in the galactic plane, is strong
evidence that much of the soft X-ray emission originates within the bounds of
the galaxy's interstellar gas.

Thirdly, there are at least three soft X-ray emitting regions of small angular
extent: Puppis-A, Vela-X (Palmieri et al., 1971 ; Grader et al., 1970) and the
Cygnus Loop (Seward et al., 1971). Three others have been reported but to

OBSER VA TIONAL DA TA

-ar

— *2I0!" - 180*- «0*— t-120* 90* -W— "-SO*^— \V

44A x-wnr wtensity

Ct*/MC

■

50<I<70

□

30<I<

□

I0<I<

-60*

-80*

Figure I.A-3. Spatial distribution of X-ray of E < 280 eV. The coordinate
system is centered at the galactic anti-center (Davidsen et al., 1972).

date have not been confirmed. The three confirmed sources are all supernova
remnants, are at small galactic latitudes, and are of a class not numerous enough
to account for the entire diffuse background. Of course, one or a few nearby
remnants of large angular extent would confuse our whole picture. But galactic
loop structures, aside from the North Polar Spur, do not appear to be strong
soft X-ray emitters. Incidently, the observation of soft X-ray emission from
near the North Polar Spur has not been confirmed by others. Only one other
observation near the Spur has been reported, but the sensitivity level is not clear
(Hayakawa et al., 1972, preprint).

NATURE OF THE LOCAL EMISSION

The nature of the local emission remains a mystery. Particularly puzzling is the
relative constancy of the intensity in the galactic plane. Near Cu = 240°; for
example, OAO-Lyman-a observations (Savage and Jenkins, 1972) show there are
very small gas column densities out to several hundred parsecs. Similarly, the
21-cm emission profiles in this region show little or no low-velocity gas. Yet the
soft X-ray intensity near Cu = 240° appears featureless. If the emission in the
plane were from a more-or-less uniformly-distributed population of stars, the
soft X-ray intensity, one would think, would be large where the local absorbing
gas density is small. Early type stars, it is true, are relatively rare in this region.

Also puzzling is the relation between the soft X-ray intensities measured in the
E < 180 eV (Boron-K edge filter) and E < 280 eV (Carbon-K edge filter) regions
(Bunner et al., 1973). X-rays of E < 180 eV are more strongly attenuated by
absorbing material. Thus in Figure I.A-6 is shown the rates in the two types of

DIFFUSE COSMIC X-RA YS BELOW 1 keV

Figure I.A-4. Spatial distribution of X-rays of E <280eV. The upper coordinate
system is centered at the galactic center, while the lower coordinate system is
centered at the galactic anti-center (Bunner et al., 1972; Williamson, 1973 and
Sanders, 1973).

detectors measured while the detectors were holding on a fixed high-latitude point
as the rocket emerged from the Earth's atmosphere. As expected, the rates are not
proportional to each other, but the Boron filter rate changes more rapidly than
the Carbon-K filter rate. Yet when these two detectors scanned about the sky
while free of atmospheric absorption, the two rates showed no systematic tend-
ency that would suggest that intensity variations are due to simple variation in
amount of absorbing material between source and detector. Apparently emission
irregularities dominate spatial absorption features. Sometimes variations in the
Carbon-K filter rates are accompanied by proportional variations in the Boron-K
filter rates. This behavior is to be expected if diffuse X-ray emission and absorp-
tion are in equilibrium along the line of sight, or if the emission is so local that
there is little (or at least constant) absorption in different directions.

OBSER VA TIONAL DA TA

500<E<IOOOeV

TRANSMISSION =EXP(-T), T = Ntr

500<E<IOOOeV

TRANSMISSION

191 196 201 206 211 TIME. SEC

250 255 260 265

270 TIME, SEC

Figure I.A-5. Counting rate of soft X-rays and X-ray transmission versus
time along the scan path (data from Bunner et al., 1972).

Lack of confirmed discrete point sources of soft X-rays (Bunner et al., 1969) and
the apparent granularity of the spatial structure of the diffuse flux (Gorenstein
and Tucker, 1972) suggest that if the source is stars of a special type, their local
space density must be large: > 10"2 (pc)"3 or more than 1 in 10 of all known
stars.

In an early publication on this subject (Bunner et al., 1969), we suggested a
population of stars with a scale height larger than that of the gas as a possible
source of the soft diffuse X-rays. The model provides the enhanced intensity
at high galactic latitudes, a source of the galactic plane emission, and requires no
extragalactic component. At energies between 0.5 and 1 keV, however, the model
predicts an enhanced intensity at intermediate galactic latitudes where absorption
by the interstellar gas has not yet dominated the effect of increased path length
through the emitting region. This enhanced intensity is not observed. The model
has been discussed in more detail by several other authors (Gorenstein and Tucker,
1972; Garmire and Riegler, 1972; Davidsen et al., 1972; Kato, 1972; Hayakawa,
1972, preprint).

Emission by the interstellar gas itself would appear to provide a reasonable model
for the origin of the diffuse X-rays in the galactic plane, because the absorption
optical depth in the plane is large at whatever longitude. X-ray emission is a very
inefficient process compared with ionization, however, and the resulting heating
of the cool interstellar medium, if the X-rays are produced in the gas, cannot be
accommodated even if a suitable charged-particle source is postulated ad hoc
(Bunner et al., 1971). A multicomponent interstellar medium requires further

DIFFUSE COSMIC X-RA YS BELOW 1 keV

10-

5 -

r-a .

•

• •

••
•

f+.

• •

ATMOSPHERIC ABSORPTION

RATE CORRELATION

ON ASCENT

DURING LOW b SKY SCAN

i i

^-B ^ ±

20 50

100

20 50

160 -284 eV RATE

100

Figure I.A-6. Counting rate of E < 180 eV X-rays versus rate of E < 280
eV X-rays (data from Bunner et al., 1973).

study as far as X-ray emitting possibilities are concerned. Emission by the inter-
stellar gas or by objects with the same spatial distribution as the gas, results in an
intensity proportional to (l-e"T), where r is the absorption optical depth. To match
the observations, therefore, an extragalactic component is required and there
results a net intensity proportional to A + Be"T. This same form of the intensity
dependence on r results from the assumption of extragalactic plus isotropic un-
absorbed components, as discussed by Davidsen et al. (1972).

EXTRAGALACTIC COMPONENT ?

Because of possible cosmological significance, there has been a persistent desire to
have at least a large portion of the high latitude diffuse soft X-ray flux be inter-
preted as extragalactic in origin. The point is simply that the lack of red-shifted
Lyman-a absorption in the spectra of quasars puts severe limits on the density of
a possible intergalactic unionized gas. Hence, it is argued that if the universe is
closed, the required mass must be in hot, ionized gas since the observed average
density of mass in the form of galaxies is small by a factor of about 60. Extra-
galactic soft X-rays would provide a possible indicator of this hot gas. Or, turning
the argument around, a demonstrated lack of extragalactic soft X-rays would put
limits on the possible density and temperature of a postulated hot intergalactic
medium (Field, 1972; Field and Henry, 1964).

The observed X-ray intensity enhancement toward the galactic poles, where the
gas density is small and expected X-ray transmission is large, suggests but by no
means demonstrates an extragalactic origin. In the first place, the sources could
be mingled with or just outside the galactic gas. In the second place, the

OBSER VA TIONAL DA TA

correlation of intensity with expected gas transmission is poor. Of course, there
are several possible causes for this poor correlation. The transmission is deduced
from 21 -cm hydrogen emission measurements, and helium, not hydrogen, is
responsible for most of the soft X-ray absorption (Brown and Gould, 1970).
There could be an unsuspected number of small unresolved cool clouds of gas,
and these would confuse both the column density measurement and X-ray
transmission estimates. These rationalizations would be comforting if we had
prior knowledge of extragalactic soft X-rays and knew there to be no high-
latitude galactic emission. But the reverse logic provides a decidedly weak case
(if any) for a hot, intergalactic medium.

We hoped our search for absorption by the gas of the Small Magellanic Cloud
(SMC) would clarify these matters. Before making the observation we decided
among ourselves that the most unsatisfactory result possible would be an X-ray
intensity that was constant as we scanned across the SMC for then, neither emission
nor absorption by the SMC would be clearly demonstrated. That, of course, is
exactly what happened (McCammon et al., 1971) as shown in Figure I.A-7.

300 100

SECONDS AFTER LAUNCH

Figure I.A-7. X-ray counting rate of X-rays (E < 280 eV) in directions near
the Small Magellanic Cloud. Solid calculated curves assume in A: absorption
by galactic and SMC gas; B: absorption by SMC gas only; C: absorption by
galactic gas only; and D: extrapolated power law spectrum extragalactic; the
rest: local origin, (McCammon et al., 1971).

Given this apparent lack of absorption by the SMC, we cannot exclude an extra-
galactic soft X-ray intensity (JQ) that is just compensated by emission from the
cloud itself. The consequences of this assumption, however, are rather interesting.
Let S be the X-ray emission rate per nucleon of stellar matter in the SMC, and let
n and n be the smoothed out and average nucleon density of stars and gas,
respectively. If emission and absorption just compensate, then

DIFFUSE COSMIC X-RA YS BELOW I keV 11

J an = Sn

o g s

where a is the X-ray absorption cross section per hydrogen atom (Brown and
Gould, 1970). The contribution to the extragalactic intensity from all galaxies
out to a distance *v c/2H is then

Jr a — n S

G 2H °

where n is the average density of galactic matter. According to Noonan (1971),
po for H°= 50 km s1 Mpc'1 is 7.5 X 10"32 g cm'3 so no is ~ 4.2 X 10"8 cm"3.
We then have

— a — n of — 1SMC

In the SMC, n /n is about 0.5; therefore, JG/JQ is about 0.8. In short, if we

attempt to save the hot intergalactic medium by supposing that the lack of
absorption by the SMC is really the result of self-emission, then the entire
supposed extragalactic soft X-ray intensity, or at least a large portion would
arise from the superposed emission from other galaxies. There is then little or
no intensity left to be accounted for by the hot gas.

If instead we suppose the emission to be somehow proportional to the gas of
the SMC and proportional to the gas in other galaxies too, with the same
emissivity , the value of JG/JQ is reduced by a factor of perhaps 10. This is
because we estimate the ratio of gas mass to star mass in the SMC is about
10 times that of other galaxies.

Figure I. A-8 shows how the SMC measurement and measurements of the diffuse
background radiation at higher X-ray energies restrict the temperature of a hot
intergalactic gas. This is essentially Figure I.A-1 of Field and Henry (1964), but
a Hubble constant H = 50 km ■ s"1 Mpc"1 has been assumed rather than 100.
The density assumed is sufficient to just close the universe (H = 1); the clumping
factor (C = <n2>/<n>2) is taken as 1 , the integration is carried out only to Z = 1 ,
and the expansion is assumed to proceed with 7 = 5/3.

As pointed out by Field (1972), the measured intensities in a real universe with
a given TQ must exceed those plotted. Because the SMC measurement falls so
near the "Big-Bang Envelope" line, it in fact (with H = 50) excludes very little-
only a band of temperatures near (2 X 106) K. On the other hand, and this is
the point I wish to emphasize, the diffuse soft X-ray measurements cannot,
taken alone, be said to provide positive evidence for a hot, dense, intergalactic
medium.

OBSER VA TIONAL DA TA

BIG BANG ENVELOPE
Zmax = l. C = l, y = 5/3, ft=l

S.S. ENVELOPE
(ANY T0)

0.1 I 10 100 I03 I04 I09 I06 I07 I08

X-RAY ENERGY (eV)

Figure I.A-8. Predicted X-ray intensities from a hot intergalactic medium with
density sufficient to close the universe (Field and Henry, 1964).

(Supported in part by NASA grant NGL 50-002-044)

REFERENCES

Brown, R., and R. Gould, 1970, Phys. Rev. D., 1, p. 2252.

Bunner, A., P. Coleman, W. Kraushaar, D. McCammon, T. Palmieri,
A. Shilepsky, and M. Ulmer, 1969, Nature, 223, p. 1222.

Bunner, A., P. Coleman, W. Kraushaar, D. McCammon, 1971, A strophys.
J. Letters, 167, p. 13.

Bunner, A., P. Coleman, W. Kraushaar, and D. McCammon, 1972,
Astrophys. J. Letters, 172, p. L67.

Bunner, A., P. Coleman, W. Kraushaar, D. McCammon, and F. Williamson,
1913, Astrophys. J., 179, p. 781.

Davidsen, A., S. Shulman, G. Fritz, R. Meekins, R. Henry, and H. Friedman,
\91 2, Astrophys. J., 177, p. 629.

DIFFUSE COSMIC X-RA YS BELOW I keV 13

Felten, J. E., 1973, X-Ray and Gamma-Ray Astronomy, Proc. oflAU
Symposium No. 55 (Madrid), H. Bradt and R. Giacconi, eds.,
D. Reidel, Dordrecht, Holland.

Field, G. B., and R. C. Henry, 1964, Astrophys. J., 140, p. 1002.

Field, G., 1972, Annual Review of Astron. and Astrophys. , 10, p. 227.

Garmire, G., and G. Riegler, 1972, Astron. and Astrophys. , 21, p. 131.

Gorenstein, P., and W. Tucker, 1972, Astrophys. J., 176, p. 333.

Grader, R., R. Hill, and J. Stoering, 1 970, Astrophys. J. Letters, 161, p. L45.

Hayakawa, S., T. Kato, T. Kohno, K. Nishimura, Y. Tanaka, and K. Yamashita,
1972, Astrophys. and Space Sci. , 17.

Kato, T., 1972, Astrophys. and Space Sci , 16, p. 478.

McCammon, D., A. Bunner, P. Coleman, and W. Kraushaar, 1971, Astrophys.
J. Letters, 168, p. L33.

Noonan, T. W., 1971, Proc. Astron. Soc. Pacific, 83, p. 31.

Palmieri, T., G. Burginyon, R. Grader, R. Hill, F. Seward, and J. Stoering,
1971, Astrophys. J. , 169, p. 33.

Sanders, W., 1973, unpublished Ph.D. Thesis, University of Wisconsin, Madison.

Savage, B. D., and E. B. Jenkins, 1972, Astrophys. J., 172, p. 491.

Seward, F., G. Burginyon, R. Grader, R. Hill, T. Palmieri, and J. Stoering, 1971,
Astrophys. J., 169, p. 515.

Williamson, F. W., 1973, unpublished Ph.D. Thesis, University of Wisconsin,
Madison.

Chapter II

A. THE X-RAY EMISSIVITY OF THE UNIVERSE:
2 TO 200 keV

Daniel Schwartz* and Herbert Gursky

American Science and Engineering, Inc.

INTRODUCTION

This paper will discuss observational results on the diffuse X-ray background
between 2 and about 200 keV. Appropriately to the sponsorship of this
Symposium by the Laboratory for Theoretical Studies, we wish to present
the results in a form suitable for theoretical discussion: namely, the volume
emissivity function B (E) [ergs/s • Mpc3 • keV emitted at energy E] . The
prescription for this is first to establish the spectral intensity I (E) [ergs/s •
cm2 • s -keV] measured at the earth, second to subtract the contribution due
to known, discrete sources, and third to unfold the equation

B (E) dV

(II.A-1)

which relates the measured intensity to the emissivity.

We may summarize the important characteristics of the diffuse X-ray back-
ground on which there is general agreement.

•

A real, cosmic X-ray background exists, which may be truly diffuse
or merely composed of discrete sources not yet resolvable. Nothing
in this paper will depend on which of those two pictures one adopts.

The diffuse X-rays are apparently isotropic over the sky, at least
to an extent which precludes a galactic origin.

All detailed theories have difficulty accounting for the production
of the measured energy into the diffuse spectrum in the sense that
they must hypothesize a rate of electron production, of heating,
or of cosmological evolution which is not otherwise observed.

"Speaker.

1 6 OBSER VA TIONAL DA TA

Strictly speaking, these three characteristics apply only to the energy range
between 2 and 40 keV where the isotropy over the entire sky has been
established by the X-ray experiments aboard the Uhuru and OSO-3 satellites.

EXPERIMENTAL PROCEDURES

The measurement of the precise spectral flux density of an isotropic diffuse
background is extremely difficult. The experimental problem is to determine,
as a function of energy, what fraction of the instrumental output is due to
internal background, where by the term "internal" we mean the output that
the instrument would have if no diffuse X-rays within the nominal bandwidth
entered the aperture. Internal background is also called "non-X-ray background,"
although in fact X-rays leaking from outside the aperture or higher energy
X-rays which interact with only a partial energy loss may both contribute to
internal background. Cosmic rays and geomagnetic particles are the primary
ultimate sources of background.

Several basic techniques have been used for estimating internal background:

• The earth, assumed to emit no X-rays, has been used as a "shutter"
and the entire instrument output obtained when the earth filled
the field of view was assumed to be internal.

• A physical shutter that is opaque to X-rays has been flown. It
either was moved into and out of place over the aperture or else
used to cover one of several identical detectors.

• Different collimator solid angles have been flown, again either by
motion of a shutter over one detector or fixed collimators over
several identical detectors.

The satellite experiments have allowed an additional technique:

• Observation of the modulation of the internal background as a
function of varying geomagnetic conditions, whereby it can be
separated from the constant isotropic X-rays.

By and large, all the above techniques are adequate to give what might be con-
sidered first -order accuracy by astrophysical standards (that is, within 25 to
50 percent errors). However, the photon counting statistics formally imply
a much higher precision; for a conservative example, a 200-s rocket flight
might count diffuse X-rays at a rate of 25 s"1 , and accumulate 5000 counts
between 2 and 10 keV. With statistical errors of only a few percent, the
following inadequacies of internal background estimation (numbered to
correspond to the techniques listed above) become apparent:

1 . Below 1 0 keV the earth can sporadically emit X-rays due to
auroral-type events. Above about 30 keV the atmospheric
albedo becomes comparable to the diffuse X-ray background.

THE X-RA Y EMISSIVITY OF THE UNIVERSE 1 7

2. X-rays can be generated by interactions in a mechanical shutter and
produce counts that would not be present when the aperture is open.

3. Data may be contaminated by diffuse geomagnetic electrons that
appear identical to diffuse X-rays. For example, an electron of about
70 ke V will on the average penetrate a 1 -mil Be window with a few
keV residual energy. However, because straggling is a dominant
effect for subrelativistic electrons, a wide bandwidth (say 50 to

100 keV) of incident electrons might be able to contribute counts
in the few keV range. These electrons are time variable, either
trapped or precipitating, and can occasionally be found even on the
L 1 magnetic shell (Schwartz, 1969). Electron fluxes far smaller
than are significant for geomagnetic studies, of the order of 0.01
(cm2-s-sr-keV)"1 at 70 keV, can contribute a few percent of the
diffuse X-ray counting rate. The existence of electrons of about
10 keV as a severe, sporadic contaminant to one-fourth keV X-rays
has been well known (Hill et al., 1970); however, the effect at higher
energies in any given rocket flight has generally been ignored.

4. A truly constant internal background component, for example,
radioactivity within the detector or vehicle, will not be modulated
as a function of geomagnetic conditions.

The reality of effects 1 and 3 as significant considerations for observations
between 7 and 40 keV was first shown by the OSO-3 experiment. Even when
the internal background is measured perfectly accurately, it may simply change
between the time it is estimated and the time when diffuse X-ray data is taken.
Such changes may be due to motion of the vehicle in space, a change in orien-
tation of the X-ray telescope axis relative to the earth's atmosphere or earth's
magnetic field, a change in the configuration of matter around the detector,
or temporal changes associated with geomagnetic activity. Table II.A-1
summarizes these background considerations, along with the principal method
used and the most likely source of remaining systematic error.

To stress the difficulty of the absolute measurement of a diffuse spectral den-
sity, we may digress to a familiar example from the study of the universal
microwave background. In radio astronomy an absolute flux is usually pre-
sented as the equivalent Rayleigh-Jeans blackbody temperature. Figure II.A-1
illustrates the derivation of the microwave temperature at 3.2 cm by Roll and
Wilkinson (1966). Briefly, that experiment used a Dicke-type radiometer
which measured the difference between an antenna horn pointed at the sky
and a cold load maintained near liquid helium temperature. The top bar
represents the measured cold load effective temperature. Each lower bar
represents the result after applying the correction listed. The key feature
here is that most of the corrections are of the same magnitude as the final
result and therefore must be known to the same precision desired for the

OBSER VA TIONAL DA TA

*T3

— < O

L. '-'

< BO

«f

-t->

„

d >»

0^ 5

■5 co

-*-"

C CO

u"2

CuOhJ

T3 S3

6 3

C/5

w w y

o o £

C/3 Od
O W

•S

•S 52

.2 *

cd o

C
O

o -S

O
-3

(U 5

•S <*>

3 a"

m* H

6 4

23 c*

CL,

O a-

3
u

Q
o

§i

•o

•B

.£5

<L>

o >

THE X-RA Y EMISSIVITY OF THE UNIVERSE

AFTER CORRECTION FOR--

SWITCH ASYMMETRY CORRECTION

HORN LOSS CORRECTION

ATMOSPHERIC RADIATION SUBTRACTED

MEASURED RADIOMETER EXCESS

ABSORPTION CORRECTION

MEASURED COLD LOAD K^^^^%^^^

I I I I I

K^^^W^%^^

pky^

MICROWAVE BACKGROUND TEMPERATURE

I I l I I I I I I L

0 2 4 6 8 10

TEMPERATURE AFTER EACH CORRECTION ('

Figure II.A-1. Derivation of the microwave temperature at 3.2 cm in the
experiment of Roll and Wilkinson (1966). Several of the corrections have
magnitude nearly equal to the final result of T = 3.0 K. The random errors
are an order of magnitude smaller than the estimated systematic effects.

microwave background. This experiment reported T = 3.0 ± 0.5 K, where
the error represents an estimate of systematic effects. This 0.5 K error
should be compared to a standard error of 0.06 K which the authors derived
due to the random errors in each correction term. In general, only such random
errors are reported for measurements of the X-ray background.

The generalizations discussed above, and the examination of the data presented
below, has led us to adopt the following point of view: Most measurements
of the flux density at various energies are reliable— they can be taken at face
value with their quoted errors and compared with other results. However,
direct measurements of a so-called "spectrum" by a single experiment are
much less reliable or useful. The unreliability results because the uncertain
systematic errors invariably are a different function of the energy than the
diffuse X-rays. Thus, one or a few data points at one end of the energy range
covered by a given experiment systematically distort the overall spectrum,
even if many other spectral points are quite accurate. The usefulness of a
spectral parameter is minimal for the following reasons: first, information
is lost by reporting a few spectral parameters instead of many flux density
measurements at various energies; second, the nonlinear least-squares fits
which must be used (due to the complicated spectral response to the typical
detectors) do not necessarily give unbiased estimates of the spectral parame-
teis; third, the procedure starts by assuming a general form for the spectrum,
such as power law or exponential shape; finally, it is not obvious how to
combine spectral parameters from two different experiments spanning slightly
different energy ranges— especially when each of those results has an esti-
mated error that excludes the other. The spectral parameters which we will

OBSER VA TIONAL DA TA

present below should be interpreted first of all as merely giving a numerical
representation of all the data, although one should certainly discuss the
physical interpretation of any spectral representation.

OBSERVATIONAL RESULTS

Figure II.A-2 presents a selection of published flux density points for the
diffuse X-rays between 2 and 200 keV. The plot gives the energy flux in

"i — i — i i i i 1 1 r

$-oso-m

0-LLL

$ -LEIDEN

J-ASE

+ -BOLOGNA

§-PRL

£-TATA

f-SACLAY

■-GSFC

ENERGY (keV)

Figure II.A-2. A selection of published energy flux measurements of the
diffuse X-ray background. Results presented only by giving spectral parameters,
and points with greater than 30 percent error estimates are excluded. The data
show a general consistency, with the high rate points around 10 keV and
150 keV possibly due to electron contamination. The slope increases with
higher energy.

keV/keV-cm2-s-sr. Results reported only by giving spectral parameters are
not included. Points with reported relative errors larger than 30 percent, and
estimates of upper limits, are also excluded. In general, only the latest results
of a given group are shown. Although the points with the smallest error bars
tend to be hidden in such a plot, we can see that the bulk of the points do fall
within a ± 50-percent error band, and therefore we may expect the precision
of the mean to be still higher. The balloon-borne measurements shown here
(except for Manchanda et al., 1972) do not contain additional so-called
"Compton scattering" corrections for reasons discussed below. The total
data suggest a gradual steepening of the spectrum from a few keV up to
100 keV; detailed analysis of several of the experiments confirms this
conclusion.

THE X-RA Y EMISSIVITY OF THE UNIVERSE 21

Rocket-borne Observations

The key feature of rocket experiments is that they generally operate in the
"cleanest" environment with regard to internal background. They are above
the secondary cosmic radiation produced in the earth's atmosphere and below
trapped particle populations. (Sporadic electron precipitation events may still
affect any one observation.) The major drawback is that the observation lasts
at most a few minutes. This usually does not allow, for example, a program
that alternates measurements of diffuse and internal background to verify
that the latter is constant.

Consider first the proportional counter observations shown in Figure II.A-2
(LLL: Palmieri et al., 1971; ASE: Gorenstein et al., 1969; GSFC: Boldt et al.,
1969; PRL: Prakasarao et al., 1971). In this energy range, 2 to 10 keV,
shielding and collimation is easily done with passive structural elements. The
fields of view used range from 20 square degrees in the LLL experiment
(shown as the eight largest diamonds between 2.4 and 8.7 keV) to 500 square
degrees by GSFC. The PRL measurement was carried out at the geomagnetic
equator; the GSFC and ASE flights from White Sands occurred at a magnetic
shell of approximately L = 1.7 to 1.8. The PRL counters were filled with a
xenon/methane mixture, the others with an argon/methane mixture. ASE and
LLL determined internal background with a rocket door closed, PRL while
looking at the earth, and GSFC by having a movable shutter that gave five
different solid angles between 0.125 and 0.17 s as well as a completely
occulted position.

Agreement among the various experiments is rather good. This may be
expected because proportional counters generally have several hundred cm
areas and because the X-ray flux is constantly increasing to the lowest ener-
gies. The signal-to-background ratios obtained were between 3 and 10 to 1.

We have omitted two results which suggested spectral line features in the back-
ground around 5 and 7 keV (Ducros et al., 1970; Henry et al., 1971). Boldt
et al. (1971), have reported an upper limit for such a feature at 7 keV of a
factor of 10 below the NRL result. This applies to an observation at galactic
latitudes from +40° to the North Pole. It is probably fair to say that with the
difficulties of establishing a continuum shape accurately, the existence of line
fluxes remains to be proven in future experiments. We may comment that
the unfolding of spectral data from a proportional counter response is by no
means trivial. Such unfolding basically depends on calculation rather than
calibration, since both the X-ray and particle spectra in space are very
different than in the laboratory.

The proportional counter measurements may be compared with a satellite
experiment of LLL (Cunningham et al., 1970; three small diamonds at 4.6,
8, and 12 keV). This involved a Nal crystal with a 0.76-cm2 • sr telescope

22 OBSER VA TIONAL DA TA

factor aboard a polar orbiting satellite. A mechanical shutter periodically
occulted the detector to allow background estimates. Only about 15 minutes
of data (apart from solar and discrete source observations) was taken before
a failure during the second day of operations.

The results of LLL (intermediate sized diamonds: Toor et al., 1970) and
Bologna (Horstman-Moretti et al., 1971) were obtained with rocket-borne
Nal counters. These detectors employed passive shielding lined with a
plastic anticoincidence scintillator to define fields of view of about 900 square
degrees. The Livermore data are noteworthy because this was the only
experiment other than OSO-3 to span a range from below 10 keV to above
40 keV. The spectral results were reported as allowing a power law fit;
however, the error bars above 30 keV are clearly large enough to be also
consistent with a considerable change in slope.

The four data points of the Bologna group (a measurement of 0.62 ± 0.04 at
52 keV is blacked out by other data points) are obtained with an ideal
technique; one of four identical detector units is blocked so that internal
background measurements are continually taken along with the diffuse
X-ray data. Again, we suggest contamination by a sporadic electron popu-
lation as the cause of the apparently high points at 90 and 1 50 keV. This
is not an unlikely occurrence at the invariant magnetic latitude of 38°
(L = 1 .6) of this observation. The 0.2-mm (54-mg/cm2) Al window would
allow electrons of roughly 100 to 400 keV initial energy to enter the Nal
volume with 50 to 200 keV residual energy, considering straggling. The
OSO-3 upper limit of 0.18 keV/(keV-cm2-ssr) at 150 keV, not shown in
Figure II.A-2, cannot otherwise be reconciled with this data.

Balloon-borne Observations

The dominant feature of the balloon-borne observations is that there exists
a significant, diffuse flux of X-rays produced in the atmosphere. These must
be separated from the diffuse cosmic X-rays by some indirect line of
reasoning. Figure II.A-3 schematically illustrates the observational situation.
The top solid curve is the counting rate that a vertically pointed telescope
with a 9° to 20° cone-angle (as used in the four experiments plotted on the
previous graph) might record as a function of atmospheric depth. The lower
solid curve is that which a shuttered detector might record and is the internal
background defined earlier. Only the Leiden-Nagoya group actually used
such a shutter; the others effectively lumped internal background along
with atmospheric.

At depths below 10 to 20 gm/cm2, the difference in the two curves is due
entirely to the atmospheric X-rays. (We intend the figure to show that the
atmospheric X-rays can have a different, although similar, dependence on
depth compared to the internal background.) Both the internal and

THE X-RA Y EMISSIVITY OF THE UNIVERSE

5 Ri

30KeV X-RAYS
APERTURE OPEN

10 100

ATMOSPHERIC DEPTH lgm/cm2'

1000

Figure II.A-3. Representation of the counting rates of a vertically
mounted, wide-aperture telescope. Diffuse X-rays cause the turn-up of
the "aperture-open" curve. One must estimate the atmospheric
contribution (dashed line) to deduce the diffuse intensity.

atmospheric background originate from the soft component of the energy
degradation of the primary cosmic rays, and show the Pfotzer transition
maximum at about 90 gm/cm2 .

The turn-up of the rates at altitudes higher than about 10 gm/cm2 is inter-
preted as the observation of X-rays external to the atmosphere. X-rays of
30 keV have a mean free path of 3.4 gm/cm2 for photoelectric absorption
compared to ceiling depths of 3 to 7 gm/cm2 attained in the various experi-
ments. The dashed curve represents an extrapolation which each experimenter
must make for the assumed behavior of the atmospheric X-rays. The
difference (Rx - R2) is then multiplied by the photoelectric attenuation at
the given ceiling depth (which may be a factor of 2 to 8 correction) to derive
the diffuse flux external to the atmosphere.

There has been some discussion (for example, Horstman and Horstman-
Moretti, 1971) that additional corrections need be applied due to single and
multiple scattering of diffuse X-rays by the atmosphere, which eventually
enter the detector aperture. This discussion is important and valuable since
such a physical process must certainly take place. However, it is not appro-
priate to apply such a correction to (Rl - R ) for the following simple
reason: Once the diffuse X-ray scatters in the atmosphere it loses its identity
and is no different than an atmospheric X-ray that might be produced by

24 OBSER VA TIONAL DA TA

electron bremsstrahlung at the exact same location. But all the atmospheric
X-rays have presumably been accounted for by the dashed line extrapolation.
One may well ask whether the dashed line is an accurate extrapolation of the
atmospheric background, but this is a very different and very important
problem.

In principle, one could study the difference between Rl and R2 as a function
of depth, and test whether it changes in the exact manner expected for
photoelectric absorption of X-rays external to the atmosphere. In practice
this is not decisive because the points below float altitude are only sampled
for a relatively brief time during ascent or descent. As (R1 - R2) becomes
smaller, the absolute error on this difference becomes larger, and it may be
that the data allows anything between zero and infinite absorption.

Strictly speaking, we might say that the true atmospheric X-ray curve could
vary considerably from the intuitively simple extrapolations used, and that
there might in fact be no diffuse X-rays at all. Returning to Figure II.A-2,
we can let the scatter of the data speak for itself in illustrating the intrinsic
accuracy which has been obtained. The lowest and highest points, by the
Tata Institute (Manchanda et al., 1972) and Saclay (Rothenflug et al., 1968)
groups, used an exponential law extrapolation. The Leiden (Bleeker and
Deerenberg, 1970) and Physical Research Laboratory (Rangan et al., 1969)
groups used a power law extrapolation for the rates versus depth. Each pair
of groups spanned magnetic shells at least from L = 1 to L = 1.7.

The detectors used in these experiments were all Nal crystals with some
combination of passive shield and plastic anticoincidence. These give
relatively high susceptibility to internal background. As the groups at
UCSD, UCB, and MIT have developed detectors with lower internal back-
ground by using 4 it active-anticoincidence techniques, they have systema-
tically tended to reduce the solid angle and concentrate on discrete source
observations.

We may suggest a prescription for obtaining a more objective determination
of the atmospheric X-ray contribution at ceiling. This is based on the concept
of a source function S (E, x) (X-rays of energy E produced (cm • s)*1 at a
depth x). This technique has been used successfully by Peterson, Schwartz,
and Ling (1973) to interpret counting rates of atmospheric 7-rays as a
function of depth. Figure II.A-4 shows the basic geometry. The function S
is strictly a convenient mathematical form, containing a few constants to be
determined. With the detector in a fixed orientation at a depth h, the source
function multiplied by the projected detector area A (0) and by the attenu-
ation exp (-//r) (where ju is the total coefficient for any interaction) is
integrated over all of space. The unknown constants in S should be deter-
mined while the detector is at large depths (h) and/or while it is oriented
downward. Then with the detector pointed upward at the float altitude,

THE X-RA Y EMISSIVITY OF THE UNIVERSE

TOP OF ATMOSPHERE

SOURCE VOLUME

DETECTOR

C(E,h)=l/2J"A(0) S(E,x)e"Mr sm0d0dr

Figure II.A-4. Geometry for calculating the contribu-
tion of atmospheric X-rays to the counting rate C
(E, h) of a detector at depth h (from Peterson et al.,
1973). An empirical volume production rate function
S (E, x) is constructed as a function of depth x. The
integral over the volume of the atmosphere gives the
contribution for a fixed detector orientation.

C (E, h) would simply be calculated and subtracted from the total output.
Physically, of course, S will contain a contribution from Compton-scattered
diffuse X-rays; however, this need not ever be considered explicitly.

OSO-3 Observations

Finally, we will discuss the data points obtained by the UCSD X-ray telescope
aboard OSO-3. These points were relatively inconspicuous in Figure II.A-2
because of their small error bars; yet they are of significance as the only case
in which a power law spectrum could not fit the data of one single experi-
ment. Because of this significance, Schwartz and Peterson (preprint) have
reconsidered the results with regard to some suggested corrections for
spaliation-induced radioactivity, fluorescence radiation from the shield, and
energy dependence of the geometry factor, and we have confirmed the
inconsistency of a power law with the OSO-3 data. The best fit of a power
law gives X2 = 20 for 3 degrees of freedom.

Briefly, the OSO-3 experiment was a 9.5 cm2 Nal crystal, actively collimated
by a Csl annulus to a 23° full-width half -maximum conical field of view.
The satellite had a 550-km altitude, 33° inclination orbit so that magnetic
shells from L = 1 to L = 2 were sampled, and the lower edge of the South

26 OBSER VA TIONAL DA TA

Atlantic trapped particle region was traversed during half of the 16 orbits per
day. The data were telemetered in six logarithmically spaced channels between
8 and 210 keV. Certain integrated and solar-pointing rates were also
telemetered.

The most serious contributor to the background was the existence of sporadic,
charged particles. Selection criteria to minimize contamination were developed.
These limited the upper threshold integral rate, required L< 1.2, and accepted
data only when pointed within the local magnetic loss cone. This caused
rejection of about 80 percent of the data.

The next most serious source of background was due to radioactivity, which
built up when inside the trapped particle regions and which then decayed
until the next traversal of the South Atlantic Anomaly. A 15-hour half-life
decay curve gave a good fit to the monitor count rates in the interval 30
minutes to 12 hours after penetrating the particle belts. The activation
coefficients derived from these monitor rates were used to correct the diffuse
counting rates over the same time span. The diamonds and upper limit in
Figure II.A-5 (taken from Schwartz and Peterson, preprint) show the effec-
tive spectrum at the Nal detector, due to radioactivity. Phenomenologically,
this spectrum is interpreted as Compton scattered 7-rays from the Mg24
daughter produced by the reaction Al27 (n, a) Na24 1 ;! nour Mg24 taking
place throughout the satellite. The solid line is a spallation spectrum measured
by Dyer and Morfill (1971), and plotted with an arbitrary normalization.
The horizontal bars integrate this spectrum over the OSO-3 energy channels
and normalize it to be consistent with the 7.7- to 12.5-keV limit. Thus the
spallation mechanism is probably not significant on this time scale.

This radioactivity correction could be made because it varied on a 1 5-hour
time scale. However, radioactivity with a half- life of a week or longer would
not decay significantly in one day and might in principle be a constant,
unnoticed contaminant of the data. By subtracting the rates when looking
at the earth from the sky rates on day 44 after launch, we show that at least
95 to 90 percent of the reported diffuse flux for the three channels from 7.7
to 38 keV cannot be contaminated by radioactivity. The points between
38 to 110 keV might require further downward correction, but this will only
accentuate the inability of a single power law to fit the data.

Examination of the detailed rates versus time after launch in the 38- to
65-keV channel, compared with the predicted build-up curve using the proton
dose by Dyer, Engel, and Quenby (1973) led us to conclude (Schwartz and
Peterson, preprint) that at most one-third of that proton dose would be the
appropriate normalization. We have increased the error bars of the upper
channels so that such a radioactivity correction (if valid) would only reduce
the quoted fluxes by 2 standard deviations.

THE X-RA Y EMISSIVITY OF THE UNIVERSE

ENERGY (keV)

Figure II.A-5. Diamonds and upper limit: Effective
spectrum of radioactivity background observed by the
0S0-3 X-ray telescope immediately after emergence
from the proton belts. Solid line: an effective spec-
trum due to spallation measured by Dyer and Morfill
(1971) as the difference in Csl crystal output
measured 86.5 minutes and 1 1 .2 hours after irradiation.
The normalization is arbitrary. Horizontal bars: The
same spallation spectrum integrated over the OSO-3
energy channels and normalized for consistency with
the measured limit at 10 keV (from Schwartz and
Peterson, preprint).

28 OBSER VA TIONAL DA TA

The most significant correction necessary to the previous 0S0-3 results has
been the allowance for K-shell X-radiation escaping from inside the collimator,
as suggested by Horstman (Dumas et al., 1973). Diffuse X-rays between 35
and a few hundred keV striking the inside of the Csl collimating annulus
would not trigger the shield anticoincidence threshold. A certain fraction of
the resultant K-escape X-rays will be emitted into the central detector,
causing a spurious contribution to the 22- to 38-keV channel. The Monte
Carlo program of J. Matteson, which has been used extensively at UCSD to
predict background rates of X-ray and 7-ray detectors, was used to calculate
an effective telescope factor (solid angle times area) for such fluorescent
X-ray events as a function of the incident photon energy above 34 keV. The
product of this telescope factor and the diffuse spectrum 2200 E"3 previously
estimated (Schwartz et al., 1970) was integrated from 34 to 210 keV. As a
result, the point at 30 keV was reduced 17 percent.

Summary

In Figure II.A-6 we attempt to summarize the most reliable data selected
according to the following criterion; the experiment either was operated
over a range of geomagnetic conditions or it incorporated some direct
means for assessing effects of electron contamination. The Livermore
rocket experiment (Palmieri et al., 1971) had a methane-filled anticoinci-
dence proportional counter over the entrance to their argon detector. This
experiment, of all the rocket and satellite observations, should be uniquely
free of charge d-par tide contamination. The ASE experiment (Gorenstein
et al., 1969) incorporated pulse-shape discrimination, which is sensitive to
relativistic electrons which may deposit only a few keV total energy but
spread out over a long path. That experiment also had one counter unit
with a 1-mil Be window and three counter units with 3-mil Be windows.
These windows would show very different transfer characteristics for the
70- to 100-keV geomagnetic electrons.

The three experiments of the Leiden-Nagoya group (Bleeker and Deerenberg,
1970) provide key evidence for the reality of a diffuse component above
40 ke V. The experiments took place at 20° , 40° , and 50° geomagnetic
latitudes. The flux densities at the various latitudes are in reasonable agree-
ment, while the inferred component of 20- to 40-keV atmospheric X-rays
is a factor of 5 higher at the northern latitude.

The solid curve shows the function

( 10E-as2for KE<23keV
I (E) (keV/ke V cm2 -sr-s) = j 1 4mM 37 , p . „ . ., 01. A-2)

/ 140E'1-37 for E > 23 keV

THE X-RA Y EMISSIVITY OF THE UNIVERSE

-r^ri — i i i i i 1 1 1 1 — i i i i 1 1 1 1

$- oso-m

O-LLL
^-ASE
^- LEIDEN -NAGOYA

I I I I I I

10' 10

ENERGY (keV)

lO-

Figure II.A-6. An attempt to select the most reliable experimental data
between 2 and 200 keV. The observations either utilized some direct means
for assessing effects of electron contamination or else operated over a range
of geomagnetic conditions. The solid curve shows the power law 10E* " ' keV/
(keV • s • cm2-sr) below 23 keV and 140E"137 above 23 keV. The dashed line
is the function 3.3 exp (-E/34.4).

and the dashed curve is

1(E) =3.3 exp (-E/34.4)

(II.A-3)
(see Cowsik and Kobetich, 1972)

The sharp break in the power law representation does not have physical
reality-this is merely a minimum parameter power law representation of the
data. The key observational conclusion of such a representation is that the
overall change in the slope is at least an exponent of 0.9. The errors in the
power law indices are roughly ±0.1 below 23 keV and ±0.15 from 30 to
100 keV. The error in the effective kT of Equation (II.A-3) is somewhat
larger, as we have arbitrarily tried to fit the data only in the 10- to 100-keV
range.

In Figure II.A-7, we wish briefly to compare with the data from a few
hundred eV to a few MeV. We have not attempted any completeness in
the higher and lower energy data. The ASE point at 270 eV was obtained
with a focusing collector and is an upper limit in the sense that Gorenstein
and Tucker (1972) argue it might all result from galactic sources. The

OBSER VA TIONAL DA TA

Iff

I — i — i i i ii ii

-i — i — i i 1 1 1 ii

-i — i i 1 1 1 M

"1 1 — I I II 14-1

10 10'

ENERGY (keV)

I0H

Figure II.A-7. The data and power-law fit of Figure II.A-6 is shown
along with a sample of measurements at higher and lower energies. The
Ranger-3 data are only an energy-loss spectrum, while for the Apollo
results the true photon spectrum has been unfolded and a correction
applied for spallation induced radioactivity.

Wisconsin upper limit is based on the absence of absorption by the Small
Magellanic Cloud (McCammon et al., 1971 , see also Chapter LA). The
Ranger-3 data (Metzger et al., 1964) represent only an energy-loss count
rate spectrum, while the Apollo-1 5 data (Trombka et al., 1973) have been
unfolded to a photon spectrum and corrected for spallation induced
radioactivity (see Chapter III. A).

THE EMISSIVITY FUNCTION

If we assume a constant emissivity B (E) per unit coordinate volume, through
which we look a distance R , then Equation (II.A-1) becomes simply

I(E)~B(E)Rr

(II.A-4)

THE X-RA Y EMISSIVITY OF THE UNIVERSE 31

Equation (II.A-4) holds to within a factor of 2 for the popular models of
Friedman cosmologies, providing there are no significant evolutionary effects
and providing that we take Rmax = 1/2 (c/Hq). We will adopt Hq = 75 km/s.
Mpc Then from Equation (II.A-2), we have

i 2.1 X 1CT26 E-°-52 keV/(cm3 • s • keV) for E < 25 keV
B(E)= \

j 2.9 X 10"25 E-1-37 keV/(cm3 • s • keV) for E > 25 keV.

For the integrated emissivity between 2 and 7 keV, B = 2.2 X 1039 ergs/
s • Mpc3. We stress that B is determined directly from the observations, and
subject to the qualifications above, it will not change significantly. The
red shift will preserve a power-law shape.

CORRECTION FOR DISCRETE SOURCES

Characteristics of the classes of discrete extragalactic sources identified in the
2U catalog (Giacconi et al., 1972) are summarized in Table II.A-2. In Figure
II.A-8 we illustrate how they modify the emissivity function. We must stress
that the spectra and total emissivities are very poorly determined for the
cases where we only have one object in the class: the Seyfert galaxy NGC
4151, the radio source Cen A, and the quasar 3C 273. We have normalized
the total power of each source according to the quantity n j and represented
its spectrum as the flat end of the range allowed by Uhuru in order to
obtain the closest agreement to the background shape. The solid curve is
the power-law representation presented earlier. The higher dashed curve
represents the subtraction of the identified extragalactic sources. The bottom
dashed curve represents a possible residual emissivity if we hypothesize that
the unidentified high-latitude sources are a new class of extragalactic object
that produces one-half the background observed from 2 to 7 keV.

INTERPRETATION AS THERMAL BREMSSTRAHLUNG

Figure II.A-9 replots the data on a semilog scale. Correction for discrete
sources allows the exponential fit to hold above 5 keV. The two solid
curves (Field, 1972) are the temperature-independent lower limits to
radiation from a hot intergalactic plasma of sufficient density to close the
universe. In the big bang model, the gas is suddenly heated to a temperature
Tq at an epoch z = 1 and cools adiabatically with an index 7 = 5/3. Field
uses a Hubble constant H = 50 km/s • Mpc. The predicted spectrum is
approximately exponential with e-folding energy k T and would be tangent
to the lower limit curve at E = (0.57) k T . The observed radiation falls a

O v ' o

factor of 2 below the minimum, and the discrepancy is worse if the

OBSER VA TIONAL DA TA

03
H

3
O

l-C

M
o

a5
X

3
O
00

c
o
u

action
of

kgrou
f

o
q

■<3-

CN
CN

* -

NO
<*>

> oo

1— 1

T— 1

i—i

v£>

CN*

/— s

CO
O

c o

i— <

■M o

'& <U

0\
CO

i/i

O ~2?

•3 oo

1— t

1 a>

Ipectrum
kT, for
-lOkeV
ange)

m
O

»n

l/J W «— 1 l-i

»— )

s_^

»— i

3 g>

e 1

.5 "ca

i—i

i-^

*— <

»-H

""•

1-4

"O wj

Cfl

■»->

3
*o O

S3 V)

g a

s ^

o a

03 V

03 *o3

■•? O0

>~> a.

3
o

C/3

*o3

S 5b

5 ^

THE X-RA Y EMISSIVITY OF THE UNIVERSE

I039

r-> I02

CD
CO

CO
CO

o
>

I035

1 I I I I I I L

MILKY
WAY

I I I I I I I I

ICf

10'
ENERGY (keV)

10'

Figure II. A-8. Volume emissivity functions: Top line
is emissivity derived from power-law representation
of Figure II.A-6; five curves for identified extragalactic
sources are derived from estimates of intrinsic lumi-
nosity and spectra based on the 2U catalog; upper
dashed curve corrects diffuse emissivity for these
sources; lower dashed curve is resulting diffuse emis-
sivity if we postulate that unidentified high latitude
sources (observed to have a spectrum with kT =
5 keV) comprise 50% of the 2-7 keV diffuse back-
ground; points at 20 keV are from UCB and UCSD
OSO-7 experiment.

OBSER VA TIONAL DA TA

n 1 1 1 1 1 1 1 r

x ASE
O LLL
□ LEIDEN- NAGOYA

• oso-n

20 40 60 80 100 120

ENERGY (keV)

Figure II.A-9. The data from Figure II.A-6 are plotted on a semilog
scale. The dashed line is the function 3.3 exp (-E/34.4). The solid lines
are "lower limits" to the emission from a hot intergalactic plasma of
sufficient density to close the universe calculated by Field (1972) with
a Hubble constant H = 50 km/s • Mpc3. The X-ray observations are at
a factor of two discordant with the big bang model.

intergalactic medium is not smooth. The disagreement implies one or more of
the following:

• The density of intergalactic plasma is 2Vz less than required for a
closed universe, or n = 2 X 10"6 particle cm"3;

• The Hubble constant is a factor 21/3 smaller, or H = 40 km/s • Mpc;

• The temperature T < 3 X 107K.

THE X-RA Y EMISSIVITY OF THE UNIVERSE 35

We hope we have not spent too much time discussing X-ray results at a 7-ray
symposium. We have tried to make the point that the X-ray observations have
a relatively high level of precision and that we can start using them to do
some interesting physics. We look for many more exciting results and new
ideas to come as study of the spectrum and of the isotropy is extended into
the MeV region.

ACKNOWLEDGMENTS

We acknowledge conversations with L. Peterson, C. Dyer, B. Dennis, and
H. Horstman regarding interpretation of the OSO-3 X-ray data. We thank
E. Kellogg and S. Murray for discussion of the Uhuru results on extragalactic
sources.

(This research was supported by NASA contract NAS 8-27973.)

REFERENCES

Bleeker, J. A. M., and A. J. M. Deerenberg, l910,Astrophys. J., 159, p. 215.

Boldt, E. A., U. D. Desai, S. S. Holt, and P. J. Serlemitsos, 1969, Nature,
224, p. 677.

Boldt, E. A., U. D. Desai, S. S. Holt, and P. J. Serlemitsos, 197 '1, A strophys. J.
Letters, 167, p. LI.

Cowsik, R., and E. J. Kobetich, Astrophys. J., Ill, p. 585.

Cunningham, C, D. Groves, R. Price, R. Rodrigues, C. Swift, and H. Mark,
1970, Astrophys J., 160, p. 1177.

Ducros, G., R. Ducros, R. Rocchia, and A. Tarrius, 1910,Astron. and
Astrophys., 7, p. 162.

Dumas, A., H. Horstman, E. Horstman-Moretti, and D. Brini, 1973, X-Ray
and Gamma-Ray Astronomy, Proceedings of IAU Symposium No. 55,
(Madrid), H. Bradt and R. Giacconi, eds., D. Reidel, Dordrecht, Holland.

Dyer, C. S., A. R. Engel, and J. J. Quenby, 1973, X-Ray and Gamma-Ray
Astronomy, Proceedings of IAU Symposium No. 55, (Madrid), H. Bradt
and R. Giacconi, eds., D. Reidel, Dordrecht, Holland.

Dyer, C. S., and G. E. Morfill, 1971, Astrophys. and Space. Sci., 14, p. 243.

Field, G., 1912, Annual Review of Astronomy and Astrophysics, 10,
Annual Reviews Inc., Palo Alto, p. 227-260.

Giacconi, R., S. Murray, H. Gursky, E. Kellogg, E. Schreier, and H. Tananbaum,
1972, Astrophys. J., 178, p. 281.

36 OBSER VA TIONAL DA TA

Gorenstein, P., E. M. Kellogg, and H. Gursky, 1969, Astrophys. 7, 156, p. 315.

Gorenstein, P., and W. Tucker, 1972, Astrophys. /., 176, p. 333.

Henry, R. C, G. Fritz, J. C. Meekins, T. Chubb, and H. Friedman, 1971,
Astrophys. J. Letters, 163, p. L73.

Hill, F. W., R. Grader, and F. D. Seward, 1970,/. Geophys. Res., 75, p. 7267.

Horstman-Moretti, E., F. Fuligni, and D. Brini, 1971, IlNuovo Gmento, 6 B,
p. 68.

Horstman, H., and E. Horstman-Moretti, 1971, Nature Phys. Set, 229, p. 148.

Manchanda, R. K., S. Biswas, P. C. Agrawal, G. S. Gokhale, V. S. Iyengar,
P. K. Kunte, and B. V. Sreekantan, 1972, Astrophys. and Space Sci, 15,
p. 272.

McCammon, D., A. N. Bunner, P. L. Coleman, and W. L. Kraushaar, 1971
Astrophys. J. Letters, 168, p. L33.

Metzger, A. E., E. C. Anderson, M. A. Van Dilla, and J. R. Arnold, 1964,
Nature, 204, p. 766.

Palmieri, T. M., G. A. Burginyon, R. J. Grader, R. W. Hill, F. D. Seward,
and J. P. Stoering, 1971, Astrophys. /., 169, p. 33.

Peterson, L. E., D. A. Schwartz, and J. C. Ling, 1973,7. Geophys. Res.,
in press.

Prakasarao, A. S., D. P. Sharma, U. B. Jayanthi, and U. R. Rao, 1971,
Astrophys. and Space Sci., 10, p. 150.

Rangan, K. K., P. D. Bhavsar, and N. W. Nerurkar, 1969,7. Geophys. Res.,
74, p. 5139.

Roll, P. G., and D. T. Wilkinson, 1966, Phys. Rev. Letters, 16, p. L405.

Rothenflug, R., R. Rocchia, D. Boclet, and P. Durouchoux, 1968, Space
Res. VIII, p. 423.

Schwartz, D. A., 1969, unpublished Ph. D., Thesis, University of California
at San Diego.

Schwartz, D. A., H. S. Hudson, and L. E. Peterson, 1970, Astrophys. 7,
162, p. 431.

Toor, A., F. D. Seward, L. R. Cathey, and W. E. Kunkel, 1 970, Astrophys. 7,
160, p. 209.

Trombka, J. I., A. E. Metzger, J. R. Arnold, J. L. Matteson, R. C. Reedy, and
L. E. Peterson, 1973, Astrophys /., 181, p. 737.

B. ATMOSPHERIC CORRECTIONS TO BALLOON
X-RAY OBSERVATIONS

H. Horstman*

University of Bologna

The group at Bologna (Brini, Fuligni, and Horstman-Moretti) has some new
results on the diffuse background between 30 and 200 keV from a second
rocket flight. On this particular flight we used detectors with different
geometric factors intending to apply the usual idea about these, that is, that
the counting rate for one detector is the effective geometric factor of the
detector times the diffuse flux with the instrumental background added on.
Two different detector shapes were used, one long and one short, with the
geometric factor of one being about twice the other. We would like to
assume that the background is the same in both detectors and that the
diffuse flux is isotropic. Then, we extrapolate the counting rate to zero
geometric factor to find the instrumental background; it is really like fitting
the difference between the counting rates of the two detectors. Because of
this, the statistical errors of the results are larger on this flight than on our
previous flight in which we used a screened detector to determine the back-
ground. The only difficulty with this method is that the background of the
two detectors is not really the same.

Ground measurements of local diffuse X-rays indicate that the longer detector
has the higher background. This longer detector has the smaller geometric
factor, and so, when we try to do the extrapolation to zero geometric factor,
we obtain an overestimate of the instrumental background. If the same
effect occurs at altitude, the present results on the diffuse flux have to be
considered as lower limits. A power-law fit gives 47 E"2'1 * °'25photons
(cm2 • s • sr • keV)"1 . See Figure II. B- 1 .

We have recently corrected our old S-l 1 results for the lead-K X-ray of the
collimator and also found a mistake in the geometric factor program which
led us to decrease the derived fluxes by 13 percent. The corrected spectrum

'Speaker.

OBSER VA TIONAL DA TA

- , PRIMARY PHOTONS

^T^^T^^

measured total

UNSCATTERED \. ^v>

PRIMARIES >v >

~ v<^SCATTERED

" N,

/ \

yy^ PRIMARIES

••

\/ PHOTONS PRODUCED

jO^ y. '

><^ IN THE ATMOSPHERE

•

j^r * s

Ss ' ' '

\ TOTAL PRIMARY

Ss ^ , ■

\ CONTRIBUTION

// "* '

V FROM

PSEUDO /y '' .' j

x MONTE CARLO

ATMOSPHERE"/^ , " s /

// ' "N ' /

// -' ' /

/^extrapolations,' /

Ss x' /

y \

sy/ s /

\ \

^ // 1 1 1 1 /! [ 1 1 1 1/ 1 II

_l_

1 1 I

\ i i i i\ i i i i

1 10

RESIDUAL ATMOSPHERE (G/CM2)

Figure II.B-1. Rough dependence on the depth for 40-keV photons
assuming a power-law dependence of the atmospheric X-rays.

is 36 E"20 * photons (cm2 • s • sr • keV)"1 . This lies somewhat above the
new results by an amount which can be explained by the different methods
used for the background determination.

Both of these results could be subject to the contamination by precipitating
electrons mentioned by Schwartz (Chapter II. A) but the agreement between
the two results is not bad. If there are electrons there, either they are
constantly precipitating or they are a small fraction of the flux observed.
I can, unfortunately, say little about possible contamination.

The present results lie slightly above and somewhat overlapping the Indian
balloon results corrected for Compton scattering and lie more on top of
Bleeker's uncorrected results. That brings up the other point that I wanted
to talk about which is the question about Compton scattering that Schwartz
brought up (see Chapter II. A). I could not let that remark go by because I
feel it contains a misinterpretation of the transport of the primary X-rays in
the atmosphere.

A TMOSPHERIC CORRECTIONS 39

If I understand Dan (Schwartz) correctly, he was saying that these Compton
corrections are "not appropriate" because once a photon is scattered it is
just like an atmospheric X-ray. Because the only means to distinguish pho-
tons is by their depth dependence, this would mean that the primary
scattered flux follows the same depth dependence as the atmospheric X-rays;
therefore, when the extrapolation is performed, the scattered diffuse primary
X-rays are also extrapolated out. It is, however, easy to see roughly what
the depth dependence of the scattered primary-photon flux is, and where
the peak intensity is reached. The results of our Monte Carlo calculations
for the propagation of the diffuse primary X-rays in the atmosphere show
that, for a given downward direction and energy, the intensity of unscattered
plus scattered photons falls off exponentially to a good approximation with
depth (exp (-jUj*.)). The unscattered photons also have an exponential
dependence on the depth (exp (-n2x)). The calculation shows that nl is
much smaller than (i . At zero gm/cm2, the flux of scattered photons is
zero; therefore the flux of scattered photons alone has a dependence on the
depth of the form

exp(-ju1x)-exp(-M2x)

At 50 to 60 keV, for example, this results in a maximum flux at about
8 gm/cm2. The intensity of the atmospheric X-rays instead peaks around
1 00 gm/cm2 , so that the two behaviors versus depth are quite different.

It should be noted that the counting rates from residual atmospheres of
15 gm/cm2 and greater are used to derive the depth dependence of
atmospheric X-rays. The peak of the scattered primary radiation lies above
those depths and cannot be included in any simple extrapolation so the n2
coefficient cannot be used. A significant amount of primary scattered
X-rays is still present at 1 5 gm/cm2 ; lower depths are really more appropriate
for the fitting if the jUj coefficient is to be used. The exact form of the
depth dependence of the atmospheric X-rays at small depths is, of course,
debatable.

My main point here was that the scattered celestial X-rays create a bump
at small depths which sits on top of the extrapolation of the atmospheric
X-rays when the standard guesses for the depth dependence are used. (See
Figure II.B-1.)

So all I am saying is that the scattered photons are in there with the ones
that come there unscattered, and you cannot get rid of them by extrapo-
lating in this way.

40 OBSER VA TIONAL DA TA

DISCUSSION

Horstman:

Have I interpreted correctly what you said this morning? (See Chapter II.A.)

Schwartz:

I disagree with exactly what you say, but yes, you are repeating what I said
this morning. (See Chapter II. A.)

Horstman:

My point was that you were stating that you don't need to use the Compton
corrections. All I am saying is that if you arrived at the correct flux at some
altitude, if you have succeeded somehow in subtracting out the atmospheric
contribution, the atmospheric X-rays, then what you have to do with the
remainder is to correct for the multiple Compton scattering.

Schwartz:

That is what I disagree with. I say once an X-ray interacts in the atmosphere,
it is an atmospheric X-ray.

Horstman:

How can you tell? It doesn't have a label on it.

Schwartz:

That is what I am saying. It does not have a label. Once it is Compton-
scattered at a place in the atmosphere, it is exactly as if an electron had
produced it there by bremsstrahlung.

Horstman:

From the calculation, it comes out that, if you look at the vertical flux of
scattered plus unscattered photons, it goes roughly exponentially. The
absorption curve varies with energy, but it is roughly exponential at any
energy.

If you can succeed in separating the celestial from the atmospheric com-
ponent, all you do is use that different absorption coefficient to try to find
the celestial component on top of the atmosphere.

Clark (Session Chairman):

I think that we have pinpointed a problem. Perhaps we should try to
resolve that at coffee.

Chapter III

A. THE MEASUREMENT AND INTERPRETATION OF
THE COSMIC GAMMA-RAY SPECTRUM BETWEEN
0.3 AND 27 MeV AS OBTAINED DURING
THE APOLLO MISSION

L. E. Peterson*

University of California at San Diego

J. I. Trombka*

Goddard Space Flight Center

A. E. Metzger, Jet Propulsion Laboratory;
J. R. Arnold and J. I. Matteson, University of California at San Diego; and
R. C. Reedy, Los Alamos Scientific Research Laboratory

INTRODUCTION

During the transearth portion of the Apollo- 1 5 and -16 missions, data on the
spectrum of the total (diffuse and discrete sources) cosmic 7-ray background
over the 0.3- to 27-MeV range were obtained (Trombka et al., 1973). An
uncollimated 7.0 cm X 7.0 cm cylindrical Nal(Tl) scintillation counter
located on a boom 7.5 m from the Apollo Service Module was used to per-
form the measurement. An analysis of the data obtained on Apollo- 1 5 is
presented here.

A major source of interference in determining the magnitude and shape of
the cosmic 7-ray spectrum can be attributed to the cosmic-ray induced
activation of the Nal(Tl) detector crystal. A Nal(Tl) crystal similar to that
used during the Apollo-15 and -16 missions was flown aboard the Apollo-17
Command Module. This crystal was returned to earth and measurements of
the induced activity were obtained. Preliminary analysis of the results are
now available (Trombka et al., (in press)). Other sources of interference with
respect to the determination of the diffuse 7-ray spectrum have also been
considered. This interference or background was due to sources aboard the

*Gpeakers.

42 OBSER VA TIONAL DA TA

spacecraft and cosmic-ray induced 7-ray emission from the spacecraft and
material surrounding the detector. Attempts have been made to correct the
measured spectrum for these background effects.

An upper limit measure of the 7-ray flux around 0.51 MeV was also
obtained.

INSTRUMENTATION

The Apollo-15 and -16 7-ray spectrometers (Alder and Trombka, 1970)
consisted of a 7.0-cm-diam X 7.0-cm-long Nal(Tl) central detector viewed
by a 7-cm (3 in.) photomultiplier. Except at the photomultiplier end, the
crystal is surrounded by a 1 -cm-thick plastic scintillator shield which detects
charged particles. The plastic scintillator is viewed by a 4-cm (1 .5 in.) photo-
multiplier and has a threshold of about 1 .0 MeV for generating an anticoin-
cidence event when interactions occur in the most optically unfavorable
location. Central detector events with no shield anticoincidence are pulse-
height analyzed into 511 channels and are transmitted at a maximum event
rate of 369 counts/s. The shield rate, the coincidence rate, and the live time
are transmitted every 0.328 s. The spectrometer and associated electronics
are enclosed in a thermal shield and mounted on a boom which could be
extended from one side of the Service Module by an astronaut. The compo-
nents carried on the boom present ~ 5 gm/cm2 averaged over all directions.
The astronaut could fully deploy the detector to 7.6 m from the spacecraft
edge or position it at intermediate distances using stopwatch timing. Further-
more, he could step the high voltage supply or disable the anticoincidence.

RESULTS

Energy Loss Spectra

The data reported here were obtained during portions of the transearth coast
of Apollo- 1 5 from about 2200 hours August 4, to 1 500 hours August 7, 1971 ,
and represent approximately four hours of operation in the extended position.
During this period the earth and moon solid angles were always less than
10'2 sr and in the fully extended boom position, the spacecraft subtended
~0.28 sr. Spectra were obtained with the detector at various boom positions,
with the anticoincidence both on and off, and with the high voltage set to
give several energy ranges up to 27 MeV. Although data were obtained over
a 0.16- to 27-MeV range, the analysis reported here is based on energy
losses > 0.3 MeV. Calibration was obtained with a Hg203 source and by
means of known, easily identifiable spacecraft background 7-ray lines. Since
long periods of time were used to obtain the data and the Command Service
Module (CSM) rotated ~ 3 rph in the ecliptic plane, counting rate anisotropics,
if they exist, were averaged out.

MEASUREMENT OF COSMIC GAMMA-RA Y SPECTR UM 43

Figure III.A-1 shows energy -loss spectra (for several important data modes)
corrected for live time, channel width, and the isotropic detector geometry
factor of 57.5 cm2. Here counts have been summed over channels consist-
ent with the detector-energy resolution which was 8.6 percent at 662 keV.
With the exception of the strong line at 0.5 1 MeV, most of the 7-ray lines
measured in board largely disappear with boom extension, leaving a con-
tinuum extending to 27 MeV, on which is superposed a number of weak
lines. Since the intensity changed only by about a factor of five while the
spacecraft solid angle changed a factor of 20, most of the count rate in the
extended position is not of spacecraft origin. From a detailed analysis of
the rates versus solid angle, we estimate ~ 6.6 X 10"3 c(cm2 • s • MeV)"1 at
2.4 MeV and ~ 1.9 X 10"3 c(cm2 • s • MeV)1 at 5 MeV are due to the
spacecraft. These are 0.1 and 0.2, respectively, of the spectrum with the
boom extended.

The flat energy-loss spectrum of 0.052 c(cm2 • s • MeV)'1 above 5 MeV with
the anticoincidence disabled in the extended position agrees with that value
expected from the shield rate of 450 c/s, from which a cosmic-ray flux of
3.50 (cm2 • s)"1 can be derived. The large ratio of cosmic-ray to photon
energy losses near 27 MeV requires effective charged-particle rejection,
which could not be measured before launch to the required accuracy.
However, preliminary results from an identical experiment on the Apollo-16
in April 1972 confirm the Apollo-15 differential energy loss spectrum below
10 MeV to within approximately 12 percent. We interpret this as indicating
that there were not systematic differences in the behavior of the instruments.

The energy loss spectrum with the anticoincidence enabled, and the boom in
the extended position is compared in Figure III.A-2 with measurements
obtained on other satellites during cislunar space flight. The Nal(Tl)
Apollo-15 and -16 detectors were identical in size to the CsI(Tl) detector
in Ranger-3 (Metzger et al., 1964) both of which differ only slightly from
the Nal(Tl) crystal on the ERS-18 (Vette et al., 1970). The 8 kg mass on
the end of the Apollo-15 boom is nearly the same as that system aboard the
ERS-18, while the Ranger-3 detector carried only ~ 3 kg. Clearly, the
present data are in good agreement with previous measurements below
~2 MeV, but are well below the 3.7- to 6.0-MeV point measured by the
ERS-18, which is apparently erroneous.

Equivalent Photon Spectra

The equivalent photon spectra, Figure III.A-3, have been obtained from the
energy loss spectra in Figures III.A-1 and -2 by using a measured response
"library" and a matrix inversion technique as described by Adler and
Trombka (1970). The 7-ray lines are separated from the continuum by using

OBSER VA TIONAL DA TA

10 r 1 1 — r

~i — i — i — i — I i i i 1 1 1 1 — i — I — i i i i

APOLLO 15
ENERGY LOSS SPECTRUM

DETECTOR INBOARD

0ETECTOR
EXTENDED

Figure III.A-1. Energy loss spectra in the 7.0-cm-diam
by 7.0-cm-long Nal (T1 ) scintillation counter measured
on Apollo-15 during transearth coast. Since the rates
decreased only a factor of about five when the detec-
tor was extended to 7.6 m, while the solid angle sub-
tended by the spacecraft decreased a factor of 20, we
interpret most of the rate in the extended position to
be associated with cosmic 7-rays. The spectrum with
the anticoincidence disabled agrees with that expected
from cosmic-rays passing through the crystal edges.

MEASUREMENT OF COSMIC GAMMA-RA Y SPECTRUM

1 1 1 1 1 — i — i — p

ENERGY LOSS SPECTRA

A RANGER JH
1 1 ERS-18

~1 1 1 1 Ml

-1 1 i Iiiii

APOLLO 15

-l_l 1 I i i i

ENERGY (MeV)

J 1 1 Iiiii

Figure III.A-2. Energy loss spectra are compared directly
with other measurements obtained outside the magneto-
sphere. These data were obtained with counters that
differ only slightly in geometry and materials.

OBSER VA TIONAL DA TA

1 1 1 1 1 1 — i — i — i 1 1 1 1 1 1 — i — i — i 1 1 1 1.

ioV

TRANSEARTH SPECTRA

PHOTON EQUIVALENT

ERS -18 (Uncorrected)

Ap 15 Photon Eq.
Continuum (corr.
for lines]

Ap 15- (Corr for
lines, Spoil
S S/C cont

Ap 15- Finol Spectrum-
(Corr for lines ,
Spoil , SC cont,
8 Atten.)

_i i i i i i 1 1 1

t.O 10

ENERGY (MeV)

100

Figure 1 1 1. A- 3. Equivalent photon spectra derived
from the Apollo-15 are shown at various stages of
data correction. First, all components due to dis-
crete 7-ray lines are removed, then the spacecraft con-
tinuum contribution and an estimate of energy losses
due to spallation nuclei are subtracted. The final
result contains a correction for absorption of local
material, assuming all energy losses at this stage are
due to an external isotropic 7-ray flux.

MEASUREMENT OF COSMIC GAMMA-RA Y SPECTR UM 47

an iterative procedure, (Trombka et al., 1970; Reedy et al., 1973). Here the
pulse-height spectrum is transformed to photon space where lines appear as
discontinuities, which can be subtracted by requiring the remaining con-
tinuum to vary slowly with energy. This procedure results in the removal of
2.5 c/s over the 0.6- to 3.5-MeV range due to lines or about 17 percent of
the energy loss spectrum and leaves a smooth equivalent photon continuum
shown in Figure III.A-3.

The following are a few comments on the determination of the measured
response library. The shape and detection efficiency of these library func-
tions strongly depends on the angular distribution of the incident 7-ray
flux. To illustrate this point, the detection probability (intrinsic efficiency)
for a 7.0-cm X 7.0-cm cylindrical Nal(Tl) detector is given in Figure III.A-4
as a function of energy for two different cases: a parallel beam incident on
the face of the crystal (the crystal axis is parallel to the beam) and an iso-
tropic distribution of 7-rays. As can be seen, there is significant difference
in the detection efficiencies over much of the energy region of interest.
The shapes of the pulse-height spectra do not seem to change quite as
radically as a function of the angular distribution of the incident flux. In
order to transform from measurement or energy loss spectra to photon
spectra, efforts were made to eliminate all background components in order
to isolate the energy loss spectra characteristic of the diffuse component.
The assumption was then made that this component was isotropic and the
transformation was then performed using an isotropic-type response library.
From a comparison of our experimental work (Trombka et al., 197 1) with
Monte Carlo calculations (Berger and Seltzer, 1972), we found that the
response library function can be calculated theoretically for any energy and
geometry needed in the analysis.

Discrete Line Spectra

The discrete line spectrum in the measured cosmic-ray spectrum can be
mainly attributed to natural radioactivity aboard the spacecraft (K-40 and
Th), proton- and neutron-induced activation in the spacecraft and materials
surrounding the detector, and activity induced in the detector itself. Using
the technique considered in the section "Equivalent Photon Spectra," the
continuous portion of the energy loss spectrum was determined, and the
continuous spectrum was subtracted from the uncorrected energy spectrum.
In this way, the energy loss spectrum characteristic of discrete lines is
determined. The results are shown in Figure III.A-5. Identification of
certain lines are also indicated. We believe that the following lines can be
identified: (1) the 0.51-MeV line due to positron annihilation; (2) the 0.63-
and 0.69-MeV lines due to proton-induced activation of the crystal pro-
ducing 124I and 126I, the 1.47-MeV characteristics of 40K; and (3) the

OBSER VA TIONAL DA TA

INTRINSIC EFFICIENCIES
2 3/4" x 2 3/4" NAI

a PARALLEL BEAM
« ISOTROPIC 4n

-i i i '''i

0.1

1.0

10.0

100

E (Mev)

Figure III.A-4. Intrinsic efficiencies as a function of energy for
a 7-cm by 7-cm Nal(T1) detector. Both parallel beam and
isotropic 7-ray fluxes are considered.

2.6-MeV line of thorium. Other lines due to thorium and (n, 7) and (n, n , y)
reactions on Mg, H, Al, 0, and Na can also be presented.

Spallation Correction

Fishman (1972) has suggested that radioactive spallation nuclei produced by
cosmic-ray interactions in the scintillation crystals might account for a large
fraction of the counting rate measured in the 1 - to 3-MeV region. Although
a direct measurement of this effect in the cosmic-ray flux is difficult and has
not been accomplished, calculations and laboratory measurements by Fishman
(1972) and Dyer (private communication) have indicated the spectra shape
and approximate magnitude of the energy loss spectrum. We have attempted
to correct the spectra of Figure III.A-3 for this effect by subtracting from the
equivalent energy loss spectrum a spallation model spectrum whose normal-

MEASUREMENT OF COSMIC GAMMA-RA Y SPECTR UM

PULSE HEIGHT SPECTRUM

I04

: ' 1

i j i | i | i | i | i | i | i | i | i | i | i | i | i | i | i | i | i | i | i | i | i | i

1 ;

5.1 Mev » CSMLINES

i »~20Kev/CHANNEL

I03

.63 Mev
" f. 69 Mev 1.28 Mev
/ /- 1.37 Mev

. V r,Mev / l47Mev
" " " \ jrt. y-i.67 Mev

""„""/ r- 2.2 Mev
." > "\"« ." *«".V» VSsfli /- 2.6 Mev

H
Z
3

I02

V»'»'v "-

»" ■" ", .

> » " "

I01

/ '*■/.":

« -1

inO

-tO.

j-L. 1 . I . 1 ,i,i,l,l,i,l,l,l,i,l,l,l,iM, l„rt„

0 20 40 60 80 100 120 140 160 180 200 220 240
10 30 50 70 90 110 130 150 170 190 210 230 250

CHANNEL
Figure III.A-5. Discrete-line energy loss spectrum from Apollo-15.

ization was a free parameter. Since spallation contributes mostly to the
energy losses in the 0.6- to 3-MeV range, the normalization was determined,
rather arbitrarily, by the criterion that the resultant photon spectrum be
relatively smooth. This was found to occur when a spallation spectrum
shape, based on the work of Fishman (1972), Dyer and Morfill (1971 ;
private communication) but approximately half their calculated intensity,
was subtracted out. As shown in Table III.A-1, this results in removal of
about 1 6 percent of the energy loss spectrum in the 0.6- to 3.5-MeV range
and a negligible amount at higher energies. Subtracting a much larger
spallation component, such as the full Dyer and Morfill value, would give no
energy loss spectrum in the 1- to 2-MeV range, while still requiring an
external photon component above 3 MeV, which is not physically possible.
Although there seems no doubt that a spallation energy loss contribution
exists, its spectral shape and intensity are only approximately known.

The spallation components are always subtracted out in energy loss space.
In an attempt to obtain experimental data on the extent of the proton-
induced activity, a Nal(Tl) crystal was flown aboard Apollo-17. The
crystal assembly was physically identical to that flown aboard Apollo-15 and

50 OBSER VA TIONAL DA TA

Table III.A-1

Composition of Apollo- 1 5 Energy- Loss Spectrum
(transearth coast, deployed position)

Energy Range

Component 0.6-3.5 MeV 3.5-9.0 MeV

7-ray lines (percent) 15.9 3.7

Spallation in Nal crystal (percent) 15.8 0.5

Spacecraft continuum (percent) 10.2 21.7

Cosmic upper limit (percent) 58.1 74.1

Total (percent) 100.0 100.0

-16 (Adler and Trombka, 1970). The assembly aboard the Apollo- 17 CSM
did not include the photomultiplier, the proton anticoincidence mantle, and
the thermal shield. The detector was a 7-cm X 7-cm right-cylindrical crystal.
A glass plate was optically sealed to the crystal. MgO was used as the optical
reflector inside the crystal assembly. This type of assembly permitted the
crystal to be hermetically sealed and allowed for a simple procedure for
optically coupling the crystal assembly to a photomultiplier tube after flight.
The crystal and reflector were enclosed in a steel jacket. An identical second
crystal assembly which was not flown was used as a control throughout the
measurement program. After splashdown, the flight (that is, activated)
crystal was returned to the recovery ship and optically mounted on a photo-
multiplier tube, and pulse-height spectra were obtained. The activated
crystal was counted in a large, steel, low-level shield. The crystal counting
started about one and one-half hours after the Command Module reentered
the earth's atmosphere. Before splashdown the control (unactivated) crystal
was optically sealed to a photomultiplier tube, and the background was
determined in the steel shield. The same photomultiplier tube was used to
count the activated and control crystal assemblies. After 30 hours of
counting aboard the recovery ship, the detector was flown back to the Oak
Ridge National Laboratory (ORNL) where measurements were continued.
This permitted the observation of the decay of the longer-lived induced
activities. Direct measurements of the induced activities were made again by
optically sealing a photomultiplier tube to the activated crystal. Indirect
measurements using both Ge(Li) detectors and a large scintillation 4ir detector
in a low-level counting system at ORNL (Eldridge et al., 1973) were performed
in order to determine the spectral distribution and intensity of the emitted
radiations. The 4n scintillation counter is divided into halves. Each half can
be operated so as to require that there be coincidence events in both halves

MEASUREMENT OF COSMIC GAMMA-RA Y SPECTR UM

before an event is analyzed and recorded (coincidence spectra) or the halves
can be operated without the coincidence requirement, and events independent
of their coincidence can be analyzed and recorded (singles spectra).

To date it has been possible to obtain qualitative identification of the
following nuclear species: 22Na (2.6 yr), 24Na (15 hr), 123I (13 hr),
124I (4 days), 126I (13 days), 128I (25 min), and 127Xe (34 days). After
suitable calibrations, quantitative concentrations of these radionuclides will
be obtained. The present results indicate that the induced activity observed
after recovery can be attributed mainly to species with half-lives of about
half a day and longer. Decay products with shorter half-lives do not make a
large contribution to the post-recovery integral count rate. This is not to
imply that there are no short half -life components. In fact, the line at
0.44 MeV is characteristic of 128I. There are a few more regions with
relatively short half-lives (in order of tens of minutes) that have not as yet
been identified.

Figure III.A-6 shows the pulse-height spectrum obtained during the first
hour and a half of counting after recovery. The spectrum has been corrected
for background by subtracting the measurements obtained with the control
crystal. Peak energies for the nuclides presently identified are indicated.
The peak positions of 123I, 124I, 126I, and 128I are displaced 27 keV due
to X-ray emission and absorption in the crystal after electron capture.

,. 0.389
* MeV I*

0.1 59 MeV '"I

,0.44MeV ,al
/ .0.51 MeV (MANY SOURCES)
' //0.603 MeV ,24l
//.0.666 MeV ,26l
\/*\ *^0.723MeV '"I ,1.27MeV»Na

>. <jS—-0.754MeV ,26l />1 .36MeV "No

\i /y 1.46MeV «K

"H, // f POSSIBLE EXCESS

•" ^Vfc,. /£,?• — I _ INACTIVATED

CALCULATED SPALLATION COUNTINUUM
(REFERENCE 5)

APOLLO 17 DATA

2.62 MeV

THORIUM LINE

POSSIBLY EXCESS

IN ACTIVATED

CRYSTAL

0.2 0.4 0.6 0.8 1.0 1.2 1.4 1.6 I.

2.0 2.2 2.4 2.6 2.8 3.0 3.2 3.4 3.6 3.8 4.0

MeV

Figure III.A-6. Proton-induced activity in 7-cm by 7-cm Nal(T1) crystal
1-1/2 hr after reentry. The background has been subtracted. Counting time
was 1800 s. The spectrum measurement started 1-1/2 hr after reentry. The
spectrum was obtained by direct internal counting of the activated crystal.

52 OBSER VA TIONAL DA TA

Measurements of the flight and control crystal carried out at the low-level
counting laboratory at the Oak Ridge National Laboratory (ORNL) prior to
flight indicated the K and Th content of the flight crystal to be slightly
higher than that for the control crystal. Thus, one would expect some indi-
cation of these elements after background subtraction. The energy identifi-
cation for 124I, I, and 24Na have been verified by measurements made
with the Ge(Li) detector and in the low-level counting system. Both energy
and half-life information have been used to determine the presence of these
nuclear species. The 123I and 128I were identified by use of the spectra
obtained on board the carrier from both energy and half-life determinations.
22 Na has been tentatively identified based on a preliminary analysis of the
data obtained by the coincidence measurements in the low-level counting
facility. 127Xe presence has been determined by the identification of
energy lines at 0.172 MeV, 0.203 MeV, and 0.375 MeV using the Ge(Li)
detector.

One factor requiring consideration was the difference in the environment
during the Apollo-15 and -16 missions compared with Apollo-17 mission.
First, the crystals aboard Apollo-15 and -16 were stowed in the Service
Module and extended approximately 762 cm away from the vehicle for
short periods of time, whereas the Apollo-17 crystal was stowed in the
Command Module for the total flight time. Thus, there was a difference in
the mass around the crystal which might cause a difference in the secondary
proton and neutron flux in the region of the stowed crystals. Secondly,
the exposure profile of the primary flux both in time and spectral distri-
butions were different. The Apollo-17 crystal passed through the near-
earth trapped proton flux twice before measurements, while the Apollo-15
and -16 detectors had passed through the trapped belts only once before
measurement. The Apollo-15 measurement of diffuse 7-ray spectrum was
made about 250 hr after lift-off, while the Apollo-17 measurements were
made some 305 hr after lift-off. It has not as yet been determined how
significant these differences are in terms of trying to infer the magnitude
of the proton-induced activity in the Apollo-15 and -16 detectors from
the Apollo-17 measurements.

