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Full text of "The First year of HST observations : proceedings of a workshop held at the Space Telescope Science Institute, Baltimore, Maryland, 14-16 May 1991"

THE FIRST YEAR 
OF HST OBSERVATIONS 





Wavelength (A) 



Proceedings of a Workshop held at the 

Space Telescope Science histitute 

Baltimore, Maryland 

14-16 May 1991 

Edited by A.L. Kinney and J.C. Blades 



ii 


SPACE 
lELESCOPE 

SQENCE 


r V ^ 


INSlllUl'E 



rw\SA 

National Aeronautics and 
Space Administration 



sJi^* 




The upper plot shows the spectrum of 3C273 in the GHRS G140L; raw data (below) and 
deconvolved data (above). The lower plot shows the spectrum of 3C 273 in the GHRS G160M; 
raw data (below) and deconvolved data (above). From the paper by R. J. Weymann. 
(See also Morris S. L., Weymann, R. J. Savage, B. D., and Gilliland, R. L. 1991, ApJ. ?>11, L21. 




SPACE 
TELESCOPE 
SCIENCE 
INSTITUTE 



rVI/NSA 

National Aeronautics and 
Space Administration 



THE FIRST YEAR 
OF HST OBSERVATIONS 



Proceedings of a Workshop held at the 

Space Telescope Science Institute 

Baltimore, Maryland 

14-16 May 1991 



Editors: 

A. L. Kinney and J. C. Blades 

Space Telescope Science Institute 

Science and Planning Division 



Published and distributed by the Space Telescope Science Institute 
3700 San Martin Drive, Baltimore, MD 21218 



The Space Telescope Science Institute is operated by the Association of Universities for Research in 

Astronomy, Inc., under NASA contract NAS5-26555 



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CONTENTS 

Preface 

PAPERS 

HST and Dense Stellar Systems 1 

I.R. King 

The Central Dynamics of 47 Tucanae 6 

I.R. King 

A Tale of Three Jets 10 

F. Macchetto and the FOC Investigation Definition Team 

Early Observations of Gravitational Lenses with the 

Planetary Camera of Hubble Space Telescope 25 

J. Kristian for the WF/PC Investigation Definition Team 

Ultraviolet Spectroscopic Studies of the Interstellar 

Medium with the Hubble Space Telescope 33 

B.D. Savage 

FOS Observations of the Absorption Spectrum of 3C 273 46 

J.N. Bahcall, B.T. Jannuzi, D.P. Schneider, G.F. Hartig, 
R. Bohlin, V. Junkkarinen 

Results and Some Implications of the GHRS Observations of the 

Lyman a Forest in 3C273 58 

R.J. Weymann 

Hot Stars and the HST 68 

R.P. Kudritzki 

GHRS Far-Ultraviolet Spectra of Coronal and Noncoronal Stars: 

Capella and y Draconis 70 

J.L. Linsky, A. Brown, K.G. Carpenter 

High Resolution UV Spectroscopy of the Chemically Peculiar 

B-Star, Chi Lupi 83 

D.S. Leckrone, S.G. Johansson, G.N. Wahlgren 

Hubble Space Telescope Optical Performance 96 

C.J. Burrows 



Introduction to the Goddard High Resolution Spectrograph (GHRS) 106 

J.C. Brandt 

Status of the Goddard High Resolution Spectrograph in May 1991 110 

D. Ebbets, J. Brandt, S. Heap 

Early Operations with the High Speed Photometer 123 

J.W. Percival, R.C. Bless, M.J. Nelson 

Early Commissioning Astrometry Performance of the Fine Guidance Sensors 131 

G.F. Benedict, W.H. Jefferys, Q. Wang, A. Whipple, E. Nelan, 
D. Story, R.L. Duncombe, P. Hemenway, P.J. Shelus, 
B. McArthur, J. McCartney, O.G. Franz, L. Wasserman, 
T. Kreidl, W.F. van Altena, T. Girard, L.W. Fredrick 

A Review of Planetary Opportunities and Observations with 

the Hubble Space Telescope 147 

R. Beebe 

Observations of Mars Using Hubble Space Telescope Observatory 161 

P.B. James, R.T. Clancy, S.W. Lee, R. Kahn, R. Zurek, 
L. Martin, R. Singer 

Decon volution and Photometry on HST-FOC Images 178 

C. Barbieri, G. De Marchi, R. Ragazzoni 



POSTERS 

FOC Images of the Gravitational Lens System G2237+0305 188 

P. Crane, R. Albrecht, C. Barbieri, J.C. Blades, A. Boksenberg, 

J. M. Deharveng, M.J. Disney, P. Jakobsen, T.M. Kamperman, 
I.R. King, F. Macchetto, CD. Mackay, F. Paresce, G. Weigelt, 
D. Baxter, P. Greenfield, R. Jedrzejewski, A. Nota, W.B. Sparks 

Reduction of PG1115+080 Images 192 

E.J. Groth, J. A. Kristian, S.P. Ewald, J.J. Hester, J. A. Holtzman, 
T.R. Lauer, R.M. Light, E.J. Shaya, and the rest of the 
WF/PC Team: W.A. Baum, B. Campbell, A. Code, D.G. Currie, 
G.E. Danielson, S.M. Faber, J. Hoessel, D. Hunter, T. Kelsall, 
R. Lynds, G. Mackie, D.G. Monet, E.J. O'Neil, Jr., D.P. Schneider, 
P. K. Seidelmann, B. Smith, J. A. Westphal 

Optical and UV Polarization Observations of the M 87 Jet 196 

P.E. Hodge, F. Macchetto, W.B. Sparks 



The Non-Proprietary Snapshot Survey: A Search for Gravitationally-Lensed 
Quasars Using the HST Planetary Camera 200 

D. Maoz, J.N. Bahcall, R. Doxsey, D.P. Schneider, N.A. Bahcall, 
O. Lahav, B. Yanny 

Faint Object Spectrograph Observations of CSO 251 204 

R.D. Cohen, E.A. Beaver, E.M. Burbidge, V.T. Junkkarinen, 
R.W. Lyons, E.I. Rosenblatt 

FOC Observations of R136a in the 30 Doradus Nebula 208 

G. Weigelt, R. Albrecht, C. Barbieri, J.C. Blades, A. Boksenberg, 
P. Crane, J.M. Deharveng, M.J. Disney, P. Jakobsen, 
T.M. Kamperman, I.R. King, F.Macchetto. CD. Mackay, 
F. Paresce, D. Baxter, P. Greenfield, R. Jedrzejewski, 
A. Nota, W.B. Sparks 

GHRS Chromospheric Emission Line Spectra of the Red Giant a Tau 212 

K.G. Carpenter, R.D. Robinson, D.C. Ebbets, A. Brown, 
J.L. Linsky 

lUE Far-Ultraviolet Spectra of Capella and y Draconis for 

Comparison to HST/GHRS GTO Observations 216 

T.R. Ay res 

Faint Object Camera In-Flight Performance 

Geometric Distortion, Stability and Plate Scale 220 

D. Baxter 

In-Flight Performance of the FOC: Early Assessment 

of the Absolute Sensitivity 225 

W.B. Sparks and the FOC IDT 

In-Flight Performance of the FOC: Flat Field Response 229 

P. Greenfield and the FOC IDT 

Background Noise Rejection in the Faint Object Spectrograph 234 

E.I. Rosenblatt, E.A. Beaver, J.B. Linsky, R.W. Lyons 

Detection of Binaries with the FGS: The Transfer Function 

Mode Data Analysis 238 

B. Bucciarelli, M.G. Lattanzi, L.G. Taff, O.G. Franz, 
L.H. Wasserman, E. Nelan 

Restoration of Images Degraded by Telescope Aberrations 245 

T. Reinheimer, D. Schertl, G. Weigelt 



Coping with the Hubble Space Telescope's PSF: 

Crowded Field Stellar Photometry 249 

E.M. Malumuth, J.D. Neill, D.J. Lindler, S.R. Heap 

Some Algorithms and Procedures Useful to Analyse HST-FOC Images 253 

C. Barbieri, G. De Marchi. R. Ragazzoni 

Deconvolution of an FOC Image Using a TIM-Generated PSF 260 

P.E. Hodge 

Rapid Deconvolution of Hubble Space Telescope Images 

on the NRL Connection Machine 264 

P. Hertz, M.L. Cobb 

On Orbit Measurement of HST Baffle Rejection Capability 267 

P. Y. Bely, D.Daou. O. Lupie 



APPENDIX 

Scheduling of Science Observations and 

Subsequent Data Processing 

Transformation: The Link Between the Proposal and the 

Hubble Space Telescope Database 270 

M.L. McCollough, H.H. Lanning, K.E. Reinhard 

Proposal Preparation by SPSS for Scheduling 

on the Hubble Space Telescope 276 

K.E. Reinhard, H.H. Lanning, W.M. Workman, III 

The Scheduling of Science Activities for the Hubble Space Telescope 281 

D.K. Taylor, K.E. Reinhard, H.H. Lanning, 
D.R. Chance, E.V. Bell. II 

The Scheduling Efficiency for the Hubble Space Telescope 

During the First Year of Operation 288 

E.V. Bell, II, K.E. Reinhard, H.H. Lanning 

Routine Science Data Processing of HST Observations 294 

D.A. Swade, S.B. Parsons, P. Van West, S. Baggett, 
M. Kochte, D. Macomb, A. Schultz, I. Wilson 



PREFACE 

This volume presents the proceedings of the meeting on the Year of First Light, held 
at the Space Telescope Science Institute in Baltimore on 1991 May 14-16. Cohosted by the 
HST Science Working Group and the Space Telescope Science Institute, the meeting took 
place at the close of the engineering commissioning period and the beginning of the 
Guaranteed Time and General Observer science programs. The goals of the meeting were 
to gather the collective experience of the scientists who were instrumental in preparing for 
the HST mission and analyzing the early science and calibration data, and to inform both 
the HST community and prospective observers of the scientific potential of the Observatory, 
even with its aberrated optics. 