The shape of the cosmic-ray-induced 7-ray pulse-height spectrum can be
divided into two parts: the discrete-line spectrum and the continuous
spectrum. The discrete-line pulse spectrum for activated nuclear species
in the crystal is produced by monoenergetic 7-rays emitted after electron
capture. The continuum for such nuclear species is produced by electrons,
positrons, positron annihilation, and 7-rays (other than those emitted after
electron capture) interacting in the crystal. If the material surrounding
the crystal is radioactive (for example, 24Na, Th, or 40K) then mono-

MEASUREMENT OF COSMIC GAMMA-RA Y SPECTR UM 53

energetic 7-rays independent of the mode of decay can be seen in the
crystal as a discrete-line pulse-height spectrum. In Figure III.A-6, the
discrete lines are indicated and the continuous distribution can be seen
underneath. The actual energy position should be moved ~ 27 keV up in
energy due to the summing of iodine-K X-ray line with the 7-ray line after
K capture.

In the Apollo-15 transearth spectrum (Trombka et al., 1973), the 124I
0.606-MeV and the 126I 0.66-MeV lines can be identified. It has been
calculated that the integrated count rate in this region above the continuum
for Apollo-15 is half of that observed in the same region above the
continuum for the Apollo-17 mission. This difference cannot be attributed
to the difference in exposure time alone. Thus, the difference in local
mass and the passage through the near-earth trapped radiation belts a
second time may be the cause of this increase.

In Figure III.A-6 the magnitude of the continuum and associated error
as predicted by Fishman (1973) is compared with the Apollo-17 measure-
ment taken aboard the recovery ship. The magnitude of the continuum
inferred from the Apollo-15 data (Trombka et al., 1973) is also shown.
Its magnitude is consistent with the Apollo-17 results if it is considered
that the discrete-line magnitude for 124I and 126I is down by a factor of
two. This also assumes that the shorter half-lived nuclides and the prompt
7-ray emission is small compared to the longer half-lived emitters. Cal-
culations (Dyer and Morfill, 1971) indicate that short half-lives may be
quite important.

Spacecraft Continuum

The following procedures were used to determine the magnitude of the
spacecraft continuum.

Spectra were obtained with the detector positioned at 183 cm, 244 cm,
457 cm and 762 cm away from the spacecraft. An effective solid angle
for each position was calculated for these positions. The discrete-line
spallation backgrounds discussed in sections "Equivalent Photon Spectra"
and "Discrete Line Spectra" were then subtracted from the energy loss
spectrum at 183 cm and 762 cm. It was then assumed that the 183-cm
spectrum characterized the energy loss spectrum of the continuous 7-rays
spectrum emitted from the spacecraft. The intensity at 183 cm is reduced
by the ratio of the effective solid angle at 762 cm to the effective solid
angle at 183 cm. This then is a first estimate of the contribution of the
spacecraft continuum at 762 cm. The spacecraft continuum contribution
is then subtracted from the residual energy loss spectrum at 762 cm and a
first estimate of the energy loss spectrum due to the diffuse component is

54 OBSER VA TIONAL DA TA

obtained. It is now assumed that the diffuse energy loss spectrum does
not depend on the distances of the detector from the spacecraft (that is,
the spacecraft occultation is ignored) and this first approximation is sub-
tracted from the energy loss spectrum at 183 cm. A second approximation
of the continuous energy-loss spectrum from the spacecraft at 183 cm is
obtained. This new continuous energy-loss spectrum is corrected for change
in solid angle to obtain its contribution at 762 cm (25 ft) and then sub-
tracted from the original residual energy-loss spectrum at 762 cm (25 ft)
in order to obtain a second approximation of the diffuse energy-loss spec-
trum. The procedure as described above is continued for another two
iterations, and it was found that the shape of the diffuse energy-loss spec-
trum did not change significantly between the last two iterations. After
the last iteration, the energy-loss spectrum was then converted to photon
spectrum. The transformation was accomplished using library functions
and efficiencies characteristics of isotropic flux distributions.

Cosmic Photon Spectrum

The photon spectrum incident on the central detector, shown in Figure
III.A-3 as a dashed line, has also been corrected for the various inter-
ferences discussed in the preceding sub-sections. The contribution of the
various components over the 0.6 to 3.5 MeV and the 3.5 to 9.0 MeV
ranges are summarized in Table III.A-1. Despite the many corrections,
about 50 to 75 percent of the energy losses cannot be accounted for by
presently understood local processes and therefore must originate externally.
Obtaining the photon spectrum incident isotropically on the spectrometer
requires a correction for local matter. Taking this to be equivalent to a
uniform shell 5.0 gm/cm2 thick of Al surrounding the Nal crystal, and
correcting for absorption, but not scattering, results in the final photon
spectrum shown in Figure III.A-3. We have assumed the photon continuum
extends as E"20 above 27 MeV; however, the result is rather independent
of this shape.

Systematic errors, which are difficult to estimate, completely dominate the
statistical uncertainties in this analysis. Correcting for the spacecraft lines
can be done to high precision. The effective solid angle for continuum
production in the spacecraft may be less certain. No correction has been
made for production in local material, which is believed to be small
(Vette et al., 1970). We estimate the equivalent photon spectrum, before
correction for spallation, to be accurate to about ± 20 percent. The
spallation correction cannot be much larger than that indicated in Figure
III.A-3. Correcting for absorption, but not scattering, results in an upper
limit to the external flux.

MEASUREMENT OF COSMIC GAMMA-RA Y SPECTR UM 55

These results can be compared to those of others who have presented
spectra at various stages of correction. The Apollo-15 photon equivalent
continuum is considerably below that determined from ERS-18, which had
no corrections for 7-ray lines, effects of local material, or spallation and
which apparently had an instrumental malfunction at higher energies. The
final photon Apollo-15 spectrum is compared with balloon and low altitude
satellite work (Golenetskii et al., 1971 ; Vedrenne et al., 1971 ; Damle et al.,
1971) in Figure III.A-7. The result of the reference is considerably above
the other work and is therefore not shown. Although the low-latitude
observations should not require a significant correction for spallation, they
do require an altitude-and latitude-dependent model to correct for cosmic-
ray produced 7-rays and, in some cases, an additional large correction for
counter efficiency.

The new results, in addition to being in reasonable agreement with the more
recent work above 1 MeV, also agree with data near 100 keV (Pal, 1973) when
extrapolated as an E*2 power law. Furthermore, the Apollo spectrum is
consistent with new data on the diffuse component near 30 MeV (Mayer-
Hasselwander et al., 1972; Share et al., 1973). Figure III.A-7 shows some of
these results, as well as at 100 MeV obtained from the OSO-3 (Kraushaar
et al., 1972, preprint).

Also shown in Figure III.A-7 is a single power law which has been suggested
(Pal, 1973) as capable of representing the total cosmic 7-ray spectrum
between approximately 0.02 and 1.0 MeV. It is clear that the derived
Apollo-1 5 spectrum is well above this extrapolation, and even though we
interpret our result as an upper limit, we do not believe that the remaining
small corrections and uncertainties can reduce the final cosmic spectrum to
the extrapolated value.

DISCUSSION

Assuming that the 7-ray fluxes are of extragalactic origin (Stecker et al., 1971)
a number of workers have attempted to account for the spectra shown in
Figure III.A-7. Compton scattering of electrons leaking from radiogalaxies
(Brecher and Morrison, 1969), red-shifted 7-rays from n° decays produced
by cosmic-ray collisions at an early epoch of the expanding universe
(Stecker, 1971), nuclear 7-rays from supernovae in distant galaxies (Clayton
and Silk, 1969), intergalactic electron bremsstrahlung (Arons et al., 1971;
Stecker and Morgan, 1972; Stecker et al., 1971) and matter-antimatter
annihilation (Stecker et al., 1971) have all been suggested. Vette et al. (1970),
in attempting to account for the ERS-18 data, fitted a model in which a n°-
decay component produced at an epoch with a red shift « 70 was super-im-
posed on a Compton scattering X-ray background. Based on the present data,

OBSER VA TIONAL DA TA

1.0

10 ;

10 :r-

^ -4

° 10

10 E

-6
10 E

I I 1 1 1 1 I I I I 1 1 1 1 1 1 I I I | HIM 1 — I I | ||

COSMIC r-RAY SPECTRA

Golenetski.etol.
Vedrenne.etal.

APOLLO 15

Mayer -
Hasselwander.et al.

osom

i i i I mil I II lllll

1.0

10
ENERGY (MeV)

100

1000

Figure III.A-7. The cosmic-photon spectrum derived
from the Apollo-15 data agrees with previous results
below 1 MeV but is well below that determined from
the ERS-18 at higher energies. Limits derived from
balloon and low-altitude satellite work, despite large
corrections for efficiency and cosmic-ray produced
7-rays, are in agreement with the Apollo results.

MEASUREMENT OF COSMIC GAMMA-RA Y SPECTR UM 57

the intensity of the flux required at very early epochs is reduced somewhat.
The final spectrum of Apollo- 15 does require an additional component
above a simple power law. A discussion of the theoretical consequences of
these results is given by Stecker elsewhere in these proceedings (see Chap-
ter IX. A).

The analysis process used here subtracts out all discrete 7-ray lines and pro-
duces a smooth continuum, as presented in Figure III.A-3. Discrete 7-rays
of cosmic origin, if they exist, would therefore be removed along with known
spacecraft and spallation contributions. Only considerable further analysis
can separate these components' and place valid limits on possible cosmic
components.

The 7-ray line near 0.51 MeV has an intensity after correction for spacecraft
production and local absorption estimated to be 3.0 ± 1.5 X 10"2 photons
(cm2 • s)"1 . The uncertainty is an estimate of the effect of systematic errors
in the correction for weak 7-ray features near this energy and for detector
efficiency and absorption. The 0.51 MeV 7-ray measured on Apollo-15
cannot originate in the spacecraft because this component decreases less
rapidly with spacecraft solid angle than the continuum. The intensity of
the line seems inconsistent with upper limits on the cosmic flux at 0.51 MeV
of < 10"2 photons (cm2 • s)"1 obtained from balloon measurements (Chupp
et al., 1970) and on Ranger-3 (Metzger et al., 1964). Since Ranger-3, which
also measured in interplanetary space, had considerably less matter locally
to the detector, it may be possible to attribute the flux to annihilation of
positrons produced by cosmic-rays or spallation j3+- decays in the local mass.
It is also possible that low-energy positrons of either solar or cosmic origin
with a flux of ~ 10"2 (cm2 • s)"1 could stop and annihilate in the inert matter
surrounding the detector. Such a mechanism has been suggested by
Stephens (private communication) and is in fact consistent with the inter-
planetary medium flux of 2 X 10"2 positrons (cm2 • s)"1 at approximately
2 MeV reported in Cline and Hones (1970). Haymes (Johnson, Harnden,
and Haymes, 1972) has reported a 7-ray line at ~ 470 keV whose intensity
is 2 X 10"3 photons (cm2 • s)"1 originating from the galactic center. The
7-ray line measured on Apollo-15 is definitely at 0.511 ± 0.012 MeV, and
the 2 a upper limit to a 7-ray at 0.47 MeV is ~ 2 X 1 0"3 photons (cm2 • s)"1 ,
based on the analysis of 4 hr of data.

ACKNOWLEDGMENTS

The work described in this paper was carried out in part under NASA
Contract No. NAS7-100 at the Jet Propulsion Laboratory, California Institute
of Technology, in part under NASA Contract No. NAS9-10070 at the Univer-
sity of California, San Diego, and in part under USAEC Contract W-7405-
Eng-36 at the Los Alamos Scientific Laboratory of the University of
California.

58 OBSER VA TIONAL DA TA

REFERENCES

Adler, I., and J. I. Trombka, 1970, Physics and Chemistry in Space,
3, J. G. Roederer and J. Zahringer, eds., Springe r-Verlag, New York.

Arons, J., R. McCray, and J. Silk, 1971, Astrophys. J., 170, p. 431.

Berger, M. J., and S. M. Seltzer, \912,Nuc. Inst, and Methods, 104, p. 317.

Brecher, K., and P. Morrison, 1969, Phys. Rev. Letters, 23, p. L802.

Chupp, E. L., D. J. Forrest, A. A. Sarkady, and P. J. Lavakare, 1970,
Planetary and Space Science, 18, p. 939.

Clayton, D. D., and J. Silk, \9 69, Astrophys. J. Letters, 158, p. L43.

Cline, T. L., and E. W. Hones, 1970, ACTA Phys., 29, Supp 1, p. 159.

Damle, S. V., and R. R. Daniel, G. Joseph, and P. J. Lavakare, 1971,
Astrophys. and Space Sci , 14, p. 473.

Dyer, C. S., and G. E. Morfill, 1971, Astrophys. and Space Sci , 14, p. 243.

Eldridge, J. S., G. D. O'Kelley, K. J. North cutt, and E. Schonfeld, 1973,
Nuc. Inst, and Methods, in press.

Fishman, G. J., 1972, Astrophys. J., 171, p. 163.

Golenetskii, S. V., E. P. Mazets, V. N. Il'inskii, R. L. Aptekar, M. M. Bredov,
Uy. A. Gur'yan, and V. N. Panov, 1971, Astrophys. J. Letters, 9, p. L69.

Johnson, W. N., Ill, F. R. Ham den, Jr., and R. C. Haymes, 1972,
Astrophys. J. Letters, 172, p. LI.

Mayer- Hasselwander, H. A., E. Pfeffermann, K. Pinkau, H. Rothermel, and
M. Sommer, 1972, Astrophys. J. Letters, 175, p. L23.

Metzger, A. E., E. C. Anderson, M. A. Van Dilla, and J. R. Arnold, 1964,
Nature, 204, p. 766.

Pal, Y., 1973, X-Ray and Gamma-Ray Astronomy, Proceedings of IAU
Symposium No. 55, (Madrid), H. Bradt and R. Giacconi, eds., D. Reidel,
Dordrecht, Holland.

Reedy, R. C, J. R. Arnold, and J. I. Trombka, 1973, /. Geophys. Res. ,
in press.

Share, G. H., R. L. Kinzer, and N. Seeman, 1973, X-Ray and Gamma-Ray
Astronomy, Proceedings of IAU Symposium No. 55, (Madrid),
H. Bradt and R. Giacconi, eds., D. Reidel, Dordrecht, Holland.

Stecker, F. W., 1971, Nature, 229, p. 105.

MEASUREMENT OF COSMIC GAMMA-RA Y SPECTR UM 59

Stecker, F. W., 1971, Nature, 229, p. 105.

Stecker, F. W., and D. L. Morgan, Jr., 1972, Astrophys. J., 171, p. 201.

Stecker, F. W., D. L. Morgan, Jr., and J. Bredekamp, 1971, Phys. Rev.
Letters, 27, p. LI 469.

Stecker, F. W., J. I. Vette, and J. I. Trombka, 1971, Nature, 231, p. 122.

Trombka, J. I., E. Eller, G. A. Osswald, M. J. Berger, and S. M. Seltzer, 1971,
USAEC Report Conf. - T10402, III, p. 111-43.

Trombka, J. I., A. E. Metzger, J. R. Arnold, J. L. Matteson, R. C. Reedy,
and L E. Peterson, 1973, Astrophys. J., 181, p. 737-746.

Trombka, J. I., R. L. Schmadebeck, M. Bielefeld, G. D. O'Kelley,

J. S. Eldredge, K. J. Northcutt, A. E. Metzger, E. Schonfeld, L. E. Peterson,
J. R. Arnold, and R. C. Reedy, Apollo- 17 Preliminary Science Report,
in press.

Trombka, J. I., F. Senftle, and R. Schmadebeck, 1970, Nuc. Inst, and
Methods, 87, p. 37.

Vedrenne, G., F. Albernhe, I. Martin, and R. Talon, \91\,Astron. and
Astrophys. , 15, p. 50.

Vette, J. I., D. Gruber, J. L. Matteson, and L. E. Peterson, 1970,
Astrophys. J. Letters, 160, p. L161.

B. INDUCED RADIOACTIVITY CONTRIBUTIONS TO
DIFFUSE GAMMA-RAY MEASUREMENTS

G. J. Fishman*

Teledyne Brown Engineering

The importance of the effects of cosmic-ray-induced radioactivity on diffuse
7-ray measurements has recently become apparent (Dyer and Morfill, 1971;
Golenetskii, 1971 ; Fishman, 1972a). In view of the new Apollo-15 results,
it is believed that a review of the physical processes involved and the derivation
of the corrections due to induced radioactivity would be in order for this
Symposium.

Induced radioactivity by protons in 7-ray detectors was first observed on the
OSO-1 7-ray detector (Peterson, 1965) and subsequently observed by other
detectors placed in earth orbit. The increased background and dead-time due
to activation reduced the sensitivity of these experiments but discrete source
observations were still possible due to the directionality of the detectors.
However, observations of the diffuse 7-ray background are difficult due to
the isotropy of the source and the various sources of background radiation
which are impossible to eliminate completely. Therefore, all extraneous
sources of background must be known and accurately accounted for in order
to derive the diffuse cosmic component.

The most reliable measurements of the diffuse component in the MeV energy
range have been made by detectors aboard spacecraft placed well outside the
trapped radiation and with small amounts of local matter. One component
of the background is the decay of radioactive spallation products formed
when primary cosmic rays interact with the scintillation detector crystal.
The counts produced by these spallation products within the detector are
unaccompanied by anticoincidence events and are otherwise indistinguishable
from counts produced by external 7-rays. The counting rate induced in a
Nal(Tl) crystal with a mass thickness of 25 gm/cm2 is

F =0.15X F (III.B-1)

c p v '

* Speaker.

62 OBSER VA TIONAL DA TA

where F = 3/cm2 • s is the nuclear active particle flux (E > 100 MeV). The
above relation is derived from a statistical treatment of the decay characteris-
tics of many spallation products and the total nuclear interaction cross section
(Fishman, 1972a). The estimated uncertainty of Equation (III.B-1) is 30 per-
cent.

The cross sections for the formation of individual spallation products are cal-
culated from semi-empirical formulae first derived by Rudstam (1966) and
recently modified and rendered more accurate by Silberberg and Tsao (1973).
In a Nal(Tl) crystal, interactions with iodine will account for over 80 percent
of the radioactive products. Table III.B-1 gives the cross section for producing
radioactive spallation products from I127 at 200 MeV, 800 MeV, and 3000
MeV. These data were provided by the NRL group (Silberberg, Tsao, and
Shapiro, private communication). Above 3000 MeV, there is no detectable
change in cross sections with bombarding energy. Using these data, it is
estimated that 70 percent of the total inelastic cross section of I127 (1260 mb)
will result in the formation of radioactive products. This is indicated at the
bottom of Table III.B-1 where the sum of the cross sections for the formation
of radioactive spallation products is assumed to be 880 mb at each energy.
Naturally, the higher bombarding energies tend to remove more nucleons from
the target nucleus. This is illustrated in Figure III.B-1 where all products with
formation cross sections greater than 10 mb from I127 are shown at various
energies.

Although the total counting rate due to induced radioactivity is fairly well
known from Equation (III.B-1), the spectrum that the counts will produce in
a detector crystal is difficult to deduce for a variety of reasons. At the higher
energies, representative of primary cosmic rays, several hundred products are
formed and would need to be considered for an accurate calculation of the
energy-loss spectrum. The decay branching ratios are not well known for
many of these products, and the cross section for formation in many cases
may be in error by as much as a factor of 2. Also, many of the nuclei have
one or more long-lived isomeric states, and there is no means to determine in
which state the spallation product will be formed. For these reasons, it is
necessary to assume a spectral shape on the basis of other data. It is assumed
that the decay spectrum would resemble the 7-ray spectrum from the decay
of a large number of mixed fission products since the atomic number and the
nuclear energy level spacings of these products are similar to that of the
iodine spallation products. This spectrum has an exponential form:

— ccexp(E/Ee) (III.B-2)

INDUCED RADIOACTIVITY CONTRIBUTIONS

i—

•—1

•■^

<L>

-i->

CTJ

-a

o
o
o
m

O
O

on t3

° £

a> o

<Ncoa\c-~-*cr\0^ocNioouoO'<t

•a -a ^ x
ro O i . co

■o^ £ f« -^ 2

m ^h (N — i ^h t-i

VO On

vou-iinTj-mmcN^Hr-^oooNO

Mm,H*hn_H-hH«&m.H-h

O O — <00-hO-^OCNO^O

^M^nTtn^i/i^ifi^hM

'-^CNCNro'^-^tin'vOvot-^r^cooo

rM VO

—

1—

CT3

J-H

OBSER VA TIONAL DA TA

where Ee is the e-folding energy. The value of E = 0.9 MeV used in the
previous paper (Fishman, 1972a) was taken from the data compiled by
Goldstein (1959). A more recent measurement of the 7-ray spectrum
from the spontaneous fission of U238 has been obtained (Sobel et al.,
1973). Their work also shows the spectra well fitted by an exponential
spectrum up to 20 MeV but with a higher e-folding energy, 1.41 MeV.
A measurement of the spectral shape produced by long-lived Nal
spallation products also yielded an exponential spectrum up to 3 MeV
with an e-folding energy of 0.6 MeV (Fishman, 1972a, b). This spectrum

'4*

1 122

I 123 "

I 124'

I 125"

602d

I 126'-

?Z6h(l32^

TO* ' -up

I 127"
no

Te 1 19"

aoas

Tef2l

Tel22

2.44

™ TjT

T«I24

12b"

Tel26

I&7I

Sb 117

Sb (18

Sb 119

Sbl20 •

SbCI"

57,25

Sbl22'-

Sbl23"

42.75

Sbl24''
2im ea 2d

Sb 125"'

Snll6

Snll7"

ScllS

24 03

SnllS-"

25O0 056

Snl20

Snl2l

Snl22

472

Snl23

1250 40m

Snl24
594

In 114"

™ " 9»

In 117"

In 1 B "

In 119

In 120

Cd 113"*
ri,\ aX

0.3 GeV

Sbll7 I Sb 118

TZW

Agl09<

40J 48* ?6O0 E4

n»«o" J' ".&■■■■

Tcalff1

Aq III'
1

In 114

Sn 116 I Snll7b

rjoW

i-ji,=
•"■>"-.

T«t20

Sn 118 I Snll9"

■am?'

TH2T

4 , .,

Snl23

3d 40n

4< '

Sbl24'- Sbl25"

0.6 GeV

Figure III.B-1. Iodine spallation products produced at various proton bom-
barding energies. All products shown have a production cross section greater
than 10 mb. Although more products are formed at higher energies, the total
cross section remains nearly constant, ~1260 mb. The segments shown are
from the Chart of Nuclides (Goldman, 1966).

INDUCED RADIOACTIVITY CONTRIBUTIONS

OBSER VA TIONAL DA TA

-J: I

if" 3 S

£ 18 ',

>
CD

*8s

-i

Ss - I

?< 1 5

8* f 1

lit"

~ Si-"''

^»aS

fife 5 8

§8 I |

sigj

— 2 i I

p%|

S -3S

|3 £

8|:I

ff>

=3 J 1

•s*";."

» - I

2s; J.°L

-*.-.?;

*9 f £

5s i "

<£^ . *

ss 8

|§ff

8.-;r

|l ■

Mi: [

*s 5 S
3* § ;

8|?|

E|"

M •>

2. 5 |

CD .
t]S - a

£ ; :

fife

S- = 5

o: t- ^

S. »

S- ! -

5|J

qb'si 3-

Q ~;r_' 3

o?" . s

'Sss2- =

OR J

2 '. s

OS i

o5""i

19 !'

gjgjj

--4

ro S

< 2 y

Eg-S»s

°s I
j2 a

h?r

""** 3

C£l '

IT • .. ■•

S.j:

S3 I

019.

■

I
i

44 .

Mi j

-5 ||

i
5g

883

M,r '!

INDUCED RADIO A CTIVITY CONTRIBUTIONS 6 7

is thought to be softer than the actual spectrum of induced radioactivity
occuring in diffuse 7-ray measurements because: (1) the measured spectrum
did not include the contribution of very short-lived products which tend to
have a harder spectrum, being further from stability (this trend is observed
in the measured spectra) and (2) the bombarding proton energy was 600 MeV,
considerably less than the average cosmic-ray energy producing activation.
The three spectra mentioned above are compared in Figure III.B-2. The
assumed exponential spectra were normalized to the total induced counting
rate of Equation (III.B-1). The experimental data were corrected to include
the short-lived products but otherwise normalized. The three spectra agree
to within a factor of three up to a few MeV and diverge thereafter. The
exponential drop-off at high energies is expected on theoretical grounds;
spallation products with high excitation energies will decay by prompt
7-emission and particle emission rather than by delayed j3-decay or internal
decay. In addition to a continuous spectrum, the true spallation product
spectrum is expected to have line features superimposed due mainly to elec-
tron capture decays of several iodine isotopes listed in Table III.B-1.

The two exponential spectra of Figure III.B-2 are directly compared with the
Apollo-15, ERS-18, and Ranger-3 energy-loss spectra in Figure III.B-3. The
important contribution of induced radioactivity in the 0.5-MeV to 5-MeV
energy range is apparent. In fact, the hard exponential spectrum exceeds the
measured energy loss from 2 MeV to 5 MeV, indicating that a softer spectrum
is required. It can also be seen that induced radioactivity has little effect on
the observed diffuse spectrum above 10 MeV from the Apollo-15 measure-
ments.

In correcting the Apollo-15 results for spallation-produced counts, Trombka
et al. (1973), required that the resulting spectrum be smooth and not dip at
intermediate energies, since such a spectrum is physically improbable. On the
basis of these assumptions, one-half of the calculated induced flux was sub-
tracted. The factor-of-two discrepancy between the calculated induced spec-
trum and that which was corrected for is within the range of errors of the
calculated induced radioactivity at the present. It is hoped that future
accelerator measurements will more accurately determine the true spallation
product spectrum so that direct and accurate corrections to the observed
diffuse 7-ray flux can be made.

OBSER VA TIONAL DA TA

10"1 -

icr2 -

in- 3 -

in-

1 1

T 1 I 1 III T 1 1-

1 MM]

1 1 1 1 1 | 1 1.

i —

— I

"^■^^

■»_

■I 1— i

*^v

T ^V

1 1 1

Ee = 0.9 MeV
Ee = 1.4 MeV

EXPER, 600 MeV PROTONS
XNaI = 25 gm Cm"2

\ \
— 1\ \
\ \

■•

I I i I i i i 1 1

\
\ \

\ ^ _

\ \ =

\ \ "

\ \ '

\ \ "

\ * -
\ *-

\ \

\ V

\ \

1 1

i i i i i il i ii

1 1 1 1 1

\ i

i i i i ill i i!

0.01

0.1

ENERGY (MeV)

Figure III. B-2. Energy-loss spectrum of induced radioactivity. The data points
are from direct measurements of 600 MeV proton-induced radioactivity in
Nal(T1), corrected for the expected contribution of unmeasured, short-lived
products (Fishman, 1972b). Also shown are two exponential spectra described
in the text, normalized to a total rate of 0.15 counts per incident high energy
proton.

INDUCED RADIOACTIVITY CONTRIBUTIONS

102c 1 — i i i i mi

— i — i i i i ii

* RANGER III

ERS-18
"S APOLLO 15

CALC.

I INDUCED

"l — i — i i i 1 1 '-i

ENERGY (MeV)

Figure III.B-3. Comparison of spectra measured in inter-
planetary space by three experiments (from Trombka et al.,
1973) with the calculated exponential spectra from Figure
III.B-2. A high-energy cosmic-ray flux of 3 cm"2 • s" was
assumed. In view of the Apollo-15 measurements, the softer
spectrum (solid line) is more likely. The estimated error
in the calculated induced spectrum is plus or minus a factor
of two at any energy.

70 OBSER VA TIONAL DA TA

REFERENCES

Dyer, C. S., and G. E. Morfill, 1971, Astrophys and Space Sci, 14, p. 243.

Fishman, G. J., 1972a, Astrophys. J., 171, p. 163.

Fishman, G. J., 1972b, Summary Report SE-SSL-1497, Teledyne Brown
Engineering.

Goldman, D. T., 1966, "Chart of the Nuclides," 9th Edition, General Electric
Company.

Goldstein, H., 1959, Fundamental Aspects of Reactor Shielding, Addison-
Wesley Publishing Company, Reading.

Golenetskii, S. V., 1971, Astrophys. J. Letters, 9, p. L69.

Peterson, L E., 1965, Geophys, Res., 70, p. 1792.

Rudstam, G., 1966, Zeits. fur Naturforschung, 21a, p. 1027.

Silberberg, R., and C. H. Tsao, 1973, Astrophys. J., Supplement No. 220.

Sobel, H. W., A. A Hruschka, W. R. Kropp, J. Lathrop, F. Reines, M. F.
Crouch, B. S. Meyer, and J. J. Sellschop, 1973, Phys. Rev. C, 7, p. 1564.

Trombka, J. I., A. E. Metzger, J. R. Arnold, J. L. Matteson, R. C. Reedy, and
L E. Peterson, 1973, Astrophys J., 181, p. 737.

C. PRELIMINARY RESULTS FROM THE FIRST
SATELLITE OF A HIGH-RESOLUTION GER-
MANIUM GAMMA-RAY SPECTROMETER:
DESCRIPTION OF INSTRUMENT, SOME
ACTIVATION LINES ENCOUNTERED,
AND STUDIES OF THE
DIFFUSE SPECTRA

G. H. Nakano*, W. L. Imhof, J. B. Reagan, and R. G. Johnson

Palo Alto Research Laboratory

Gamma radiation from terrestrial and extraterrestrial sources has been investi-
gated with a high-resolution lithium-drifted germanium, Ge(Li), spectrometer-
cryogen system flown on board a low-altitude, spin-stabilized, polar-orbiting
satellite (1972-076B) launched on October 2, 1972. The application of large
germanium spectrometers in 7-ray astronomy provides the high-energy resolu-
tion required to facilitate the detection of monoenergetic nuclear 7-rays and to
search for sharp structural features in the diffuse background spectrum. Sources
giving rise to these types of 7-radiation in terms of nucleosynthesis were re-
viewed by Clayton at this Symposium (Chapter XI. A) and 7-ray emission pro-
duced in solar flares also was discussed by Forrest (Chapter VI.A) and Ramaty
(Chapter XI.C). Although lithium-drifted germanium detectors have been
flown on a few occasions in high-altitude balloon experiments (Jacobson, 1968;
and Womack and Overbeck, 1970), this is the first time high-resolution de-
tectors of this type have been flown on a satellite. In this paper we present
a brief description of the instrument and discuss some very preliminary re-
sults obtained from earth orbit. In Chapter III.D we shall discuss some of
the important backgrounds encountered in the satellite flight. These two
papers represent a brief review of topics presented at the Annual Meeting
of the American Geophysical Union (Nakano et al., 1973; Imhof et al., 1973).

The important features of the spectrometers are shown as a cross-sectional
view in Figure III.C-1. In the flight system we employed solid cryogen
coolers, which provide obvious design advantages in the zero-G environment.

* Speaker.

OBSER VA TIONAL DA TA

HOUSING FOR THERMAL LINK
3

PHOTOMULTIPLIER

Figure III.C-1. A cross-sectional view of the impor-
tant features of the Ge(Li) spectrometer. Each
instrument weighed 83 kg (163 lb).

The 50 cm3 Ge(Li) detector, with an active area of 15 cm2, is maintained at
cryogenic temperatures by a copper thermal link which is coupled to the
cooler consisting of sufficient CO. , about 1 5 kg (35 lb), to provide a one-
year lifetime. The operating temperature of the detector (130 K) is somewhat
higher than the liquid nitrogen temperature normally used in the laboratory
(Nakano and Imhof, 1971). The instrument is collimated to ±45° by a large
tungsten shield weighing ~ 20 kg ( ~ 45 lb) which provides a minimum of
30 gm/cm2 shielding outside the viewcone. A plastic-scintillator anticoinci-
dence system completely surrounds the shield except for a small access
port for the thermal link and is viewed by four photomultiplier tubes. Gamma-
ray pulses from the Ge(Li) detector corresponding to energy losses ranging
from 40 keV to ~ 2.8 MeV are analyzed by the 4096 pulse-height analyzer
with an overall systems resolution of 3.5- to 4.0-keV full width at half-
maximum (FWHM) (1 .33-MeV Co60). The output digital addresses from the
pulse-height analyzer are stored on an onboard tape recorder which affords
data coverage on a worldwide basis. A maximum rate of 1625 addresses/s
can be recorded, and the counting rates are sampled once every 32 millisec-
onds.

HIGH-RESOL UTION GERMANIUM INSTR UMENT 73

The geometry of the experiment is illustrated schematically in Figure III.D-1
of the following paper (Imhof et al., 1973). It is important to note that the
satellite was launched into a noon -midnight, sun-synchronous orbit (inclination
~98.4°) with apogee and perigee of 761 km and 736 km, respectively. The
vehicle has a spin period of 5 s and is magnetically torqued to maintain the
spin vector perpendicular to the orbit plane so that in daylight (descending
node) the sun is always viewed once per spin period. Two identical 7-ray
spectrometers were mounted antiparallel to each other and offset ±15° from
the vehicle's equatorial plane. Unfortunately, one of the instruments failed
at launch. The other instrument performed successfully for about 10 days
and then suffered a serious degradation in gain and energy resolution; however,
after a few weeks it made a rather remarkable recovery and subsequently
operated at about 90 percent of its original gain with somewhat degraded
resolution of about 10 keV as compared to the original 3.5 keV at the 59.6-
keV line. The following discussion is confined to data obtained during the
first 10 days of operation when the resolution was 3.5 to 4.0 keV over the
entire spectrum.

A representative pulse-height spectrum is shown in the top section of Figure
III.C-2 where the data, integrated over all spin angles, are summed from
several low background passes at low and midlatitudes. In addition to the
59.6- keV 7-ray line from the Am241 in-flight calibration source, several
discrete 7-ray peaks are observed in the data. Some of these lines have been
identified and are associated with isomeric transitions induced in the german-
ium sensor itself (Womack and Overbeck, 1 970) and with short-lived radio-
activations produced in the tungsten shielding. In Table III.C-1, the more
prominent lines present in the spectrum are tabulated with their probable
production modes. The most prominent line, always present in the data, is
the 51 1-keV electron-positron annihilation radiation, a significant portion
of which apparently are produced on the satellite as indicated by the com-
paratively high count rate and by the lack of substantial intensity modula-
tions with spin angle. The overall pulse-height spectrum is characterized
by a rather hard continuum which is again a manifestation of locally produced
high-energy background.

In spite of difficulties due to the high-continuum background and large
statistical errors, the diffuse background 7-rays of cosmic origin can be
detected, particularly in the lower energy portion of the observed pulse-
height spectra. In a preliminary analysis the data from selected low back-
ground passes in the equatorial region were grouped into four spin quadrants.
The bottom two spectra presented in Figure III.C-2 correspond to data
taken when the look direction of the detector is in the upward quadrant and
in the downward quadrant, respectively. It should be noted that below 200

OBSER VA TIONAL DA TA

1.0

54 SO 93 139 166 198

JluklV ktV k.V k.V k.V
ill67 1 1 11
k.V

479

k.V
1

511 686

k.V k.V
1 1

*"*smhw''**^^

««*Wj*

V'.^>^^^r^^««5^w^K^^^;

0 01

SUM OVER ALL ANGLES
LOW AND MID- LATITUDES

0001

■»"•'-' ■—-'«* «■*«*». «

CHANNEL NUMBER

.J.,'.",

- !!=::rSH:3::3-a! DOWNWARD OUADRANT Sir: '-'jHEiSast!!
-EQUATORIAL

.... 311 . "

k.V
i

IKfe:!]

■ "*•

:| HI H] j

y ■ ft:

CHANNEL NUMBER

Figure III.C-2. A representative pulse-height spectrum where the data, inte-
grated over all spin angles, are summed from several low- and midlatitude
passes (top section). Pulse-height spectrum from related low-background
passes in the equatorial region when the look direction of the detector is in
the upward direction (middle section) and with downward direction (bottom
section).

keV in the continuum portion of the spectrum between peaks, the counting
rates appear to be greater in the upward viewing spectrum than in the down-
ward spectrum. This feature is confirmed when counts in the corresponding
channel intervals are summed to improve statistics and are directly compared.
A detailed analysis of the data, taking proper account of the shielding and
detection efficiencies, has just begun. At high energies it is more difficult
to obtain definitive measurements of the diffuse background since the gen-
eral backgrounds are relatively high.

HIGH-RESOL UTION GERMANIUM INSTR UMENT

Table III.C-1
Gamma-Ray Peaks Observed (Preliminary)

Source

54keV

Ge73m from Ge72 (n, 7) Ge73m

59.6 keV

Am in-flight calibration source

67keV

Ge73m from Ge72 (n, 7) Ge73m

93keV

139 keV

Ge75m from Ge74 (n, 7) Ge75m

186 keV

198 keV

Ge71mfromGe70(n,7)Ge71m

479 keV
511 keV

Re187fromW186(n,7)W187 -»

Id
Positron/electron annihilation

Re187

686 keV

Re187fromW186(n,7)W187 -~

Re187

Our spectrometer design was not optimized to study discrete stellar sources,
and its sensitivity to photons is quite limited above a few hundred keV.
Nevertheless, by making detailed analyses of the angular distributions and
spectral variations with geomagnetic latitude, we hold some prospects of
investigating the diffuse background spectrum, of detecting 7-ray line
emissions from solar flares and, perhaps, of searching for positron annihila-
tion radiation coming from the direction of the galactic center.

(Supported by the Office of Naval Research, the Advanced Research
Projects Agency, and the Lockheed Independent Research Program.)

REFERENCES

Imhof, W. L., G. H. Nakano, R. G. Johnson, and J. B. Reagan, 1973, Trans
American Geophys. Union, 54, p. 435.

Jacobson, A. S., 1968, Thesis, University of California at San Diego.

Nakano, G. H., and W. L. Imhof, 1971 , IEEE Trans. Nucl. Sci., NS-18,

p. 258.

Nakano, G. H., W. L. Imhof, J. B. Reagan, and R. G. Johnson, 1973, Trans.
American Geophys. Union, 54, p. 435.

Womack, E. A., and J. W. Overbeck, 1970,7. Geophys. Res., 75, p. 181 1.

. PRELIMINARY RESULTS FROM THE FIRST
SATELLITE OF A HIGH-RESOLUTION
GERMANIUM GAMMA-RAY SPECTRO-
METER: BACKGROUNDS FROM ELECTRON
BREMSSTRAHLUNG AND FROM

ELECTRON-POSITRON
ANNIHILATION

W. L. Imhof*, G. H. Nakano, R. G. Johnson, and J. B. Reagan

Lockheed Palo Alto Research Laboratory

In continuation of the previous talk we shall now consider in more detail some
of the backgrounds encountered in the first satellite flight of a lithium-drifted
germanium spectrometer. This is a brief summary of a recent presentation at
the annual meeting of the American Geophysical Union (Nakano et al., 1973;
Imhof et al., 1973). We have just seen that in the flight data several 7-ray lines
are observed and that these have been attributed to isomeric states produced by
cosmic rays interacting in the instrument. Fortunately, the intensities of these
lines are rather low. By far the most prominent background line experienced
is that at 51 1 keV. Smooth backgrounds attributable to electron bremsstrahlung
are also commonly encountered in the satellite measurements. The bremsstrah-
lung backgrounds can be divided into two basic classes: (1) those arising from
radiation belt electrons stopping in the vicinity of the spectrometer, and (2)
bremsstrahlung produced by electrons precipitating into the earth's atmosphere.
The latter phenomenon can be a significant source of background even when the
satellite is thousands of kilometers away from the radiation belts. The geometry
for observing two of these common sources of background, the 51 1-keV peak
and the bremsstrahlung continuum, is illustrated schematically in Figure III.D-1.
This drawing shows that 51 1-keV 7-rays are produced in the atmosphere by
cosmic rays at all latitudes and that bremsstrahlung associated with electron
precipitation can be observed from a satellite even at fairly low latitudes,
although it is a more prevalent background at higher latitudes.

* Speaker.

OBSER VA TIONAL DA TA

ELECTRONS

COSMIC RAY

Figure III.D-1. Schematic illustration of the satellite geometry. The viewing
cone of the spectrometer is shown for observing the bremsstrahlung associ-
ated with electron precipitation at high latitudes.

The aforementioned backgrounds are best illustrated by presenting some typical
energy spectra measured during the satellite flight. The spectra shown in Figure
III.D-2 were taken at a variety of geomagnetic latitudes and with the spectrom-
eter in various look directions. In the top section is shown a typical spectrum
associated with electrons precipitating into the atmosphere. When the 7-ray
spectra are compared with the electron spectra that are measured with an elec-
tron detector on the same satellite (at a time when the satellite passes directly
through the environment of the precipitating electrons), the shapes are consis-
tent with the bremsstrahlung calculations of Berger and Seltzer (1972; 1973).
At the time the bremsstrahlung spectrum in Figure III.D-2 was taken, the
satellite was over the South Pole of the earth and no radiation belt particles
were in the immediate vicinity. The bremsstrahlung spectrum produced by
trapped electrons when the satellite is actually in the outer radiation belt is
often harder in spectral shape than that associated with precipitating electrons.
This is attributable to the fact that the energy spectrum of the trapped elec-
trons is frequently harder than that of the precipitating electrons. However,
even the hardest bremsstrahlung spectra observed do not present a serious
background for observations above a few hundred keV. For example, when
the satellite is in the heart of the outer radiation belt the 5 1 1 -ke V peak can
usually be seen with little interference.

HIGH-RESOLUTION GERMANIUM BACKGROUNDS

SATELLITE AT MID LATITUDE
VIEW BREMSSTRAHLUNG FROM
PRECIPITATING ELECTRONS

■Hi w - ■ - \u n

•^v-A,

>V**«^,y««.s,.

*■'•' '^VT^t.'wi

1.0

8 w

I 0.01

As*

'.: ■rf^V^Vf^effi.^i. :i,«i.t-.-.'v'y-v-

' '*': TH^:\*'i;-".*H;'

•"■<**

;->'>,^.- '',_>•■•■. ; -,.y^.,/' .•» '

Q001U

CHANNEL NUMBER

Figure III.D-2. Some selected energy spectra measured during the satellite
flight. In the top section is shown a spectrum associated with electrons pre-
cipitating into the atmosphere. The bottom three sections contain spectra
measured over the earth's polar cap regions and near the Equator.

80 OBSER VA TIONAL DA TA

The bottom three sections of Figure III.D-2 contain spectra measured over
the earth's polar cap regions and spectra obtained near the Equator. Over
the polar caps the 5 1 1 -ke V line intensity is significantly higher when the
spectrometer is pointing downward. This increase in counting rate, when
viewing in the downward direction, is attributed to electron/positron
annihilation radiation produced in the atmosphere by cosmic-ray interactions.
From a preliminary analysis of the data, the counting rates and latitude varia-
tions are found to be consistent with balloon observations of the 5 1 1-keV
intensities. At low latitudes where the 51 1-keV production rate is weaker,
the contribution from the atmosphere is less evident in the data. It should be
noted that the continuum background is also weaker at low latitudes. We
are presently studying the data to see if there is an increase in the 5 1 1-keV rate
when the satellite is in the outer radiation belt, as a possible indication of
the existence of trapped positrons. At this early stage of the data analysis
only an upper limit has been found for the fluxes of trapped positrons. In
any event, their possible contribution to the background is negligible.

For unfolding the background contributions it is a great advantage to have
the spectrometer placed on a spinning satellite. This is clearly illustrated in
Figure III.D-3 where the counting rate profiles of the 7-ray spectrometer and
of an electron spectrometer on board the same satellite are shown. The top
section represents data taken during a complete orbit of the satellite, including
two polar cap crossings and four outer radiation zone crossings. Both the
electron and 7-ray spectrometers respond significantly in the outer belt
regions. In addition, on this orbit that occurred at a time of high geomagnetic
activity, significant increases in the 7-ray counting rates were observed over
the polar caps and Equator-ward of the belts. One of the regions where
the 7-ray counting rates were enhanced is shown in the middle section of the
figure on a more expanded time scale. Here one can clearly see the pronounced
modulation in counting rate. Some individual spin profiles, summed over six
spins to improve statistics, are shown in the bottom section of the figure. With
a careful analysis of the data, the bremsstranlung source distributions can be
unfolded from the data (Imhof et al., 1973). Likewise, when analyzing the
data for 7-rays originating from other sources, the bremsstranlung contribu-
tions can often be eliminated with the selection of data at desirable positions
and look directions.

HIGH-RESOLUTION GERMANIUM BACKGROUNDS

Hffi

551

A»XS.£-»-A»5<0g 'SAVa-VWWW9

■^

-i

u
c

■o

U
QJ

•i-'

03
3

■a

*->

-C

■f-1

■5

c
o

•I-1

o
*->

QJ
03

+-'

+->

-Q

iEo8

82 OBSER VA TIONAL DA TA

The results presented here are taken from a very preliminary analysis of the
data acquired in the first satellite flight of a high resolution Ge(Li) 7-ray
spectrometer. This flight has demonstrated the practicability of flying such
a system on a satellite and has provided much information on the backgrounds
encountered. These data will represent an important basis for designing future
7-ray spectrometers for satellite usage.

(Supported by the Office of Naval Research, the Advanced Research
Projects Agency, and the Lockheed Independent Research Program)

REFERENCES

Berger, M. J., and S. M. Seltzer, 1972, J. Attn. Terr. Phys., 34, p. 85.

Imhof, W. L., G. H. Nakano, R. G. Johnson, and J. B. Reagan, 1973, Trans.
American Geophys. Union, 54, p. 435.

Nakano, G. H., W. L. Imhof, J. B. Reagan, and R. G. Johnson, 1973, Trans.
American Geophys. Union, 54, p. 435.

FURTHER CONSIDERATIONS OF
SPALLATION EFFECTS

Give Dyer*

Imperial College

I just want to reinforce what Dr. J. Fishman has said by presenting a few
results which we arrived at independently at Imperial College, London. We
undertook an investigation to estimate the effects of South Atlantic Anomaly
Traversals on the future UK-5 hard X-ray telescope that was primarily aimed
towards discovering the effects of the activation of the central detector crystal
by trapped protons.

We were able to see from the results that cosmic-ray effects would also be
important, and these I'll mainly talk about here.

We based our estimates both on the Rudstam formula and, as J. Fishman has
pointed out, also on an extensive number of irradiation experiments.

Figure III.E-1 summarizes the results of an accelerator experiment in which
we irradiated the UK-5 central crystal which was 5 cm long and 3.4 cm
diameter of Csl with 155-MeV protons.

Figure III.E-1 shows the energy -loss spectra that we obtained soon after
irradiation.

As Dr. Fishman pointed out, it is important to measure these decays over a
wide range of times. He mentioned down to 10 microseconds.

We measured the activation after one minute, which is the quickest we could
get the crystal out of the beam and optically seal it on our photomultiplier
tube; thus, we have in fact experimentally measured decays down to shorter
half-lives than Fishman was able to, and perhaps we have to apply less of a
unit to the estimated correction due to extrapolation back to these short
half-lives.

* Speaker.

OBSER VA TIONAL DA TA

COUNTS

keV1

100

1000

ENERGY (keV)

Figure III.E-1. Pulse-height spectra of the proton-induced activation of 3.4-
by 5-cm Csl (UK-5 central detector) as a function of time after low-dose
irradiation. Shortest period after irradiation is 1 min and the longest period of
time is 2 hr.

CONSIDERA TIONS OF SPALLA TION EFFECTS 85

These spectra all show the peak on the left, which is the K-capture peak at
about 35 keV. A second major feature is found at around 170 to 200 keV,
depending on the time after irradiation that the spectral measurement is
made, and it is due to 7-emissions from a number of isotopes, which will be
listed.

The time scale of the decays goes from 1 min on the highest curve in Figure
III.E-1 down to 2 hr on the bottom curve. Also shown are features just above
400 keV and maybe just above 600 keV. You can see these features more
clearly in Figure III.E-2.

The other major feature is the j3+- continuum of decays, which is the important
feature when it comes to correcting the Apollo results, and this gives the shape
shown around 1 MeV.

As you can see, the spallation continuum is quite hard to start with, but it
decays away quickly because of the shortness of the half-lives of the j3+- decays.
It drops at about 3 to 4 MeV very steeply.

Figure III.E-2 also shows the longer half-life decays. We obtained these by
giving a higher dose to a second crystal. You really need to expose the crystals
to two different dosages to cover the wide range of decays from a minute to
several hundred days. These half-life decays run from Curve P at the top of
Figure III.E-2 (taken 7 hr after irradiation when the rate was down) down to
the very bottom curve which is going into the background about 200 days
after irradiation.

You can see that the j3+- emitters have really decayed away quickly, and we
are left with quite a number of important line features at 35 keV and around
200 keV, 400 keV, and 600 keV.

I will not go into the isotopes for those, because I can provide lists to interested
individuals.

Figure III.E-3 shows the results of computations based on the Rudstam for-
mula for the number of isotopes produced in such a crystal and due to
different types of energetic particles. We used a typical inner-belt spectrum,
we assumed monoenergetic 155-MeV protons to compare with the experiment,
and we also assumed cosmic rays of 2 GeV in energy where the spallation
cross-section becomes independent of energy.

The short-term decays are listed in Table III.E-1. You can see a number of
|3+- emitters. There are none longer than 2 hr so that they decay very quickly.
Table III.E-2 shows the longer half-life decays, mainly electron captures.
One can see decays at around 200, 400, and 600 keV and at long half -lives
of around 150 days, which can build up activity over a satellite history. One
can observe the rates building up if the detector is flown for a sufficiently long
time, of the order of several hundred days.

OBSER VA TIONAL DA TA

100 1000

ENERGY (keV)

Figure III.E-2. Pulse-height spectra of the UK-5 detector as
a function of time after high-dose irradiation. Shortest per-
iod after irradiation is 1 min; the longest period is several
hundred days.

CONSIDERA TIONS OF SPALLA TION EFFECTS

FLUX

-2 -1

cm ' s

sr1

keV1

10"1

10:

10 4 -

10 '

101

° OSO-3 uncorrected

x OSO-3 corrected by Schwartz

i i Estimated Activity for OSO-3

detector (min. after 2 weeks)

• Ranger-3, uncorrected

1 ' ERS-18, uncorrected

i 1 Estimated activity (after 60 days C.R.

flux of 3.0 cnvV > 40 MeV).

. ...I

102 103

ENERGY keV

10"

Figure III.E-3. Uncorrected measurements of the diffuse X-ray
spectrum and estimated corrections for induced activity.

OBSER VA TIONAL DA TA

Table III.E-1
Short-Term Decays

Isotope

Decay Mode and
Energy (MeV)

Half- Life
TtA (mins)

Predicted Numbers Produced

Inner Belt

(ii)

1 55MeV

(iii)
Cosmic Rays

55 Cs

(J+ 1.97
/T 0.442

3199

3183

576

55 Cs

(3+ 3.0 (70%)
2.5 (30%)

E.C. (25%)

0.460 (20%)
0.285 (2Q%)

8623

8448

3063

ssCs

E.C.

y 0.406

360

11959

11190

5743

55 Cs

(3+ 3.8 (82%)
E.C. (8%)
0.426 (7%)

1.6

13947

12158

8834

Aei25

y 0.187
0.056
0.243

E.C.

1080

4324

3431

3529

Aei23

t? 1.7
0.148

120

6584

4068

8385

'123

E.C.

7 0.160

780

8340
12411

7929
12124

3322
3492

122

0+ 3.0

11935
17598

10918
16695

6477
6806

'|21

/3+ 1.2
0.21

14395
21015

12345
18876

10367
10894

nTe

127

p^0.70

560

6000
9060

5,sb-8

p* 3.1

3.5

1577
2229

967
1479

2305
2422

S1sb117

EC.
7 0.161

168

2609
3648

1376
2104

4641
4877

5,*>„<,

0+ 2.4
1.3

3592
4969

1601
2449

7646
8035

"s»»u

E.C. (71%)

657
865

105
198

2847
2992

CONSIDERA TIONS OF SPALLA TION EFFECTS

Table III.E-2
Long-Term Decays

Isotope

Energies of 7-rays (MeV)
and Branching Ratio

Half-Life
r t/j (days)

Predicted Numbers Produced

(i)
Inner Belt

(ii)
155MeV

(iii)

Cosmic Rays

55 Cs

EC. 0.670

6.2

6000

55 Cs

E.C.

1716

1659

212

S5Cs

E.C. 0.380

1.3

5507

5508

1410

^Cs

E.C. 0.406(80%)

0.25

11959

11190

5743

Ae!3l

0.163

Aei29

0.196 jO.040

367

Aei27

E.C. 0.370(40%)
0.203 (60%)

1505

1408

723

Aei2S

E.C. 0.1 87; 0.243

0.7

4324

3431

3529

54 Xe

Aei23

0.148

0.08

6584

4068

8385

126

E.C. (55%) 0.386
(34%); 0.650

(33%)

6000
9060

3 1

M26

E.C.

2899
4382

2750
4206

575
604

'124

E.C. (70%) 0.605

(95%)

5159
7742

4964
7591

1469

1544

S3,
'l23

E.C. 0.160

0.5

8340

12411

7929
12124

3322
3491

52 j m
1 125

0.1 10; 0.035

136
206

129
197

27
28

By m
1 123

0.089; 0.1 59

104

685
1020

652
996

273
287

By m
lei21

0.082;0.214

154

2564
3743

2199
3362

1846
1960

2Te

121

0.570 (87%);
0.506(13%)

14395
21015

12345
18876

10367
10894

BTe119

E.C.

4.5

6082
8693

4264
6520

7176
7541

s,sb122

0.566 (66%)

2.8

71
105

65
99

31
40

5,sb120

E.C. 0.089; 0.199

397

574

312
477

371
389

slSb

E.C.

1.66

834
1192

585
894

984
1034

"Sn™,

0.159-.0.162

284
397

150

229

505
530

90 OBSER VA TIONAL DA TA

Notice that there are no j3+- decays in Table III.E-2. None of the |3+ emitters
which are produced have half-lives of the order of days. These isotopes have
very short half-lives.

Experimental data of this sort enables you to plot decay rate against time
after radiation. If you normalize these data to a typical cosmic-ray flux,
taking 3 particles/cm2 • s as the intensity (because this is the magnitude of
the flux quoted for the time of the ERS-18 measurement), plot all these
decay curves, and find the area under them, you can estimate the activity
built up after a certain time after radiation. Using this method, the activity
build-up for up to 60 days after radiation was obtained. This is the ERS-18
time scale, and we thus came up with the correction shown in Figure III.E-3
for the measurement obtained during that flight.

The left end of Figure III.E-3 shows corrections that apply at the lower
energy end of the spectrum and will not be discussed here. At the high-
energy end, in the region of 1 MeV, I show the ERS-18 points in the usual
manner, depicting a flattening and an apparently erroneous channel that
goes from 4 to 6 MeV. And I show our estimates of the build-up activity
correction due to the cosmic-ray flux. I believe this is the sort of correction
which Jack Trombka has tried taking away, and I'm a little puzzled as to
why it's a factor of 2 too high.

Dr. J. Fishman says these calculations are good to perhaps 20 percent, and
I don't think they are off by a factor of 2, particularly in our case as we
obtained measurements up to 3 a minute after irradiation. Here there is
no arbitrary factor depending on the theoretical computations of the half-
life distribution. We simply took the experimental data and found the area
under the decay curves. As Dr. Fishman pointed out, what might be missing
in this procedure is half-lives from do jus to 1 min, which would raise these
estimates even more and might produce a few decays in excess of 4 MeV.
I don't know that these 10-£is to 1-min half-lives would be too important,
but they might extend a correction to above 4 MeV where the present
estimate of the j3+- energy deposition falls over. Also, I believe that
Apollo- 15 gave the same energy-loss spectrum as the ERS-18 results, and
my correction in the 1- to 2-MeV channel is around 50 percent— I'm not
sure how you say that a 50 percent correction produces a hole in the
spectrum if subtracted from the total spectrum. Perhaps that's a good
note to end on and start the discussion.

DISCUSSION

Shapiro:

Dr. Dyer and Dr. Fishman both use a reasonable cosmic-ray flux, but as was
pointed out, the energy even in the case of 600-MeV protons is quite a bit

CONSIDERA TIONS OF SPALLA TION EFFECTS 91

below the average cosmic-ray energy. In Dyer's case, it is roughly a factor
of 20 below the average, and Fm wondering what influence this might have
upon the estimated corrections?

Dyer:

Well, that is a matter of concern. I'd like to see flight detectors actually
irradiated and that type of technique used to do the corrections. I think
this is the only way of doing it. You need to measure three types of
radiation: short half -life radiation, then radiation with a half-life of a
minute to several hours, and finally radiation with a half -life of several hours
to several days.

I wouldn't think it would lower the correction, having a more energetic
beam striking. I think you still get a very large number of |3+- emitters
produced, and I don't see how they are going to be a factor of 2 different
from the 155- or 600-MeV radiation.

Member of the Audience:

The cross sections for the production of those products above about 600 MeV
are reasonably constant, aren't they?

Shapiro:

Well, I'm more disturbed about the 155 MeV calculations.

Member of the Audience:

One-hundred and fifty-five is low. Above the 600 MeV or so, as Jerry
Fishman said, the cross sections seem to be reasonably constant or at least
they are slowly varying.

Shapiro:

How much difference would it make whether you use Rudstam's original
formula or the improved Silverberg corrections?

Member of the Audience:

In that connection my guess is that the effects of the heavier nuclei would be
more than proportional to their number, and I wonder if you have considered
that?

Dyer:

That if you got a higher energy irradiated flux you produce larger nuclei?

92 OBSER VA TIONAL DA TA

Vette:

No, a-particles in the cosmic rays.

Dyer:

No, Fve not considered that. I wouldn't consider it so important, less than
1 percent.

Vette:

As I recall also, your (Dyer) calculations would indicate you would see a
buildup running at least 60 days after launch, or something like that?

Dyer:

Yes, it would be confined to the lower energy end of the 7-ray spectrum in
the 1 00- to 200-ke V region where you could detect this sort of buildup. As
I pointed out, these 0+- emitters decay with a half-life of less than 2 hr,
and therefore, you don't see any buildup after 2 hr in the 1-MeV region.
So maybe experimentally you can get a handle on the buildup by looking for
the buildup in the 1 00-ke V region after a few days and use that factor to
give you a measure of the j3+- correction.

Vette:

As I recall from the ERS-18 results, there were no time changes after a
couple of days. There were some changes which might have been due to
activation, but there was really nothing after the first 2 or 3 days. There
seemed to be some problems in correlating some of the calculations with
some of the observations.

Forrest:

On the OSO-7 we also observed some of these induced 7-ray lines. I think
all of the features are predicted, but some of the predicted features are not
seen. This may indicate that some of the predictions are overestimated.

Dyer:

Are you talking about cesium iodide?

Forrest:

We have both on OSO-7: the sodium iodide detector and the cesium iodide
shield.

CONSIDERA TIONS OF SPALLA TION EFFECTS 93

Trombka:

What if we have, for instance, a microsecond dead time after each proton event,
how would that change your predicted values? In other words, on Apollo- 1 5
and -16, we're shut off because of the anticoincidence mantle for certain
periods after a proton event.

Dyer:

For how long is it?

Metzger:

In most cases on Apollo, it was a period of about 12 or 13 j/s.

Dyer:

Well, as Jerry Fishman mentioned, this gives the lower threshold above which
you are interested in measuring the half-lives. You want to measure the half-
lives in excess of 10 ;us and find the error in the decay curve down to the
lower limit. Our present data has been obtained for times down to 1 min.
The mystery is what happens in the 10-fis to 1-min region. There might be
additional decays there.

Pieper:

Just as a matter of opinion, I don't think you're going to get much of anything
that is going to deposit more than 10 MeV in the detector at a time later than
a dozen microseconds, or something of that sort. Shorter times of less than
1 us are possible, particularly if you do have an anticoincidence shield. But
there aren't that many highly excited states in the kinds of isotopes you're
going to find that are nearer your target nucleus, even reasonably near on the
periodic chart.

You have to get down to the carbons and oxygens and the low Z-materials
that have high excited states, and you don't form very many spallation
products.

Dyer:
Right.

Fishman:

Right. What I had in mind was |3-7 cascades where you step through perhaps
three or four different nuclei in rapid succession within a microsecond or so.
I admit this is very unlikely, but it doesn't take much between 10 and 20 MeV
to subtract from the measured fluxes.

94 OBSER VA TIONAL DA TA

Metzger:

Unless that happens beyond the dead time of the instrument, the instrument
does not know that it has happened.

Can we go into the high-energy region?

Vette:

Yes, I want to open up the general discussion on observations, experimental
problems, techniques, and then we might come back to the problems dis-
cussed by Dyer and Fishman.

I just want to make one point. If the spallation effect could possibly be taken
as the cause of the lower energy feature in the 1-MeV region, then there
remains a problem of the higher-energy region, 10 to 20 MeV, that Apollo
has seen. The question does arise as to what occurs at the open end of the
system which does not have the anticoincident plastic around it. I was
wondering whether you have considered the cosmic-electron flux that could
pass through that open end.

I made some quick estimations the other day and I came up with a number
of something like 25 percent of the events observed in the Apollo-15 experi-
ment could be due to cosmic-ray electrons just going through that open end
of the detector. This is basing the estimation on the spectra of Simnet and
McDonald.

In your paper (Peterson and Trombka, Chapter III.A) you did not mention
your anticoincidence efficiency, but you did mention that you had a threshold
of 1 MeV for detecting particles going through the anticoincidence shield.
Have you determined that your anticoincident efficiency is better than one
part in 500 or one part in a thousand. Without such rejection efficiencies, you
would obtain an apparent 7-ray intensity in the 20-MeV region.

Peterson:

I think the problem of rejection efficiency must be more like one part in 50
or one part in 1 00. You can determine that if you look at the energy-loss
curve given in our paper.

Trombka:

Yes, with the anticoincidence on and off one gets a feel for the rejection
efficiency. The rejection efficiency curves presented in our paper (Chapter III.A)

CONSIDERA TIONS OF SPALL A TION EFFECTS 95

were obtained with the anticoincidence shield both on and off during flight.
These results were compared with ground studies of the rejection efficiency
for the system as a function of energy.

In terms of the matter of the secondary electrons, there are two effects. In
the first place, there is a glass plate in front of the scintillation detector. I
don't know how many of the electrons will be stopped, but it should be an
effective shield. Again, the problem of the dead time of the instrument should
be considered because it would be shut off most of the time when such a
cosmic-ray effect would occur.

Peterson:

I haven't worked out the numbers, but there's a fair amount of material in
back of the scintillation detector so the electrons first of all have to pass
through about 20 gm/cm2, which means we're talking about 50-MeV elec-
trons penetrating, probably. I don't know what the integral flux of 50-MeV
electrons is, but it has to be on the order of ~ 10/cm2 • s.

Member of the Audience:

I think I'm confused about the factors of 2 in the magnitude of the spallation
effect discussed by Trombka. If the spallation cross sections don't change
very much between several hundred MeV and a few GeV, then most of the
cosmic rays are around several hundred MeV and modulation effects could
easily, I would have thought, give a factor of 2 in the magnitude of spallation
combination. Now, is that not true?

Vette:

I think all of these are normalized to the observed cosmic-ray flux at about
50 MeV or so.

Member of the Audience:
At what time?

Vette:

At the time of the measurement.

Trombka:

I think that the problem lies in the comparison of the results obtained on
Apollo-1 5 as compared with Apollo-17. On Apollo-17, we looked at the
longer lived induced spallation lines. What we saw during the transearth
measurements on Apollo-15 were rather distinct lines in the 0.57- to 0.7-MeV

96 OBSER VA TIONAL DA 7>

region, which we attributed to spallation products. These lines were observed
in the measurement made on the Apollo-17 detector after recovery. The
difference in the induced activity calculated using the intensity in 0.57 to
0.6 MeV on Apollo-17 with that which I observed on Apollo- 15 was a factor
of two or three. That is, Apollo-17 intensity was higher than that observed
on Apollo-15.

I realize the environment around the crystal on both flights was somewhat
different. The Apollo-17 crystal was stored in the Command Module while
the Apollo- 1 5 crystal was in the Service Module. There is a difference in
exposure time also: the 250 hr on Apollo-25, which is when we measured th(
transearth spectrum, as compared to 300 hr on Apollo-17. Thus, you know
there is not a factor of 2 in that time difference.

F. HEAO GAMMA-RAY ASTRONOMY
EXPERIMENTS

A. Metzger*
Jet Propulsion Laboratory

I will report information which is perhaps not known to all the theorists, and
particularly the foreign guests, at this Symposium, namely, two 7-ray experi-
ments in the 0.1- to 10-MeV region that had been planned as part of the initial
HEAO program.

The first experiment is the combined UCSD-MIT experiment, based on a design
originally proposed by Larry Peterson and his collaborators at the University
of California at San Diego (UCSD) and recently modified to include aspects of
an experiment proposed by Walter Lewin of the Massachusetts Institute of
Technology (MIT). It will be flown, hopefully, on the first satellite of the
reconstituted HEAO program in 1976 or 1977.

This experiment has been designed to measure and map the cosmic 7-ray
spectrum. A schematic of the instrument is shown in Figure III.F-1 . It is a
large scintillation detector system that can obtain data in a number of modes.

The center detector is a 12.7-cm (5-in.) diameter by 7.6-cm (3-in.) long sodium
iodide crystal viewed by a single photomultiplier tube. Around the central
detector is an annular shield of cesium iodide for anticoincidence rejection of
both charged particles and 7-rays. The full-width at half maximum field of view
of the central detector is 40°. A phoswich of cesium iodide is located between
the sodium-iodide crystal and its photomultiplier tube to provide rejection from
the rear. The purpose of this is to provide a directional system with maximum
sensitivity for the diffuse spectrum based on efficient suppression of the 7-ray
flux entering outside the field of view. The annular shield is divided into two
halves, and the capability exists for them operate as a pair spectrometer in
coincidence with the central detector.

* Speaker. Q7

OBSER VA TIONAL DA TA

HEA O GAMMA-RA Y ASTRONOMY EXPERIMENTS 99

Around the inner system is a circular grouping of six detectors which are of two
types. One type is shown with cross-hatching and is designed to be a low-energy
detector system— there are two of these. Each one is a sodium iodide scintil-
lator, 12.7 cm (5 in.) in diameter by 0.95 cm (0.375 in.) thick. They are
designed to cover a range of roughly 30 to 300 keV and also have the phoswich
configuration to provide rejection in the backward direction.

The field-of-view of each low-energy detector is 1 .5° by 20° and that is achieved
with a passive-slat collimator.

The four detectors of the other type positioned on the circle are designed to
more accurately localize point sources than the central detector, so that the
field-of-view of each of these is 20°. These sodium iodide crystals are 7.6 cm
(3 in.) in diameter by 7.6 cm (3 in.) in length in order to cover the same energy
range as the central detector. They also have the phoswich configuration.
Outside the circle of six detectors is an external anticoincidence shield, which
serves the circle of six as well as the central detector. So the anticoincidence
shield is of varying thickness for the different detectors, with maximum
effectiveness for the central detector.

There is sophisticated logic to allow this detector system to be commanded in
any number of ways. For example, all of the scintillators around the central
detector can be programmed to function in anticoincidence with the central
detector. The reduction in background should be very significant in this case.
It will be possible to position a blocking crystal above any of the seven primary
detectors. In addition, the entrances to each of these detectors will be covered
at all times by a thin sheet of plastic scintillator in order to remove electrons
which might otherwise be mistaken for 7-rays.

The sensitivity of the system has been calculated as sufficient to detect a
source one-thirtieth of the Crab's emission at 0.3 MeV, and one-third of the
Crab's emission at 3 MeV.

The second experiment, one of the instruments accepted for the original
HEAO-B mission, was a high-resolution 7-ray spectrometer that utilized
solid-state detectors. This experiment was proposed by Bud Jacobson at the
Jet Propulsion Laboratory (JPL). It has how been deferred, hopefully to the
third launch of the revised program, although the payload of the spacecraft
has not been set as yet. A schematic of this instrument is shown in Figure
III.F-2.

This is an early version of the instrument. There have been a few changes, but
the basic arrangement is the same. The system contains four Ge(Li) solid-state
detectors. If the instrument can be kept at the present scale, each of the Ge(Li)
detectors will have a volume of some 60 cm3 and a surface area of 16 cm2. The

100

OBSER VA TIONAL DA TA

specified resolution is 2.5 keV close to 1 MeV, which means a resolution some
40 to 50 times better than one can expect to get from a sodium iodide
detector.

Ge(Li) GAMMA RAY DETECTOR

INPUT FET'S
AND BIAS FILTER
PACKAGE

1-1/2 in. DIA DOME-
FACE PM TUBES
(14 REQ'D)

SOLID CRYOGEN
REFRIGERATOR

Ge(Li) CRYSTALS

46 mm DIA, 60 cc VOL

(4 REQ'D)

PROPERTIES

2
EFFECTIVE AREA = 64 cm
FIELD OF VIEW = 30 FWHM
SOLID ANGLE = 0.21 ster 2
TELESCOPE FACTOR = 13 cm - ster

SHIELD ISOTROPIC 2

GEOMETRY FACTOR = 1160 cm

Figure III.F-2. HEAO Ge(Li) detector and refrigerator.

The solid-state detectors will be surrounded by a thick anticoincidence mantle
of CsI(Na), which will be in two parts, covering the sides and rear, somewhat
like a clam shell in configuration. Above the cluster of solid-state detectors
will be a CsI(Na) collimator with holes drilled to permit access to the solid-
state detectors. That is what defines the field-of-view which will be 30° full-
width at half maximum, equivalent to a solid angle of 0.21 sr.

HEAO GAMMA-RA Y ASTRONOMY EXPERIMENTS

101

Cryogenic cooling is needed. The refrigerator presently planned is a two-stage
sublimation unit using solid methane and ammonia. The arrangement of the
heat-transfer system has been changed to exit from the rear instead of the
side in order to improve thermal performance and simplify the mechanical
design.

The electronics and command capability will be designed so that the system
can function as a total absorption spectrometer, a sum-coincidence spectro-
meter, and a pair spectrometer.

The calculated line sensitivity at a 3-a level as a function of energy is shown by
the solid line in Figure III.F-3. These calculations have been carried considering
the background contributions from the earth's albedo, the cosmic 7-ray
spectrum, and an estimate of what the spacecraft is likely to produce. The
sensitivities calculated are well below upper-limit predictions of line intensities
for /--process 7-ray from the Crab Nebula.

Ge(Li) DETECTOR LINE SENSITIVITY AT A 3<r LEVEL

0.1

ENERGY (MeV)
Figure III.F-3. Ge(Li) detector sensitivity to line spectra from a point source.

102

OBSER VA TIONAL DA TA

Figure III.F4 shows the capability of the instrument in terms of predicted
7-ray line fluxes from supernova as calculated by Professor Clayton (see
Chapter XI. A). The fluxes are based on a supernova occurring at a distance of
1 megaparsec, which means that a supernova of the estimated intensity could
be detected at a distance of 9 megaparsecs or, alternatively, that at a distance
of 1 megaparsec. Such a supernova could be seen for 108s or for several
years after its occurrence.

o
ci-

10-* -

10"° -

io-' -

10"

1 1 1 1 1

Ni56(0. 812 c

MeV)^/-\ d-106pc
y A-^ M(Ni56) = 0.14Mq

-3

// \X\

If \ N \ DETECT0R

/ \ \\ SENSITIVITY
/ \ \\ AT 0.812 MeV.

\ \\ X

-4

/ 4
/ /

/ 1
/ /

/ / y

/ / /
" / / /

/\ \ \\

-5

\\ \V ^ — — (0.511

\\ \\ MeV)

^V48 H \\

i / \

(1.31 MeV) V \\

-6

\ M

\ \ L

\ l\ \]^Co56(0.847MeV)

-7

111 ~~

So**

rCr48 \ I

11.16 MeV),

1 (0.31 MeV)| 1 p — -^

\ 1

1 "^

-8

-9

;

' 1 ' 1

10;

10'

10*

io1

t (sec)
Figure III.F-4. Supernova 7-ray fluxes.

Chapter IV

A. RECENT OBSERVATIONS OF COSMIC
GAMMA-RAYS FROM 10 MeV TO 1 GeV

Gerald H. Share*

Naval Research Laboratory

INTRODUCTION

Radio astronomy was born in the 1930's when Karl G. Jansky (1932; 1933)
discovered a "steady hiss-type static of unknown origin" which he concluded
"is fixed in space, that is, that the waves come from some source outside the
solar system." The source was in the direction of the center of the galaxy.
From further observations, Jansky demonstrated that radio emission is also
observed, but with diminished intensity, when other regions of the Milky Way
passed within the field of view of his antenna. Some 30 years later the newest
branch of astronomy was born when a detector on board the OSO-3 satellite
found that 7-ray photons 1016 times more energetic than the radio waves were
also emitted from the plane of the galaxy (Clark, Garmire, and Kraushaar,
1968). However, the similarity in the early histories of these two disciplines
stops right there. Whereas Jansky discovered extraterrestrial radio emission
while studying the arrival direction of thunderstorm static, the discovery of
cosmic 7-rays came after more than a decade of intensive investigation by
various laboratories.

In this paper, I shall discuss recent observations of cosmic 7-rays made subse-
quent to the discovery of energetic photons from the galactic plane. An
extensive review of the field prior to 1971 has been compiled by Gal'per et al.
(1972; also Fazio, 1973; and Pal, 1973). I shall treat three main areas under
current investigation: (1) 7-ray emission from the plane of the galaxy, with
emphasis on observations made in the vicinity of the galactic center; (2) 7-ray
emission from the Crab Nebula and its pulsar; and (3) diffuse 7-radiation.

GAMMA RADIATION FROM THE PLANE OF THE GALAXY

The OSO-3 telescope measured detectable intensities of 7-radiation emitted along
the galactic equator at all galactic longitudes. These measurements are sum-
marized in Figure IV.A-1 , taken from a final report on the observations

* Speaker.

103

104

OBSER VA TIONAL DA TA

90 <in<!50

5 3r

270<in<330

210 < £a<270

PX^l

90 -60 -30 0

30 60 90 -90 -60 -30
Galactic Latitude (degrees)

30 60 90

Figure IV.A-1. Variation of the counting rate of cosmic 7-rays observed
from 0S0-3 as a function of galactic latitude for successive 60° intervals of
galactic longitude.

(Kraushaar et al., 1972). The variation in counting rate of the instrument is
shown as a function of galactic latitude for six 60° intervals of galactic longitude.
For comparison the authors have indicated by the histogram the expected rates,
assuming that the radiation originated in collisions of cosmic-ray nuclei with
interstellar gas. The galactic distribution of gas was obtained from 21 -cm
measurements of atomic hydrogen. The agreement between the expected
intensity and their observations is good, with the exception of the region near
the galactic center. In this region, they found that the measured intensity was
significantly above the calculated value. Because the radiation appeared to be
associated with diffuse emission from the plane, they expressed it in terms of
an equivalent line intensity (7/cm2 -s-rad) for an apparent width of ±1 5 in
latitude. For longitudes 30° < 2U < 330°, they measured an average integral
intensity of (3.4 ± 1.0) X 10"5 7/cm2 -s-rad for energies above 100 MeV;
whereas in the vicinity of the galactic center, they found a broad maximum
along the plane with an intensity of (1 .1 ± 0.3) X 10"4 7/cm2 -s-rad.

As the angular resolution of the detector of OSO-3 was about ±15°, the width
of the apparent band of emission in directions away from the galactic center
could have been almost entirely due to instrumental effects. However, the
broad maximum in intensity, observed along the galactic equator in the
direction of the center, could not be attributed entirely to instrumental effects.

GAMMA-RA YS FROM 10 MeV TO 1 GeV 105

Ogelmann (1969) suggested that the distribution of 7-ray emission from the plane
could be accounted for by the distribution of known X-ray sources, assuming
that they emitted photons with a hard spectrum, « E"2 in differential intensity.
This suggestion could not be tested in greater detail by the OSO-3 detector
because of its limited angular resolution.

Initial measurements at higher angular resolution were made predominantly in
the Northern Hemisphere. Most of these instruments employed multiplate
spark chambers as their prime detector, which permitted angular resolutions
better than ±3°. In some early reports, evidence was presented for emission of
7-rays from the plane of the galaxy in the vicinity of Cygnus (Valdez and
Waddington, 1969; Frye and Wang, 1969; and Hutchinson et al., 1969).
However, these measurements were of marginal statistical significance and,
furthermore, indicated an intensity considerably above the revised intensity
measured on OSO-3 (Kraushaar et al., 1972).

The higher intensities observed in the direction of the center of the galaxy
prompted balloon expeditions to the Southern Hemisphere by various groups.
Using a wire spark chamber with magnetic-core readout, the group at Goddard
Space Flight Center investigated the galactic center region with an estimated
angular resolution of ~2° at 100 MeV. Their instrument was a prototype
version of the SAS-B 7-ray telescope which was launched late in 1972. From
a balloon flight conducted over Australia in 1969, Kniffen and Fichtel (1970;
also Fichtel et al., 1972) confirmed the high 7-ray intensity in the vicinity of
the galactic center (-25° 02u < +20°). Their results are summarized in
Figure IV.A-2, where they have summed their data in 2° and 6° bands of
latitude. On comparing the observed distribution with what they would have
expected for atmospheric 7-rays, they found about a four standard-deviation
excess within ±6° of the galactic equator. The measured "line intensity"
> 100 MeV, (2.0 ± 0.6) X 10"4 7/cm2-s-rad, is in agreement with that obtained
from OSO-3. Fichtel et al. (1972), also set an upper limit on the galactic flux
emitted between 50 MeV and 100 MeV. This limit led them to conclude that
at least 50 percent of the galactic flux comes from the decay of 7r°-mesons
produced in cosmic-ray collisions. They also searched for possible point sources
in this vicinity and were unable to detect any at a sensitivity of about 3 X 10"5
7/cm2-s above 50 MeV.

However, three other groups using balloon-borne instruments sensitive to
photons > 100 MeV have failed to detect diffuse emission from the galactic
plane near the galactic center. The first group, a collaborative effort between
Case Western Reserve University and the University of Melbourne, has reported
results from a series of three balloon flights over Australia, during an investi-
gation of 7-rays in the Southern Hemisphere (Frye et al., 1971a). Their
investigation was performed with a multiplate spark chamber, and data were
recorded on photographic film. They estimate their angular resolution to be
~2° averaged over a typical spectrum for energies > 100 MeV. The intensity

106

OBSER VA TIONAL DA TA

h-
2

>-
<

UJ
CO

1.0

</>

r*-

~i 1 — | — i 1 r

25 < in < +20

i_L

Z) '

K 1.5
o

1.0

a^iE

J I L

-25 -15 -5 0+5 +15 +25

Figure IV.A-2. Ratio of observed line intensity of
> 100 MeV 7-rays to expected background intensity for
-25° < £u < + 20°, plotted as a function of galactic latitude
bu (from Fichteletal., 1972).

of 7-rays observed during these flights is shown plotted against the sine of
galactic latitude in Figure IV.A-3, where the bin widths have been corrected for
exposure and atmospheric contributions. Events specified as "R" refer to those
exhibiting a straight single track emerging from one of the conversion layers

GAMMA-RA YS FROM 10 MeV TO 1 GeV

107

'■ FLIGHT
12 h PAIRS
100 MeV

£. 55^^51

20
16

FLIGHT n

- PAIRS, Ey>
100 MeV

- PAIRS, Ey > 100 MeV

3 FLIGHTS
COMBINED

S 12

E «
u

w 20

>
u

S i

d 40

o 30

40
30 h
20

*-?,

WEIGHTED
MEAN

PAIRS, Ey >
100 MeV

FL
. COMBINED
PAIRS
100 MeV

3 FLIGHTS
COMBINED.
R's, Er >
100 MeV

SS H^=l^^H^^=uitF^^

3 FLIGHTS
COMBINED
PAIRS +R*»
Ey>
100 MeV

SAME EXCEPT
ZENITH ANGLE __,
<20» J^

\ >I»2.0*I0-4 y's
v cm*2 tec"' rod"'

— 4—4

-.64 -48 -.32 -.16 O .16

SIN (GALACTIC LATITUDE, b )

.32 .48 .64 .80

Figure IV.A-3. Variation in 7-ray intensity scanned across the galactic equator
near the galactic center by Frye et al. (1971a). The dashed curves in parts G
and H represent the intensity reported by Fichtel et al. (1972).

in the spark chamber. The summed data for the three flights are shown in parts
G and H of the figure and are compared with the enhancement expected along
the galactic equator, based on the intensity reported by Fichtel et al. (1972).
With the sensitivity of these measurements, it is difficult to explain why the
galactic emission was not detected.

Another observation, which has recently been published, was performed by the
group at Minnesota (Dahlbacka et al., 1973). They used an instrument incor-
porating a nuclear emulsion stack as a converter for the 7-rays and a narrow-gap
spark chamber to identify the proper events in the emulsion. With this tech-
nique an angular resolution better than 1° at energies > 100 MeV can be
achieved. The region of the galactic center was investigated during a balloon

108

OBSER VA TIONAL DA TA

flight over Australia in 1970. The number of events observed as a function of
galactic latitude near the galactic center is shown in Figure IV.A-4. The upper
plot was derived from measurements made on events located in the emulsion
stack, whereas the lower plot was obtained from measurements of the spark
chamber photographs (~ 3° resolution). The expected numbers of events are
shown by the dashed curves, assuming that the events are atmospheric in origin.
The distributions do not provide any evidence for emission from the galactic
plane, although the upper limits set by the observations are not inconsistent with
the intensities reported by Kraushaar et al., (1972) and Fichtel et al. (1972).

Q
UJ
>

to
m
o

UJ
GO

10
8
6
4
2
0
25
20
15
10
5
0

I ' 1 '

NUCLEAR
EMULSION

1 ' 1 '

f '

^ '

" '

— "i

1 • 1 ■ 1 '

■

—

_ SPARK
CHAMBER

-^.

—

1,1.

■15 -12 -9 -6 -3 0 3 6 9 12 15 18
GALACTIC LATITUDE bn

Figure IV.A-4. A histogram of the number of 7-ray events in strips parallel to
galactic plane reported by Dahlbacka et al. (1973). The upper histogram is for
events found in the emulsion and the lower one is for events observed in the
spark chamber. The dashed curves represent the expected shape for no excess
of emission from the galactic plane.

GAMMA-RA YS FROM 10 MeV TO 1 GeV 109

The third group, from the University of Southhampton (Browning, Ramsden,
and Wright, 1972), has reported evidence for point sources of 7-ray s along the
galactic plane near the center. They claim that these sources can account for
the apparent diffuse intensity observed from the plane, and furthermore, that
there is no residual diffuse intensity after the sources are subtracted. I shall
return to these results later.

The above discussion indicates that there still appears to be some disagreement
between the various experiments. Two recent measurements, made at energies
significantly below those we have discussed, have helped to clarify the situation.
Both were made over Argentina in the late fall of 1971 during the expedition
"Galaxia 71 ." The first was performed by H. Helmken and J. Hoffman of the
Smithsonian Astrophysical Observatory using a large area gas Cerenkov counter
which employed a plastic scintillator as the converter for photons above 15 MeV.
Although the instrument has good rejection properties for various backgrounds,
it suffers from its relatively poor angular resolution, ~ 30° full width at half
maximum (FWHM). This requires that in searching for continuous emission
from a possible 7-ray source, measurements must be made both on and off the
source in order to determine the background level. From two balloon flights,
Helmken and Hoffman (1973a) have reported that they detected a 3.8 a excess
from the direction of the galactic center. Due to their detector's broad angular
resolution, they were unable to determine whether the excess came from point
sources near the center, or whether it could be attributed to emission from
along the galactic plane.

The other experiment was performed by R. L. Kinzer, N. Seeman, and myself
at the Cosmic -Ray Laboratory (Chief Scientist, M. M. Shapiro) at the Naval
Research Laboratory (NRL). (A detailed description of this experiment will
be published in Astrophysical Journal and can also be found in the Proceedings
of the 13th International Cosmic Ray Conference.) Our experiment was similar
in design to that flown by the Minnesota group; it incorporated a stack of
nuclear emulsions with a wide-gap spark chamber in order to unambiguously
identify the 7-ray interaction, as well as to provide an angular resolution of
~ 1.5 . The difference between this instrument and the one flown by the
Minnesota group resides in its energy range. Whereas the Minnesota detector
had a threshold energy of about 100 MeV, our instrument had a low-energy
threshold near 10 MeV and was relatively insensitive to photons ^,200 MeV.
The lower threshold was attained by design features which restricted the
amount of material between the spark chamber and nuclear emulsion stack,
reducing the scattering of the particles considerably and permitting low-energy
electrons to be followed back into the emulsion.

110

OBSER VA TIONAL DA TA

The NRL experiment was flown to an atmospheric depth of 2.5 g-cm"2 and was
pointed in the direction of the galactic center. The distribution of 7-rays as a
function of galactic latitude was obtained from a partial analysis of events
located in the stack of emulsion and is shown in Figure IV.A-5. Plotted are

-12 -6 0 6 12 18

GALACTIC LATITUDE (deg.)

Figure IV.A-5. Distribution of observed 7-rays within (a) 3° and (b) 1° bands of
galactic latitude for 320° < fiu < 40° as reported by the NRL group. The
curves are normalized to the observed events for lb I > 6 and represent the
distribution expected for 7-rays of atmospheric origin.

the number of 7-rays observed as a function of galactic latitude for 3 - and 1 -
intervals. The curves superimposed on the histogram were normalized for
lb11 1 > 6° and show the expected number of events, assuming the 7-rays were
entirely of atmospheric origin. Evident is a significant excess of events within

GAMMA-RA YS FROM lOMeV TO 1 GeV 111

± 3° of the galactic equator; 32 events were observed whereas only 1 3 were
expected. The probability of randomly obtaining this excess of events is less
than 1 0"5 . The distribution of 7-radiation along the plane appears to be
considerably narrower (~ 3° wide) than measured by either the OSO-3 or
Goddard detectors.

From the measurements which I have discussed above, an integral spectrum for
7-rays emitted along the galactic equator in the vicinity of the galactic center
can be constructed. This spectrum is shown in Figure IV.A-6. There is good
agreement between the intensities measured by Kraushaar et al. (1972), and
Fichtel et al. (1972), near 100 MeV. As mentioned earlier, the upper limit set
by Minnesota is consistent with these measurements. Plotted at 15 MeV are
the integral fluxes determined from the NRL observations for two assumed
emission spectra, 7r°-decay from cosmic-ray collisions with interstellar gas and
a power-law representative of Compton collisions of high-energy electrons on
starlight and microwave radiation. Due to its design, the NRL instrument is
more sensitive to lower energy photons; therefore the estimated flux for a
power-law spectrum is lower than that for the harder 7r°-spectrum. Shown by
the dashed lines are extrapolations of these measurements to higher energies.
Within the uncertainties, our measurements and those at higher energies indicate
that the 7r°-mechanism can account for the observed emission; however, as
shown by the dotted-dashed curve, a spectrum with equal contributions from
both 7T° and power-law production mechanisms provides a better fit to the
observations. The flux measured by Helmken and Hoffman, if attributed
entirely to emission from the plane, is higher than our observations and requires
a much larger contribution from Compton collisions or bremsstrahlung.

The upper limit set by Frye et al. (1969), is in apparent contradiction with the
other observations above 100 MeV, assuming that the emission comes from a
narrow band along the galactic equator. This upper limit is consistent with
our measurements at lower energies only for a fairly steep energy spectrum.
However, preliminary spectral information obtained from our data appears
inconsistent with such a steep spectrum.

SUGGESTED POINT SOURCES OF GAMMA RAYS IN THE VICINITY OF
THE GALACTIC CENTER

Frye et al. (1969), reported the first evidence for emission from a point source
in the vicinity of the galactic center. The source was designated Sgr 7-I and was
reported to have been observed on each of three balloon flights (Frye et al.,
1971a). The combined statistical significance for all three observations was

112

OBSER VA TIONAL DA TA

10 -

T 5X10

i
a

X 10 -

X -5

ID 5X10

-5

1 1 1 I I I I

I I I I I I

■ Fichtel et a I. (1972)

T Dahlbacka et al.(l973)

V Bennett et a I. (1972)

D Frye et al. (1971 a)

A Kraushaar et al. (1972)

7T°aiE"

• o Share et al. (1973)

X Helmken 4- Hoffman
(1973)

J I I I i I i

J l i i i i i

20 50 100 200

ENERGY (MeV)

500

Figure IV.A-6. Measurements of the flux of 7-rays from the galactic plane
near the center of the galaxy. The NRL measurements are given for three
assumed spectra and are extrapolated to higher energies.

about four standard deviations. Subsequently, this group reported the obser-
vation of three additional sources, designated as G7 2+3, G7 341+1, and
Libra 7-I. The first two had a combined significance of about 4a over back-
ground, after data from all three flights were summed. The third source was
observed with a significance of 60 during one of their flights, but had not
been observed by them during an earlier exposure. Table IV.A-1 summarizes
the data on these possible sources. Other possible sources in the vicinity of
the galactic center have been reported by the group in Southhampton
(Browning et al., 1972); however, their evidence is of marginal statistical
significance. Data on these possible sources, as well as one mentioned by
Dahlbacka et al. (1973), are also given in the table.

The region about the galactic center was investigated with the NRL telescope
for emission of 7-rays with energies > 1 5 MeV from point sources. A

GAMMA-RA YS FROM lOMeVTO lGeV

113

/— \ „

CO #

' — 1

«-H

> — i

2 A

1— I

^^-^

V3 *

* O

t*»

«-h

«*

06 A

> 8°

«> <N

S E

o S

V©

A ">

x 2

E >S

O in

+1 ^

^ V

.3 =

0
00

-3 jo

I— c

*pi

o
O

o
CN

o
O

in
On

o
CN

■*

T-H

e>-

/"»>

■^-

t>>

+ 7

*""*

m ^fr

— i O

<*~ **■>

~— <

e>>

?■»

+ +

<* ^F

u< 0O

'+3

CN -^

po m

I— 1

<— i

00
CO

<- X

o a

£ ^

Je^

3
o

"ci

C3 O

<l>

*^ « , N

"e3

■4-J

t«

Q ja -J

t-l
O

S 3 «•

c <

~T J |

114

OBSER VA TIONAL DA TA

galactic map of the arrival directions of the observed 7-rays is shown in Figure
IV.A-7. There is a concentration of events along the galactic equator between
350° and 360° in longitude, but limited statistics preclude the possibility of
attributing, with certainty, this concentration to one or more point sources.
However, if it were due to two equally intense point sources, their estimated
fluxes above 15 MeV would each be ~ 6 X 10"s 7/cm2-s. This same region is
known, however, to contain an enhanced columnar density of atomic hydro-
gen (see for example, Garmire and Kraushaar, 1965) and therefore might be
expected to exhibit an increased emission of 7r° -decay 7-rays resulting from
collisions with high-energy cosmic rays.

4 0--

• ••

1 1 1 I 1 i 1
320°

■ 340 r • **) •••• ••*

./-20°:

-40"

%i 1 1

»•

•20" •

40"

Figure IV.A-7. Galactic map of arrival directions of 7-rays reported by the
NRL group. The RMS uncertainty in arrival direction is shown by the open
circles. Regions within the dashed curves had relative exposures > 75 percent
and > 50 percent.

GAMMA-RA YS FROM 10 MeV TO 1 GeV

115

None of the locations listed in Table IV. A- 1 for possible 7-ray sources shows a
significant concentration of events in Figure IV.A-7 (excluding Libra 7-I). A
map of events obtained from a separate exposure to Libra 7-I is shown in
Figure IV.A-8. Again, there is no evidence for an excess in the direction of the
suspected source. These exposures, therefore, failed to confirm the existence
of any of the suspected sources. Upper limits (95 percent confidence level)
placed on their intensities > 15 MeV are given in the table. Limits placed on
the fluxes above 10 MeV, also shown in Table IV.A-1, were derived from a
broad resolution survey (~ 10°) using only measurements from the NRL spark
chamber. These limits indicate that if the sources are real, they must either be
variable or their differential emission spectra must be significantly harder than
a power-law in energy « E"2 .

a>
-o

LU
Q

1 1 '

1 ' 1 ' 1 ' 1

•

■

—

• •

•

—

/^TTn

-20

• • ^

X^-Libra y~\

•

•
•

-40

. 1 ■

1 . 1 . 1 . 1

200 220 240 260

RIGHT ASCENSION (deg)

Figure IV.A-8. Map of arrival directions of 7-rays observed in a search by the
NRL group for the variable source Libra 7-L

116 OBSER VA TIONAL DA TA

O'Mongain (1973; see also Hearn, 1969) has recently studied the statistical
methods employed in analyzing data for sources of 7-ray emission. He con-
cludes that in many cases authors have underestimated the probability that
the suspected sources could have been generated by statistical fluctuations.

THE CRAB NEBULA AND ITS PULSAR

The Crab Nebula has been a target of 7-ray investigations for many years.
However, prior to the discovery of the pulsar near the center of the Nebula,
these investigations had failed to detect a significant signal from the Crab.
Upper limits to the continuous emission above 1 00 MeV were placed at about
2 X 10"5 7/cm2-s (see for example, Frye and Wang, 1969).

The existence of the pulsar gave 7-ray astronomers an added dimension to
investigate. Assuming that a large fraction of the energy emitted by the Crab
was pulsed, then measurements performed at ~ 1 ms resolution would
benefit from the reduced background. In 1969, about one month after the
observed "glitch" in the pulsar frequency, our group at NRL searched for
emission of pulsed 7-rays above 10 MeV during a balloon flight over Texas
(Kinzer et al., 1971a). The initial study was performed at about 10° resolution
and provided evidence that pulsed 7-rays were emitted in phase with the
optical peaks. Results from this study are shown in Figure IV.A-9, where
the time of arrival of events originating < 1 0° from the Crab are plotted in
part (a) against the pulsar's optical phase; for comparison, the time of
arrival of "background" events (> 10° from the Crab) is shown in part (b).
The evidence was of marginal statistical significance and prompted a more
detailed study of the data at higher angular resolution using the stack of
emulsions actively incorporated into the design of the telescope. The direc-
tions of ~ 50 percent of the events occurring close to the times of arrival of
both the primary and secondary optical peaks were determined to within
about 2° from measurements in the emulsion; however, there was no signifi-
cant concentration near the Crab (Kinzer et al., 1971b). This apparent
disagreement with our earlier suggestion could be explained, however, as
being due to the differing energy thresholds of the two samples of data.
Indeed, a subsequent study of only low-energy events observed in the spark
chamber confirmed the evidence for pulsation and furthermore, indicated
that the pulsed emission at 7-ray energies may exhibit substructure with
widths of ~0.5 ms (Kinzer et al., 1973).

This suggestion of emission at the lower 7-ray energies prompted Albats
et al. (1972), to alter their telescope in order to permit 7-rays with energies
as low as 10 MeV to be detected. Their results from an exposure to the Crab
are shown in Figure IV. A- 10 for 7-rays with energies between about 10 and
30 MeV. Two samples of data are shown which have slightly different
selection criteria. Both exhibit a striking excess within about 1 ms of the

GAMMA-RA YS FROM 10 MeV TO 1 GeV

117

30 -

c
15
w 20

Q 140

120 -

100 -

1 1 1 1

a) < 10°

J primary optical
I peak

i i

1 interpulse peak

n n

-» 1

> •— '

u " " kn"

b) > 10°

U1!^

■V-FI-fL

■"i

i_i

►

1 1 1 1 1

10 - '-,

10 15 20

Phase (-msec)

Figure IV.A-9. Number of 7-ray events > 10 MeV observed by Kinzer et al.
(1971), relative to the optical phase of NP0532. (a) Events pointing within
10° of the Crab; (b) events pointing outside 10° from the Crab. The dashed
lines give the mean numbers (N) of 7-rays/time-bin and the errors shown are

±Vn.

primary radio peak. Conspicuous by its absence, however, is any evidence of
a pulse in the vicinity of the secondary radio peak. This is to be concentrated
with measurements in the 100- to 400-keV region shown in Figure IV.A-11
and obtained by Kurfess (Kurfess and Share, 1973). In this lower-energy
domain the secondary peak and interpulse region between the primary and
secondary pulses contribute a substantial fraction of the X-rays emitted by
the pulsar. The primary X-ray peak is found to occur within 0.5 ms of the
primary optical peak. This suggests that the radiation emitted, from the
radio band up to the high energy X-ray band, originates from a region no
greater than about 150 km in extent; this distance is about 10 percent of
the radius of the speed-of-light cylinder.

118

OBSER VA TIONAL DA TA

50 -

2 40

RADIO POSITION OF
NP0532 MAIN PULSE = 3.7 MS

MAIN

— | (PULSEf- |— 1NT|R_ ,P8ULSE-^1 ^-BCKGND 0-2,19-33-

3-6

J\ -

[L \W

^BCKGND = 26.7

— BCKGND= 14.1

2 4 6 8 10 12 14 16 18 20 22 24 26 28 30 32
PHASE, MS

Figure IV.A-10. Number of 7-ray events 10 to 30 MeV observed by Albats et
al. (1972), within 15° of the Crab plotted relative to the radio phase of
NP0532.

The close relationship in the phase of the primary peak appears to
persist up to photon energies near 1 GeV and perhaps higher. Recent
results from an experiment performed by the group at Cornell are shown
in Figure IV.A-12 (McBreen et al., 1973). The measurements were made
at energies above ~ 200 MeV using a gas Cerenkov counter having a
sensitive area of about 45,000 cm2. In the energy range above 700 MeV,
significant peaks were observed at both the location of the primary optical
peak and secondary peak. In addition, the peak coincident with the primary
optical pulse appeared to have an intrinsic width ~ 0.7 ms. This is
narrower than has been observed at optical and X-ray energies. Similar
structure is also apparent in the lower energy range between 240 and 700 MeV,
but is less significant statistically. The authors point out the possible exis-
tence of pulse structure in the interpulse region between the main and
secondary peaks. Additional evidence for structure outside of the main peaks
was reported by our group at NRL (Kinzer et al., 1973).

GAMMA-RA YS FROM 10 Me V TO 1 GeV

119

392-

390-

388

w 386

384

540

i 1 1 1 1 1 1 1 1 i i r

PRIMARY OPTICAL
PULSE

Ann J

rifl

I I ' I I I I I 1 1 1 L

538 -

536

534

£ 532

530

528

1 1 1-

PRIMARY OPTICAL.
PULSE

i 1 r

i 1 1 r

20 30 40

I CHANNEL ^0.517 MSEC

Figure IV.A-11. The X-ray "light curves" for photons from
the Crab Pulsar between 100 and 400 keV observed by Kurfess
(1971) during two balloon flights on (a) Oct. 10, 1970, and
(b) Oct. 21, 1970.

Although questions remain concerning the shape of the pulsation and
possible variability, evidence is mounting supporting the existence of
7-ray pulsations from the Crab. In order to illustrate the compelling
nature of the evidence, I have summed in phase the 1-ms resolution data

120

OBSER VA TIONAL DA TA

120

100

80-

S. 60

c
o>

2 <

«_ 20-

1 ' I
— 3<r

I'M1'

-i — i — i — i — i — I — p

240-700 MeV

mJ^U

10-

M.P
I

SP
I

13.4 msec

-5a

^r^

i—i >700 MeV

4in4J

□

pi ■ ■ i ■ ■ i

l^F

J_L

I ■ I

J_l_

9 12 15 18 21 24 27 30 33

Time in Units of Period/33

Figure IV.A-12. Phase histograms of two independent samples of 7-ray events
observed with the Cornell 4.5 m2 Cerenkov telescope (McBreen et al., 1973).
The events in the upper histogram originated within 2° of the Crab Nebula,
while those in the lower histogram within 1° of the Crab. The arrival times of
the optical main pulse and secondary pulse are shown. The indicated back-
ground levels were derived from the events recorded in the intervals 0 to 9 and
24 to 33 ms.

of NRL, Case -Melbourne, and Cornell. This summation is shown in Figure
IV.A-13 where the data have been combined in 3-ms bins centered on the
main optical peak. The ratio of the average number of events in 3-ms bins
in the pulsed region to the average number in the background region is
1.30 ± 0.08. Furthermore, the bin centered on the main optical peak stands
more than seven standard deviations above the background level.

Measurements of the intensity of pulsed 7-rays are summarized in Figure
IV.A-14. The dashed line represents an extrapolation of a power-law fit to
X-ray observations of the total emission from the Crab Nebula. The low-
energy data, up to a few MeV, come from measurements with large area Nal
crystals or plastic scintillators. At higher energies visual techniques using
spark chambers were employed, with the exception of the recent measure-
ments by Helmken and Hoffman (1973b) and McBreen et al. (1973), in

GAMMA-RA YS FROM 10 MeV TO 1 GeV

121

Primory Optical Peok

Summed data: NRL > 10 MeV

Case-Melb. 10-30 MeV
Cornell > 700 MeV

Pulsed bins

Background bins

5 6 7 8 9 10
PHASE (~3ms/bin)

Figure IV.A-13. Summed phase histogram of 7-ray observations of the Crab
Pulsar taken from Figures IV.A-9, 1 1, and 12. The original data were plotted
at 1-ms resolution but are summed here in 3-ms bins in order to display
the broad features of the observations.

which gas Cerenkov counters were used. In contrast to their measurement
between 10 and 30 MeV, the higher energy measurement of Albats et al.
(1972), does not show a significant pulse within 1 ms of the main radio
peak; it does, however, show an excess in the broad pulsed region. Our
upper limit plotted at 40 MeV comes from the emulsion analysis (Kinzer et
al., 1971b). The upper limit above 100 MeV previously reported by the
Saclay-Palermo-Milan collaboration (Leray et al., 1972) has been superseded
by a recent measurement giving evidence for pulsed emission above 20 MeV
(Parlier et al., 1973). It is apparent from the mixture of upper limits (2a)
and claimed observations, that the sensitivity of the individual experiments
require about an order of magnitude improvement in order to permit detailed
studies of the Crab Pulsar.

Observations in the 100-MeV region by the Cornell group (McBreen et al.,
1973) indicate that the total emission of the Crab Nebula is consistent with
the power- law shown in Figure IV.A-14. This suggests that about half of the
0.1- to 1-GeV emission from the Crab Nebula comes directly from the pulsar.
In the 10- to 100-MeV region only upper limits or marginal evidence for

122

OBSER VA TIONAL DA TA

10'

uj 10'

▲ Greisen et al. (1973)
o Kurfess (1971)

• Orwig et al. (1971)
^~h Hillier et al. (1970)

■ Kinzer et al (1973)
A Albats et al (1972)
V Kettenring et al. (1971)

♦ Helmken et al. (1973)

D Parlier" et al (1973)
▼ Browning et al. (1971)

\
T \

1 \

• — *— -

10* 10' \0 10

PHOTON ENERGY (keV)

Figure IV.A-14. Measurements of the time-averaged pulsed intensity of
NP0532. The straight line represents an extrapolation of a power-law fit to
the total emission spectrum of the Crab at X-ray energies.

continuous emission from the Crab Nebula have been obtained (Frye and
Wang, 1969; Kinzer et al., 1971c; and Parlier et al., 1973). These limits are
consistent with the power-law extrapolation and also suggest that the pulsed
emission represents a large fraction of the total emission from the Crab Nebula.

DIFFUSE COSMIC GAMMA-RADIATION

One of the most difficult areas of experimental 7-ray astronomy is the inves-
tigation of the primary diffuse radiation. The non-visual detectors, such as
Nal and Csl crystals, which are used at low energies, are susceptible to various
backgrounds. These backgrounds can be caused by inefficiencies in anticoin-

GAMMA-RA YS FROM 10 MeV TO 1 GeV 123

cidence counters, as well as by radioactive buildup from proton spallation and
neutron interactions in the crystal and surrounding material (Pal, in press;
Kasturirangan and Rao, 1972; Dyer and Morfill, 1971; and Fishman, 1972). At
energies above 10 MeV, where both "non-visual" counter telescopes and "visual"
spark-chamber telescopes have been employed, background contamination is
still a problem. Inefficiencies in anticoincidence counters, which reject the
intense fluxes of charged particles, can be a major problem in counter tele-
scopes (Valentine et al., 1970). Although spark-chamber telescopes are
capable of discriminating against this type of background, they may be suscep-
tible to other more subtle forms, for example, local production of 7-radiation.
In addition, detectors flown on balloons within the atmosphere, or on low
orbiting satellites, must contend with the secondary atmospheric 7-radiation.

However, evidence continues to be compiled indicating the existence of a
general diffuse glow of photons from the keV region up to energies of a few
hundred MeV. A power law in energy is capable of fitting the general shape
of the spectrum up to about 1 MeV, but there are suggestions of some
departures from this spectrum. These departures include a possible steepening
in the spectrum near 40 keV (Schwartz, Hudson, and Peterson, 1970) and a
possible flattening above 1 MeV (Trombka et al., 1973).

In this section, I shall summarize the measurements made at energies above
10 MeV. Until recently, only upper limits to the intensity of the isotropic
component of cosmic 7-rays had been reported (Clark, Garmire, and Kraushaar,
1968; Frye and Wang, 1969; Bratolyubova-Tsulukidze et al., 1970; Valentine,
Kaplon, and Badhwar, 1970; Kinzer et al., 1971c). Further analysis of the data
from OSO-3 has convinced Kraushaar et al. (1972), that the residual rate which
their detector observed in directions away from the galactic plane was due to
cosmic 7-radiation. The fact that this residual rate remained constant over a
wide range of geomagnetic cutoff rigidities, and therefore charged particle
intensities, was an important consideration in the conclusion of Kraushaar et
al. (1972). Their detector also provided an indication that the spectrum of the
radiation was softer than the spectrum from either the horizon of the earth or
from the galactic plane, both believed to arise predominantly from n°- decay
7-rays.

A recent measurement from within the atmosphere using a balloon-borne
telescope has led to the suggestion by the group at the Max Planck Institut
(Mayer-Hasselwander et al, 1972) that the intensity of diffuse 7-rays in the
vicinity of 30 to 50 MeV is considerably above an extrapolation made between
X-ray data and the 100-MeV observation of Kraushaar et al. (1972). The
detector flown by the Max Planck group incorporated a multiplate spark cham-
ber with magnetic core readout. During two balloon flights over Texas in 1971,
their detector measured the intensity of 7-rays as a function of atmospheric
depth. These measurements are plotted in Figure IV.A-15 and provide

124

OBSER VA TIONAL DA TA

RESIDUAL ATMOSPHERE [gem"2]
10° 101 102

101

10°

10"1 —

10-

t 1 — i — i i i n | 1 1 — i i ii I ii r

JULY 2, 1971

101

-I 101

10°

J I I 1 1 1 1 1

J ' I I 1 1 1

10"

10°

101

RESIDUAL ATMOSPHERE [gem"

102

Figure IV.A-15. Counting rates of electron pairs as a function of residual
atmosphere observed during two balloon flights conducted by the Max Planck
Institut over Texas. The full lines are fits to the data deep in the atmosphere
and represent the growth of secondary 7-rays. The dashed curves are fits to
all the data obtained at depths < 50 g-cm'2, assuming the presence of an
extraterrestrial component of 7-rays.

evidence for a departure from the expected growth curve of atmospheric 7-rays.
By extrapolating the measurements made between ~ 50 g-cm"2 and ~ 2 g-cm"2
to the top of the atmosphere, the authors found a residual rate over 10 standard
deviations above zero. There were some differences in the absolute intensities
measured during the two flights; in addition, a fairly large uncertainty of about
0.5 g-cm"2 was present in the measurement of the atmospheric depth. However,

GAMMA-RA YS FROM 10 MeV TO 1 GeV 125

the authors did not feel that these uncertainties affected their conclusions
concerning the existence of a cosmic diffuse component. They also presented
evidence that the spectrum of this component was appreciably softer than the
atmospheric spectrum. This conclusion was reached on the basis of measure-
ments made on the distribution of the opening angles of pairs observed in the
spark chamber. However, the observed increase in the average opening angle
appears to occur abruptly at depths less than about 3 g-cm"2 and is therefore
suspicious.

During the NRL balloon flight over Argentina in 1971 , an investigation was also
made of the growth of atmospheric 7-rays as a function of depth in an attempt
to establish the existence of the primary diffuse component. Advantage was
taken of the increased cutoff rigidity (1 1.5 GV), which reduced the intensity of
secondary radiation. The data are shown in Figure IV.A-16, where the counting
rate of electron pairs is given in the left ordinate and the estimated intensity of
vertically incident 7-rays is shown on the right. Data obtained over Texas
(R > 4.5 GV) are also displayed for comparison. A linear extrapolation of the
data over Argentina gave evidence for a residual rate above the atmosphere which
was about 3.5 a above zero (Share, Kinzer, and Seeman, 1972; and preprint 1972).
An upper limit obtained from our Texas data (Kinzer et al., 1971c) is consistent
with this residual rate.

Due to the difficulties in making measurements of this kind, we made a detailed
investigation of various possible sources of background which might have simu-
lated this residue. Among those investigated were local sources for producing
the residual photons, such as the pressure vessel enclosing the system, and atmos-
pheric 7-rays incident from the horizon. From our investigations we concluded
that these sources were not likely to have contributed appreciably to the residue.
There is however, another source of background that can account for the residual
rate. In order to understand this background, we need to examine the NRL
telescope.

A schematic drawing of the NRL telescope is shown in Figure IV.A-17. Down-
ward 7-rays are detected after they convert in a stack of nuclear emulsion and
produce either Compton electrons or electron pairs which trigger the propor-
tional counter (P) and two scintillators (B) without the presence of an accom-
panying particle in any of the anticoincidence scintillators (A). The absorption-
Cerenkov counter (C) restricts detected 7-rays to those below ~ 200 MeV; it
also rejects about 50 percent of upward moving 7-rays converting in the Plexiglas
block (C) and producing upward-moving low-energy electrons which can also
trigger the telescope. These remaining upward-moving electrons are a likely
source for the residual rate of 7-rays which we observed. However, as I mentioned
above, only events appearing to be downward -moving electron-pairs were used in
our growth curve. How then can these upward-moving electrons simulate down-
ward-moving pairs? If the electrons are of low energy, they can be scattered

126

OBSER VA TIONAL DA TA

i i

III!,

, [ , ,

i i , , , , , |

i.ii

i i i i | y i i i | I I i

i | i i i i |

—

;

/ T "

4.5

GV~\ /

S \-ll.5 GV

/l/

i , , ,

.lii

llllllll

. ■ . . 1 ■ .. . ■ 1 ... ,

- 04 3

- 03 <

10 15 20 25 30 35 40 45

ATMOSPHERIC DEPTH (g-crrf2)

Figure IV.A-16. Vertical intensities of 7-rays 10 < E < 200 MeV at rigidities
> 4.5 GV and > 1 1 .5 GV as determined by the N R L group from the counting
rates of "electron pairs" observed in its wide-gap spark chamber as a function
of atmospheric depth. The lines are least-square fits and the errors shown are
statistical. (Not shown is the rate 100 ± 13/min observed at 55 g-cm'2 for
R >4.5 GV.)

appreciably in the emulsion and then emerge in the downward direction; the
event would then appear to be a downward pair of low energy.

Another source for these low-energy electrons which can enter the detector's
geometry is the splash albedo from the atmosphere. These electrons can pass
through the space between the active walls of the spar', chamber and the anti-
coincidence cup surrounding the Plexiglas block. They will be detected and
appear as downward pairs if they are scattered back out of the emulsion and
have sufficient energy to reach the bottom coincidence counters (B).

We estimate that the combined rate from both of these types of events, which
imitate downward electron-pairs, can contribute appreciably to our residual rate
of pairs above the atmosphere. For this reason, we have concluded that our
measurement must be interpreted only as an upper limit to the true diffuse
7-ray intensity.

GAMMA-RA YS FROM 10 MeV TO 1 GeV

127

-A »nTTTTTT Illllllllllllllll ^

-B i 1

"w//(W»/»\. - - V/^//k//^//<(;

10 cm.

Figure IV.A-17. Drawing of the detector used by the NRL group showing an
electron pair in the wide-gap spark chamber (S.C.). (A) plastic anticoincidence
counters; (E) emulsion stack 650 cm2 X 1.25 cm; (P) multiwire proportional
counter; (B) two plastic coincidence counters; (C) absorption Cerenkov

-2

counter of clear Plexiglas (15 g-cm" ). Cerenkov light from up-coming parti-
cles is reflected by (R) onto phototubes (not shown) imbedded in the block.

The measurements of diffuse 7-radiation above 1 MeV are summarized in Figure
IV.A-18. The solid line represents an extrapolation of the fit of X-ray data to a
power-law spectrum (Kasturirangan and Rao, 1972), while the dotted-dashed
curves represent the uncertainty in this extrapolation. Measurements above
10 MeV are typically obtained over a wide range in energy; this range is shown
by the dashed lines, and the points have been plotted at the median energy
photon detected for an assumed E"2 spectrum. The data above 1 MeV from
ERS-18 (Vette et al., 1970) were found to have been in error and have been
superseded by measurements from Apollo-15 (Trombka et al., 1973; see also
Trombka and Peterson, Chapter III. A). The measurements from Apollo-15
indicate that the energy spectrum of low-energy 7-rays flattens above about

128

OBSER VA TIONAL DA TA

10 —

10'

1 I I I I I I 1 1 — I — I I I I I I 1 1 — I — I I I

T Bratolyubova-Tsulukidze et a I. (1970) "^

A Frye and Wang (1969)

□ Kraushaar et al. (1972)

X Mayer-Hasselwander et al (1972)
Valentine et al. (1970)
Our Results Share et al. (1973)

Damle et al (1971)
Golenetskii et al (1971)
Vedrenne et al. (1971)
Vette et al (1970)
Daniel et al (1972)

j i i i i i i

J i i i i i i i

j i m i i i i

10 100

PHOTON ENERGY (MeV)

Figure IV.A-18. Measurements of diffuse cosmic 7-radiation. Energy ranges
for observations > 1 0 MeV are shown, and the fluxes are plotted at the median
energy photon detected for an assumed E"2 spectrum.

500 keV; above 1 MeV their measured intensities are still higher than the upper
limits reported by Golenetskii et al. (1971), and by Daniel, Joseph, and Lavakare
(1972).

The intensity reported by Mayer-Hasselwander et al. (1972), at higher energies
appears consistent with the data from Apollo-15. However, there may be a
systematic error in the intensity given by Mayer-Hasselwander et al. (1972).
They report that their measurement of the atmospheric 7-ray intensity is about
60 percent of the value calculated by Beuermann (1971); however, measurements
by other groups indicate that the calculated flux may be too low (Fichtel,
Kniffen, and Ogelmann, 1969; and Seeman, Share, and Kinzer, 1973). This

GAMMA-RA YS FROM 10 MeV TO 1 GeV 129

suggests that the primary diffuse intensity reported by Mayer-Hasselwander et al.
(1972), might therefore be low by about a factor of two.

The upper limit determined by our measurement over Argentina, although con-
sistent with the reported intensities, suggests that the flux of diffuse 7-rays near
30 MeV is lower than reported by either Trombka et al. (1973), or Mayer-
Hasselwander et al. (1972). The fluxes reported by these authors are considerably
above a power-law spectrum fit to both the X-ray observations and the 100 MeV
measurements of Kraushaar et al. (1972). This had led to suggestions that an
additional component may be needed to explain the results from 1 to 50 MeV.
Theoretical models for generating this additional component have been recently
summarized by Silk (preprint, 1973), Stecker (1973), and Strong et al. (1973).
Further discussion can also be found in other sections of this volume.

FUTURE OBSERVATIONS

Gamma-ray astronomy has finally emerged as an observational science. However,
as is apparent from this summary of recent measurements, an improvement in
sensitivity is required in order to permit more detailed investigations. The new
generation of satellite detectors, ESRO's TD-1A and COS-B and NASA's SAS-2,
represent the first step in providing the increased sensitivity. This is primarily
due to the longer observation periods and lower 7-ray background intrinsic in
satellite observations.

These detectors should be able to measure the energy spectrum of the diffuse
radiation >30 MeV and to begin to investigate its spatial isotropy. They should
also have the sensitivity to verify the existence of the various possible point
sources of 7-rays reported from balloon-borne observations and, furthermore,
to study their energy spectra and to establish whether or not they are variable.
There is also little doubt that these detectors will be able to investigate emission
of diffuse 7-radiation from the galactic plane and to map its distribution at
resolutions of ~ 3°. These measurements of high-energy photons from the
galactic disk, like the ones made 25 years earlier in the radio band, will substan-
tially further our knowledge of the distribution of matter, magnetic field
strengths, and cosmic-ray fluxes in the galaxy.

Continued work at balloon altitudes should be encouraged, especially in the
light of the reduced funding for "expensive" satellite programs. These balloon-
borne instruments should be designed with improved resolution in energy,
angle, and timing in order to help compensate for the atmospheric background
and to permit continued investigation of periodically pulsing objects such as
the Crab Pulsar. Improved sensitivity for balloon-borne detectors should follow
naturally from the development of high-altitude super-pressure balloons and
from observations conducted at high geomagnetic cutoff rigidities.

130 OBSER VA TIONAL DA TA

REFERENCES

Albats, P., G. M. Frye, Jr., A. D. Zych, 0. B. Mace, V. D. Hopper, and
J. A. Thomas, 1972, Nature, 240, p. 221.

Bennett, K., P. Penengo, G. K. Rochester, T. R. Sanderson, and R. K. Sood,
1972, Nature, 238, p. 31.

Beuermann, K. P., 1971,/. Geophys. Res., 76, p. 4291.

Bratolyubova-Tsulukidze, L. I., N. L. Grigorov, L. F. Kalinkin, A. S.

Melioransky, E. A. Pryakhin, I. A. Savenko, and V. Ya. Yufarkin, 1970,
Acta Phys., 29, Suppl. 1, p. 123.

Browning, R., D. Ramsden, and P. J. Wright, 1971, Nature Phys. Sci, 232,
p. 99.

, 1972, Nature, 238, p. 138.

Clark, G. W„ G. P. Garmire, and W. L. Kraushaar, 1968, Astrophys. J.
Letters, 153, p. L203.

Dahlbacka, G. H., P. S. Freier, and C. J. Waddington, 1973, Astrophys. J.,
180, p. 371.

Damle, S. V., R. R. Daniel, G. Joseph, and P. J. Lavakare, 1971, Astrophys.
and Space Set, 14, p. 473.

Daniel, R. R., G. Joseph, and P. J. Lavakare, 1972, Astrophys. and Space
Sci., 18, p. 462.

Dyer, C. S., and G. E. Morfill, 1911, Astrophys. and Space Sci., 14, p. 243.

Fazio, G. G., 1973, X-Ray and Gamma-Ray Astronomy, Proceedings of
IAU Symposium No. 55 (Madrid), H. Bradt and R. Giacconi, eds., D. Reidel
Dordrecht, Holland.

Fichtel, C. E., R. C. Hartman, D. A. Kniffen, and M. Sommer, 1972,
Astrophys. J., 171, p. 31.

Fichtel, C. E., D. A. Kniffen, and H. B.Qgelmann, 1969 , Astrophys. J.,
158, p. 193.

Fishman, G. J., 1972, Astrophys. J., 171, p. 163.

Frye, G. M., Jr., P. A. Albats, A. D. Zych, J. A. Staib, V. D. Hopper,
W. R. Rawlinson, and J. A. Thomas, 1971a, Nature, 231, p. 372.

, 1971b, Nature, 233, p. 466.

Frye, G. M., Jr., J. A. Staib, A. D. Zych, V. D. Hopper, W. R. Rawlinson,
and J. A. Thomas, 1969, Nature, 223, p. 1320.

Frye, G. M., Jr., and C. P. Wang, 1969, Astrophys. J., 158, p. 925.

GAMMA-RA YS FROM 10 MeV TO 1 GeV 131

Gal' per, A. M., V. G. Kirillov-Ugryumov, B. I. Luchkov, and O. F. Prilutskii,
1972, Soviet Phys. Uspekhi, 14, p. 630.

Garmire, G., and W. L. Kraushaar, 1965, Space Set Rev., 4, p. 123.

Golenetskii, S. V., E. P. Mazets, V. N. Ill'inskii, R. L. Aptekar', M. M. Bredov,
Yu. A. Gur'yan, and V. N. Panov, 1971, Astrophys. J. Letters, 9, p. L69.

Hearn, D., 1969, Nucl Inst, and Meth., 70, p. 200.

Helmken, H. F., and J. A. Hoffman, 1973a, Nature Phys. Set, 243, p. 6.

, 1973b, Proc. 13th Int. Conf. on Cosmic Rays,

Denver (in press).

ffillier, R. R., W. R. Jackson, A. Murray, R. M. Redfern, and R. G. Sale,
1970, Astrophys. J. Letters, 162, p. LI 77.

Hutchinson, G. W., A. J. Pearce, D. Ramsden, and R. D. Wills, 1969,
Acta Phys., 29, Supp. 1, p. 87.

Jansky, K. G., 1932, Proc. IRE, 20, p. 1920.

, 1933, Proc. IRE, 21, p. 1387.

Kasturirangan, K., and U. R. Rao, 1972, Astrophys. and Space Sci, 15, p. 161

Kettenring, G., H. A. Mayer-Hasselwander, E. Pfeffermann, K. Pinkau,

H. Rothermal, and M. Sommer, 1971, Proc. 12th Int. Conf. Cosmic Rays,
Hobart, 1, p. 57.

Kinzer, R. L., R. C. Noggle, N. Seeman, and G. H. Share, 1971a, Nature,
229, p. 187.

Kinzer, R. L., N. Seeman, and G. H. Share, 1971b, Proc. 12th Int. Conf.
Cosmic Rays, Hobart, 1, p. 45.

, 1971c, Proc. 12th Int. Conf. Cosmic Rays, Hobart,

l,p. 51.

Kinzer, R. L., G. H. Share, and N. Seeman, 1973, Astrophys. J., 180, p. 547.

Kniffen, D. A., and C. E. Fichtel, 1970, Astrophys. J., 161, p. LI 57.

Kraushaar, W. L., G. W. Clark, G. P. Garmire, R. Borken, P. Higbie, C. Leong,
and T. Thorsos, 1972, Astrophys. J., 177, p. 341.

Kurfess, J. D., 1971, Astrophys J. Letters, 168, p. L39.

Kurfess, J. D., and G. H. Share, 1973, Nature Phys. Sci., 244, p. 39.

Leray, J. P., J. Vasseur, J. Paul, B. Parlier, M. Forichon, B. Agrinier,

G. Boella, L. Maraschi, A. Treves, L. Buccheri, A. Cuccia, and L. Scarsi,
1972, Astron. and Astrophys., 16, p. 443.

132 OBSER VA TIONAL DA TA

Mayer-Hasselwander, H. A., E. Pfeffermann, K. Pinkau, H. Rothermel, and
M. Sommer, 1972, Astrophys. J. Letters, 175, p. L23.

McBreen, B., S. E. Ball, Jr., M. Campbell, K. Greisen, and D. Koch, 1973,
Preprint, to be published in Astrophys J.

Ogelmann, H., 1969, Nature, 221, p. 753.

O'Mongain, E., 1973, Nature, 241, p. 376.

Pal, Y., 1973, X-Ray and Gamma-Ray Astronomy, Proceedings of IAU
Symposium No. 55 (Madrid), H. Bradt and R. Giacconi, eds., D. Reidel,
Dordrecht, Holland.

Parlier, B., B. Agrinier, M. Forichon, J. P. Leray, G. Boella, L. Maraschi,
R. Buccheri, N. R. Robba, and L. Scarsi, 1973, Nature Phys. Set,
242, p. 117.

Schwartz, D. A., H. S. Hudson, and L. E. Peterson, 1970, Astrophys J.,
162, p. 431.

Seeman, N., G. H. Share, and R. L. Kinzer, 1973, Proc. 13th Int. Conf. on
Cosmic Rays, Denver (in press).

Share, G. H., R. L. Kinzer, and N. Seeman, 1972, Bull A.P.S., 17, p. 524.

, 1973, Proc. 13th Int. Conf. on Cosmic Rays,

Denver (in press).

Stecker, F. W., 1913, Nature Phys. Set, 241, p. 74.

Strong, A. W., A. W. Wolfendale, and J. Wdowczyk, 1973, Nature, 241,
p. 109.

Trombka, J. I., A. E. Metzger, J. R. Arnold, J. L. Matteson, R. C. Reedy,
and L. E. Peterson, 1973, Astrophys. J., 181, p. 737.

Valdez, J. V., and C. J. Waddington, 1969, Astrophys. J. Letters, 156,
p. L85.

Valentine, D., M. F. Kaplon, and G. Badhwar, 1970, Acta Phys., 29,
Supp. l,p. 101.

Vedrenne, G., F. Albernhe, I. Martin, and R. Talon, 1971, Astron. and
Astrophys., 15, p. 50.

Vette, J. I., D. Gruber, J. L. Matteson, and L. E. Peterson, 1970,
Astrophys. J. Letters, 160, p. LI 61.

B. REPORT ON GAMMA-RAY ASTRONOMY

RESULTS OBTAINED IN EUROPE SINCE THE

IAU SYMPOSIUM NO. 55

K. Pinkau*

Max Planck Institut

INTRODUCTION

Since the IAU Symposium No. 55 in Madrid in 1972 (Pal, 1973; Fazio, 1973),
little progress has been made in obtaining new results on celestial 7-rays in
Europe.

The 7-ray experiment S-133 on board ESRO's TD-1 satellite worked through
its first period of operational life from March 1972 to October 1972, when the
satellite went into hibernation. The experiment was activated in February 1973
for a second all-sky scan. For this second scan, the trigger counter thresholds
have been raised. It is hoped to thereby increase the 7-ray energy required to
trigger the experiment, thus providing a kind of "two-color" all-sky scan in
7-rays in combination with the 1972 data.

Data analysis of this experiment has been very slow and tedious and is not as
yet in such a state that first results could be presented. This is, in part, due to
the fact that TD-l's tape recorders failed after the First two months of oper-
ational life, and the tapes of the very good real-time coverage provided by
ESRO were slow in arriving. A more serious problem, however, was the severe
background problem encountered. This requires that all spark chamber images
be visually inspected, and this work has as yet not been Finished.

In what follows, the results on measurements of the diffuse flux and on the
Crab Pulsar NP 0532 are updated. The various reports on point sources dis-
covered are, in the author's opinion, of a preliminary nature and require con-
firmation by independent measurement with good statistics.

DIFFUSE FLUX

The present status of 7-ray measurements concerning the diffuse flux is well
illustrated by Figure rV.B-1 (taken from Trombka et al., 1973), where the

''Speaker.

133

134

OBSER VA TIONAL DA TA

= i i i |iiiii i i i |iiiii — i i i | him — i i 1 1 hi
COSMIC T - RAY SPECTRA

Golenelski.etol.
Vedrenne ,etol.

APOLLO 15

Moyer-

Hasselwander , et of

aJ_

i i I "ill

oson

i i i Inn

I ' ' I"-"

1.0

10
ENERGY (MeV)

1000

Figure IV.B-1. The cosmic photon spectrum derived
from the Apollo-15 data agrees with previous results be-
low 1 MeV but is well below that determined from the
ERS-18 at higher energies. Limits derived from balloon
and low altitude satellite work, despite large corrections
for efficiency and cosmic-ray produced 7-rays, are in
agreement with the Apollo results. (Trombka et al.,
1973, also Chapter 1 1 1. A)

GAMMA-RA Y ASTRONOMY RESULTS IN EUROPE 135

results of Vedrenne et al. (1971), and Mayer-Hasselwander et al. (1972), are
compared with the results of Golenetskii et al. (1971), OSO-3 (Kraushaar et al.,
1972), and Apollo-15 (Trombka et al., 1973). (See Peterson and Trombka,
Chapter III.A.)

Apart from OSO-3, the results seem to indicate that these authors find diffuse
7-ray fluxes in excess of the 25 X 103 (E/l keV)"2,1 spectrum proposed by Pal
(1973). It appears to be too early to speculate in detail about the physical
significance of this still rather uncertain result. If all these findings are confirmed,
the diffuse 7-ray spectrum would exhibit a shoulder below 100 MeV, as pointed
out by Pal (1973). It is interesting to note that 7-ray production through the
7T° -process at various red shifts in the past should integrate up to just such a
shoulder. (See Stecker, Chapter IX.A.)

In this context, a remark concerning the analysis of 7-ray data appears jus-
tified. In the domain where pair production is dominant (^ 20 MeV), 7-ray
astronomy experiments are triggered by the diverging beam of electron-positron
pairs that are created close to the trigger-telescope. Multiple scattering causes
these electrons to diverge, and the solid angle of such an instrument is not well
defined.

Furthermore, if P is the probability that one of the two electrons triggers the
instrument, the total triggering probability will be 1 - (1-P)2, thus causing a
significant enhancement of the probability that a 7-ray incident at large zenith
angles will actually trigger the counter because one of the electrons was scat-
tered into the sensitive cone of the telescope.

These considerations show that the energy-angle response function (a(E,0)) of a
7-ray counter telescope cannot be separated into one function of energy and
another one dependent on angle only, as was assumed in the case of the OSO-3
data analysis (Kraushaar et al., 1972). Rather, the effective solid angle £2 defined
as

$2tt sin 9 a(E,0) dd

f2 =

a(E,0 = 0)

will remain a function of energy, increasing with decreasing energy. This has the
interesting consequence that the ratio of line flux factor to isotropic flux factor
G^g/G. as defined by Kraushaar et al. (1972), will depend upon energy and
thus on the assumptions on the line flux and isotropic flux energy spectra,
respectively. This has to be borne in mind when comparing the results of OSO-3
of the galactic plane emission with that of high galactic latitudes.

CRAB PULSAR NP 0532

In recent months, two results have been published that appear to establish the
Crab Pulsar spectrum in the 10- to 100-MeV energy range. They are the meas-
urements of Albats et al. (1972), and of Parlier et al. (1972). Figure IV.B-2

136

OBSER VA TIONAL DA TA

Figure IV.B-2. Crab Nebula spectrum (from Parlier et al., 1972).

shows the Crab Pulsar spectrum as presented in the paper of the Saclay -Milan-
Palermo group (Parlier et al., 1972, see there for the references. Measurement
point No. 20 is that of Albats et al., 1972).

These results are significant in two respects:

• First, the ratio of the continuous to pulsed flux from the Crab is about
a factor of 6 at 0.1 MeV and this decreases to less than half that value
at 20 MeV. Indeed, all the flux >20 MeV could be pulsed.

• Secondly, while the interpulse appears to be dominant in the low
energy 7-ray domain (Kurfess, 1971), both Albats et al. (1972),
and Parlier et al. (1972), claim that the main pulse is dominant in
their results. It would certainly be very interesting to study, with
good statistics, the transition between these two different results

in the 1- to 10-MeV region. (See also Chapters IV. A, IV.C and V.A.)

CONCLUSION

Gamma-ray astronomy has been, and still is, a slowly developing branch of
science. This is due to the very great experimental difficulties. Furthermore,
7-ray observations cannot be carried out from very simple, or small, spacecraft.
It appears that the development of the field has also been slowed down by the
comparatively large amount of time lost in the effort to obtain access to
satellite space.

GAMMA-RAY ASTRONOMY RESULTS IN EUROPE 137

REFERENCES

Albats, P., G. M. Frye, A. D. Zych, O. B. Mace, V. D. Hopper, and J. A. Thomas,
1972, Nature, 240.

Fazio, G. G., 1973, X-Ray and Gamma-Ray Astronomy, Proceedings oflAU
Symposium No. 55, (Madrid), H. Bradt and R. Giacconi, eds., D. Reidel,
Dordrecht, Holland.

Golenetskii, S. V., E. P. Mazets, V. N. Il'inskii, R. L. Aptekar, M. M. Bredov,
Yu. A. Gur'yan, and V. N. Panov, 1971 , Astro. Letters, 9, p. 69.

Kraushaar, W. L., G. W. Clark, G. P. Garmire, R. Borken, P. Higbie,
C. Leong, and T. Thorsos, 1972, Astrophys. J., Ill, p. 341.

Kurfess, J. D., 191 1 , Astrophys. J. Letters, 168, p. L39.

Mayer-Hasselwander, H. A., E. Pfeffermann, K. Pinkau, H. Rothermel, and
M. Sommer, 1912, Astrophys. J. Letters, 175, p. L23.

Pal, Y., 1973, X-Ray and Gamma-Ray Astronomy, Proceedings of

IAU Symposium No. 55, (Madrid), H. Bradt and R. Giacconi, eds., D. Reidel
Dordrecht, Holland.

Parlier, B., B. Agrinier, M. Forichon, J. P. Leray, G. Boella, L. Marashi,
R. Buccheri, R. Robba, and L. Scarsi, 1972, to be published in Nature
Phys. Sci.

Trombka, J. I., A. E. Metzger, J. R. Arnold, J. L. Matteson, R. C. Reedy, and
L. E. Peterson, 1913, Astrophys. J., 181, p. 737.

Vedrenne, G., F. Albernhe, I. Martin, and R. Talon, \91\,Astron. and
Astrophys., 15, p. 50.

C. PRELIMINARY RESULTS ON SAS-2

OBSERVATIONS OF >30 MeV

GAMMA RADIATION

D. A. Kniffen,* C. E. Fichtel, and R. C. Hartman

Goddard Space Flight Center

INTRODUCTION

It was Morrison (1958) who first pointed out that the low interaction cross
section of the high-energy 7-ray makes it a unique and valuable medium for
obtaining information on many of the major energy transfers which take
place in the universe. Furthermore, its charge less state allows the information
to be related to the regions in which the processes are occurring. In papers
presented at this conference, Stecker, Ginzburg, and Clayton have pointed
out that the spectra obtained from the observations of energetic 7-radiation
may provide most important information concerning a number of astro-
physical problems. These problems include the study of the distribution
of high energy nuclei in the universe in space and time, the possible existence
of antimatter on a universal scale, the origin of the > 50 MeV galactic
emission observed by Kraushaar et al. (1972), and other phenomena unique
to large scale astrophysical bodies. In addition, the field of high-energy
7-ray astronomy provides an opportunity to extend our knowledge of the
electromagnetic phenomena for diffuse and discrete source X-ray emission
to high energies.

Within our own galaxy, high-energy 7-rays speak directly to the presence
of energetic protons within discrete sources and in the galaxy as a whole
through the broadly peaked but distinctive spectrum of 7-rays produced
by the high-energy nucleons interacting with other nucleons. In this way,
the cosmic-ray distribution throughout the galaxy may be studied as well
as the high-energy particle gas surrounding individual objects from which
cosmic rays have come. The picture that emerges will significantly aid in
the understanding of the dynamics of our galaxy and the origin of energetic
charged particle cosmic rays.

"Speaker.

139

140 OBSER VA TIONAL DA TA

Beyond our galaxy 7-ray observations serve as an indicator of conditions
existing in the cosmological past. In an expanding model of the universe,
the density of matter was much greater in the past than it is observed to be
in the present epoch. Two of the processes expected to be the most likely
producers of 7-radiation on the universal scale are nuclear interactions of
energetic cosmic radiation with the intergalactic gas and nucleon-antinucleon
annihilation. Both processes produce a characteristic 7r°-decay 7-ray spec-
trum in the rest frame, but the energy is degraded by the cosmological red
shift caused by the expansion of the universe. Hence, 7-ray astronomy can
address itself directly to the subject of cosmology.

Also expected to be important contributors to 7-ray production are the
electromagnetic interactions important in X-ray astronomy, including the
interactions of energetic electrons with matter (bremsstrahlung), with
cosmic photon fields (Compton scattering), and with magnetic fields
(magnetic bremsstrahlung).

Within discrete stellar objects, in addition to these mechanisms, there are
other processes unique to the objects which may produce detectable levels
of 7-radiation. Examples of such possibilities are the radioactive decay of
the nucleosynthesis products as they are explosively ejected in supernovae
(Clayton, 1973) and short intense burst of energetic photons emitted in
the hydromagnetic shock wave following a stellar collapse (Colgate, 1968).
The detection of 7-rays and the determination of their spectral characteris-
tics during such events would provide most important clues to the validity
of the theories which predict them.

The potential significance of 7-ray observations has led a large number of
groups to develop a variety of detectors for the search of this rare photon
in a very high background of energetic charged particle cosmic rays. The
first unambiguous positive observations of extraterrestrial 7-rays above a
few 10's of MeV was made by Kraushaar et al. (1972), with their OSO-3
7-ray detector, launched in 1968. This pioneering experiment measured
a general diffuse flux and an enhanced emission from the galactic disk of
7-radiation above 50 MeV. Theoretical models for the origin of these
observational fluxes have been difficult to obtain because of the limited
angular and spectral resolution of the OSO-3 experiment. Share (Chapter
IV.A) has reviewed other results obtained from a large number of detec-
tors flown from balloons and satellites. Positive observations have been
obtained for the diffuse flux, the galactic disk emission, and a large number
of discrete sources, but conflicting evidence between experiments in some
cases and marginal statistics in others has left a generally uncertain picture
with the possible exceptions of pulsed 7-ray emission from the Crab Nebula
Pulsar NP 0532 and the galactic plane emission.

SAS-2 OBSER VA TIONS 141

In March of 1972, the first of the second generation of satellite 7-ray
experiments was launched aboard the ESRO TD-1. The experiment
consisted of a nine-deck vidicon spark chamber 7-ray telescope. On
November 15, 1972, the SAS-2 was launched into orbit with a larger
32-deck magnetic core digitized spark chamber. These instruments should
provide the sensitivity and angular and spectral resolution with the inherently
low background of a satellite experiment needed to address many of the
important questions in 7-ray astronomy.

In this paper we will give a description of the SAS-2 detector and present
some of the preliminary results we have obtained.

EXPERIMENT

Figure IY.C-1 is a schematic view of the SAS-2 telescope, a 32-deck spark
chamber with a scintillator-Cerenkov counter charge d-par tide triggering
telescope and a large plastic scintillator anticoincidence dome surrounding
the entire experiment. Each spark chamber module is separated from the
next by a 0.03 radiation length tungsten-pair production plate. The tungsten
plates serve as scattering plates for the electrons following their formation,
allowing the energy of each electron and hence of the incoming 7-ray, to
be determined by analysis of the multiple scattering. This information is
also used to obtain a weighted bisector of the pair for determining its
arrival direction in spark chamber coordinates. A large number of thin
plates are used so that the electron pair can be clearly identified and the
arrival direction of the 7-ray can be accurately measured. The signature
required for a trigger of the spark chambers is for a particle to pass
undetected through the anticoincidence dome and to pass simultaneously
(within about 500 ns) through the two elements of one of the four scin-
tillator-Cerenkov charged particle telescopes. This coincidence triggers
the application of high voltage across the spark chambers and initiates the
readout system.

Figure 1Y.C-2 shows a photograph of a single wire grid module containing
two planes of 200 wires each on opposite sides of the frame. The wires
within a plane are parallel and orthogonal to the wires on the opposite
plane. Each grid wire is threaded through a ferrite core contained on a
shelf on the side of the frame. Two additional wires are threaded through
each core to readout those set during an event. As a spark breaks down
along the ion path remaining along the trajectory of a charged particle,
the current flows along one or more affected wires in each plane of the
grid, setting one or more cores. The readout of such set cores thus pro-
vides the coordinates of the charged-particle passage through that modular
deck.

142

OBSER VA TIONAL DA TA

UPPER SPARK
CHAMBER

SCINTILLATOR

LOWER SPARK

CHAMBER

CERENKOV

COUNTER

ASSEMBLY

(4 UNITS)

ELECTRONICS
BOXES (4)

GUARD

SCINTILLATION

COUNTER

LIGHT PIPES (4)

PHOTO-
MULTIPLIERS (8)

PH0T0MULTIPLIERS(4)

SAS-B GAMMA RAY EXPERIMENT
Figure IV.C-1. Schematic diagram of the SAS-2 7-ray spark-chamber telescope.

If the distribution of set cores is plotted separately for each of the two
orthogonal planes, a picture is obtained such as that shown in Figure IV.C-3,
which is a reproduction of a 1 6-mm microfilm frame of the two orthogonal
views of a 7-ray pair production event. The scale for the vertical axis is
compressed by a factor of three relative to the horizontal so incoming angles
are exaggerated.

The flight unit was given a preflight calibration at a tagged photon facility
established for this purpose at the 170-MeV electron synchrotron at the
National Bureau of Standards in Gaithersburg, Maryland. The beam pro-
vides monoenergetic photons selectable over the 30- to 150-MeV energy
interval. A very extensive calibration is currently underway using the
essentially identical flight spare experiment unit. Until this calibration is
complete, the results must be considered preliminary and flux and intensity
values should be considered to be no better than about a factor of 1.5.

The characteristics of the telescope include an area of about 540 cm2, a
solid angle of about 1/3 sr, and an asymptotic high energy pair production
efficiency of 29 percent. Timing accuracy of about 1 to 2 milliseconds
allows a search for periodic emission. Arrival directions for 100-MeV
7-ray s can be measured to about two degrees at 100 MeV. The energy
threshold is about 30 MeV, although it is not sharp. Differential energy
measurements can be made on 30- to 200-MeV 7-rays, and integral fluxes
obtained for > 200-MeV 7-rays. A more detailed description of the
experiment has been given by Derdeyn et al. (1972).

SAS-2 OBSER VA TIONS

143

Figure IV.C-2. Photograph of a single wire grid spark-chamber module with
two planes of 200 parallel wires. The direction of the wires in one plane is
orthogonal to that in the opposite plane. Each wire is terminated through a
ferrite core.

The experiment was launched as the sole experiment aboard the second of
the Small Astronomy Satellites (SAS-2) on November 15, 1972. The
orbital trajectory is essentially equatorial and approximately circular at
a height ranging from 440 to 610 km above the earth's surface. Figure
IV.C-4 gives an artist's concept of the telescope, surrounded by a gold
colored thermal blanket, sitting atop the spacecraft control section. The
satellite is spin stabilized with magnetic torquing of commandable electro-
magnets against the earth's magnetic field providing steering to any
selectable point on the sky. Attitude is determined by a magnetometer-
sun sensor combination and, to more precision, by a star sensor which is
capable of determining the telescope pointing direction to about a quarter
of a degree, thus allowing the directions of the 7-rays in spark chamber
coordinates to be transformed into celestial and galactic coordinates.

144

OBSER VA TIONAL DA TA

25 50 75 100 125 150 175

YEAR

DAY

MIN

SEC

MSEC

520

~": : " : ' . „i ' : : . : . . <-.. *.*.:*. .

-+ »- — EE — EE

■ ■1 1

100

125

150

175

100

125

150

175

— . —

0 25 50 75 100 125 150 175

Figure IV. C- 3. A microfilm plot of an event that presents two orthogonal
views of the digitized trajectory of a pair of electrons produced by a 7-ray
interacting with one of the tungsten plates between the 32 spark-chamber
modules. The x's and y's denote cores set due to the passage of charged
particles in the two orthogonal views. The vertical axis is compressed by
nearly a factor of 3 relative to the horizontal, causing angles to be over-
emphasized.

The viewing program has been chosen so as to examine each portion of
the sky with about a one week exposure, with early emphasis on those
regions of the sky expected to be most interesting in 7-rays. Figure
IV.C-5 gives a view of the sky with the regions of the sky examined to
date with the 1/3 sr field-of-view denoted by the cross-hatched area. Second-
week exposures have already been obtained on the galactic center region as
well as the anticenter, Crab Nebula region as denoted by the double cross-
hatched area.

DATA ANALYSIS AND REDUCTION

SAS-2 data are recorded at a one-k bit/s rate on redundant onboard
continuous-loop tape recorders. Once per orbit the recorded data

SAS-2 OBSER VA TIONS

145

Figure IV.C-4. An artist's concept of the SAS-2 in orbit. The experiment
surrounded by a thermal blanket sits atop the spin-stabilized spacecraft.
Attitude is controlled by magnetic torquing.

Figure IV.C-5. Regions of the sky in galactic coordinates viewed by SAS-2
through May 21, 1973. The cross-hatched regions are those viewed during
the first 5 weeks after launch.

146 OBSER VA TIONAL DA TA

are transmitted at a 20 k bit/s rate to a tracking station loaded near Quito,
Ecuador. Real-time data taken before and after the recorder playback is
used to correlate the spacecraft clock with the station clock. This provides
time in the data stream that is accurate to better than 2 ms in absolute time.

The data stream contains, in addition to spacecraft time, the spark-chamber
event data, experiment and spacecraft control section housekeeping data
(counter rates, voltages, currents, temperatures, pressures, and so forth),
and aspect data from a digital solar aspect sensor, two fluxgate magnetome-
ters, and an N-slit star sensor. Three orbits of data per day are transmitted
via transmission links directly to the Goddard Space Flight Center (GSFC)
to determine the aspect for the purposes of planning any necessary maneu-
vers. Maneuvers of the spin-stabilized spacecraft are accomplished by
command of electromagnet torquing coils which provide fields that interact
with the terrestrial field to provide maneuvering rates of up to about 5.0° a
minute. Analog magnetic tapes of the remaining orbits are shipped to GSFC
where time is correlated and the data placed in proper time sequence with any
overlapping data eliminated. The magnetometer, sun sensors, and star sensors
are used to determine the spark-chamber pointing direction to an accuracy
of about 0.25°.

Analysis of the spark-chamber data is made by an automatic pattern recog-
nition device designed to recognize the readout patterns produced as
7-rays interact in the tungsten plates to create electron-positron pairs
(Figure IV.C-3). An alternate mode for analysis of the event data is made
by interactive editing of the events with a graphics display unit. The results
available at the present time have been obtained using the latter mode.

Events selected for editing have been carefully chosen to ensure that no
ambiguities will be introduced into the measured fluxes by misidentification
of spark chamber events. The selection is based on the following criteria:
(1) only intervals that contain data taken when the spark chamber axis
points away from the earth are chosen for analysis; (2) only 7-ray pair
events are selected; (3) events that can masquerade as pair events as a
result of interactions in the material of the side walls of the spark chamber
are rejected for analysis; (4) events that set cores in the top grid are
rejected to provide a veto for the rare events which form in the small
amount of material between the coincidence counter and the spark cham-
bers; and (5) 7-rays arriving at very large angles with respect to the detector
axis are not included in the analysis. Edited events are automatically
processed to obtain the energy and chamber arrival direction of each
observed 7-ray according to procedures developed in the analysis of
balloon data as described by Fichtel et al. (1972). The directional infor-
mation is then combined with attitude and orbit data to provide the 7-ray
arrival direction in celestial, galactic, geographic, and geomagnetic coordinates.

SAS-2 OBSER VA TIONS 147

Events with zenith angles greater than 90° with respect to the outward radius
vector of the satellite position are rejected from further consideration for the
celestial analysis, safely avoiding the terrestrial horizon which lies at zenith
angles greater than 1 10.0°.

The sensitivity of the telescope to each region of the sky is determined by an
automatic analytic program that checks against all status conditions that
affect sensitivity. In addition, the accumulation is made differentially in
time in order to include instantaneous detector live time and to exclude
those portions of the sky occulted by the earth.

RESULTS

The results can be classified by three categories: diffuse, presumably extra-
galactic 7-rays coming from regions of the sky not associated with the
galactic plane; 7-radiation from the galactic plane; and discrete sources of
energetic 7-rays.

Diffuse 7-Radiation

For the regions of the celestial sphere, which we have examined thus far,
there seems to be a weak, but finite component of high-energy 7-rays which
exists for regions away from the galactic center. OSO-3, even with its much
smaller sensitivity 1 .6 (cm2 • sr efficiency) compared to about 30 (cm2 • sr
efficiency) for SAS-2 above 100 MeV, also indicated a finite, apparently
constant diffuse flux for regions of the sky which were far enough from
the galactic plane that no portion of the relatively wide angle of the OSO-3
detector (~35°) overlapped the galactic plane. From observations that
SAS-2 has made, it now appears that in the region -20° < Cn < +20°, and
bjj > 0 , the flux reaches a constant level at least for bn > +1 5° . The data
reported here come from the region of the sky centered at (Cu = 0, bn = +25°).
The diffuse energy spectrum is presented in Figure IV.C-6. Notice that the
spectrum is quite steep, steeper than other 7-ray spectra observed on SAS-2
or the ear Uer balloon work of the Goddard group (for example, Fichtel
et al., 1972), including data on the galactic center region, the Crab, and the
atmospheric secondary spectrum, upward or downward. The integral flux
above 100 MeV is (3.9^) X 10"5/(cm2 • sr • s) consistent with the OSO-3
result of (3.0 ±0.9) X 10"5/(cm2 • sr • s) averaged over all regions of the sky
(Kraushaar et al., 1973). (Value corrected according to private communica-
tion with G. W. Clark, 1973.) The OSO-3 experiment did not measure the
energy spectral shape.

148

OBSER VA TIONAL DA TA

Figure IV.C-6. The diffuse 7-ray spectrum measured
by SAS-2 for regions of the sky analyzed with
Ib.J > 20°. See text for a discussion of the specific
region. For the present, a factor of uncertainty of
1.5 should be attached for each point.

The Galactic Plane

SAS-2 has confirmed the high intensity of 7-rays coming from the galactic
center region. The emission region extends along the galactic plane for at
least 60° to 70° centered about the galactic center and is no wider than
9° full-width half-maximum for 100 MeV 7-rays and could be narrower,
since there is still a final correction to be applied to the SAS-2 attitude
data. The average intensity level for this region is about 1 .2 X 10" /
(cm2 • rad • s), to which an uncertainty factor of 1 .5 is attached until the
SAS-2 calibration is complete. Whereas the average energy spectrum
from this region is much harder than the diffuse radiation, the number of
7-rays between 30 and 60 MeV relative to the number above 100 MeV is
inconsistent with a pure n° -decay component. Apparently, there are
other components with softer spectra. Because the SAS-2 aspect has not
yet been determined with sufficient accuracy, at present the SAS-2 data
would allow either a diffuse radiation or a sum of point sources for the
soft component; however, there would have to be several (at least six)

SAS-2 OBSER VA TIONS 149

point sources or there would have been a greater nonuniformity than was
observed.

Discrete Sources

High-quality attitude data are not yet available for a detailed study of
discrete sources with SAS-2. However, a positive flux is detected for the
Crab Nebula on the basis of an analysis of a sixth of the data available.
A complete analysis of the data combined with accurate attitude data
will allow a study of the energy spectrum of the Crab Nebula emission
and the possible periodic pulsations from NP 0532.

No evidence is obtained for 7-ray emission from Sco X-l, with 95-percent
confidence limits based on about a fourth of the data of 1 .7 X 10"6/
(cm2 • s) for 7-rays above 40 MeV and 1 .0 X 10'6/(cm2 • s) above 100 MeV.
A full analysis of a typical one-week exposure will allow 95-percent
confidence limits of about 2 to 3 X 1 0"7/(cm2 • s) for sources for which
no positive indication is obtained.

DISCUSSION

Diffuse Radiation

Figure IV.C-7 shows that the isotropic 7-radiation for |bnl > 20° exhibits
an enhancement relative to the single extension of the power-law spectrum
valid in the X-ray region from 1 to 20 MeV and then a rapid decrease in
intensity in the region from 40 to 200 MeV with an apparently reasonably
smooth curve through the entire 7-ray region. Until more SAS-2 data from
many regions of the sky have been analyzed with the full angular resolu-
tion, it is not possible to say that the radiation is truly uniform over the
sky and uniform on a fine scale also. However, it seems a plausible
hypothesis to assume that the regions examined thus far by SAS-2 are
representative and to consider the possible origin of the radiation. There
is, of course, the possibility that radiation is the sum of many, many weak
sources of unknown origin. However, there are at least two other possi-
bilities: one that the radiation comes from diffuse electrons interacting
with matter, photons, or magnetic fields, and the other is that the 7-rays
are of cosmological origin.

With regard to the diffuse electron possibility, bremsstrahlung seems
unlikely because, at an energy where an increased slope would be expected,
1 to 1 0 MeV, due to an increasing rate of energy loss, the inverse is
observed. For both synchrotron and Compton radiation, a power-law
electron energy spectrum leads to a power-law photon spectrum, but
with a different slope. The observed photon spectrum would imply a

150

OBSER VA TIONAL DA TA

10"

I0"a r

I Mll|

I I I I Ml|

1 1 1 ■ I ■ 1 1 1

TOTAL THEORETICAL

APOLLO 15 DATA

10"

♦ GOLENETSKI ttol(l97l)
-oVEDRENNE et at (1971)

• MAYER-HASSELWANDER

et al,(l972)
APOLLO I5'TR0MBKA

•t ol,(l973)
SAS-I^THIS WORK

SAS-I

E DIFFUSE y-RAY SPECTRUM

J ' I I I ml I I l l l nil JLl ' I I I ill

0.1

Ey (MeV)

10'

ioJ

Figure IV.C-7. Diffuse radiation observed by several
experiments. Also shown is the linear extrapolation of
the X-ray data (solid line) and the spectrum predicted
by Stecker et al. (1971), for 7-rays produced by the
decay of neutral pions resulting from cosmic-ray
interactions with interstellar matter in the cosmo-
logical past.

similarly shaped parent electron spectrum which would have features that are
at least as sharp. There is no reason to expect such a spectral shape for
diffuse electrons, although there is no experimental knowledge of the
electron spectrum in the relevant energy range. Further, especially in
the synchrotron case, the intensity seems too high to be consistent with
reasonable estimates of the interstellar parameters.

SAS-2 OBSER VA TIONS 151

Of the pure 7-ray cosmological hypotheses, there are two of which the authors
are aware that seem to be possible candidates. They are the cosmic ray
interstellar matter interaction model and particle -antiparticle annihilation
in the baryon-symmetric big-bang model. In both theories, the resulting
7-ray spectrum, which is primarily due to 7r°-decay, is red shifted substan-
tially due to the expansion of the universe. These theories are discussed by
Stecker et al. (1971), and Stecker in these proceedings (Chapter IX. A).

In an expanding model of the universe, the density of matter is much
greater in the cosmological past than it is observed to be in the present.
However, since the 7-radiation produced reaches us from large distances,
the energy of the photons is degraded by the cosmological red shift caused
by the expansion of the universe. One curve developed by Stecker (1969)
involving red shifts up to about 100 is shown in Figure IV.C-7. The
theoretical curve is seen to agree with experimental data reasonably well.

An alternative attempt to explain the 7-radiation through red-shifted 7-rays
from 7T° -decay arises from the big-bang theory of cosmology with the prin-
ciple of baryon symmetry. Harrison (1967) was one of the first to propose
a model of this type. Omnes (1969), following Gamow (1948), considers
a big-bang model which is initially at a very high temperature and density,
and then shows that, if the universe is baryon symmetric, a separation of
matter from antimatter occurred at T > 30 MeV. The initial phase separation
of matter and antimatter leads ultimately to regions of pure matter and pure
antimatter containing masses of the size of galaxy clusters. Stecker, Morgan,
and Bredekamp (1971) have predicted the 7-ray spectrum which would be
expected from annihilation at the boundaries of such clusters from the
beginning of their existence. This spectrum is very similar to the one shown
in Figure IV.C-7 in the energy range for which data exists, and is not
included in the figure for that reason.

Galactic Plane Radiation

Since the final-attitude data did not exist for SAS-2 at the time this article
was written, discussion of the galactic center region must be limited to a
summary of the broad features observed by SAS-2: (1) The enhancement of
the galactic radiation in the region of the galactic center observed by OSO-3
is confirmed; (2) It is 60° to 70° in length along the plane and no more than
9° wide; (3) The energy spectrum is not a pure it0 -spectrum, but rather it
also contains an enhanced flux below 70 MeV relative to that expected from
a 7r° -spectrum; and (4) The enhancement is not due just to a few point
sources, although it could, of course, be due to a large number of point sources.

152 OBSER VA TIONAL DA TA

Discrete Sources

A discussion of the significance of discrete sources must await further data
analysis; however, the sensitivity of SAS-2 should allow detailed study of a
number of discrete sources and should allow us to place upper limits to the
flux of objects with little or no emission that are almost two orders of
magnitude lower than the existing upper limits.

REFERENCES

Colgate, S. A., 1968, Canadian J. Phys., 46, p. S476.

Derdeyn, S. M., C. H. Ehrmann, C. E. Fichtel, D. A. Kniffen, and R. W. Ross,
1972, Nucl. Inst. andMeth.,98, p. 557.

Fichtel, C. E., R. C. Hartman, D. A. Kniffen, and M. Sommer, 1972,
Astrophys. J., 1 71 , p. 31 .

Gamow, G., 1948, Phys. Rev., 14, p. 505.

Golenetskii, S. V., E. P. Mazets, V. N. Il'inskii, R. L. Aptekar', M. M. Bredov,
Yu. A. Guryah, and V. N. Panov, 197 '1, Astrophys. J. Letters, 9, p. L69.

Harrison, E. R., 1967, Phys. Rev. Letters, 18, p. L101 1.

Kraushaar, W. L., G. W. Clark, G. P. Garmire, R. Borken, P. Higbie, V. Leong,
and T. Thorsos, 1913, Astrophys. J., Ill, p. 341.

Morrison, P., 1958,7/ Nuovo Cimento, 7, p. 858.

Mayer-Hasselwander, H. A., E. Pfeffermann, K. Pinkau, H. Rothermel, and
M. Sommer, 1972, Astrophys. J. Letters, 175, p. L23.

Omnes, R., 1969, Phys. Rev. Letters, 23, p. L38.

Stecker, F. W., 1969, Astrophys. J., 157, p. 507.

Stecker, F. W., D. L. Morgan, and J. Bredekamp, 1911, Phys. Rev. Letters,
27, p. L1469.

Trombka, J. I., A. E. Metzger, J. R. Arnold, J. L. Matteson, R. C. Reedy,
and L. E. Peterson, 1973, Astrophys. J., 181, p. 737.

Vedrenne, G., F. Albernhe, I. Martin, and R. Talon, \91\,Astron. and
Astrophys., 15, p. 50.

Chapter V

A. OBSERVATIONS OF HIGH-ENERGY
GAMMA RAYS

G. G. Fazio*

Smithsonian Astrophysical Observatory and Harvard College Observatory

INTRODUCTION

At energies above 101 * eV, the predicted fluxes of cosmic 7-rays from discrete
sources are so small (< 1 m2 day" l ) that it becomes impractical to measure
these fluxes with detectors in high-altitude balloons and satellites. However,
this radiation can be observed, indirectly, with ground-based instruments.
When high-energy 7-rays strike the earth's atmosphere, they generate a shower
of particles that in turn emit a burst of Cerenkov light. In principle, a ground-
based observer, using rather simple apparatus, can record the intensity and
direction of either the particles or the Cerenkov light, or both, and hence
determine the energy and arrival direction of the incident 7-ray photon. In
the energy region between 101 1 and 101 3 eV, the shower particles are absorbed
in the atmosphere; therefore, only the Cerenkov light technique can be used.
It is this energy region in which most experiments are done and about which
I will describe our recent results. Let me begin by first describing the instru-
mentation used in these experiments because it has bearing on the interpreta-
tion of the results.

INSTRUMENTATION

In the absence of sufficiently strong sources of cosmic 7-rays, the design of
experiments must be based on theoretical models of the properties of the
Cerenkov light generated by air showers. These studies have been done by
Zatsepin and Chudakov (1962), Zatsepin (1965), Long (1967), Rieke (1969),
and Bosia, Maringelli, and Navorra (1972). In general, at 101 1 eV these
light bursts of 3-ns duration have angular diameters of the order of 0?5 but,
when viewed away from the shower axis, are elongated and displaced in angle
toward the shower maximum. The number of Cerenkov photons per unit area
at ground level is rather constant within 130 m of the shower core and falls
rapidly beyond this radius. Therefore, a light detector with sufficient sensitivity

■Speaker. ^

154 OBSER VA TIONAL DA TA

will be able to detect showers over an area of 5 X 104m2 . Thus, the principal
advantages of this technique are the combination of a large sensitive area with
good angular resolution.

There are also disadvantages. The basic one is that there is no equivalent of
the anticoincidence counter to remove those showers generated by charged
cosmic-ray particles; these are at least several hundred times more numerous.
Another disadvantage is that the technique can be used only on clear, moon-
less nights.

To obtain the maximum possible light-collecting power and hence the mini-
mum possible threshold energy, the Smithsonian Astrophysical Observatory
(SAO) constructed a light reflector of 10-m aperture, mounted on an alt-azimuth
antenna positioner. The reflector is located at the 2300-m level of Mount
Hopkins, Arizona (Figure VA-1). Based on theoretical calculations, the re-
flector, when used as a 7-ray detector pointed to the zenith, has a threshold
energy of 9 X 101 ° eV, an effective sensitive area of 1 .3 X 104 m2 , and an
effective angular resolution of 1°. The maximum shower count rate at the
zenith is about 400 min" 1 . At angles away from the zenith, the threshold
energy and the effective collecting area increase. There is no energy resolution
in the integral counting mode other than this method of varying the threshold
energy. Attempts are being made to achieve energy resolution by measuring
the intensity of each Cerenkov light burst.

The primary cosmic radiation generates an isotropic background source of
Cerenkov light bursts; hence, a 7-ray source can be distinguished by an increase
in the number of showers detected in the direction of the suspected source.
Two observational techniques have been used to detect this anisotropy: the
drift -scan mode and the tracking mode. In the drift-scan mode, the reflector
is aligned 20 to 30 min of right ascension ahead of the suspected source; the
earth's rotation then sweeps the field-of-view of the detectors over the source.
Many drift scans must be accumulated on each object, since the expected
anisotropy is less than 1 percent. Although the drift -scan has advantages in
terms of stability and ease of operation, it is most inefficient because less
than 20 percent of the observing time is spent on the source. To increase the
on-source observing time, the tracking mode was developed. In this mode, two
phototubes are located at the focus of the reflector and separated by 4? 2. The
reflector then tracks the source in such a manner that one phototube views the
source while the other phototube views the background shower rate "off
source. Every 10 min, the fields-of-view are reversed. In this mode, approxi-
mately 90 percent of the time is spent observing the source.

The major limitation to the sensitivity of these experiments is the isotropic
cosmic-ray background owing primarily to proton-initiated air showers (P-EAS).
Several groups have sought to distinguish 7-ray -initiated showers (G-EAS) from

OBSER VA TIONS OF HIGH-ENERG Y GAMMA RA YS

155

Figure V.A-1. The 10-m optical reflector located at the 2300-m level of

Mount Hopkins, Arizona.

P-EAS by making use of subtle differences in the light distributions from the
two types of showers (Tornabene and Cusimano, 1968; O'Mongain et al.,
1968).