At the time of the meeting, routine science observations were being scheduled on 
HST for four of the six scientific instruments. Only the High Speed Photometer and the 
Fine Guidance System (as used for astrometry) were still in their early commissioning 
activities. As a result, we were able to invite and solicit presentations covering many 
astronomical fields and representing many of the HST capabilities. The presentations 
included topics such as planets, hot stars, cool stars, chemically peculiar stars, stellar 
systems, galactic phenomenon, gravitationally lensed quasars, and the absorption systems 
observed in nearby quasars. 

Because HST is a unique and complex observatory, we had also invited presentations 
and poster papers on the spacecraft performance, including the improved understanding of 
the Optical Telescope Assembly and the resultant imaging quality. Posters on the planning 
and scheduling system, which are included as appendices in this volume, show how science 
observations are executed with the telescope. 

The meetings Scientific Organizing Committee would like to thank all participants 
in the three day meeting and, in particular, the contributors to this volume. The Editors, 
Anne Kinney and Chris Blades have collected an outstanding compilation of scientific and 
technical manuscripts. As these offer excellent examples of the scientific capabilities of the 
Hubble Space Telescope, we are very grateful for the Editors' efforts and are pleased to 
distribute these early scientific results to the astronomical community. Readers are also 
directed to the dedicated editions of the Astrophysical Journal Letters (Ap. J. Lett. 369, 
No. 2 and 377, No. 1) and several papers describing the performance of the scientific 
instruments (Greenfield et al. 1991 SPIE 1494, p. 16 for FOC, Harms, R., et al. 1984, 
Instrumentation in Astronomy V., p. 410 for FOS, and Lauer, T.R. 1989, PASP, 101, p. 
445 for WF/PC). 

Scientific Organizing Committee: 

Dr. Riccardo Giacconi, Space Telescope Science Institute (ST Scl) 

Dr. Albert Boggess, Goddard Space Flight Center (GSFC) 

Dr. H. S. (Peter) Stockman, ST Scl 

Dr. David Leckrone, GSFC 

Dr. Colin Norman, ST Scl 

Dr. Michael Fall, ST Scl 



HST AND DENSE STELLAR SYSTEMS 



Ivan R. King 
Astronomy Department 
University of California 
Berkeley, CA 94720 

U.S.A. 



1. INTRODUCTION 

In this discussion of the capabilities of HST in observing dense stellar systems, I will 
take up globular clusters and galaxies, but not active galactic nuclei, since the AGN's 
will be discussed in another part of the workshop. Actually my talk will have three 
parts: globular clusters, nearby galaxies, and distant galaxies. 

We of course have four cameras available. The WFC has a pixel of 100 milliarcseonds 
(mas) and a field size of 160 arcsec. The PC and the FOC f/48 each have a pixel of 45 
mas; for the PC the field size is 67 arcsec, and for the FOC f/48 it is only 22 arcsec. 
(Its extended-field capability is not usable here, because the reduced counting capacity 
in that mode cannot accommodate the range of magnitudes that we encounter.) The 
FOC f/96, with its pixel of 23 mas, was designed to critically sample the diffraction- 
limited HST image down to about 520 nm. At shorter wavelengths it would have been 
somewhat undersampled. But its field is only 11 arcsec. 

(I am not going to say anything about the f/288 mode of the FOC, partly because its 
field is so tiny and partly because the spherical aberration makes its use very limited.) 

In the actual aberrated images the FWHM of the core is about 65 mas, a value that 
does not depend much on wavelength, so the f/96 does have sampling that is adequate 
to use the resolution that the sharp image core gives us. We have in fact lost about a 
factor of 1.5 in the visible, relative to a diffraction-limited image, but we nevertheless 
have nearly the resolving power that HST was intended to have. But what we have 
really lost is about 5/6 of the light, which is out in the halo, where it is of no use to us. 
Even worse, in dense stellar systems this light spreads over the neighboring objects that 
we would like to measure, thereby increasing the background in an unpleasant way. We 
lose not only the 2 magnitudes represented by only 1/6 of the light being in the core, 
but we probably lose another magnitude because of this spread-out light. 

I should mention some other differences between the cameras, too. The two most 
important are the dynamic range and the PSF. The WF/PC CCD's have a quite large 
dynamic range, whereas the IPCS detectors of the FOC go non-linear, and then satu- 
rate, at count rates of only a couple of tenths up to 1 count/pixel sec. On the other 

1 



hand, the PSF in an FOC image is independent of position, whereas in both the WFC 
and the PC a central obscuration in the re-imaging optics causes the PSF to vary quite 
significantly with position in the image. (A third difference, the readout noise of the 
WF/PC, is almost always of no importance when imaging dense stellar systems.) 

2. GLOBULAR CLUSTERS 

Globular clusters are one of the places where HST has a real advantage, because 
of its large scale and its high resolving power. But at the same time we have a small 
field (especially if we want the highest resolving power, which is achieved with the FOC 
f/96). Globular clusters are big things, and we can sample only small parts of them. 

The superior resolving power of HST allows us to study faint stars at the centers 
of globular clusters; at ground-based resolutions these are literally covered up by the 
merged outer envelopes of the images of brighter stars. Prior to launch I had calculated 
that the FOC f/96 should be able to see stars to the HST limit even at the center of 
47 Tucanae, which is one of the densest globulars known. This may indeed be possible 
after COSTAR is in place. But at the present time we are frustrated by the overlapping 
halos of the bright stars, which saturate almost all the area of a visible-light image. The 
PC is of course not subject to this limitation, but those overlapping halos still keep us 
from the faint images. 

But even in a cluster as dense as 47 Tuc we are able to work in the ultraviolet and 
avoid saturation. Thus Mike Shara will be reporting on the discovery of blue stragglers 
in the center of 47 Tuc (later published as Paresce et al. 1991). 

Another application of resolving power is in the clusters whose cores have collapsed 
down to a radius too small to resolve reliably from the ground. An example is M15, for 
which Tod Lauer will report PC observations later in this symposium. (See also Lauer 
et al. 1991a.) 

One problem that has always interested me is the degree of equipartition of energy 
that has been reached in globular clusters. As the stars encounter each other and 
their velocity distributions relax, the tendency is for each stellar type to take on a 
velocity dispersion that is inversely proportional to the square root of its mass. The 
relative velocity dispersions should manifest themselves as different core radii for the 
different species. This is an impossible problem from the ground, because the cores are 
impenetrable, but with HST I look forward to taking it on. 

Another problem — unfortunately postponed to the COSTAR era — is the study of 
horizontal branches in the globular clusters of M31. They are all at the same distance 
modulus and can therefore be compared directly with each other. Observation of a 
sufficient number of them should clear up the question of the dependence of absolute 
magnitude on metallicity, and may go a long way toward solving the "second-parameter 
problem." 

I should also mention here another headache that the spherical aberration has 
created for us when we try to do stellar photometry in globular clusters. A photometry 
program such as DAOPHOT groups stars in sets of mutually overlapping images, and 
does an iterative fitting of the individual images within each group. But with our present 
PSF, all the stars in a cluster field make a single horrendous group. In my view, the 
computing time then becomes prohibitive. My solution has been to do a deconvolution 
of the image by Fourier techniques, which preserves photometric integrity, and then 
measure the deconvolved image. This approach has its own problems, though. Prior 
to deconvolution, bright stars with saturated centers have to be laboriously subtracted 



out, one by one, as do stars whose halos run beyond the edges of the image. And 
deconvolved images have statistical problems that are far from trivial. 



2. THE MILKY WAY AND NEARBY GALAXIES 

2.1 The Milky Way 

Both the FOC team and the WF/PC team are doing studies of Baade's Window, 
in the Galactic bulge. In the latter case, higher-latitude fields are also being studied. 
Here again we have a problem that benefits greatly from the resolving power of HST. 
Terndrup's (1989) pioneering ground-based study of Baade's Window was stopped by 
crowding around 19th or 20th magnitude; HST should go faint enough to study the 
main-sequence turnoff, below 22nd magnitude. 