However, the most successful experiments in distinguishing the origin
of the air shower have been performed by Grindlay (1971a; 1971b). He
has presented evidence that he has been able to distinguish the Cerenkov
light from the electrons at the maximum of the shower's electromagnetic
cascade (height hmax) from the Cerenkov light owing to the unscattered,
penetrating shower "core" of predominantly pions, muons, and secondary
electrons. These latter particles would be present only in P-EAS. The

156

OBSER VA TIONAL DA TA

technique uses two searchlight-mirror detectors operated in coincidence
mode and separated by 70 m, with each mirror rotated inward from the
source direction by an angle 6 so that each is pointed at the shower max-
imum (Figures V.A-2 and -3); for a 7-ray energy of 1012 eV, hmax = 6.2
km and 6 = 0?35. A third mirror system is used to detect the penetrating
shower core (h = 3.5 km, d = 0°65). Because the light from the lower
component is relatively nearby, it is rich in the ultraviolet component and
hence can be distinguished from light at the shower maximum. A G-EAS
is registered only when light from shower maximum is not accompanied
by light from the lower level.

EAS MAXIMUM,
~ 330 g/cm2

Figure V.A-2. Simplified geometry of the detection technique
used by Grindlay to reject proton-initiated extensive air showers.

OBSER VA TIONS OF HIGH-ENERG Y GAMMA RA YS

157

Figure V.A-3. Photograph of the series of 1 .5-m searchlight mirrors used by
Grindlay at Mount Hopkins, Arizona.

With this technique, Grindlay has been able to reach an average rejection
ratio of 70 percent. For 7-ray energies >5 X 101 i eV, the combined effects
of P-EAS rejection and increased angular resolution have made possible an
order-of-magnitude increase in sensitivity over mirrors of the same size
used in the normal modes. The drift-scan mode was used in these experi-
ments, and the complicated pointing geometry permitted only 5 percent
of the operating time on source. Recent experiments using the 10-m reflec-
tor and a 1 .5-m searchlight mirror on an alt-azimuth antenna positioner per-
mitted operation in the tracking mode and a considerable increase in operat-
ing time on source.

Grindlay, in cooperation with Prof. R. Hambury Brown's group at the
University of Sydney, has converted the two 6.6-m reflectors of the stellar
interferometer at Narrabri, Australia, for use as atmospheric Cerenkov light
detectors (Figure V.A-4). P-EAS rejection was obtained with a second photo-
tube located in an off -axis position on one of the reflectors, and the reflec-
tors were operated in the tracking mode. Several discrete sources in the
Southern Hemisphere were investigated for the first time in 1972.

158

OBSER VA TIONAL DA TA

Figure V.A4. One of the 22-foot reflectors at Narrabri, Australia.

OBSERVATIONS AND RESULTS

Since 1968, the 10-m reflector has been used to search for cosmic 7-rays
from more than 27 discrete sources, including supernova remnants, pul-
sars, X-ray sources, magnetic variables, radio galaxies, and quasars. With
the exception of the Crab Nebula, none of these sources was detected
(Weekes et al., 1972). For 7-ray energies greater than 2 X 1011 eV, the
upper limits to the flux were of the order of 10"10 photon cm"2 • s . It
takes approximately 10 hours of observation on source to reach these
limits. An extrapolation of the X-ray spectrum of some of these sources
would indicate a 7-ray flux in excess of this value. Simple Compton-
synchrotron models for producing 7-rays in radio sources also predict
fluxes above this value. Where enough information is known about these
radio sources, the upper limits place important constraints on the source
parameters, particularly the average magnetic field in the source.

Although other groups in the past have reported evidence for discrete
7-ray sources in this energy range (for example, Stepanian, Vladimirsky,
and Fomin, 1972), we have investigated these same sources and have
found no evidence of 7-ray emission.

OBSER VA TIONS OF HIGH-ENERGY GAMMA RA YS 159

During 1972, Grindlay's observations with the Narrabri reflectors in the
Southern Hemisphere yielded preliminary evidence for 7-rays from several
sources. No radiation above 2 X 1011 eV from the discrete source Sgr A at
the galactic center nor from several of the 100-MeV 7-ray sources was reported
by Frye et al. (1971). These results are very tentative and further observations
have recently been made from April through June 1973.

The Crab Nebula is a very special case. Observations with the 10-m reflector
over the years 1969 to 1972 indicate an average flux of 7-rays of 4.4 ± 1 .4
X 10"1 1 photon-cm"2 • s"1 with energy above 2.5 X 101 ! eV at the 3.1 a
level (Fazio et al., 1972). This flux corresponds to an emission of 6 X 10
ergs/s, which is significantly less than the X-ray emission of the Nebula.
However, the 7-ray flux may vary with time, and the most significant flux
(1.21 ± 0.24 X 10"10 photon-cm"2 • s"1) may occur 60 to 120 days after a
major spin-up of the Pulsar NP 0532. This increase was observed on three
different occasions, and if only the flux in these intervals is used, the effect
is at the 5 o level. The total 7-ray energy observed on each occasion was
~ 104 1 ergs, an energy approximately equal to the energy of the pulsar
spin-up.

The average 7-ray flux detected can be explained easily by a Compton-syn-
chrotron model of the Crab Nebula, in which the 7-rays are produced by
Compton scattering of relativistic electrons on their own synchrotron radia-
tion (Gould, 1965; Rieke and Weekes, 1969;Grindlay and Hoffman, 1971).
The primary unknown variable in this theory is the magnetic field in the
Nebula; and, hence, a measurement of the 7-ray flux is a way of determining
the average magnetic field. Figure VI.A-5 shows the current data along
with the exact Compton-synchrotron model of Grindlay and Hoffman (1971).
The data are best fitted with a value of (Bj> = 2.5 X 10"4 gauss in the uniform
field model of the Nebula and by a value of B^ = 10"3 gauss for a model
based on a 1/r field, where Bm is the value of the field at the inner edge
of the first wisp surrounding the pulsar.

The variability of the 7-ray flux is more difficult to explain. In a Compton-
scattering process, the electrons have too long a lifetime. The synchrotron
process requires electrons of the order of 101 7 eV in a field of 10"3 gauss;
these electrons have a lifetime of ~ 103 s, and hence this process seems more
feasible. The number of 101 7-eV electrons required is small compared to the
total number of electrons injected into the Nebula. The 60-day delay and
the 60 to 120-day duration may be due either to a time delay in the electron-
acceleration process, for example, in the wisps, or, assuming that the particles
stream out from the pulsar, to the light -travel time in the geometry of the
region where the synchrotron radiation is produced.

160 OBSER VA TIONAL DA TA

Because of this possible variation in the 7-ray flux from the Crab Nebula, it
is important that the SAS-2 experiment monitor the 100-MeV flux from the
Nebula for time variations. If, indeed, the 10-m reflector has detected a con-
tinuous flux of 7-rays at 1011 eV from the Crab Nebula, it becomes particu-
larly interesting to determine to what extent this flux is pulsed. Grindlay
(1971c), using the proton-shower-rejection technique, first reported a
pulsed flux of 7-rays from NP 0532 based on 42 drift scans in January 1971.

Ninety-nine additional scans were obtained during November and December
1971, which also showed evidence of a pulsed effect (Grindlay, 1972).
Later, it was discovered that the November and December data were analyzed
with the wrong period, owing to a computer-program error. Reanalysis
of this data resulted in even more significant evidence of having detected
a pulsed effect from NP 0532. The sum of all 141 drift scans exhibited
a 4.2 a effect, but the primary and interpulse both appear to be 1.7 ms
early with respect to the corresponding optical peaks. These data corres-
pond to a pulsed flux above 6 X 1011 eV of 8 ± 6.5 X 10"12 photon
cm"2 -s .

Grindlay repeated the observations in 1973 by using the tracking mode to
increase the observing time on the source. Preliminary analysis of the data
again showed evidence of pulsed 7-rays, but the primary peak of the radiation
may be delayed in phase from the optical pulse by 1.7 m (Grindlay et al.,
1973). This repeated positive effect is most interesting, and it is still possible
that Grindlay, using the proton-shower-rejection technique, has detected a
pulsed 7-ray flux from NP 0532, but the present results do not give unam-
biguous proof.

Helmken et al. (1973), by using data on the Crab Nebula obtained with the
10-m reflector, have analyzed the arrival times of air showers for over 200
hours of these observations; the arrival times were recorded to a precision
of 200 jits. An analysis of the data by use of the optical pulsar period and
phase revealed no statistical excess at the primary pulse of the interpulse.
A typical upper limit to the flux at 1.8 X 1011 eV for a 1.3-ms bin width
and E"11 spectrum was 1 .4 X 10"11 eV-cm^-s^-eV'1. Upper limits to the
flux were also obtained at energies of 3 .0 X 1 01 1 e V and 4.7 X 1 0 1 1 e V
(Figure V.A-5).

The lower energy X- and 7-ray data are best fitted by a curve of the form
1.0E"1'1. Extrapolated to the 101 1 eV energy region, the curve lies 2
orders of magnitude above the current upper limits. The extrapolated flux,
if true, would be verified in less than 20 minutes of observations. Thus, an im-
portant result of this work is that the upper limits to the 7-ray flux indicate
a major break in the pulsar spectrum between 1 and 100 GeV. When the

OBSER VA TIONS OF HIGH-ENERGY GAMMA RA YS

161

McBREENet_g|..(l973)
(PULSED)

FAZIO et_o[. (1972)

CONTINUOUS
HELMKEN et gl. (1973)
(PULSED)

GRINDLAY et_gl (1973)
(PULSED)

GRINDLAY (1972)
(PULSED)

UCD- HARWELL (1972)
.(PULSED)

C-S MODEL

B~l/r
JX^/Bj^ = I x id"3 GAUSS

UNIFORM FIELD
<BJ> = 2.5xl0"4 GAUSS

9 10 II

LOG PHOTON ENERGY (eV)

Figure V.A-5. Graph of recent results of the pulsed and continuous flux from
the Crab Nebula. The solid lines represent the Compton-synchrotron model
computed by Grindlay and Hoffman (1971) for a uniform magnetic-field
model and for a field proportional to 1/r.

previously reported positive continuous flux is taken with the present upper
limits to the pulsed flux, it places an additional upper limit of 30 percent
on the ratio of the pulsed-to-continuous component at 101 1 eV. This is a
reversal of the trend at lower energies.

The University College, Dublin-Harwell group (N. Porter, private communica-
tion, 1972) also have evidence for a periodic flux of 2 X 101 2 • eV 7-rays
from the pulsar, but, again, the primary pulse is not in phase with the optical
pulse. If real, the effect would correspond to a flux of 2 X 10"1 2 photon-cm"2

FUTURE EXPERIMENTS

It is particularly important to continue observations on the Crab Nebula for
two reasons: (1) To determine if an increase in the continuous 7-ray flux
is associated with the pulsar spin-up; and (2) to determine if Grindlay 's
technique of rejection of proton-induced showers has detected a pulsed flux
from NP 0532.

1 62 OBSER VA TIONAL DA TA

The next priority would go to observation of the sources seen in the
100-MeV to 1-GeV region with balloon-borne detectors, for example, the
sources reported by Frye et al. (1971), and with the SAS-2 and TD-1
satellite detectors.

The sensitivity achieved in the current experiments has been the result
of many hours of observation on a limited number of sources. It is still
possible that there exist sources of detectable intensity that were not
included in our survey or that were not observed in other regions of the
spectrum. Therefore, Weekes et al. (1972), have proposed an all-sky
survey of the Northern Hemisphere. Very few observations have been
made in the Southern Hemisphere.

It is also possible that the 7-ray sources examined are time variable, which
makes verification even more difficult. Delays between balloon-borne
7-ray detector flights are of the order of 6 months. One advantage of
the atmospheric-Cerenkov light technique is that immediate observations
can be made on a suspected source. In view of this, we ask groups that
have discovered a possible source of cosmic 7-rays to communicate the
information to us as soon as possible.

In all the above programs, an increase in detector sensitivity would be
most helpful. In theory, the proton-shower-rejection technique used by
Grindlay should significantly increase the sensitivity. Hence, it appears
that the design of any future detectors should use this technique. One
possibility is the construction of a second large reflector near the 10-m
reflector at Mount Hopkins; another would be to lower the 7-ray
threshold energy (E ) of the 10-m reflector. The current reflector mount
could support a second 10-m reflector. Since E <* A"1''2, where A is the
area of the reflector, doubling the area would reduce the threshold energy
only by a factor of 0.7, but additional reductions could be made by in-
creasing the frequency bandwidth and operating in the coincidence mode.

Continued studies of the structure of the Cernekov light bursts produced
in air showers are also necessary to maximize the efficiency of present
detectors. For example, N. Porter has suggested that the geomagnetic
field can have important effects on the angular distribution of Cerenkov
light from extensive air showers.

ACKNOWLEDGMENTS

The results described in this paper, obtained with the 10-m reflector,
have been produced through the efforts of many people. Trevor C.
Weekes has been primarily responsible for the operation of the 10-m
reflector with the assistance of Ed Horine. The pulsar data analysis was

OBSER VA TIONS OF HIGH-ENERGY GAMMA RAYS 163

done through the very patient efforts of Henry Helmken. George Rieke
and Eon O'Mongain assisted in many of the observations and in the inter-
pretation of the data. Jonathan Grindlay performed the Compton-
synchrotron model calculations.

REFERENCES

Bosia, G., M. Maringelli, and G. Navorra, 1972, II Nuovo Cimento, 9B,
p. 201.

Fazio, G. G., H. F. Helmken, E. O'Mongain, and T. C. Weekes, 1972,
Astrophys. J. Letters, 175, p. LI 17.

Frye, G. M., P. A. Albats, A. D. Zych, J. A. Staib, V. D. Hopper, W. R.
Rawlinson, and J. A. Thomas, 1971, 12th Int. Conf. on Cosmic Rays,
Hobart, Tasmania, paper OG-24.

Gould, R. J., 1965, Phys. Rev. Letters, 15, p. 577.

Grindlay, J. E., 1971a, II Nuovo Cimento, 2B, p. 119.

, 1971b, Smithsonian Astrophys. Obs. Spec. Rept. , No. 334.

, 197 lc, Nature Phys. Set ,234, p. 153.

, 1972, Astrophys. J. Letters, 174, p. L9.

Grindlay, J. E., H. F. Helmken, T. C. Weekes, G. G. Fazio, and F. Boley
1973, Proc. 13th Int. Conf. on Cosmic Rays, Denver, in press.

Grindlay, J. E., and J. A. Hoffman, 1971, Astrophys. J. Letters, 8, p. L209.

Helmken, H. F., G. G. Fazio, E. O'Mongain, and T. C. Weekes, 1973,
Astrophys. J., in press.

Long, C. D., 1967, Ph. D. Thesis, National Univ. of Ireland.

O'Mongain, E. P., N. A. Porter, J. White, D. J. Fegan, D. M. Jennings,
and B. G. Lawless, 1968, Nature, 219, p. 1348.

Rieke, G. H., 1969, Smithsonian Astrophys. Obs. Spec. Rept. No. 301.

Rieke, G. H., and T. C. Weekes, 1969, Astrophys. J., 155, p. 429.

Stepanian, A. A., B. M. Vladimirsky, and V. P. Fomin, 1972, Nature
Phys. Set, 239, p. 40.

Tornabene, H. S., and F. J. Cusimano, 1968, Canadian J. Phys., 46, p. S81.

1 64 OBSER VA TIONAL DA TA

Weekes, T. C, G. G. Fazio, H. F. Helmken, E. O'Mongain, and G. H.
Rieke, 1972, Astrophys. /., 174, p. 165.

Zatsepin, G. T., and Chudakov, 1962, Soviet Phys.-JETP, 15, p. 1 126.

Zatsepin, G. T., 1965, Soviet Phys.-JETP, 20, 459.

Chapter VI

A. OBSERVATIONS OF GAMMA-RAY EMISSION
IN SOLAR FLARES

D. J. Forrest*, E. L. Chuppf, A. N. Suri, and C. ReppinJ

University of New Hampshire

INTRODUCTION

The primary purpose of this paper is to review the observations of solar flare-
associated 7-rays. Some preliminary discussion of the features of the measure-
ments will be given even though the full interpretation of the measurements
is not complete, as far as understanding the physics of solar flares is concerned.
The observations discussed here were first presented at the NASA Symposium
on High Energy Phenomena on the Sun (Chupp et al., 1972), and a more
detailed report has been published (Chupp et al., 1973).

The University of New Hampshire 7-ray detector, which is situated in the
wheel section of the OSO-7 spacecraft, has been described in detail by
Higbie et al., 1972. Briefly, it consists of a 7.6-cm by 7.6-cm Nal scintillator
surrounded by and in anticoincidence with an active Csl shield. It is cali-
brated by a gated radioactive source (Forrest et al., 1972) twice each orbit
and has an energy resolution of ^8 percent FWHM at 662 keV. Two inde-
pendent pulse-height spectra covering the energy range 0.3 to 9 MeV are
simultaneously accumulated over an 180-s interval: one in the solar direction
and one in the antisolar or background direction. An auxiliary 7.9-cm2 Nal
X-ray detector is also included in the instrument. It covers the energy range
7.5 to 120 keV in four energy bands, and a complete X-ray spectra is taken
every 30 seconds.

* Speaker.

fAlexander von Humboldt and Fulbright-Hayes Scholar on leave at the Max Planck
Institute for Extraterrestrial Physics, Munich, Germany.

$On leave from the Max Planck Institute for Extraterrestrial Physics, Munich,
Germany.

165

166 OBSER VA TIONA L DA TA

GAMMA-RAY OBSERVATIONS

Figure VI .A- 1 shows the counting rate versus time in several energy bands
covering the range 7.5 keV to 8 MeV observed during the 3B Ha flare that
began at ~ 0621 UT on August 4, 1972. Also shown is the radio burst at
19,000 MHz as reported in UAG-21 (Lincoln and Leighton, 1972). The
rates in the energy interval 7.5 to 120 keV are from the X-ray detector,
and the rate in the 0.35-to 8-MeV interval is from the central 7-ray detector.
As can be seen, OSO-7 was eclipsed by the earth before the event was over,
but according to the radio burst, most of the flash phase was observed. The
time correspondence between the radio, X-ray, and 7-ray continuum is self-
evident. Figure VI.A-2 shows some of the same rates on an expanded time
scale. All of the rates were observed to increase above their pre flare values
at 0621 ± 1 UT. Although the lower energy channels quickly reached their
instrumental saturation level, the two higher energy channels did not. These
channels indicated that the rates continually increased over a 200-s interval
and then appeared to level off until the eclipse at 0632 UT. The pulse-height
spectrum that was observed in the time interval 0623 to 0632 UT is shown
in Figure VI.A-3. As can be seen, there is an increase in the energy continuum
that extends above 3 MeV and two pronounced photopeaks at 0.5 and
2.2 MeV in the solar quadrant. The two peaks at 1.17 and 1.33 MeV are
leakage peaks from the onboard Co60 calibration source. The two peaks at
0.5 and 2.2 MeV have been interpreted as resulting from positron annihilation
at 511 keV and neutron capture in hydrogen at 2.23 MeV. The time history
of the counting rates in the two photopeaks are shown in Figure VI.A-4
where the 60- to 1 20-keV rate (X-ray Channel 4) is reproduced for com-
parison. Although the photopeak counting statistics in the individual 180-s
scans are not sufficient to determine a detailed time history, it can be seen
that the rates in the photopeaks follow the time history of the high-energy
continuum quite closely.

The counting rate observed in association with the 3B Ha flare that started
at ~ 1517 UT on August 7, 1972, is shown in Figure VI.A-5. Also shown
is the radio flux at 15,400 MHz (Solar-Geophysical Data, Report No. 342,
February 1973). The OSO-7 spacecraft was in eclipse during the flash phase
of the flare and no continuum X-rays with energies greater than 120 keV
were observed after the spacecraft came out from eclipse at 1538 UT.
However, evidence for line emission at 0.5 and 2.2 MeV was observed
between 1538 and 1547 UT.

The time-averaged fluxes at these two energies observed on August 4 and 7
are given in Table VI.A-1. Also given are the fluxes observed at 4.4 and
6.1 MeV on August 4. These latter two lines have been interpreted as
arising from C12* and O16*.

GAMMA-RA Y EMISSION IN SOLAR FLARES

167

EAK OF BURST
OFF-SCALE

19.000 MHz
SLOUGH, U.K.

AUGUST 4, 1972
OSO-7

-NIGHTTIME

SATUI

"^^v^Mr^,

tmm^f RWf** WH^

0800 0900

UNIVERSAL TIME

Figure VI.A-1. Counting rates and radio flux versus time for the flare on

August 4, 1972.

INTERPRETATION

The 7-ray lines observed from the flare on August 4, namely, at 0.51, 2.23,
4.43, and 6.13 MeV (from positron annihilation, neutron capture on hydro-
gen, and excited states of C12 and 016, respectively) have been predicted to
be the most intense lines based on known cross sections, solar abundances,
and assuming nuclear interaction of the energetic solar particles with the
solar atmosphere (Lingenfelter and Ramaty, 1967). The ratio of the observed
lines are those predicted to result from a spectrum of energetic solar particles

168

OBSER VA TIONAL DA TA

I09

Z
3
O

t 1 r

SATURATION

AUGUST 4, 1972
0S0-7

0615

0620 0625

UNIVERSAL TIME

0630

Figure VI.A-2. Counting rates on an expanded time scale for the flare on

August 4, 1972.

GAMMA-RA Y EMISSION IN SOLAR FLARES

169

240,

200-

160

0) 100-
0.

20-

0S0-7

Solor Flore Gamma Ray Spectrum
(0.4-3.1 Mev)
August 4,1972

Solor Quod.

1.17 Mev] _ 60
l.33MevJ (-°

-2.2 Mev

nji^^MnlV^

J 25 ' 40 ' 60 '~&0 ' 100 ' 120 ' I40J i60 180 200

Channel Number

Figure VI.A-3. Pulse-height spectra recorded in the time interval 0623
0632 UT on August 4, 1972.

170

OBSER VA TIONAL DA TA

10'

103

^ 102

Aug 4, 1972 OSO - 7

— X-ray Channel # (60-120 kev)
• 0.5 Mev
a 2.2 Mev

4 h

6:00' :05 :10 :15 :20

Universal Time

1 I I

2.0

1.6

1.2

0.8

0.4

:25

:30

:35

Figure VI.A-4. Time history of the photopeak counting rate on August 4, 1972.

of the form

N(>P) = N0exp(-P/P0)

where the characteristic rigidity PQ is in the range 60 to 80 MeV. The spectrum
of energetic particles measured on satellite detectors near 1 AU between
August 4 and 8 agree with this spectral shape (Ramaty and Lingenfelter, 1972).
However, there is at least one reference to a ground level effect from this flare
(Pomerantz and Duggal, 1973). If this is true, then at least a portion of the
energetic solar particles must have had a much higher characteristic rigidity.
The absolute intensity of the observed 7-ray line fluxes, however, is much lower

GAMMA-RA Y EMISSION IN SOLAR FLARES

171

S2QQ00f

I5P00

2-10000

O 5000

AUGUST 7, 1972
0S0-7

W5"

1600
UNIVERSAL

Figure VI.A-5. Counting rates and radio flux versus time for the flare on
August 7, 1972.

(by a factor of 102 to 103) than was predicted from a flare of this magnitude.
The intensity of the 7-rays is based mainly on the solar atmospheric density
in the region where the particles interact and the number of energetic parti-
cles accelerated and released. In the past the only estimate of the total
number of particles accelerated and released was based on measurements near
1 AU and model-dependent extrapolations back to the sun.

If the observed 200-s rise time of the very hard X-ray continuum is inter-
preted to be the time history of the rate of nuclear reactions producing the
positrons and neutrons, then several other interesting results can be derived.
First, unless the electrons and protons are accelerated and stored very high
in the atmosphere and what we are seeing is the dumping of these particles
into the denser lower atmosphere, then the time scale for converting some
form of potential energy into the kinetic energy of relativistic particles is
also 200 s. Second, a study of the rise time of the 2.2-MeV line flux indi-
cates that the neutrons must have been captured in a region where the density

172

OBSER VA TIONAL DA TA

3
00

O
O

>-< —

.2 6

o u

C/2 go

c3 o

C &*

.2 D

03 •<

—
.2P rt

</>

<L>

.2
'o

U
5S

c/i

-o

O
X

d
+i

*«*

O MD CO
O w

cs)

rr*

f""<

l~->

bO
3
<

<l>

GAMMA-RA Y EMISSION IN SOLAR FLARES 1 73

was *2X 1017 protons/cm2 (Reppin et al., 1973). That this region would
be in the photosphere is expected since the neutrons, being uncharged, can
easily escape from the region where they are produced to the higher density
regions where they can be slowed down and captured. Also, the observed
risetime of the 511-keV line cannot be more than 200 s. This fact, together
with the known cross section for annihilation, implies that the density in
the region where the positrons annihilate must be greater than 2 X 1 01 1
electrons/ cm3. A study of the line width of the 511-keV line observed on
August 4 has lead to an upper limit for the temperature in this same region
of « 7 X 106 K (Dunphy et al., 1973). Because the positron is charged, it
is reasonable to assume that the positrons are trapped in the region where
they are produced, and that the above temperature is indeed an upper limit.

The observations of line emission on August 7, after the flash phase was over,
were expected because of the % 200-s annihilation time for the positrons
(some of the positrons are produced from radioactive isotopes with long
half-lives) and the 100-s capture time for neutrons in the photosphere.

CONCLUSION

The solar flare 7-ray observations reported here appear to be in general
agreement with models and calculations proposed by Lingenfelter and
Ramaty, 1967. Further study of these observations together with other
observations of the same flare at other wavelengths and observations of
particles should lead to a rather specific acceleration and interaction model
for these flares. Of particular interest are the reported He3 measurements
(McDonald et al., 1973). He3 in the intensities reported must have been
produced in the same sort of nuclear reactions that produced the 7-rays.

(Supported by NASA Contract NAS 5-11054)

REFERENCES

Chupp, E. L., D. J. Forrest, and A. N. Suri, Proceedings of Symposium on
High-Energy Phenomena on the Sun, Greenbelt, Maryland, September 28-
30, 1972, R. Ramaty and R. G. Stone, eds., in press.

Chupp, E. L., D. J. Forrest, P. R. Higbie, A. N. Suri, C. Tsai, and P. P.
Dunphy, 1973, Nature, 241, p. 33.

Dunphy, P., E. L. Chupp, D. J. Forrest, and A. N. Suri, 1973, Bull
American Phys. Soc, 18, p. 695.

Forrest, D. J., P. R. Higbie, L. E. Orwig, and E. L. Chupp, \912,Nucl
Inst, and Meth., 101, p. 567.

Higbie, P. R., E. L. Chupp, D. J. Forrest, and I. U. Gleske, 1972, IEEE
Trans. Nucl Sci. NS-19, No. 1, p. 606.

1 74 OBSER VA TIONAL DA TA

Lincoln, J. Virginia, and Hope I. Leighton, 1972, World Data Center A,
Report UAG-21 ; U. S. Department of Commerce, NOAA, Environ-
mental Data Service, Ashville, North Carolina.

Lingenfelter, R. E., and R. Ramaty, 1967, High Energy Nuclear Reactions
in Astrophysics, B. S. P. Stern, ed., Benjamin, New York.

McDonald, F. B., B. J. Teegarden, J. H. Trainor, W. R. Webber, and
E. C. Roelof, 1973, Bull American Phys. Soc, 18, p. 697.

Pomerantz, P., and S. P. Duggal, 1973, Nature, 241, p. 33.

Ramaty, R., and R. E. Lingenfelter, Proceedings of Symposium on High
Energy Phenomena on the Sun, Greenbelt, Maryland, September 28-30,
1972, R. Ramaty, ed., in press.

Reppin, C, E. L. Chupp, D. J. Forrest, and A. N. Suri, 1973, Proceedings
13th Int. Conf. on Cosmic Rays, Denver, in press.

Chapter VII

A. ENERGY SPECTRA OF COSMIC
GAMMA-RAY BURSTS*t

T. L. Cline and U. D. Desai

Goddard Space Flight Center

R. W. Klebesadel and I. B. Strong

Los Alamos Scientific Laboratory

INTRODUCTION

The occurrence of intense, several-second duration bursts of 0.1- to 1.2-MeV
cosmic 7-rays, recently found using multiple Vela satellites (Klebesadel et al.,
1973), has been confirmed with measurements from the IMP-6 satellite.
Observations regarding times of occurrence, photon flux, and temporal and
spectral characteristics of the bursts are entirely consistent. In particular,
since the IMP-6 instrument incorporates a hard X-ray detector with active
particle rejection and full-time omnidirectional particle intensity monitoring,
the results fully confirm and establish the hard X-ray or 7-ray nature of the
incident flux.

Detailed differential energy spectra were obtained with the IMP-6 for six of
the eight known events occurring during the March 1971 to September 1972
lifetime of the instrument. All of these are multiple-pulse events, with several
seconds separation between distinct pulses of one or two seconds duration.
The pulse spectra do not obey single-index power laws in energy, but can be
simply represented by exponentials in photon flux throughout the 100- to
1200-keV region. The characteristic energies at maximum intensity appear
to cluster near 150 keV, with indications that departures from this value can
be interpreted as circumstantial, due to attenuation when the source is at
great angles from the detector axis. These burst pulses appear to ride on a
softer component that exhibits a longer decay-time constant, and has a
characteristic exponent near 75 keV. There is no evidence for monoenergetic
line structure in the several hundred-keV region, or for marked changes in
the spectrum with time during a single pulse. Size spectra can be estimated

*Post-Symposium observational paper, see Introduction.
fPublished in Astrophysical Journal Letters, October 1, 1973.

1 75

1 76 OBSER VA TIONAL DA TA

to predict the frequences of occurrence of smaller events for both a galactic
model (for example, a new class of 7-ray flare star) and an extra-galactic
model (for example, supernovae). In either case, the total emission is below
the value currently obtained for the diffuse celestial X-ray background and is
unlikely to account for any of its spectral features.

INSTRUMENTATION

The IMP-6 satellite was launched on March 14, 1971 , into an elliptic orbit
with an initial apogee of over 200,000 km. Gamma-ray monitoring was
provided on a nearly continuous basis, except for passes every 4.14 days
through the magnetosphere, lasting several hours each. The detector was
in operation from launch until May 2, 1971, and again for the period from
June 9, 1971, to September 27, 1972. The instrument used consisted of a
6-cm (2.25-in.) diameter by 3-cm (15 -in.) thick CsI(Tfi) crystal, entirely
surrounded by a thin plastic scintillator for particle rejection, viewed by a
single PM tube. In addition to full-time monitoring of the rates of total
intensity, particle intensity and 7-ray intensity, and energy spectra of incident
7-rays were measured by a 14-channel analyzer with simultaneous storage in
all channels. The spectra were accumulated for one half of the time, for
each * 6.3-s period from sun rise to sunset on the detector, determined by
the optical aspect. This 50-percent duty cycle resulted in missing several of
the very brief 7-ray bursts. The spectral accumulation times were fixed at
% 5.1 s so that the * 6.3-s lifetimes were asynchronously split into 2 or
3 intervals of shorter durations, making possible more than one spectral
determination during some of the pulses. The gain of the system was cycled
through four positions with changes at « 1-week intervals for purposes of
in-flight calibration, so that some of the bursts happened to be observed
with a 69- to 1 150-keV dynamic range and some with a 53- to 880-keV
range. The primary purpose of this 7-ray detector was used as a coincident
annihilation spectrometer incorporated in a positron detector. The secondary
objective was that of a solar-flare monitor, and it was in this mode of opera-
tion that these unexpected 7-ray bursts were observed.

DATA OBSERVATION AND ANALYSIS

The times of occurrence of 7-ray bursts observed with multiple Vela satellite
coincidences were used to identify coincident increases in the IMP-6 7-ray
intensity. Six of the eight Vela events were observed well above the omni-
directional background, the others being missed because of the 50-percent
detector duty cycle. It is possible that other events, of intensity too low to
exceed the Vela threshold triggers, may also be observable with the IMP-6
instrument. Figure VII. A-l shows the response of the IMP 7-ray detector to
the event of June 30, 1971. During a several- second interval, the counts in

COSMIC GAMMA-RA Y BURSTS

177

IMP-6 EVENT OF 30 JUNE 1971

1000

500

0
30,000 -

20,000

10,000

«=

5,000

4,000

3,000

2,000
1,000

yP ANALYZER
(140-475 keV)

62800 62900 63000 63100 63200
SECONDS (U.T.)

63300

Figure VII.A-1. The response of the detector to a 7-ray
burst, as indicated by the plastic anticoincidence (P), the Csl
7-ray detector (7P), and several channels added to give the
140- to 475-keV photon rate, where the 7-ray energy response
is maximized. Each point samples two differential energy
spectra.

the plastic scintillator (P) surrounding the 7-ray crystal increased by about
50, while the neutral counts in the crystal (7P) simultaneously increased by
about 18,000. Pulses satisfying the 7-ray logic were fed to a multichannel
analyzer, from which the outputs of three channels, added to provide the
flux of 140- to 475-keV photons, indicated an increase during one « 5-s
interval of nearly 5000 counts from a total omnidirectional and secondary
background of about 400 counts. This illustrates the remarkable intensity
of the bursts and shows that the response is entirely consistent with that
of hard X-rays or 7-rays.

The times of occurrence and various properties of all Vela/IMP events
observed during the IMP-6 experiment lifetime are listed in Table VII.A-1.
The temporal structures of the observed bursts, known from the Vela results,
were compared in order to determine over which intervals in the event struc-
tures the IMP spectra were obtained. Since the IMP-6 satellite was spinning,

178

OBSER VA TIONAL DA TA

Table VII. A- 1.

Characteristics of 7-ray Burst Spectra
(Exponential fits in dn/dE provide IQ in units of
photons cm"2 • keV1 , and EQ in units of keV, both of which
have systematic uncertainties depending on relative look angle.)

Event

Burst

Look Angle

Mar. 15, 1971

2nd Max

1.9

156

Includes source

(a ^50°, 5 =-30 ±10°)

Mar. 18, 1971

Decay of 1st

1.8

Spins through source

Jun. 30, 1971

1st Max

0.7

276

Source below satellite
horizon

Jun. 30, 1971

2nd Max

5.5

142

Spins through source

Jun. 30, 1971

Decay of 2nd

0.7

110

Spins through source

Jan. 17, 1972

Decay of 1 st

0.10

138

Source position undeter-
mined

Jan. 17, 1972

2nd Max

0.35

184

Source position undeter-
mined

Jan. 17, 1972

Decay

0.11

170

Source position undeter-
mined

Mar. 28, 1971

1st Max

0.50

128

Source near or below
horizon

Mar. 28, 1972

2nd Max

0.55

176

Source position undeter-
mined

May 14, 1972

1st Max

0.8

166

Includes source

(a ^175°, 5% +77°)

May 14, 1972

2nd Max

0.8

152

Includes source

an analysis was also made for each burst to determine in which direction the
detector was facing, relative to the source, at the moment each spectrum was
obtained. Each of the six events was observed by the Velas to have at least
two distinct pulses of up to a few seconds duration, separated by intervals of
several seconds. The time resolution of the IMP 7-ray detector (^ 2.5 s)

COSMIC GAMMA-RA Y BURSTS 1 79

permitted obtaining individual spectra for many maximum-intensity pulses
and, for some cases, two separate spectra of a given several-second pulse.
(Vela data show that a given maximum-intensity pulse can contain a variety
of fast-time variations (Klebesadel et al., 1973); these are necessarily averaged
over in the IMP spectra.) Figure VII.A-2a shows photon-number spectra,
dn/dE, for several of the bursts, as sampled on a 6.3-s half-spin basis. It
indicates that, in this energy region, relatively good fits to these raw data
are obtained to exponentials of the form dn/dE = I exp (-E/E ) photons
cm"2 keV"1 burst"1. The IQ and EQ values are listed in the table, along with
the relative look direction accuracies. The directions of origin of the six
events are known with widely varying accuracy; but, in the case of the June 30
event, it is known that the first spectrum was recorded when the source was
below the satellite horizon of the detector. Thus, the harder (250-keV)
spectrum may be accounted for by attenuation of the lower energy photons
in the metal surface of the satellite. If that is also the case of the March 28,
1972, event, then all the pulses are consistent with 150-keV spectra. Two
of the six events (March 15, 1971, and May 14, 1972) have unambiguously
known source directions that are not far from the center of the field of view,
and these are definitely 150-keV spectra. In addition, the March 18, 1971,
event and a number of decays of the other events, not listed, are consistent
with softer spectra, suggesting that a slower-time constant soft component
is present in addition to the 1 50-keV peaks. Figure VII. A- 2b (insert) shows
the energy spectrum, or power spectrum averaged over the pulse burst
duration, Edn/dE, of an event for which the source direction was known
to be in the view direction. It is seen that the energy output is a maximum
in this region. This may indicate that the photons released from the source
objects are essentially 7-ray in nature, not composing X-ray distributions
with spectral tails in the 7-ray region. If much softer X-rays are emitted in
the primary burst, they most likely undergo relatively greater absorption
near their region of origin.

DISCUSSION

It is clear that the observed 7-ray bursts represent an entirely novel form of
cosmic energy release. The durations of the individual pulses are typically
1 to a few seconds, and the separations between pulses in a given burst are
up to 20-odd seconds. The temporal structure might therefore be compared
to that of solar flares, but with time scales one to two orders of magnitude
shorter, suggesting a conceivable source origin of 7-ray flare stars (see also

180

OBSER VA TIONAL DA TA

I E

2 •'

h-
o

2: .01

.01

100 r 14 MAY 1972*1

I0r

^>-

1 00 1 000

Ey(keV)

+ 30 JUNE 1971 #1
O30 JUNE 1971*2
• 18 MAR 1971*2

+ 15 MAR 1971 *l
028 MAR 1972*1
• 17 JAN 1972 *2

I I I I L

I I I I

300 600 900 1200
Ey (keV)

Figure VII.A-2. (a) Number spectra dn/dE, of several
bursts, selected for the greatest variety of responses. The
harder spectra are interpreted as due to attenuation of the
incident beam by the satellite material in cases where the
source was below the detector horizon, (b) The energy
spectrum, E dn/dE, of a directly observed event is shown
in the insert.

Stecker and Frost, Chapter XIII. A). The 150-keV energy spectra, including
the one known case of the May 14, 1972, event which has a power law
from 10 to 100 keV (Wheaton et ah, preprint), contain too much emission
in the X-ray region to fit « 150 keV blackbody spectra. However, the
spectra contain too little emission in the lower energies to be compared to
the typical, steep X-ray spectra, having index of «-3 or more, of most hard
solar flares and many celestial X-ray sources. For those pulses which were
observed with sufficient temporal resolution to obtain more than one spec-
trum per pulse, there is no evidence for changes in the characteristic energy
during its extent (not illustrated). Further, there is no evidence for line
structure in this energy region. It is possible, however, that great improvements

COSMIC GAMMA -RA Y BURSTS 181

in energy and time resolution might show fine-scale spectral variability with
a variety of monochromatic lines, which average out over 2-s summations.

An integral size spectrum can be constructed, assuming a power law with
index -1.5, normalized to 6 or 8 events per 1 .5 years with sizes greater than
10"4 erg • cm"2 for the energy region above 100 keV. Since the 18 known
events have source directions compatible with isotropy (Strong and
Klebesadel, preprint) rather than with, for example, galactic plane clustering,
the source objects must either have distances in the tens to hundreds of pc if
galactic, or have distances of greater than several Mpc if extragalactic in
nature. Thus, this size spectrum can be normalized for these two models in
order to obtain predictions of the frequencies of occurrence of smaller events.
In the case of extragalactic sources, for example, 7-ray rich and optically poor
supernovae or other large collapsing objects, a summation of all emissions up
to cosmological distances produces a total isotropic background intensity
which is below the presently observed diffuse cosmic background in this
energy interval. Thus, an extragalactic origin cannot be ruled out. Further,
if all sources have spectra with « 150-keV exponentials, then the total cosmic
spectrum will not extend into the several-MeV region with sufficient intensity
to explain the bump in the diffuse cosmic background observed (Trombka
et al., 1973) at those energies.

REFERENCES

Klebesadel, R. W., I. B. Strong, and R. A. Olson, 1973, Astrophys. J. Letters,
182, p. L85.

Trombka, J. I., A. E. Metzger, J. R. Arnold, J. L. Matteson, R. C. Reedy,
and L. E. Peterson, 1973, Astrophys. J., 181, p. 737.

SECTION 2
THEORY

Chapter VIII

A. THE ASTROPHYSICS OF THE DIFFUSE
BACKGROUND OF X-RAYS
AND GAMMA RAYS

Ramanath Cowsik*
University of California

INTRODUCTION

Studies in the field of X-ray and 7-ray astronomy have given rise to new insights
into the structure and composition of our galaxy, the intergalactic space, and the
universe itself. Since there have been many comprehensive reviews (Silk, 1970,
and preprint; Felten, 1972) on the subject, we will describe here only our views
on the origin of the diffuse X-ray and 7-ray background and some of the astro-
physical implications of such a radiation background. In the same spirit, detailed
references to all existing literature is not made, and one can refer to the com-
prehensive reviews for this purpose.

The range of energies that is of interest here extends from ~102 eV to ~108 eV,
over six decades, and a variety of processes contribute to the generation of a
diffuse background. In order to make statements about the distribution of the
sources of the radiation background, we appeal primarily to the angular dis-
tributions of the radiations about us. Considerations based on plausibility of
models of origin and on minimizing the energy requirements in the sources sup-
plement the classification of the sources either as galactic or extragalactic. Our
views on the origin of the various components are summarized in Table VIII. A-l .

Finally, in the section on the 100-MeV 7-ray flux, it is shown that the measured
7-ray fluxes at ~ 100 MeV from the galactic disk place a rather stringent upper
limit on the energy density of any background at submillimeter wavelengths.

"Speaker.

185

186

THEORY

Table VIII.A-1
Origin of the Diffuse X-Ray and 7-Ray Background

Energy Range
-250 eV - 2 keV

Process

Thermal brems-
strahlung

Source Region

Discussed In
Section:

Diffuse X-Ray

Our galaxy and the
external galaxies

Compton scat-
tering of the
2.7 K photons

Intergalactic space

~2 keV - 200 keV

Thermal brems-
strahlung

Intergalactic space

The 2 - 200 keV

~0.2MeV-10MeV

Compton scat-
tering of ~104
K photons

Cosmic ray sources
in the galaxy

0.3-3 MeV 7-rays

> 100 MeV

Compton scat-
tering of star-
light

Central regions of
galaxy (extended
source at the
galactic center)

Gamma-rays in
100-MeV Range

7T°->27

Galactic disk (line
source)

7T°+27 ?

Galactic halo?
(isotropic background)

SOME IMPORTANT MECHANISMS FOR GENERATION OF X-RAYS AND
7-RAYS

Thermal Bremsstrahlung

A high temperature plasma emits X-rays mainly through free-free transitions.
Here the electrons that have a thermal energy distribution emit bremsstrahlung
photons in the field of the ions. This process is weakly dependent on the exact

DIFFUSE BA CKGROUND OF X-RA YS 187

chemical composition of the plasma and the rate of X-ray emission by an optically
thin plasma is given by (Hayakawa, 1969)

1 e /mc

p,_ (E ) = — r — a., en

HffV x) 6n3 he th e IkT

1/2

[Eznzgff(Z,T,Ex)]

x- exp (-Ex/kT) (VIII.A-1)

= 0.81 X 10-12 ne2 T"% geff X ± exp (.E /kT)

For a plasma of solar composition the effective Gaunt factor (geff) is approxi-
mately equal to unity. Equation (VIII.A-1) integrates easily to yield a cooling
time

r« 1.96 X 1011 T1/2 /ne s (VIII.A-2)

In galaxies, clouds of hot plasma can be created continuously, for example by
supernova explosions. These will cool continuously emitting radiation. Then,
at any time there will be an equilibrium distribution temperature of these clouds
extending up to T , the maximum temperature of generation of these clouds.
If all clouds are created at T and they cool mainly through the free-free
process, then the integrated emission of all the clouds can be approximated by

p(Ey~ exp (-Ex/kTmax) (VIII.A-3)

1
Notice that this is steeper than the single temperature case by a factor — , in-
dicating that there is less emission at high energies. x

Besides free-free emission there would be free-bound and bound-bound transi-
tions which will lead to sharp edges and lines in the emitted spectrum depending
on the elemental abundances in the plasma.

Decay of Neutral Pions

Neutral pions are produced in the interaction of nuclear cosmic rays with ambient
matter, and these pions decay almost instantaneously to two 7-rays. This sub-
ject has been studied extensively by Stecker (1971) and in Figure VIII.A-1 we
show the spectrum of 7-rays generated through this process. It is very flat in
the region of ~70 MeV and has a spectral slope identical to the cosmic-ray beam
at high energies.

188

THEORY

Compton Scattering of Thermal Photons

The importance of this process under astrophysical conditions has been made
clear by the work of Morrison and his coworkers (see for example Brecher and
Morrison, 1969). In this process a highly relativistic electron scatters a low-
energy thermal photon into the X-ray energy region. Cowsik and Kobetich
(1971 ; 1972) have made a detailed calculation of this process; this calculation
is briefly outlined below.

Under most astrophysical conditions the spectral distribution of background low-
energy photons can be taken to be the Planck function

K(e) =

8tt

hV exp(e/kT)-l

(VIII.A-4)

,-26

p-p + p-a+a-p + a-a **

(H> --— ^

^^ \

^\ \

\ \

\ \
\ \

\ \

o-"

III 1 1 1 1 1 1 1

\ \

. 1 \ \

10"'
E,(GeV)

Figure VIII.A-1. The integral spectrum of 7-rays generated by cosmic-ray
interactions with interstellar matter. Nuclei heavier than helium do not
contribute significantly (Stecker, 1970).

DIFFUSE BA CKGROUND OF X-RA YS 189

The angular distribution of these photons is isotropic, that is

d cos 0

= constant (VIII.A-5)

The exact expression for the differential Compton-scattering cross-section is
quite involved. However, simplication occurs because the mean energy of
the X-ray generated by this process is usually much smaller than the energy
of the electron involved in the scattering. Accordingly, the differential
scattering cross-section for the emission of an X-ray photon of energy E in
a collision of an electron of energy e with a photon of energy e integrated
over the angular distribution of the incoming and outgoing photons becomes
(Hayakawa, 1969)

l\ c (~~2\ c 2

■n (mcz) Ev (mc*) E

- -r
4

a (E,e,E )=-r2 \ _ 2-x - ^ *- +

x' 4 e E3e2 E E3e

(VIII.A-6)

4E (W) E 8Ee

-^Cn / x + r-2

E 4E2e mc2

On making the substitutions

.2 „ BE

B = (mc2) , D = vrr 2/4 and x = — f— (VIII. A-7)

e 4E2e

this reduces to (Blumenthal and Gould, 1970)

8DB

a = — ( 1 + x - 2x2 + 2x Cn x) (VIII. A-8)

c Eze

Consider now a delta-function spectrum of electrons, 5(E-E ), generating
X-rays by Compton scattering against a thermal photon field (Equation
VIII.A-4).

This X-ray spectrum is shown in Figure VIII. A-2. In view of the fact that the
universal thermal background of microwave photons is the most relevant to
the discussion of the isotropic component of the X-ray background, the plot
in Figure VIII. A-2 corresponds to T = 2.7 K. The spectral shape for any
other temperature is obtained by simply sliding the same curve by a factor
T/2.7° along the X-axis on a log-log graph.

190

THEORY

The most important feature that is to be noticed in this figure is that the emis-
sion by electrons of single energy is over a very wide bandwidth, extending over
a factor of 20 in X-ray energies, at half-maximum. Because of this large band-
width, any kink or peak or other spectral feature in the electrons is smeared out
over an extremely broad energy region of the X-ray spectrum. Apart from the
broad bandwidth, the mean energy of the X-ray depends quadratically on the
electron energy. This relationship further contributes to the smoothing of the
X-ray spectrum relative to the electron spectrum.

The Compton scattering of 2.7 K photons in the intergalactic space by cosmic-
ray electrons leaking from galaxies could lead to an important contribution to
the X-ray background (Brecher and Morrison, 1969). In view of this, we

I I 1 I I Mil

I I I I I I III

Mill

10"

10"
X

10"

, a = 28*108 MeV

Figure VIII.A-2. The X-ray flux emitted in collisions of electrons of energy E
with the 2.7 K photon field is plotted as a function of X-ray energy. Notice
that the full width at half maximum is a factor of ~20 wide; also, the mean
X-ray energies related quadratically with the electron energy. These effects
tend to yield an X-ray spectrum that is very much smoother than the electron
spectrum. This fact was used to show that Compton scattering of the 2.7 K
photons is not a significant source at ~30 keV (Cowsik and Kobetich, 1972).

DIFFUSE BA CKGROUND OF X-RA YS 191

calculated the spectral shape of the electrons in the intergalactic space using
the radio data of Lang and Terzian (1969). The expected X-ray spectrum is
shown in Figure VIII. A- 3 marked as L. . The normalization of this curve is
arbitrary.

DIFFUSE X-RAY FLUX BELOW A FEW keV

Since the early observations by Bowyer et al. (1968), there has been a sub-
stantial progress in our understanding of the diffuse flux at ~250 eV. The
observations and related theoretical considerations are reviewed compre-
hensively by Silk (preprint).

The comparison of the Compton X-ray flux from the intergalactic space
(L. in Figure VIII. A-3) with the observed data indicates that this process
may contribute significantly to the background below a few keV. However,
since the normalization of this curve is somewhat arbitrary it is reasonable to
expect that only a part of the observed flux indeed arises through this
process. In view of the fact that our galaxy emits significantly in this band-
width one may expect that the diffuse background is generated as a super-
position of emission of all the galaxies in the universe. This suggestion was
first made perhaps by Silk (1970) and has had much experimental confirmation
due to the observation of several extragalactic sources using the Uhuru satellite
(Gursky et al., preprint, Giacconi et al., preprint). We show in Table VIII.A-2
(taken from a preprint of Silk) the contribution of various types of extra-
galactic objects to the X-ray background at ~2 keV.

Summing up the last column of the table shows that the sources contribute
significantly at ~2 keV. What is the spectrum of emission to be expected?
If we try to fit a thermal bremsstrahlung spectrum to individual sources,
the maximum temperature that is encountered in these extragalactic sources
is ~2 X 107 K. Following our discussion concerning Equation (VIII.A-2)
the cooling time of a plasma at this temperature is

1.96X 1011 (2X 107)1/2 ,, -

t * - — « 1015 s * 3 X 107 yr

taking n « 1 . This cooling time is much smaller than 1 /Hq so that there
would be a broad temperature distribution in the temperature of the plasma
leading to a spectral shape as given by Equation (VIII. A-3). Viz

p(E ) exp (-E /kT )

rv \/ g 2 r v x' max7

192

THEORY

io1

,(3) SUM OF 1. AND 2

D
><

10-5

I al, 1968
ol, 1969 o,b

(in, el ol, 1969

>l ol , 1969

el ol, 1969
O Hayakawa, el ol, 1969
▼ 8unner, el ol , 1969

• Bowper, el ol, 1968

• Toor, el ol, 1970
V Cunmngnom, el ol, 1970
A Scliwortz, el ol, 1970

D Rolhenllug el ol , 1968

O Sleeker and Deerenburg. 1969

A Melzger, el ol, 1964

» Kosluri Rongen. 1971

• Hoymes, el ol, 1969

• Pelerson, 1967
0 Veltle, el ol, 1970

• Domle, el ol, 1977
Golenelslii, el ol., 197)

• Sfiore, el ol, 1970

• Frye and Wong, 1969

• Clod, el ol, 1968 ; Gon

• Cnupp, el ol, 1969

GALACTIC

7 -RAY MODEL

Q INVERSE COMPTON, L|,"""\

(D SUM OF I. AND 2

10 10~3 10-2 10" 1 10 102

X-ray energy Ex (MeV)

Figure VIII.A-3. The X-ray energy flux is plotted as a function of X-ray energy
and compared with experimental data. Curve 1 is our calculation of the inverse
Compton scattering; curve 2 is calculated using the galactic 7-ray model of
Cowsjk (1971); and curve 3 is the sum of the two contributions. The experi-
mental data for X-ray energies E < 0.17 MeV and E > 10 MeV were taken
from the review paper by Silk (1970) and for 0.17 MeV< E < 10 MeV were
taken from Damle et al. (1971), and Golenetskii et al. (1971). The enhanced
emission at 2 keV < Ex < 200 keV is attributed to a hot (3 X 108K) intergalactic
gas.

DIFFUSE BA CKGROUND OF X-RA YS

Table VIII.A-2

Contribution of Identified Extragalactic Sources to the

Isotropic X-Ray Background*

193

Local space

Flux nL

4ttH

- o
(keV/cm • s-sr)

Class

Source Lx(2-10 keXOf8

density n
(N=0.03 Mpc-3)

Small

LCM

4X 1038

galaxies

SMC

1 X 1038

Adopted

2X 1038

ION

1.9

mean

Normal

M31

3X 1039

galaxies

Our

galaxy §

5 X 1039

Adopted
mean

4X 1039

3.8

Radio

galaxy

Cen A

8X 1041

Iff3 N

0.74

Seyfert
galaxy

NGC4151

2X 1041

0.02 N

3.8

AbeUI

clustersf % Centaurus

4X 1043

Virgo

1.5 X 1043

Adopted
mean

3 X 1043

2X 10'5N

0.57

Abell II

clustersf

Coma
Perseus

5X 1044
1 X 1045

Abell 2256 1 X 1045

Adopted
mean

8X 1044

5X 10-"6 N

3.8

Quasar

3C273

7X 1045

3X 10"8 N

0.18

*Data are taken from the Uhuru catalogue (Giacconi et al., 1972); H is set equal to

SOkm-s"1 Mpc"1.
fOniy those sources identified with clusters and known to be extended are included.
tThe Centaurus and Virgo clusters are not in Abell's 1958 catalogue; however, they

approximately correspond to Abell's richness class I.
§ Estimated X-ray luminosity of our galaxy (Seward et al., 1972).

194 THEORY

with T « 2 X 107 K. This spectral shape fits excellently the results below

max r r j

a few keV. Therefore, we conclude that free-free emission from the various
extragalactic sources would contribute significantly to the X-ray background
below few keV. However, the spectral shape is too steep to contribute sign-
ificantly at higher energies.

THE 2 TO 200 keV REGION AND POSSIBLE THERMAL BREMS-
STRAHLUNG OF THE INTERGALACTIC GAS

Investigating the possible origin of the X-rays in this energy band we noticed
(Cowsik, 1971) that thermal bremsstrahlung of a tenuous plasma at a tempera-
ture of 3 X 108 K had the right spectral form to fit the observations. An
emission measure,y*n 2 dfi « 1.3 X 1017/cm5, was required to give the
observed intensities. Assuming no clumping (that is, n independent of C)
and taking

= « 1028cm

one gets

n «3X 1CT6 cm"3 %

cnt 8yrGmH

for H = 50 km/s-Mpc. Therefore we suggested the possibility of a hot inter-
galactic plasma as a possible source of this background (Cowsik, 1971 ; Cowsik
and Kobetich; 1972). The thermal bremsstrahlung fit to the experimental data
after plausible subtractions of other emission mechanisms below 2 keV (see
previous section) and above 200 keV (see following section) is shown in
Figure VIII.A-4.

That a hot intergalactic medium could be the source in this region has been in-
dependently pointed out by Field (1972). Of course, the idea of a hot inter-
galactic medium is not new. It has been discussed in the context of continuous
creation of matter in the form of neutrons in the steady-state cosmology by
Gold and Hoyle (1959) and by Gould and Burbidge (1963). However,
Petrosian and Ramaty (1972) have provided arguments based on excessive
production of hard X-rays through the radiative decay of the neutron that
continuous creation of matter as neutrons is forbidden by X-ray observations.
The X-ray spectral observations in the region of 2 to 20 keV yield merely the
temperature and emission measure of the radiating plasma. Therefore, the
question arises as to whether the emission indeed comes from a hot inter-
galactic medium or from hot gas in various galaxies.

DIFFUSE BA CKGROUND OF X-RA YS

195

v 10-'-

io-1-

O Hayokawa, et al., 1969

• Toor, et al, 1970

V Cunningham, et al., 1970

A Schwartz, et al., 1970

a Rothenflug, et al., 1968

O Bleeker and Deerenburg, 1969

A Mefzger, ef al., 1964

I I

25 50 75 100 125

X-ray energy Ex (keV)

150

175

Figure VIII.A-4. The difference between the observed energy flux and the
calculated flux (see Figure VIII.A-3) in the energy interval 2 keV < Ex <
200 keV is plotted as a function of X-ray energy. The line represents the
thermal bremsst rah lung emission for a hydrogen plasma at 3.3 X 10 K.
The line of sight integral Tn N dl for this emission is 1.3 X 10l7/cm5. If
one assumes no clumping andj dl = 10 cm, one gets Ne « N « 3 X 10"
/cm3. Such a density is adequate to close the universe if H =55 km/sMpc.

We believe that there are indeed reasons that indicate that the hot intergalactic
medium is the most plausible explanation of this emission. The high degree of
isotropy as measured by Schwartz (1970) has been analyzed by Silk (preprint)
to show that one needs at least 107 sources in the sky to yield the required
degree of isotropy. This means that a reasonable fraction of the galaxies should
contain hot plasma at 3 X 108 K. In order that the spectrum of emission is
not transformed by free-free cooling (Equations VIII. A-2, and -3)

1.96 X 1011 (3X108)1/2 . J

>jj =3X 1017s

(VIII.A-9)

196 THEORY

which yields n 10"2/cm3. Even if one takes n ~ 10"2, one finds that about
5 percent of all visible galactic matter should be at a temperature of 3 X 108K
in order to generate an emission measure of 1.3 X 1017 cm"5. Firstly, most
of the mass of the galaxies is concentrated in stars (temperature ~ 10 K),
with gas contributing to less than 10 percent of the total mass. The galaxies
would definitely be unable to contain gravitationally such a large amount of
hot plasma.

There is a second argument in favor of a critical mass density existing in the
form of a hot intergalactic plasma. This argument essentially invokes the in-
tergalactic medium as a heat sink for the energy released during the synthesis
of heavy elements in supernovae exploding in the galaxies. It has been pointed
out that in our galaxy with a mass of ~10n solar masses one needs ~109
supernovae to generate the heavy elements. How much energy is released in
this process? We may estimate that about one solar mass collapses to a neutron
star in each event yielding an energy of G(l/2 M 0)/R. Thus the energy
generated per unit mass of the galaxy

QGMQ2 1 GMft ,.

u = 109 — • — n = a « 2 X 1018 erg/g

R 1011 M© 100R

using R « 106 cm. This released energy is not seen as electromagnetic radiation
in any frequency band, and this energy must therefore have gone into the
kinetic energy of matter; 5 X 109 ergs/gm corresponds to a temperature of
~1010 K. This shows that we need to have approximately 100 times as much
intergalactic matter as in the galaxies to absorb this energy so that the mean
temperature of the universe may not be too high. With a hot intergalactic
plasma of critical density (equal to ~50 times the mass density contributed
by the galaxies) one has a hot universe at ~3 X 108 K.

0.3 TO 3 MeV GAMMA RAYS

Stecker, Morgan, and Bredekamp (1971) have attempted to explain this flux of
7-rays as due to annihilation of matter and antimatter at z * 100 in a baryon
symmetric universe. However, neither the absolute intensity nor the isotropy
of this radiation has been established. Therefore, we wish to investigate here
the possibility that this radiation could be of local, galactic origin. In fact,
there are indications in the cosmic-ray electron spectrum that such emission
could be taking place from our galaxy. Before discussing this galactic source
in detail we must emphasize that the burden of proof lies with experiments.
Should they show that the radiation is indeed isotropic, one has to give away
the galactic model, which is discussed below.

The cosmic-ray electron spectrum is well measured in the region of ~100 MeV
to ~100 GeV; it is shown in Figure VIII. A- 5 after correcting for solar

DIFFUSE BACKGROUND OF X-RA YS 197

modulation effects. The spectrum below a few GeV has a spectral index of
~1 .6, but steepens to an index of ~2.6 above a few GeV. This is not the
true electron spectrum that is injected into the interstellar space by the
cosmic-ray sources but is contaminated by interstellar secondaries generated
by the nuclear component of cosmic rays. The positron flux gives a very good
estimate as to the amount of this contamination. After subtracting the
secondaries, the spectrum of electrons injected by the sources is shown in
Figure VIII.A-6. Since the processes of cosmic-ray acceleration are electro-
magnetic in nature, one may safely assume that the spectrum of electrons
accelerated by the sources is a simple power law with an index of 2.6 similar to
that of the nuclear component. The difference in the energy between the
accelerated spectrum and the injection spectrum must have been radiated
away. If part of this radiation is due to Compton scattering against optical
frequency photons then one obtains a 7-ray luminosity (Cowsik, 1971)

UEJ

^~(/3-lWl6 E7e)

c P-i

3E\p-1

*'fj7)

t± (VIII.A-10)

p-i v

K"(fj?n

0-1

|3, S, EH, ET and p are constants derived from the electron spectrum and
e^3 X 10"6 MeV is the assumed mean energy of the optical photon. The in-
tensity as seen by a detector having isotropic response is shown in Figure VIII.
A-7. One notices that the spectrum is insensitive to the parameter p. The
general shape of the curve is essentially determined by |3^2.6, the spectral
index of the cosmic-ray electrons at low and high energies, respectively.
Despite this elegant fit to the data, one has to wait for measurements of
angular distribution in this energy region before the galactic nature of the
MeV 7-ray fluxes can be established.

GAMMA RAYS IN THE 100-MeV RANGE

The pioneering work of Clark, Garmire, and Kraushaar (1968) established
the existence of a line source coincident with the galactic disk with an
enhancement around the galactic center and a possible isotropic component.
The measured intensity in the direction of the center is ~1047/cm2;S-r,ad,
the line source elsewhere is about a third of this ~2 to 3 X 10s 7/cm2 • s • rad.

198

THEORY

10"

_i 1 1 1 1 i i i i i i 1 1 1

Pio-3

-^ 10"

10"

10"
5x10"

i Mill r

TOTAL ELECTRONS

I I I II

J I I I I I I

J L

005 01

10 10

Energy (GeV)

Figure VIII.A-5. Electronic component of cosmic rays in the
interstellar space. The secondary electron flux generated by
the nuclear component is normalized using positrons.

DIFFUSE BA CKGROUND OF X-RA YS

199

s= 10-

I MINI

1.0

Energy (GeV)

Figure VIII.A-6. The injection spectrum of primary cosmic-
ray negatrons. Hatched area indicates uncertainties in the
estimate. The predictions of the model are f or p = 2 (.-.-.-.)

and p = 3 ( ); E = 0.5, 0.7, 1 .0 GeV starting from top.

Since it is the difference between the accelerated power-law
spectrum ~E"2-6 and the injection spectrum that governs the
7-ray intensities the uncertainty in the 7-ray fluxes are with-
in a factor of ~2.

200

THEORY

100 200

Energy (MeV)

Figure VIII.A-7. A well defined background 7-ray flux (absolute normalization
within X2) is predicted by model. For E7 «1 MeV, there is a large flux of
intergalactic origin. The model predicts correctly the primary gamma-ray slope
and intensity that would generate the experimental response shown (Anandetal.,
1969; data from Silk, 1970).

DIFFUSE BACKGROUND OF X-RA YS

201

The line source has been explained by Stecker (1969b) as due to production
and subsequent decay of neutral pions by cosmic rays. If the same mech-
anism should yield the enhancement of the intensities near the center, one
needs a substantially large enhancement of gas density near the central
regions of the galaxy. There is no evidence, direct or indirect, for such an
enhancement. On the other hand, there is evidence that the density of stars
increases considerably towards the galactic center. The density distribution
of stars as a function of galacto-centric distance oo, is shown in Figure VIII.
A-8 (Perek, 1962). The increase in mass distribution of stars towards the
center is 1/cj3 and would be that of the distribution of starlight. With such

-5

cokpc

Figure VI 1 1. A-8. The mass distribution in the galaxy in units of Mq/pc3 is
shown as a function of cylindrical coordinates centered at the galactic cen-
ter (taken from Perek, 1962).

enhanced starlight density, the Compton scattering of the cosmic-ray elec-
trons of these photons would provide an intense 7-ray source. Preliminary

202

THEORY

calculations of the angular distribution expected through this process is shown
in Figure VIII.A-9 (Cowsik and Hutcheon, 1971). The actual calculations
yielded only 70 percent of the intensity towards the galactic center as due to
this process. If one adds the line source due to 7r° -> 2? decay contributing
~30 percent with a flat dependence on galactic longitude (Stecker, 1969b)
then the emission from the galactic disk can be explained completely.

GALACTIC LONGITUDE

Figure VIII.A-9. The gamma-ray intensities that are calculated using starlight
distribution implied by Figure VIII.A-8 are compared with results of Clark et al.
(1968). The preliminary theoretical estimates are multiplied by ~1.4 and then
averaged over the aperture of the detector. It is seen that the Compton scattering
of starlight contributes negligibly beyond ~60° galactic longitude. Beyond this
point the n0^-2y process discussed by Stecker (1969) dominates and should be
added to the Compton fluxes to make a detailed fit to the observations.

What is the spectrum of 7-ray generated through the Compton scattering of star-
light? The cosmic-ray electron spectrum has a spectral slope of j3j^l .6 below
~3 GeV and a slope of (3 «2.6 above ~3 GeV. The maximum and the mean
energies of the scattered photons in the Compton process are given by

and

7max \m

4 /E V

<v=i(-)

(VIII.A-11)

DIFFUSE BACKGROUND OF X-RA YS 203

The spectral slope of Compton 7-rays will be al = (j3j- 1)1/2 = 0.3 at low
energies and a = (j3 -1)1/2 « 0.8 at high energies. The typical 7-ray energy
at the transition that takes place can be calculated by using Equation
(VIII. A- 1 1), with E « 3 GeV, and the mean energy of the starlight photon
~3 eV (corresponding to a temperature of 104 K).

E 4

7-transition «—

3X 109

5X 105
150MeV

X 3eV

This means that below ~150 MeV the spectral slope would be ~0.3 and as such
would be very difficult to distinguish from a ir° source. We now wish to em-
phasize the inevitability of the existence of a Compton source of 7-rays in the
central regions of our galaxy. There is evidence both for the existence of high
density of starlight near the galactic center (as discussed before) and for the
relativistic electrons through their synchrotron emission; thus the Compton
scattering must occur leading to significant 7-ray flux from the region of the
galactic center. Detailed spectral shapes and sky maps due to this process will
be presented at the Cosmic Ray Conference at Denver by Sullivan and Cowsik.

The isotropic component at ~100 MeV has an intensity of ~10~57/cm2-s-sr.
At this moment it is hard to pinpoint a precise source for this radiation. We
wish to add that the existence of an extended halo to our galaxy may con-
tribute significantly to this flux, and also that a truly extragalactic component
in this energy region cannot be excluded.

THE 100 MeV 7-RAY FLUX AND A LIMIT ON THE ENERGY DENSITY
IN THE SUBMILLIMETER BACKGROUND

As pointed out in the previous section, the line source of 7-rays away from the
galactic center can be completely explained as due to the decay of 7r°'s produced
by nuclear cosmic rays (Stecker, 1969b). This means that any other source of
7-rays must be very weak indeed. One such could be provided by the existence
of intense submillimeter radiation which would then be scattered to 7-ray
energies by cosmic-ray electrons. Thus, one may use the 7-ray fluxes to put
stringent limits on the microwave background. This is done in Table VIII.A-3
taken from Cowsik (1972). From this table it is clear that the energy density in
any radiation background over and above the universal thermal background at
2.7 K should be less than 0.6 eV/cm3 . In Figure VHI.A-10 this limit is shown
in comparison with observations at microwavelengths.

204

THEORY

~co
o

CD
CO

^3 10

id,7h

PENZIAS and WILSON (1965)
MUEHLNER ond WEISS (1969, 1971)
THADDEUS (1970), B0RT0L0T etal (1969)
CORNELL

PRINCETON
BLAIR etal (1971)

THIS WORK

■2 7°K BLACKBODY

O/"

I 10

Wavelength cm

10'

Figure VMI.A-10. Measured background radiation fluxes are compared with that
expected from a blackbody at 2.7 K. The 7-ray fluxes measured by Kraushaar
et al. (1972), put a stringent limit on the intensities allowable at submillimeter
wavelengths. In plotting our upper limit of ~0.6 eV/cm3 we have assumed that
the background radiation has a bandwidth equal to that of the detector of the
Cornell instrument.

SUMMARY

Thus it appears that one has a reasonable explanation for a good part of the
diffuse X-ray and 7-ray background that is observed over six decades in
energy. Thermal sources seem to dominate up to an energy of ~200 keV.
Angular distribution measurements are essential in choosing between galactic
and universal models for the intensities in the MeV region. The source of
100 MeV 7-rays from the disk and galactic center seem to be well under-
stood as due to the decay of neutral pions and Compton scattering of star-
light, respectively. These observations put a stringent limit on the energy
density in any possible radiation background at submillimeter wavelengths.

I cannot close this review more effectively than by making a call for all
experimentalists in the field to measure the angular distribution of photons
in the MeV range, which is of very great astrophysical and cosmological
importance, for it relates either to the cosmic-ray sources in our galaxy or
to annihilation of antimatter in baryon symmetric cosmologies.

DIFFUSE BACKGROUND OF X-RA YS

205

Table VIII. A-3
Gamma-Rav Fluxes at E7MOO MeV: Theory and Experiment*

Region of Sky Scanned

Source

60°<l<30o,b = 0°

b = tt/2 (halo)

(disk)

Flux (cm2s rad)-1

Flux (cm2s sr)-1

Experiment
p+H^7r°^27

e+ (2.7°)^e+7

(3.4 ±0.6)X 10*
>2.6 X 105
>1.1 X 10s

(3 ± 0.4) X 10"5
3.7 X lO"6
1.1 X 10"5

Residual

<0.9X 10s

2.3 X 10-5

e+e (sub-mm)->e+7

4.5 n . X 10-8

4.4n,X lO-8
ph

Maximum number den-
sity of sub-mm quanta
Corresponding energy
density

n . <200 cm"3

p (sub-mm) = N
<0.25 eV

n . <500 cm-3
ph

p<0.6 eV/cm3

*The radio disk is assumed to extend up to ~1 kpc above the galactic plane in making the
theoretical estimates. Because of the Gaussian response of the detector with angles, the
expected counting rates increase more slowly than that proportional to the assumed thick-
ness of the disk. Note that the estimates from the halo direction are uncertain and are to
be given much lower weight.

ACKNOWLEDGMENTS

I want to thank Professor P. Buford Price and the members of his group for their
active interest in this work.

(Supported in part by NASA grant NGR 05-003-376.)

REFERENCES

Anand, K. C, G. Joseph, and P. J.Lavakore, 1969, Proc. Indian Acad. Set, 71,
p. 225.

Blumenthal, G. R., and R. J. Gould, 1970, Rev. Mod. Phys., 42, p. 237.

Bowyer, C. S., G. B. Field, and J. Mack, 1968, Nature, 223.

Brecher, K., and P. Morrison, 1969, Phys. Rev. Letters, 23, p. 802.

Clark, G. W., G. P. Garmire, and W. L. Kraushaar, 1968, A strophys. J. Letters,
153, p. L203.

Cowsik, R., 1971, Proc. 12th Int. Conf. on Cosmic Rays, 1, p. 334.

206 THEORY

., 1972, Nature Phys. Sci, 239, p. 41.

Cowsik, R., and I. D. Hutcheon, 1971, Ibid, 1, p. 102.

Cowsik, R., and E. J. Kobetich, 1971, Ibid, 1, p. 38.

, 1972, Astrophys. J., 177, p. 585.

Damle, S. V., R. R. Daniel, G. Joseph, and P. J. Lavakaro, 1971, Proc. 12th
Int. Conf. Cosmic Rays, 1, p. 84.

Felten, J. E., 1972, X-Ray and Gamma Ray Astronomy, Proc. oflAU
Symposium No. 55 (Madrid), H. Bradt and R. Giacconi, eds., D. Reidel,
Dordrecht, Holland.

Field, G. B., 1912, Ann. Rev. Astron. and Astrophys. , 10, p. 227.

Gold, T., and F. Hoyle, 1959, Paris Symp. on Radio Astronomy, IAU
Symposium No. 9, Bracewell, ed., p. 583.

Gould, R. J., and G. R. Burbidge, 1963, Astrophys. J., 138, p. 969.

Hayakawa, S., 1969, Cosmic Ray Physics, John Wiley and Sons, New York,
p. 609.

Kraushaar, W. L., G. W. Clark, G. P. Garmire, R. Borken, P. Higbie, G. Leong,
and T. Thorsos, 1972, Astrophys. J., 177, p. 341.

Lang, R. R., and Y. Terzian, 1969, Astrophys. J. Letters, 3, p. L29.

Perek, L., 1962, Adv. in Astron. and Astrophys. , 1, p. 165.

Petrosian, V., and R. Ramaty, 1912, Astrophys. J. Letters, 173, p. L83

Schwartz, D. A., \910, Astrophys. J., 162, p. 439.

Seward, F. D., G. A. Burginyon, R. J. Grader, T. Palmieri, and J. Stoering,
191 \, Astrophys. J., 169, p. 515.

Silk, J., 1970, Space Set Rev., 11, p. 671.

Stecker, F. W., 1969a, Astrophys. J., 157, p. 507.

, 1969b, Nature, 224, p. 870.

, 1971, Cosmic Gamma Rays, Mono Book Corp., Baltimore.

Stecker, F. W., D. L. Morgan, and J. Bredekamp, 1971, Phys. Rev. Letters,
27, p. LI 469.

DIFFUSE BA CKGROUND OF X-RA YS 207

DISCUSSION

Member of the audience:

What is the evidence for the lifetime being 1014s for the electrons as distinct
from nuclei?

Member of the audience:

You may say that Daniel's measurement indicating a steepening of the spectrum
at 100 GeV may be taken as evidence, or you may say that electrons at 200 GeV
propagate with the nuclei and you get a lifetime similar to that. I'm sure you
are quite an expert in the field, and you know the answer.

Member of the audience:

There's a possibility but no proof positive at all.

Member of the audience:

We cannot prove anything, we can only give plausible statements.

Member of the audience:

Would you take issue with him? (Laughter)

Member of the audience:

Give or take an order of magnitude.

Stecker:

I'm not quite sure how you can calculate such a detailed 7-ray spectrum
(Figure VII. A-7) when you took such crude numbers.

Cowsik:

I just calculated this using 1-eV photons and the spectrum that I get is what I
have shown in Figure VIH.A-7. It's not detailed in the sense that it may have
fluctuations or errors in it which may be a factor of two or something like that.
From today's discussion, I conclude that we do not know the 7-ray flux to that
kind of an accuracy.

Stecker:

You say there are enough 1-eV photons and cosmic-ray electrons to account
for the total flux, and that it would be anisotropic?

208 THEORY

Cowsik:

Yes. In MeV/MeV units, this would turn out to be 3 X 10"2, or something
like that.

Stecker:

These are Compton interactions you are talking about?

Cowsik:
Yes.

Stecker:

Because there are many calculations going back to those of Morrison?

Cowsik:

There is an essential difference. I'm not taking all the electrons in the galaxy
and putting them inside a typical radiation density that exists in the galactic
disk and asking what would be the spectrum that comes up. I'm doing
something different (see text).

Something may be happening in the sources in which the originally accelerated
power law is reduced. This can happen, let us say, if we can take an energy
dependent leakage lifetime from the source region that can easily kill the
lower-energy electrons more efficiently than the high-energy electrons. In
fact, recent cosmic-ray data have evidence that indicates that even the nuclei
in cosmic rays may have been stored in the sources.

Member of the audience:
I think it goes the other way.

Cowsik:

One has to discuss that. I don't want to discuss it in detail, but certainly the
model that I'm using is allowing the electrons to be killed by this factor.

Vette:

What you are talking about here for the 7-ray production, is it in the galaxy?

Cowsik:

What I'm talking about here is in the sources.

DIFFUSE BACKGROUND OF X-RA YS 209

Member of the audience:

The source is assumed to be inside of the galaxy?

Cowsik:
Yes.

Vette:

But you wouldn't expect this to be isotropic?

Cowsik:

Yes, this won't be isotropic. This will show enhancement in the galactic disk,
but I do not know how much. The experimenters will have to comment on
what kind of limits one can place on the anisotropy of MeV radiation. (See
paper of Kniffen et al., Chapter IV.)

Stecker:

You're talking about essentially a large anisotropy because all the 7-rays are
coming from sources in the galactic plane, and I'm wondering if any of the
observational people here might mention their angular resolution and what
their upper limit is on their anisotropy.

Metzger:

The slide I showed from Apollo- 16 shows a clear demonstration of anisotropy
and the background of the energy range up to about half a MeV. We don't
have the statistics yet to look beyond that with the Apollo- 16 data. We
have looked with Apollo- 15 data and we see anisotropy in a couple of cases
corresponding to two of the five or six fixed position measurements which
are made for the X-ray spectrometer during transearth coast, but we have not
yet to our satisfaction pinpointed the source of that anisotropy.

Stecker:

What's the number on the anisotropy?

Metzger:

At the most, 5 percent above the mean.

Stecker:

And you show just where the Crab Nebula and the galactic center were?
What would be your guess of the average anisotropy if you subtracted it?

210 THEORY

Metzger:

I'd say less than 10 percent.

Vette:

Steve White, do you have any comments on the anisotropy from any of your
measurements?

White:

Well, unfortunately, we had only one flight and we dedicated only a small
amount of time to making the 7-rays in that one flight. We were looking
primarily in the direction of the antigalactic center during that time although
we had a little time devoted to looking outside, so we can't say anything.

Steigman:

Didn't you fellows see a marked anisotropy in the galactic plane?

Vette:

Oh, yes, at 100 Me V.

Member of the audience:
That's at high energy.

Vette:

You're not talking about 100 MeV, you're talking about 1 MeV.

Cowsik:

One- to 10-MeV range.

Vette:

There really hasn't been a good definitive measurement in the directional
aspect of it. That's why I think some of the detectors we heard about today,
particularly the Compton telescopes, do offer some possibility of making
measurements there we haven't made today.

(See also related discussions by Kniffen et al., Chapter IV, and Stecker,
Chapter IX.)

Chapter IX

MECHANISMS FOR PRODUCTION OF THE
DIFFUSE GAMMA-RAY CONTINUUM
RADIATION

F. W. Stecker*

Goddard Space Flight Center

BASIC MECHANISMS

The basic mechanisms expected to be important in the production of cosmic
7-radiation were suggested by Morrison in a classic paper in Nuovo Cimento in
1958. They are Compton interactions with low-energy photons, bremsstrahlung
interactions, cosmic-ray induced tt° production, and matter-antimatter annihi-
lation. Of these four mechanisms, the first two involve cosmic-ray electrons
and are electromagnetic processes, whereas the last two involve nucleons,
mainly protons, and are strong interaction processes. Above 511 keV the
7-radiation from matter-antimatter annihilation arises mainly from the decay
of 7r°-mesons produced in the annihilation process, so that the kinematics
involved in the last two processes is similar. Because this paper will be
concerned mainly with diffuse continuum radiation rather than line radiation
or radiation from point sources, the discussion here will be restricted mainly
to the above four processes. (For a treatment of the theory of the production
of cosmic-line radiation, see Clayton, Chapter XI. A.)

COMPTON INTERACTIONS

The most astrophysically significant role which Compton interactions are
expected to play in cosmic 7-ray production involves the interactions of
relativistic cosmic-ray electrons with low-energy photons of the universal
2.7-K microwave blackbody radiation field. The microwave photons have an
average energy near 10"3 eV and a number density of ~400/cm3, considered
to be uniformly distributed throughout the universe. Compton interactions
with cosmic-ray electrons can then produce 7-rays with typical energies of

<E >=10"372eV (IX.A-1)

^Speaker. 211

212 THEORY

where 7 = (E /rn c2) is the Lorentz factor of the cosmic-ray electron. Thus
a 50-GeV electron with a Lorentz factor of ~10 will typically produce
7-rays of energy ~10 MeV through Compton interactions with 2.7-K photons.
We can define the "spectrum" of 7-rays from a single Compton interaction as
the normalized probability distribution of 7-rays of energy E expected to be
produced by an electron of energy E . Such a spectrum turns out to be flat
and rather broad around the average energy <E > (Heitler, 1954; Jones, 1965).
Because <E > <*72 , the spectrum of 7-rays produced by a power-law cosmic-
ray electron spectrum of the form K E "re will also have a power-law form
K E ~r7, but with T = (r + l)/2. in fact, for interactions with a blackbody
spectrum of low-energy photons at temperature (T), the Compton-generated
7-ray spectrum is given in photons/(cm2 -s-sr-GeV) by

I(E ) = 6.22X10-21L[10"2-962ref(r)] KT<re+5>/2E -<re + D/2

7 e e y

(IX.A-2)

where L is the path length (cm) over which production occurs and E is in
GeV with the factor f(r ) ~ 1 given by Ginzburg and Syrovatskii (1964).
For example, if r, ~ 2.6, then T ~ 1.8 for galactic cosmic-ray electrons.
Because the 2.7-K blackbody radiation is believed to be universal, Compton
interactions have been invoked to explain the cosmic X-ray background
spectrum where the observed r ~ 2 and to set limits on the metagalactic
cosmic-ray electron intensity wM to show that it must be much less than
the galactic value, that is, wM << wG (Felten, 1965; Gould, 1965;
Fazio, Stecker, and Wright, 1966; Felten and Morrison, 1966; Cowsik,
Chapter VIII.A; Ginzburg, Chapter X.A).