2.2 The Magellanic Clouds 

The study of the R136 cluster in 30 Doradus, in the LMC was, I believe, the first 
scientific achievement of HST. (For a more recent treatment, see Weigelt et al. 1991.) 
Also, we shall hear during this Workshop from Francesco Paresce about the resolution 
of the expanding shell around Supernova 1987A. (See also Jakobsen et al. 1991.) 

Here is another pair of systems in which high-resolution problems abound: star 
clusters, dense regions such as the bar of the LMC, etc. 

2.3 The M31 Group 

The bulge of M31 is a fascinating place, and a number of us will be delving into 
it. I hope to be able to see the low-metal-abundance Baade giants quite close in to the 
center, and thus determine the core radius of the M31 halo. And I know that Mike 
Rich will be studying these same stars in the near infrared. 

I have already mentioned the globular clusters, but I should also note that when 
they are observed with the FOC, parallel exposures with the WFC will study the M31 
halo population better than ever before. 

Perhaps the most exciting problem associated with these galaxies is the centers of 
M31 and M32, both of which are very dense. First, direct imaging of the center of 
M32 will be very interesting. In an observation that was never published, Stratoscope 
saw a steady rise of brightness to a peak at least as sharp as its resolving power of 0.2 
arcsec (Schwarzschild, private communication). Even more provocative, however, are 
the central rotations, whose angular velocity surpasses ground-based resolving power 
(Kormendy 1987, 1988). After COSTAR is installed, the long-slit spectrograph of the 
FOC will be able to tell just how high these angular velocities are. Simple dynamical 
reasoning shows that the central mass densities are directly proportional to the angular 
velocities. 

And there is the problem of the origin of the ultraviolet light. At the moment of 
this writing, colleagues and I are struggling to interpret what we see in a far-UV image 
of the center of M31. 

Finally (but I doubt it, really), there are problems of the disk population. M31 is 
a very crowded object, and all sorts of studies of the disk population will benefit from 



the resolving power of HST. 

2.4 Other Nearby Galaxies 

Most of the observations I have mentioned are in the future. But we will hear today 
from Tod Lauer about the discovery of a sharp nucleus in NGC 7457. (See also Lauer et 
al. 1991b.) And other not-too-distant galaxies are amenable to all sorts of new studies. 

One problem that has interested me in particular is the cores of elliptical galaxies. 
At the distances of the nearest giant ellipticals, cores such as those of M31 and M32 
would be completely unresolved; and, in fact, Schweizer (1976) has suggested that many 
ellipticals that seem to have barely resolved cores are indeed in this situation. 1 am 
personally relieved, however, to find that in at least one Virgo elliptical HST imaging 
has shown the core radius to be just the 2 arcsec that I had estimated from ground-based 
images (King 1978). 



3. DISTANT GALAXIES 

For more distant galaxies our main concern is morphology. Here the resolving 
power of HST is crucial. The diameters of distant galaxies depend somewhat on the 
cosmology, but the important point is that they never get very small, so that HST 
resolving power can be used to study morphology at any distance. The only drawback 
at large distances is the dimming factor of (1 + z)\ which holds in all cosmologies. This 
can make the needed exposure times quite long, especially because of the long focal 
ratios of our cameras. 

We have carried out one test on middle-distance galaxies, among the Scientific 
Assessment Observations. Single-orbit exposures were made with the FOC f/48 and 
with the WFC. The results (King et al. 1991) included a successful study, with the 
FOC, of a galaxy at magnitude 20.5 and good morphological assessments, with the 
WFC, of galaxies ranging from 17.5 to 19.8. Deconvolutions were made in all cases. 
One important conclusion, however, was that for really good results many orbits of 
exposure would be needed. 



4. CONCLUSION 

Contrary to many misguided early prophecies, imaging with HST is far from dead. 
For globular clusters and galaxies, a number of quite interesting results have already 
been obtained, and many more are in the offing. 



REFERENCES 

Jakobsen, P., Albrecht, R., Barbieri, C, Blades, J. C, Crane, P., Deharveng, J. M., 
Disney, M. J., Kamperman, T. M., King, I. R., Macchetto, F., Mackay, C. D., 
Paresce, F., Weigelt, G., Baxter, D., Greenfield, P., Jedrzejewski, R., Nota, A., 
Sparks, W., Kirshner, R. P., and Panagia, N. 1991, Ap. J. (Letters), 369, L63. 

King, I. 1978, Ap. J., 222, 1. 

King, I., Stanford, S. A., Seitzer, P., Bershady, M. A., Keel, W. C, Koo, D. C, Weir, 
N., Djorgovski, S., and Windhorst, R. A. 1991, A. J., 102, 1553. 



Kormendy, J. 1987, in Structure and Dynamics of Elliptical Galaxies (I.A.U. Symposium 

127), ed. T. de Zeeuw (Dordrecht: Reidel), p. 17. 
Kormendy, J. 1988, Ap. J., 325, 128. 
Lauer, T. R., Holtzman, F. A., Faber, S. M., Baum, W. A., Currie, D. G., Ewald, S. 

P., Groth, E. J., Hester, J. J., Kelsall, T., Light, R. M., Lynds, C. R., O'Neil, E. J., 

Schneider, D. R, Shaya, E. J., and Westphal, J. A. 1991a, Ap. J. (Letters), 369, 

L45. 
Lauer, T. R., Faber, S. M., Holtzman, F. A., Baum, W. A., Currie, D. G., Ewald, S. 

P., Groth, E. J., Hester, J. J., Kelsall, T., Kristian, J., Light, R. M., Lynds, C. R., 

O'Neil, E. J., Shaya, E. J., and Westphal, J. A. 1991b, Ap. J. (Letters), 369, L41. 
Schweizer, F. 1976, Ap. J. SuppL, 31, 313. 
Paresce, F., Shara, M., Meylan, G., Baxter, D., Greenfield, P., Jedrzejewski, R., Nota, 

A., Sparks, W., Albrecht, R., Barbieri, C., Blades, J. C., Crane, P., Deharveng, J. 

M., Disney, M. J., Jakobsen, P., Kamperman, T. M., King, 1. R., Macchetto, F., 

Mackay, C. D., and Weigelt, G., 1991, Nature, 352, 297. 
Terndrup, D. M. 1988, A. J., 96, 884. 
Weigelt, G., Albrecht, R., Barbieri, C, Blades, J. C, Crane, P., Deharveng, J. M., 

Disney, M. J., Jakobsen, P., Kamperman, T. M., King, I. R., Macchetto, F., Mackay, 

C. D., Paresce, F., Baxter, D., Greenfield, P., Jedrzejewski, R., Nota, A., and Sparks, 

W. 1991, Ap. J. (Letters), 378, L21. 



THE CENTRAL DYNAMICS OF 47 TUCANAE 



Ivan R. King 
Astronomy Department 
University of California 
Berkeley, CA 94720 
U.S.A. 



Abstract. The blue stragglers recently discovered in 47 Tucanae have an excess con- 
centration in the very center of the cluster that is statistically significant at a marginal 
level. Studies of the distribution of the ordinary stars (mainly near the main-sequence 
turnoff) may imply the presence of a population of massive remnants near the center 
of the cluster, although more extensive observations will be needed to confirm such a 
hypothesis. 



1. INTRODUCTION 

In an ingenious use of engineering data — an HST focal run that happened to be 
made, at ultraviolet wavelengths, at the center of of the globular cluster 47 Tucanae — 
Paresce et al. (1991) reported the discovery of 21 blue stragglers in a central field that 
spanned only 44 arcsec. Blue stragglers had been seen at the centers of globular clusters 
before (Nemec and Harris 1987, Nemec and Cohen 1989, Auriere et al. 1990), and a 
suspicion was growing that they might be concentrated there; but the high resolving 
power of HST gave the first opportunity to examine the center of a dense cluster — 
and there they were. It is an easy guess that the blue stragglers concentrate to the 
center because they have a greater mass than that of the dominant main-sequence 
stars. What I want to do here is to pursue this reasoning further, and to examine some 
of the statistics of their detailed density distribution — which will lead to a more general 
excursion into the dynamics of the cluster center. 

Whenever one finds a particular stellar species concentrated to the center of a 
globular cluster, the presumption becomes strong that the stars are there because they 
have a higher mass, have relaxed into equipartition, and therefore have lower velocities. 
In principle one ought to be able to use their degree of central concentration to estimate 
their mass. In practice this is not so easy, because of small numbers, because of the 
smallness of the field that we have covered, and because of some complications that I 
will discuss below. Unfortunately 1 began looking, nt this ])roblem only the week before 
the workshop [and have not been able to return In it since], so this is going to be only 
a sketchy discussion. 