ELECTRON BREMSSTRAHLUNG

Bremsstrahlung interactions are expected to take place between cosmic-ray
electrons and interstellar and intergalactic gas and may be significant in pro-
ducing low-energy 7-rays and X-rays both in the galaxy and in intergalactic
space (Cowsik, Chapter VIII.A). The probability distribution spectrum for
7-rays from bremsstrahlung of a cosmic-ray electron of energy Ee is quite flat
and may be approximated by

f(E7lEe)s

E -1 for 0 < E < E

e -v e

(IX.A-3)
0 otherwise

DIFFUSE GAMMA-RA Y CONTINUUM RADIA TION 21 3

so that the 7-ray production spectrum is given by

. f _, Ie(>E7)

Ib(E7) = X'1 / dr (p (T)> ^ (IX. A-4)

J o f

where p is the matter-density of the gas in g/cm3 and X is the average radiation
length for interstellar matter and is ~65 g/cm2. We may also write Equation
(EX. A-4) in terms of the atomic density (n, in cm"3) and the path along the
line of sight (L in cm) so that

Ie(>E7)
Ib(E7) = 3.4 X 10"26 nL- (IX.A-5)

It follows from Equation (IX.A-5) that for bremsstrahlung from cosmic-ray
electrons following a power-law spectrum K E ~re, T = F (for relativistic
electrons) so that, in general, the 7-ray spectrum from bremsstrahlung is
steeper than that from Compton interactions.

NEUTRAL PION DECAY

We next discuss the 7-radiation from the decay of 7r°-mesons produced by
cosmic-ray interactions between high-energy nucleons and gas nuclei in
interstellar and intergalactic space. This process has received the most
attention because it now appears that 7r°-decay 7-rays from cosmic-ray
interactions may account for almost all of the 7-radiation above 1 00 MeV
observed in our galaxy (Fichtel et al., 1972; Clark et al., 1968; Stecker, 1969a;
Stecher and Stecker, 1970; Cavallo and Gould, 1971; Ginzburg, Chapter X.A),
because it has long been recognized as an important process for cosmic 7-ray
production (Morrison, 1958; Pollack and Fazio, 1963; Ginzburg and
Syrovatskii, 1964), because it has the most difficult spectrum to calculate
theoretically (Hayakawa et al., 1964; Dilworth et al, 1968; Stecker, 1970;
Cavallo and Gould, 1971 ; Levy and Goldsmith, 1972) and because various
theoretical calculations are somewhat contradictory. I do not intend to break
the tradition, in fact, I hope to help resolve here some of the contradictions
that have arisen among the theoretical calculations.

Figure IX. A- 1 shows the type of 7-ray spectra obtained from the decay of
7r°-mesons with various simple energy spectra f(E ). Typically the spectrum
is flat near m c2/2 ~ 70 MeV, and symmetric about this value on a logarithmic
energy plot. These characteristics can easily be shown from the kinematics of
7r°-decay (Stecker, 1971a) and they will not be repeated here. Figure IX.A-2
shows how a typical 7r°-decay 7-ray spectrum can be built up from an arbitrary

214

THEORY

f(Eir)x8(ET-E0;

f(E,r)xS(Eir-E1)

El>Eo

f(ET) =

(E0-EJ for E^E^E,

'B A' ,w" A
0 OTHERWISE

F(Er)

inE.

^nE.

'A I

■A2

JnE.

Figure IX.A-1. Gamma-ray spectra from the decay of neutral pions for
various simple pion energy distributions (v = m c2/2).

pion-energy spectrum, and that the spectrum always has a maximum at
~70 MeV. Figure IX.A-3 shows the differential 7-ray spectrum obtained by
Stecker (1970), illustrating the various expected characteristics. Figure
IX.A-4 shows a comparison of the integral spectra obtained by Stecker (1970)
and Cavallo and Gould (1971) normalized to compare the shapes obtained.
The wiggles in the spectrum represent artifacts of the assumed pion-production
models and should not be taken too seriously. The shapes of the two spectra
are in good agreement and probably represent an accurate approximation to
reality within the uncertainty indicated by the wiggles.

DIFFUSE GAMMA-RA Y CONTINUUM RADIA TION

215

v JnE?

Figure IX.A-2. Idealized superposition of 7-ray spectra from the decay
of pions having various energy distributions (v = m c2/2)

Ey(GeV)

Figure IX.A-3. The calculated differential production spectrum of 7-rays
produced in cosmic-ray interactions in the galaxy based on the "isobar (i)-
plus-fireball (f)" model of Stecker (1970).

216

THEORY

QJ*r)
Q o(>0.01)

0.1

STECKER (1970)

CAVALLO AND GOULD (1971)

0.01

0.1
Ey(GeV)

Figure IX.A-4. A comparison of the shapes of the integral galactic pion-decay
energy spectra calculated by Stecker (1970) and Cavallo and Gould (1971).
The total production rate is normalized to unity.

The largest discrepancy between the various calculations is in the total 7-ray
production rates calculated by various workers. These rates are compared in
Table IX. A- 1.

Pollack and Fazio (1963) and Dilworth et al. (1968) obtained total 7-ray
production rates per hydrogen atom which would be equivalent to roughly
1.1 X 10"25s"1 and 1.0 X 10"25s"1 respectively, for energies above 100 MeV.
This corresponds to the quantity

Q ,7r°(>100MeV) = 4 7rI (> 100 MeV)/<nL>

(IX.A-6)

Pollack and Fazio used the observed cosmic-ray spectrum at the earth for
their calculations. Stecker (1970) used a demodulated cosmic-ray spectrum
to estimate the galactic cosmic-ray spectrum and, for this reason, obtained a
slightly higher value of 1.3 X lO"25^1 for Q (see following discussion). From
the OSO-3 satellite observations, Kraushaar et al. (1972) obtained an upper
limit for Q of 1.6 X lO"25^1 ; recently Stecker (1973) obtained a theoretical
upper limit of ~1.5 X 10"25s"1, assuming a maximum solar demodulation
effect to obtain a maximum galactic cosmic-ray spectrum as deduced by

DIFFUSE GAMMA-RA Y CONTINUUM RADIA TION

217

Vi 2

c S

° 2

C/3 "J3

^ 2

JS +3

<3J

£j

<+J

Cm
ft

•^

•5

■—

E
s

O £

'03

1/)

^j-

*-J

,_,

, i

,__,

^_M

,__i

,_i

<L>

I/O

PL)

^-^

/"^

/-^

S^pX

s-^

-o

(U
l-l

1— 1

•a

3
o
O

"rt

T3
C

r-H

t-H

<u
o

3
cd

218

THEORY

Comstock et al. (1972). Thus, all of the above values for Q n0 are basically
consistent. The value obtained by Cavallo and Gould (1971) appears to be
somewhat high compared with the others, but a value of ~1.3 X 10"25 s"1
falls within their 30-percent error bracket. It is the author's opinion that a
value of 1 .3 X 10"25 s"1 is close to a "best value" for Q 0. The value of
Levy and Goldsmith (1972) is a factor of ~2.5 higher and requires some
discussion.

Figure IX. A- 5 shows an up-to-date summary of the accelerator data on total
cross section (a) times multiplicity (f) for neutral pion production in p-p
interactions for energies up to ~1500 GeV shown as a function of kinetic
energy (T) (Stecker, 1973). These data are well approximated by the broken
power law

%(T)?o(T)^'

■n n

10-25T7.64cm2

0.4 <T< 0.7 GeV

(IX.A-7)

.4 X 10-27Ta53cm2 T > 0.7 GeV

0.1

^CHARLTON.

DAO. et al (1973K
971, 1972) -~^_, ^»-^'^^\^

B0GGILD, et al (1971)
MEUSSINOS, et al (1962)
PICKUP, et al (1962)

et al (1972)
-DODD. et al (1961)
•CAMPBELL, et al (1973)

EUHOFER.
et al (1971. 1972)

EISNER, et al (1964)
HUGHES, et al (1956)
BATSON. et al (1959)
BARNES, et al (1961)

MESHCHERIAKOV, et al (1956)

0.1

T(GeV)

Figure IX.A-5. Cross section times multiplicity for neutral pion production in
p-p interactions as a function of incident kinetic energy (from Stecker, 1973).

DIFFUSE GAMMA-RA Y CONTINUUM RAD I A TION 219

as the reader can verify from the figure. Taking the cosmic-ray spectrum
I (T) = 0.15 T2-2 cm"2 s"1 sr"1 GeV1 used by Levy and Goldsmith (1972),
the total 7-ray production rate from p-p interactions is given by

cLyH = 87r / dTI(T)a0(T)?0(T)

r 0.7 /• °°

= 3.77X10"25 / T5-44dT + 3.17X 10"26 / T167dT

J 0A J 0.7

= 0.66X lO"25^1
Adding in the contribution from p-a, a-p, and a-a interactions in the galaxy
brings the total production rate per hydrogen atom up to ^ 10"25 s"1 . There
is, of course, some uncertainty in the assumption of the true demodulated
galactic cosmic-ray spectrum as distinguished from that observed at the earth.
However, using the upper limit to the demodulated cosmic-ray spectrum given
by Comstock et al. (1972), an upper limit of (1.51 ± 0.23) X 10"25 s1 on the
7-ray production rate is obtained, with the error bracket reflecting the experi-
mental error in the accelerator data on af . The above value is consistent with
the value of 1.6 X 10"25 s"1 given by Kraushaar et al. (1972), which also
represents an upper limit since it does not take account of the additional
contribution from cool H and H2 which may be adding to the observed flux.

Why then is there such a large discrepancy between the results presented here
and those obtained by Levy and Goldsmith? The answer appears to lie in the
difference between assumptions on the total cross section for n° -production
as a function of energy and the multiplicity (f) assumed. While we have
chosen to rely on measurements from accelerator experiments, Levy and
Goldsmith adopt a theoretical multiplicity law based on the scaling hypothesis
which may hold above 100 GeV. This logarithmic multiplicity law has some
empirical support in the cosmic-ray measurements above 70 GeV cited by
Levy and Goldsmith, but is contradicted in other cosmic-ray measurements
so that the situation at high energies is not as yet clear (Sreekantan, 1972).
The logarithmic multiplicity law, based on the scaling prediction, depends on
arguments that hold asymptotically in the high-energy limit and that do not
appear to be valid below 50 GeV. However, they may begin to be valid within
the 50 to 300 GeV energy range, as evidenced by data obtained at the accelera-
tor facilities at Serpukhov and Batavia (Slattery, 1972).

Figure IX.A-6 shows a solid-line fit to the data given in Figure IX.A-5 in
comparison with the dashed line that shows the product o^q^q, based on
the assumptions of Levy and Goldsmith for proton kinetic energies greater

220 THEORY

p
T 10

_>* ^"JE— -

,--'''

*j*^

\~~~~

•* DATA POINTS

FIT TO DATA POINTS

LEVY AND GOLDSMITH

ASSUMPTIONS

7 *

, , 1

■

i i i i i i i 1 j 1

10 100

T(GeV)

Figure IX.A-6. Comparison of accelerator data from Figure IX.A-5 with the
assumptions made by Levy and Goldsmith (1972).

than 1 GeV. The Levy-Goldsmith assumptions show a reasonable fit to the
data above 100 GeV where the scaling prediction may hold. However, below
100 GeV the dashed curve is, in all cases, above the data points. Figure
IX.A-7 shows the 7r°-production function for pp interactions given by the
product a 0f 0I , based on the data given in Figure IX.A-5. This figure
shows clearly that almost all of the 7T°-mesons produced in cosmic-ray pp
interactions involve cosmic-ray energies between 1 and 10 GeV. Figure
IX.A-8 shows the integral 7-ray production function that is proportional to
the integral of the curve shown in Figure IX.A-3, and is defined such that
q (pp) (<oo) corresponds to the total 7-ray production rate/hydrogen atom/s
from pp interactions alone (the number q H) given in the previous approxi-
mate calculation. It can be seen from Figure IX.A-8 that only 10 percent of
the 7-ray production occurs in interactions involving protons below 1 GeV
and perhaps another 10 percent occurs in interactions above 30 GeV. This
means that (1) because cosmic-ray modulation effects are only important
below 1 GeV, the uncertainty in the true cosmic-ray spectrum due to modu-
lation effects produces only a small uncertainty in the total calculated 7-ray
production rate, and (2) the uncertainty in the exact form of the pion-
multiplicity law f (T) above 30 GeV produces little uncertainty in the total
7-ray production rate. Indeed, 90 percent of the 7-ray s are produced in
interactions below 30 GeV where the form of the multiplicity law used by

DIFFUSE GAMMA-RA Y CONTINUUM RADIA TION

221

T(GeV)

Figure IX.A-7. Differential neutral pion production function from p-p

interactions.

Levy and Goldsmith does not hold. Figure IX.A-8 also shows that the median
proton energy for 7r°-production is ~3 GeV. If we compare the values of
a 0f o used by Levy ^d Goldsmith with those used here (Figure IX.A-6) at
the median it0 -production energy of 3.3 GeV, we obtain a ratio of 2.5 which
just corresponds to the ratio between the values for the total 7-ray production
rate given by Levy and Goldsmith (1972) and the author (Stecker, 1970). Thus,
the discrepancy between the two values is accounted for. The conclusion is
that the Levy-Goldsmith value appears to be too high because it is based on an
asymptotic multiplicity law that does not hold in the energy range where at
least 90 percent of the 7-ray s are produced.

HIGHER ENERGY DECAY PRODUCTS

Mesons and hyperons are also produced in strong inelastic nucleon-nucleon
interactions at somewhat higher energies and their important decay modes

222

THEORY

/.1U

~90% ____

-27

6.10

- 80%/

-27

5.10

-27

2. 4.10

~ 50%/

QtPP)(<T)-8Jlp(T)^(T)v(T)dT

± -27

t ^ 3.10

-27

2.10

- 20%/

-27

1.10

-10%/

— -^ 1

i i

0.1

T(GeV)

10'

10J

Figure IX.A-8. Integral 7-ray production function from the decay of neutral
pions produced in p-p interactions.

leading to 7-ray production are summarized in Table IX.A-2. In addition,
nucleon resonances can be formed which lead to decay chains involving
7T°-mesons in particular. These processes have been discussed in detail by
Stecker (1971a), with particular regard to the 7-ray spectra produced. In
particular, it is found from accelerator measurements that hyperons and baryon
resonances formed in p-p interactions tend to carry off a roughly constant
fraction (~60 percent) of the energy of the incident proton; from this it can
be shown that the resulting 7-ray spectra from the decay of these excited
baryon states maintain the same power-law form as the incident cosmic-rays
at higher energies: if I =K E"rcr, then I N+ Y~K E "rcr. In particular,
if 7T°-mesons are produced by a process leading to a multiplicity law ?aE* ,
and given an average energy <*Eb where b = 1 - a, the resulting 7-ray spectrum
has the form (at high energies)

I (E ) = K E ■

(IX.A-9)

where

ro=-

(rcr+b)-(a + i)

(IX.A-10)

DIFFUSE GAMMA-RA Y CONTINUUM RADIA TION

223

i-^

■3;

i— i

•*-»

so.

■*

pci

IT)

—

00 ^

V£>

.5 ft5

l-H

x! w

l-H

-■H

^—^

;£

•£

2*-

+
+1

+
+

O -H

o <s

O <N

fc=

224 THEORY

and

[r p-(a+l)]/b

2nLK a0x0

K = (IX. A- 1 1 )

(r +b)-(a+l)

For the decay of 2°-hyperons, the spectrum is given by

-r

I oOV=K.oE7 " «XA-12)

2" ' 2

with

"cr I 2 iT*A

K =-— - nLa n I x

M2-M2 \(rcr"1)

zo 2r_ WS0(V m2 / (IX.A-13)

cr \ M^

(M denoting the mass of the particle subscripted) and for nucleon resonances
(isobars)

\&J-*pf* (IX.A-14)

with

2KCTRjnL (r _1}

cr ^fcc-r

Kj= (2Xi) CT l^cr (IX.A-15)

where %, is typically 10"1 - 10"2 (Stecker, 1971a). The relevant data for
hyperons and isobars are given in Tables IX.A-3 to IX.A-5. Table IX.A-6
shows the relevant data for the fireball models of pion production (Stecker,
1971a) and the resultant differential 7-ray spectra at high energies are shown
in Figures IX.A-9 and IX. A- 10. The scaling hypothesis predicts a logarithmic
increase in pion multiplicity with energy, but the resultant form of the 7-ray
spectrum at high energies should be close to the result given in Figure IX. A- 10.

NUCLEON-ANTI NUCLEON ANNIHILATION

Gamma-rays from the decay of 7r°-mesons produced in nucleon-antinucleon
annihilations have spectral characteristics typical of pion-decay 7-rays: a
maximum at m c2/2 ~ 70 MeV and a nearly flat spectrum in the vicinity of

DIFFUSE GAMMA-RA Y CONTINUUM RADIA TION

225

* —

op -a

55 o

z'- "N

o
d

no ft

o >-*

PL, U

E. §

do .y

W3 O

13 <35

C«

£3 in

3 2

7> O

l-H

z 5

E
£73

—

Is"-

3 CO

<L>

1> —

Z _;

+ n

^3 S

^ .2

W O

T3 +3

3 a

S &

l-J «5

/— N

c 2

r-»

■*

226

THEORY

Tt-

1— 1

era

T-H

1 1

i i

<-H

T-t

^■^

, <>

i "1

<<t

N«- '

i— •

%->

1 "°.

i "*.

<-^

ss Sec
(mb)

tJ-

•<t

1 NO

1 no

■*

/*-s

r-~

«fr

""J-

oq in co no

<— i

<— 1

co <-< d d

- £

g2^

o .- — »»

in in

3 § >

oq in

o c <u

(N ^t \D vd

t-^

•— <

•-H

DIFFUSE GAMMA-RA Y CONTINUUM RAD I A TION

227

in
<

o O

" $

■*-< «— i

s z

to —

£> ^c

Q t

c -

o ©

o -

£ <

•^

>» GO

03 s_^

C/5

5 d

"2 °

'So

O -r;

a> > — -

C X)

-a &o

W E

B tfl

-d w

£ o

= •-

3 u

60 ,_

|«2

3 o

-f—

-i—

1-4

0 g
s §

O o

T3 73

f-H

d -S

2 «

O O

CO otf

if-

GO Q

w >

0 >

0 **

Cu 0

£ •«

> .2

"* c

•i §

•2 o

o =

o .d

Energy (
in P-P Cc

■*

M 0

3 3

T3 ■—
O X
- —

0- M

C-;

5 S.

0 j:

s >

0 11

•a 3

s °

^— v

«s 0

as <a

00
0

'"",

S £

** 0

•O 4)

a> "O
E «

3 —

to •*-•

2 <«

S °

'Z C

/— V
00
00

0
+

•3 a-

M °

N— /

O —

b f

-j—

/~s

*-^

228

THEORY

c/a

(/5

-o
o

<L>

cd
H

l-H

°;3
Z U

Z a,

13
O

o
6

,__!

cd
X>

cd
X)
o>

■—

i
o

£ «i

3e>

8,2

td C

V °
cd +3

I 5

S o

O £

-J

DIFFUSE GAMMA-RA Y CONTINUUM RADIA TION

229

C
V

—I r-

N*(1.410)

N#( 1.688)

tt° (ONE-FIREBALL
MODEL)

(TWO-FIREBALL

MODEL)

N*(1.410)

N*(1.688)

LOG Ey (GeV)

10 y

Figure IX.A-9. Calculated 7-ray spectra from various secondary
particles produced in galactic cosmic-ray interactions (Stecker, 1971).

230

THEORY

>
a>
(3
i

o
a>

-26
-28
-30
-32
■34
-36
-38
-40
-42
-44
-46
-48
-50
-52
-54

-56

1 1

III!

i i

r=2.9

\ <r>=3.o

r=3.2

1 1 1

1 1

■ i i i

r=2.6 \

i ■

• N

-2 -1 0 1 2 3 4 5 6 7 8 9 10
LOG Ev(GeV)

io y

Figure IX.A-10. Total calculated galactic 7-ray production spectrum from
cosmic-ray interactions (Stecker, 1971).

the maximum which is symmetric on a log E plot about the point m^c /2.
However, if the annihilations are assumed to occur near rest in the laboratory
system (that is, in the universe) the spectrum is bounded between a maximum
7-ray energy of ~919 MeV and a minimum energy of about 5 MeV. This is
because the maximum energy given to a 7r°-meson occurs in the three particle
annihilation

p + p

7T+ + n +1T°

(IX.A-16)

DIFFUSE GAMMA-RA Y CONTINUUM RADIA TION

231

and is 923 MeV. (Two-particle annihilations involving 7r°-mesons being for-
bidden by selection rules involving conservation of G-parity (Stecker, 1971a).)

Frye and Smith (1966), using accelerator data, and independently Stecker
(1967, 1971a), using a theoretical pion-production model in p-p annihilation,
have calculated the resultant 7-ray spectrum from p-p annihilation at rest.
There is excellent agreement between the two calculations, and the resultant
spectrum, on a logarithmic energy plot, is shown in Figure IX.A-1 1 .

10-210-s

j 1 '''''

j 1 1 1 i 1 1 1

-2

-1

Ey(GeV)

Figure IX.A-1 1. Normalized local differential 7-ray spectrum from p-p

annihilation at rest.

THE COSMO LOGICAL GAMMA-RAY BACKGROUND

We now turn to a discussion of the isotropic 7-ray background spectrum
which is expected to be of cosmological origin. Figure IX.A-1 2 shows schema-
tically the results of recent observations of this background spectrum by

232

THEORY

Trombka et al. (1973), Mayer-Hasselwander et al. (1972), Share et al. (1973),
and Kraushaar et al. (1972) (see also Chapters III.A, and IV.A, B, C).

10'

3 -

2 _

10 -

^ 1-

* 10"11-

E
u

> io-2

-3 _

-4 _

5 _

1 1 1 1 1

\ OSO-3

\ranger-apollo 15

\ APOLLO 15

APOLLO 15\

MPI \

NRL \

\SAS2

\0S0-3

1 1 1 1 \

10"

10_1 1

Ey(MeV)

10 10'

Figure IX.A-12. Recent observational results on the cosmic
7-ray background spectrum.

Results from the OSO-3 detector in the 10- to 100-keV energy range
have shown that the background radiation in this range is isotropic to better
than five percent over angular scales of 10 degrees (Schwartz, 1970). In the
energy range between 0.2 and 4 MeV, Damle et al. (1972) have found evi-
dence for the isotropy of the diffuse background flux. Above 50 MeV, the
results from SAS-2 and OSO-3 (Share and Kniffen et al, Chapters IV.A and B)
indicate that there is a relatively hard component of 7-radiation of galactic
origin, and a true diffuse extragalactic background component observed at

DIFFUSE GAMMA-RA Y CONTINUUM RADIA TION 233

high galactic latitudes that is soft ('vE-3 ) and that connects smoothly with the
Apollo data below 30 MeV (Figures IX.A-12, DC.A-13 and IX.A-14). The
evidence would thus seem compelling that the spectrum represented in Figure
IX.A-12 is of extragalactic origin and is therefore not consistent with the
galactic-origin hypothesis suggested by Cowsik in Chapter VIII.A.

Because of the cosmological aspects relating to studies of the diffuse isotropic
7-ray background, it is necessary to discuss the physics of 7-ray production in
past epochs; such radiation may be reaching us today from distances of the
order of ~15 billion light years. According to big-bang cosmology, the uni-
verse was in a smaller, denser state in the distant past and has been continually
expanding. This general expansion has caused all electromagnetic radiation
to be Doppler shifted to the red (that is, to longer wavelengths which implies
lower energies). The red shift is usually designated by z = a\/A.

This red shift implies that a spectrum of 7-rays, for example, from 7r°-decay
(either from annihilation or cosmic-ray interactions), that has a maximum at
~70 MeV locally, would have that maximum shifted to a lower energy if such
radiation were produced at an epoch corresponding to a significant red shift.
To find the total spectrum expected to be observed, we must integrate over
all red shifts where 7-rays were being produced and weigh the integration with
various factors of the quantity (1+z) (for a complete discussion, see Stecker,
1971a, Chapters 9 to 14). Also, for z > 100, Compton interactions between
7-rays and intergalactic gas may result in energy loss for the 7-rays so that, in
general, an integrodifferential transport equation involving E and z must be
solved in order to obtain the expected total 7-ray spectrum resulting from
high red shift processes such as matter-antimatter annihilation (Stecker et al.,
1971). Absorption processes such as pair-production mechanisms involving
intergalactic gas and 2.7 K blackbody photons eliminate 7-rays with large red
shifts from various parts of the observed spectrum. Gamma rays arising from
any pion-decay process at cosmological distances contribute significantly to
the isotropic background only above 1 MeV, because 7-rays at lower energies
have been red-shifted by a factor of < 70. Such a red shift corresponds to an
epoch when the universe was opaque to 7-rays and absorption effects were
important. The basic equation to be solved, the cosmological photon transport
(CPT) equation, is of the form

/*£(E)

— + — [-E H (z)jr] = 3 (E, z) - k(z)ab + dE' k(z)sc (E|E V(E')

3t 3E

(IX. A- 17)

234

THEORY

10'

-2

01
UJ

h-
o

UJ
CO

I I I I ""I I I I I I IM| 1 1 I I I lll| 1 1 I I INI

TOTAL THEORETICAL

'ZuAV-100

APOLLO 15 DATA

-4

z
o

X
Q_

-5

► GOLENETSKI etal (1971
roVEDRENNE et al (1971)
Z • MAYER-HASSELWANDER
I etal, (1972)
- APOLLO I5=TR0MBKA
et al,(l973)

SAS-n = THIS WORK

-6

SAS-I

E DIFFUSE /-RAY SPECTRUM

0.1

10
Er (MeV)

10'

10*

Figure IX.A-13. Comparison of the observed background with a two-com-
ponent model involving the production and decay of neutral pions produced
in intergalactic cosmic-ray interactions at red shifts up to 100. (See also
Kniffen et al., Chapter IV. C).

DIFFUSE GAMMA-RA Y CONTINUUM RADIA TION

235

I I I Mil) 1 1 I I I III) I 1 I I I lll| 1 1 I I I II

> icr3

UJ
CO

CVJ

z
o

I-
o

\0~

10"

APOLLO 15

TOTAL THEORETICAL

► GOLENETSKI et al (1971)
kroVEDRENNE0t ol (1971)

• M AYER-HASSELWANDER

et al,(l972)
APOLLO I5'TR0MBKA

el al,(l973)
SAS-I'THIS WORK

DIFFUSE /-RAY SPECTRU

SAS-E

■ ' ' ' ■ u

I I 1 I III

_l ■ ■ ' ■ ■■

0.1

10'
Ey (MeV)

10'

10*

Figure IX.A-14. Comparison of the observed background spectrum with a
two-component model involving the matter-antimatter hypothesis as dis-
cussed in the text.

236 THEORY

where E is the photon energy, and kab and ksc are the photon absorption
and scattering rates (which are a function of z because the intergalactic gas
density is assumed to scale as (1+z)3 because of the expansion of the universe).
The script quantities for the 7-ray intensity and production rate

^(E,z)=(l + zV3I(E,z)
and (IX.A-18)

^(E,z) = (l+z)-3Q(E,z)

are quantities comoving with the expansion, defined so that their red-shift-
density dependence cancels out. £(E) is an upper limit on the scattering
integral defined by the Compton process and H(z) is the Hubble parameter
which, in terms of the Hubble constant HQ , is given by the relation

H(z) = HQ(l+z) [1 + £l]Vz (IX.A-19)

where 12 is the ratio of the mean-gas density in the universe to the density
needed to close the universe gravitationally. The term

%j by

— =-(l+z)H(z)— (IX.A-20)

at oz

and the second term in Equation (IX.A-17) expresses the energy loss of the
7-ray s because of the expansion red shift.

COSMO LOGICAL SPECTRUM FROM MATTER-ANTIMATTER
ANNIHILATION

Equation (IX.A-17) can be solved for cosmological models involving the
annihilation of nucleons and antinucleons in a bary on-symmetric universe
(Stecker et al., 1971; Stecker and Puget, 1972;Puget, 1972; Chapter XV. A).

Between ~5 and ~50 MeV, Equation (IX. A- 1 7) reduces to a power-law form
I (E) ex E'r ANN (Figures IX.A-15 and IX.A-16) with the value for TANN
estimated by Stecker and Puget (1972) to be ~2.5 <rANN < ~3.5.

ABSORPTION EFFECTS-THE GAMMA-RAY WINDOW

In the vicinity of ~1 MeV and below, absorption effects due to Compton
scattering become important and cause the spectrum to bend over as shown
in Figures IX.A-15 and IX.A-16. Figure IX.A-17 shows the critical red shift
for absorption of 7-radiation as a function of observed energy. At lower
energies, absorption is due to Compton interactions with intergalactic matter;

DIFFUSE GAMMA-RA Y CONTINUUM RAD LA TION

237

on
O

>-
U£
«C

I—

100

Ey (MeV)

Figure IX.A-15. The cosmological 7-ray spectrum from matter-
antimatter annihilation calculated by solving the CPT equation numeri-
cally for O = 1. The solid line represents the complete solution. The
other curves represent the effect of neglecting the absorption and
scattering (transport) terms in Equation (IX.A-17).

in the intermediate range absorption is due to pair-production interactions with
intergalactic matter (Arons and McCray, 1969; Rees, 1969). At the higher
energies absorption is due to pair -production interactions with blackbody
photons (Fazio and Stecker, 1970). There is a natural "window" between
M MeV and MO GeV which is optimal for studying cosmological 7-ray s.
Absorption effects come in below 1 MeV and above 10 GeV.

COSMOLOGICAL SPECTRUM FROM COSMIC-RAY PION DECAY

Figure IX.A-13 shows a two-component model normalized for a best fit to
the observations involving the production of intergalactic 7-rays from cosmic-

238

THEORY

-1

-2

10"3

T71 r~r
\

7 ^

n0=10 cm ° *

nn=10"5cm 3

J 1 1 1 1 1 1 1 1

-1

Er(MeV)

1 1 1 1 1

Figure IX.A-16. The effect of absorption of 7-rays at high red
shifts by the protogalactic gas.

ray interactions with intergalactic gas producing 7r°-mesons out to a maximum
red shift of 100 (Stecker, 1969b, c, 1971b). The cosmic rays may be produced
in protogalactic sources (protars). Three problems arise with this explanation:
(1) even with a relatively steep assumed cosmic-ray spectrum (yE~2J) the
bulge in the theoretical spectrum may be too large to fit the observations,
although this discrepancy may not be too serious considering observational
uncertainties; (2) large amounts of energy are needed in cosmic-rays at high

DIFFUSE GAMMA-RA Y CONTINUUM RADIA TION

239

i r

'iimiii iiiiiiiii iniiiii

n0=lCr7cm"3

y+p — e+ + e" + P

n0 = 10"5cm"3

WINDOW

y+y — e++e"

»-4 ia-3 lr,-2 Ir>-1

10" 10"3 10"' 10"1 1 10 10' 10° 1(T 10° 10° 10' 10

Ey(MeV)

Figure IX.A-17. The critical red shift for absorption of 7-radiation as a
function of observed 7-ray energy.

red shifts;* and (3) the maximum red shift for cosmic-ray production (ZMAX)
is a free parameter chosen to fit the observations. The matter-antimatter
annihilation hypothesis does not suffer from the above mentioned problems.
The parameter ZMAX does not enter into the theory; annihilations occur at
all red shifts and the 1-MeV flattening is an absorption effect as discussed
earlier. The transport Equation (IX.A-17) was solved to determine the exact
form of the spectrum. Energy considerations do not present a problem.
Another advantage of the theory is that it arises as a natural effect in a
cosmology such as that suggested by Omnes (Omnes, Schatzman, and Puget,
Chapters XIV.A, B, and XV.A).

*In a recent private discussion between the author and P. Morrison, it became apparent
that the energy problems may not be too great with this (protar) hypothesis if, indeed,
spinars existed at such red shifts of about 70 to 100 (Stecker, 1971b). If it is considered
that each spinar produces approximately 1062 ergs over a time scale of 107 to 108 years
(Morrison, 1969), a time comparable to the Hubble time at these red shifts, then at
most 20 percent of the presently observed galaxies are needed to have arisen from this
early spinar state in order to provide the cosmic-ray energy needed to account for the
diffuse 7-radiation above 1 MeV. At a red shift of about 70, the free-fall time for
forming spinars from gas clouds is comparable to the Hubble time. This may provide a
natural upper limit to the red shift (zMAX), for primordial cosmic-ray production in
the spinai model. (It should, however, be noted that such spinars may arise in other
ways (Stecker, 1971b) and that they may now be a class of moribund objects unrelated
to galaxies as we see them now.)

240 THEORY

COSMO LOGICAL ANNIHILATION SPECTRUM COMPARED WITH
OBSERVATIONS

Figure IX.A-14 shows a detailed comparison of the annihilation -hypothesis
spectrum with present observations assuming TANN = 2.5 (see discussion of
Stecker and Puget, 1972). The two-component model shown presents an
excellent fit to the observational data.

COSMOLOGICAL COMPTON MODEL

Several other models of isotropic 7-ray production have been put forward
recently. One suggestion is that the whole spectrum in the 10"3- to 102-MeV
range is due to Compton interactions of intergalactic electrons with the
universal blackbody radiation (Felten, 1965; Gould, 1965; Hoyle, 1965;
Fazio, Stecker, and Wright, 1966; Felten and Morrison, 1966). In its most
recent version, Brecher and Morrison (1969) have attempted to explain the
observed spectral features, namely, the steepening in the spectrum at M-0 keV
and flattening above 1 MeV, using the Compton hypothesis. The Brecher-
Morrison spectrum is shown in Figure IX.A-18, superimposed on the data-curve
of Figure IX.A-12. The fit is reasonable except at the extreme high- and low-
energy ends of the energy range. However, Cowsik and Kobetich (1972) have
recently recalculated the Brecher-Morrison spectrum using a true blackbody
target spectrum and a more realistic energy distribution for Compton-scattered
photons (rather than the 5 -function approximations used by Brecher and
Morrison). The result is a smearing out of the spectral features of the Brecher-
Morrison model into a smooth power-law spectrum. Other problems with the
Brecher-Morrison model stem from the fact that in order to get a large enough
flux generated, electrons are required to leak out of normal galaxies in a time
much shorter than the M07y deduced for protons on the basis of cosmic-ray
isotropy measurements.

COSMOLOGICAL BREMSSTRAHLUNG MODEL

Another hypothesis that attempts to account for the whole photon spectrum
is the electron-bremsstrahlung hypothesis. Figure IX.A-19, which compares
the spectrum generated by this process with the observations, shows an excel-
lent fit with the theoretical spectrum based on calculations by Arons, McCray,
and Silk (1971) below 1 MeV, and Stecker and Morgan (1972) above 1 MeV.
The break at ~3.5 MeV is due to energy loss by cosmic-ray electrons inter-
acting with the 2.7 K blackbody radiation. Unfortunately, we again have severe
energetic problems with this process, bremsstrahlung being an inherently
inefficient 7-ray generating mechanism. Another problem lies in getting
galaxies to leak low-energy nonrelativistic electrons at a fast enough rate.
Assuming this could be done, an electron spectrum would be distorted by
heating the intergalactic medium to 109 K. The problems with this mechanism

DIFFUSE GAMMA-RA Y CONTINUUM RADIA TION

241

IO4
103

l\\ 1 1 1

102

- \

? 1

•
CO

" 10"1

E
o

V-

OBSERVED "\\

-i-M

—

x 10'2

10"3

io-4

IO"5

i r»-6

1 1 1 1

V -

10"

io-

10"

Ey(MeV)

10'

Figure IX.A-18. Comparison of the observed background spectrum
with the Brecher-Morrison model.

have been discussed by Setti and Rees (1970), Prilutskii and Rozental
(1971), and Cowsik and Pal (1971). (See also Cowsik, Chapter VIII.A.)

RELATIVISTIC THERMAL SOURCES

One additional mechanism for producing a second component of 7-radia-
tion was suggested by Sunyaev (1970), namely, thermal bremsstrahlung from
relativistic electrons in a 20-MeV plasma in such objects as the nuclei of

242

THEORY

10'

10;

10'

„ 10-1

* 10"2

10"

BREMS

10"

Er(MeV)

Figure IX.A-19. Comparison of the observed background spectrum
with the electron bremsstrahlung model as discussed by Arons et al.
(1971) and Stecker and Morgan (1972) with a spectral break at
EB = 3.5 MeV as discussed by Stecker and Morgan.

Seyfert galaxies. This, of course, immediately presents the problem of having
enough Seyfert galaxies to account for the observed flux. However, a much
more serious problem with the fundamental physics of the mechanism has
been pointed out by Prilutskii et al. (1971). They notice that in order to
contain the hot relativistic plasma, a magnetic field is required of a strength

DIFFUSE GAMMA-RA Y CONTINUUM RADIA TION 243

such that

H2
— >nekT (IX.A-21)

87T

In that case, the ratio (R) of the electron energy loss rate from synchrotron
radiation to that from bremsstrahlung is of the order of

aTc(H2/87r)(kTe/mec2)2

R« >«^kT/mc2) >>l (IX.A-22)

aaTc nekTe

for a relativistic plasma where kT > mec2 and a"1 * 137 (aT is the Thomson
cross section). In fact, for a plasma of temperature TM v given in MeV,

R~500T2MeV (IX.A-23)

Thus, in an optically thin plasma, the synchrotron loss rate is the dominant
loss term in the energy-equilibrium equation determining the equilibrium
electron spectrum. This will ensure that the electrons have a nonthermal
spectrum and produce nonthermal radiation. In addition, the 20-MeV cutoff
in the electron spectrum suggested by Sunyaev (1970) will not exist. The
details of the argument are further described by Prilutskii et al. (1971).

SUMMARY- INTERPRETATION OF PRESENT OBSERVATIONS

It is the opinion of the author, based on the previous discussion, that the
most promising theoretical interpretation of the unexpected increase in the
observed background flux of 7-radiation above 1 MeV, at present, is that this
radiation has arisen from the annihilation of nucleons and antinucleons,
primarily at high red shifts, on the boundaries between regions of matter and
antimatter (Stecker et al., 1971 ; Stecker and Puget, 1972; Omnes, Schatzman,
Puget, Chapters XIV.A and B and XV.A). This conclusion is, of course,
conditional upon future observations and theoretical investigations.

The arguments presented in the paper of Steigman (Chapter XIV.C) put
restrictions on bary on-symmetric cosmologies, but nonetheless are not in
conflict with the particular cosmological model discussed here and in the
papers of Omnes, Schatzman, and Puget (Chapters XIV.A, B, and XV.A).

Tables IX.A-7 and IX.A-8 summarize some of the significant aspects and
spectral attributes of the various mechanisms important for the production
of the diffuse cosmic 7-radiation. The last column in Table IX.A-7 lists the
cosmic domains where the various mechanisms probably play an important

244

'■*->

" 1

o o-

'^3 >>

3 *>

00 »

'oo

2 g

2 S

o °

j>>

a, rt

■si

'lo
oo
O

—

Cd ""•

o 3

42 -°

3 ^

CO CJ

go,

oo -3

•<->

2 1

■4—* cd

r3 3

C<3

w £

•*->

+-»

>»

H— »

h->

-1— »

l/l

T3
0)

c
o

>
^3

C 2

C/3

° s

a 2
43 x

H->

a> a,
5 ho

3
e1 oo

Cd cd

o-
cd

4 g

T oo

•~-

.a o

H
cd

B 1

ft «

U o

1
oo

'■*->

•t-J

H->

i—

+->

ccj

i-i

■6

h-j

—3 OO

•a c
2 o

00 "£J

C
O

"3 /— N

o <^

O C

tp o

c
o

'Eh

o .2

cq .a

J2 2
3 a)

■*->

£.s

.2 S

H ^

DIFFUSE GAMMA-RA Y CONTINUUM RADIA TION

245

^ cu

cd
.3
cu

co
O

( cu

'5b

cu £
rT co

<H— 1

3
o

'co

■3 S

3 «

cu 0

cu 3
X) c

00 «i
cu O

cu
3
<u

Spectrum power-law bet
~ 5 and +* 50 MeV, turn

CkH

c/3

13
cu

'ob

"3
E

Power-law roughl
I(E )~E"2
Relation between

co
■M

X
CD

CJ
t-

o
Eh

o
O

-t—
c

S3
O
EX

<+-
o

cu
in

-3
cu

E
■|

■a

-o
a

o
in

>
o

—

iximum
a log E^
roughly
w

'C
cu

cd
j3

3
i—

a,
oo
cu

er-law roughly

)~e;2

tion between

/— s
co *— '

3 +
o u

3
O

_CJ

« >-w

W *

53
O
3

ID
Eh

ly flat near mi
symmetric on

h. Power-law
E"3 above a fe
[red MeV.

3
cu

3
O

o
cu

5 BS

O
On

2 cd "t-j
cd cu ^

S Z cd

5^ ) ?
cd I 3

kH 3

00^ ^

°*,<«-.

cd
Of

r-v

E °

T3 cd

cosmic-ray
high photc

CU
'(JO

<u
3

co
«B

'ob
cu

CU
'-t->

4-1

'Si

3
3
O

cd
o

lH
CU

3
O

'E,

"ob

J3 -S

etric aroun
ph (for dec

"ob

at typical
ergies
:ed toward

•Si

i
■—

*->

cd
1

V
Id

-a

o
c

T3-

cd cd
I-®

w >>

cd
00 o

£S3„

E ^^3
"3 oo --

'ob

3
cd

O <u

— -1 "3

-3 c -a
.3 <" s

cd
>>

■a

<L>

Cd fc.

3 c°
O w

+j cd

r2 53

"ca

Un 04

tlC

i-U

tLi O

a
o

'-J3

V- >

O kH

3 cu

■M

*o ii

k-
CU

O cd

Oh .§

3
O

CO ^-v

O co
O 3

53 C

2 «!

i_i

'cl

<J .2

OH JH

-C

cu
cu

CO
CO

1—

"cd

1-4

to O

jd w
"a3 <u

13 S
is cs
3 E

cu ~

Z .3

246 THEORY

role. The results from OSO-3 and SAS-2, as summarized in these proceedings
in the papers of Share and Kniffen et al. (Chapters IV.A and C), indicate that in
the energy range above 50 MeV, there is a distinct hard component of galactic
origin and a much softer, high galactic latitude component of extragalactic
origin. The galactic component appears to be predominantly (that is, greater
than 50 percent) of 7r°-decay origin and therefore is small relative to the steep
extragalactic component much below 50 MeV. The extragalactic component
fits onto the Apollo data (see Peterson and Trombka, Chapter III.A) below
30 MeV, so that all indications are that the flux below 30 MeV is over-
whelmingly extragalactic. Because the galactic flux is much harder above
100 MeV than the extragalactic flux, the galaxy stands out well above the
extragalactic background at these energies. However, below 30 MeV, the
galaxy becomes relatively dim and blends into the background as only a
small perturbation. These conclusions are contrary to the galactic origin
hypothesis for 0.2- to 10-MeV 7-radiation discussed by Cowsik in Chapter
VIII.A, but at present appear to be more consistent with recent satellite
observations as presented at this conference.

If the galactic disk component of 7-radiation is primarily of 7r°-decay origin,
I will stand by my previous arguments (Stecker, 1969a; Stecher and Stecker,
1970; Stecker, 1971a, Chapter 8) that the OSO-3 measurements of
Kraushaar et al. (1972) and those obtained by SAS-2 (Kniffen et al.,
Chapter IV. C) indicate that there may be a substantial amount of molecular
hydrogen in the galaxy. This is implied by my recent calculations of the
7-ray production rate (Stecker, 1973) which confirm my earlier calculations
of 1 .3 ± 0.2 X 10"25 s"1 (Table IX.A-1), but are now on a much more solid
basis. Forthcoming results from the SAS-2 and Copernicus satellites should
settle the question in the near future.

In the galactic center region, the flux should be somewhat softer than in the
disk as a whole because of a significant component from Compton interac-
tions (Stecher and Stecker, 1970; Stecker, 1971a, Chapter 8). Preliminary
observational results suggest that this is the case (Share, and Kniffen et al.,
Chapters IV. A and C). Again, here we await the final results from SAS-2.

The note of anticipation is appropriate here because it seems that at the time
of this first international 7-ray astrophysics Symposium, we are on the
threshold of a new era of observational 7-ray astronomy.

REFERENCES

Arons, J., and R. McCray, 1969, Astrophys. J. Letters, 158, p. L91.
Arons, J., R. McCray, and J. Silk, 1971 , Astrophys. J. , 170, p. 431 .
Brecher, K., and P. Morrison, 1969, Phys. Rev. Letters, 23, p. 802.

DIFFUSE GAMMA-RA Y CONTINUUM RADIA TION 24 7

Cavallo, G., and R. J. Gould, 1971,77 Nuovo Cimento, 2 B,p. 77.

Clark, G. W., G. P. Garmire, and W. L. Kraushaar, 1968, Astrophys. J.
Letters, 153, p. L203.

Comstock, G. M., K. C. Hsieh, and J. A. Simpson, 1972, Astrophys. J. , 173,
p. 691.

Cowsik, R., and E. J. Kobetich, \91 2, Astrophys. J. , 177, p. 585.

Cowsik, R., and Y. Pal, 197 1 , Cosmic Ray Astrophys. , Tata Institute Press,
Bombay.

Damle, S. V., R.R Daniel, G.Joseph, and P.J.Lavakare, 1912, Nature, 235, p. 319.

Dilworth, C, L. Maraschi, and G. C. Perola, 1968, II Nuovo Cimento, 56 B,
p. 334.

Fazio, G. G., and F. W. Stecker, 1970, Nature, 226, p. 135.

Fazio, G. G., F. W. Stecker, and J. P. Wright, 1966, Astrophys. J. , 144, p. 61 1.

Felten, J. E., 1965, Phys. Rev. Letters, 15, p. L1003

Felten, J. E.,andP. Morrison, 1966, Astrophys. J., 146, p. 686.

Fichtel, C. E., R. C. Hartman, D. A. Kniffen, and M. Sommer, 1972,
Astrophys. J. , 171 , p. 3 1 .

Frye, G. M., and L. H. Smith, 1966, Phys. Rev. Letters, 17, p. L733.

Ginzburg, V. L., and S. I. Syrovatskii, 1964, Origin of Cosmic Rays,
New York, Macmillan.

Gould, R. J., 1965, Phys. Rev. Letters, 12, p. L511.

Hayakawa, S., H. Okuda, Y. Tanaka, and Y. Yamamoto, 1964, Prog. Theo.
Phys. (Japan) Suppl. 30, p. 153.

Heitler, W., 1954, Quantum Theory of Radiation , Oxford Press, London.

Hoyle, F., 1965, Phys. Rev. Letters, 15, p. L131.

Jones, F. C, 1965, Phys. Rev., 137, p. B1306.

Kraushaar, W. L., G. W. Clark, G. P. Garmire, R. Borken, P. Higbie, C. Leong,
and T. Thorsos, 1912, Astrophys. J. , 177, p. 341 .

Levy, D. J., and D. W. Goldsmith, 1912, Astrophys. J. , 111, p. 643.

Mayer-Hasselwander, H. A., E. Pfefferman, K. Pinkau, H. Rothermel, and
M. Sommer, 1972, Astrophys. J. Letters, 175, p. L23.

Morrison, P., 1958, II Nuovo Cimento, 1, p. 858.

248 THEORY

Morrison, P., 1969, Astrophys. J. Letters, 157, p. L75.

Pollack, J. B., and G. G. Fazio, \963, Phys. Rev. , 131, p. 2684.

Prilutskii, O. F., Y. P. Ochelkov, I. L. Rozental, and I. B. Shukalov, 1971,
Izvestia Akademii Nauk SSSR, 35, p. 2453.

Prilutskii, O. F., and I. L. Rozental, 1971, Aston. Zh, 48, p. 489.

Rees, M. J., 1969, Astrophys. J. Letters, 4, pp. L61 and LI 13.

Schwartz, D., 1970, Astrophys. J. , 162, p. 439.

Setti, G., and M. J. Rees, 1970, Non Solar X- and y-Ray Astronomy, I A U
Symp. No. 37, (Rome), L. Gratton, ed. D. Reidel, Dordrecht, Holland,

p. 352.

Share, G. H., R. L. Kinzer, and N. Seeman, \973, X-Ray and Gamma-Ray
Astronomy, Proc. of IAU Symp. No. 55, (Madrid), H. Bradt and
R. Giacconi, eds. D. Reidel, Dordrecht, Holland.

Slattery, P. \972, Phys. Rev. Letters, 29, p. LI 624

Sreekentan, B. V., 1972, Space Sci Rev. , 14, p. 103.

Stecher, T. P., and F. W. Stecker, \910,Nature, 226, p. 1234.

Stecker, F. W., 1967, Smithsonian Astrophys. Obs. Spec. Report. ,261.

, 19 69a, Nature, 222, p. 865.

, 19 69b, Astrophys. J., 157, p. 507.

, 1969c, Nature, 224, p. 870.

, 1970, Astrophys. and Space Sci. , 6, p. 377.

, 1971a, Cosmic Gamma Rays, NASA SP-249,

(Washington, D.C., U.S. Gov't. Print. Off.).

, \97lb, Nature, 229, p. 105.

, 197 3, Astrophys. J., 185, p. 499.

Stecker, F. W., and D. L. Morgan, \97 2, Astrophys. J. , 171, p. 201 .

Stecker, F. W., D. L. Morgan, and J. Bredekamp, 1971 ,Phys. Rev. Letters,
27, p. L1469.

Stecker, F. W., and J. L. Puget, 1972, Astrophys. J., 178, p. 57.

Sunyaev, R. A., 1970, Pis'ma Zh. Eksp. Teor. Fiz. 12, p. 381.

Trombka, J. I., A. E. Metzger, J. R. Arnold, J. L. Matteson, R. C. Reedy, and
L. E. Peterson, 1973, Astrophys. J. , 181, p. 737.

Chapter X

A. GAMMA-RAY ASTRONOMY AND
COSMIC-RAY ORIGIN THEORY

V. L. Ginzburg*

P. N. Lebedev Physical Institute

The science of 7-ray astronomy will yield entirely new information that cannot
be obtained by optical, radio, or X-ray astronomy and which will be important
for the entire study of high-energy astrophysics, including the astrophysics of
cosmic rays and the problem of their origin. Indeed, only 7-ray astronomy
allows us to study the nuclear component of cosmic rays far from the earth.
(We will refer to the nuclear component here as "cosmic rays" and the electron-
positron component as "relativistic electrons.")

Before the present 7-ray observations, we had only indirect knowledge about
the cosmic rays far from the earth; this knowledge was obtained mainly by
radio observations. The radioastronomical data, as is well known, enable us to
obtain the form of the relativistic electron spectrum, but the spectrum itself
and the corresponding energy density of the electrons (we,) can be deduced
only by making an additional assumption about the strength of the magnetic
field (H) in the radiating region. To estimate the energy density of the cosmic
rays (w ), we have also to assume a relation between wct and we . In fact, it
is usually assumed that they are proportional, that is,

Wcr=KrWe=KH'1(H2/87r)V ^A^

Here w = w V, W = w V, and (H2/87r)V are respectively the energy of the
cosmic rays, the relativistic electrons, and the magnetic field in the source of
volume V, Kf = (wcr/we) and kh = (H2/87iwcr).

Thus, from radio astronomy observations (and also knowing the distance to
the source (R), we can determine the quantities wct, we, and H only by fixing
the values of Krand kh . Near the earth, Kf ~ 100, and in quasi-equilibrium

*Presented in absentia.

249

250 THEORY

conditions, probably k ~ 1 . These values are usually assumed, but in doing
this two far-reaching assumptions are made. In nonstationary sources of cosmic
rays, it is entirely possible that «H < 1 or even « 1 . Close to strong sources
of infrared and optical radiation it may turn out that Kf » 100 because
electrons undergo rapid energy loss. It is possible that in some cases, if mainly
electrons are accelerated, k « 100.

' r

It is, in principle, possible to use radio and X-ray data together to determine the
magnetic field strength (H), itself (or the quantity n kh), if the radio emission
mechanism is synchrotron radiation and the X-radiation is produced by inverse
Compton scattering of the same relativistic electrons in a known radiation
field. But, here too, we cannot find the energy of the cosmic rays w directly
without assuming the values of k or k„ .

A vital question has not yet been answered concerning the energy density of
cosmic rays w M = wM in the metagalaxy (or the metagalactic region close
to the galaxy). Metagalactic models for the origin of cosmic rays are still being
discussed (Setti and Woltjer, 1971 ; Burbidge and Brecher, 1971 ; Shklovskii,
1971) and are sometimes even considered preferable to galactic models for the
origin of cosmic rays. In the metagalactic models, wM «* wG and wG = wcr G~
10"12 erg/cm3, is the energy density of cosmic rays at the earth and, we may
also assume, in a considerable part of the galaxy. I have previously given my
views on the origin of cosmic rays on many occasions (Ginzburg, 1970, 1971 ;
Ginzburg and Syrovatskii, 1964, 1967, 1971). I feel that the metagalactic
models are much less likely than the galactic models of the origin of cosmic rays.
The main arguments rely on energy considerations and are also connected with
7-ray observations. However, these and other arguments are not yet conclusive,
especially in regard to local metagalactic models in which wM « wG only in a
restricted region in the vicinity of the galaxy.

Since we assume fewer relativistic electrons in the metagalaxy than at the
earth (Ginzburg, 1970), in the metagalactic models far from the galaxy,
k » 100. It is also hard to doubt that in intergalactic space kh « 1 since
for «H ~ 1 , HM ~ 5 X 10"6 oe. Therefore, we cannot rely on radio and X-ray
data to determine the cosmic-ray intensity in remote regions of the galaxy
and in radiogalaxies and determine the validity of the metagalactic models; it
is necessary to find a new, independent method. Such a method is provided by
7-ray astronomy (see for instance Ginzburg and Syrovatskii, 1964, 1965; Clark,
Garmire, and Kraushaar, 1968, 1970; Fazio, 1968; Stecker, 1971 ; Cavallo and
Gould, 1971 ; Fichtel et al., 1972).

Protons and nuclei in cosmic rays collide with protons and nuclei of inter-
galactic and interstellar gas. As the result of these collisions, various particles
are produced. Of particular importance here are the secondary it0- mesons and
2°-hyperons which quickly decay to produce 7-rays. The probabilities and

COSMIGRA Y ORIGIN THEOR Y

251

kinematics of all the essential reactions are fairly well known (Stecker, 1971 ;
Cavallo and Gould, 1971) and enable us to calculate the spectrum of 7-rays
with an accuracy which is entirely sufficient from the point of view of
cosmic-ray-origin theory (see Stecker, Chapter IX.A). The integral flux of
7-rays from a discrete source is given by the expression

7(>E7) dft

5 X 1023 (al )M photons/cm2 s (X.A-2)

where SI is the solid angle subtended by the source, R is the distance to the
source (cm) and M is the mass of gas in the source in grams. The chemical
abundances in the source are assumed to be the same as the common
abundances of the elements (especially in the case of He) and thus the average
mass of a gas nucleus is taken to be 2 X 10"24g. The value for (aIG)E7=100 MeV
is taken from Figure X.A-1 to be 10"26 s"1 Sr"1 as given by Stecker (1971).
Therefore

5X lO-3 M(w fwn)

F(>E)= 2 CT G photons/cm2s (X.A-3)

-26
10

-27

. . 1

1 1 1 1 1 1 1 1 1 1

1 1 1 1 1 1 1 \

10
Ex(MeV)

Figure X.A-1. Integral 7-ray production rate (from Stecker, 1971

252 THEORY

where w is the cosmic-ray energy density in the source, assuming that the
form of their spectrum is the same as that observed near the earth. Within the
limits of this approximation, for sources like the galaxy where nonionized atomic
hydrogen predominates, M =1 .2 MHI, where MHI is the mass of neutral atomic
hydrogen.

The spectrum of 7-rays from 7r°-decay is concentrated mainly in the energy
range above 50 to 1 00 MeV (where the 7-rays do not originate in highly red-
shifted sources). (See Figure X.A-1 and Stecker, Chapter IX.A.) For 7-rays
from pion decay, we find

F (E >50MeV)-F (E > 100 MeV)

|=_i1j: 1 1—1 1 = 0.12 (X.A-4)

F (E > 100 MeV

7V 7

In the case of bremsstrahlung radiation from relativistic electrons with the
spectrum I (E) = KE"2-6, £ = 2.03, and for the case of synchrotron radiation
or inverse Compton scattering from relativistic electrons £ = 0.74. Thus,
spectral measurements of the 7-ray flux allow us, in principle, to distinguish
between the various production processes and establish the "nuclear" nature
of the 7-radiation. Once this is done, measurements of the flux allow us to
determine the quantity w /wG in the source. Here we have assumed that
the cosmic-ray spectrum in the source is similar to the spectrum observed
near the earth. This determination, even by the method given above, would
represent an important step forward and, I feel, would be a very important
achievement for high-energy astrophysics.

I wish to illustrate my remarks with two examples of the potential for 7-ray
observations of specific astronomical objects, viz., the Magellanic clouds and
the galactic center. Observations of the Magellanic clouds provide a potential
test for the local metagalactic origin model as well as other metagalactic models
of the origin of cosmic rays. If wM « wG ~ 10"12 erg/cm3 the metagalactic
models can be discarded (Ginzburg, 1972). The Large Magellanic Cloud (LMC)
and the Small Magellanic Cloud (SMC) distances and neutral-hydrogen masses
are approximately equal and are given by (Bok, 1966).

R(LMC) = 55 kpc, R(SMC) = 63 kpc,

Mm (LMC) = 1.1 X 1042g, Mm (SMC) = 0.8 X lO^g.

Therefore, if w = w^ ,

' cr G

F7,LMc(>100MeV)-2X1()"7'

F SMC(> 100 MeV) = 1 X lO-7 photons/cm2s (X.A-5)

COSMIC-RA Y ORIGIN THEOR Y 253

It is important here to note that the fluxes given above follow immediately for
any metagalactic model because for these models, by definition, for the
Magellanic clouds as well as for the galaxy, the role played by their internal
cosmic-ray sources is unimportant and therefore wM « wG « wLMC » WSMC-

For the galactic models, on the contrary, there is no reason to expect the
above quality to hold. Even assuming similar activity of cosmic-ray sources
in our galaxy and the Magellanic clouds, it is probable that wG > wLM€
> wSMC because of the smaller sizes of the clouds and the correspondingly
more rapid escape of cosmic rays from them. Besides, in our galaxy there
is apparently a strong central source of cosmic rays (which will be presently
discussed), but in the clouds there is probably no such source.

Thus, if the metagalactic models are valid, the flux from both Magellanic
clouds should be <; 3 X 10"7 photons/ cm2 s. (Any additional nonnuclear
sources of 7-radiation in the clouds would only serve to increase the flux.)

I now turn to the important question of 7-radiation from the region of the
galactic center. Such radiation has already been observed (see elsewhere in
these proceedings). Using the values given by Clark, Garmire, and Kraushaar,
(1970), and Fichtel et al. (1972), we find

F (E > 100 MeV) = (3 - 10) X 105 photons/cm2 s (X.A-6)

On the basis of spectral measurements (Fichtel et al., 1972) and from several
indirect observations, it seems likely that we are observing 7-rays from the
galactic center region which were produced by cosmic rays and are the products
of the decay of ir° -mesons. Accepting this interpretation, we shall draw
several conclusions (Ginzburg and Khazan, 1972). By inserting the result
(Equation X.A-6) into Equation (X.A-3), we conclude that the galactic-center
region contains a cosmic-ray component of total energy

Wc = wcVc « (3 - 10) X 1066 (wG/nc) ~ (3 - 10) X lO54/^ erg (X.A-7)

taking R = 10 kpc. If we assume that the central source is larger than 300 pc,
we cannot assume that the gas density is much greater than ~ 1 cm" . (If
Lc ~ 1021 cm, Vc ~ 1063 cm3 and Mc ~ 2 X 1039nc ~ 106 ncM© where M©
is the mass of the sun. If n, ~ 10 cm"3, Mc ~ 107M@ , which is probably an
upper limit for an area of this size.) For nc ~ 1 cm"3 , it follows from
Equation (X.A-7) that Wc ~ (3 -10) X 1054 erg, which is only an order of
magnitude smaller than the total energy of cosmic rays in the galaxy
(Ginzburg, 1970; Ginzburg and Syrovatskii, 1971).

On the other hand, a result of the order of 1055 erg is obtained from an
analysis of astronomical data indicating that there was an explosion in the

254 THEORY

region of the galactic nucleus approximately 107 years ago (Oort, 1971;
Van der Kruit, 1971). A similar number for the energy of cosmic rays produced
in an explosion of the galactic nucleus was used in Ginzburg and Syrovatskii
(1964).

If the size of the central 7-ray source is less than 200 to 300 pc, then nc can be
greater than 1 0 cm"3 . We then obtain a smaller estimate for Wc from
Equation (X.A-7), but the intensity of cosmic rays Icr c = Ic is not diminished.
For example, if n, = 10 cm-3 and Vc = 1063 cm3,, Wc ~ (3 - 10) X 1053 and
lJlG = Wc/WG ~ (3 - 10) X 102 . It seems that it would be rather difficult
to confine cosmic rays within a smaller volume for 107 years. Therefore the
value of W ~ 3 X 1053 erg would seem to represent a lower limit and it is more
likely that W <; 3 X 1054 erg. If this is the case, the central cosmic-ray source
would be essential from the point of view of the total energy balance of cosmic
rays in the galaxy. The average power of injection would be Uc ~ Wc/Tc ^
1 040 erg/s with T = 107 yr. The number is of the same order of magnitude
as the total power of injection used in the galactic-origin models (Ginzburg and
Syrovatskii, 1964; 1970).

If future measurements confirm the existence of a central galactic 7-ray
source of 7r°-decay origin, then we will have one more important argument
against the metagalactic models for the origin of cosmic rays, since our own
galaxy will then prove sufficient to supply a considerable part of the observed
cosmic rays as opposed to other galaxies and quasars which would be the pre-
dominant source of cosmic rays in the metagalactic models. This would be true
without even taking into account the production of cosmic rays in supernovae
and pulsars. (In fact, I feel that the role of supernovae is essential.) The
assumption of metagalactic sources for cosmic rays will thus become superfluous

REFERENCES

Bok, B. J., 1966, Ann. Rev. Astron. and Astrophys., 4, p. 95.

Burbidge, G.,and K. Brecher, 1971, Comm. Astrophys. and Space Sci, 3, p. 140.

Cavallo, G,and R. J. Gould, 197 1, Nuovo Gmento, B2, p. 77.

Clark, G. W., G. P. Garmire, and W. L. Kraushaar, 1968, Astrophys. J. Letters,
153, p. L 1203.

, 1970, Bull. American Phys. Soc. , 15, p. 564.

Fazio, G. G., 1968, Ann. Rev. Astron. and Astrophys., 5, p. 481.

Fichtel, C, R. Hartman, D. Kniffen, and N. Sommer, 1972, Astrophys. J., 171,
p. 31.

Ginzburg, V. L., 1970, Comm. Astrophys. and Space Phys., 2, p. 1.

COSMIC-RA Y ORIGIN THEOR Y 255

., 1971 , Proc. 12th Intl. Conf. on Cosmic Rays, Hobart, Tasmania.
., 1972, Nature Phys. Sci. , 239, p. 8.

Ginzburg, V. L., and Ya. M. Khazan, \912,Astrophys. Letters, 12, p. LI 55.

Ginzburg, V. L., and S. I. Syrovatskii, 1964, The Origin of Cosmic Rays,
New York, Pergamon Press.

, 1965, Uspekhi. Fiz. Nank, 87, p. 65.

, 1967, Radio Astron. and the Galactic System, Proc. IAU Symp.

No. 31, H. van Woerden, Academic Press, New York.

, 197 '1, Proc. 12th Intl. Conf. on Cosmic Rays, p. 53.

Oort, J. H., 1971 , Les Noyau des Galaxies, Pontifical Academia Scientiarum,
p. 321.

Setti, G., and L. Woltjer, 1971 , Nature Phys. Sci., 231, p. 57.

Shklovskii, I. S., 1971, Astron. Tsirkulyar SSSR, 661, p. 1.

Stecker, F. W., 1971 , Cosmic Gamma Rays, NASA SP-249, U. S. Government
Printing Office. Washington, D. C.

Van der Kruit, P. C, 1971, Astron. and Astrophys., B, p. 405.
DISCUSSION

Ramaty:

As far as the galactic center is concerned, what Prof. Ginzburg said is quite
clear, but the cosmic rays at the earth are probably not coming from the
galactic center. I suppose that what Prof. Ginzburg was trying to do here
was take up the argument for galactic origin of cosmic rays. But that
question is not necessarily going to be solved by understanding the origin of
the 7-rays from the galactic center.

Steigman:

I fully agree with what Reuven Ramaty has just said and the point is, you
really do not know what is causing the galactic center source. It could be an
enhancement of the density in the galactic center, which is likely to be the
case. We do know that the nonthermal radiation background is rather uniformly
distributed throughout the galaxies, so some cosmic rays are not produced
predominantly in the galactic center.

256 THEORY

I would also like to ask a question of Stecker that is related to all of this: The
Copernicus results seem to indicate a large amount of molecular hydrogen in
interstellar space, perhaps as much as the atomic hydrogen which is indicated
by the 2 1 -centimeter observations. A factor of 2 or so increase in the gas
density would seem to bring the results for the 7-ray production rate per
hydrogen atom below what Cavallo and Gould have suggested and even
below the rate Stecker suggested. Does Stecker have any comments about
that?

Stecker:

Yes. I'm glad you asked that. A factor of 2 was exactly what Ted Stecher and
I said was needed in order to explain the OSO-3 measurements, and therefore
we did postulate a significant amount of molecular hydrogen and gave argu-
ments for it a couple of years ago {Nature, 226, p. 1234, 1970; see also Nature,
222, p. 865, 1969).

With regard to the number of cosmic rays at the galactic center and the gas
at the galactic center, let me add that in the same paper we estimated that we
could only explain about half of the flux seen in the direction of the galactic
center on the basis of an increased gas density. On this basis I would agree
with Prof. Ginzburg and also Dr. Ulmer, who is here and did some thesis work
on this. It would seem that there may well be an enhancement in the cosmic-
ray flux toward the galactic center.

Cowsik:

Concerning the source from near the galactic center, one point seems to be
interesting to note. The number density of stars as we approach the galactic
center increases rather quickly; locally it increases as R"3 and below a distance
of about half a kiloparsec to a kiloparsec from the galactic center, it seems to
level off. If one considers the distribution of stars and if one considers the
nonthermal background and some reasonable value, in fact an upper limit
on the magnetic flux that can be there, then one knows exactly what the
electron density is. It's not substantially higher or lower than what is evident
at the earth. In fact, it is about the same. And you know the photon density
because you know the starlight density.

If you take these electrons and scatter them, you can calculate the flux of
7-ray s that you will get. They are of the right order of magnitude and do
have the right distribution of 7-rays towards the galactic center as seen by
Clark. Of course, above this center source one needs the uniform source, which
can only come by cosmic-ray prior production.

COSMIC-RA Y ORIGIN THEOR Y 257

Stecker:

Here we should point out that there are strong observational reasons now that
the galactic center does have a hard spectrum above 100 MeV and has to be
primarily of 7r°-decay origin above this energy. This is deduced from the work
of Fichtel, et al. (Astrophys. J., 171, p. 31 , 1972) that Prof. Ginzburg referred
to, so I think this is fortunately one of the things we do not have to argue
about from a theoretical point of view anymore. The flux is primarily of pion
decay origin, and I think we'll hear more about it later. (See papers of Share and
Kniffen, Chapters IV. A and IV.C.)

Cowsik:

I'm just commenting that the flux from the galactic center goes up approxi-
mately as the star density increases.

Stecker:

But by the same argument, we know the gas density and the dust density go
up toward the galactic center.

Cowsik:

It goes as 1/R.

Vette (Session Chairman):

Let's carry on this one in the coffee break.

B. GALACTIC GAMMA RAYS: MODELS
INVOLVING VARIABLE COSMIC-
RAY DENSITY

A. W. Strong, J. Wdowczykf and A. W. Wolfendale*

University of Durham

MODELS INVOLVING VARIABLE COSMIC-RAY DENSITY

It is well known that the variation of the 7-ray flux around the whole galactic
plane (Kraushaar et al., 1972) cannot be explained on the simple model of n0-
production in cosmic-ray interactions with the interstellar gas, if the cosmic-
ray distribution is assumed uniform and the observed distribution of neutral
hydrogen is taken as the only significant gas component (See for example,
Clark et al., 1970). This model gives roughly the correct intensities away from
the galactic center, but does not reproduce the observed increase by a factor
of about three toward the center.

It is possible that point sources are responsible, but models involving supernova
remnants as sources appear to be inadequate (de Freitas Pacheco, 1973). If
the mechanism is predominantly it0 -decay, as indicated by the results of
Fichtel et al., (1972), then other ways of producing the central increase
include the presence of large amounts of molecular hydrogen or an increase in
the cosmic-ray density toward the center.

DETAILS OF THE MODELS USED AND RESULTS

We have investigated the last possibility for two particular models of the
variation of cosmic-ray density with position in the galaxy. It seems likely that
there will be a correlation between mean magnetic field strength (H) and the
cosmic-ray density, if the cosmic rays are generated within the galaxy. If
they arise in sources of high field strength HQ » H, then we might expect

'On leave from the Institute of Nuclear Research, Lodz, Poland.
*
Speaker.

259

260 THEORY

(Woltjer, 1965), from Liouville's theorem, that the cosmic-ray density will
be proportional to H. Alternatively, if there is equipartition of energy
between cosmic rays and magnetic fields, then cosmic-ray density will be
proportional to H2.

Thielheim et al., (1971) have used their model of the galactic magnetic field
(Thielheim and Langhoff, 1968) to predict the distribution of synchrotron
radiation in the galaxy, assuming the cosmic-ray electron flux is proportional
to H, and they find that it is consistent with the observations at 400 MHz.

For the variation of H with distance (R) (kpc) from the galactic center, we
have taken the radial part of their model, that is,

Hcxexp (^ I -£-)}

1 100

The data on line-of-sight distribution of neutral hydrogen from 21 -cm surveys
was used to calculate the weighted column density

NH(8,b) = / w(p, 8) nH (p, 2, b) dp (X.B-2)

where p = distance from the sun, and w is a weighting factor given by (1) w = 1,
that is, constant cosmic-ray density (equivalent to metagalactic origin); (2)
w = H/Hq, where HqIs the field at the sun given by Equation (X.B-1); and
(3)w = (H/H0)2.

The resulting line fluxes were calculated assuming a rectangular response and
adopting the yield function given by Cavallo and Gould (1971):

-25 [+6o

j(fi, > 100 MeV) =— k / NH (£,b) db cm"2-^1 rad'1 (XB"3)

4tt J_q

where 0 = 15° and k is a constant to allow for unseen components of the
interstellar gas (such as H2), and the possibility that the observed cosmic-ray
flux at the earth may not be representative of the local mean flux. The
results of these weightings are shown in Figure X.B-1, for k=l. For w=l and
w = H/Hq, there is insufficient increase towards the center to fit the observa-
tions. For w = (H/Hq)2 the fit is quite good for k = 1.5, as shown in Figure
X.B-2. The value k > 1 is indicated by recent observations of significant
amounts of H, in dense clouds in the interstellar medium.

The advent of detectors with better angular resolution (such as that aboard
SAS-2) should allow an improved assessment to be made.

GALACTIC GAMMA RAYS: MODELS

261

180 160 140 120 100 80 60 40 20 0 340 320 300 280 260 240 220 200 180 160
GALACTIC LONGITUDE f1

Figure X.B-1. Gamma-ray line fluxes calculated from Equation (X.B-3),with
0=15° and k = 1 . The curves are for weightings (1 ) w = 1 , (2) w = H/HQ,and
(3)w=(H/H0)2.

180

160 140 120 100 80 60 40 20 0 340 320 300 280 260 240 220 200 180
GALACTIC LONGITUDE f1

Figure X.B-2. Line fluxes observed by Kraushaar et al. (1972), after sub-
traction of diffuse background, and prediction of model with w = (H/Hq)
and k = 1.5.

Note: This account differs in several respects from that presented at the
Symposium and is to be taken as superseding that account. Stecker
(Chapter IX.A) has used the most recent accelerator data to obtain a value
for the yield above 100 MeV of 1 .3 X 10"25 s"1 per H atom. This implies
k=2.1 for Figure X.B-2. This value is still plausible for the reasons
stated above.

262 THEORY

REFERENCES

Cavallo, G., and R. J. Gould, 1971, Nuovo Gmento, 2B, p. 77.

Clark, G. W., G. P. Garmire, and W. L. Kraushaar, 1970, IAU Symposium
No. 37, p. 269.

Fichtel C. E., R. C. Hartmann, D. A. Kniffen, and M. Sommer, 1972,
Astrophys. J., 171, p. 31.

de Freitas Pacheco, J. A., 1973, Astrophys. J. Letters, 13, p. L97.

Kraushaar, W. L., G. W. Clark, G. P. Gamire, R. Borken, P. Higbie, C. Leong,
and T. Thorsos, 1972, Astrophys. J., 177, p. 341.

Thielheim, K. O., and W. Langhoff, 1968, /. Phys. A: Gen Phys., 1, p. 694.

Thielheim, K. 0., H. J. Kuchoff, W. H. Steib, and G. Wenner, 1971, Proc.
12th Int. Conf. Cosmic Rays, Hobart, Tasmania, 7, p. 2612.

Woltjer, L., 1965, Stars and Stellar Systems, Galactic Structure, V, Chicago,
p. 531.

Chapter XI

A. PROSPECTS FOR NUCLEAR-GAMMA-RAY
ASTRONOMY

Donald D. Clayton*

Rice University

INTRODUCTION

Each new astronomy has provided us with new types of information. Radi-
ations of vastly differing wavelengths tend naturally to have their origins in
differing physical processes of emission, so that the different astronomies
record, by and large, differing types of events. The enrichment of astronomical
knowledge is obvious. If history is any reliable guide, we can expect to detect
7-ray lines emitted during the electromagnetic deexcitation of nuclei. Their
observation will confirm that excited states of nuclei are being produced, and
the fluxes and spectra will identify the specific nuclei and their rate of
excitation. Because extreme physical circumstances are required for the
production of excited nuclei at low densities where they can be seen, unique
information about the source regions will be obtainable.

In this paper prospects for two sources of 7-rays from outside the solar
system are considered. Both radioactive decay and inelastic collisions produce
nuclei in excited states. As Rutherford emphasized from the beginning, the
radioactivity would have all passed away were it not being continually re-
plenished. Therefore, radioactive 7-ray sources in space will be associated
with events of nucleosynthesis— probably supernova explosions of some type.
The fluxes and spectra will depend on the yield of radioactive nuclei, their
7-ray emission lines, and their half-lives. Inelastic collisions with high-energy
cosmic rays are probably not important sources as far as nuclear-deexcitation
7-rays are concerned. The average high-energy fluxes are known to be too
small. The best prospect here is for much larger fluxes of MeV particles,
especially near the source regions. My attention will fall outside the solar
system, thereby intentionally passing over the sun, moon, and planets as
interesting special sources.

^Speaker.

263

264 THEORY

EXPLOSIVE NUCLEOSYNTHESIS

The idea that the common intermediate-mass nuclei are synthesized during
their explosive ejection (Arnett and Clayton, 1970) from stars, rather than
before it, has one extremely important observational consequence. Several
abundant nuclei are ejected in the form of radioactive progenitors, and their
decay outside the star can clarify many unproven hypotheses concerning
nucleosynthesis. Specifically, if the 7-ray lines from radioactivity in supernova
ejecta and in the accumulated background of the universe can be detected
(and the anticipated fluxes are promising), it will be possible to:

• Prove supemovae eject new nuclei and measure the supernova yield,

• Prove nucleosynthesis occurs during the explosion rather than prior
to it,

• Measure the supernova structure by the profiles of the lines and
their Compton tails,

• Discover Galactic supernova remnants,

• Demonstrate that nucleosynthesis is occurring today in the universe
and measure its average rate today in the isotropic background,

• Determine whether the average rate of nucleosynthesis has been
relatively constant or peaks in the distant past,

• Gain additional information about the average density in the universe,
and

• Evaluate evolving versus steady-state cosmologies.

That is a lot to promise; if it is correct, these observations will be as enter-
taining and profound as other great experiments in astronomy, such as the
solar neutrino experiment and the microwave background experiment, for
example. My object will be to outline these possibilities as a guide to the
chances of successful detection.

The Radioactive Species

The most abundant species having a radioactive progenitor is ^Fe. Bodansky,
Clayton, and Fowler (1968) showed that ejecta in the process of silicon
burning resemble the solar abundances between A = 28 and A = 57 if they
contain roughly equal amounts of 28Si and ^Ni. This result suggested that
several prominent nuclei, primarily ^Ca, 48Ti, and 56Fe were rejected as
radioactive ^Ti, 48Cr, and 56Ni respectively. Clayton and Woosley
(1969) strengthened that result by showing that if the silicon burning
had occurred slowly enough for j3 decays to raise the neutron excess
to a value for which 56Fe itself could be ejected during silicon burning,
implausible overabundances of key species would result. They further

PROSPECTS FOR NUCLEAR- GAMMA-RA Y ASTRONOMY 265

strengthened the case for ^Ni by showing that something similar to an e-
process centered on ^Ni would also synthesize otherwise troublesome 58Ni,
especially if the free-particle densities were somewhat in excess of their equil-
ibrium values. Clayton, Colgate, and Fishman (1969) used these discoveries
to make the first estimates of the importance of ^Ti, ^Cr, and ^Ni to the
7-ray astronomy of young supernova remnants. Because of the centrality of
the 56Ni versus 56Fe argument, Hainebach, Arnett, Woosley, and Clayton
(1973 preprint) have pursued the evidence favoring 56Ni even further. They
show that two- or three-component e-processes with differing neutron enrich-
ments (and with freezeout corrections) overwhelmingly select 56Ni produc-
tion when asked to produce the solar abundances by superposition. I think
the evidence now makes it virtually certain that 56Fe was ejected dynamically
from the synthesizing events as 56Ni. The preference for low-77 solutions
[Arnett and Clayton (1970); Arnett (1971); Hainebach, Arnett, Woosley, and
Clayton (1973 preprint)] in explosive burning of carbon, oxygen, and silicon
and continuity arguments strongly suggest that 44Ca and 48Ti were also
ejected as 44Ti and 48Cr. The solar mass fractions of these species, their half-
lives, and the prominent 7-ray lines emitted during their decay are included in
Table XI.A-1. The 56Co->56Fe decay should be, because of its rich spectrum,
high abundance, and 77-day half-life, the single most important radioactive
decay for 7-ray astronomy. It remains possible, however, that a less abundant
product may prove to be easier to detect if the exploding remnants remain
opaque too long.

Clayton (1971) discovered that a significant fraction of ^Ni was probably
synthesized as radioactive ^Fe, with r„ = 3 X 105 yr, or perhaps as ^Co,
with t1/2 = 5.26 yr. In either case, 7-rays of 1.17 MeV and 1.33 MeV are
subsequently emitted. The arguments for and against ^Fe synthesis are
complex and by no means certain. About 1 percent of ^Ni could be
synthesized by arresting about half of the Cr seed at ^Cr (which decays to
60 Fe) in the rapid neutron-induced reactions on seed nuclei during explosive
carbon burning (Howard, Arnett, Clayton, and Woosley 1971 ; 1972).
Several to fifty percent of ^Ni may have been synthesized as ^Fe directly
from ^Fe-seed nuclei in the same event. Clayton (1971) has made the in-
triguing observation in this regard that only ^Ni is abundant enough to have
absorbed the ^Fe seed in explosive carbon burning, thereby suggesting that
much of the iron seed has been arrested at 60Fe. Because of the strong
(p,n) flows during high-temperature carbon burning, it also seems plausible
that a percent or so of the ^Ni is due to ^Co nuclei ejected in the explosion.
Although ^Co synthesis should be less efficient than ^Fe synthesis, it may
nonetheless be more important in young remnants because of its favorable
half-life, which is long enough to assure transparency yet short enough to have
a detectably high decay rate. Without going into the matter further here, I

266 THEORY

let p be the percentage (fraction X 100) of ^Ni nuclei synthesized as ^Fe
nuclei and p' be the percentage synthesized as ^Co, and I expect

1>P60(%)<50

0.1<P6O(%)<5

I note here that Clayton (1971) did not explicitly include ^Co in his
considerations. However, there do appear to be circumstances in which the
7-rays due to ^Co synthesis could, for many years, exceed those due to
synthesis of all other nuclei.

The r-process synthesizes many heavy radioactive nuclei, which are expected
to have unfortunately small yields. Clayton and Craddock (1965) considered
the flux expected from supernova remnants if the r-process yield were great
enough for the "californium hypothesis" of Type I light curves to be correct.
In particular, they calculated the expectations of the Crab Nebula in that
regard. There is a large range of half-lives present in initial transbismuth
debris, however, so their conclusions on the 920 year-old Crab (that the
strongest line should be no greater than 10^ cm'2 -s"1) would require re-
calculation for remnants having different ages and distances. The main
problems with this idea would seem to be that it requires the /--process to be
concentrated in relatively rare events in order that these nuclei not be greatly
overproduced, and that there seems to be no compelling reason to associate
the Type I light curves with radioactivity. I therefore currently hold little hope
for this 7-ray source, although additional clarifying remarks will be made
later.