2. THE CENTRAL CONCENTRATION OF THE BLUE STRAGGLERS 

Not only are the blue stragglers concentrated relative to the outer parts of the 
cluster; they are centrally concentrated even within our small field. Of the 21 stars, 9 
fall within the f/96 field, which is only 1/4 of the total. Unfortunately this is not what 
I would call a strongly significant result. In a binomial distribution where each object 
has a 1/4 probability of faUing in the inner field, the probabiUty that 9 or more will do 
so at random is 0.056. 

Another test is their cumulative radial distribution. Here I wanted to compare 
with various distribution laws and did not have time to calculate the areas of annuh 
that did not He completely within the f/48 field; so I considered only the stars out to 
the largest circle around the Meylan center that hes completely within the field. (The 
Paresce et al. paper gives three alternative centers; I beheve that the Meylan center 
is a good estimate of the true center, because of the way in which he determined it. 
But I don't want to go into those details here, and the exact choice of center probably 
does not matter a great deal either.) The largest such circle had a radius of 217 pixels; 
unfortunately it left me with only 14 of the blue stragglers. 

The distribution of Ught near the center of a cluster as concentrated as 47 Tuc is 
well represented by the formula 

f - l^ (1) 

where / is surface brightness and the core radius re for 47 Tuc is a httle over 25 arcsec. 
(In the present units, I used 610 pixels.) I converted this into a cumulative distribution 
function and compared with the blue stragglers by means of a Kolmogorov-Smirnov 
test. The probabiUty of a deviation at least this large, at random, came out 0.053 — 
again a quite marginal result. 

But intuitively the c.d.f. for the blue stragglers really does look different, so I 
asked the question, "with what mass group does it agree best?" For modehng the 
distribution of groups of various stellar mass, I adopted a simple line of reasoning. 
Deep in a potential well a relaxed velocity distribution is quite close to Gaussian, and 
its density distribution is consequently quite close to a Boltzmann law, in which spatial 
densities are proportional to exp( — mC/), where m is the mass of a star in the group and 
U is the potential. This behavior has the consequence that for the relative distributions 
of two mass groups, 

71^2 oc (n-mi)"^^''"' , (2) 

where the n's represent spatial number densities. 

Distributions of the form of the one in Eq. (1) have the property that the spatial 
distribution has the same form, except that the power of the quantity in the denominator 
becomes larger in magnitude by 1/2. For example, the projected distribution in Eq. 
(1) corresponds to the spatial distribution 

If we take Eq. (1) to represent the turnoff stars (since the red giants and the subgiants 
have nearly the same mass, and all these stars together are responsible for nearly all 
of the hght), and if we take the turnoff mass as the unit of mass, then we quickly find 



that the projected distribution of stars of mass m is 

/o 



1 + {r/rcy\ 



(4) 



What I did then was to compute these curves for various values of m, integrate them 
into c.d.f.'s, and compare them with the c.d.f. of the blue stragglers. No statistical tests 
here; just, which one gave the best fit? The answer came out 7 times the turnoff mass. 
This seemed a priori such an implausible result that I decided to look further. 



3. THE DISTRIBUTION OF TURNOFF-MASS STARS 

It occurred to me to ask, are the ordinary stars distributed "right"? That is, do 
they fit the curve expressed by Eq. (1)? The answer was resoundingly that they do not. 
In this case there was no question of small-number statistics; I had 369 stars in the 
circle of radius 217 pixels. The Kolmogorov-Smirnov coiiiparison gave a probability of 
5 X lO-''. 

What is wrong here? We had always considered 47 Tuc to be a showcase-model of 
a smooth, relaxed cluster, with no suspicion of 0.0005-probability shenanigans going on 
at its center. 

One possibihty, of course, is that the data are wrong. We know that our flatfielding 
isn't yet perfect, the star-detection thresholds might not have been everywhere uniform, 
and we observed through some pretty weird bandpasses that might be doing something 
funny to us. 

But if we accept the data, there is a hne of reasoning that leads to an interesting 
interpretation of what we see. First, note that no one else has ever had a chance to look 
at this phenomenon before. Ground-based density distributions near the center of 47 
Tuc have had to depend on photoelectric photometry of the integrated light of all the 
stars in any area that we study. This light is dominated by that of the red giants, and 
they are so few in number that the Hght distribution is statistically very much noisier 
than the distribution of the much more numerous faint stars that HST allows us to 
study. For the stars we see a resounding statistical effect, but in the distribution of 
integrated Ught it would not show up at any sigitificant level at all. 

Furthermore, the distribution that we see makes good physical sense. A cluster 
like 47 Tuc must contain a number of massive remnants — the presence of more than 10 
pulsars is ample evidence of that — and because of their mass they should congregate at 
the center. They create an additional potential well there, and Boltzmann tells us that 
that deepening of the potential should increase the density in the center. 

To make this more quantitative I computed some models. They have the observed 
mass function of 47 Tuc (as verified by a projection of the model at the radius at which 
the mass function was determined), and they have the additional property that I can 
arbitrarily add massive remnants, in the equipartition density distribution that their 
mass implies. They do indeed concentrate very strongly to the center. For a plausible 
value, I took their mass to be 1.4 Mq. I haven't tried very many models, nor fitted 
them in great detail, but it looks as if a proportion of about 2-3% of the stars as 
massive remnants will cause the visible stars to have a distribution such as the one that 
we observe. 

Such a model would also make a siguificniil change in the fitting of the radial 
distribution of the blue stragglers. With part of their concentration at the very center 



explained by the extra potential well due to the massive remnants, the distribution of 
the blue stragglers would no longer imply such implausible individual masses. 

But let me emphasize again that this discussion is based on a single field, observed 
in unconventional color bands, and possibly subject to systematic errors. Whereas I 
beheve that the reasoning given above is sound, and that massive remnants should be 
expected a priori, I would be very reluctant to make any strong assertions until we 
have a larger radial range in 47 Tuc, and, hopefully, other clusters observed at HST 
resolution. 



4. CONCLUSIONS 

It would be tempting to foUow this hne of reasoning, in order to find out how many 
massive remnants he hidden at the center of 47 Tucanae, and then to use the model to 
determine from their distribution the masses of the blue stragglers. But at that point 
we would be out on the end of a long, shaky hmb of reasoning. For the time being, 
I would prefer to regard this as a provocative non-result, to be taken up again when 
better data become available. What we see and deduce may be real and correct; but 
on the other hand, the moral of the story may be that a person who squeezes too hard 
on rough data will only end up by hurting his hands. 



REFERENCES 

Auriere, M., Ortolani, S., and Lauzeral, C, 1990, Nature, 344, 638. 

Nemec, J. M., and Cohen, J. G., 1989, Ap. J., 336, 780. 

Nemec, J. M., and Harris, H. C, 1987, Ap. J., 316, 172. 

Paresce, F., Shara, M., Meylan, G., Baxter, D., Greenfield, P., Jedrzejewski, R., Nota, 

A., Sparks, W., Albrecht, R., Barbieri, C., Blades, J. C., Crane, P., Deharveng, J. 

M., Disney, M. J., Jakobsen, P., Kamperman, T. M., King, I. R., Macchetto, F., 

Mackay, C. D., and Weigelt, G., 1991, Nature, 352, 297. 



A TALE OF THREE JETS 



F. Macchetto* 

and the FOC Investigation Definition Team 

Space Telescope Science Institute 

3700 San Martin Drive 

Baltimore, Maryland, 21218 

USA 

* Associated with the Space Science Department of ESA 



1. INTRODUCTION 

The study of the optical counterparts to radio jets wiU be the subject of an intensive 
investigation program with the European Space Agency's Faint Object Camera first 
proposed 10 years ago (Macchetto 1981, Miley 1981). It has been known for a long 
time that radio jets are widespread in active galaxies (e.g. Miley 1980, Bridle h Perley 
1984) and that they must play a fundamental role in galaxy activity and in the transport 
of energy from the nucleus to the radio-emitting lobes (Rees 1971, Blandford & Rees 
1974). The jet radiation is beheved to be synchrotron emission, but there are stiU only a 
few cases where the optical counterparts of radio jets have been detected (e.g., Butcher, 
van Breugel & Miley 1980, Fraix-Burnet, et.al., 1991). These extragalactic optical jets 
have a number of unfavorable observational characteristics that make it very difficult 
to understand their physical nature. The jets are rare, very faint, relatively small and 
are embedded in the bright stellar background of the parent galaxy. 

Nevertheless, since the optical emission appears to be the continuation of the radio 
synchrotron spectrum, optical observations are essential to define the physical parame- 
ters and constrain the jet emission models. Observations with the Faint Object Camera 
allow many of the observational disadvantages to be overcome. The much greater spatial 
resolution provided by the FOC, compared to even the best ground based optical tele- 
scopes, means that the relative contrast of the jet to the underlying stellar background 
is improved by at least an order of magnitude. In addition observations in the ultravio- 
let provide another order of magnitude improvement in contrast, since the brightness of 
the starhght falls rapidly into the UV, while the steepness of the jet emission is typically 
flat (a ~0.5). 

Furthermore, high spatial resolution observations in the visible and ultraviolet, offer 
the possibiUty to determine the precise location where particle acceleration occurs, 
since the electron Ufetime for optical emitting electrons is extremely short. Finally, 
comparison of radio and optical morphologies obtained with similar resolution wiU 

10 



allow the investigation of confinement mechanisms and diffusion processes within the 
relativistic plasma. 

To date, the FOC has observed at high spatial resolution the optical counterparts to 
the radio jets in PKS 0521-36; 3C66B and M87. The results show unexpected features 
and point at different physical mechanisms at work. In the next section, I wiU tell the 
tale of these three jets. 



2. PKS 0521-36 

One of the most prominent radio and optical jets is that found in the elliptical 
galaxy PKS 0521-36, a relatively isolated radio galaxy at a redshift z = 0.055 which 
also harbors a bright V = 16 BL Lac nucleus and extended optical line emission. De- 
tailed spectroscopic and morphological studies have been carried out by Danziger et 
al. (1979, 1983); Cayatte k Sol (1987); and Boisson, Cayatte, & Sol (1989). Recently 
Sparks, Miley, & Macchetto (1990) reported optical polarization measurements of the 
jet and nucleus, which confirmed the expected high polarization if the emission is due 
to synchrotron radiation. 

Images of PKS 0521-36 were obtained on 1990 August 25 with the FOC using the 
F430W and F320W filters and f/96, 512 x 512 mode (Paresce 1990) with a correspond- 
ing pixel size approximately 0".022. Pointing was defined using the radio VLBI position 
obtained by Morabito et al. (1986). 

The galaxy has a very bright, V ~ 16, BL Lac nucleus which severely tests the 
abihty of HST and the FOC to detect faint structure in the vicinity of bright objects. 
Accurate knowledge of the point-spread function appropriate to the data is essential 
to remove the halo around the nucleus which arises from the presence of spherical 
aberration in the HST primary mirror. 

Various techniques were tried to remove the effects of the nucleus, from straightfor- 
ward substraction of a scaled pointspread function (PSF), direct Fourier deconvolution, 
maximum entropy deconvolution, and Lucy's (1974) iterative deconvolution technique. 
A combination of point source subtraction and Lucy's method gave the best results. 
Specific PSF observations were obtained of the star BPM 16274, a UV flux standard 
(Bohlin et al. 1990; Turnshek et al. 1990). 

An example of the resultant deconvolutions is shown in Figure 1. This uses Lucy's 
(1974) iterative deconvolution technique which constrains the result to be positive. 

The VLA contour data as published by Keel (1986) is shown superposed on the 
FOC data in Figure 2. It is immediately obvious that the FOC data has a resolution 
very similar to that of the VLA data, but shows considerably more morphological 
information that the ground-based optical data. 

The FOC image shows a bright knot located ^ I'.'S to the NE and clearly resolved 
as in the VLA data. The width of the knot is ^ 0!'8. Beyond this bright knot, the jet 
has approximately constant surface brightness and a morphology similar to the VLA 
image with a total length of 6!'5. The jet is also resolved in width, 0(6 wide in the 
fainter regions of the jet, with Uttle or no evidence of structure on a scale of < 0. 1. 
The FOC data appears to show more flux than the VLA data in the region at sUghtly 
larger radius from the nucleus but close to the southern tip of the knot. 

There is a second component seen in the deconvolution at 0!'36 from the nucleus, 
Figure 2. It is visible in the raw data, although by no means as clearly as in the de- 
convolved image. Such a source would be completely within the core of the highest 
resolution VLA contour map published by Keel. We beHeve we have detected a pre- 

11 




1. The jet in PKS 0521-36 as observed by the FOC. An elliptical model of the galaxy 
has been subtracted. The jet is approximately 6.5 long. 



12 




2. The VLA contour data of Keel (1986) is shown superposed on the FOC data. Note 
the excellent agreement of the two data sets. 



13 



viously unseen inner jet structure at a distance corresponding to 300 pc, Ho = 75 km 
s-^ Mpc-^ 

FOC observations of the nucleus of PKS 0521 - 36 seem to have resolved it into a 
bright nucleus and an inner jet extension. Presumably, the optical polarization mea- 
sured by Sparks, Miley & Macchetto (1990) comes from both these components. The 
fact that the nuclear polarization is large and perpendicular to the jet direction (Sparks 
et al. 1990; Angel k Stockman 1980; Bailey, Hough & Axon 1983; Bridle et al. 1986) 
is consistent with what has been observed in other quiescent blazars and quasars. The 
second component close to the nucleus may be an inner extension of the jet, since it 
lies on the geometrical projection of the jet toward the nucleus. However, since this 
extension falls entirely within the VLA radio core, there is no independent way of con- 
firming its existence. If real, it could be a site of electron acceleration along the jet due 
to transverse shocks (Drury 1983; Blandford & Eichler 1987). 

The large bright knot further along the jet is a clear counterpart to the radio knot. 
The radio and optical polarization position angles (Keel 1986; Sparks et al. 1990) 
suggest a magnetic field aligned along the jet direction. This knot is unresolved in the 
optical polarimetric measurements, but the general sense of the magnetic field is stiU 
along the jet direction at that position. The bright knot is clearly an important site 
where particle acceleration is occurring. 

The general optical morphology of the jet does not exhibit any significant degree of 
dumpiness even at FOC resolution. 

Using the standard formula of Rybicki & Lightman (1979), we derive a mean Mfetime 
for the electrons: 

^1/2 = i^-^:^ y^ 

where B is the magnetic field in Gauss and 7 is the Lorentz factor. The mean 
distance for electron difussion is 

^1/2 - ^^2 l^P^ 

where 1/ = I2OB7 /10~ is the cut-off frequency. With typical value B ^ 10~ , the 
electron difussion distance is D ^ 200 pc in the optical and D as 100 kpc in the radio 
regioii. The corresponding hfetime for the optical electrons is t^M ~ 600 yr. 

This imphes that there must be continuous acceleration along the jet of the electrons 
responsible for the optical emission, since electron diffusion from the bright knot could 
not account for the observed optical extent. 



3. 3C 66B 

3C 66B is a relatively nearby bright radio source associated with a 13th magnitude 
galaxy at a redshift of 0.0215 (Matthews, Morgan, &: Schmidt 1964). The galaxy Hes in 
a small group close to the cluster AbeU 347, at a distance of 86 Mpc. At this distance 
an angular scale of OCl corresponds to a projected linear size of 41 pc. 

The radio source has been the subject of several comprehensive studies (e.g., Northover 
1973; Miley & van der Laan 1973; van Breugel &; Jagers 1982; Leahy, Jagers, &: Pooley 
1986). Its structure is intermediate between that of an edge-darkened double and a 
head-tail morphology, indicating that the morphology may be affected by motion of 
the parent galaxy through an ambient medium (Miley et al. 1972). When mapped 

14 



at sufficient resolution, the radio emission from the nuclear "head" apears jetlike and 
one-sided. An optical counterpart of the jet in 3C 66B was detected by Butcher et al. 
(1980) and further studied by Fraix-Burnet et al. (1989b). 

Images of 3C 66B were obtained with the Faint Object Camera (FOC) on 1990 
September 23 using the F320W filter in the f/96 512 x 512 mode (Paresce 1990) and 
reduced as in Macchetto et al. (1991). The filter bandpass is about 900 A FWHM 
centered at 3360 A. The pixel size after correction for geometric distortion is approxi- 
mately 0'.'022, ^ 10 pc, and the field size approximately 11" ^ A kpc. Two exposures 
of 1500 s duration were made with this filter, both in fine lock, resulting in a tracking 
accuracy of 0!'007 rms. 

The galaxy was clearly visible in the raw data, but in order to investigate the jet 
alone, a model of the underlying galaxy hght was substracted from the deconvolved 
image (Fig. 3) The model was obtained by fitting eUipses to isophotes and creating 
an image with exactly elliptical isophotes having the same parameters as those of the 
galaxy. This process removed some flux from the nucleus and is responsible for the 
residual "arc" opposite the bright knot in the jet, which slightly distorted the fitting of 
the eUipse and, therefore, led to an incomplete galaxy subtraction opposite it. 

In order to compare our HST image with the best VLA map of 3C 66B, we smoothed 
the deconvolved and galaxy-subtracted optical image with a circular Gaussian function 
having cr = 3.1 binned pixels (FWHM ^ 0"35; Fig. 4a). Figure 4b shows a contour 
map of the "A-configuration" VLA map of 3C 66B obtained at 6 cm wavelength by 
Leahy et al. (1986). 

Several conclusions can be drawn from these figures. 

1. The similarities between the jet as seen in the "raw" and deconvolved images gives 
us confidence that the deconvolution process is not introducing any severe artifacts. 

2. The similarities between the smoothed optical image and the radio map imply 
that on a scale of 0!'3 the radio and optical emission have very much the same 
radio-to-optical spectral index all along the jet. 