Typical Supernova Yield

In the absence of more certain knowledge, I will assume a simple model of
galactic nucleosynthesis in supernovae. Arnett and Clayton (1970) and,
more specifically, Arnett (1971) have described the conceptual framework
more accurately; however, my aim is only to extract typical numbers for the
typical supernova event. Let the explosively synthesized nuclei be coproduced
in the same abundance ratios that we find in the solar system in identical
supernova events occurring at the galactic rate

NSN=Ret/TR (XI.A-1)

Fowler (1972) finds that TR - 4 X 109 yr and galactic age AG = 12 X 109 yr
are not unreasonable caricatures of r-process nucleosynthesis (which I take
here to characterize all explosive nucleosynthesis). Taking a current supernova

PROSPECTS FOR NUCLEAR-GAMMA-RA Y ASTRONOMY

267

i-H

s— • /— V

cn o

O CN

-1 +

cn'

»T>

o" w

v^-^

NO CM

cn °.

B .2

- CN

^— N.

^•^

/<~\

60 -g

/— ^

©^

*— »

t"-

3 H3

**~s

C^ /-^

«/">

^^^

^3 S

NO ro

■^

T3 C;

w —

»— <

\^/

00 g?

*T w

■*

^*-^

3 O

CN no

r-;

<L>

"cL °-

*-;

»— i

•-3 cfl

c <u
2 ^

~ en

/"— s

O /-?

U-) ,^-v

O <tf

oo m

i-H

3 b

CN w

Tf NO

_i vo

oo r-

oo in

»— 1

«— <

y— i

d -h

d -5

v*H

*»■

u--

«t

</>

'-,

< — i

••h

■— 1

«™H

— H

>— •

to.

l/i

t-^

I-H

H
>>

•a

^ l_

2 ^

-o

NO
X CN

1 '

cn u-i

ex o*

1 "3,

o
U

.„,

<D O

£ 1

u
Oh

8 S

rr>

«-H

i-H

— H

>— i

<— 1

«-H

»"H

*""'

00 p

268

THE OR Y

rate N°SN(today) = 0.25 yr"1 then gives R = 0.5 yr"1 . The initial supernova
rate would, with these particular parametric values, have been twenty times
greater.

Let the average yield of the typical event be such that its product with the
total number of events prior to the birth of the sun shall have produced a
galactic mass having solar composition. The total number of such events is

NSN =

NSNdt = N°SNTR

1-e

-WT

,ag/tr

(XI.A-2)

where t@ is the time of solar formation (approximately 7 X 109 yr). The
number of events is nearly exponential in AQ/TR and multiplied by TR if
TR < t@, as seems likely. With the specific choice of parameter values taken
above, AG/TR = 3 and the number of events would have been NgN =
1.7 X 109.

If the mass of the galaxy is 1.8 X 1011 Mq (Schmidt, 1965) and the mass
fraction of iron in the sun is Xq = 1.3 X 10"3 (Cameron, 1968), and if the
average composition of the galaxy at that time was solar, the galaxy would
have contained 2.3 X 108 Mq of ^Fe. The average yield for each of the
1.7 X 109 contributing events would have been

2.3 X 108 M0 of ^Fe

M„N(56Fe) = — 2 =0.14M@/SN

SN 1.7 X 109SN events ®

(XI.A-3)

The corresponding number of ^Fe atoms per event is

»33

Y,N(*Fe) =

0.14(2.0 X 10JJ)(6.0X 10")
56

= 3.0X 10

(XI.A-4)

which would have been ejected initially as ^Ni atoms. These numbers for
several interesting abundances formed explosively as radioactive progenitors
are shown in Table XI.A-1.

It is not difficult to question the appropriateness of many of the assumptions
leading to this estimate. However, my point of view is that the simplest
reasonable argument is the most appropriate one for gearing our expectations.

Table XI.A-1 shows the total yield of ^Ni to be Y^^Ni) = 4.4 X 1052
atoms/supernova. According to the earlier discussion, the yields of ^Fe and
^Co are evaluated as

PROSPECTS FOR NUCLEAR-GAMMA-RA Y ASTRONOMY 269

YSN(60Fe) = 4.4X 1050 Pfi0

(XI.A-5)

YSN(60Co) = 4.4X1050p'60

The yield of 238U under these assumptions is listed in Table XI.A-1 only as an
example of transbismuth r-process yield rather than as a nucleus of particular
importance for 7-ray astronomy. Indeed, Clayton and Craddock (1965) found
that the most important nuclei for the Crab were likely to be M9Cf and 214Bi.
Nonetheless it is instructive to note that this "typical 238U yield" is about
four orders of magnitude too small for that required for the californium
hypothesis of the light curve. If the latter hypothesis is correct, the /--process
will have to have occurred in events about 1 04 times less numerous than the
typical supernovae we are considering in this section. Whereas this is possible,
it suggests that all Type I events are not r-process events, in which case the
original hypothesis loses its "raison d'etre."

Typical Line Fluxes

If species z decays with mean lifetime t (z) = 1/X , and if each decay is
accompanied by g. photons of type i, then the flux of those 7-rays at the
earth due to a nearby supernova is

xzYSN<z> -u

F.=g. — e z (XI.A-6)

1 ' 4;rR2

where R is the distance to the supernova and t is the time since its detonation.
This formula neglects attenuation due to absorption or scattering in the source
and therefore, will be correct only for times which are long enough so that
the expanding remnant has become transparent to 7-rays.

Using information from Table XI.A-1 one obtains

IIV If)4

F.(*Ni) = g. ; e**8-8*) cm*-!1 (XI.A-7)

1 * R2 (kpc)

F.(*Co) = g.2fX1°3 e«/llld> cm2^1 (XI.A-8)

1 R2 (kpc)

F-C48 V) = g. - 26 e<W»» cm"2 s1 (XI. A-9)

'R2(kpc)

270 THEORY

2 1 X 10"2

F.^Ti) = g.— e"(t/69yr> cm**1 (XI. A- 10)

1 ' R^ (kpc)

F.(»Fe) = 2J X 10"7 P60 e-(t/4.3 X 105yr) cm-2.s-l (XI ^ j}

^R2(kpc)

F.rCo) = qiL6X1°2p60e-t/7.6 yr)cm-2.s-i (XI.A-12)

R2(kpc)

Several of these fluxes are shown in Figure XI .A- 1 as a function of time. The
supernova itself has been placed at R = 103 kpc to emphasize that the A = 56
lines may even be observable from supernovae in other galaxies. A supernova
in M31 , for example, would present a ^Co line flux above the detectable level
of about 10"* cm"2-s_1 for more than a year. These lines show a rise time
rather than a pure exponential decay, because a specific model was adopted
by Clayton et al. (1969), for the transparency of the expanding supernova.
They took a rather optimistic (in light of recent nucleosynthesis theory) model
-a 0.5-M@ ball of iron expanding at 1 .7 X 109 cm/s so that the product of
mean density times radius is

p(t) R(t) * 8 X 1013 t^gm-cm"2- (XI.A-13)

which falls below 10 gm-cm"2 (a rough estimate of the optical depth for 7-rays)
for t > 3 X 106 s. Thus, at best, the lines will be poorly visible for the first
month. Even then it is clear that Compton scattering will have a serious
effect on the 7-ray spectrum near those times when. they begin to emerge.
Brown (1973) has calculated this effect for some special cases similar to those
considered by Clayton et al. (1969). Figure XI.A-2 shows one of his results
when 3.5-MeV and 1 .25-MeV lines are emitted isotropically from a depth of
18.6 gm-cm"2 within an iron sphere of radius 37.2 gm-cm"2 . The total
mass of such a sphere depends upon its metric radius, of course, so with R(t)
= 1.7 X 109t we find that the mass whose line spectrum corresponds to
Figure XI.A-2 is

M(Fig. XI.A-2) = lQ6sj M@ (XI.A-14)

Therefore the total mass of that example could be any reasonable multiple of
a solar mass at time of order a few months. The question of the mass of
layers over-lying the CO core at the time of detonation is even more uncertain,
but it will clearly be worthwhile to evaluate dynamic models of 7-ray opacity
for exploding massive evolved stars. For the time being I wish only to
emphasize that whether the ^Ni lines emerge at all (they did in Figure XI.A-1)
depends on the structure and dynamics of the exploding object. Ideally we

PROSPECTS FOR NV CLEAR- GAMMA- RA Y ASTRONOMY

271

10-2

Ni56(0.812/
MeV)

d = 106 pc
M(Ni56)=0.14Mo

(sec)

Figure XI.A-1. Prominent medium-lifetime 7-ray line fluxes as a function
of time from a distant (d=106 pc) supernova ejecting 0.14MQ of 56Ni
and 2.0 N 10-4MM of 44Ti. The early growth reflects the increasing trans-
parency of an expanding model (Clayton et al., 1969)

may one day watch these and the ^Co lines rise to peak intensity before
beginning their decay, and the rise time of these fluxes will be a crucial
measure of the structure of the exploding object. The 270-day 57Co lines
from 57Ni progenitors may also play an important role in this problem
(Clayton, 1973), although I have not included them here due to their relatively

272

THEORY

0.05 0.1

0.5 1.0
Energy (MeV)

Figure XI.A-2. Effect of Compton scattering on 3.5-MeV lines (solid
histogram) and 1.25-MeV lines (dashed histogram) emitted isotrop-
ically from a point at a depth of 18.6 gm-cm"2 from the surface of an
Fe sphere of radius 37.2 gm-cm"2 (Brown, 1973).

low energies (136 keV and 122 keV). This astronomy will allow us to measure
that structure somewhat analogously to the way neutrino astronomy has
allowed us to measure the interior of the sun— and probably with similar sur-
prises.

The 1.16-MeV line emitted subsequent to the decay of the ^Ti could be quite
strong in several present galactic remnants, and will surely emerge even if the
A = 56 lines should happen not to get out. In this sense the ^Ti synthesis may
prove to be extremely important. The real need, of course, is for the galaxy
to arrange a visible supernova, preferably after (if ever) instruments like
HEAO-B are operational. The A = 48 lines, on the other hand, seem likely to
be of no special importance, because they are both weaker and shorter lived than
the 56Co lines.

PROSPECTS FOR NUCLEAR-GAMMA-RA Y ASTRONOMY 2 73

The ^Co lines have not been entered on Figure XI.A-1, but a comparison of
Equations (XI.A-1 2) and (XI.A-1 0) show that they are comparable to those of
44 Ti for about 10 years if p ' is around unity (that is, about 1 percent of ^Ni
is due to synthesis of ^Co, which requires about 2 percent of ^Fe seed to
reside at ^Co at completion of explosive carbon burning). Remnants through-
out the galaxy (R < 20 kpc) should ultimately prove detectable for a decade.

The ^Fe lines (actually the same as the ^Co lines but with a much longer half-
life) are also not shown in Figure XI.A-1 . They are a special case due to the
long 3 X 105 yr half-life, which ensures that many radiating remnants exist
but they may have large angular size due to the long time available for dispersal.
For the flux to exceed a detectable 10"4 cm"2^'1 requires R < 160 pc if
10 percent of ^Ni is due to synthesis of ^Fe. A circle of 160 pc radius con-
stitutes about 10^ of the area of the galactic disk and should thus contain one
of the approximately 104 supernovae that should have occurred during the
lifetime of ^Fe. However, the size of a remnant 10 years old might cover a
significant fraction (even half!) of the sky for an event about 100 pc away;
therefore, simple on-source-off-source differences will have to be measured
with this in mind. The radiation from such sources seems more likely to
appear as a general galactic background. The general flux from a wide angle
containing the galactic center would be

F60(galactic) ^3X10^ cm"2 s"1 (XI.A-1 5)

if pm is about 10 percent. This is also about the same as the average flux from
the galaxy due to the 44Ti lines (Clayton, 1971), but in this case the actual
flux depends on the details of the positions and times of the last few galactic
supernovae.

As a very crude estimate of transbismuth fluxes, I will assume that every trans-
bismuth species is synthesized with a yield Y^ equal to that listed for 238U
in Table XI.A-1 . There are so many different half-lives in the /•-process ejecta,
moreover, that one may roughly assume that, whatever the age of the remnant,
there exists one 7-producing nucleus with a half-life approximately equal to
the age of the remnant. This species produces the largest flux. In this case
Equation XI.A-6 becomes

F.(X.=f !) = q. - -r (XI.A-1 6)

which has the approximate value

_ 1.25 X 10"5 „ ,

F = — cm^-s"1 (XI.A-1 7)

r R2 (kpc) t (yrs)

It is obvious that these fluxes will not commonly be observable unless the
r-process is restricted to much rarer events, thereby raising the yield of each

274 THEORY

event. This conclusion, stated earlier, renders this particular prospect unlikely.
Clayton and Craddock (1965) took a yield four orders of magnitude greater to
provide radioactive power for the Crab light curve and were thereby able to
calculate marginally detectable lines from the Crab. Equation (XI.A-17) yields
only Fr % 10"8 cm"2^"1 from the Crab, and is probably a more realistic estimate
The site of the r-process is so poorly understood, however, that a great surprise
would come as no shock.

The Universal Background

One need only appreciate that the average galactic luminosity due to radio-
active 7-rays has been 3 X 1040 erg-s"1 to realize that their contribution to
the isotropic background radiation may be significant. The cosmological
principle allows us to estimate their flux very easily. Taking H = 55 km s"1
Mpc"1 (Sandage, 1972), the observed universal density of matter is p =
1 .7 X 10"31 gm-cm"3 (Oort, 1958). If the average mass fraction of ^Fe is
X@ (56Fe) = 1.3 X 10"3, this corresponds to 2.2 X 1034 gm-cm-3 of ^Fe.
Consequently, the average iron number density in the observed universe is

nC^Fe) = 2.3 X 10'12 cm"3 (XI. A- 18)

The flux of these 7-rays/ sr is (Clayton and Silk, 1969)

7=7gn(KFe) (XI.A-19)

where gY is the number of 7-rays emitted per ^Fe nucleus synthesized. The
value of g^ is 2.8 for only ^Co decays and g = 4.9 if both 56Ni and 56Co
decays are used. Taking the latter value yields

= 2.7 X It)'2 cm-2- sr"1 -s'1 (XI.A-20)

To emphasize the size of this flux, Clayton and Silk (1969) pointed out that it
is as large as the total integrated universal background at photon energies in
excess of 300 keV. Clearly it must be an important component of that back-
ground unless the A = 56 lines do not escape from their sources. Because this
estimate is based on the observed mass density, it will be proportionately
greater if the universe contains "hidden matter" that has also synthesized ^Fe.

The simple density argument does not determine the frequency distribution of
the photons comprising the flux in Equation XI.A-19. For those ^Fe nuclei
synthesized early in the universe, the associated 7-rays will now have been
considerably red shifted. It is just this feature that allows the spectrum to
carry a wholly new astrophysical datum; that is the red shift distribution in the
7-ray spectrum measures the distribution of the ages of ^Fe nuclei. Hidden
in it is the chronological account of the rate of nucleosynthesis.

Let f(t) be the rate per unit of cosmic time at which ^Fe nuclei were (and
are being) synthesized. Let it be normalized such that

PROSPECTS FOR NUCLEAR-GAMMA-RA Y ASTRONOMY 2 75

f(t)dt = l (XI.A-21)

so that f(t)dt is the fraction of all 56Fe nuclei that were synthesized at cosmic
time (t) in the interval dt (t = cosmic time today). It follows that f(t)dt is
also the fraction of A = 56 7-rays whose travel times are t - t in the interval
dt. In any standard cosmological model the travel time t - t is some
function of the red shift (z). Thus f(t) = f(z) and f(t) (dt/dz)dz becomes the
fraction of the photons having the red shift z in the interval dz. The 7-ray
source function per unit time per unit ^Fe nucleus per unit energy interval in
the rest frame is just

P(E,t) = 2 P. (E,t) = 2 g. 5 (E-E.)f(t) (XI.A-22)

i i

where the sum is over the lines of type i emitted with rest energy E. emitted

with rest energy E. at the rate of g. 7/56Fe nucleus synthesized. The differential

flux today due to 7-rays of type i is

92F. c _, f ° R(t ) [R(t ) ]

■ ^—n^Fe) —2-?. — ^-E,t dt

dEdn 4n R(t) ' I R(t)

^nC^Fe) J (l+z)P. [o+z)E,t]-^dz

(XI.A-23)

where E is the energy today of the photon and R(t) is the scale factor of the
universe. (See for example McVittie, 1965). Because

E R(t) 1

E; R(t ) 1+z

(XI.A-24)

the integral over cosmic emission time can also be expressed as an integral over

received energies:

R(t ) dE
dt = —-2- — (XI.A-25)

R(t) E

This integral is easily done due to the 5-function nature of P. to give

32F. c g. nC^Fe) R(t )

!_ =_fl_l 1 -1-qL f(t ) (XI.A-26)

dEdSl 4ir E. R(tE) E

where the time t£ is the solution of Equation (XI.A-24). In the Friedman dust
models one has (Weinberg, 1972)

R(tJ E. T E."l '/>

_e: = _L_H l-2q +2q -± (XI.A-27)

R(tF) E ° L 4° %EJ

276 THEORY

so that Equation (XI.A-26) reads

32F. „ g. n^Fe) I E. T2

L. = £-JL± J f(t_) l-2q+2o» (XI.A-28)

dEdtl 4tt E.Ho 'E L ° ° EJ

Clayton and Silk (1969) evaluated the flux in a simpler form for the two cases
where

R(t)at1/X (XI.A-29)

They are the low-density universe (q = 0, y = 1) and the Einstein-de Sitter
universe (q = Vi, y = 3/2). In those cases Equation (XI.A-24) can be explicitly
solved for tF and, furthermore, the factor involving q simplifies:

*% _c

9E3£2 47T

,n(56Fe) f /EVl/Ef1

It is straightforward to confirm with Equation (XI.A-28) or Equation (XI.A-30)

that ,_

32F

L dE =i- g. nC^Fe) (XI.A-31)

3E912 4tt '

as required by photon conservation.

The spectrum due to each line is characterized by a step at the rest energy

(£!>_} <L**!fl.C« ) (XI.A-32)

\dEan/F=F 4tt eh

E=E.
1

that is directly proportional to the average rate of nucleosynthesis today in the
universe. Detection of the series of correlated rest edges will confirm that
nucleosynthesis is still occurring and measure its present rate f(tQ). Each rest
edge is followed at immediately lower energies by identical red shifted continua,
whose shape and extent depend upon the cosmological model and the history
f(t) of galactic nucleosynthesis. It is interesting to note that the ratio f(to)/HQ
is a ratio of characteristic times: 1/H is approximately the age of the universe
and l/f(t ) would be the time required to synthesize XQ (^Fe) at a constant
ratef(to)°

Some simple profiles for the ^Co line of 1 .24 MeV are shown in Figure XI.A-3.
If nucleosynthesis has occurred within galaxies at a constant rate up to the
present time (t ), and since it began at some time (t*), then f(t) = (tQ-t*)"

PROSPECTS FOR NUCLEAR-GAMMA-RA Y ASTRONOMY

277

Figure XI.A-3. Differential flux due to a single line (56Co at 1.24 MeV).
Models of constant galactic synthesis of ^Fe over a period of 7 X 109 yr are
shown on the left, and models of exponentially decreasing nucleosynthesis are
shown on the right. This rather short duration of galactic nucleosynthesis
was chosen only for ease of comparison, so that it could fit in the age of the
Einstein-de Sitter universe with H =75 km/s/Mpc. The low-density universe
is shown as a solid line and the Einstein-de Sitter as a dashed line. The steady-
state-universe line profile is dotted on the left figure. The line 10"2 E"2 is also
shown to indicate the approximate level of the observed diffuse background.
(Clayton and Silk, 1969).

between t* and t and is zero elsewhere. The left half of Figure XI.A-3 shows
that case from Clayton and Silk (1969), who took tQ- t* = 7 X 109 yrs so
that it could fit easily within the Einstein-de Sitter universe based on Hq=75
km/s/Mpc. The right half of Figure XI.A-3 shows this line profile for ex-
ponential nucleosynthesis f(t) = A exp[-X (t-t*)] , where A is a normalization
constant and X = 1/TR from Equation (XI.A-1). The two choices of X
shown there give different relative strengths to present-day nucleosynthesis
in comparison with the initial galactic rates. The rest edges are still detectable
here, but smaller than for the case of constant nucleosynthesis.

278 THEORY

It is worth noting here that these figures are applicable to Sandage's (1972)
value H = 55 km/s/Mpc if one only increases t - t* by the factor 75/55,
giving the more reasonable t - 1« - 9.6 X 109 yr, and if the value of the flux
is reduced by the factor (55/75)2 . The latter comes about because n^Fe)
a Hq2 and f(t) ex Hq if we require t - t* a H _1 . It is clear that the flux at

this rest edge may well be comparable to the isotropic background, whose
approximate value is shown for comparison. The model TR = 4 X 109 yr and
AG = 12 X 109 yr used in estimating the typical supernova yield resembles
the curve labeled X = (2 X 109 yr)"1 in Figure XI.A-3. Its rest edge is the
smallest shown— about 15 percent of the observed background. Such small
edges would go undetected unless observers design detectors and use data
reduction methods designed to extract the steps from the continuous back-
ground.

The steady -state universe, shown as a dotted line in Figure XI.A-3, affords a
somewhat different problem. To maintain a constant iron density requires
a creation rate

C = 3Hn (^Fe) = constant (XI.A-33)

so the 7-rays are created at the rate g. C. The age distribution of ^Fe nuclei
is no longer given by the galactic production function f(t), because galaxies of
all differing ages coexist. The density of nuclei having age t - t in the interval
dt is simply

dN(to - 1) = SHn^Fe) e^Mdt (XI.A-34)

Both results follow directly from the fact that the scale factor for the proper
distance between comoving-coordinate points is

^o) = eH(t0-t)=1+z (XI.A-35)

R(t)

where z is the red shift of a photon whose travel time is t - t. Since Equation
(XI.A-35) is also the ratio of the rest energy E. to the received energy E,
Equation (XI.A-34) is easily rewritten as a distribution in energy of photons of

tyPe': « /E\'dE

dNi(E) = 3g]n(56Fe)lYl — (XI.A-36)

and the differential flux is, as before,

(XI.A-3 7)

PROSPECTS FOR NUCLEAR-GAMMA-RA Y ASTRONOMY 279

The flux is independent of both the Hubble constant and the details f(t) of
galactic production, and the spectrum is proportional to E2 up to the rest
edge E.. This spectrum is the dotted one in Figure XI.A-3. It is of interest
to note that Equation (XI.A-32) also gives the correct answer in this case, for
the size of the rest edge, if the present production rate f(t ) is replaced by
3H according to Equation (XI.A-33):

/a2F. \ „ Sg-n^Fe)

a L_ =£_J1___. (XI.A-38)

\3E9ft /e=e. 4tt E.

In setting this rest edge equal to those of the evolving universe in Figure XI.A-3
we have been somewhat arbitrary, because f(t )/H « 2 for the evolving
models in the left side of the figure, whereas the steady state gives the slightly
larger value 3. However, the average proper density n(56Fe) could also differ
slightly from the value inferred from the solar composition, but not much,
because the average galactic age (3H y1^ 6 X 109 yr was also approximately
the age of our galaxy when the solar system formed.

The main point of the steady -state cosmology is that the strong rest edge and
the (E/E.)2 spectrum remain even if the galactic production function f(t) were
strongly peaked in the past, as in the evolving cosmologies in the right. If the
lines emerge unscattered from the sources, a strict steady-state universe will
have very stong rest edges, similar to saw teeth.

Figure XI.A-4 illustrates the entire A=56 spectrum for the Einstein-de Sitter
case. Two points need be made: (1) the rest edges are clearly more prominent
in the case of constant galactic nucleosynthesis than they are in the e"2 -
exponential case; but (2) the general shape of the continuum feature produced
is quite similar for the two cases. The Einstein-de Sitter universe requires 30
times more mass than has been observed in galaxies, but Figure XI.A-4 assumes
that only the observed galaxies contain 56Fe. If nonvisible matter has under-
gone nucleosynthesis, the spectrum normalization would have to be increased.
One already sees that it cannot be increased very much, and I tentatively
conclude that the density of 56Fe does not exceed the observed density by
more than a factor of two unless the A=56 lines are trapped in their sources.
The fascinating thing about the Apollo-1 5 points of Trombka, Metzger, Arnold,
Matteson, Reedy, and Peterson (1973) is the way they show positive curva-
tures near 400 keV. This suggests a multisource spectrum, and it is quite
conceivable that the radioactivity spectrum may be significant in the overall
shape. Certainly the changes of second derivative will, if they remain after
further experimental scrutiny, be important keys to the origins of this
spectrum. The radioactivity spectrum may be less visible if the exploding
source remains opaque for several months. Compton scattering as extensive
as that in Figure XI.A-2 would remove at least half of the photons from the

280 THEORY

rest frequency at the source and redistribute them at energies of 0.5 MeV or so.
If this source function were employed in Figure XI.A-4, the rest edges would
be smaller by a factor of two or so, and the whole high-energy slope would be
diminished in importance. At present no firm conclusion can be made because
the Nal(Tl) scintillator aboard Apollo-15 had not the energy resolution to
detect structure like that in Figure XI.A-4. Nonetheless, such structure should
be detectable. In Figure XI.A-4, the age of this universe is 1 1 .8 X 109 years,
and nucleosynthesis in galaxies began at t=2 X 109 years, corresponding to
z=2.5. The spectrum has series of rest -frequency edges and red shifted continua.
The rest edges, which are calculated without Compton scattering in the source,
are smallest for nucleosynthesis peaked in the early galactic history. Photons in
the radioactivity background are significant, but higher-energy -resolution
observations will be needed to extract the presence of detailed structure.
Although the density required for %=lA with HQ=55 km/s-Mpc is Pc=5.9 X 10"30
gm-cm"3, this figure assumes that only the galaxies, with density pG=0.028 pc,
contain 56Fe. This calculation (Clayton and Ward unpublished) is thus a
lower limit to the anticipated 7-ray density.

Discussion

It is within scientific grasp to learn the answers to many or all of the questions
about nucleosynthesis enumerated in the Introduction. What is needed is a
7-ray telescope with high-energy resolution, moderately good angular
resolution, and long operation times outside the earth's atmosphere while
reponsive to ground command. Of primary importance is energy resolution
of a few percent or better to extract lines from continua and to detect rest
edges in the universal continuum. Because the rest energies of the 7-rays and
their relative production rates are known from laboratory studies, relatively
sophisticated data analyses can be performed; one could sum the counting
rates just before and just after each rest edge, for example, and compare the
decrement with that at arbitrary energies in the spectrum. T.^e angular
resolution is needed to identify specific radiating objects ^supernovae). As
far as I know, the best type of instrument for accomplishing these two needs
would be similar to the one I described at the NASA X- and 7-Ray Committee
Study of November 1965— a honeycomb of parallel holes drilled through
actively collimating Csl or Nal with solid-state (for example, Li-drifted Ge)
7-ray detectors at the bottom of each hole.

Operation outside the earth's atmosphere is necessary in order to reduce the
emission background of the earth's atmosphere and its opacity. Ground
command will be necessary for viewing different objects and for extracting
the isotropic component. Last, but by no means least, we need nature's
cooperation in presenting us with a new galactic supernova and preferably
a visible one, although an invisible one could be immediately recognized by
a large increase of the A-56 Hnes (see Equations (XI.A-7) and (XI.A-8)).

PROSPECTS FOR NUCLEAR-GAMMA-RA Y ASTRONOMY

281

LJ
T3

Figure XI.A-4. The composite ^NH^Co-^Fe 7-ray spectrum in a specific
Einstein-de Sitter universe. The solid line is an exponentially decreasing rate
of nucleosynthesis/galaxy =e"2 of the initial rate today. The dashed line is a
constant rate of nucleosynthesis/galaxy. For comparison, the dotted line is
the background spectrum observed on the Apollo-15 spacecraft by Trombka
etal., (1973), (heavy solid dots are data points).

Good observation of at least one supernova is needed to measure what fraction
of a 7-line emerges unscattered from their source, for without this calibration,
the interpretation of the universal background will remain insecure; with a little
bit of luck, the entire science of explosive nucleosynthesis will gain a firm

282 THEORY

observational footing from these very special photons. Like all photons, they
tell us that an electromagnetic deexcitation occurred ; unlike any other photons,
they alone tell us that a new nucleus was just born.

INELASTIC SCATTERING

When particles collide with energies in excess of those of their nuclear excited
states, nuclear excitation is possible by the process of inelastic scattering. Let
us not consider here the interesting problems of fast particles impinging on
special dense objects like the earth's atmosphere, the moon, or the sun's outer
atmosphere. Gamma rays from all three sources have been observed, but I
will be concerned with radiation from outside the solar system. I also wish
to set aside 7-rays from the surfaces of stars and collapsed objects, although
both may present some observable sources. By design I will restrict myself
to some remarks concerning the interstellar medium and its interaction with
fast particles— either a general cosmic-ray flux or special regions of high flux
near sources of fast particles. The point to be made at once is that "fast
particles" rather than "cosmic rays" may be a more appropriate nomenclature,
because the largest cross sections and the largest fluxes may be found in the
region of several MeV.

Let F. be the 7-ray flux at the earth within a solid angle 12 due to collisions

A + B -> A* (E.) + B

(XI.A-39)

A*(E.)^A + 7(E.)

The center-of-mass kinetic energy before the collision is E, and after the
collision is E-E.. Let the cross section for this process be designated by o1AB
(E). In practice, one of these particles will be a nearly stationary constituent
of interstellar or circumstellar plasma, and its number density NA (x,t) is a
function of time and place ; the other particle will be regarded as the fast one
with flux 0 (x,t) that is also a function of time and place. Then

=— I r'2NA 0B (E) dAB (

'v(n)

where the integral is over center-of-mass energies E>E. and V(£2) is the volume
of interstellar gas viewed by the solid angle 12. I will suppress the time
dependence, although the arguments are evaluated at t -r/c if the flux is
measured at t. Euclidean geometry is consistent with the assumption that the
only 7-ray lines of this type we are likely to see come from the galaxy. For

PROSPECTS FOR NUCLEAR-GAMMA-RA Y ASTRONOMY 283

an infinitesimal pencil of directions d£2, NA(x) will be constant over the
volume element dV = r2 d!2dr so that the differential flux is

Na(*)*bW

J E. J o

dfi
dF.= —

and if the position dependence of the fast-particle flux is ignored, this
integral becomes a product of the integrated number of particles per unit area
along the line of sight times the integral over the energy of the cross section
times the flux.

One thing to notice immediately is that there should be another term involving
the product NL 0A(E) in the integrand, and, if the chemical composition
of the gas were identical to that of the fast particles at fixed velocity, both
integrals would give the same number of 7-rays. In the second case, however,
the energy of the received 7-rays may be significantly Doppler-shifted if the
particles A were moving at significant fractions of c. (I will not concern my-
self at all with truly relativistic velocities where pion production dominates
7-ray considerations and where, in any case, the fluxes are too small to
produce detectable low-energy 7-rays.) The Doppler broadening in the second
case might make the lines harder to resolve.

Let us make an order-of-magnitude estimate in order to define the ballpark.
Imagine a telescope viewing the galactic disk. Let the solid angle £2 contain
interstellar gas equal to p percent of the galactic mass (the total interstellar
gas being about p =10 percent) in the galactic disk. Let pA and p'B be the
percentages of interstellar particles and of fast particles having identities A
and B respectively, so that NA = (pA/100)N and 0B = (pB/100)tf>. (Through-
out, I have chosen to express unknown parameters in percentages in order
that they have expected values nearer unity in resulting expressions.)
Assuming the emissivity of the disk to be nearly uniform means that the
flux is comparable to the value it would have if the emission within V(£2) were
all from the galactic center, about 10 kpc away. For that case

Fj (galaxy) * pg (pApB + pBpA) 1 .8 X 1011 0 o\B (mb) cm"2 -s1 (XI.A-42)

where 0 (>E.), the total flux above threshold, and a (mb) is the average cross
section in units of millibar ns. For example, if we consider that the 6.1-MeV
radiation from 160 has a proton cross section of about 100 mb above_its
threshold (effectively about 8 MeV), and if the flux above 8 MeV is 0(>8)
h 50 cm"2 s"1 , a not unreasonable extrapolation from observations above 30
MeV, then with p16 = 0.07, pH = 90, pH = 90, p j6 = 0.3 (Cameron, 1968);
Snapiro and Silberberg (1970) give

284 THEORY

P16 (galaxy) « 3 X 10"6 cm'2*"1

if the gas in V (£2) is p = 1 percent of the galactic mass. This corresponds
to a total production of about 1041 7-rays-s"1 from the entire galactic gas, in
rough agreement with the estimate of Fowler, Reeves, and Silk (1970), and
a flux of about 3 X 10"5 cm"2^'1 with an omnidirectional counter. I do not
want to argue this as the best calculation of the emissivity of the galactic disk.
My point is that line fluxes of order 10*5 cm'2^'1 will be expected from 160 ,
and that this detection will be only marginally possible. That is, this prospect
lies near an uncertain edge of detectability.

What basic features of the estimate could be plausibly altered to obtain a
larger 7-ray flux at the earth? One idea would be a discrete source nearby.
However, one readily calculates that if a mass mMQ of gas concentrated at a
distance d(pc) is irradiated, the 7-ray flux at earth is

0.8 X 1012 _
F. (source) = ma (mb) 0 (pApB + pfipA) (XI.A-43)

so that for d = m = 1 , the 160* 7-ray flux, for example, would be F 6 (discrete)
= 3 X 10"9 0(>8 MeV). Thus the flux from a discrete source can hardly be
much greater than that of the disk as a whole, unless the fast particle flux 0
is very much greater (for example 0 > 105 cm'V1 ) than in the general cosmic
radiation. This might occur for a short period following a supernova explosion,
or it might occur for a long period around a rapidly rotating collapsed object.

Another attractive idea is that fast-particle flux 0(E) could be a very steep
function of energy. The solar modulation is thought to be (Goldstein, Fisk,
and Ramaty, 1970) so severe for E < 30 MeV/nucleon that measurements at
earth give little insight into the general interstellar flux. If we should be so
fortunate that 0(E) a E"nwith n>2, a great deal of special information could
be extracted from sources. The high fluxes will give observable counting rates,
and the steep energy dependence will produce informative threshold-dependent
features. We may perhaps even expect this near the sources, because Braddy,
Chan, and Price (1973) have found that big solar flares produce a very steep
spectrum having n ^ 3.7 with preferential acceleration of heavy ions. Of
course solar flares are not the origins of cosmic rays, but let us make do with
what we have and suppose that, like flares, the acceleration mechanisms for
cosmic rays also produce a steep low-energy spectrum. If it is as steep as n=2,
the nuclei having low excitation thresholds can be excited more strongly than
more abundant nuclei having higher excitation thresholds. This point is
illustrated in Figure XI.A-5, which shows the relative abundances of cosmic-ray
nuclei (assuming terrestrial isotope ratios) as a function of the excitation energy

PROSPECTS FOR NUCLEAR-GAMMA-RAY ASTRONOMY

285

102i-

3
UJ

Li?

101

s
s

OBS

i i

1 v
I !

10«

i ill i i ii

■ 0

,24

N14

Sj28

#Be9

^Cr*2

o.i

1.0
E(A*) (MeV)

Figure XI.A-5. Relative abundances of cosmic-ray nuclei (Shapiro and
Silberberg, 1970) plotted as a function of the energy of their first excitation
level. Terrestrial isotope ratios have been assumed. The energy range of the
line feature observed from the galactic-center region (Johnson et al., 1972) is
indicated. No observations have been made above 0.93 MeV (Fishman and
Clayton, 1972).

E. = E(A*) of their first excited states. There is a general positive slope of
approximately E.+1 in these abundances, which reflects only the fact that the
most abundant nuclei tend to be the most stable, and those in turn tend to
have the largest excitation energies.

First consider an example of how 7-ray astronomy could measure the exponent
n in the fast-particle spectrum. The 14N nucleus has excited states at 2.31 and
3.94 MeV with "effective thresholds" of about 3.3 and 4.9 MeV (to allow
the outgoing proton at least 1 MeV to beat the Coulomb barrier). The

286 THEORY

excitation of the 3.95-MeV level results (96 percent ) in a cascade of 1.6-and
2.3-MeV 7-rays, whereas the excitation of the 2.31 MeV level results only in
the 2.3-MeV 7-rays. Thus the relative 7-ray fluxes should be

F , £.3E* °™ CE) dE

-M % + 1 (XI.A-44)

Fl6 JJ9E-na^95(E)dE

I can foresee a need for tabulations of nuclear cross sections of this type;
however, if we only assume for simplicity that the ratio of these two averaged
cross sections is near unity, then the n dependence is proportional to the ratio
of fast particle fluxes:

F,_ /4.9\n"1

-^oc + 1 = 3.8 for n = 3.5

F \3 3/

ri.6 yj'J'

whereas the corresponding ratio would be near two if there were a deficiency
rather than an excess of MeV particles. One could also compare the first
excited-state lines of 12C and 160, but then the abundance ratio in the source
would be an unknown. If many different lines of many different species can
be observed, an interesting picture of the abundances and energy spectrum
could be assembled. A related type of problem was extensively discussed by
Lingenfelter and Ramaty (1967) for the case of solar flares, and many of
their conclusions pass directly to extra-solar 7-ray astronomy. Now that some
of these solar-flare 7-rays have been seen by OSO-7, we may expect further
clarifications on prospects for the future of galactic astronomy. For fast particles
more energetic than 10 MeV, the most prominent astronomical lines should
be the pair-annihilation line, the 2.23-MeV radiative neutron capture by
hydrogen, and the excited states of C, N, 0, and Ne nuclei. Fowler, Reeves,
and Silk (1970) emphasized for these particles, however, that the 7-ray flux
is limited by the requirement that the accompanying spallation reactions not
overproduce Li, Be, and B abundances. They find that the rate of production
of 12C and 160 7-rays is less than 10"26s"1 per interstellar H-atom. However, this
limit is an average over time and place and could be greatly exceeded in limited
regions for limited times.

Figure XI.A-5 shows those nuclei that will be most excitable by low-energy
fast particles, so let us turn our attention to the possibility that large fluxes
of particles with E < 5 MeV may be common. In addition to the solar-flare
observations to motivate this hypothesis, we have the fact that if the HI
regions are heated by fast particles, they must surely lie in the MeV region.
If the fast-particle spectrum is steeper than E"2 , the nuclei with low excited
states are excited more frequently than the more abundant nuclei. Of all
these, 7Li is anomalous in that its cosmic-ray abundance is very much

PROSPECTS FOR NUCLEAR-GAMMA-RA Y ASTRONOMY

287

greater than the general line through Figure XI.A-5. In a steep fast -particle
spectrum, the 478-keV line of 7Li* should be the most prominent if the
fast particles have the relative abundances of the cosmic rays. The peculiar
fact was used by Fishman and Clayton (1972) in their attempt to account
for the spectral feature observed toward the galactic center by Johnson,
Harnden, and Haymes (1972). They point out that a 432-keV 7-ray due to
7Li (p,n) 7Be* (432) should accompany the main 7Li radiation with about
one-third its intensity. Their fit to the data of Johnson et al. (1972), is shown
in Figure XI.A-6. The computed solid curve is a simple power-law continuum
plus the 7Li doublet feature at an intensity comparable to the observed
counting rate. The fit is basically quite good, so the explanation could be
correct. Fishman and Clayton (1972) p = 2 percent and p ' = 0.08 percent

102

10°

50 60 70

CHANNEL NUMBER

Figure XI.A-6. The curve shows the shape of the 7 Li-inelastic-scattering
feature superimposed on a smooth power-law continuum (Fishman and
Clayton, 1972). The profile was computed for an energy resolution equal to
that of the detector used by Johnson et al. (1972). Because of the limited
resolution, the line at 432 keV due to 7Li (p, n) 7Be* is not physically
separated from the line at 478 keV due to 7Li (p, p') 7Li*, which is three
times stronger. The histograms are the data of Johnson et al. (1972), with
their energy channels summed in groups of three adjacent channels to reduce
statistical fluctuations. The consistency of their feature with the one proposed
by Fishman and Clayton (1972) is evident.

288 THEORY

in Equation (XI.A-42), to conclude that if the radiation was from this much
gas toward the galactic center, one gets 0 « 5 X 104 cm"2-s_1 between about
2 and 10 MeV/nucleon. This is a much larger flux than one is accustomed to
think of in cosmic rays, but it is not out of line with E"3'5 spectra like those
of a solar flare; with n = 3.5 one has 0(>2) « 900 0 (>30). The large-energy
density of over 100 eV cm"3 would create dynamic instabilities were it a
galactic-wide phenomenon, however, so it must exist instead in bottles of high-
flux regions. Nonetheless, our calculation requires the amount of irradiated
gas to be large enough to obtain the observed flux, so there are strong astrophy-
sical problems here. Another problem is that the high 7Li abundance is usually
assumed to be spallogenic from high-energy cosmic rays, so that large low-
energy fluxes of this nucleus might not be expected from that point of view!
My philosophy is that observational 7-ray astronomy is quite capable of
teaching us the truth in these matters, so elaborate models for or against this
particular explanation may not be appropriate at present. We also need much
better evidence of the cosmic-ray flux at the solar system, because the
Goldstein, Fisk, and Ramaty (1970) calculations show that the particles at
earth's orbit having 30 MeV/nucleon are three to four orders of magnitude
less abundant than their 50 to 100-MeV progenitors at the boundaries of the
solar system. Perhaps Pioneer- 10 will give us badly needed facts on this
modulation problem.

(Partially supported by the National Science Foundation GP-18335)

REFERENCES

Arnett, W. D., and D. D. Clayton, 1970, Nature, 227, p. 780.
Arnett, W. D., \91\, Astrophys. J., 166, p. 153.

Bodansky, D., D. D. Clayton, and W. A. Fowler, 1968, Astrophys. J. SuppL,
16, p. 299.

Braddy, D., J. Chan, and P. B. Price, 1973, Phys. Rev. Letters 30, p. L669

Brown, R. T., 1973, Astrophys. J., 179, p. 607.

Cameron, A. G. W., 1968, Origin and Distribution of the Elements, L. H. Ahrens,
ed., New York, Pergamon Press, p. 125.

Clayton, D. D., 1971, Nature, 234, p. 291.

, 1973, Science, in press.

Clayton, D. D., and W. Craddock, 1965, Astrophys. J., 142, p. 189.

Clayton, D. D., S. Colgate, and G. J. Fishman, 1969, Astrophys. J., 155, p. 75.

Clayton, D. D., and J. Silk, 1969, Astrophys. J., 158, p. L43.

PROSPECTS FOR NUCLEAR-GAMMA-RA Y ASTRONOMY 289

Clayton, D. D., and S. E. Woosley, 1969, Astrophys. J., 157, 1381.

Fishman, G. J., and D. D. Clayton, 1972, Astrophys. J., 178, p. 337.

Fowler, W. A. 1972, Cosmology, Fusion, and Other Matters, F. Reines, ed.,
Univ. of Colorado Press, Boulder, p. 67.

Fowler, W. A., H. Reeves, and J. Silk, 1970, Astrophys. J., 162, p. 49.

Goldstein, M. L., L. A. Fisk, and R. Ramaty, 1970, Phys. Rev. Letters, 25,
p. L832.

Howard, W. M., W. D. Arnett, D. D. Clayton, and S. E. Woosley, 1971 ,
Phys. Rev. Letters, 27, p. LI 607.

, 1972, Astrophys. J., 175, p. 201.

Johnson, III, W. N., F. R. Harnden, and R. C. Haymes, 1972, Astrophys. J.,
Letters, 112, p. LI.

Lingenfelter, R. E., and R. Ramaty, 1961, High Energy Nuclear Reactions in
Astrophysics, B. Shen, ed., New York, W. A. Benjamin, Inc.

McVittie, G. C, 1965, General Relativity and Cosmology, Urbana, Univ.
Illinois Press.

Oort, J. H., 1958, Solvay Conference on Structure and Evolution of the
Universe, Brussels, R. Stoops.

Sandage, A., 1912, Astrophys. J., 178, p. 1.

Schmidt, M., 1965, Galactic Structure, A. Blauw and M. Schmidt, eds.
Chicago, University of Chicago Press, p. 528.

Shapiro, M. M., and R. Silberberg, \91Q,Ann. Rev. of Nuclear Science, 20,
p. 323.

Trombka, J. I., A. E. Metzger, J. R. Arnold, J. L. Matteson, R. C. Reedy, and
L. E. Peterson, 1913, Astrophys. J., 181, p. 737-746.

Weinberg, S. 1972, Gravitation and Cosmology, New York, John Wiley, p. 495.

B. POSITRONIUM FORMATION RED SHIFT
OF THE 511-keV ANNIHILATION LINE

M. Leventhal*

Bell Laboratories

Energetic positrons stopping in a gas, in principle, can annihilate from either
the free state with an electron bound in a gas atom or first capture an electron
and form the hydrogen-like positronium atom before annihilating. Stecker
(1969) and Ramaty, Stecker, and Misra (1970) have made the point that
positronium formation may be of importance in astrophysics. The cross sec-
tions for positron energy loss and annihilation processes in gases are such that
annihilation does not take place until the positron has slowed down to eV-
type energies (Deutsch, 1953). The cross section aF for free annihilation with
the emission of two antiparallel 5 1 1-keV 7-rays (Dirac, 1930) is many orders
of magnitude smaller than the positronium formation cross section ap
(Massey and Mohr, 1954) and one might think at first that it can be neglected.
For 20-eV positrons incident on atomic hydrogen ap « 10"23 cm2 and
ap « 10"16 cm2. However, in dense gases, this is not the case because posi-
tronium formed with kinetic energy greater than its binding energy of 6.8 eV
can ionize in a collision with a gas atom before annihilating. Thus the positron
may eventually be slowed down to energies below the threshold ET = (I-6.8)eV
for positronium formation, where I is the first ionization potential of the
stopping gas. For positron energies below ET , only free annihilation can occur.
The relative cross sections for the various elastic, inelastic, and annihilation
processes involved are such that, for typical gases at atmospheric pressure, the
fraction of positrons annihilating from the positronium state (f) is in the
range 20 to 50 percent. Numerous laboratory experiments have demonstrated
this large positronium formation fraction (Green and Lee, 1964). As the
pressure is reduced to the extremely dilute situation found in many astro-
physical situations, f may approach 100 percent. An estimate of the atomic
density (N) below which f would approach 100 percent can be obtained

*Speaker.

291

292 THEORY

from the expression

N=-

where r is the time required for a positronium ionizing collision, a. is the
cross section for positronium ionization in a gas collision, and v is the posi-
tronium velocity. Ifr> 1.4 X 10"7 s, the ortho-positronium annihilation
lifetime (see below), free annihilation should be suppressed. Using the
characteristic values of o} = 7 X 10"17 cm2 and v = 2 X 108 cm • s calculated
for positronium incident upon atomic hydrogen (Massey and Mohr, 1954)
N is found to be of order 101 5 atoms • cm"3. Since densities as large as
1015 atoms • cm"3 are found only in the vicinity of condensed objects, a
predominantly positronium annihilation spectrum is expected in many astro-
physical situations. (The possible presence of intense UV fields, which might
photoionize the positronium within its annihilation lifetime, has been ignored
in the above discussion. Also, it has been assumed that the temperature
T ^ 10s K°, that is, the hydrogen plasma, is predominantly neutral atoms.)

Positronium may be formed with its spins aligned antiparallel, the singlet-para
state with spin = 0, or with the spins parallel in the triplet-ortho state with
spin = 1 . On a purely statistical basis, the ortho state is expected to be formed
three times as often as the singlet state because of the threefold degeneracy in
the azimuthal spin quantum number M = 1, 0, -1. Singlet positronium
annihilates in 1.3 X 10"10 s yielding two antiparallel 51 1-keV 7-rays. Triplet
positronium annihilates in 1 .4 X 10"7 s yielding three 7-rays in a plane. The
probability per unit energy range PT(E) of finding a triplet 7-ray with a
particular energy in the interval from 0 to 5 1 1 ke V is given by (Ore and
Powell, 1949)

2 I E(m-E) 2m(m-E) / m-E
PT(E)= -^ '- + — K- In'

(772 - 9) ( (2m-E)2 E2 \ m

2m(m-E)2 /m"E\ /2m-E>
(2m-E)3 \ m / \ E

where m is the electron rest mass energy of 51 1 keV, and the distribution has
been normalized to unity. A plot of PT(E) is given in Figure XI. B-l. The
positronium annihilation spectrum Pp(E) can then be constructed by adding
the singlet and triplet spectra, Pp(E) = 0.75 X 3 X PT(E) + 0.25 X 2 X Pg(E)
where the weightings reflect the relative amounts of triplet and singlet formed
and the number of 7-rays each gives. In principle, the singlet spectrum Pg(E)

POSITRONIUM FORMA TION

293

DO
OH
<

600

0 200 400 600
PHOTON ENERGY ( keV

CD
<

25 kev

RESOLUTION

0 200 400 600
PHOTON ENERGY ( keV )

Figure XI.B-1. Plots of the detected positronium annihilation function
Pp(E) for various energy resolutions (FWHM) of the 7-ray detector. A
Gaussian width of 3.2 keV has been assumed for the singlet annihilations. The
broken curve in the upper left-hand box is a plot of the pure, triplet
annihilation function PT(E).

should be a normalized delta function centered at 51 1 keV. However, the
finite motion of the positronium atom will give a Doppler broadening in all
real cases. Pp(E) can in principle be modified by charge exchange collisions
with gas atoms, converting triplet to singlet positronium before annihilation.
The cross section for this process (Massey and Mohr, 1954) is of the same
order of magnitude as a.. Hence again, for densities of order 101 5 atoms • cm"'
or smaller, an unmodified Pp(E) is expected.

Finally, to obtain the observed positronium spectrum P'p(E), an instrumental
rounding of Pp(E) must be performed because most 7-ray detectors in the
energy range of interest have poor energy resolution. Assuming a Gaussian
resolution function one obtains

(E-E')2

P'P(E) =

Jo °^

Pp(E')dE'

294

THEORY

where the full width at half maximum of the convoluted function is 2.35a.
The function P'p(E) has been numerically evaluated on a computer. Plots of
P'p(E) for various width resolution functions are given in Figure XI.B-1. PS(E')
has been taken as a normalized Gaussian of width 3.2 keV corresponding to
the Doppler broadening for 20-eV positronium. However, it is important to
point out that the spectra are insensitive to the Doppler width of Pg(E)
because the resolution functions are in general much broader. Clearly, posi-
tronium annihilation yields a spectrum red-shifted from 5 1 1 keV and of
asymmetric shape. The apparent red shift of the peak versus resolution is
plotted in Figure XI.B-2.

UJ
Q.

Q
UJ

cr.

z
o

Q-
<

100 120 140 160

RESOLUTION OF DETECTOR ( k eV )

Figure XI.B-2. A plot of the apparent peak position of the detected
positronium annihiU
resolution (FWHM).

positronium annihilation function Pp(E) as a function of detector

A feeble 476 ± 24-keV feature, now 5.3 standard deviations above the 7-ray
continuum from the galactic center region, has recently been detected
(Johnson, Harnden, and Haymes, 1972; Johnson and Haymes, 1973). Two
interpretations of the feature have been presented: the first, that it is the
51 1-keV positron annihilation line red-shifted by the gravitational potential

POSITR ONIUM FORMA TION 295

at the surface of a neutron star (Borner, Cohen, and Ramaty, 1972; Guthrie
and Tademaru, 1973); and the second interpretation is that the feature is a
nuclear 7-ray emitted when the 7Li nuclei in the cosmic rays scatter inelastically
from the interstellar gas (Fishman and Clayton, 1972). Since both proposals
are far from conclusive, we think it worthwhile to consider the possibility
that positronium annihilation is being detected.

The feature from the galactic center region was detected with a fitted energy
resolution of 86 keV (Johnson and Haymes, 1973). According to Figure
XI.B-2, a positronium annihilation spectrum would have appeared peaked at
490 keV, which is consistent with the observation. Further support for the
positronium hypothesis would be provided by detecting the characteristic
asymmetric shape of the feature. Unfortunately because of (1) the poor
statistical quality of the available data, (2) the fact that the feature sits on a
large, sloping background continuum, and (3) the large energy resolution of
the detectors employed, conclusions about the shape of the feature would at
present be dubious. Additional data acquisition with 7-ray detectors of higher
resolution should eventually answer this question.

Additional support for the hypothesis might come from the detection of
positronium line radiation (Leventhal, 1970). Implicit in the above discussion
was that annihilation took place from the ground atomic state of positronium.
Some fraction of the positronium should be formed in excited states (Massey
and Mohr, 1954) which optically decay before annihilating. The expected
spectrum is identical to that of atomic hydrogen with all wavelengths multi-
plied by two because of the reduced mass factor. However, interstellar
extinction in the plane of the galaxy will greatly hinder such measurements
(Becklin and Neugebauer, 1968).

If positronium annihilation is indeed being observed, it is interesting to
speculate on the origin of the positrons. Since the 7-ray telescope employed
by Johnson et al. (1972) had an acceptance cone of 24°, a large fraction of
the galactic disk was observed. Positrons produced as secondaries in the cos-
mic rays should be stopping and forming positronium in the galactic gas. An
order of magnitude calculation of the flux at earth due to cosmic-ray posi-
tronium has been made. Assuming (1) the cosmic-ray positron flux as meas-
ured at earth (Fanselow, Hartman, Hildebrand, and Meyer, 1969) is uniform
throughout the galaxy, (2) that these positrons are contained within the
galactic gas until they stop, and (3) the galactic gas consists of a disk 20 kpc
in diameter by 0.25 kpc wide containing one hydrogen atom • cm" , a flux at
earth is found which is several orders of magnitudes smaller than the observed
flux of 1 .8 ± 0.5 X 10"3 photons cm"2 • s"1 . Hence we conclude that the posi-
trons must be generated by some other mechanism. If the source were at the
galactic center, D = 10 kpc, 7.3 X 1042 positronium annihilations per second
would be required, yielding a source luminosity of ~1037 ergs • s"1 in the
annihilation radiation alone.

296 THEORY

ACKNOWLEDGMENTS

I wish to thank my Bell Labs colleagues, K. Jefferts, L. Lanzerotti, S. McCall,
P. M. Platzman, J. Tyson, and R. Slusher for their enthusiastic and construc-
tive comments, R. Fulton for computer assistance, W. Johnson and F. Harnden
for helpful discussions of their experiment, M. Ruderman for bringing the
problem to my attention, and G. Steigman for helpful comments.

REFERENCES

Becklin, E. E., and G. Neugebauer, 1968, Astrophys. /., 151, p. 145.

Borner, G., J. M. Cohen, and R. Ramaty, 1972, Bull. American Asst. Soc. ,
4, p. 410.

Deutsch, M., 1953, Prog. Nucl. Phys. , 3, p. 131.

Dirac, P. A. M., 1930, Proc. Camb. Phil. Soc. , 26, p. 361.

Fanselow, J. L., R. C. Hartman, R. H. Hildebrand, and P. Meyer, 1969,
Astrophys. J. , 158, p. 771.

Fishman, G. J., and D. D. Clayton, 1972, Astrophys. J. , 178, p. 337.

Green, J., and J. Lee, 1964, Positronium Chemistry, Academic Press, New
York and London.

Guthrie, P., and E. Tademaru, 1973, Nature Phys. Sci., 241, p. 77.

Johnson, W. N., F. R. Harnden, and R. C. Haymes, 197 '2, Astrophys. J.,
172, p. LI.

Johnson, W. N., and R. C. Haymes, 197 3, Astrophys. J. , in press.

Leventhal, M., 1970, Proc. Nat. Acd. Sci , 66, p. 6.

Massey, H. S. W., and C. B. O. Mohr, 1954, Proc. Phil. Soc. London, A67,
p. 695.

Ore, A., and J. L.Powell, \949, Phys. Rev., 75, p. 1696.

Ramaty, R., F. W. Stecker, and D. Misra, 1970,/. Geophys. Res., 15, p. 1 141.

Stecker, F. W., 1969, Astrophys. Space Sci. , 3, p. 579.

C. NUCLEAR GAMMA RAYS
FROM SOLAR FLARES

R. Ramaty*

Goddard Space Flight Center

INTRODUCTION

Solar 7-ray line emissions at 0.5, 2.2, 4.4, and 6.1 MeV were detected during
the flare of August 4, 1972, by a 7-ray monitor flown on OSO-7 (Chupp et al.,
1973). Line emissions at 0.5 and 2.2 MeV were also detected on August 7,
1972, but only upper limits could be set on the 4.4- and 6.1-MeV lines from
this flare. In previous papers (Lingenfelter and Ramaty, 1967; Cheng, 1972),
the theory of nuclear reactions in solar flares was treated in detail and predic-
tions were made as to the expected fluxes of 7-rays and high-energy neutrons
at earth from such reactions at the sun. Following the discovery of Chupp
et al. (1973) we have reviewed and updated these calculations including more
recent nuclear cross sections. The results of these calculations were also pre-
sented by Ramaty and Lingenfelter (1973a).

By comparing the predicted emissions with the observations, we can show that
the observed lines at 0.5, 2.2, 4.4, and 6.1 MeV are produced, respectively, by
positron annihilation, deuterium deexcitation following neutron capture on
hydrogen, and the deexcitation of the first nuclear levels of C12 and O16.
Furthermore, from the comparison of the calculated and observed line inten-
sities we can deduce the spectrum of the accelerated particles at the sun inde-
pendent of the assumed interaction model. The total number of accelerated
particles required to produce the observed line emission, however, does depend
on this model. In the subsequent treatment we shall use two limiting models:
a thick-target model in which the accelerated particles move from the flare
region downward into the sun, undergoing nuclear interactions as they slow
down in the solar atmosphere; and a thin-target model in which the spectrum
of accelerated particles is not modified during the time in which the nuclear
interactions take place. The latter model assumes that either the total path
length traversed by the particles at the sun is small in comparison with their

* Speaker.

297

298 THEORY

interaction length, or that the particle energy loss from ionization and nuclear
interactions is just balanced by energy gains from acceleration.

We have recalculated the fluxes of various 7-ray lines expected from accelerated
particle interactions in solar flares: at 0.51 1 MeV from positron annihilation,
at 2.23 MeV from neutron capture on hydrogen, and at 1.63, 1.99, 2.31, 4.43,
5.5, and 6.14 MeV from deexcitation of nuclear levels in C, 0, N, and Ne.
These calculations are based on an ambient solar composition given by
H:He:C:N:0= hlO'^JX 10"4: 10"4:9.2X 1CT4 (Cameron, 1967) and an
accelerated particle population consisting of protons with spectrum

N(P) = P0-» exp (-P/P0). (XI.C-1)

Here P is rigidity, P is a characteristic rigidity which we treat as a free parame-
ter, and N(P) is the total number of protons per unit rigidity. In the calcula-
tions, N(P) is normalized at the sun to one particle of rigidity greater than
zero.

For the thick-target model, the yield of secondaries from a particular type of
interaction is given by

Qs = T7 I dPN(P) J dx' a(P') exp (-(x-x')/L), (XI.C-2)

where t? is the number of target atoms per gram of solar material, a is the
cross section as a function of rigidity, and x and L are the stopping range and
nuclear interaction length of protons of rigidity P (both measured in g • cm'2).
This yield has the units of secondaries per incident proton of rigidity greater
than zero.

For the thin-target model, the instantaneous production rate of secondaries
from the same type of interaction is

n c I

J c

q =nc dPN(P)0a(P) (XI.C-3)

where n is the number density of the ambient solar material in cm"3, and c |3
is the velocity of a proton of rigidity (P). The units of q are secondaries per
second per proton of rigidity greater than zero.

NUCLEAR GAMMA RA YS FROM SOLAR FLARES

299

PHOTON PRODUCTION

Line emission at 0.51 1 MeV is produced from the annihilation of positrons.
The principal source of positrons in solar flares are nuclear reactions of
accelerated particles with the ambient solar atmosphere. These reactions pro-
duce 7r+-mesons and a variety of radioactive isotopes which decay by positron
emission. The main positron emitters, their formation reactions, threshold
energies, half-lives and maximum positron energies are given in Table XI.C-1.
Except for the reactions N14 (p, n) 014 (for example, Andouze et al., 1967)
and N14 (p, d) Cn (Jacobs et al., 1972), the cross sections for these reactions
were given in Lingenfelter and Ramaty (1967). The resultant positron yields
are given in Figures XI.C-1 and XI.C-2 for the thick- and thin-target models,
respectively.

IO"2p

10 =

DC -4

LJ 10

»-

O
0.

10 -

1 1 1

y/ 5

N(P>0) = l

_+/

/c^^

sss552 -=

k£ ^^

«-****

jC^^

'AT

IIUi 1

i i i

50 100 150 200 250 300
P0(MV)

Figure XI.C-1. Yield of positron emitters at the sun
for the thick-target model.

300

k-l

2 >

■*

«~j

■<*

l-H

r-H

•— 1

-C

'g'

.2 o

^~v

'cT

3 O

c£

O-

"^-^

'•— '

v_^

<»

■*

v©

s-»

PL,

<4-H ^^

Cm '2

"c3 c

a: w

C/i

mum

Ener

(Me

1— 4

1— <

"-5

<"H

o
Q

+
+

"ca

CO.

ca.

G
cd

„

l/l

00
t

z
t

1/)

CQ.

NUCLEAR GAMMA RA YS FROM SOLAR FLARES

301

T
O
Id
CO

q:
uj

H
UJ

i i

n=l

1 1

io"7

N(P>0)=l

io-

^j!L-

IO"'9

=■

io20

io2'

1^-"

1 II l

1 1

i i

0 50 100 150 200 250 300
P0(MV)

Figure XI.C-2. Production rate of positron emitters
at the sun for the thin-target model.

The intensity of the 0.51 1-MeV line depends on the number of positrons that
annihilate at the sun. For the thick-target model it is reasonable to assume
that all the flare-produced positrons will ultimately annihilate at the sun.
But, for the thin-target model, it is possible that a significant fraction of the
positrons will escape from the sun before they annihilate. The determination
of this fraction, however, is beyond the scope of the present paper and we
defer the discussion on this for future research. For simplicity, in the subse-
quent calculations we assume that all flare -produced positrons annihilate at
the sun, but keep in mind that this assumption may lead to too many 0.51 1
photons in the case of the thin-target model.

The annihilation radiation yield at 0.51 1 MeV also depends on the mode of
positron annihilation. Positrons annihilate either directly with a free electron
or in a bound state of positronium. In the latter case, the annihilations

302 THEORY

proceed from a 1 S state leading to two 0.5 1 1-MeV photons, or from a 3S
state leading to a 3-photon continuum. Because the 3-photon continuum does
not contribute to line emission at 0.51 1 keV, and because the probability of
forming the 3S state is three times that of forming the 1S state, positronium
formation and its annihilation produces on the average only one-half of a
0.51 1 photon per positron as compared to two photons for free annihilation.
This point was discussed in detail by Stecker (1969) for positron annihilation
in interstellar space. However, in the much higher density of the solar atmo-
sphere we expect that only a small fraction of the positrons will annihilate
from the 3S state of positronium. This is expected because the collision
frequency of positronium with ambient electrons can be sufficiently high to
either dissociate the positronium or to cause a transition to the l S state
before the annihilation of the 3S state. Also, the rate of positronium forma-
tion in a plasma is greatly reduced in comparison with positronium formation
in a neutral medium. Thus, if the flare region is highly ionized, free annihi-
lation will dominate over positronium annihilation (Ramaty and Lingenfelter,
1973b).

By allowing two photons per positron at the sun, the time-integrated 0.51 1-
MeV line intensity at earth for the thick-target model is

F(0.51 1) = — [Qs (7r+) + Qs (Cn) + Qs (N13) + Qs (014) + Qs (O15)]
2ttR2

(XI.C-4)

where the various Q 's are given in Figure XI.C-1. The units of F are photons
cm"2 per proton of rigidity greater than zero. For the thin-target model we
obtain an equation similar to Equation (XI.C-4), with F and Q replaced by
the instantaneous rates 0 and q . The units of 0 are photons cm"2 • s"1 per
proton of rigidity greater than zero. The resultant 0.51 1-MeV photon fluxes
are given in Figures XI.C-3 and XI.C-4, for the thick- and thin-target models,
respectively.

Next we consider the 2.23-MeV line resulting from the deexcitation of
deuterium following neutron capture on hydrogen. The principal neutron-
producing reactions and their cross sections were discussed by Lingenfelter
and Ramaty (1967, and references therein). To these reactions we have
added the reaction N14 (p, n) 014 mentioned previously which is important
for low-energy protons (< 1 5 MeV). The neutron-producing reactions are
summarized in Table XI.C-2. The total neutron yields for the thick- and thin-
target models are then obtained by using the appropriate cross sections in

NUCLEAR GAMMA RA YS FROM SOLAR FLARES

303

z
o

I-
o

50 100 150 200 250 300
P0(MV)

Figure XI.C-3. Time-integrated photon fluxes at earth in the
thick-target model.

Equations (XI.C-2) and (XI.C-3). Having obtained the neutron yields, we
have to consider the propagation of neutrons in the solar atmosphere
because 2.23-MeV photons are not produced if the neutrons escape from
the sun or decay in the solar atmosphere.

Neutron propagation in the solar atmosphere is determined predominantly
by scattering (elastic and inelastic) and capture by protons. The total
neutron-proton scattering cross section a is essentially constant at 20 barns
for neutron energies from 1 eV to 10s eV and then drops to about 0.1
barn at 100 MeV (Hughes and Schwartz, 1958). The capture cross section
is inversely proportional to velocity and is given by oc = 2.2 X 10"6 (F ,
where o is measured in barns and c/3 is the velocity of the neutron. Since
for all energies of interest o » a , the neutron mean-free path in the solar

304

THEORY

100 150 200 250 300
P0(MV)

Figure XI.C-4. Instantaneous photon fluxes at earth in the
thin-target model.

atmosphere is determined principally by neutron-proton scattering. Using
the columnar density of the solar atmosphere as given by Allen (1963), we
find that a 10-MeV neutron moving radially outward from the sun will
probably escape from the solar atmosphere before making even one collision
if it is produced in a region of density less than 1017 cm"3. This density
corresponds to a depth of about 300 km below the base of the chromosphere.
Because solar flares occur in the chromosphere or corona, it is reasonable to
assume that for both interaction models all upward-moving neutrons are going
to escape from the sun. The downward-moving neutrons on the other hand
decay or are captured, depending on whether the capture time is greater or
smaller than the half-life of the neutrons. From the capture cross section
(ac), the capture time (t ) is given by

t = (c0n<x >

l ~

1.5 X 1019
n(cm"3)

(XI.C-5)

NUCLEAR GAMMA RA YS FROM SOLAR FLARES

305

Table XI.C-2
Neutron-Producing Reactions

Reaction

Threshold Energy
MeV/Nucleon

H1 (p, mr+) H1

292.3

He4 (p, pn) He3

25.9

He4 (p, 2pn) H2

32.8

He4(p, 2p2n)H1

35.6

C12(p,n...

19.8

N14(p,n...

6.3

016(p,pn...

16.5

Ne20 (p, pn . . .

17.7

As can be seen, t is independent of energy and is inversely proportional to
the density (n) of the region where the neutrons interact with the ambient
medium. From the previous discussion it follows that downward-directed
neutrons probably will not collide with the ambient gas until they reach the
photosphere where n is of the order 1017 cm"3. Thus t = 1 50 s, and since
the half-life of the neutron is 720 s, it follows that most of the neutrons will
be captured before they decay. Furthermore, since the neutron capture occurs
at a columnar depth of about 1024 cm"2 or 1.6 g • cm"2 while the stopping
range of a 2.2-MeV photon is about 25 g • cm"2, all the upward-moving photons
resulting from deuterium deexcitation will escape from the sun. Assuming
isotropic production of neutrons in the interaction region, the 2.23-MeV line
intensity at earth for the thick- and thin-target models, respectively, are

F(2.23) = Qs (n)

4ttR2 2

(XI.C-6)

0(2.23) = qs (n)

4vrR2 2

(XI.C-7)

where Q (n) and q (n) are the total neutron yields as calculated from
Equations (XI.C-2) and (XI.C-3).

306 THEORY

The calculations of the other intensities at 1.63, 1.99,2.31,4.43, 5.5, and
6.14 MeV in Figures XI.C-3 and XI.C4are straightforward. Unlike the 0.51 1-
and 2. 23 -MeV lines, these emissions are prompt, that is, the excited states or
secondary products decay by photon emission in a time scale much shorter
than any characteristic time of the flare process.

Line emission at 4.43 and 6.14 MeV results from the deexcitation of the first
nuclear levels of C12 and O16, respectively. The intensity of the 6.14-MeV
line is the same as calculated by Lingenfelter and Ramaty (1967). The inten-
sity of the 4. 43 -MeV line is about 50 percent greater than in Lingenfelter and
Ramaty (1967), because it is possible to produce C12(443) by the spallation
of O16, a process which was neglected in that paper. The calculations for this
process, given in Figures XI.C-3 and XI.C-4, are based on cross sections meas-
ured by Zobel et al. (1968).

Radiation at 2.31 MeV corresponds to the first nuclear level of N14. The
2.31 -MeV line is produced by the direct excitation of the first and second
levels of N14. The latter is at 3.94 MeV and it deexcites 96 percent of the
time through the first level, thereby producing a 1 .63 -MeV photon in addition
to the 2.31 -MeV photon. In addition to direct excitation, the first level of
N14 can also be populated by the decay of O14, which is produced by the
reaction N14 (p, n)014. O14 0 decays 99.4 percent of the time to N14(2-31),
and hence each j3-decay is accompanied by a 2.31 -MeV photon. As mentioned
previously, the 1 .63-MeV line is formed by the deexcitation of N14(3-94). In
addition, this line is also produced by the deexcitation of the first level of
Ne20. Finally the 1 .99-MeV line results from the deexcitation of C11 which
is formed from the spallation of C12 . The cross section for this process is
given by Zobel et al. (1968).

The principal reason for showing the 1 .99- and 2.31 -MeV lines is that in
studies with detectors of poor energy resolution, these lines could, in principle,
be confused with the deuterium deexcitation line. As can be seen from
Figures XI.C-3 and XI.C4, however, the C11 (199) -MeV line is negligible in
comparison with the deuterium line. On the other hand, the 2.31 -MeV line
from N14 could compete with the 2.23-MeV line at low P 's. If all the
emission at ~2.2 MeV is from N14, however, the intensities of the 4.43-MeV
and 6.14-MeV lines should be about the same as that of the 2.2-MeV line. As
will be discussed below, this was not the case for the August 4, 1972 event.
Therefore, for this flare at least, we conclude that the 2. 2 -MeV line is pro-
duced almost entirely by neutron capture.

We have summarized in Table XI.C-3 the principal mechanism leading to line
emission in the solar atmosphere. In addition to the lines given in Figures
XI.C-3 and XI .C4, we have listed in Table XI.C-3 lines at 5.2 MeV from O15
and N15 and at 7.12 MeV from O16. From the cross sections of Zobel et al.

NUCLEAR GAMMA RA YS FROM SOLAR FLARES

307

<~>

G-

T3
O

o'
<s

<L>

z— v

<3-

ft
o

"Dh

z—v

^~^

Z—N

^-^

z—v

z- V

/"~ v

z— N

i—

— .

~Oh

s— '

"—■^

N~-/

Nw/

s— '

*** '

s— ^

<fr

«*

*£>

—

+->

<L>

c
o

ti-
ed

c
o

■*-»

'.£

H—»

'■+-'

+-I

•43

+■>

X
o

■*->

h— »

+->

-t-"

■8

3
C

-o

•a

•n

■sT

^_J

•-^

1—

©^

£,

iri

°lo

IT)

irt

<L>

S-

1-4

C jT^

u >

r-l

hoton
(Me

*— *

1— 1

v£>

'-H

•-H

IT)

308

THEORY

(1968), the intensities of both these lines should be approximately 50 percent
of the 6.14-MeV line intensity.

DISCUSSION

Let us now compare the results of our calculations with the observations of
Chupp et al. (1973) for the August 4, 1972, flare. We defer the discussion
on the August 7, 1972, observations, since the 7-rays from this flare were
observed only after the flare maximum and hence they require a more detailed
treatment of the time dependence of the 7-ray intensities.

In Figures XI.C-5 and XI .C -6, the shaded areas represent the time-averaged
ratios of the measured 0.5-, 4.4-, and 6.1-MeV line intensities to the measured
2.2-MeV line. The curves represent the calculated ratios as functions of the
characteristic rigidity (PQ), for the thick-target model in Figure XI.C-5 and
the thin-target model in Figure XI.C-6. As can be seen, the calculated 4.43
and 6.14 curves are strong functions of PQ.

50 100 150 200 250 300
P0(MV)

Figure XI.C-5. Relative line intensities for the thick-
target model.

NUCLEAR GAMMA RA YS FROM SOLAR FLARES

309

5 10 -

Figure XI.C-6.

150 200
P0(MV)

Relative line intensities for the thin-
target model.

Therefore, the comparison of these ratios with the measurements allows
us to deduce the value of the characteristic rigidity (P ). We find that for
both models PQ has to be in the range 70 to 80 MV. This range of P's
should be compared with values of P as obtained from charged-particle
observations near earth.

According to Bostrom et al. (1972), the peak proton intensities after the flare
of August 4, 1972, were: j(>10MeV)= 106 cm"2 s_1,j(> 30 MeV) =
2.6 X 10s cm-2 s'1 , and j (> 60 MeV) = 8 X 104 cm'2 s'1 . From these inten-
sities we can deduce the local proton density (u) in the interplanetary medium.
If u(P) °c exp (-P/PQ), and if the protons are nonrelativistic

u(>P) = j(>P)(mc/e)(P + P0)

-l

(XI .C -8)

From Equation (XI.C-8) and the integral intensities j (> P) given above, we
can calculate values of u(> P) for various P's. For P = 6.5 MV we obtain

310

THEORY

proton densities as follows:

u(>10MeV) = 1.53X 10"4, u(> 30 MeV) = 2.67 X 10'5,and
u(>60MeV) = 6.13X 10"6.

These numbers are plotted in Figure XI .C -7 as functions of P. The straight
lines are exponentials in rigidity with PQ = 60 MV and PQ = 70 MV, and they
bracket the observed proton densities. The charged-particle observations at
the peak of the proton event are therefore consistent with a characteristic
rigidity of about 60 to 70 MV. Furthermore, this characteristic rigidity
remains essentially the same for the remainder of the particle event (J. King,
private communication).

-s

100 200 300 400

P(MV)

Figure XI.C-7. Proton densities in the interplanetary
medium from local particle measurements.

NUCLEAR GAMMA RA YS FR OM SOLAR FLARES 31 1

The similarity between the proton spectrum as observed near earth and the
proton spectrum at the sun as deduced from the 7-ray observations, seems to
imply that, at least for the August 4, 1972, flare, both the escape of particles
from the flare region and their propagation in the interplanetary medium are
essentially independent of energy.

Let us now calculate the total number of protons at the sun. In the case of
the thick -target model this calculation is independent of the ambient density
(n) but requires the knowledge of the time-integrated photon flux from the
flare. According to Chupp et al. (1973), the average intensity of the 2.2-MeV
line from 0616 UT to 0632 UT on August 4, 1972, was 0.22 photons cm"2 s"1 ,
and thus the total photon flux in the 6-min time interval was 80 photons cm"2 .
This is only a lower limit because OSO-7 went into earth eclipse at 0632 UT
before the termination of the 7-ray event. The total 7-ray flux, however, was
probably not much larger, since as indicated by the microwave data (Toyo-
kawa Observatory, private communication), the acceleration of the charged
particles probably ceased at 0635 UT, 3 min after the eclipse of OSO-7.

For a flux of 80 photons cm"2 and PQ ranging from 60 MV to 70 MV, we
find from Figure XI.C-3 that N ( > P) is between 1 .4 and 6 X 1 033 . These
numbers should be compared with the number of protons released from the
flare based on measurements of the proton flux near earth. This number can
be obtained from the local densities shown in Figure XI.C-7 if we assume that
the particles fill some volume (V) in the interplanetary medium to that density.
A conservative estimate of this volume would be to assume that it is that of a
cone of opening angle 30° with vertex at the sun and height 1 .5 AU. This is
consistent with the fact that the August 4 flare was located at 1 5° N latitude
on the sun and that some charged particles were observed on the Pioneer-10
space probe at 2 AU from the sun (B. J. Teegarden, private communication).
The volume (V) is therefore ~ 1039 cm3, and from Figure XI.C-7, N (> P) is
7 X 1035 and 1 .5 X 1036 for PQ equal to 70 MV and 60 MV, respectively.
When compared with the values of N(P> 0) deduced above from the 7-ray
observations for the thick-target model, we see that possibly not more than
about 1 percent of the flare -accelerated protons could have escaped downward
into the sun. However, we should note that this conclusion is strongly depen-
dent on our estimate of the total number of protons at the sun, based on the
local proton observations. Clearly a more detailed understanding of the
source function of flare protons in the interplanetary medium is required.

In the case of the thin-target model, the instantaneous photon flux at earth
directly determines the emission measure at the sun, that is, the product
nN (P > 0). From Figure XI .C 4 we find that the observed flux of 0.22 pho-
tons cm"2 • s"1 implies that nN is 1 to 2 X 1045 cm"3 for PQ ranging from
70 MV to 60 MV. Then, using the values of N ( P > 0) as deduced from the
proton flux near earth, we get (n) is 1 .5 to 3 X 109 cm"3. As before, note

312 THEORY

that this density is subject to the uncertainty in the total number of protons
at the sun obtained from the local observations.

Let us finally evaluate the total energy in the flare-accelerated protons that
produce 7-rays at the sun. For the spectrum given in Equation (XI.C-1) and
for nonrelativistic protons, the average proton energy is P 2/Mc2. Therefore,
for the thick -target model, the energy in the 1 .4 to 6 X l(r3 protons as
deduced above is 1.2 to 3.7 X 1028 ergs.

For the thin-target model, both the instantaneous 7-ray production rate and
the instantaneous proton energy loss rate depend on the product of the
ambient density (n) and the total number of protons at the sun. The ratio
of these two rates, therefore, is independent of both n and N and depends
only on the spectrum of the accelerated particles, that is, PQ in our calcula-
tions. This point was first made by Lingenfelter (1969). For the proton
distribution in Equation (XI.C-1), normalized to one proton of rigidity greater
than zero, the instantaneous energy loss rate (W) to ionization, excitation and
nuclear interactions in ambient hydrogen of unit density is given in Table
XI.C-4. As can be seen, for PQ's between 60 to 70 MV, W is very closely equal
to 1 .7 X 10"18 ergs • s"1 . Using the values of Nn given above, we see that 7-ray
production by flare-accelerated protons is accompanied by the dissipation of
1 .7 to 3.4 X 1027 ergs • s"1 . For the time interval of 6 min during which 7-rays
were observed, the total dissipated energy is 0.6 to 1 .2 X 1030 ergs.

This energy is comparable to the total of ~ 2 X 1030 ergs emitted in H-a by
the flare of August 4, 1972, (H. Zirin, private communication) and is consis-
tent with the suggestion of Gordon (1954) that the optical energy from flares
results from ionization losses of accelerated particles. On the other hand, in
the thick -target model the ionization losses of the 7-ray producing protons is
only a few percent of the observed optical energy; in this case additional
energy loss may be expected from electrons.

Table XI .C4

Energy Loss Rate of the Proton Distribution in
Equation (XI.C-1) in Solar Material of n = 1 cm"3

P0(MV)

W(ergs-1)

2X 10'18

1.9 X 10'18

1.8 X 10"18

1.6 X 10"18

100

1.5 X 10-18

120

1.4 X 10'18

200

1.2 X 10_18

300

1.1 X 10"18

NUCLEAR GAMMA RA YS FR OM SOLAR FLARES 313

REFERENCES

Allen, C. W., 1963, Astrophysical Quantities, London.

Audouze, J., M. Ephere, and H. Reeves, 1967 , High-Energy Nuclear Reactions
in Astrophysics , B. S. P. Shen, W. A. Benjamin, eds., New York, p. 255.

Bostrom, C. O., J. W. Kohl, and R. W. McEntire, 1972, The Solar Proton
Flux - August 2-12, 1972, The Johns Hopkins University, Applied Physics
Laboratory, Silver Spring, Maryland.

Cameron, A. G. W., 1967, Origin and Distribution of Elements, L. H. Ahrens,
ed., Pergamon Press, London.

Cheng, C. C, 1912, Space ScL Rev. , 13, p. 3.

Chupp, E. L., D. J. Forrest, P. R. Higbie, A. N. Suri, C. Tsai, and
P. P. Dunphy, 1973, Nature, 241 , p. 333.

Gordon, I. M., 1954, Dokl. Akad. NaukSSSR, 96, p. 813.

Hughes, D. J., and R. B. Schwartz, 1958, Neutron Cross Sections,
Brookhaven National Laboratory, Upton, New York.

Jacobs, W. W., D. Bodansky, J. M. Cameron, D. Oberg, and P. Russo, 1972,
Bull. American Phys. Soc, 17, p. 479.

Lingenfelter, R. E., and R. Ramaty, 19 '67 ', High-Energy Nuclear Reactions
in Astrophysics, B. S. P. Shen, and W. A. Benjamin, eds., New York, p. 99.

Lingenfelter, R. E., 1969, Solar Phys. , 8, p. 341 .

Ramaty, R., and R. E. Lingenfelter, 1973a, High Energy Phenomena on the
Sun, Conference Proceedings, R. Ramaty and R. G. Stone, eds.,
GSFCX-693-73-193,p. 301.