3. On the scale of the HST resolution, the jet of 3C 66B is filamentary. Filamentary 
structure has not previously been seen for optical jets, although it has been observed 
on a scale of 0''l for the M87 jet in the radio (Owen et al. 1989), (and now with 
the FOC, see Section 4) 

4. Two distinct "strands" appear in the HST image. These can be traced from > 3. 7 
(1.5 kpc) from the nucleus out to a distance of 7!'6 (3 kpc), where they disappear 
into the noise. 

5. The separation between the strands varies between about 0'.'3 and 0!'4, that is, 
about 150 pc. 

6. The strands appear to undergo sharp "kinks" at distances of 2'.'5 (1.0 kpc) and 6"2 
(2.5 kpc) from the nucleus. 

7. An additional small bright feature is visible in the HST image off the jet, 5.5 (2.3 
kpc) to the north of the nucleus. It is unclear whether this is radiation from a 
compact region within the 3C 66B galaxy or whether it is due to an unrelated 
object. It may be related to the "blue knots" found by Fraix-Burnet et al. (1991). 

It is interesting to compare our observations of 3C 66B with the high-resolution radio 
data of Owen, Hardee &; Cornwell (1989) and our FOC high resolution observations of 

15 




3. Deconvolved and galaxy-subtracted image of the jet of 3C 66B. The jet is 7% (3 
kpc) long. The FOC resolution of 0!'l corresponds to 40 pc at the distance of the 
gcdaxy. 



16 




- 42 46 



- 42 45 56 



42 46 4 



— 42 46 



-^ 42 45 56 



2 20 2.50 



2 20 2.00 



Contour map of 3C 66B jet smoothed with a Gaussian "beam" with <t = 3.1 
pixels (0!'44 each). The contours are at intervals of 1.12 counts from 1.12 to 11.2 
counts and intervals of 3.73 counts from 14.92 to 37.3 counts, corresponding to the 
contours of Leahy et al. 1986 for a nominal flux conversion and spectral index a 
= 1.1. Note that the spatial registration and scale are only approximate for the 
HST data. The coordinates' R.A. and Decl. axes are from Leahy et al. (bottom) 
The 5 GHz VLA map at comparable resolution to the optical image of (top) (from 
Leahy et al. 1986). 



17 



M87 (next section). The M87 jet also shows complex filamentary structure, with limb 
brightening over most of its length. However, it appears to be more sharply bounded 
than the jet in 3C 66B. 

Although, unlike M87, 3C 66B is not a cD galaxy at the center of a rich cluster, 
there are several similarities between M87 and 3C 66B. Their total radio luminosities 
are comparable, but the jet in 3C 66B is sUghtly longer and can be traced out to about 
3 kpc from the nucleus, a factor of 2 farther than in M87. Both jets appear filamentary 
and edge brightened. However, in 3C 66B, although further measurements are needed 
to confirm the fainter features, the double-stranded filaments appear to be embedded in 
a broader structure, which is not accurately cohnear with the filaments. Also, although 
the 3C 66B jet is filamentary, the pitch angle of the filaments is quite different from that 
in the filaments in M87 - we do not see the tightly wound hehx that the M87 shows. 

A very basic question posed by these observations concerns constraints on the re- 
gions where the synchrotron-radiating electrons are accelerated. Does the presence of 
optical radiation along the jet imply that locahzed particle acceleration is required, or 
could the electrons be accelerated in the nucleus and transported to the sites of the 
radiation? 

The average age of the radiating synchrotron electrons, with typical values for the 
magnetic field strength and cut-off frequency, is less than 1000 yr for the optical elec- 
trons, while the light travel time from the nucleus to the end of the observed optical jet 
is about 10 yr. Hence, at first sight, locahzed particle acceleration would be required. 

However, as discussed later for M87, in situ acceleration is not needed in a two fluid 
model, (Pelletier and Roland 1989). In this case, a relativistic flow of electron-positron 
plasma moves in a channel through the jet. The external non-relativistic jet of electron- 
positron plasma carries most of the mass and kinetic energy. A mixing layer occurs at 
the boundary of these two fluids, and synchrotron radiation can be produced at this in 
interface. 

The other new features present in our observations are the sharp bends or "kinks" 
in the double-stranded filaments. The origin of these kinks is unclear. The fact that 
they are mimicked in more than one filament suggests that they are not due to an 
instabihty mode in an individual filament. They may well trace out irregularities in 
the ISM of the galaxy and/or be due to time-dependent variations in the power of the 
nuclear machine responsible for producing the jet. 

Another possibihty is that these bends and kinks are the result of observing a jet 
which contains magnetic field modes which are at an angle to the jet. Models for force- 
free magnetized jets have been calculated by Konigl Sz Choudhuri (1985) containing 
two magnetic field orientations, one parallel to the jet axis and the other a double 
hehx twisted field. Both our observed filamentary edge-brightened configuration and 
the presence of bends are at least consistent with their model. 

4. M87 

The EO galaxy M87 harbours the prototypical and most studied example of an 
optical jet. First observed by Curtis in 1918, it remained httle more than a curiosity, 
until Baade and Minkowski studied it in 1954 and first used the term "jet" to describe 
the sequence of optical knots extending to about 20" from the nucleus. Since then, the 
jet has been observed at radio (Owen, Hardee &: CornweU, 1989, Biretta, Stern & Harris 
1991), optical (de Vaucouleurs & Nieto 1979, Keel 1988, Fraix-Burnet, Le Borgne & 
Nieto 1989) and X-rays wavelengths (Schreier, Gorenstein & Feigelson, 1980). 

18 



The radio and optical morphologies and polarization structure of the jet are similar 
(Schlotelburg et al, 1988) to within the resolution Umits of the ground based obser- 
vations. These results are best explained by emission from synchrotron radiation. In 
addition, the emission detected at x-ray wavelengths in the jet region also suggests that 
the synchrotron spectrum extends to high frequencies. 

The origin of the optical continuum emission in radio jets can best be tested in 
M87, since it is the only jet which shows both a spectrum break and x-ray emission. 
Observations in the ultraviolet will help determine the exact frequency at which the 
break occurs. This, in turn, determines the value of the magnetic field and turbulences. 

Optical and ultraviolet polarization observations at high resolution when compared 
to equivalent radio data will help determine the precise location and nature of the 
acceleration process, since the lifetime of these electrons, and therefore, their travel 
distance, is very small. 

Images of M87 centred on the nucleus and positions along the jet were obtained with 
the FOG, utiHzing the foUowing filters in the f/96 512 x 512 modes: F120M, F140W, 
F220W (direct and with polarizers POL0,POL60,POL120), F430W (with polarizers 
only) and F501N. Additionally, a zoomed acquisition image was made with the F372M 
filter. The exposures were made in fine lock, with an expected tracking accuracy of 
0!'007. The F120M image has low signal and is very noisy, but all other images of the 
nucleus were useful; adequate signal strength to give a significant impression of the jet 
were obtained only with the F140W, F220W, F372M and F430W filters. The F220W 
images are particularly valuable, because of the high contrast of the jet against the fight 
from the stellar population of the galaxy which is very weak in the ultraviolet region. 

The F220W exposures were corrected using the Lucy (1974) deconvolution tech- 
nique to produce the final deconvolved images with minimum beam size 0"2 FWHM. 
The results of this restoration are shown in Figure 5. This figure shows side-by-side the 
deconvolved FOC and VLA data. As is easily observed, two FOC frames (each of 11" 
X 11") are needed to cover the jet which extends for over 20" in length. 

The complex structure and the wide range of intensities make it very difficult to 
display all the features in a single picture. Furthermore, the value of the relative 
intensity of the ultraviolet and radio data is arbitrary and was chosen only to show the 
most prominent features. Detailed intensity comparisons and determination of spectral 
indices as a function of position along the jet will be carried out in the near future. 

The FOC observations demonstrate for the first time that the radio and ultraviolet 
brightness distribution is generally the same over a scale of about 0. 1 or about 10 pc 
(M87 presents 78 pc arcsec" at an assumed distance of 16 Mpc). 

The FOC data shows that all the prominent optical knots (A,B,C, etc) have now 
been fully resolved and show the same remarkable structure as the radio data. (See 
Owen, Hardee and CornweU, 1989, also for the knot notation). The jet is Hmb bright- 
ened, shows very well-defined edges along its conical structure with an opening angle 
of ~ 6.5" and has a tight filamentary structure. The appearance of this structure indi- 
cates that the filaments are wrapped around the jet with pitch angles of about 30" - 40" 
between the nucleus and knot I. At knot A, and between knots A and B, the filaments 
are more tightly wrapped with a pitch angle between 80" and 90" and decreasing from 
knot A towards knot B. The pitch angle may increase again towards knot C. 

SHces taken across the jet at different locations show prominent hmb brightening 
similar, but not identical, to the radio data. Detailed comparisons are beyond the scope 
of this paper, but this general agreement provides strong evidence in favor of the two 
fluid model for this jet. (Pelletier & Roland, 1988, Owen, Hardee & Cornwell 1988). 