, 1973b, Conference Papers, Thirteenth International Conference

on Cosmic Rays, Denver, Colorado, paper 1 34.

Stecker, F. W., 19 69, As trophy s. Space Sci. , 3, p. 479.

Zobel, W., F. C. Maienschein, J. H. Todd, and G. T. Chapman, 1968,
Nucl. Sci. and Eng. , 32, p. 392.

DISCUSSION

White:

Could you elaborate a little bit more on the difference between the experi-
mental observation and the theoretical calculations?

314 THEORY

Ramaty:

In my first calculation with Richard Lingenfelter, we took certain numbers
published by Bill Webber who calculated the total number of protons at the
sun using a three-dimensional isotropic diffusion model in the interplanetary
medium. Just by doing that, but propagating the particles in a reasonably
small solid angle, something like 30°, we get a reduction of a factor of 50 in
the total number of protons at the sun.

For this particular flare of August 4, we get reasonable agreement if I assume
that only about 10 percent of the protons seen at the earth are coming from
the sun. In this case, according to Frank McDonald, the observed particles
at earth were in great part accelerated by flare -produced shock waves.

White:

Can you definitely ascribe it to the number of protons observed, rather than
the number of g/cm2 traversed by the charged particles at the sun?

Ramaty:

I think I can, but of course there are two models: the thick -target model,
which is completely independent of the number of g/cm2 , and the thin
target, for which the relevant parameter is the ambient density in the flare.

What you get from the 7-rays is a product, a total number of particles from
the sun times the ambient density: the so-called emission measure.

Chapter XII

A. ULTRA-HIGH ENERGY GAMMA RAYS

A, W. Strong, J. Wdowczyk,* and A. W. Wolfendale

University of Durham

INTRODUCTION

The problem of the origin of the cosmic radiation is well known. Although
there are many potential sources in the galaxy such as novae, super novae,
pulsars, and the galactic nucleus, there is great difficulty in explaining the
observed isotropy of the radiation above 1017 eV where the galactic magnetic
field is not strong enough to randomize the particle directions and consider-
able anisotropics should result. The suggestion of an extragalactic origin for
all the radiation demands a rather high-energy density of this component
throughout the universe (y 1 eV • cm"3), but if only particles above 1015 eV
or so come predominantly from outside the galaxy this difficulty disappears.

If these very energetic primaries are indeed of extragalactic origin, their
interaction with the radiations in space, principally the relict radiation
(Penzias and Wilson, 1965, and later papers), becomes a process of impor-
tance. The protons will lose energy by way of interaction with this radia-
tion and produce, successively as the energy is increased, e+e" pairs and pions.
In turn, the energy loss would be expected to show up as an increase in slope
of the energy spectrum of protons recorded at the earth.

There is the well known increase of slope at ^ 3 X 1015 eV and it is con-
ceivable that this arises from just such interactions, principally e+e" produc-
tion, at early stages of the evolution of the universe when the relict radiation
was at a higher temperature than its present value of 2.7 K (Hillas, 1968;
Strong et al., 1973a,b). The latter authors have calculated the flux of 7-rays
expected to result from such interactions and, although they are hardly of
'ultra -high' energy, they will be considered briefly later because their origin
is similar to that of the much higher 7-rays which are the main subject of
this work.

"Speaker.

315

316 THEORY

If cosmic-ray protons did not start to be produced until comparatively late
stages in the evolution of the universe when the temperature of the relict
radiation was close to its present value, e+e" production would cause a reduc-
tion of primary proton intensity by a factor ^ 3 above 'v- 3.1018 eV and
pion production would cause a reduction starting at ^ 5.1019 eV and reaching
a factor ^ 100 above ^ 2.1020 eV. Such a reduction does not appear to have
been detected experimentally (although at the energies in question, the num-
ber of events detected, that is, extensive air showers, is small and errors of
energy determination are not negligible). A possible way out of the problem
is to assume that the production spectrum of primary protons above 1019 eV
is flatter than that at lower energies so that after attenuation the spectrum at
the earth has roughly the same slope as that below 1019 eV. Although this
idea is perhaps improbable, it is by no means impossible. Clearly in this case,
there will be significant energy going into pions and electrons and eventually
7-rays, and these will be, in principle, detectable. The details to be described
in the following sections are taken largely from the work of Wdowczyk et al.
(1972).

INTERACTION OF PROTONS WITH THE RELICT RADIATION

The Attenuation Length for Photomeson Production

The interaction process has been considered by a number of authors starting
with Greisen (1966) and Zatsepin and Kuzmin (1966), and an accurate analy-
sis has been given by Stecker (1968). Stecker (1968) has summarized
experimental data on the total photomeson production cross section and
inelasticity for high-energy protons in the relict radiation as a function of
7-ray energy in the proton rest system and used the data in the derivation of
X (X = (K n a) 'l where K is the inelasticity of the interaction, n is the

a v a v P 7 eff p ■* 7

photon density, and a is the meson production cross section). The values of
X derived in this way are given in Figure XII. A- 1 . Kuzmin and Zatsepin (1968)
and Adcock (1970, private communication) have derived values of the inter-
action length (X.) as a function of E , and these values are close to what would
be expected from Stecker's results.

Energy Spectrum of Protons

Wdowczyk et al. (1972), considered two limiting forms for the energy spectrum
of protons at the earth, as shown in Figure XII. A-2. The higher intensities
(A) come from the work of Linsley (1963) and the lower spectrum (B) is that
due to Andrews et al. (1971). Although the latter seems more probable,
results for both will be given so that interpolation (or extrapolation) may be
made if results are needed for some other spectrum.

UL TRA-HIGH ENER G Y GAMMA RA YS

317

kin1

(cm) 10

1 r

STARLIGHT

& I.R.

l 1 1 1 1 r-! r

Pr (S)
bb

Ey(eV)

Figure XII.A-1 . Interaction length against photon energy for collisions of
energetic photons with photons of the various photon fields (e+e~ production
only). The full lines represent the calculations of Wdowczyk et al. (1972),
and the dotted lines those of Stecker (1971). Also shown (top right-hand
corner) is the attenuation length for protons in the relict radiation from the
work of Stecker (1968).

Energy Spectrum of 7-Rays on Production

In Wdowczyk et al. (1972), use was made of Stecker's data to calculate the
production spectra of 7r°-mesons, and in turn that of 7-rays, with the results
shown in Figure XII.A-2. An important datum is the total energy going into
7-rays: for A this is

7 X 10'25 eV • cm"3 s"1 from 7r°-mesons, together with
4 X 10"25 eV • cm"3 • s"1 from e+e" pairs.
For spectrum B the corresponding figures are

6 X 10"26 eV • cm"3 • s"1 from 7r°-mesons and
*v» 4 X 10"25 eV • cm"3 • s'1 from e+e" pairs.

(The spectra A and B are very similar in the energy region where pair pro-
duction is important.)

318

THEORY

E
o

-60

-62-

-64 -

-66"

-68
10

\ V

—

_Wa

\ "

■

10
Ev(eV)

-39

-I -41
10

~\ -43
10

-45

-47

E
u

Figure XII.A-2. Alternative primary proton spectra jA and jB and the
corresponding 7-production spectra, W(E ), from Wdowczyk et al.
(1972).

With spectrum A and assuming a residence time for photons in the
universe of T = 13 X 109y = 4 X 1017 s, the integrated 7-ray intensity is
7 X 102 eV • cm"2 • s"1 • sr"1 , and the energy density is 3 X 10'7 eV • cm"3 .

This energy density can be compared with that in the proton spectrum at the
earth above 5 X 1019 eV of ~ 3 X 10"10 eV • cm"3 . The large disparity is

ULTRA-HIGH ENERGY GAMMA RA YS 31 9

because of the fact mentioned in the first section of this paper: that the proton
spectrum on production must be much flatter than that observed. Essentially,
the detected protons come from within one attenuation length (y 3 X 1025 cm
at a mean energy of 3 X 1020 eV) whereas 7-rays come from the whole of the
universe (y 1028 cm).

The production spectra in Figure XII.A-2 are close to those given by Stecker
(1973, and private communication) for the same alternative proton spectra.

PROPAGATION OF 7-RAYS THROUGH THE UNIVERSE

Extragalactic Radiation Fields

Summaries of the radiation fields and the corresponding interaction mean-free
paths for 7-7 collisions have been given by a number of authors, notably Gould
and Schreder (1967a, b), Stecker (1971), and Wdowczyk et al. (1972). There
is, understandably, agreement for the relict radiation, but a small disparity for
starlight and IR, and a large disparity for the radio background. A comparison
is made in Figure XII. A- 1 between the results of Wdowczyk et al. (1972) —
full line-and that of Stecker (1971 -an updating of Gould and Schreder,
1967a, b)— dotted line. Of importance for the propagation of 7-rays above
1020 eV is the difference in the radio background, and this needs to be con-
sidered. The problem of measuring the isotropic component of the radio
background at the earth is severe and experimental differences are rather great.
Wdowczyk et al. (1972), used the measurements of Clark et al. (1970), which
show a fall off in intensity at photon energies below 10"8 eV, and insofar as
these measurements are later than those used by the other authors, their
analysis will be used here. It is interesting to note that if the shorter interac-
tion lengths are valid for E > 1020 eV, then the result will be a reduction in
the intensity of such 7-rays at the earth but an increase in intensity at energies
just below this value.

The Interaction Process

The generated high-energy photons will interact with the photons of the
various radiation fields to produce electron pairs, 7 + 7 -> e+ + e". Muon pairs
can also be produced if the energy is high enough (E > 1019 eV for relict
radiation) and the mean-free paths for muon-pair and electron-pair production
eventually become equal; however, in the energy region where a significant
effect would occur, the radio background takes over.

The interaction process has been examined in some detail by Bonometto and
Lucchin (1971), Allcock and Wdowczyk (1972), and by Wdowczyk et al.
(1972). The authors have pointed out the important fact that at high-photon
energies the angular distribution of the electrons peaks in the forward and

320

THEORY

backward directions so that in the laboratory system one of the electrons takes
an increasingly large fraction of the energetic photon energy. In the absence
of an extragalactic magnetic field of magnitude above ^ 10"11 G (see
Wdowczyk et al., 1972), the electrons produced will collide with the relict
photons and produce further energetic photons by the inverse Compton
effect; in this way a 7-ray cascade will be built up.

First Generation 7-Spectrum

Whether or not magnetic fields are present, the 7-ray spectrum of Figure

XII. A-2 will be generated and a first generation spectrum L (E ) = 1/47T

W(E ) A(E ) will be formed. This spectrum has been calculated in Wdowczyk

et al. (1972), for proton spectra A and B with the result shown in Figure

XII.A-3.

,-39 , , .

1 1 I I

-40

CO
CVJ

E 10

-41

-42

- 10"43L

SPECTRUM A.

,18

>19

1010 101"
Ey(eV)

Figure XII.A-3. First generation production spectra from Wdowczyk et al.
(1972). If the mean extragalactic magnetic field exceeds 10~10 G or so,
cascading will be inhibited and these will be the spectra of y-rays above
'N/ 1018 eV.

UL TRA-HIGH ENER G Y GAMMA RA YS 321

Cascading in the Absence of Galactic Magnetic Fields

The cascading problem is one of some complexity, and the only calculations
reported to date appear to be those in Wdowczyk et al. (1972). An order of
magnitude estimate of the upper limit to the intensity at low energies (where
most of the energy will eventually lie) can be obtained from energy conser-
vation. For example, if starlight were to be disregarded (that is, X. for star-
light > 1028 cm) then, very roughly, the intensity would have an average
value below 1014 eV (Figure XII.A-1) of I ( < 1014 eV) given by

14
101H eV

EI (E ) d E ^ 7 X 102 eV • cm"2 • s1 ■ sr"1

y c v y' y

for proton spectrum A; (where EQ « 10 eV); that is

Ic (E < 1014 eV) % 10"25 cm'2 • s"1 • sf1 • eV1

Similarly, in the presence of considerable starlight but no radiation causing
attenuation below 1011 eV, we would expect

Ic(E < 1011 eV) £ 10"19 cm"2 -s'1 -sr'1 -eV1

again for proton spectrum A.

The diffusion equations were solved in Wdowczyk et al. (1972) giving the
energy spectra shown in Figures XII .A-4 and XII. A-5. It can be seen that the
intensities at low energies are not inconsistent with what would be expected
from the remarks in the previous paragraph.

Cascading in the Presence of Galactic Magnetic Fields

The presence of significant fields causes the electrons produced in 7-7 colli-
sions to lose energy by synchrotron radiation and thus give rise to 'low'
energy 7's rather than to transfer most of their energy to a single photon by
Compton interactions. The problem was considered in Wdowczyk et al. (1972),
and results were given for what might be a reasonable field: < B ) = 10"9 gauss.
Not surprisingly, perhaps, the synchrotron spectrum so derived is rather similar
to that from cascading, below 1015 eV. However, a note of caution is neces-
sary because of the effect of the field on the proton spectrum. It is possible
to envisage a situation where the proton spectrum is higher elsewhere and the
7/p ratio at the earth would be correspondingly higher.

322

THEORY

10
10
10
10
10'
10"
10"

l_ V X-RAYS

-9

-13

i io-2-

£ 10-33

io-37
io-41 -
io-45-

i — i — i — i — i — i — i — I — I — i — i — i — i — I — i — r

DIFFUSE

PRIMARY SPECTRUM A.

SYNCHROTRON v
RADIATION

r's FROM

p-ybb
INTERACTIONS

FIRST
GENERATION -\

i i i I I I I I I I I I I I I L

IO6 IO8 io10 io12 io14 io16 io18 io20

Ey(eV)

Figure XII.A-4. Gamma-ray spectra over the whole energy
range, from Wdowczyk et al. (1972). The diffuse X-ray
spectra summarized by Ipavich and Lenchek (1970) are also
shown, as is primary proton spectrum A.

SUMMARY OF PREDICTIONS CONCERNING 7-RAYS ABOVE 101Z eV

The intensities of 7-rays shown in Figures XII.A-4 and XII .A-5 (for the case of
< B ) < IO"1 1 G) have been used to give the 7/p ratios shown in Figure XII.A-6.

It will be noticed that the ratios are approaching measurable fractions at
energies above IO19 eV. Of particular interest is the peak in the region of
2 X IO19 eV which comes from the effect of the transition from domination
by relict radiation to that by the radio background at this energy (Figure
XII.A-1). An approximate analysis was made in Wdowczyk et al. (1972) of

ULTRA-HIGH ENERGY GAMMA RA YS
lO"1

323

-5

-9

i — i — i — i — i — i — i — i — i — i — r

PRIMARY SPECTRUM B.

Ey(eV)

Figure XII.A-5. Gamma-ray spectrum for primary proton
spectrum B, from Wdowczyk et al. (1972). (approximate -
relaxed from Figure XII.A-4).

the upper limit that can be set on this ratio from studies of extensive air
showers (7-initiated showers would be poor in muons compared with proton-
initiated showers). It can be seen that so far the experimental limit is signi-
ficantly higher than the maximum predicted ratio. However, there are suffi-
cient uncertainties in our knowledge of extragalactic parameters to make it

324

THEORY

0.1

0.01

0.001

APPROX
EXPERIMEI
UPPER
LIMIT

.18

1st.

GENERATION,

TOTAL, B./

.19

10*~ 10" 10"

E (eV)

Figure XII. A-6. 7/p ratio from Wdowczyk et al. (1972)

possible that detectable fluxes of very energetic 7's do exist, and it is urged
that systematic searches be made. One point that should be stressed in this
connection is the possibility of a nonuniform radio background; this could
produce a transition effect which would concentrate 7-ray energy in a parti-
cular region to an even greater extent than in the present case and give rise to
a much higher ratio.

GAMMA- RAYS IN AN EVOLVING UNIVERSE

As remarked in the Introduction, there is the possibility that the kink in the
proton spectrum at 3 X 1015 eV is connected with electron-pair production
on the relict radiation at early epochs. The 7-rays expected from these inter-
actions may allow constraints to be put on models in which the primary
spectrum above 1014 eV is of extragalactic origin; this topic therefore has
relevance to the ultra high-energy 7-region.

UL TRA-HIGH ENERGY GAMMA RA YS

325

Strong et al. (1973a, b), have examined the problem in detail and their
derived 7-spectrum at the earth is given in Figure XII. A-7. As will be appre-
ciated, although the total energy in the spectrum will be constant (it is
'W .7 X 105 eV • cm"2 • s"1 • sr'1 ), the spectral shape will depend on the energy

>
a>

10"

10-V

-12

w 10-13

-15

■f- OSO-3 (Clark etal., 1971)

♦* COSMOS-208 (Bratolubova Tsulukidze et al., 1970)
_rk- PROTON 2 (Bratolubova • Tsulukidze et al., 1970)
COSMOS 1 63 (Golenetskii et al., 1971)

I 1 ERS-18 (Vetti et al., 1970)

lllllll Mayer- Hasselwander etal., 1972
I Daniel etal., 1973

10"

10'

10 10

Ey(eV)

10"

Figure XI I. A-7. Comparison of observed and predicted isotropic
flux of 7-rays. The predicted intensities are from the work of Strong
et al. (1973), with (1) representing the more probable situation.

326 THEORY

density of extragalactic starlight. The 7-ray spectra in Figure XII.A-7 corres-
pond to different assumptions about the variation of starlight density with
epoch.

The experimental situation is not clear. There appear to be a number of
intensities below the expected spectra and if these are correct then the
suggested origin of the proton spectrum kink is not correct (although this
does not preclude the very energetic protons of this energy being of extra-
galactic origin). However, the recent measurements of Mayer-Hasselwander
et al. (1972), are in good agreement with the prediction.

A firm conclusion cannot be made at this stage; therefore, although with new
measurements being made at the present time, this problem, at least, should
be solved rather soon.

ACKNOWLEDGMENTS

The authors wish to thank the Science Research Council of the U.K. for the
provision of research grants to support this work.

REFERENCES

Allcock, M. C, and J. Wdowczyk, 1972, IlNuovo Cimento, 9, p. 31 5.

Andrews, D., D. M. Edge, A. C. Evans, R. J. Reid, R. M. Tennent,

A. A. Watson, J. G. Wilson, and A. M. Wrey, 1971, Proc. 12th Int. Conf.
on Cosmic Rays, Hobart, Australia, 3, p. 995.

Bonometto, S. A., and F. Lucchin, 1971 , Letters Nuovo Cimento, 2, p. 1299.

Bratolubova-Tsulukidze, L. I., N. L. Grigorov, L. F. Kalinkin,

A. S. Melioransky, E. A. Pryakhin, I. A. Savenko, and V. Ya. Yufarkin,
1970, Acta Phys. Hung. , 29, Suppl. 1 , p. 123.

Clark, G. W., G. P. Garmire, and W. L. Kraushaar, 1971, Proc. 12th Int. Conf.
on Cosmic Rays, Hobart, Australia, 1, p. 91.

Clark, T. A., L. W. Brown, and J. K. Alexander, 1970, Nature, 228,
p. 847.

Daniel, R. R., G. Joseph, and P. J. Lavakare, 1973, X-Ray and Gamma-
Ray Astrophysics, IAU Symposium No. 55, (Madrid), H. Bradt and
R. Giacconi, eds., D. Reidel, Dordrecht, Holland.

Goletskii, S. V., 1971 , Astrophys. J. Letters, 9, p. L69.

UL TRA-HIGH ENERGY GAMMA RA YS 32 7

Gould, R. J., and G. P. Schreder, l967a,Phys. Rev. , 155, p. 1404.
, 1967b, Phys. Rev., 155, p. 1408.

Greisen, K., 1966, Phys. Rev. Letters, 16, p. L748.

Hillas, A. M., 1968, Canadian J. Phys., 46, p. S623.

Ipavich, F. M., and A. M. Lenchek, 1970, Phys. Rev.,D2, p. 266.

Kuzmin, V. A., and G. T. Zatsepin, 1968, Canadian J. Phys., 46, p. S617.

Linsley, J., 1963, Proc. Int. Conf. on Cosmic Rays, Jaipur, 4, p. 77.

Mayer-Hasselwander, H. A., K. Pfeffermann, H. Pinkau, H. Rothermel and
M. Sommer, 1972, Astrophys J. Letters, 175, p. L23.

Penzias, A. A., and R. W. Wilson, 1965, Astrophys. J., 142, p. 19.

Stecker, F. W., 1968, Phys. Rev. Letters, 21, p. L1016.

, 1971, Cosmic Gamma Rays, Mono Book Corp., Baltimore.

, 1973, Astrophys. and Space Sci. , 20, p. 47.

Strong, A. W., J. Wdowczyk, and A. W. Wolfendale, 1973a, Nature, 241, p. 109.

, 1973b,/. Phys. A., in press.

Vette, J. I., D. Gruber, J. Matteson, and L. E. Peterson, 1970, Astrophys. J.
Letters, 160, p. LI 61.

Wdowczyk, J., W. Tkaczyk, and A. W. Wolfendale, 1972,7. Phys. A. , 5, p. 1419.

Zatsepin, G. T., and V. A. Kuzmin, 1966, Zh. Eksperim. Teor. Fiz. Letters,
4, p. LI 14.

Chapter XIII

A. A COMPARISON OF THE RECENTLY
OBSERVED SOFT GAMMA-RAY
BURSTS WITH SOLAR BURSTS
AND THE STELLAR SUPER-
FLARE HYPOTHESISE

F. W. Stecker and K. J. Frost

Goddard Space Flight Center

Recently, Klebesadel, Strong, and Olsen (1973) reported the exciting discovery
of 7-ray bursts having a typical duration of the order of seconds and typical
photon energies of the order of hundreds of keV. This observation has now
been confirmed by Cline, Desai, Klebesadel, and Strong (1973) using the
detector aboard the satellite IMP-6 (Chapter VILA).

Predictions of 7-ray bursts from superndvae have been made by Colgate (1968),
but there are several difficulties in interpreting the observed bursts as origina-
ting in supernovae. In particular, the observed bursts have typical durations
of the order of seconds with multiple bursts being common. They also appear
to have soft exponential spectra with photon energies in the range of 150 to
250 keV. They have been observed to occur frequently with no apparent
correlation with observed supernova events.

In contrast to the observed events, Colgate (1968) has predicted that 7-ray
bursts from supernovae would have durations of the order of 10"5 s and hard
power-law energy spectra with a characteristic energy of about 2 GeV. It
thus appears to us possible that the observed bursts do not originate in super-
novae, and that alternative possibilities for the origin of these bursts should
be explored. We suggest here the alternative possibility that these outbursts
are simply giant versions of the X-ray bursts typically seen in solar flares.

The observed 7-ray bursts bear a strong resemblance in many respects to the
solar X-ray bursts observed recently with a 2-s time resolution (Frost, 1969;
Kane, 1969). Figure XIII.A-1 shows a representative nonthermal solar X-ray
burst. This burst is dominated by two impulsive spikes, each about 10 s in
duration. If a burst such as this were emitted by a star other than the sun,

*Post-Symposium theoretical paper, see Introduction,
f Published in Nature Physical Science, October 1 , 1973.

329

330

THEORY

1400

ft 100°
CO

O
O

600

200

-30S

SOLAR X-RAY

BURST
15-250 Kev.
FEB. 2, 1969

eMf

i i i i i

¥w

i i i

0508 0513 0518 0523 0528

U.T.

Figure XI I I.A-1 . A solar X-ray burst observed on OSO-5 with a time struc-
ture similar to that observed for the nonsolar bursts.

then only the narrowest parts of the burst might be detected above the back-
ground noise. Such bursts would appear to be shorter than solar bursts as is
the case with the recently observed nonsolar bursts. Thus there may be little
intrinsic difference between the time-scales of solar bursts and the suggested
stellar bursts at the source. In both cases, the time scale is much longer than
that predicted for supernovae.

The spectral characteristics of the nonsolar bursts have been measured by
Cline et al. (1973; see also Chapter VILA). These spectral data from IMP-6
are found to be well described by an exponential spectrum of the form
I oc e"E'E° with E being between 150 and 250 keV for a typical initial burst.
Subsequent bursts in multiple-burst events appear to be softer with E ~
100 keV. The spike component of solar X-ray bursts could also fit an expo-
nential energy spectrum with E ~ 100 keV (Frost, 1969). The multiple spike
characteristics seen in the nonsolar bursts are commonly seen in solar X-ray
bursts as well.

GAMMA-RA Y BURSTS AND SUPERFLARE HYPOTHESIS 331

We therefore consider it generally plausible that these bursts are caused by
the bremsstrahlung of electrons accelerated to high energies in a stellar flare
event. Assuming the acceleration to depend only on the strength of the
effective field seen by the electrons, and not on electron energy, the final
energy of the electron will be determined by the time the electron spends in
the field. If we assume this acceleration time to be collisionally determined, the
average time being T, the probability (P) of an electron being accelerated for
time (t) is given by the distribution

dP/dt = r1 et/T (XIII.A-1)

The spectrum of accelerated electrons would then be of the form

I(E)dE oc (dE/Eo)e"E/Eo (XIII.A-2)

where the mean acceleration rate is given by the constant EQ/T and EQ is the
average electron energy. The resulting photon spectrum should then also
approximate an exponential form. The above considerations are fairly general
and it appears that they may be applicable to both solar and nonsolar bursts.

We conclude that the time scale, mean photon energy, and energy spectrum
shape (therefore possibly the acceleration mechanism) for both the solar and
nonsolar bursts are strikingly similar. There is so much similarity that it is a
bit surprising considering that there is such a wide variation of surface condi-
tions among the various stars in the galaxy. There does however appear to be
one important difference. The nonsolar bursts that have been observed, which
presumably must be both the closest and strongest of the nonsolar bursts,
have a much greater intrinsic intensity than their solar counterparts. The
strongest solar flares could have a total energy of ~ 1032 erg (Bruzek, 1967).
The bursts seen by Klebesadel et al. (1973) involve an energy flux of ~ 10"5
-10"4 erg/ cm2. Denoting this flux by e, the X-ray energy at the source is
given by

fts 27rR2e (XIII.A-3)

assuming the source flare radiates into 2n sr. Assuming e = 3 X 10" erg/ cm ,
a source at a distance R = 10 pc would have a typical total X-ray energy
£ 2s 2 X 1035 erg and a corresponding total energy of 1038 to 1039 erg. A
stellar burst of the type hypothesized here would then involve the acceleration
of ~ 106 to 107 times more electrons than a strong solar flare. We may
speculate that such an event might involve a star with a magnetic field strength
~ 103 times larger than the sun. Such fields may not be uncommon, particu-
larly in stars earlier than FO, although the observational establishment of these
fields is difficult and often impossible (Babcock, 1960). In addition, common

332

THEORY

white dwarf stars may have surface fields up to 3 X 107 G (Ostriker, 1970) so
that they may be likely sources for these bursts. It seems reasonable to assume
that such stars as are likely to produce the observed bursts should be near
enough so that no concentration toward the galactic plane should be expected.

The stellar flare hypothesis immediately lends itself to various observational
tests. Possible observational consequences are: (1) repetitions of the bursts at
the same position; (2) simultaneous radio bursts at the same position; and
(3) 7-ray lines at 0.51 MeV (positron annihilation), 2.23 MeV (n+p->d-Hy),
4.4 MeV (C12 ) and 6.1 MeV (O16 ) as have been seen in strong solar flares
(Chupp et al., 1973, see also Chapter VI. A). These lines may be present
because the flare can accelerate protons as well as electrons so that various
nuclear reactions may occur in the flare.

If the stellar flare hypothesis is verified, it may imply a significant source of
low-energy cosmic-rays in the solar neighborhood (and throughout the galaxy),
depending on the frequency and intensity of the flares.

The authors wish to thank Drs. T. Cline and R. Ramaty for valuable discus-
sions and T. Cline for communicating his data to us prior to publication.

REFERENCES

Babcock, H. W., 1960, Stars and Stellar Systems, IV, Univ. of Chicago Press,
Chicago.

Bruzek, A., 1967, Solar Physics, J. Xanthakis, ed., Interscience Pub. Co.,
London.

Chupp, E. L., D. J. Forrest, P. R. Higbie, A. N. Suri, C. Tsai, and P. P. Dunphy,

1913, Nature, 241, p. 33.

Cline, T. L., V. D. Desai, R. W. Klebesadel, and I. B. Strong, 1973,
As trophy s. J. Letters, in press.

Colgate, S. A., 1968, Can. J. Phys., 46, p. S476.

Frost, K. J., 1969, Astrophys. J. Letters, 158, p. L159.

Kane, S. R, 19 '69, Astrophys. J. Letters, 157, p. LI 39.

Klebesadel, R W., I. B. Strong, and R. A. Olsen, 197 3, Astrophys. J. Letters,
182, p. L85.

Ostriker, J. P., 1970, Acta Phys. , 29, Suppl. 1 , p. 69.

SECTION 3
COSMOLOGY

Chapter XIV

A. MATTER-ANTIMATTER COSMOLOGY

R. Omnes*

Universite de Paris

INTRODUCTION

Antimatter is quite a relevant subject for a meeting dealing with cosmic 7-rays
because annihilation is an important potential source of hard photons. There-
fore I am glad to have this occasion to report upon some recent work concern-
ing the possible existence of antimatter on a large scale.

The starting point of these investigations was an attempt to understand the
origin of matter as being essentially analogous to the origin of the background
thermal radiation. This background radiation is probably a remnant of a prior
situation where the universe was hot and space was much more compact than
it is now. It was noticed long ago (Gamow, 1948; Alpher, 1948; Alpher,
Bethe, and Gamow, 1948; Alpher and Herman, 1948a, 1948b, 1949, 1950,
1951, 1953; Alpher, Herman, and Gamow, 1948, 1949; Alpher, Follin, and
Herman, 1953; and Alpher, Gamow, and Herman, 1967) that, according to
general relativity, an isotropic universe had to pass through conditions where
the temperature at early times was very high (this is the hot, big -bang cosmo-
logy). When the temperature was somewhat higher than 100 MeV, thermal
radiation contained all kinds of elementary particles including, among others,
nucleons and antinucleons. It is therefore tempting to wonder whether matter
is a remnant of these particles. Preliminary investigations of this question
showed however that a sufficiently efficient separation could not come from
statistical fluctuations (Goldhaber, 1956; Zel'dovitch, 1965). More precisely,
if we introduce the basic parameter 77 = N/N which is invariant under expan-
sion (where N is the present particle density of matter and N the density of
thermal photons), one finds that statistical fluctuations give (Zel'dovitch, 1965)

t?< 10"18
whereas the observed value is

n= 108 to io"10 (xiv A-i)

*Speaker.

335

336 COSMOLOGY

Therefore, some mechanism of separation between matter and antimatter,
more efficient than mere fluctuations, had to be found.

A few years ago it was noticed that nucleon-antinucleon interactions at inter-
mediate energy (less than 1 GeV) could produce such a separation (Omnes,
1969). The basic idea is the following: According to the mesonic theory of
nuclear forces (Ball, Scotti, and Wong, 1969; Aldrovandi and Caser, 1972), it
turns out that the S-wave scattering of nucleons and antinucleons is repulsive,
that is, the scattering lengths are positive. The same result is also found from
a phenomenological analysis of nucleon-antinucleon interactions (Bryan and
Phillips, 1968). This important feature can in principle be checked experi-
mentally by measuring the energy of X-rays emitted by the protonium (p-p)
atom (Caser and Omnes, 1972) with enough precision. This experiment is
now under way at CERN (Backenstoss, private communication). If this
effective low-energy repulsion between nucleons and antinucleons turns out
to be correct, it could in principle induce a separation between nucleons and
antinucleons among the particles constituting thermal radiation at high
temperatures.

This hypothesis has been analyzed theoretically, by using a variety of different
models (Omnes, 1969, 1970, 1972a, 1972b; Aldrovandi and Caser, 1972; and
Cisneros, 1973), with the following conclusion: Separation could indeed be
the result of a phase separation occurring above a critical temperature which
is of the order of 300 MeV. Some approximations had to be made in all of
these models so that this conclusion can only be considered as tentative. I
shall not deal in detail with this question here.

A detailed study of the universe evolution when the temperature drops from
300 MeV to 30 keV has been performed recently (Aldrovandi et al., unpub-
lished) with the following results:

• The parameter t? is essentially stabilized after a period of intense
annihilation at T = 1 MeV at the value

r7> 109 (XIV.A-2)

(only a lower limit could be obtained).

• The system of matter and antimatter constitutes an emulsion (that
is, a three-dimensional maze). The size of such an emulsion can be
characterized by the ratio L between a large volume V and the area
of the matter-antimatter boundary S enclosed in this volume

V
L = <— ) (XIV.A-3)

MA TTER-ANTIMA TTER COSMOL OG Y 337

L is equal to 104'5 cm when T = 1 MeV and a volume L3 contains at
that time a mass of matter of the order of 1013 g.

• Neutrons are lost by annihilation around 1 MeV so that there is no
helium formation at this stage.

Here again, I shall not deal with the details of this analysis.

COALESCENCE

I come now to the first subject of this talk which is to show how annihilation
along the matter-antimatter boundary can induce important fluid motions by
the effect of which the emulsion size L will grow tremendously during the
radiative period. This effect has been called coalescence (Omnes, 1971a, b, c;
Aldrovandi et al., 1973).

First, let me stress that this effect is relevant for any antimatter model of the
universe. Even if the separation effect described above were but a theoretician's
dream, coalescence would still be the essential feature of a model where matter
and antimatter would be given in separate regions in the initial conditions at
time zero (Harrison, 1968, 1970) or any other conceivable model. This inves-
tigation has been carried out by Aldrovandi, Caser, Puget, and Omnes (to be
published).

The basic idea of the model is the following. Along the matter-antimatter boun-
dary, annihilation produces high-energy particles: photons, electrons, and posi-
trons. These particles, together with secondary particles which they put into
motion by collisions, carry their momentum to the fluid which is made of
matter (or antimatter) and radiation over some distance X. Let us consider the
case where the boundary has a curvature radius R and X « R. Because as
many particles generated by annihilation are going towards matter as are going
towards antimatter, the pressure they exert on both sides of the boundary is
inversely proportional to the area of the effective surface where they are stopped.
These areas are proportional to (R + X)2 and (R - X)2 so that a discontinuity
pressure [p] appears along the boundary:

[p] = 2 pa - (XTV.A-4)

where p is the annihilation pressure carried by the high-energy particles.

Equation (XIV .A-4) is of a well-known type: it is essentially the Laplace -Kelvin
formula which gives the discontinuity pressure associated with a surface tension
with coefficient

a = 2p X (XIV.A-5)

338 COSMOLOGY

so that we expect it to reduce the boundary area, that is, according to Equa-
tion (XIV. A-3), to increase L. This is coalescence.

Perhaps it should be mentioned at this stage that the amount of matter or
antimatter connected within the emulsion is infinite despite the finite value of
L (that is, you can go to infinity by staying within the maze) (de Gennes et al.,
1959; Broadbent and Hammersley, 1957). As a result, coalescence is but an
unfolding of the boundary.

THE THEORY OF COALESCENCE

The details of the analysis look a bit different when the temperature is respec-
tively larger or smaller than 100 eV, because (in the first case), high -energy
photons have a small mean-free path (because of the reaction 7 + thermal
photon -»• e+e"). Below T = 100 eV, this mean-free path becomes larger than
L so that primary photons (due to annihilation of ir° -mesons) do not contribute
to coalescence. I shall restrict myself to this last case.

In order to quantitatively treat the coalescence effect, one performs an analysis
of the transfer and thermalization of particles. High-energy particles as well as
thermal photons and matter electrons are described by a set of Boltzmann
equations.

Primary high-energy electrons (generated by if -*■ n~ -> e") first give their
momentum to thermal photons by Compton scattering. Thereby they give
rise to X-rays which carry the momentum. These X-rays travel along a distance
X = (N Oj)'1 and give all their momentum to some electrons by Compton
scattering. After this first collision they travel a distance much longer than L
before being thermalized so that their energy is homogeneously distributed
over the emulsion and does not affect the motion of the fluid. The second-
generation electrons transfer their momentum, partly to thermal photons
through Compton collisions (say a fraction £ of this momentum) and partly
to the matter plasma by Coulomb collisions (that is, a fraction 1 - if).

It is a somewhat trivial but tedious exercise to compute the spectra of the
particles and to write the Boltzmann equations that describe these processes.
Once we have these Boltzmann equations, we can write hydrodynamical equa-
tions by taking, as usual, the first few moments of the particles' distributions.
For the plasma we get an equation of motion which is

pm | = (1-|) 1- 'A-**-. vp (XIV.A-6)

where p is the plasma mass density. J is the momentum density of X-rays.
P is the density of momentum of thermal photons so that the third term in

MA TTER-ANTIMA TTER COSMOLOG Y 339

the right-hand side represents the momentum given by radiation to the plasma.
The second term represents the inverse effect: The plasma transfers its momen-
tum to photons within a time of drag rD . One has

'.-V? T°=7? T» (XIVA-7)

The last term in Equation (XTV.A-6) represents the effect of the plasma pres-
sure that is negligible in general except near the boundary where annihilation
creates a loss in particles.

The equations for thermal photons need not be written here because, even
when they are written in a form involving hydrodynamical motion plus
diffusion, they are still ugly. Let us note only that for distances larger than
\Q (the thermal photons mean-free path), the system plasma + radiation
behaves like a unique fluid obeying the equation of motion:

dV -» E J

p = - v — + — +17

dt 3 K *

V2V +- V(V V) (XIV.A-8)

= —

Here E is the energy density of thermal radiation (E/3 is the pressure) and 17
is the viscosity coefficient

(XIV.A-9)

Equation (XTV.A-8) must be supplemented by an equation for the energy
transfer which is

3E - / - - 4 - \

— +V Jc-D VE+-EV =0 (XIV.A-10)

These equations of motion must be completed by a boundary condition that
gives the pressure discontinuity across the boundary. The basic idea has
already been given and the passage from kinetic equations to discontinuity
follows the lines provided by the kinetic theory of surface tension. One gets

4 J(°)\) a

1 1 =_ (xrv.A-ii)

3 R R

340 COSMOLOGY

THE ANNIHILATION RATE

The annihilation pressure p or the momentum density c J(0) are determined
by the rate of annihilation at the boundary. To compute it, one can use
Equation (XTV.A-6). Essentially what happens is that annihilation creates a
dip in the plasma density which tends to be filled under the effect of the plasma
pressure gradient V pm while the corresponding flow of matter is slowed down
by the drag of plasma against radiation. Thus,

J(0) = N

kT To

47rm t
p

mpc (XTV.A-12)

As a result, it is found that only a small fraction of matter (< 10"2) is annihi-
lated during the coalescence period so that r? does not decrease appreciably.

THE RATE OF COALESCENCE

Given the hydrodynamical equations together with the boundary condition
(Aldrovandi and Caser, 1972), one can compute the rate of change of L with
time. In fact, the extreme geometrical complexity of the emulsion can be
turned into advantage by averaging the equations over a large volume V. This
leads to a simple equation for the variation of L:

■• <*(t) „

L = — L"2 (XIV.A-13)

Pit)

which can be explicitly solved, taking into account the rate of change of a and
p with time which is due to expansion. The result is again quite simple, namely

16 a .

L3= t2 (XIV.A-14)

5 p

It is found therefore that L increases with time, which is coalescence.

For numerical purposes, one can compute the mass (M) contained within a
typical volume (V) of the emulsion. We have chosen for V the average volume
which is seen from an interior point, which turns out to be given by

V = 8ttL3 (XIV.A-15)

so that

M=pmV (XIV.A-16)

MA TTER-ANTIMA TTER COSMO L OGY 341

For T = 3000 K, t = 1013 s (the conventional end of the radiative period), one
gets

M=1043g (XIV.A-17)

that is, a galactic mass.

THE SIZE OF INHOMOGENEITIES: GALAXY FORMATION

We have found that coalescence quite naturally generates inhomogeneities of
matter and antimatter that can be the origin of galaxies.

In fact, things are not that simple; because of annihilation, matter is kept
ionized in the symmetric universe a much longer time than in the conventional
hot big-bang model, so that coalescence can still go on during this long recom-
bination period and generate much higher masses. Moreover, drag becomes
less effective, which tends to increase the rate of coalescence. Furthermore,
the viscosity becomes much smaller so that turbulence can be generated.

The study of this long recombination period is still incomplete. I believe Puget
and Stecker will say more about it in this Symposium (Stecker and Puget, 1972;
Chapter XV.A). The masses of matter will be larger than before, that is,
in the range of mass of clusters and matter will have a turbulent motion (that
is, the kind of situation first envisioned by Ozernoy and Chibisov (1970). The
main difficulty ordinarily found with turbulence (that is, its dissipation at the
end of the radiative period (Peebles, 1972; Dallaporta, 1972)) is much reduced
here since turbulence would be generated by the coalescence motion itself.

ANTIMATTER AND 7RAYS

Coming back to the subject of this meeting, it is interesting to consider the
consequences of this model as far as 7-ray detection is concerned. For the
sake of the argument, we shall consider the Stecker-Puget model where whole
clusters are made of only one type of matter. These clusters are born from the
largest eddies generated by coalescence.

In such a case (Stecker and Puget, 1972;Steigman, 1971, 1972) annihilation
on the boundaries of clusters is too weak to be detectable at the present level.
Apparently, the only detectable 7-rays come from early annihilation and could
be seen in the isotropic background around 1 MeV after being red-shifted
(Stecker et al., 1971). This effect will be described in a communication by
Stecker (see Chapter IX. A).

342 COSMOLOGY

WHAT IS THE EVIDENCE FOR ANTIMATTER?

Except for the 1-MeV bump in the X-ray background, the present model has
behaved somewhat like a hat from which a rabbit was drawn: the correct
amount of matter in the universe has been computed; Steck^r and Puget claim
that the model gives the right kind of turbulence (that is, the right size for the
largest eddies and the right velocities) to agree with the parameters of clusters
and galaxies (namely, their mass and angular momentum), so that it gives
exactly those cosmological parameters which up to now had been hidden in
the initial conditions. Furthermore, the model has also shown a remarkable
knack for embodying past objections (Stecker et al., 1971) and using them
for progress: the hat is still being brushed but the rabbit is alive and well
(Aldrovandi et al., unpublished; Omnes, 1971a, 1971b, 1971c; Aldrovandi et
al., 1973).

However, one feels quite frustrated to find how difficult it is to show experi-
mentally the existence of antimatter.

I am now going to describe briefly one conceivable type of consequence. It
concerns a possible mechanism for the activity of quasars and Seyfert galaxies
which is yet far from being properly analyzed. In fact, I only mention it here
because of its possible relevance to 7-ray astronomy, and my excuse for
releasing it too early will be the occasion provided by this meeting.

AN ECUMENIC MODEL OF QUASARS

Many models of quasars have already been proposed. Some are as follows:
(Zel'dovitch and Novikov, 1971; Burbidge and Burbidge, 1967; Schmidt, 1969;
Robinson, Schild, and Schucking, 1963)

• Quasars have been tentatively identified with supermassive stars
(Hoyle and Fowler, 1963; Fowler, 1964). The main difficulty for
this theory comes from the star temperature which is too low for
nuclear energy to be produced efficiently. One must therefore appeal
to rotational energy, but this raises difficult problems of conversion
(Fowler, 1966; Rosburgh, 1965; Bisnovatyi-Kogan et al., 1967;
Bardeen, 1966; Wagoner; 1969).

• A nonrelativistically rotating supermassive star tends to collapse
rapidly. This has led to a variety of models for quasars where some
stabilization is provided by rotation (Fowler, 1966; Rosburgh, 1965;
Bisnovatyi-Kogan et al., 1967; Bardeen, 1966; Wagoner, 1969),
turbulence, or magnetic fields (Layzer, 1965; Ozernoy, 1966). These
last two agents are good stabilizers, but turbulence should be contin-
uously generated by a process which, to my knowledge, has not yet

been found.

MATTER-ANTIMATTER COSMOLOGY 343

• Several models of quasars identify them with star clusters (Gold et al.,
1965;Ulam and Walden, 1964;Woltjer, 1964; Miller and Parker,
1964; Spitzer and Saslaw, 1966). For our purpose, the basic aspect
of this class of models is the importance attributed to collisions.

• One has suggested antimatter as an efficient source of energy for
quasars (Teller, 1966; Burbidge and Hoyle, 1956; Ekspong et al.,
1966). Here the difficulty is to propose a specific structure for the
matter-antimatter system (Aldrovandi et al., unpublished; Schatzman,
1970). One must also be aware of the limitations imposed to annihi-
lation by the observation of high-energy 7-rays (Clark et al., 1968;
Steigman, 1969).

I shall briefly describe another model for quasars that has been suggested by the
matter-antimatter symmetric cosmology. Because it reconciles many features
of already existing proposals, it might be called an ecumenic model. A con-
venient consequence of this model is that most relevant calculations have
already been published in the literature.

Let us now state the model. A supermassive star £ made of antimatter is
located within the nucleus of a matter galaxy. Energy is generated by the
annihilation of accreting matter and impinging stars. Heat being thus produced
in a stochastic manner, large temperature differences are produced between the
regions where annihilation is taking place and the average temperature. Turbu-
lent convection is therefore continuously generated. On the other hand, high
magnetic fields are expected.

There are reasons derived from our cosmological model to expect the occur-
rence of such a peculiar object. It is conceivable (even though, not yet quite
clear or necessary) that, by effect of the coalescence motions, some amount of
antimatter may be trapped within matter. The general characteristics of
coalescence as described above show that the mass of this inclusion cannot be
too small as compared to a galactic mass, say M ~ 108 M @. The contraction
of such a mass of antimatter will take place after recombination as a conse-
quence of annihilation pressure (that is, the high-energy electrons and positrons
produced by annihilation communicate their momentum to matter and anti-
matter if there is a magnetic field. Such a strong boundary pressure can induce
contraction (Sunyaev and Zel'dovitch, 1972). In this way we expect that a
supermassive star such as 2 could be produced.

It may be that 2 has a hard early life and there are a few unsolved problems
concerning this period. It is necessary that stabilization by turbulence or
magnetic fields occur very soon after the birth of 2 to avoid collapse and this
point has not been clarified (although, in this model, we expect galaxies to
contract at the same epoch as 2 because of the same mechanism). Also, one
does not know why £ should stand in the galactic nucleus; perhaps its large

344 COSMOLOGY

mass could serve to start the initial condensation of the galaxy, or its motion
in the galaxy could lead it to the center either by gravitational effects
(Spitzer, 1971) or because of a specific viscosity generated by annihilation
(Lequeux, private communication).

Assuming the existence of such an object, we will now show that it behaves
in many ways like a quasar. For the sake of definiteness, we shall consider an
object 2 with mass 108 M0 with a radius R = 1 pc. situated at the center of
a galactic nucleus. We shall use data concerning our galaxy for the environment
density so that most of the accreting matter is probably in the form of stars
(Rougoor and Oort, 1960). One finds that 2.2 stars (with a solar mass) are
entering into 2 every year with a velocity 1 000 km/s. The average particle
density <N ) in 2 is 109 antiprotons per cm3, and the average mass density (p>
is 10"15 g/cm3.

The characteristics of 2 are well known (Zel'dovitch and Novikov, 1971). Its
density profile is that of a polytrope with index n = 3. The temperature T is
related to the density p by

M \ 1/6
T=1.97X107K( p1/3 (XIV.A-18)

cg.s.

The thermal luminosity is given by

M
M

Lth = 1.3X1038 [Tr- erg/s (XIV.A-19)

However, it should be pointed out that this value for L . can be overestimated
if large magnetic fields contribute to the pressure near the surface. If it were
left to itself, 2 would start gravitational collapse when it reaches a critical
state corresponding to a central density and a radius

/ M \ "7/2
PC = 2X1018 g/cm3 (XIV.A-20)

\M0 /

Rc (M/M0 = 108) = 3 X 10'2 pc. (XIV.A-21)

(if one assumes 2 to be made of pure hydrogen).

/ M

t =-

-° > 109

—

Lth

\Me

MA TTER-ANTIMA TTER COSMO L OG Y 345

The energy of 2 is then independent of its mass

Ec = -4 X 1054 ergs (XIV.A-22)

An important quantity is the evolution time of 2, which can be quite small if
2 is not stabilized otherwise, namely

-l

yr (XIV.A-23)

2 is heated by infalling stars which begin to annihilate when they penetrate
antimatter. Their initial velocity is (GM/R)/2 = V. The star surface is heated
by annihilation. An energy flux is produced which is essentially given by
4> = VN mc2 where m is the proton mass. (When V is reduced, this flux
becomes of the order of V N mc2 where V is the local sound velocity.) The
cascade of thermalizing particles has been analyzed in another context (Aldro-
vandi et al., 1973). First, X-rays are produced by the products of annihilation
(7, e1) via pair production, Compton effect, and the reactions 7 + thermal
photon -> X and e + thermal photon -> X. These X-rays are thermalized by
Compton effect later. Large quantities of energy are accumulated near the
surface of the star where the particle density is much larger than N. Two cases
are possible which have only been analyzed grossly, and both lead to the same
result: either strong convective motions take place which blow off the star
envelope, or energy is transported by diffusion over a distance of the order of
the star radius. In that case, the local temperature becomes larger than 10 K.
Once again, this leads to a blowing off of the envelope by evaporation.

The long-distance transport of energy in 2, in which the main pressure is
radiative, will take place through shock waves. These shock waves will leave
a complicated pressure distribution resulting into turbulence. Altogether, the
annihilation process appears to be rather complicated and violent and it is very
difficult to analyze it in detail. The only simple relation which can be derived
comes from energy balance:

2 7rR2nVM0c2= <L> (XIV.A-24)

Here n is the star density around 2 and (L> the average luminosity, in general
higher than Lth. Many parameters are free here, so that it is no surprise that
the highest known quasar luminosities are easily obtained.

One will not detect the original products of annihilation. Gamma-rays pro-
duced by 7T°- mesons will be stopped in a short distance by several processes

346 COSMOLOGY

(pair production interactions with protons and electrons, Compton effect,
pair production by collision with thermal photons) so that this model does not
contradict the limit set upon annihilation by 7-ray astronomy (Clark et al.,
1968).

The most difficult question that is raised by this model is to describe the kind
of average equilibrium which will take place in 2. It is only locally heated by
annihilation in a random way and the energy is carried mostly by turbulence
and shock waves. It would obviously be essential to analyze this kind of
process and see what limitations can be imposed on the radius (by star pene-
tration) and on the encounter frequency (by the evolution time of 2). We
have not yet done this work because we were not able to master the problems
of transfer which are involved. Let us note only a favorable circumstance:
strong fluid motions should be continuously generated, which would tend to
stabilize 2 (Layzer, 1965; Ozernoy, 1966).

Another unsolved problem concerns magnetic fields. It is a general consensus
that annihilation can produce large magnetic fields, although only preliminary
studies of this effect have been made (Schatzman, 1970; Peyraud, 1971; Aly,
to be published). Large-scale magnetic fields can also be present in 2 since its
origin or they can be produced by relative motions (including differential
rotation).

Things are complicated by the violent events which the model predicts. Too
much local energy generation can result in instabilities, ejection of antimatter,
rejection of matter, even disruption of 2 (considering the small value of Ec).
However, 2 will not suffer fragmentation (Montmerle, 1971). Despite the
nightmarish character such a system may have for a theoretician, it does not
look incompatible with what is observed.

An important new feature of this kind of model concerns the lifetime of
quasars. Typical values of M = 108 M 0 and ( L > = 1046 ergs/s give a lifetime
t s 109 years. Values of M/M0 up to 102 times higher are still compatible
with the model. This shows that quasars have been active since the origin of
galaxies. Therefore, the highest red shifts of quasars provide important cos-
mological information. Furthermore a strong evolution towards decay is
predicted with the right order of lifetime (Schmidt, 1970).

To conclude, let us now list what relations can be made between the model
and observations.

• Evolution (Schmidt, 1970)

• The validity of the "Christmas tree" behavior for compact radio
sources (Kellerman, 1972; Dend, 1972). The individual flashes
corresponding here to a new star or a new cloud annihilating.

MA TTER-ANTIMA TTER COSMOLOG Y 347

• The analogy between Seyfert galaxies and quasars. In this model,
the difference is only quantitative. All quasars should be in a galaxy
(Kristian, 1972), even if it is only a dwarf one, as one would expect
if the ratio between the masses of matter and antimatter is not far
from one.

• The ejected matter, in the form of dust and gas, has a stellar composi-
tion. Such ejected matter constitutes the atmosphere of 2, which
agrees with the characteristics of the emission lines (Burbidge and
Burbidge, 1967).

• Multiple absorption red shifts are probably due to gas ejected by
radiation pressure and quenched by line-locking (Wampler, 1972).

• Infrared emission might be due to synchrotron emission by annihi-
lation electrons in a high magnetic field (Low, 1970), but most
probably it is due to external dust. Indeed, such dust should be
abundant near a region where stars explode.

• The origin of extended radio sources frequently associated with
quasars and of the cosmic electrons radiating in these sources would

be explained in this model as previously suggested by Layzer (1965) and
Ozernoy(1966).

Important observations to test the model might come from X-ray and 7-ray
observations of quasars with a low-density central star, if it turns out that
enough lower-energy 7-rays from annihilation can escape. A cutoff in the
energy of these 7-rays could be seen at a value related to the temperature
existing in the annihilation region.

REFERENCES

Aldrovandi, R., and S. Caser, 1972, Nucl Phys., B38, p. 593.

, 1972, Nucl Phys., B39, p. 306.

Alpher, R. A., 1948, Phys. Rev., 14, p. 1577.

Alpher, R. A., H. A. Bethe, and G. Gamow, 1948, Phys. Rev., 73, p. 803.

Alpher, R. A., J. W. Follin, and R. C. Herman, 1953, Phys. Rev., 92, p. 1347.

Alpher, R. A., G. Gamow, and R. Herman, 1967, Proc. Nat. Acad. Sci. , 58,
p. 2179.

Alpher, R. A., and R. Herman, 1948a, Phys. Rev. , 74, p. 1737.

, 1948b, Nature, 162, p. 774.

, 1949, Phys. Rev., 75, p. 1089.

348 COSMOLOGY

., 1950, Rev. Mod. Phys., 22, p. 153.

., 1951, Phys. Rev., 84, p. 60.

., 1953, Ann. Rev. ofNucl. Sci., 2, p. 1.

Alpher, R. A., R. Herman, and G. Gamow, 1948, Phys. Rev. , 74, p. 1 198.

, 1949, Phys. Rev., 75, p. 3321.

Ball, J. S., A. Scotti, and D. Y. Wong, 1969, Phys. Rev., 142, p. 1000.

Bardeen, J. M., and S. P. S. Hnand, 1969, Astrophys. J. , 144, p. 953.

Bisnovatyi-Kogan, G. S., Ya. B. Zel'dovitch, and I. D. Novikov, 1967,
Ast.Zh., 44, p. 525.

Broadbent, S. R., and J. M. Hammersley, 1957, Proc. Comb. Phil Soc. , 53,
p. 629.

Bryan, R. A., and R. J. N. Phillips, 1968, Nucl. Phys. , B5, p. 201 .

Burbidge, G. R., and E. M. Burbidge, 1967, Quasistellar Objects, Freeman,
San Francisco.

Burbidge, G. R., and F. Hoyle, 1956, Nuovo Cimento, 4, p. 558.

Caser, S., and R. Omnes, 1972, Phys. Letters, 39B, p. L369.

Cisneros, A., 1973, Phys. Rev.,Dl, p. 362.

Clark, G. W., G. P. Garmire, and W. L. Kraushaar, 1968, Astrophys. J., 153,
p. 203.

Dallaporta, N., and F. Lucchin, 1972, Astron. and Astrophys. , 19, p. 123.

de Gennes, P. G., P. Lafore, and J. P. Millot, 1959, /. de Physique et le
Radium, 20, p. 624.

Dend, W., 1972, Communication at the 6th Texas Symp., New York, in press.

Ekspong, A. G., N. R. Yamdagni, and B. Bonnevier, \966, Phys. Rev. Letters,
16, p. L564.

Fowler, W. A., 1964, Rev. Mod. Phys., 36, pp. 545, 1 104.

, 1966, Astrophys J., 144, p. 180.

Gamow, G., 1948, Phys. Rev., 14, p. 505.

, 1948, Nature, 162, p. 680.

Gold, T., W. I. Axford, and E. C. Ray, 1965, Advances in Astron. and
Astrophys. , 3, Z. Kopal, ed.

Goldhaber, M., 1956, Science, 124, p. 218.

MA TTER-ANTIMA TTER COSMOLOG Y 349

Harrison, E. R., 1968, Mon. Nat. Roy. Astro. Soc, 142, p. 129.

, 1970, Commun. Math. Phys., 18, p. 301.

Hoyle, F., and W. A. Fowler, 1963, Nature, 197, p. 533. (also inM N R.
A.S., 1963, 125, p. 169.)

Kellerman, K., 1972, Comm. at the 6th Texas Symp., New York.

Kristian, J., 1972, Comm. at the 6th Texas Symp., New York.

Layzer, D., 1965, Astrophys. J. , 141, p. 837.

Low, F. J., 1970, Astrophys. J. Letters, 159, p. L173.

Miller, R. H., and E. N. Parker, 1964, Astrophys. J., 140, p. 150.

Montmerle, T., 1971, Thesis, Paris.

Omnes, R., 1969, Phys. Rev. Letters, 23, p. L38.

, 1970, Comments and Addenda, Phys. Rev., 1, p. L723.

, 1971a, Astron. and Astrophys. , 11, p. 450.

, 1971b, Astron. and Astrophys. , 10, p. 228.

, 1971c, Astron. and Astrophys. , 15, p. 275.

, 1972a, Phys. Rep., 3C, p. 1.

, 1972b, Proceedings of the IAU Meeting on Dense Matter,

Boulder.

Ozernoy, L. M., 1966, Astr. Zh., 43, p. 300.

Ozernoy, L. M., andG. V. Chibisov, 1970, Soviet Ast A. J., 14, p. 915.

Peebles, R., 1972, Comments Astrophys. Space Phys.

Peyraud, N., 1971 , Astron. and Astrophys.

Robinson, I., A. Schild, and E. E. Schucking, ed., 1963, Quasistellar Sources
and Gravitational Collapse, Univ. of Chicago Press.

Rosburgh, I. W., 1965, Nature, 207, p. 363.

Rougoor, G. W., and J. H. Oort, 1 960, Proc. Nat. Acad. Sci. U. S. , 46, p. 1 .

Schatzman, E., 1970, Phys. and Astrophys. , CERN, Geneva.

, 1970, CERN Lectures.

Schmidt, M. 1969, Ann. Rev. of Astron. and Astrophys. , p. 527.

, 1970, Astrophys. J., 162, p. 371.

350 COSMOLOGY

Spitzer, L. J., 1911, Nuclei of Galaxies, O'Connell, ed., North-Holland.

Spitzer, L. J., and W. C. Saslaw, 1966, Astrophys. J. , 143, p. 400.

Stecker, F. W., D. L. Morgan, Jr., and J. Bredekamp, 1971, Phys. Rev. Letters,
27, p. L1469.

Stecker, F., and J. L. Puget, 1972, Astrophys. J. , 178, p. 57.

Steigman, G., 1969, Nature, 224, p. 477.

Sunyaev, R. A., and Ya. B. Zel'dovitch, 1912, Astron. and Astrophys. , 20,
p. 189.

Teller, E., 1966, Perspectives in Modern Physics, R. E. Marshak, ed.,
Interscience, New York.

Ulam, S. M., and W. E. Walden, 1964, Nature, 201, p. 1202.

Wagoner, R. V., 1969, Ann. Rev. Astr. Astrophys., 7, p. 553.

Wampler, D., 1972, Comm. at the 6th Texas Symp., New York.

Woltjer, L., 1964, Nature, 201, p. 803.

Zel'dovitch, Ya. B., 1965 , Advances Astron. Astrophys. , Z. Kopal, ed., 3,
p. 241.

Zel'dovitch, Va. B., and I. D. Novikov, 197 1 , Relativistic Astrophysics,
University of Chicago Press.

B. THE DEUTERIUM PUZZLE IN
THE SYMMETRIC UNIVERSE

B. Leroy, J. P. Nicolle, and E. Schatzman*

Observatoire de Meudon

In our present understanding of the model of the symmetric universe, we are
led to the following picture proposed by Omnes (1972)t -

• Separation era, during which a partial separation between baryons
and antibaryons takes place, at t < 10"5 s or kT > 350 MeV
(critical temperature of the phase transition);

• Annihilation era (t < 1600 s, 3 50 MeV >kT> 25 keV). At the
end of the annihilation era, the annihilation pressure becomes
efficient to produce the coalescence;

• Coalescence era (1400 s < t < 106 years; 25 keV > kT > 1/3 eV).

At the end of the coalescence era, the mean-free path of the products
of the annihilation become comparable to the size of the emulsion.

During the annihilation era, the size of the emulsion is governed by the
diffusion of nucleons. If we look more closely at the situation, we can see,
as shown by Steigman (preprint) that the main process is the diffusion of
neutrons (at least, as long as there are neutrons). The only way in which
neutrons can be kept in the emulsion is by neutron electron scattering.
However, this can last only as long as there are blackbody electrons. As soon
as the temperature drops below 0.5 MeV, the number of free electrons, which
goes like 1032-3TMeV3 10"(a25/TMev)> decreases very quickly. If the size of
the emulsion is large enough the neutrons are kept until nucleogenesis takes
place around T 'v 0.1 MeV. If the size of the emulsion is too small, the
neutrons are lost (they annihilate at the boundary of the emulsion), and no
nucleosynthesis can take place. From the analysis of the diffusion process,
this seems to be the case.

* Speaker.

jFurther references can be found in the paper quoted.

351

352 COSMOLOGY

In the following we shall consider how the present abundance of deuterium
can be used as an independent proof that no nucleosynthesis has taken place
and therefore that neutrons were lost before nucleosynthesis. The argument
is the following:

1. we consider the nucleosynthesis during the radiative era;

2. we estimate the relevant cross sections;

3. we estimate the maximum abundance of 4He at the end of the
nucleosynthesis era; and

4. we solve the diffusion problem in order to get an estimate of the
rate of loss of the neutrons. This leads to a correction factor to the
rate of formation of 4He. An estimate of this correction factor, to
match the maximum abundance of 4He, leads to an estimate of the
maximum size of the emulsion.

NUCLEOSYNTHESIS DURING THE RADIATIVE ERA

Let us consider the reactions taking place between nuclei and antinuclei at the
boundaries of the emulsion. Let us assume that we have only protons and
alpha particles. The reaction

"p + 4He -> pions + ( 3 nucleons

D + 1 nucleon
3He or 3T

leads mainly to the production of nucleons and the destruction of a-particles.
Half of the nucleons produced are destroyed in flight in the regions of anti-
matter, either in NN reactions or in Na reactions (Figure XIV.B-1). Let us
call R the probability of the reaction NN in flight and (1-R) the probability
of the reaction Nam flight. Neglecting provisionally the production of
deuterium and tritium, we have the following expressions for the rate of
reaction.

— =-(ov)

dt pp

_PP + <ov>p_ [~1-|-R+-|(1-R)] P«

r3 3 9 "1

pa + <ov> R+-(l -R)

aa L2 2 2 J

+— < ov )-a p

DEUTERIUM PUZZLE IN THE SYMMETRIC UNIVERSE

353

OTHERS

Figure XIV. B-1. Schematic representation of the annihilation
at the interface and the secondary reactions taking place in
flight. Only the main reactions have been plotted following
the aN reaction at the interface. A symmetric figure would
have to be drawn for the aN reaction.

dot

— = - (aw)- pa + < ov > -
^ pa r pQ!

--(l-R)l pa

+ <ov) -

l--(l-R)

By taking a proper average over space, and assuming <p)=<p>, <a>=<a),
calling X and n the ratios

< av >- < av > -

pa aa
X = ; JU =

<™>Pp <0V>pp

we obtain

3 33

dp p3 + 2Xp2a + pa2 (-5X2 +— n) X/za3

4 ' 4

da Xp2a + pa2 (/i + X2) +- a3X/i

354 COSMOLOGY

In a similar way, we can consider the rate of production of deuterium. Esti-
mating that the most important part in the balance equation for the deuterium
arises from the deuterium production, we obtain, B being the branching ratio
in the pa reaction

Apa 5

dD -^-B(p+-aA)

2 2

3 33

dp p3 + 2Xp2a + pa2 (-5X2 +—{i) Xjua3

4 4

It results from the experiments of Barkas et al. (1957), that in the reaction p -
nucleus, 1.3 pion on the average is absorbed in the nucleus, out of the average
5 pions produced in the annihilation.

After annihilation, we are left with an 3He or a 3T in excited states. We shall
assume that the final nuclei left are in the same ratio as observed by
Zaimidoroga (1965, 1967) for the pion capture by 3He. According to the
summary given by Koltun (1969), we have the following ratios:

77- + He3 -> H3 + 77° 15.8 ±0.8%

7T-+He3^H3+7 6.9 ±0.5%

7r" + He3->p + 2n 57.8 ±5.4%

7f + He3^D + n 15.9 ±2.5%

7f + He3->D + n + 7 3.6 ±1.2%

7r' + He3 ->p + 2n + 7 ?

To summarize briefly these data, assuming that the tC can do the same to H
as to 3He, we shall accept the following branching ratios:

p + a^3X + 7T-+3N + 7r 60%

p+a->3X + 7T^N + 2D + 7r 20%

p + a^3X + 7T-*3Y + 7r 20%

From these data, we conclude immediately that very little 4He must have been
left at the beginning of the radiative era, otherwise a too large abundance of
2D would have been produced.

ESTIMATE OF THE CROSS SECTION AND RATES OF FORMATION

The < ov ) includes both the nuclear part and the effect of the convergence of
the wave function for capture at low energy. We have

DEUTERIUM PUZZLE IN THE SYMMETRIC UNIVERSE 355

2irn

av= o , v.
nucl

-27m

1 -e

with n = Z Z e2/hv. If we compare the two cross sections, and calculate the
ratio

(ov>-
pa

<ov> -
pp

we have to include the effect of the charge of the a, the effect of the relative
mass and of the relative velocity in the collision.

Assuming, as already suggested by Schatzman (1970),

we obtain

nucl

3.5

nucl

<ov

In the same way, we obtain, as an estimate,

^56

(ov> _
pp

With these values, we get the main contribution to the rate of change of the
abundances (for small values of a and D),

doc X a

which gives

a = % (p/p0)x

and

dD XB / a v

— px"

dp 2 \p V

356 COSMOLOGY

which gives

B /p\x

We see that deuterium is built up, whereas a-particles are destroyed. Assuming
that we start with zero deuterium, we have

B
Do -%

where the origin is taken at the end of the nucleosynthesis. The final concen-
tration (observed at the present time) gives

P Po P Po Po

from which we derive :

%_2_ _P_
Po B Po

If we take the value of 5 at the surface of the earth, as given by Urey et al.
(1932), and Craig (1961), 5 =s 2.10"4 and with (2/B) a 10, we obtain

% P

— a2.10'3 —
Po Po

If we take the protosolar gas value of Geiss and Reeves (1972), 5 = 3.10" ,
we obtain

% P

— =3.10-4 —
Po Po

The ratio (p/p0) is the annihilation ratio between 0.1 MeV, and 1/3 eV. From
the recent work of Aldrovandi, Caser, Omnes, and Puget (1973), it is quite
clear that most of the annihilation has taken place already by T = 25 keV, and
we cannot expect (p/p0) to be very small. For further calculations, we shall
take (p/p0) = 0.1, or (%/p0) = 2.10"4, which represents a depletion factor A

DEUTERIUM PUZZLE IN THE SYMMETRIC UNIVERSE 35 7

at the end of the nucleosynthesis, compared to the results of Wagoner, Fowler,
and Hoyle (1967), of the order of 10"3.

This confirms entirely what has been announced earlier, that is to say that
there is very little 4He left at the end of the epoch of nucleosynthesis.

RATE OF LOSS OF THE NEUTRONS AND 4He FORMATION

In order to get an idea of the rate of loss of the neutrons, we shall consider
the diffusion with a time dependent diffusion coefficient to the surface of a
sphere with a radius growing with time.

The equation of diffusion,

Id 2 30_30
r2 3r 3r 3t

with 0 = 0 at r = a(t), can be solved in the following way:
Introducing r = x a(t), 0 < x < 1 , dr = D dt/a2 , we have

13 2 30 _ 30
x2 3x 3x 3r

A solution is 0 = sin 7rx/x exp (-7T2r), from which we derive the time scale of
depletion by diffusion towards the boundary

dn \ 7r2D

dt Km a2
The equation of conservation of the neutrons becomes

dn w2D n p
dt a2 r t

n p

We are concerned with the last phase of nucleosynthesis, for kT < 1 MeV,
for which r increases very quickly to infinity. If we simplify the equation
of formation of the a's to a pure neutron capture process we obtain

— = < av > pn
dt P"F

358

COSMOLOGY

and the number of a's at the end of the nucleogenetic period is

<OV>pnPn0 eXP

a =

We shall simplify the whole problem by assuming that the depletion factor A
can be estimated by the quantity

A =

n2D

The average A is obtained in the following way. We calculate the amount of
helium formed from the temperature Tl where the rate of destruction 4He
(7, n) 3He becomes negligible (Tl =0.8 MeV). The concentrations p and nQ
are proportional to the expansion factor to the minus cube, and we can write

T3 dT exp

T3dT

From the estimate of the integral, and writing a = aQ T""M v, it is possible
to get an estimate of a . The result is not very sensitive to the value of n.
With n = 17/6 (corresponding to the rate of growth during the coalescence
period), a diffusion coefficient

D=10815TM -5'2 I0<a25^)

MeV

we obtain for A = 10" a maximum value

a0<10469for5 = 2.10^

a0<104-64for5 = 3.10-s

If we consider the formation of 3He and if we take the abundance ratio
3He/H as 10"5, we obtain aQ < 10460. This is quite compatible with the
diffusion length. For a sphere

LD = ir

Ddt

4.25

MeV

-9/4

DEUTERIUM PUZZLE IN THE SYMMETRIC UNIVERSE 359

CONCLUSION

From this short discussion, we see that the low abundance of deuterium is
some sort of proof that the neutron loss has actually taken place before the
beginning of the nucleogenesis.

We can then assume either that the diffusion length actually determines the
size of the emulsion, and it seems quite possible that the abundance of the
a's was vanishingly small at the end of the nucleogenesis, or that the abundance
of deuterium and other light elements results from the nucleogenesis. It then
leads to a determination of the size of the emulsion during the nucleogenesis.
In fact, a small amount of coalescence before the end of the annihilation
period would be enough to increase the size of the emulsion beyond the
diffusion length and put the two determinations in complete agreement.

A final comment is interesting to make: Since the beginning of the theory of
the symmetric universe, a number of criticisms have been made, which have
been met with success, one after the other. Just like a puzzle, the pieces have
been found to adjust to each other. In the present case, one has the feeling
that the new piece has just matched a hole between two pieces. This gives
great confidence for the future of the model.

REFERENCES

Aldrovandi, R., S. Caser, R. Omnes, and J. L. Puget, 1973 , Astron. and
Astrophys. , in press.

Barkas.W. H.etal., 1957, Phys. Rev., 105, p. 1037.

Craig, H., 1961, Science, 133, p. 1833.

Geiss, J., H. Reeves, 1972, Astron. and Astrophys. , 18, p. 126.

Koltun, D. S., 1969, Adv. in Nucl. Phys., 3, p. 149.

Omnes, R., 1972, Phys. Reports, 3 C, p. 1.

Schatzman, E., 1970, Phys. and Astrophys. , CERN lectures.

Urey, H. C. et al., 1932, Phys. Rev., 40, p. 1.

Wagoner, R. V., W. A. Fowler, and F. Hoyle, 1967, Astrophys. J., 148, p. 3.

Zaimidoroga, 0. A. et al., 1965, Sov. Phys.-JETP, 21, p. 848.

Zaimidoroga, 0. A. et al., 1967, Sov. Phys.-JETP, 24, p. 1 1 1 1.

C. ANTIMATTER IN THE UNIVERSE?

Gary Steigman*

Yale University

INTRODUCTION

In several previous papers you have heard of the development of a cosmological
model that is symmetric in the sense that exactly half the particles in the
universe are, in fact, antiparticles. You have also heard of some of the
observational consequences of such a model, particularly as they relate to
7-ray astronomy. The conclusion the previous speakers have reached is that
it is possible to build such a model without violating the many constraints set
by observation. I am much less convinced than they are of this conclusion
and have in the past addressed myself to some of the problems posed by a
"symmetric" cosmology. Although I think we are all agreed that this subject
is in a rather early stage of development and that there are many as yet
unsolved problems, the subject is sufficiently important to justify our
continuing interest in it.

In these remarks, I wish to adopt an approach that is different from that of
the previous speakers. Rather than asking if a symmetric cosmological model
can be constructed that is consistent with observations, I wish to ask the
question, "If the universe does indeed contain equal amounts of matter and
antimatter, how would we know about it?" There are several straightforward
ways in which antimatter could signal its presence to us, and I shall discuss
them shortly. As we shall see, there is no evidence whatever for large amounts
of antimatter in the universe. From that we may reach one of two conclusions.
Either the universe is not symmetric, or, if it is, the ubiquitous antimatter
prefers to remain clandestine (see Puget, Chapter XV .A for relevant discussion).
If indeed we adopt the latter conclusion, then the limits set by observations

"Speaker.

jol

362 COSMOLOGY

set severe restraints on the possible cosmological models. The conclusion that
appears to emerge is that matter and antimatter must be separated on the
scale of clusters of galaxies if the universe really is symmetric. Much of what
I am going to present has already appeared in print so I shall limit myself to
a general discussion, omitting the details which may be found in the original
papers. (Steigman, 1969;Steigman, 1971;Steigman and Strittmatter, 1971;
Steigman, 1972).

DIRECT EVIDENCE

In principle it is easy to detect the presence of antimatter. You travel to
where you suspect a concentration of antimatter, put your detector down
(the most rudimentary device will do), and watch. If your detector disappears
then you better get out of there fast; you have detected antimatter. Seriously
though, just such experiments have in fact been performed within the solar
system via the manned flights to the moon and the unmanned probes to
Venus. Now we know, as we suspected with very good reason, that the moon
and Venus are made of ordinary matter. Even before the days of space flights
we had pretty good reason to believe the solar system was all made of
ordinary matter; the solar wind which sweeps out from the sun past the
planets acts as a probe just as our detector would.

Unfortunately, we are not likely to learn very much about a sizable part of our
galaxy by this method. However, we are fortunate that rather than having
to travel around ourselves there are obliging particles which come to us: the
cosmic rays. Now, unfortunately, the cosmic rays give us no information
about their sources because (except for the very highest energy cosmic rays)
they are tied to the magnetic field and do not travel in straight lines.
Therefore, we cannot be sure of what region of space we are sampling when we
examine the cosmic rays. However, we can be certain that, despite extensive
searches, no antinucleus has ever been found in the cosmic rays. Now, at
some level (~1 part in > 104) we would expect to detect secondary antiprotons
in the cosmic rays. The secondary production of antihelium or heavier anti-
nuclei in collisions between the cosmic rays and the interstellar gas will be
down by many orders of magnitude. These antinuclei would provide, if
detected, clear evidence that somewhere in the galaxy (universe?) there
were large amounts of antimatter. Evenson (1972) has set limits to the
fraction of helium nuclei which are antihelium. No antihelium nucleus
has been found and at the 95 -percent confidence level he finds a fractional
limit for the rigidity range 1 to 10 GV, of 1 X 10"3 and for the range 10 to
25 GV, of 8 X 10"2. For heavier antinuclei, a limit at the 95-percent
confidence level has been set by Golden et al. (1973, private communication),
for rigidities 4 to 125 GV, of 5 X 10"3, by Buffington et al. (1972), for
rigidities < 33 GV of 2 X 10^ . In the range 33 to 100 GV their limit is
2X 10"2.

ANTIMA TTER IN THE UNIVERSE? 363

As I emphasized, we cannot be sure where the observed cosmic rays come
from. From the ratio of light (Li, Be, B) nuclei to medium (C, N, 0) nuclei
we know that the cosmic rays must be able to travel several hundred parsecs
in a few million years. So the cosmic rays we sample probably come from
a volume whose typical dimension is roughly a few hundred parsecs. They
may in fact come from a much larger volume. The isotropy of the cosmic
rays, the smoothness of the distribution of galactic, nonthermal, radio
emission, and the relative constancy of the cosmic ray flux at earth over periods
as long as 4.5 billion years all indicate the cosmic rays we observe fill a
volume comparable in size to and perhaps even greater than our galaxy. The
lack of antimatter in the cosmic rays gives us good evidence that every second
star in our galaxy is not made of antimatter. Indeed, the limits on antinuclei in
the cosmic rays are already so low that even if a small fraction (perhaps one
percent or so) of them were extragalactic in origin, they would be telling us
that very few, if any, extragalactic systems could be made of antimatter.

In summary, the cosmic rays provide us with the only practical means of
sampling the universe outside our solar system. The evidence is straightforward:
no antinuclei have ever been found in the cosmic rays. Therefore, some region
of space contains very little, if any, antimatter. Unfortunately, we cannot
be certain which region of space it is.

INDIRECT EVIDENCE

When matter and antimatter meet, they annihilate. The annihilation products
are typically pions; there are roughly 5 to 6 charged and neutral pions in
a typical annihilation. The charged pions decay into muons with the emission
of a muon neutrino; the neutral pions decay most often into two 7-rays. The
muons themselves decay into electrons (and positrons) with the emission
of both an electron neutrino and a muon neutrino. The end products of a
typical annihilation are high-energy electron-positron pairs, 7-rays, and two
kinds of neutrinos. We can therefore hope to learn of the presence of
antimatter indirectly by detecting the products of its annihilation with
ordinary matter. The electron -positron pairs will probably not travel very
far from where they are created either because they will be tied to magnetic
fields or because they will scatter on any photons present (starlight, infrared,
blackbody, and so forth) and lose energy rapidly. Furthermore, we know
there exist mechanisms for accelerating electrons and positrons to high
energy in any case (pulsars). Hence the electron-positron component of
annihilation is not likely to provide us with any unambiguous information
about the presence of antimatter.

Neutrinos, of course, are very difficult to detect. As a result, large fluxes are
required; hence, the limits one might set are not very interesting. A major

364 COSMOLOGY

fraction of the matter in the universe would have to be annihilating before a
detectable flux of neutrinos would be produced. If that were the case there
would be other, more immediate consequences. Of course, a strong, nearby
source (for example, the galactic center) might produce a detectable flux of
neutrinos, but there too we would expect other, more obvious effects (for
example, 7-ray emission). For a discussion of these questions see Steigman
and Strittmatter, 1971.

Finally, we come to the 7-rays produced in annihilation. It is of course most
appropriate that they be discussed at this conference. A typical annihilation
produces a spectrum of 7-rays extending from several tens of MeV to several
hundred MeV. On the average, 3 to 4 7-rays are produced per annihilation.
Observations of ~ 100 MeV 7-rays then enable us to place limits on the amount
of contemporaneous annihilation.

The OSO-3 observations (Kraushaar, et al., 1972) of ~100-MeV 7-rays indicates
a galactic component superimposed upon an isotropic, presumably extragalactic
component. From their results, we can draw the following conclusions
(Steigman, 1969; Steigman, 1971 ; Steigman, 1972). If there is a cool, neutral,
intergalactic gas that is symmetric, its density could be no larger than n ~ 10"11
cm"3 . I remind you that the average density of matter in galaxies is ^ 1 0"7
cm"3 ; hence such a cool, intergalactic gas would constitute a minor component
of our universe. For a hot, ionized, intergalactic gas we find that if it is
symmetric, then its density must be low (< 10"9 cm"3). If, in fact, there is a
hot, intergalactic gas whose density is close to the critical density, the fraction
of it which could be mixed matter and antimatter would be less than one part
in 108. Thus, either such a gas is not symmetric, or it maintains very well
separated regions of matter and antimatter. While on the subject of inter-
galactic gas, it is worth pointing out that the Coma cluster of galaxies has been
detected as an X-ray source (Gursky et al., 1971) whose spectrum is interpreted
as thermal bremsstrahlung radiation from a hot intracluster gas. If this inter-
pretation is correct, then from the lack of 7-rays from Coma, we can say that
less than one part in 104 of that gas is antimatter.

The observations of the galactic, 7-ray component indicates an annihilation
rate per interstellar hydrogen atom of less than 10"25 s"1 . If, in fact, these
7-rays are interpreted as annihilation products, we can set the following limits
on the antimatter component in the galaxy: If the annihilation occurs in
interstellar clouds, less than one particle in 1016 is an antiparticle; the annihi-
lation occurs in the intercloud medium, the limits are less than one in 1012.
Indeed, it is worth pointing out that an antiparticle will only survive ~30 years
in an interstellar cloud and ~300,000 years in the intercloud medium; both
times are very short compared to the age of the galaxy (~1010 years). Hence,
it is clear that any model that requires the galaxy to be symmetric must find

ANTIMA TTER IN THE UNIVERSE? 365

an extremely efficient mechanism which keeps large amounts of matter and
antimatter very well separated over long periods of time. The most straight-
forward interpretation is of course that the galaxy probably contains no
macroscopic amounts of antimatter.

Finally, a word about 7-ray sources. There have been no detections of ex-
tragalactic 7-ray sources at about the level of ~10'5 photons/cm2 /s. If we
wish to use annihilation as an energy source for some of the more spectacular
extragalactic objects (for example, QSOs, Seyfert galaxies, radio galaxies, and
so forth) then we predict that they would be 7-ray sources. The lack of
detections of any of them as sources sets severe restraints on such models.
Either annihilation has nothing to do with these sources or, somehow, the
7-rays are absorbed at the source. This latter suggestion is not unreasonable.
However, it should be remembered that twice as much energy is released in
7-rays as in electron-positron pairs in a typical annihilation. Then we must
inquire into the effect on the source if these 7-rays are to be absorbed. Will
the absorption result in reradiation in another part of the spectrum? Can
such a model be made consistent with all observations?

CONCLUSIONS

We have been discussing the means of detecting the presence of antimatter in
the universe. We have seen there are several, straightforward, observational
tests and all have, thus far, proved negative. The most straightforward inter-
pretation of these results is that the universe is, in fact, not symmetric. Of
course, it is possible the universe is symmetric, but the matter and antimatter
are well separated from each other. Choosing between these two possibilities
must, of course, be a personal decision. Perhaps, in making this decision, we
should all bear in mind a quotation which sits, framed, on the desk of William
A. Fowler at Caltech. He attributes it to, "Proverbs for Graduate Students,
1933." It reminds us that, "The terrible tragedies of science are the horrible
murders of beautiful theories by ugly facts."*

*In the discussion following my talk, D. Clayton of Rice suggested that we search for
the evidence of annihilation by looking for the 1-GeV 7-ray line formed when nucleons
and antinucleons annihilate directly into two 7-rays. This purely electromagnetic channel
should occur but only very infrequently when compared to the strong interaction channels
via mesons. A rough estimate indicates only one in ~10 to 10 annihilations will be of
the two 7-type. The two 7-annihilation has been searched for, unsuccessfully, in several
experiments (Gursky et al., 1971 ; P. Nemethy, 1973, private communication). As a
result I do not expect a detectable ~l-GeV annihilation line even if all the observed
~ 100-MeV 7-rays are from annihilation.

366 COSMOLOGY

REFERENCES

Buffington, A., L. H. Smith, G. F. Smoot, L. W. Alvarex, and M. A. Wahlig,
Nature, 236, p. 335.

Evenson, P., 197 '2, Astrophys. J. , 176, p. 797.

Gursky, H., E. Kellogg, S. Murray, C. Leong, H. Tananbaum, and R. Giacconi,
197 '1, Astrophys. J. Letters, 167, p. L81.

Kraushaar, W. L., G. W. Clark, G. P. Garmire, R. Borken, P. Higbie, V. Leong,
and T. Thorsos, 1972, Astrophys. J. , 177, p. 341 .

Steigman, G., 1969, Nature, 224, p. 477.

, 1971, Proc. Int. Sch. Phys. Enrico Fermi, R. K. Sachs, ed.,

Academic Press, Course XLVII, p. 373.

, 1972, Cargese Lectures in Physics, 6, E. Schatzman, ed.,

Gordon and Breach.
Steigman, G., and P. A. Strittmatter, 1971 ,Astron. and Astrophys., 11, p. 279.

Chapter XV

A. GAMMA-RAY BACKGROUND SPECTRUM
AND ANNIHILATION RATE IN THE
BARYON-SYMMETRIC BIG-
BANG COSMOLOGY

J. L. Puget*

Observatoire de Meudon

INTRODUCTION

The negative results of the search for antimatter nuclei in cosmic rays imply
that if there is symmetry between matter and antimatter in the universe, each
kind must be gathered in separated regions of galaxy or galaxy-cluster size. In
such a case, in order to try to get experimental information on the problem of
baryon symmetry on a cosmological scale we have to rely mostly on the obser-
vation of annihilation products. Among the annihilation products are 7-rays
and neutrinos that have very long mean-free paths. Neutrinos especially can
reach us from dense regions; Steigman and Strittmatter (1971) used upper
limits on the neutrino flux from space to put upper limits on the annihilation
in Seyfert galaxies. Nevertheless, for the diffuse background due to annihila-
tion on a cosmological scale, 7-rays are the best test available because they are
easier to detect than neutrinos.

Two kinds of 7-rays are produced in matter-antimatter annihilation; 0.511-MeV
7-rays from positron annihilations (4.81 per annihilation); 70-MeV 7-rays from
7T° decay (3.4 per annihilation). The number of 7-rays of each kind is roughly
the same, and to compare them as a possible source of information on the
annihilation rate we must look at their absorption cross section and also at
the background due to other sources.

The absorption cross sections are respectively 10"25 and 1.8 X 10 cm for
0.5 and 70-MeV 7-rays. In a dense universe there is a "window" between
1 MeV and 10 GeV in which 7-rays observed might come from a red shift of
about 100 (see Stecker, Chapter IX. A). The X-ray background between

^Speaker.

368 COSMOLOGY

40 keV and 1 MeV can be represented by a power law with a spectral index
2.1 , so it is more likely to detect the 70-MeV annihilation 7-rays than the
0.5 MeV 7-rays.

These considerations prove that the best direct experimental test for presence
of antimatter on a cosmological scale lies in observations of the 7-ray back-
ground spectrum between 1 and 70 MeV.

EXPERIMENTAL DATA AND RED-SHIFTED GAMMA-RAYS FROM
ANNIHILATION

It has been shown by Stecker, Morgan, and Bredekamp (1971) that the excess
of 7-rays observed above 1 MeV could be explained by annihilation of 7-rays
coming from high red shift. They computed the spectrum (see Stecker,
Chapter IX.A) using a simple theoretical model for the annihilation rate
dependence on the red shift

6.36

* = * (1 + z)

v v,o v

(where ^ is the annihilation rate per unit volume and z is the red shift) and
they chose the constant ^ to fit the data. That led them to the conclusion,

J v,o

already found by Steigman (1969), that matter and antimatter cannot be

mixed up in equal quantities in intergalactic space with a density larger than

10"12 cm'3.

The most recent data (see Chapters III. A and IV.C) show very good agreement
with the spectrum computed by Stecker et al. (1971), and leads us to a detailed
discussion of the annihilation rate for red shifts lower than 100. The theoreti-
cal spectrum below 70 MeV but above ~ 5 MeV (where absorption is negligible)
is a power law with an index (m - 3.5) if the annihilation rate is written in the
form

* =* (l+z)m

v v,o v '

for

12 = _2_=1
ncrit

(n . is the so-called critical density of the Einstein-de Sitter model) and
(m - 3) for n = 0.

ANNIHILA TION RA TE IN THE BIG-BANG COSMOLOG Y 369

To get a good fit of the data one needs a spectral index of the order of 3 which
means that m must be such that

6 ^ m ^ 6.5

so one can consider that the annihilation rate will fit the data if it falls in the
range

*v = 10-34 * 0.5 (1 + z)6.2S ± 0.25 §-l cm-3 (XV.A-1)

MATTER-ANTIMATTER COSMOLOGY: THEORY

In the recent years, a baryon-symmetric cosmology has been developed in the
framework of the big bang theory of the universe and is summarized in these
proceedings by Omnes and Schatzman (Chapters XIV. A and B). In this model,
matter and antimatter separate at an early stage and at the end of the coales-
cence period (which coincides with the recombination time) forming an emul-
sion of characteristic size given by

L = 5 X 1029 (1 + z) 6 cm (1 + z > 600)

The fluid motions induced by the coalescence process on a scale of the order
of L reach a velocity

V~-=8.3X 1011 (1+zT1-34 cms'1 (l+z<600)
t

I want to discuss now what could happen in such a model after recombination
(which takes place around 1 + z ~ 600) in order to discuss the problem of the
annihilation rate. The following theory has been worked out by Stecker and
Puget (1972). The evolution of the characteristic dimension L of the emulsion
as a function of red shift is plotted in Figure XV.A-1. At the time we wrote
our original paper (Stecker and Puget, 1972), the theory of coalescence in the
radiative period had not yet been completely worked out, and we developed a
simple model in terms of cloud collisions to put upper and lower limits on L.
Recent work (Aldrovandi et al., preprint) allows us to plot the value of L up
to the recombination red shift, and the corresponding fluid velocities induced
by coalescence (Figure XV.A-2). One can compute the Reynolds number
corresponding to those coalescence motions and see that large-scale turbulence
is generated near recombination.

In a matter-antimatter symmetric big bang, the annihilation electrons and posi-
trons produce a large flux of X-rays by interaction with the cosmic blackbody

370

COSMOLOGY

Figure XV.A-1. The different lengths relevant to
the problem plotted as a function of red shift. XQ is
the mean free path of thermal photons; X (10%) the
mean free path of X-rays corresponding to 1 0 percent
ionization rate; and X (50%) the mean free path of
X-rays corresponding to 50 percent ionization rate.

photons, and these X-rays tend to keep the matter ionized longer than in a
nonsymmetric big bang. Furthermore, the recombination occurs very gradually
and ionization remains high near the boundary regions, as shown on Figure
XV.A-3. The viscosity which was determined by the radiation field drops
to the kinematic viscosity which is 10 orders of magnitude lower when
matter (or antimatter) becomes neutral and decouples from the radiation
field. The large-scale fluid motions then become supersonic. In order to
compare the parameters of the annihilation-generated turbulence with the
parameters of primordial turbulence used by Ozernoy et al. (1970) and
Ozernoy (1971), in their theory of galaxy formation, we have neglected the
remaining ionization after a red shift of ~ 600 in a first step. We find a good

ANNIHILA TION RA TE IN THE BIG-BANG COSMOLOG Y

371

Figure XV.A-2. The velocities relevant to the problem plotted as a function

of red shift.

agreement taking account of the uncertainties in the theory of generation
of turbulence.

I want to underline here the differences between the symmetric model and
the nonsymmetric one. Dallaporta and Lucchin (1972, preprint) have shown
that it is likely that a primordial turbulence will be dissipated before recom-
bination. In our model, turbulence is generated near or even during recom-
bination, so this problem disappears. The question of dissipation during the
phase of supersonic turbulence (before galaxy formation) and after galaxy
formation might also be a very serious one as shown by Silk (1972, preprint).
In the original model we just assumed for simplicity that no coalescence at
all takes place after z ~ 600. In fact, a source of motion exists. The ioniza-
tion near the boundary shown on Figure XV.A-3 which is due to photoioni-
zation collisions implies that the momentum carried away by these X-rays is
transmitted to the matter with a mean-free path which is of the order of
the width of the ionized region. We are in a case where the annihilation
pressure generates a surface tension of the type discussed by Omnes and
coworkers (see Omnes, Chapter XIV. A). This surface tension, which induces
coalescence during the radiative period, will also take place here and even if
the corresponding increase of size is negligible (which is certainly true for
low z as we shall see later), the fluid motions induced will compensate the
dissipation of kinetic energy, at least partially.

372

COSMOLOGY

[P]
PI

l+z=3.2

IONIZATION DUE
TO X RAYS

IONIZATION DUE TO r RAYS y >

"~FOR 1+z=320" "\

DISTANCE FROM THE BOUNDARY (cm)

Figure XV.A-3. The ratio of the proton density to the neutral
hydrogen density given for three values of the red shift as a
function of the distance from the annihilation layer.

The theory of the galaxy formation period, which includes such phenomena
as galaxy formation from the density fluctuations induced by shocks in the
supersonic turbulence generated at recombination time, formation of clus-
ters by the breaking up of the emulsion into separate clouds, and production
of magnetic fields on the boundaries between matter and antimatter, is
obviously a very complicated problem and it is not possible at this point to
rely on a complete theory of this period to discuss the annihilation rate.

I shall now change my point of view and, keeping in mind the general pic-
ture, give a detailed discussion of the annihilation rate based on the consis-
tency of our arguments with the observations on one hand, and on the

ANNIHILA TION RA TE IN THE BIG-BANG COSMOLOG Y 373

elements of our theory which have been worked out so far on the other
hand.

ANNIHILATION RATE AT Z « 1

As we have seen, the theory does not tell us if the regions of matter and
antimatter are of a galaxy cell size or of a galaxy -cluster cell size, so I shall
consider both hypotheses. If dense clusters contain as much matter as
antimatter there will be several sources of annihilation. I shall consider
these sources without going into the details, considering only the conclu-
sion we shall be lead to.

Intergalactic Gas

Observations of diffuse sources of X-rays in 20 rich clusters show that a
hot intergalactic gas, containing about as much mass as the galaxies them-
selves, must exist in clusters. This intergalactic medium must form an emul-
sion of matter and antimatter and, considering the magnetic fields produced
on the boundaries, the diffusion can be slowed down to a level such that
the annihilation rate does not exceed the value given by Equation (XV.A-1).

Galaxies (or antigalaxies) and Intergalactic-Gas

The velocities of galaxies (or antigalaxies) in a rich cluster are large (up to
103 km/s), and the crossing time for a galaxy is smaller than the age of
the universe, so a galaxy could be surrounded by matter or antimatter with
equal probability. Accretion of intergalactic gas on large galaxies will pro-
duce an annihilation rate

where M is the mass accreted by all the galaxies in one cluster per year.
Therefore, M must be smaller than 10"4 M@ so as not to conflict with the
annihilation rate given by Equation (XV.A-1). This value seems too small.

GALAXY-ANTIGALAXY COLLISIONS

Detailed study of galaxy-antigalaxy collisions have shown that the annihi-
lated mass is probably of the order of magnitude of MA with

MA =MT

A T c

1.25 X 10
* - =

v,gg T

374 COSMOLOGY

where r is the average collision time for one cluster, V is the relative
velocity, MT the total interstellar gas mass of the galaxy; any evaluation
of the collision time gives r «. 1018 which means again that the annihi-
lation rate from such a process would produce more 7-rays than observed.

If we consider cluster size regions, the annihilation takes place only on
boundary regions and, even for a dense intergalactic gas, magnetic fields
slow down the diffusion enough to bring the annihilation rate below the
rate given by Equation (XV.A-1) (Puget, 1971).

In conclusion, we shall make the hypothesis that clusters and groups of
galaxies are of matter only or antimatter only; this gives us the present
value of L:

L0=2.5X 1025cm

Considering that for low z, L is changing only with the expansion of the
universe because, for coalescence to take place, the fluid motions must
be such that

Vf>Vexp=7=4X1070+Z)'/'

V is shown in Figure XV.A-2 from which it is clear that no significant
coalescence can take place for z < 200 because the expansion velocity is
then much larger than the maximum fluid velocity which we can expect.

We shall use L = LQ (1 + z)"1 up to (1 + z) ~ 200 and L = 5 X 1029
(1 + z)"17/6 for 1 + z > 200. (We must nevertheless keep in mind that this
last relation has not been fully justified for 200 < (1 + z) < 600 when the
regions far from any boundary are neutral.)

ANNIHILATION RATE FOR (1 + z)< 100

There is some observational evidence that cluster formation occurs at
rather low red shifts. In our picture, the depression of density on boundary
regions becomes deeper and larger and eventually gravity overcomes expan-
sion and bound clusters are formed. We shall neglect this process here
because other processes like ionizing radiation from quasars or young
galaxies for z < 3 also modify the picture.

Let us study the motion of the plasma. For that purpose we need to find
how the anisotropy and the temperature gradient affect the motion of the
plasma. Physically, due to the importance of Thomson collisions of the
electrons of the plasma with blackbody photons and with the X-rays and
7-rays produced in annihilation, we examine the motion of the plasma on

ANNIHILA TION RA TE IN THE BIG-BANG COSMOLOGY 375

each side of the annihilation layer at distances much smaller than the mean-
free path of thermal photons. Technically, we write the Boltzman equation
for the photons and integrate it to get the equations of momentum conser-
vation and energy conservation, to which we add the equations of motion
of the plasma. These three equations have four unknown quantities: the
temperature gradient, the anisotropy of the photon distribution, the
velocity, and the density of the plasma. We can eliminate the first two in
order to get an equation of motion of the plasma which has to be combined
with the continuity equation,

3n nv f L NV n - nQ

— nv + NV (1 - §) -e-u - v2 — - — dx = 0 (XV.A-2)

dt \ 3x rD I LrQ nQ

where v is the plasma velocity, n the plasma density,

n0 = < n > , N = n (x = 0) , V = v (x = 0) , u =

A C

£ is the fraction of the momentum of the X-rays which is transmitted to
the blackbody photons, v is the thermal velocity of the plasma, and

"radiation

k =

^matter

During the radiative period, the second, third, and fifth terms of this
equation of motion of the plasma are negligible for distances smaller than,
or of the order of

SS = n/2t^v2 t

D s

which is the distance over which the density gradient extends. (rD is the
characteristic time for slowing down of charged particles by the radiation
field.) The equation is then a simple diffusion equation and the solution
for n is:

n(x,t) =

V 47TTI,V2 t
D s

376
and

NV = (n v)

n0Vs/TDV/2

z=0

2,J¥\t

COSMOLOGY

n(x) and v(x) are plotted on Figure XV.A4.

v/v

1 >v^_i 1 1 1 1 1 1

b£

5£

ior

n/n0

.5n0

1 1 1 1 1 1 1 1 1

lor

Figure XV.A-4. The density and velocity of the
plasma given as a function of the distance from the
boundary. The unit for the v scale is v = v (rD/7rt)

The annihilation rate is then given by

* =

= 7.3Xl(T29(l+z)61/12

The fifth term of Equation (XV.A-2) corresponds to the anisotropy of the
photon field inducing a heat flow which dissipates the excess energy left
by X-rays and 7-rays in the regions where n is larger than the average
density (nQ). For x ».^, it is the dominant term for the motion of the
plasma, but it does not affect the annihilation rate in a noticeable way.

ANN MIL A TION RA TE IN THE BIG-BANG COSMOLOG Y

377

When (1 + z) becomes lower than 1 .4 X 103 , the third term in Equation
(XV.A-2) which is the flux of momentum from X-rays to the plasma,
becomes as large as the fourth term which is the pressure gradient of the
plasma. The annihilation rate then is given by

rD\ * e

-A'

^ =— n„ v —

v L

0 s

2JtT

L8k

(1 -Erf A)+ 1

with

A =

r t

NV 1
no Ao

(4^Dvs2t)

-V4

where

r x

Erf(x)

e-u du.

(Aldrovandi et al., 1973 preprint)

For % small enough to be negligible compared to 1 , A is in fact almost a
constant = 4.2

and

NV ^lT^_Y -A2

n0 vs 3^ \ t /

Thus,

^ =1.7X 10'32 (1 + z)73/i:

This solution is valid down to (1 + z) ~ 600. Below that value it breaks down
for two reasons having opposite effects:

• L might increase more slowly than (1 + z)-17^6 due to recombination;

• The mean-free path of the X-rays produced by the annihilation-
generated electrons and positrons, which was equal to X, becomes
much shorter due to the large photoionization cross-section as
shown on Figure XV.A-1.

378

COSMOLOGY

The equation of motion of the plasma in the vicinity of the boundary remains
the same because of the ionization due to X-rays (see Figure XV.A-3)* but
the momentum left in the plasma per unit time and unit volume is now pro-
portional to X"1 instead of X^1 . (Xx is taken equal to the distance from the
boundary for which n /nH = 1 .) Figure XV.A-4 shows the density and
velocity of the plasma as a function of distance from the boundary in units
ofvc=vs(rD/rrt)*.

The equation giving the annihilation rate has to be solved numerically. The
result is shown in Figure XV.A-5*, which gives the annihilation rate as a
function of red shift compared with the rate given by Equation (XV.A-1).

-10

-14

i to

■E

-18

/ J

-22

/ /

—l

^//

-26

1—

~ /w

XlEUTRAL

PLASMA

*•

-30
10

-34

- / PERIOD

PERIOD

* i i

• * i

2 3 <

10 10 10 10

d+2

)

Figure XV.A-5. The annihilation rate is given
as function of red shift.

Considering the uncertainties in these calculations, the agreement is as good
as can be expected. The major uncertainties affect the rate for the range
(1 + z) < 5 and cannot affect the 7-ray spectrum between 1 and 1 5 MeV

*The results given on Figures XV.A-3 and XV.A-5 are from preliminary numerical
evaluations; exact numerical computations will be published later.

ANNIHILA TION RA TE IN THE BIG-BANG COSMOLOGY 379

very much. Furthermore, the 7-ray flux must fall off for energies above
70 MeV and below 1 MeV, so the theoretical spectrum is quite well defined.
If the good agreement of this spectrum with the data is confirmed by
future measurements, a way of checking this model will be to look at
angular fluctuation of the background as a function of energy. The 7-rays
observed at the energy E come mostly from a red shift (1 + z) ~ 70/E (Meyr.
and the angular fluctuations will be related to L(l + z).

REFERENCES

Ozernoy, L. M., G. V. Chibisov, 1970, Astr. Zh. , 47, p. 749 (English
translation in Soviet Astro. A. J. , 14, p. 91 5.

Ozernoy, L. M., \91 \ , Astrophys. J. Letters, 7, p. L201.

Puget,J.L., 197 1, Nature, 230, p. 173.

Stecker, F. W., D. Morgan, and J. Bredekamp, 1971, Phys. Rev. Letters,
27, p. L1469.

Stecker, F. W., and J. L. Puget, 1972, Astrophys. J. , 178, p. 57.

Steigman, G., 1969, Nature, 224, p. 477.

Steigman, G., and P. A. Strittmatter, 197 1, Astron. and Astrophys. , 11,
p. 279.

B. DISTORTION OF THE MICROWAVE BLACK
BODY BACKGROUND RADIATION IMPLIED
BY THE BARYON-SYMMETRIC COSMOL-
OGY OF OMNES AND THE GALAXY
FORMATION THEORY OF
STECKER AND PUGET

F. W. Stecker*

Goddard Space Flight Center

J. L. Puget

Observatoire de Paris

One consequence of the baryon-symmetric cosmological model of Omries is the
continuing annihilation of matter and antimatter throughout all stages in the
evolution of the universe. This annihilation can cause a distortion in the
microwave blackbody spectrum from a purely thermal spectrum because of
deposition of annihilation energy at red shifts less than 104 and particularly
at red shifts less than 103 . The theory of this distortion was first discussed
by Zel'dovich and Sunyaev (1969; see also Sunyaev and Zel'dovich 1970a,
1970b). They show that because of the varying evolution of the optical
depth of the universe to radiation at various wavelengths and because the
Compton process conserves photon number and does not lead to pure thermal-
ization, two different distortions arise in the blackbody spectrum. Distortions
in the Rayleigh-Jeans (<* v2) portion of the spectrum are due to energy deposi-
tion at red shifts between 104 and 103 (Zel'dovich, Illarionov, and Sunyaev,
1972). Distortions in the Wien portion of the spectrum (<*e-v) are due to
energy deposition at lower red shifts after the cosmic gas cools to its atomic
state and thermalization does not take place as efficiently.

In order to quantitatively estimate the expected distortions, we define the
parameter

^Speaker.

381

382 COSMOLOGY

■/

Ae(t)

— -fdt (XV.B-1)

which is a measure of the maximum fraction of the energy density in the
radiation that contributes to the nonthermal part of the microwave back-
ground.

In Equation (XV.B-1), e(t) is the energy density in the blackbody radiation as
a function of time (or red shift z, where t = t(z)).

It then follows that

* (z)M c2 Ht
v p dt (XV.B-2)

e (z) dz

where ^ (z) is the annihilation rate function discussed in the main paper by
Puget (Chapter XV.A).

For the red shift range 600 < z < 104, the annihilation rate is given by

*v(z) = 1.7 X 10"32 (1 + z)6+1/12 cm"3 s"1 (XV.B-3)

(see Puget, Chapter XV.A). The resulting value of qR_j affecting the Rayleigh-
Jeans part of the spectrum is then

qRJ = 1.2X 10"2 (XV.B-4)

which is, in fact, an upper limit because part of the energy goes into large-scale
fluid motions. The corresponding distortion in the Rayleigh- Jeans part of the
blackbody spectrum is of the order of 2 percent, well below the observational
uncertainties of about 20 percent in the wavelength region greater than 1 cm.

For the red shift range z < 600, we will adopt the annihilation rate fitting the
7-ray observations (that is, the largest value consistent with the present obser-
vations (see Puget, Chapter XV.A)).

#v(z) a 10"34 (1 + z)6-25 ciTfV (XV.B-5)

The resulting value of q affecting the Wien part of the blackbody spectrum is

q S6X10"5 (XV.B-6)

DISTOR TION OF THE BLA CKBOD Y BA CKGROUND 383

This may be related to the parameter y used in the calculations of Zel'dovich
and Sunyaev, since

y = q/4 = 1.5 X 10"5 (XV.B-7)

This is well below the observational upper limit on y set by Zel'dovich and
Sunyaev of 0.15.

In fact, we expect more distortion than indicated by Equation (XV.B-7)
because of dissipation of turbulence created at higher red shifts which feeds
energy into the microwave background below z = 600. To estimate this
effect, we have made a more detailed numerical calculation of the mean gas
temperature as a function of red shift (Stecker, Puget, and Bredekamp, in
preparation) and used the relation given by Zel'dovich and Sunyaev

f kTe

yT = noao cH^1 / dz __ll (1+z)1'2 (XV.B-8)

where n is the present mean gas density in the universe, taken to be 3 X 10"6
cm , a is the Thomson cross section, and H is the Hubble constant where
H0"1 s°6X 1017s. We then obtain from Equation (XV.B-8) a value of

yT ss 2 X 1 0^ « 0. 1 5 (XV.B-9)

Sunyaev and Zel'dovich discussed the problem of blackbody distortion due
to antimatter annihilation, but they estimated the annihilation rate without
taking account of limitations due to annihilation pressure on the boundary
regions between matter and antimatter. They therefore overestimated the
annihilation rate by a large factor (see Stecker and Puget, 1972).

Our conclusion is that the annihilation rate for our model of galaxy formation
(Stecker and Puget, 1972), while large enough to provide the turbulence
needed to explain galaxy formation (Stecker and Puget, 1972; Aldrovandi,
Caser, Omnes, and Puget, preprint), does not produce a distortion in conflict
with present observations of the microwave blackbody radiation.

REFERENCES

Aldrovandi, R., S. Caser, R. Omnes, and J. L. Puget, (Preprint LPTHE 73/5,
to be published).

Stecker, F. W., and J. L. Puget, 1972, Astrophys. J., 178, p. 57.

384 COSMOLOGY

Sunyaev, R. A., and Ya. B. Zel'dovich, 1970a, Astrophys. and Space Set,
7, p. 3.

, 1970b, Astrophys. and Space Set , 9, p. 368.

Zel'dovich, Ya. B., A. F. Illarioniv, and R. A. Sunyaev, 1972, Zh. Eksp. Teor.
Fiz., 62, p. 1217.

Zel'dovich, Ya. B., and R. A. Sunyaev, 1969, Astrophys. and Space Sci , 4,
p. 301.

SECTION 4

FUTURE DIRECTIONS

IN GAMMA-RAY ASTRONOMY

Chapter X VI

A. A PANEL DISCUSSION ON
THE FUTURE DIRECTION OF
GAMMA-RAY ASTRONOMY *

Giovanni Fazio

Harvard and Smithsonian Astrophysical Observatories

Carl Fichtel

Goddard Space Flight Center

Glenn Frye

Case Western Reserve

Kenneth Greisen

Cornell

Albert Metzger

Jet Propulsion Laboratory

Evry Schatzman

Paris Observatory

Floyd Stecker

Goddard Space Flight Center

Jacob Trombka

Goddard Space Flight Center

The following discussion was convened at the final session of the NASA Inter-
national Symposium and Workshop on Gamma-Ray Astrophysics to sum up the
present status of the field and discuss its future. The remarks below are based
on a transcript of this discussion. They are free and informal. Because exten-
sive editorial work was necessary, the editors apologize for any misinterpreta-
tion of remarks that may be present.

"The panel consisted of the following members.

387

388 FUTURE DIRECTIONS IN GAMMA-RA Y ASTR ONOMY

FICHTEL:

One thing that has impressed me is that, indeed, we do now have an observa-
tional science. I think we were in an awkward position a few years ago when
there was really very little to talk about. I think we now not only have actual
results that are coming forward from many different areas, but also see that
there is a tremendous wealth of data that should be forthcoming in the near
future. I am not going to give a speech, so I am going to stop in just a moment,
but I will say I do think there can be a tremendous interaction between
experimenters. I know Klaus Pinkau's group and ours always talk in terms of
trying to exchange information very quickly so we can use this to feedback
and look at things. And I certainly am willing to try to work with the people
doing balloon experiments.

I know, for example, there is tremendous interest in very high energy region
of about a GeV, and we will work with Ken Greisen on research in this region.
We have a very great stimulus now, we will see the field expand, although
perhaps not quite so rapidly as X-ray astronomy, but certainly very rapidly.
And now I think the theoreticians deserve a chance.

FAZIO:

I would like to say a few words about the importance of getting the SAS
(Small Astronomy Satellite) results. I think a good part of the future of
7-ray astronomy depends on what comes out of this satellite. That is, whether
it becomes a growing, ever more fruitful area or whether it just dies greatly
depends on what is produced. So I would like to make a call to get out the
data as quickly as possible since many of the balloon groups are very dependent
on it for their future planning.

FICHTEL:

Needless to say, we are working very-extremely-hard (with the SAS data).
I am certainly willing to cooperate on an informal basis to let people know our
tentative results, if they want to take their chances on tentative results.

FAZIO:

I realize, just as anybody in this field does, the complications of trying to add
up data over several weeks. Indeed most of the time it is very, very difficult.

FICHTEL:

That kind of thing is possible; however, it is difficult to run the sensitivity
programs. They take a while. So for example, we may know something is
there but be reluctant for a while to say what it is.

PANEL DISCUSSION 389

FAZIO:

Yes, I fully agree.

SCHATZMAN:

I think that generally speaking, the theoreticians are very eager to see the
improvements in 7-ray astronomy, and by this I mean an extension of the
range of energies which are covered by the observations, particularly toward
higher energies. Also important is increasing the resolving power and precision
and maybe increasing the energy resolutions.

A couple of years ago I became interested in 7-rays produced in supernova
explosions (a subject discussed by Clayton).* If there were the production of
elements, spallation products of these atoms would be 7-radioactive and quite
a few of them would decay with a short radioactive period of the order of
seconds or days. Thus we have to consider the possibility of detecting a 7-ray
flash at the time of an explosion. This can be considered as feasible if the next
supernova does not explode too far away from us. If I remember the numbers
properly, if an explosion takes place at a distance of less than 10 megaparsecs,
we should have the sensitivity to detect the flash, assuming that the proper
instrumentation is in space at the time of the outburst and that it looks in the
proper direction. But the trouble is that at the time the explosion becomes
visible, it is likely that most of the radioactive elements will have already decayed
because the time between the beginning of the bursts and the observation of
maximum light can be several days. So at least in the case of a supernova of
Type I, which is one with a slow maximum, it is likely that by the time it is
optically visible, the 7-ray flash has disappeared.

But on the other hand, as far as nucleosynthesis is concerned, the issue can be
extremely important because presently the amount of nucleosynthesis which
takes place in supernovae is a matter of speculation. If we look for the 7-rays
coming from radioactive products produced by nucleosynthesis, they could be
observable.

TROMBKA:

Along these lines, maybe I am going to put some people on the spot. I think
one of the major developments that have to come about is the development of
solid-state detectors of high enough detection efficiencies in order to see the
line fluxes. We heard some reports on what is going on. Could we ask for a
little further information? What seems to be the direction of the Lockheed
group? What is going on in terms of the development of larger detectors and

"See Chapter XI. A.

390 FUTURE DIRECTIONS IN GAMMA-RA Y ASTRONOMY

the development of intrinsic detectors so that the mechanical problems
involved in flying these detectors will be easier?

NAKANO:

Intrinsic germanium is the way to go, as far as not having to maintain solid-
state detectors at cryogenic temperatures. But right now the intrinsic-ger-
manium detectors are not nearly as large as the drifted detectors. What is
needed is more detector sensitivity and some way to gang these detectors
together, something very similar to what Bud Jacobson of Jet Propulsion
Laboratory is doing. I would say, if one really wants the biggest sensitive
area one can get, that would mean using drifted detectors, at least for now.
But they are making strides on the intrinsic detectors. The biggest one that I
know of is a cylindrical detector of about 30 cm3 , but perhaps in a year or
two they will catch up in development to the other detectors.

FAZIO:

Is there a supernova burst mode in SAS-2?

FICHTEL:

Yes. The problem is a little different. I assume when you (Schatzman)
spoke of a flash of 7-rays, you meant nuclear lines. There is an interesting
concept suggested by Colgate that if the hydromagnetic shock-wave model
is indeed correct, there will be a photon pulse of very high energy 7-rays that
will come out which is very, very short— a tiny fraction of a second. We looked
at some of the experimental ways you might detect such an event. It turns
out with the atmospheric fluorescence experiments and the mode we suggested
for SAS originally, there is not much chance for detection of this flash. But
if you go to somewhat lower energies, I think there is a real chance to detect
these flashes, because you can possibly see at least ten times as far as the
Virgo cluster. You then have a reasonable chance to see such an event.

Bob Hartman and Mike Sommer came up with an interesting concept where
you have very large scintillators which can detect several pulses in a very short
time with very fast time resolution. This will in fact be on a balloon flight that
will be launched from Palestine, Texas, in the near future. This would be
another very interesting phenomenon for someone to look for in the flash
phase because this kind of flash is uniquely associated, as Colgate has shown,
with the cosmic-ray hydromagnetic origin theory. If you don't see one of
these flashes, then Colgate can't be right.

PANEL DISCUSSION 391

GREISEN:

The thing that I have been impressed with at this meeting is that despite fears
in the past that the cosmic 7-rays, like the cosmic rays themselves, might be
almost a featureless waste (a continuum) that one could not extract precise
knowledge from because of the lack of detailed features, it seems not to be so.
There seems to be an abundance of features and any single mechanism for
production of the 7-rays does not seem to be able to account for the whole
energy range. There seem to be a number of bumps that add up to produce
what in the first approximation is only a smooth spectrum, but on closer
examination has structure. That has also reinforced the dictum of the
Astronomical Society, which they have been insisting on over the years, that
it really is important to investigate the full range of the electromagnetic spec-
trum. The results observed, even in opposite extremes of the spectrum, tend
to be coupled through theoretical models very tightly with each other and
impose constraints on each other.

I think that in spite of the fact that each of us experimenters has a special
interest, and has to have one (one has an obligation to push for development
of opportunities for his particular little area), we should all remember the
importance of uniting as a group to support this multispectral concept. While
the great value of this has been shown, in particular in close looks at the low-
energy 7-rays, particular questions were also raised and fascinating hints were
shown in the data at lower energies that suggest it is also important to look
at the higher energy 7-rays, not only at the discrete sources, but also at the
continuum.

If, for instance, a substantial part of the radiation around 100 MeV comes
from annihilation processes, we know that that spectrum peaks at around
100 MeV and cuts off sharply beyond that, as mentioned by Floyd Stecker
earlier (see Chapter IX.A). If cosmological origin is important, it was made
clear that one can predict other features that should show up as distortions
of this spectrum.

Well, I have a feeling that we are almost marking sort of a birthday of 7-ray
astronomy. It is true it is not exactly the birthday today. There was a satellite
launched a little while ago, the SAS-2. There have been some forward steps
by means of balloons, but really the planned shuttle program should open up
a new era. I don't think this can be over-emphasized because one has to put
the emphasis on the word "should." It is by no means obvious that it will
open a new area, because the whole program might be spent in nonscientific
enterprises, but there is a possibility that we really could have many scientific
opportunities as long as money is provided for performing experiments.

The whole shuttle program will only be justified if there are many missions and
each mission should carry something worthwhile. And the worthwhile things

392 FUTURE DIRECTIONS IN GAMMA-RA Y ASTRONOMY

to get out in space are the things that will see parts of the universe that can't
be seen from the ground or to do experiments that you couldn't do down
here.

I think that high-energy astrophysics, ultra-violet, X-ray or 7-ray astronomy,
is not the only important thing to do on the shuttle, but it's one of the most
important. So I hope we gain opportunities in the near future to make this
whole subject achieve a real stature and not remain forever in an infant stage,
where it is now.

STECKER:

As a theoretical type, I would like to second your remarks about the bumps
and distortions in the continuum spectrum. It is the bumps and wiggles that
keep us in business. In this regard, I think one of the theoretical emphases
would be on getting better energy resolution at all energy ranges so we can
see and study the bumps and wiggles, particularly in regard to what they say
about the cosmological origin of the diffuse background.

This wasn't brought out, but if those 7-rays are cosmological, then there may
be another absorption effect at around 10-GeV energy.* We haven't done
too much with that energy range yet, but I hope we will keep extending our-
selves up into that energy range, and also I hope that some of these theoretical
points that are brought out from time to time by people like us will be
noticed by experimentalists, to the extent that it might help direct the inves-
tigations. As an example of this, I can cite the Apollo-15, -16 and -17 results
that were reported on.

In that detector experiment the primary purpose was to gather information
about lunar 7-rays and the detector did an excellent job, but the detector
was also kept on in transearth orbit to look at the cosmic background in a
theoretically critical energy range and we found some very surprising features
which may have exciting implications, which we heard about from Professors
Omnes, Schatzman, and Puget.

FRYE:

I would like to address my remarks a little more particularly to the time-gap
between SAS-2 and the shuttle program. And this naturally is when balloon
observations must be done from the top of the atmosphere. With many of us
interested in this energy region and this type of observation, I think it is a
worry to all of us that there is going to be an observational hiatus when the
SAS-2 experiment is over. I am assuming that it will operate for one or two
years.

"See Fazio and Stecker, Nature, 226, 135 (1970).

PANEL DISCUSSION 393

I would like to make a couple of comments about the balloon technique— what
one might be able to do in this period until the shuttle is operational.

One point is that in the very interesting region from 10 to 30 or 40 MeV, with
regard to the diffuse measurements, the flux that is reported for these measure-
ments is as large as the atmospheric background if you are at the order of
1 g/cm2 . So there is not as strong an argument for being on a satellite in this
energy region as one has for other energy intervals. One still has the very long
observation time from the satellite, but any experiment that looks at the Crab
in this region at the very minimum is going to have to contend with this back-
ground.

Therefore, one avenue that I think people should look at, and really take a long
look at, is trying to build larger-area detectors. With a large collecting area you
have better statistics per unit time. To take advantage of the advances in
balloon technology, either higher altitude or super-pressure balloons have
potentiality for many days of observation. Also, with respect to the Crab in
particular, one does need the very accurate sub-millisecond timing, so that
until again something of either the HEAO (High Energy Astronomy Observa-
tory) or shuttle capability comes along, these observations are going to have
to be made from balloons.

GREISEN:

I would like to ask a question. It is obvious down in the energy range related
to the HEAO 7-ray apparatus experiment utilizing solid-state detectors where
one is looking for particular energy transitions, that one really needs very fine
energy resolution. But in your (Stecker) statement calling for better energy
resolutions at higher energies, I know we need better ones because there is
essentially none there now. But I have the feeling that all the models that
were proposed failed to produce any very sharp features. The bumps were
broad bumps and so on. And I think it would be somewhat misleading if it
got to be stated as a requirement for these experiments that there be fine
energy resolution. Do you see any real need for energy resolution better than
something like a factor of two, or say 50 percent at energies above 30 MeV?

STECKER:

It is true that above 30 MeV you probably expect broad features, but we have
seen several sets of theoretical curves, and the curves have different curvatures.
And one would like, from the theoretical point of view and this is of course
somewhat academic, to determine as exactly as possible a good spectral shape
because, as was brought out in this conference, log -log spectra aren't always
straight lines. The shape of the spectrum can then become very important in
distinguishing one theoretical model as opposed to another, even when you
don't have a sharp feature .

394 FUTURE DIRECTIONS IN GAMMA-RA Y ASTRONOMY

GREISEN:

Yes, but you don't need to have very good determination of the energy for
each event in order to derive the shape if you study and understand your
instrumental response.

FRYE:

Certainly, if I could interject, within the 7r°-region of 100 MeV or so, one
would like at least 20-percent resolution to see if there is any feature that one
can ascribe to that.

FICHTEL:

I would like to completely support Floyd (Stecker), because I think we are
finding even now that if you are going to really measure a spectral fall-off,
for example, you really need something more like 10- or 20-percent energy
resolution rather than a factor of two if you are going to see a sharp fall-off
and see at what energy it occurs. So I think that argument is certainly true.
You are really going to have to do what Floyd (Stecker) wants.

STEIGMAN:

Another way to get the shape of the spectrum is not necessarily to have very
good resolution in any particular energy, but to extend the energy range over
which you can observe. Is it harder to extend the energy range or to build
better resolution detectors?

FICHTEL:

One of the problems in extending the energy range is a problem of sensitivity.
In other words, if you can measure the energy better, you don't need so many
photons, and ultimately we run out of photons at some energy.

GREISEN:

I think the question from the audience (Steigman) can't be answered in
general. It is different with each different apparatus and every different part
of the energy range. But ordinarily, the methods available for measuring
energy are very difficult to make precise. If one observes the scattering of
electrons in 30 plates of a spark chamber, one has a very limited sample of a
random distribution, and the mean is not very accurately determined.

I think it is well enough determined for these purposes, because that dip that
you are looking for is not so extremely sharp. At very, very high energies other
methods start to become available for measuring energy, either based on scat-
tering or something like transition radiation. But always it comes down to

PANEL DISCUSSION 395

detecting a small number of samples of a random variable. And if one really
felt pushed to make a 10-percent energy resolution throughout the energy
range, you might close down the whole business, Floyd, all the way.

STECKER:

I don't think 10 percent here is necessary.

GREISEN:

On the other hand, when you can use an intrinsic detector and stop a low-
energy particle, then it is possible to get 2-percent resolution, or something
like that.

METZGER:

If I could comment on the same subject, in the energy region that I was dis-
cussing in terms of the two HEAO 7-ray instruments, there is distinct contrast,
and one doesn't get something for nothing. With solid-state detectors one gets
the resolution and identifies the feature, hopefully, but at the same time, at
least in terms of present technology, the efficiency of the detection is much
less. So the discovery of the feature is not as well done with a solid-state
detector as with the more efficient sodium-iodide system.

Of course, this means that there is a logic in flying a HEAO mapping experi-
ment of the Peterson type utilizing Nal(Tl) detectors, first, to determine
where the promising areas to look might be. Later, one can study these areas
by making use of solid-state detectors with the high-resolution experiment
proposed by Bud Jacobson of JPL.

STECKER:

I want to throw out another related question to the experimenters. On Monday
we had a long discussion about problems with the intrinsic detector-produced
radiation in sodium-iodide crystal detectors, particularly in the region around
1 MeV, which is important theoretically for cosmological reasons for such models
as the matter-antimatter cosmology. Therefore, it would be nice from the
theoretical point of view to minimize as much as possible these detector
background problems. Although there were quite a few comments about past
and present experiments, I would like to hear some comments about minimizing
this problem in the future.

METZGER:

You couldn't have picked a tougher energy range.

396 FUTURE DIRECTIONS IN GAMMA-RA Y ASTRONOMY

TROMBKA:

Part of the minimization will be achieved after an understanding of the intrinsic
radiation. You run into other problems also. What you usually want is high
Z-type materials in order to stop the 7-radiation, to get good efficiency, and
so forth. In doing so, you are hurting yourself immediately in terms of the
spallation problem. I think the problem will have to be overcome by obtaining
a real understanding of the cosmic-ray -induced spectrum and then seeing
whether one can see keys within the spectrum itself, in order to interpret the
magnitude of this effect.

PIEPER:

Jack, it seems to me absolutely remarkable that the four experiments that have
been done in that region, the ERS, Ranger, and the two experiments on
Apollo, all agree in terms of the spectra of counts/cm2 -s"1 in pulse-height space.
All fall on top of each other with the exception of one single point, and the
problem really is to translate that spectrum from pulse-height space into pho-
ton space and that is where the experiment doesn't need to be done again.
What needs to be done is to learn how to interpret it.

TROMBKA:

I think one of the major efforts for Al (Metzger) and myself will be the study
of the spallation problem in detail. To this end, we did the Apollo-17 experi-
ment and hope to continue on Skylab with some other materials to get a
better understanding of it.

Another area I think that we have to work on, partially on HEAO, is deter-
mining the isotropy of the 7-ray flux. Along these lines, we at Goddard are
looking at the use of small satellites. It turns out, even with the spallation
products building up, one can look at isotropy with rather simple techniques.
I know Dr. Pieper and I are interested in this and some people from France
are also interested in this. There are a number of experiments that can be done
relatively simply. One can do rather meaningful experiments rather simply and
hopefully inexpensively. Our idea, and hopefully some of the others will
express theirs, is to use a pancake-shaped detector. If you remember from my
talk, the counting efficiency of a detector is strongly dependent on angular
distribution of the incident flux. By rotating a properly shaped detector in
an anisotropic field, one gets a significant variation in the count rate which
reflects the isotropy or anisotropy of the flux. Our calculations indicate that
a crystal of about 5 cm in height by 20 cm in diameter can detect anisotropics
of a little less than one percent. In order to minimize problems due to spalla-
tion and induced activity, the spacecraft can be surrounded so that the total
spacecraft mass is inside of an anticoincidence mantle.

PANEL DISCUSSION 397

SCHATZMAN:

In connection with the anisotropics there is the question of the angular fluc-
tuations in intensity. What seems critical for the antimatter cosmology is the
angle of 1/10 of a degree, and this is a fairly small angle. I understand that
presently you can expect to detect deviations over one or two degrees, but
this doesn't seem to bring any important information if the cosmological
model is correct. Further information could be obtained with better angular
resolution.

METZGER:

At what minimum energy would that angular resolution be required to be
useful?

SCHATZMAN:

It turns out, due to the fact that the angle for a given length depends on the
distance due to relativity effect, that the angular fluctuation you can expect
at high red shifts is constant. So from z > 2, let's say for example, to z = 100,
it is 1/10 of a degree.

MEMBER OF THE AUDIENCE:

I just wanted to make a comment on that point. For, let's say, 50 to 70 MeV
annihilation 7-radiation in the type of cosmological model we have discussed,
the type of fluctuation we can expect is a one percent fluctuation.

WHITE:

May I emphasize that in the region from 1 to 1 0 MeV, to detect 7-rays one
should make use of the process with the highest cross section that one has, and
that of course is the Compton process. And one should make use of the ma-
terial which is best for taking advantage of that cross section; therefore, one
would want to go to a low Z-material rather than a high Z-material.

Secondly, in order to get rid of the background, one can do a number of
things. One can, of course, make the angular resolution as good as possible.
This gets rid of background problems. One can also surround the detector
completely with anticoincident scintillators in order to get rid of charged-
particle background. One can get rid of the backward-going 7-rays by being
sure of the anticoincidence in backward direction. One can get rid of neutrons
by using time-of-flight techniques.

So there are a number of things one can do to get rid of background. In
addition, if one wants to get the total 7-ray energy, one can go farther than
people have gone today. One can use a two-Compton scatterer, and one can

398 FUTURE DIRECTIONS IN GAMMA-RA Y ASTRONOMY

make the second scatterer so big that the total 7-ray energy is deposited in it
and thus measure the total energy of the 7-ray. One then has the scattered
electron in the scatterer, so one obtains the total energy of the 7-ray. Then
one doesn't have the uncertainties that we discussed in the double-scatter
experiment. One can also get very high efficiencies by making the detector
very large. Therefore, in the region that you are talking about, near an MeV,
I recommend two-Compton-scatter telescopes very highly.

FRYE:

I might add to this that if one does make the first scatter in a visual detector
where you see the recoil, then you completely determine the kinematics and
this produces better angular resolution than one has in a system where you
just get the energy deposition. And we have such a system. I don't have any
data on it yet.

RAMATY:

I would like to make two comments. There is another astrophysical process
which leads to 7-ray lines which, even though not as efficient as nucleo-
synthesis, has an advantage that we are almost sure it is going to occur. This
process is induced by accelerated charged particles on ambient nuclei.

The other thing I would like to say concerns solar flares for which Floyd
Stecker said that everything is understood. It is true that the mechanism for
7-ray production in flares is understood, but what is important in solar flares
is to get an understanding of the magnitude of the solar flare. Only for one
of the flares, a big flare, were 7-ray lines seen. And I remind you that the
flux that was seen was about 0.1 photons cm"2 s"1 , which is much higher
than the fluxes of 10"5 cm"2 s"1 to be studied in future experiments. If the
threshold or sensitivity for detection of 7-rays is going to be lower, I think
one can in principle learn quite a lot about the mechanism of flares. Although
in this conference the main emphasis was on galactic and extragalactic
phenomena, I think that this point is worth mentioning.

STECKER:

I didn't mean to imply that everything was understood. I just meant the basic
physical processes were understood as opposed, for example, to some of the
other production processes discussed.

RAMATY:

I think in the case of the flares, we can take off from there and start under-
standing flares, because we don't really have to argue too much about the
radiation process.

PANEL DISCUSSION 399

STEIGMAN:

In other parts of the electromagnetic spectrum, theoreticians are somewhat
constrained by another piece of information that we have and that is
polarization measurement. Is there any hope whatever to get polarization
measurements in the several MeV energy range?

FRYE:

Well, it is possible in principle from the spark chamber where you observe the
pair and measure the polarization of the pair. So, in principle, the informa-
tion is there. It depends on the detector and the fluxes that you have. This
is one of the things that all of us must keep in the back of our minds, but I
think we are one order of magnitude away from this now, when fluxes are
just being confirmed.

FAZIO:

You need an awful lot of photons from a known source to get enough statis-
tics to do the polarization. Really that is what it comes down to.

GREISEN:

Could I comment that the angle that everyone sees so far on the pairs isn't
the initial opening angle, but the angle created by the multiple Coulomb
scattering, which has nothing to do with the polarization. So that one is
forced to observe the pair very close to the vertex and that means really
extremely close. The problem is partly one of maintaining even emulsion.
It turns out, depending on what energy you are interested in, you have to
invoke some length of track if you want to measure an angle of some precision,
such as 100 jum and in that length of track, already the scattering exceeds the
original opening angle in some cases.

The problem is difficult and almost certainly drives one to an instrument that
has low detection efficiency. That is because you have to observe the pair
before it crosses much matter and so it is hard to have a high conversion
efficiency. So it is a problem of counting rate. I can't overstress the impor-
tance of being able to do experiments that have enough observation time and
a large enough detection area of the instrument to make possible these finer
details. It seems to me something like the time I remember on top of Mount
Evans, with clouds all around, and it looked as though the whole Rocky
Mountains had only two peaks. There is that whole mountain range down
below, but one needs have a little more vision to see it.

I think the sights should be set not on just barely being able to tell the source
is there, but getting enough information about it to try to approach a kind of

400 FUTURE DIRECTIONS IN GAMMA-RA Y A STR ONOMY

question like polarization, which is one of the hardest. There are some that
are easier, but still seem hard enough. But it is true that currently, in the
MeV range, one of the hardest problems isn't just the counting statistics, but
background problems.

However, I want to emphasize that above 50 or 100 MeV and on up, there is
a lot of promise. But the action is in attaining a better sensitivity. That is
of the order of 10"7 cm-2 s"1 . There are a few sources which may give 10"5,
as possibly the closest source at the best time, but to pursue this subject, we
have to be able to measure down in that range and you have to do very little
calculating to see what that amounts to. If you have a 103 cm2 detector with
an efficiency of 10 percent, so it is effectively 100 cm, you will get one count
in a day with that sort of source. The sort of number of counts that is needed
to extract detailed information is thousands. So clearly, that is not a very
promising instrument. One needs a larger instrument and long observing times.

FAZIO:

Again, worrying about the future of the balloon-borne aspects of the field,
Glenn (Frye) commented that there would be a gap between SAS-2 and the
time when the shuttle program begins. Now, at one time there was a HEAO
experiment planned with a larger area spark chamber. Could I ask what is the
situation on that now, Carl?

FICHTEL:

Yes. Very succinctly, we are essentially in the same position as the low-energy
experiment of Bud Jacobson of JPL. The high-energy 7-ray experiment which
has 10 or 20 times the sensitivity of SAS and good energy resolution is, as is
Jacobson's experiment, one of the candidate experiments for HEAO-C, and a
selection has not yet been made on the new mini -HEAO.

It could indeed provide, of course, a very important next step because as
everyone has said in this conference, what we need to do next is get more
sensitivity and energy resolution in this range and to determine really what
the shuttle experiment should be— whether the shuttle experiment should,
for example, concentrate on angular resolution, details of energy, or indeed,
continue to an even higher-sensitivity survey. So, as with the low-energy 7-ray
experiment, it's in limbo right now.

FAZIO:

To get down into the sensitivity region where you are talking about thousands
of counts on a source, really requires one or two orders of magnitude better
sensitivity over what we are doing right now.

PANEL DISCUSSION 401

GREISEN:

Let me describe the situation a little bit. In our recent flight where we did see
the pulse from the Crab, we looked at it for a few hours and we got something
like 30 events from each of the peaks. That's not very much for telling pulse
structure. But if we were in a satellite environment, our background that we
have to contend with would be down by two orders of magnitude and a
reasonable time of observation, instead of being 3 hours, could be 300 hours;
it could be increased by two orders of magnitude. There are then four orders
of magnitude gained in the combination of time and background, which is two
orders of magnitude anyway in signal-to-noise ratio. One could have not only
good photon statistics, but the noise could be smoothed out in that time, so
that one would really be able to see fine details of pulse structure, even at high
energies.

I received a letter from John Bachall at Princeton, as soon as they got a pre-
print of our paper saying that 'that's fine for a starter, but can you measure
the polarization— because for these high-energy 7-rays from the pulsar, the
polarization would tell the mechanism of production?' It turns out that it is
not unfeasible. Instead of looking for orientation of the polarization plane,
there is a very nice method that is in use at Stanford and Cornell at high ener-
gies with electrons. This involves the use of graphite. At very high energies,
because the recoil momentum of the nucleus when the pair production occurs
is so low when the nucleus is in a crystal, pair production is strongly inhi-
bited for some planes of polarization.

They can use that method to produce a beam that has a polarization as high as
20 or 30 percent by selective absorption, and so one could use that device
too. It's very efficient, actually, for measuring polarization. It would be
possible to measure the background spectrum very, very quickly around 1 GeV
or 100 MeV, but it would also be possible to measure it all the way to 101 1 eV,
which means well beyond the type of cutoff that Floyd Stecker was discus-
sing. This could be done if one could have an instrument of the type that has
several square meters of sensitive area and if one could be free to take data,
not just for a couple of hours after a balloon launch once a year, but steadily.

MEMBER OF THE A UDIENCE:

I think there's one area that we may have slighted in this discussion, and it's
probably been the one area in which the strongest evidence has existed for
7-radiation at energies above 50 MeV, namely the 7-radiation from the galactic
plane. Now, we all know it's there, or at least we think it's there. The question
is, let's say it's there and SAS has a look at it. Where is it? What fine resolu-
tion do we obtain? What do we learn about cosmic-ray densities and matter
densities there? I think that is a very important subject which we are forgetting

402 FUTURE DIRECTIONS IN GAMMA-RA Y ASTR ONOMY

about, so maybe we should consider questions about the angular resolution
we can obtain in the future in studying the galactic plane.

FICHTEL:

Clearly, there should be a significant structure to the flux from the galactic
plane. We already know that there is a broad distribution— I think the word
"diffuse" should be avoided, because there is no implication at this point that
it is only diffuse. There is clearly a hard component, as well as a soft, and it
will be extremely important to find out exactly what this is. And I am fairly
certain that one is going to want to have a finer angular resolution than we
presently have.

In fact, I suspect that indeed SAS, as far as the plane is concerned will answer
many questions. It is also going to only whet one's appetite for a very fine
angular resolution experiment of the future. Sometimes there is the conflict
between sensitivity and fine angular resolution. For something as intense as
the galactic plane as we now know it to be, we could indeed back off on sensi-
tivity in order to pick up the finer angular resolution. I think indeed, this
would be one of the next steps for the future.

FRYE:

If one is really trying to look forward to payloads that might be put on the
Space Shuttle, I think it is quite obvious from the series of discussions here
that there is a strong indication that in another decade there isn't going to be
any one 7-ray detector that will have the energy range, time resolution,
capability for polarization, and so forth. It's going to take a divergence, really,
in the design of the various instruments to cover these measurements. I wish
I could hope that there were going to be many shuttles and that we could
design instruments that would have these capabilities. I would hope that the
scientific community would be able to generate some scientific pressure and
make our case with the people who will eventually decide what goes on board
the shuttle.

SHAPIRO:

I would remind you that according to present ideas, if there are to be any
shuttles, there will be many shuttles. And that means, of course, that this is
not so much the question as the other one that's already been raised: How
well physicists and astrophysicists will be prepared? How well they can be
prepared, as a result of adequate support in the intervening years, to take
proper advantage of at least some of those many shuttle flights that are
projected.

PANEL DISCUSSION 403

FAZIO:

Maury (Shapiro), you said "many." We are going to have to hope they are
going to fund many.

MEMBER OF THE AUDIENCE:

The whole basis of having a shuttle is that you can make it "cheap" per pound
if you had many. Some of us, of course, have reservations as to how many
can be used scientifically, since we don't know in detail about it, or even
broadly perhaps, about any of the other projected uses. We have no way of
knowing whether there may be ample justification on other grounds for
having many (shuttles). I should think that scientists could easily be embar-
rassed by the frequency of opportunities.

GREISEN:

It's something like the possibility of building a glorious new art museum to
house art treasures and then not having any money left to buy any of the art
treasures.

METZGER:

The shuttle has, among its capabilities, the ability to take things into orbit
and then leave them there for extended periods of time. And it seems to me,
that of the various modes that the shuttle offers, this is the one with the most
attraction in our part of the spectrum, because, as you pointed out, time is
of the essence.

GREISEN:

Shouldn't we propose a resolution to that effect?

METZGER:

But one other reason why that is a very useful way to go is that it allows
several experiments to be up for a long period of time together. One has
many promising experiments over this wide energy range, but simply flying
them in sequence is not going to tell as much as the ability to combine them
simultaneously.

GREISEN:

In that connection, I would like to point out that even if the shuttle experi-
ments are semi -independent, they can work together better than some present
experiments can. For instance, if the analysis of data from the SAS-2 should
indicate something very interesting, and if they propose to those of us who

404 FUTURE DIRECTIONS IN GAMMA-RA Y ASTRONOMY

do, to look at it with baUoon experiments, it will take us an awful long while
(at least for our group— I don't know about Glenn Frye's)— to get ourselves
together and be able to fly a balloon. I presume it would take a lot less time
than that to reorient— that is, to point-some satellite which had instruments
we were observing with all the time. Presumably within minutes the satellite
could be ready to look at another object.

FICHTEL:

That's correct, except for the time scale. The SAS-2 magnetic tracking is
rather slow and usually takes a few orbits (to reorient). Of course, if it's
close by, maybe 60° or so, we could do it in one orbit and you are talking
about hours rather than minutes. So, if anybody sees something that you
think is interesting, we would be most happy to hear about it or look at it.
We certainly would let you know, as we said, if we think we see something.

MEMBER OF THE A UDIENCE:

The panel expressed a concern about the gap between SAS-2 and HEAO. I
think this point cannot be stressed enough. In particular, I am sure that some
of you are very familiar with the tragic road that HEAO has gone. HEAO is
by now almost a four -letter word and nothing else. The priority that the
7-ray experiments have on HEAO is extremely low. On the third mission,
they are considered, but that's about it. And they are competing directly
with cosmic-ray experiments. So, as far as HEAO is concerned, I wouldn't
give one penny in that basket. I am very pessimistic.

If I can get to the shuttle for just a second. The shuttle has the capability
of putting one million pounds per year in orbit, but there is no research and
development money available to conduct some of these experiments. So by
the time (1985 or so) you put a million pounds per year in orbit, you don't
know what to put up, because you don't have any money left to do anything.

Now what surprises me is that nobody mentioned that the Europeans are
working very hard on 7-ray experiments, which should be the logical thing
in between SAS and HEAO. In fact, they are very fortunate that HEAO
isn't going, or at least not going as fast as we think it should. Is there
anybody here in the audience who can tell us about COS-B?

SOMMER:

I could tell you just about what the COS-B would be. That is why I wanted
to speak, because nobody mentioned it. Well, you may have seen some things
about COS-B in the literature. As you may know, it is a European collabora-
tion between France, Germany, and Italy—

PANEL DISCUSSION 405

LEWIN:
And Holland.

SOMMER:

And Holland. And it is supposed to have quite a high sensitivity for sources.
For example, it should be able to detect fluxes like 10"7 cm"2 • s"1 . I am not
involved with it, but I know a little bit about it. They are making some cali-
brations in Hamburg and up to now everything seems to be going along quite
well. It's supposed to be launched in 1975, 1 think. So this is quite soon and
it should be a good link between the SAS-2 experiment and the presumed
shuttle experiment.

MEMBER OF THE A UDIENCE:

We took the COS-B payload to CERN and we investigated the background
properties. The data at the moment are just being evaluated. We plan to
make a balloon flight with this COS-B payload late this year. The program
is to launch COS-B in February 1974.

TROMBKA:

Do we have some COS-B sensitivity data?

FAZIO:

I just happened to look up the Madrid (IAU Symposium) discussion that we
had on it. You might correct me if it has been updated. I notice here that
the threshold is around 30 MeV. It is a wire spark chamber with an energy
calorimeter. The sensitive area is about 576 cm2. I think it may have been
about 600. The area-solid angle factor is about 70 cm2 • sr. The energy reso-
lution will be about 50 percent at 100 MeV. The satellite will be spin stabilized
and placed in a highly eccentric orbit. This was an important thing, I thought,
that was different. The main advantages of this eccentric orbit are the reduc-
tion of earth albedo and the reduction of radiation-belt effects, minimum
occultation by the earth, longer observing times, and adequate ground station
coverage .

MEMBER OF THE A UDIENCE:

Does it have very good time resolution for the Crab?

MEMBER OF THE A UDIENCE:

Time resolution criteria were set up at a time when one did not know that the
Crab Pulsar existed, so one could not include the time resolution which enables
periods to be seen of this order.

406 FUTURE DIRECTIONS IN GAMMA-RA Y ASTRONOMY

FAZIO:

Another experiment that we really didn't hear very much about was the TD-1 ,
which is up now. I have some more figures on the sensitivity of the TD-1 . It
has a sensitive area of 1 30 cm2 , an area-solid angle factor of 28 cm2 • sr, and
an angular resolution of about 3°. What is it for SAS-2?

FICHTEL:

SAS is about 5 1 2 cm2 .

FAZIO:

And COS-B is about 600, so you can see how the area is increasing.

GREISEN:

I wonder, Carl (Fichtel), whether you could show some rough diagram of the
angular resolution for SAS-2?

FICHTEL:

It is roughly around 1 .7° at 100 MeV, and it gets better at higher energies
and poorer at lower energies. There are some calibration points that are not
complete.

GREISEN:

Another question: At about what energy does it cease to be possible to get
an energy measurement by observing the scatter?

FICHTEL:

It just doesn't fade away. It is about 25 percent at the heart of the energy
range from 50 to 70 MeV. By the time you get up to 150 MeV, you are down
to possibly the factor of two energy resolution you mentioned.

TROMBKA:

Time does not permit us to continue the discussion. I would like to thank
all the panel members, the participating panelists, and all of the speakers this
week. To me it has been an extremely exhilarating symposium.

We thank you.

LIST OF AUTHORS

Index

Adcock, C, 316

Adler, I., 42, 43, 50

Albats, P., 116, 118, 121, 135, 136

Aldrovandi, R., 336, 337, 340, 342,

343,345,356,369,383
Allcock, M. C, 319
Allen, C. W., 304
Alpher, R. A., 335
Aly,J. J., 346
Anand, K. C, 200
Andouze, J., 299
Andrews, D., 316
Arnett, W. D., 264-266
Arnold, J. R., 41,279
Arons,J. R., 55,237,240,242

Babcock, H. W., 331
Bachall, J., 401
Backenstoss, G., 336
Badhwar, G., 123
Ball, J. S., 336
Bardeen, J. M., 342
Barkas, W. H., 354
Becklin, E. E., 295
Berger, M. J., 47, 78
Bethe, H. A., 335
Beuermann, K. P., 128
Bisnovatyi-Kogan, G. S., 342
Bleeker, J. A. M.,24, 28
Blumenthal, G. R., 189
Bodansky, D., 264
Bok, B. J., 252
Boldt, E. A., 21
Bonometto, S. A., 319
Borner, G., 295

Bosia, G., 153
Bostrom, C. O., 309
Bowyer, C. S., 191
Braddy, D., 284

Bratolyubova-Tsulukidze, L. I., 123
Brecher, K., 55, 188, 190, 240, 241,

250
Bredekamp, J., 151, 196, 368, 383
Brini, D., 37
Broadbent, S. R., 338
Brown, R. H., 157
Brown, R. T., 10, 11,270,272
Browning, R., 109, 112, 113
Bruzek, A., 331
Bryan, R. A., 336
Buffington, A., 362
Bunner, A., 3-9

Burbidge, E. M., 342, 343, 347
Burbidge, G. R., 194, 250, 342, 347

Cameron, A. G. W., 268, 283, 298
Caser, S., 336, 337, 340, 356, 383
CavaUo,G., 213, 214, 216-218,

250,251,256,260
Chan, J., 284
Cheng, C. C, 297
Chibisov, G. V., 341
Chudakov, A. E., 153
Chupp, E. L., 57, 165, 297, 308,

311,332
Cisneros, A., 336
Clark, G. W., 40, 103, 123, 147, 197,

202,213,250,253,256,259,

343, 346
Clark, T. A., 319

407

408

GAMMA-RA Y ASTROPHYSICS

Clayton, D. D., 55,71, 102, 139, 140,
21 1, 263-266, 269-271, 273, 274,
276, 277, 280, 285, 287, 295, 365,
389

Cline, T. L., 57, 175,329,330

Cohen, J. M., 295

Colgate, S. A., 140, 265, 329, 390

Comstock, G. M.,218,219

Cowsik, R., 29, 185, 188, 190, 192,
194, 197, 202, 203, 207-210, 212,
233,240,241,246,256,257

Craddock, W., 266, 269, 274

Craig, H., 356

Cunningham, C, 21

Cusimano, F. J., 155

Dahlbacka, G. H., 107, 108, 112, 113
Dallaporta, N., 341,371
Damle, S. V., 55, 192,232
Daniel, R. R., 128,207
Davidsen, A., 3, 4, 5, 6, 8, 9
Deerenberg, A. J. M., 24, 28
de Freitas Pacheco, J. A., 259
de Gennes, P. G., 338
Dend, W., 346
Derdeyn, S. M., 142
Desai, V. D., 175,329
Deutsch, M.,291
Dilworth, C, 213, 216, 217
Dirac, P. A.M., 291
Ducros, G., 21
Duggal,S. P., 170
Dumas, A., 28
Dunphy, P., 173

Dyer, C. S., 26, 27, 48, 49, 53, 61, 83,
90,91-93, 123

Ekspong, A. G., 343
Eldridge, J. S., 50
Engel, A. R, 26
Evenson, P., 362

Fanselow, J. L., 295

Fazio, G. G., 103, 133,153, 159,
212,213,216,217,237,240,
250, 387-389, 390, 392, 399,
400, 403, 405, 406

Felten,J. E.,3, 185,212,240

Fichtel, C. E., 105-108, 111, 128,
139, 146, 147,213,250,253,
257, 259, 387-390, 394, 400,
402, 404, 406

Field, G. B., 3, 9, 11, 12,31,34,
194

Fishman, G. J., 48, 49, 53, 61, 62,
64,68,83,90,91,94,97, 123,
265, 285, 287, 295

Fisk, L. A., 284, 288

Follin, J. W., 335

Fomin, V. P., 158

Forrest, D. J., 71,92, 165

Fowler, W. A., 264, 266, 284, 286,
342, 357, 365

Frost, K. J., 180,329

Frye, G. M., Jr., 105, 107, 111, 113,
116, 122, 123, 159, 162,231,
387, 392, 394, 398-400, 402, 404

Fuligni, F., 37

Gal'per, A. M., 103

Gamow, G., 151,335

Garmire, G. P., 4, 8, 103, 1 14, 123,

197,250,253
Geiss, J., 356

Giacconi, R.,31, 191, 193
Ginzburg, V. L., 139, 212, 213, 249,

250,252-257
Gold, T., 194, 343
Golden, 362
Goldhaber, M., 335
Goldman, D. T., 64
Goldsmith, D.W., 213, 217-221
Goldstein, H., 64
Goldstein, M. L., 284, 288
Golenetskii, S. V., 55, 61, 128, 135,

150, 192,234,235
Gordon, I. M., 312

INDEX- LIST OF A UTHORS

409

Gorenstein, P., 8, 21,28, 29
Gould, R.J. , 10, 11, 159, 189, 194,

212-214, 216-218, 240, 250, 25L

256,260,319
Grader, R., 5
Green, J., 291
Greisen, K., 316, 387-388, 391, 393-

395,399,401,403,406
Grindlay, J. E., 155-157, 159-162
Gursky, H., 15, 191,364,365
Guthrie, P., 295

Hainebach, K., 265

Hammersley, J. M., 338

Harnden, F. R., 57, 287, 294

Harrison, E. R., 151,337

Hartman, R. C, 139,295,390

Hayakawa, S., 6, 8, 187, 189, 213

Haymes, R. C, 57, 287, 294, 295

Hearn, D., 116

Heitler,W., 212

Helmken, H. F., 1 09, 1 1 1 , 1 20, 1 60

Henry, R. C, 9, 11, 12, 21

Herman, R. C, 335

Higbie,P. R., 165

Hildebrand, R. H., 295

Hill, F. W., 17

Hillas, A. M.,315

Hoffman, J. A., 109, 111, 120, 159,

161
Hones, E. W., 57
Horine, E., 162
Horstman, H., 23, 28, 37, 40
Horstman-Moretti, E., 22, 23, 37
Howard, W. M., 265
Hoyle, F., 194, 240, 342, 343, 357
Hudson, H. S., 123
Hughes, D. J., 303
Hutcheon, I. D., 202
Hutchinson, G. W., 105

Illarionov, F., 381
Imhof,W. L.,71,73, 77,80
Ipavich, F. M., 322

Jacobs, W. W., 299

Jacobson, A. S., 71, 99, 390, 395, 400

Jansky, K. G., 103

Jenkins, E. B., 6

Johnson, R.G.,7 1,77

Johnson, W. N., Ill, 57, 285, 287,

294, 295
Jones, F. C, 212
Joseph, G., 128

Kane, S. R., 329

Kaplon, F., 123

Kasturirangan, K., 123, 127

Kato, T., 3, 8

Kellerman, K., 346

Khazan, Ya. M., 253

King, J., 310

Kinzer, R. L., 109, 116-118, 121-

123, 125,128
Klebesadel, R. W., 175, 179, 181,

329, 331
Kniffen, D. A., 105, 128, 139, 209,

210,232,234,246,257
Kobetich, E. J., 29, 188, 190, 194,

240
Koltun, D. S., 354
Kraushaar, W. L., 3, 55, 103-105,

108,111, 114, 123, 129, 135,

139, 140, 147, 197,204,216,

217,219,232,246,250,253,

259,261,364
Kristian, J., 347
Kurfess, J. D., 117, 119, 135
Kuzmin, V. A., 316

Lang, R. R., 191
Langhoff, W., 260
Lavakare, P. J., 128
Layzer, D., 342, 346, 347
Lee, J., 291
Leighton, H. I., 166
Lenchek, A. M., 322
Lequeux, J., 344
Leray,J. P., 121

410

GAMMA-RA Y ASTROPHYSICS

Leroy, B.,351

Leventhal, M., 291,295

Levy, D. J., 213,217-221

Lewin, W., 97, 405

Lincoln, J. V., 166

Ling, J. C, 24

Lingenfelter, R. E., 167, 170, 286,

297,299,302,306,312,314
Linsley, J., 316
Long, CD., 153
Low, F. J., 347
Lucchin, F., 319, 371

Manchanda, R. K., 20, 24

Maringelli, M., 153

Massey, H. S. W., 291-293, 295

Matteson, J. L, 28,41,279

Mayer, Hasselwander, H. A., 55, 123,

128, 129, 135, 150, 232, 234, 235,

326
McBreen, B., 118, 120, 121
McCammon, D., 10, 30
McCray, R., 237, 240
McDonald, F. B., 173,314
McVittie, G. C, 275
Metzger, A. E., 30, 41, 43, 57, 93, 94,

97, 209, 210, 279, 387, 395-397,

403
Meyer, P., 295
Miller, R. H., 343
Misra, D., 291

Mohr, C. B. O., 291-293,295
Montmerle, T., 346
Morfill, G. E., 26, 27, 49, 53, 61, 123
Morgan, D. L., Jr., 55, 151, 196, 240,

242, 368
Morrison, P., 55, 139, 188, 190, 208,

211-213,239,241

Nakano,G. H.,71,72,77, 390
Navorra, G., 153
Nemethy, P., 365
Neugebauer, G., 295
Nicolle, J. P., 351

Noonan, T. W., 11
Novikov, I. D., 342, 344

Ogelmann, H., 105, 128

Olsen, R. A., 329

Omnes, R., 151, 239, 243, 335-337,

342,351,356,369,371,381,

383, 392
O'Mongain, E. P., 116, 155
Oort, J. H., 254, 274, 344
Ore, A., 292
Ostricker, J. P., 332
Overbeck, J. W., 71,73
Ozernoy, L. M., 341, 342, 346,

347, 370

Pal, Ya. A., 55, 103, 123, 133, 135,

241
Palmieri, T. M.,5, 21,28
Parker, E. N., 343
Parlier, B., 121, 122, 135, 136
Peebles, R., 341
Penzias, A. A., 315
Perek, L, 201
Peterson, L. E., 24-27, 41, 61, 94,

95,97, 123, 127, 135,246,279
Petrosian, V., 194
Peyraud, N., 346
Phillips, R. J. N., 336
Pieper, G. F., 93, 396
Pinkau,K., 133,388
Pollack, J. B., 213, 216, 217
Pomerantz, P., 170
Porter, N., 161,162
Powell, J. L., 292
Prakasarao, A. S., 21
Price, P. B., 284
Prilutskii, O. F., 241-243
Puget, J. L., 236, 239, 240, 243,

337,341,342,356,361,367,

369,374,381-383,392

Quenby, J. J., 26

INDEX- LIST OF A UTHORS

411

Ramaty, R., 71, 167, 170, 173, 194,
255, 284, 286, 288, 291 , 295, 297,
299, 302, 306, 314, 398

Ramsden, D., 109

Rangan, K. K., 24

Rao,U. R., 123, 127

Reagan, J. B. ,7 1,77

Reedy, R.C., 41, 47, 279

Rees, M. J., 237, 241

Reeves, H., 284, 286, 356

Reppin, C, 165, 173

Riegler, G., 4, 8

Rieke,G. H., 153, 159

Robinson, I., 342

Roll, P. G., 17, 19

Rosburgh, I. W., 342

Rothenflug, R., 24

Rougoor, G. W., 344

Rozental, I. L., 241

Rudstam, G., 62, 83, 85

Sandage, A., 274, 278

Sanders, W., 4, 6, 7

Saslaw, W. C, 343

Savage, B. D., 6

Schatzman, E., 239, 243, 343, 346,

351,355,369,387,389,390,

392, 397
Schild, A., 342
Schmidt, M., 268, 342, 346
Schreder,G. P., 319
Schucking, E. E., 342
Schwartz, D. A., 15, 17, 24-28, 38-40,

123, 195,232
Schwartz, R. B., 303
Scotti, A., 336
Seeman,N., 109, 125, 128
Seltzer, S. M., 47, 78
Setti, G., 241,250
Seward, F. D., 5, 193
Shapiro, M. M., 62, 90, 91, 283, 285,

402, 403
Share, G. H., 55, 94, 103, 1 17, 125,

128, 140, 232, 246, 257

Shklovskii, I. S., 250

Silberberg, R., 62, 63, 283, 285

Silk, J., 3, 55, 129,185, 191, 192,
195, 200, 240, 274, 276, 277,
284,286,371

Slattery, P., 219

Smith, L. H., 231

Sobel, H. W., 64

Sommer, M., 390, 405

Spitzer, L. J., 343, 344

Sreekantan, B. V., 219

Stecher, T. P., 213,246, 256

Stecker, F. W., 55, 57, 129, 135,
139, 150, 151, 180, 187, 188,
196,201-203,207-218,221,
222,224,229-231,233,236-
240, 242, 243, 246, 250-252,
256,257,261,291,302,316,
317,319,329,341,342,367-
369,381,383,387,391-395,
398, 401

Steigman, G., 210, 243, 255, 341 ,
343,351,361,362,364,367,
368, 394, 399

Stepanian, A. A., 158

Stephens, S. W., 57

Strittmatter, P. A., 362, 364, 367

Strong, A. W., 129, 259, 315, 325

Strong, I. B., 175, 181,329

Sullivan, J. D., 203

Sunyaev, R. A., 241, 243, 343,
381,383

Suri, A. N., 165

Syrovatskii, S.' I., 212, 213, 250,
253,254

Tademaru, E., 295
Teegarden, B. J., 311
Teller, E., 343
Terzian,Y., 191
Thielheim, K. O., 260
Toor, A., 22
Tornabene, H. S., 155

412 GAMMA-RA Y ASTR OPHYSICS

Trombka, J. I., 30, 41-43, 47, 50, 53, Zirin, H., 312

67, 69, 90, 93-95, 123, 127, 129, Zobel, W., 306

133-135, 150, 181, 232, 234, 235,

246, 279, 281, 387, 389, 396, 406
Tsao, C. H., 62, 63
Tucker, W., 8, 29

Warn, S. M., 343
Ulmer,M. P.,256
Urey, H. C, 356

Valdez, J. V., 105

Valentine, D., 123

Van der Kruit, P. C, 254

Vedrenne, G., 55, 135, 150, 234, 235

Vette, J. I., 43, 54, 55, 92, 94, 95,

127,208-210,257
Vladimirsky, B. M., 158

Waddington, C. J., 105
Wagoner, R. V., 342, 357
Walden, W. E., 343
Wampler, D., 347
Wang, C. P., 105, 116, 122, 123
Ward, R. A., 280
Wdowczyk,J., 259, 315-324
Weekes, T. C, 158,159, 162
Weinberg, S., 275
Wheaton, W. A., 180
White, S., 210, 313, 314, 397
Wilkinson, D. T., 17, 19
Williamson, F. W., 4, 7
Wilson, R.W., 315
Wolfendale, A. W., 259, 315
Woltjer, L., 250, 260, 343
Womack, E. A., 71,73
Wong, D. Y., 336
Woosley, S. E., 264, 265
Wright, P. J., 109,212,240

Zaimidoroga, O. A., 354
Zatsepin, G. T., 153,316
Zel'dovich, Ya. B., 335, 342-344,
381,383

<4J.S. GOVERNMENT PRINTING OFFICE: 1973-730-732/66-1-3

QB460.G35

3 5002 00402 9158

Gamma-ray astrophysics; a symposium held

DATE DUE

0s COH02 7/srt —

460 *

G35 ~

AUTHOR

Stecker.

1 1 1 L t

Gamma-ray astrophysics, ~~

DATE DUE

BORROWERS NAME

*/fvf/jf-

' jS.S*4£u^

r i • .—

~?tzr \ ■

460
G35

•

-IONAL AERONAUTICS AND SPACE ADMINISTRATION
WASHINGTON. DC. 20546

OFFICIAL BUSINESS
ALTV FOR PRIVATE USE »300

POSTAG

NATIONAL

SPACE

AND FEES PAID
ERONAUTICS AN

MINISTRATION

SPECIAL FOURTH-CLASS RATE
BOOK

POSTMASTER :

If I'ndellverable (Section 158
Postal Manual ) Do Not Return

"The aeronautical and space activities of the United States shall he
conducted so as to contribute . . . to the expansion of human knowl-
edge of phenomena in the atmosphere and space. The Administration
shall provide for the widest practicable and appropriate dissemination
of information concerning its activities and the results thereof."

— National Aeronautics and Space Act of 1958

NASA SCIENTIFIC AND TECHNICAL PUBLICATIONS

TECHNICAL REPORTS: Scientific and
technical information considered important,
complete, and a lasting contribution to existing
knowledge.

TECHNICAL NOTES: Information less broad
in scope but nevertheless of importance as a
contribution to existing knowledge.

TECHNICAL MEMORANDUMS:
Information receiving limited distribution
because of preliminary data, security classifica-
tion, or other reasons. Also includes conference
proceedings with either limited or unlimited
distribution.

CONTRACTOR REPORTS: Scientific and
technical information generated under a NASA
contract or granr and considered an important
contribution to existing knowledge.

TECHNICAL TRANSLATIONS: Information
published in a foreign language considered
to merit NASA distriburion in English.

SPECIAL PUBLICATIONS: Information
derived from or of value to NASA activities.
Publications include final reports of major
projects, monographs, data compilations,
handbooks, sourcebooks, and special
bibliographies.

TECHNOLOGY UTILIZATION
PUBLICATIONS: Information on technology
used by NASA that may be of particular
interest in commercial and other non-aerospace
applications. Publications include Tech Briefs,
Technology Utilization Reports and
Technology Surveys.

Details on the availability of these publications may be obtained from:

SCIENTIFIC AND TECHNICAL INFORMATION OFFICE

NATIONAL AERONAUTICS AND SPACE ADMINISTRATION

Washington, D.C. 20546

Internet Archive Audio

Featured

Top

Images

Featured

Top

Software

Featured

Top

Books

Featured

Top

Video

Featured

Top

Mobile Apps

Browser Extensions

Archive-It Subscription

Save Page Now

Full text of "Gamma-ray astrophysics; a symposium held at NASA Goddard Space Flight Center, April 30 to May 2, 1973"

See other formats