19 




5. Left. Deconvolved FOC image of the M87 jet. Two FOC images are needed to 
cover the length of the jet (~ 20"), hence the small gap in the spatial coverage. 
Right. Deconvolved VLA image kindly supplied by F. Owen and J. Biretta. 
The images show the remarkable agreement of the ultraviolet and radio observa- 
tions. 



20 



In this scenario, the jet consists of a cone around which one or more bright filaments 
are wrapped. The optical and radio emission comes mostly from a surface layer in which 
these filaments are embedded. The synchrotron lifetimes in the ultraviolet are typically 
only of the order of 100 yr, corresponding to light travel times of about 30pc. This is 
comparable to the width and presumably the thickness of the optical and radio strands. 

The low emissivity in the jet's core indicates that the energetic particles are not 
suffering significant synchrotron losses. This is also compatible with the model in which 
the jet's core has a relatively low magnetic field. In this case, the high-energy particles 
can be produced in the central black-hole and propagate along the jet's interior in a 
low magnetic field region, thereby suffering only modest synchrotron losses. As they 
diffuse across the jet and into the high magnetic field boundary layer, they can produce 
the optical emission without the need for in situ acceleration. 

Several mechanisms can be invoked to explain how the emission from this boundary 
layer is produced. Non-linear evolution of synchrotron instabilities has been proposed by 
Bodo et al, 1991, to explain the formation of filaments in jets out of equipartition, such 
as that in M87, where the energy of the relativistic electrons exceeds that of the magnetic 
field. They find that in a plasma subject to constant heating, after an initial phase in 
which the instability growth rate follows the Unear model, the instabiUty reaches a 
quasi equilibrium state on timescales of the order of several synchrotron timescales. 
This mechanism can explain the formation of filaments of enhanced emission observed 
in the lobes and jet of M87. 

Optical synchrotron emission can also be produced through a diffusive shock accel- 
eration mechanism (Fraix- Burnet, 1991). This process is so efficient that requires the 
magnetic field turbulence to be quite low. The source of energy of this turbulence could 
be the kinetic energy of the jet which can be transferred to the magnetic field (or to 
the plasma) through the interaction of the jet with the interstellar medium. 

The answer to which of the competing mechanisms and scenarios are at work and 
the determination of the relevant physical parameters will have to wait further detailed 
analysis of the FOC optical, ultraviolet and polarization data. 



5. CONCLUSIONS 

The study of the optical counterparts to radio jets with the Faint Object Camera on 
board the Hubble Space Telescope has already produced new and unexpected results. 

The jet in PKS 0521-36, which is the most distant, has been fuUy resolved. Because 
of its length, magnetic field configuration and optical morphology, it seems to require 
reacceleration sites for the optical electrons. These could well be provided by shocks at 
the site of the brighter knots observed. 

For both 3C 66B and M87, we have observed, for the first time at optical wave- 
lengths, a filamentary structure which is similar to the radio data. In this case, we 
conclude that the emission comes from a boundary layer where the filaments and the 
strong magnetic field are located. 

One of the most puzzUng, but fundamental, results that must be explained by 
models of particle acceleration is why within an object, over more than five decades 
of frequency and a large variation of physical conditions, the old and young electrons 
have similar spatial distributions, although large differences are observed from object 
to object. 

Our FOC observations of PKS 0521-36, 3C 66B and M87 show that, even with its 
present aberration problems, the HST is a uniquely important instrument for studying 

21 



synchrotron jets. Future observations with the HST of other extragalactic jets should 
provide fundamental information about the nature of collimated activity in galaxy 
nuclei. 



6. ACKNOWLEDGEMENTS 

The FOC/IDT members are: Rudolf Albrecht, Cesare Barbieri, David Baxter, J. 
Chris Blades, Alec Boksenberg, Phil Crane, Jean Michel Deharveng, Michael J. Disney, 
Perry Greenfield, Peter Jakobsen, Robert Jedrzejewski, Theo M. Kamperman, Ivan 
R. King, F. Duccio Macchetto, Craig D. Mackay, Antonella Nota, Francesco Paresce, 
William Sparks, and Gerd Weigelt. 

We wish to thank Drs. F. Owen and J. Biretta for making available to us the VLA 
radio data for M87, and Dr. Keel for the PKS 0521-36 data. 



22 



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Bailey, J., Hough, J.H., & Axon, D.J., 1983, M.N.R.A.S., 203, 339 

Baade, W. k Minkowski, R. 1954, Ap. J., 119, 221 

Biretta, J. A., Stern, C.P. & Harris, D.E. 1991, A. J., 101, 1632 

Blandford, R.D., & Eichler, D., 1987, Phys. Report, 154, 1 

Blandford, R.D., k Rees, M.J., 1974, M.N.R.A.S. 169, 395 

Bodo, G., Ferrari, A., MassagUa, S., Rossi, P., Shibata, K., k Uchida, Y. 1991, Astr. 

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BohUn, R.C., Harris, A. W., Holm, A.V., k Gary, C. 1990, Ap. J. Suppl., 73, 413 
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Bridle, A., k Perley, R.A. 1984, Ann. Rev. Astr. Ap., 22, 319 
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Butcher, H.R., van Breugel, W.J.M., k Miley, G.K. 1980, Ap. J., 235, 749 
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A. Ferrari k A.G. Pacholczyk (Dordrecht:Reidel), 131 
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M.A.C., k di Sergo AHghieri, S. 1991, A. J., 101, (1), 88 
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Leahy, J. P., Jkgers, W., k Pooley, G.G. 1986, Astr. Ap., 156, 251 
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Macchetto, G. Miley, k M. Tarenghi (Noordwijk:ESA), 15 
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Macchetto, G. Miley, k M. Tarenghi (Noordwijk:ESA), 9 
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269 
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91, 546 

23 



Northover, K.J.E.D. 1973, M.N.R.A.S., 165, 369 

Owen, F.N., Hardee, P.E., k CornweU, T.J. 1989, Ap. J., 340, 698 

Paresce, F. 1990, FOC Instrument Handbook, Space Telescope Science Institute 

PeUetier, G., & Roland, J. 1988, Astr. Ap., 196, 71 

Rees, M.J. 1971, Nature, 229, 312 

Rybicki, G.B., k Lightman, A. P. 1979, Radiative Processes in Astrophysics (New York: 

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Schl6telburg, M., Meisenheimer, K., & R6ser, H.J. 1988, Astr. Ap., 202, L23 
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Sparks, W., Miley, G., k Macchetto, F. 1990, Ap. J., 361, L41 
Turnshek, D.A., Bohlin, R.C., Williamson, R.L., Lupie, O.L., Koornneef, J., k Morgan, 

D.H. 1990, A. J., 99, 1243 
van Breugel, W.J.M., k Jigers, W. 1982, Astr. Ap. Suppl, 49, 529 



24 



EARLY OBSERVATIONS OF GRAVITATIONAL LENSES WITH 
THE PLANETARY CAMERA OF HUBBLE SPACE TELESCOPE 



Jerome Kristian 
Carnegie Observatories 

for the 

Wide Field/Planetaiy Camera 
Investigation Definition Team " 



* S.P. Ewald, E.J. Groth, J.J. Hester, J.A. Holtzman, T.R. Lauer, R.M. Light, D.P. 
Schneider, E.J. Shaya, W.A. Baum, B.G. Campbell, A.D. Code, D.G. Currie, G.E. 
Danielson, S.M. Faber, J.G. Hoessel, D. Hunter, T. Kelsall, C.R. Lynds, G. Mackie, D.G. 
Monet, E.J. O'Neil, Jr., P.K. Seidelmann, B. Smith and J.A. Westphal. 



During Cycle 0, we obtained preliminary data for six "classic" gravitational lens systems, as 
part of the integrated WFPC IDT observing program. The lens observations were intended 
as a reconnaissance and evaluation program, to assess, with real data, what can usefully be 
done with the degraded point spread function (PSF) and to decide whether more extensive 
HST observations would be productive. 

All observations were made with the Planetary Camera, in the team's standard V and I 
bands (F555W and F785LP). Two filters were used because the quasar images and the 
lensing galaxies have very different colors; the V band emphasizes the quasar images, and 
the I band the galaxies. Exposure times were chosen so as to not saturate the brighest 



25 



images, and also to fit within single dark-side orbits, to maximize spacecraft efficiency. 

The general overall goals of our lens program are (1) To test the validity of gravitational 
imaging theories by observationally characterizing the lens systems as completely as possible 
[the mass distribution and redshift of the lensing galaxy and the intensities and geometry of 
all of the quasar images, including the expected but so far undetected (n+ l)th images], and 
comparing the data with theoretical models. (2) To use individual lens systems as tools to 
obtain cosmologically interesting information - small and large scale galaxy mass 
distributions and mass/light ratios, HO, qO, etc. 

The results of the Cycle observations are rather encouraging. Useable data were 
obtained for all six systems, with at least modest new results for four of them, and it appears 
that many of the original objectives of the program can eventually be achieved, albeit with 
a considerable increase of exposure time. Lens observations are one of the few deep space 
programs suited to the telescope in its present condition: the most important information 
lies in the morphology and relative brightness of the images, rather than in their absolute 
magnitudes, which have so far proved difficult to measure with HST; there is structure on 
angular scales which requires HST resolution, but the structure is not so complex as to 
confuse the data with badly overlapping PSFs and the like; and the lens images are for the 
most part relatively bright, so the signal to noise is great enough to allow at least partial 
correction of the PSF. 

The data are still being analyzed by the IDT; the results and the data themselves will be 
published elsewhere. Below we give just a few comments on each of the objects. 



26 



Q0957 + 561 

This, the first lens discovered, is still not well understood in spite of (or perhaps because of) 
the considerable amount of data available, both in the radio and optical. Optically, it has 
been known since the earliest days to consist of at least 3 components: two quasar images 
separated by 6 arc seconds and a lensing galaxy very close to one of the images. 

The new HST data show the system in much better spatial detail. The three images are 
completely resolved for the first time, and the galaxy is bright enough that, even with the 
relatively short exposures, a rough core radius can be obtained, which should further 
constrain the theoretical models. 

Deeper HST observations will provide detailed information on the structure of the lensing 
galaxy and, with luck, the location and intensity of the 3rd image, both of which will be of 
importance in finally understanding the system. 



PG1115 + 080 

PG1115, the "triple quasar", was the second gravitational lens discovered (Weymann et al 
1980, Nature 285, 641), and the first and so far most interesting to be observed in our 
program. Over the decade following its discovery, some details of the system were found 
from the ground in a remarkable series of observational tours de force, using speckle 
techniques and direct imaging in extremely good seeing (see, e.g.. Christian et al 1987, Ap.J. 
312, 45, and references therein). Following a suggestion based on early imaging data and 
theoretical models, it was found that the brightest of the three components was itself double; 
the lensing galaxy was detected (although there is some disagreement in the literature as to 
its location); and variability between components, on a time scale of months, was seen. 



27 



The morphology was revealed in its full glory with the first raw HST pictures. The system 
is completely resolved, including the two bright components (separation 0.5 arc sec) and the 
lensing galaxy's location and brightness is easily and unambiguously seen in the raw data. 
The quality of the data, the accuracy and convincingness of the results, and the ease of 
obtaining them compared with the ground-based results, which are the product of 
formidable skill, expertise, dedication and labor, are a dramatic illustration of the power of 
HST for this kind of observation. 

This system also provided a useful test case for various techniques to correct for the point 
spread function, in order better to see the lens galaxy. It was found that, at least in this 
case, the most effective technique was to remove the quasar images by subtracting a 
properly scaled bright image taken near the same time. This was found to be better than 
using theoretical PSFs or deconvolutions of various kinds, although in practice it was rather 
difficult to do (cf. the poster paper by E.J. Groth at this Workshop). 

The 1115 system is relatively simple, as the known lenses go, and shows promise of being 
completely understood. Modeling of the Cycle data by Schneider and Kristian suggest that 
the observed galaxy is capable by itself of producing the observed images, and what is 
required for a complete understanding of the system is a better knowledge of the distance 
and the mass distribution of the lens galaxy, and the location and brightness of the expected 
5th image. Further HST observations will produce the second and third of these three data. 
The distance will have to await spectroscopy, perhaps with the FOS, but in the meantime, 
plausible and useful estimates can be made from the brightness and structure of the galaxy. 

PG1115 is at present perhaps the best candidate for estimating the Hubble constant from 
measurements of time delays among the images. It and 0957 (which can be done from the 
ground), are the overall best candidates, and 1115 is better in some ways. It is simpler 
(0957 is complicated by the presence of a rich cluster, whose lensing effects are substantial 
and poorly known) and appears to be well on its way to being completely understood, and 
there are more images, which provide much tighter observational constraints and less 
ambiguity in matching intensity changes between images. Ground-based observations show 



28 




}^^^rm^^Mm 



PG1115 - upper left, "raw" data; upper right, Lucy deconvolved; lower left, "raw" with 
quasar images subtracted; lower right, Lucy deconvolved with stretch set to display image 
cores. 



29 



variability in 1115 on time scales of months, which is consistent with theoretical predictions 
and convenient for measurement. The system is too compact and geometrically complex for 
accurate photometry from the ground, but the Cycle data show that such measurements 
can easily be done with HST. 



1634.9 + 267 

This is a faint 4 arcsec quasar pair, which has been suggested as a lensed system on the basis 
of detailed similarities in the spectra, although no lens has been detected. 

The Planetary Camera data shows two point sources vkdth a possible very faint 3rd image 
whose reality remains to be verified. 



MG2016+112 

The lens system MG2016+112 was discovered in 1984 (Lawrence et al, Science 223, 66). 
Since then it has been observed extensively both in the visible and radio, mostly by 
Schneider and a host of distinguished colleagues in the MG lens survey group, who pushed 
ground-based techniques to their limits. It is now known to be composed of three radio 
sources and involves perhaps as many as 7 optical objects, 4 of them galaxies or extended 
emission regions. 

Schneider et al saw evidence for a 3rd QSO image, which is one of the few even suggested 
in the literature, but it has never been confirmed. 

Both ground-based and HST observations are difficult because of the small angular scale 
and faintness of the system. In spite of its observational difficulties, however, it is important 
because it provides a very stringent test of lensing theories. 



30 



Because of the faintness of the objects, our hmited Cycle data do not go deep enough to 
add anything new or substantive to the existing data, but they indicate that longer exposures 
will be able to do so. The two lens components are seen clearly in the V band, and faintly 
in the I band. Remarkably, the I band data also show D, one of the two potential lensing 
galaxies, at a redshift of 1. This is likely the most distant galaxy detected by HST in its early 
days. 



2237 + 0305 

Also known as the "Huchra lens" or the "Einstein cross", this famous system is composed of 
a distant QSO aligned almost precisely with the nucleus of a nearby galaxy. The galaxy 
splits the QSO into four images arranged in a rough square 1 arcsec on a side. A fifth 
image has been reported on the basis of ground-based data in a remarkable recent paper 
by Racine (A.J. 102, 454, 1991). 

The system is completely resolved by the Cycle PC observations. Unfortunately, the 
exposures are too short to definitively confirm or rule out the reported 5th image, although 
deeper exposures will be able to do so. 



2345 + 007 



There is an ongoing argument about whether or not this double quasar is a lens system. No 
lensing galaxy has been seen from the ground, although two groups have reported 
(somewhat different) small-scale structure in one of the images. The absence of an 
observable lens has led to a suggestion of imaging by dark matter. The system is of 
considerable interest because the 2 spectra have several Ly-alpha absorption systems in 



31 



common, which makes it a potentially powerful probe of Ly-alpha clouds. 

The Cycle PC data do little to resolve the uncertainty. They show only a pair of point 
sources, with no sign of any structure or other objects. The new data definitely rule out the 
presence of structure in either of the images such as has been reported earlier. The failure 
to detect a lensing galaxy was expected from the much deeper ground-based data which 
exists. 



32 



ULTRAVIOLET SPECTROSCOPIC STUDIES 

OF THE INTERSTELLAR MEDIUM 

WITH THE 

HUBBLE SPACE TELESCOPE 

Blair D. Savage 
Department of Astronomy 
University of Wisconsin 

475 N. Charter Street 

Madison, WI 53706 

1. INTRODUCTION 

The study of the interstellar medium (ISM) is concerned with answering such questions as: 
l.What is the three dimensional distribution of gaseous matter in the galaxy ? 2.What is the 
composition of that gas ? 3. What is the relationship between interstellar gas and dust ? 4. 
What are the physical conditions in the different phases of the interstellar medium? 5. What 
physical processes control these conditions ? 6. How does the ISM participate in basic 
Galactic processes ? 

Satellite spectrographs make it possible to seek answers to these questions using the very 
significant diagnostic power of UV spectroscopy. Previous UV spectroscopic observations of 
the ISM were mostly obtained with the Orbiting Astronomical Observatory (OAO) series of 
satellites including OAO-2 and OAO-3 (the Copernicus satellite) and more recently with the 
International Ultraviolet Explorer (lUE) satellite. For a review of the OAO results and in 
particular those of the Copernicus satellite see Spitzer & Jenkins (1975) and Spitzer (1988). 
The lUE results relevant to the ISM are reviewed in a number of papers found in Kondo et 
al. (1987). 

With the launch of HST a new suite of instruments has become available for the study of 
the ISM. These include two imaging cameras (WF/PC and FOC) and two spectrographs (FOS 
and GHRS) . The first HST imaging results are found in Astrophysical Journal Letters (1991 
March 10). The imaging results for SN 1987A (Jakobsen et al. 1991) and the Orion Nebula 
(Hester et al. 1991) reveal that even though HST imagery is significantly impacted by the 
spherical aberration, important imaging science related to the ISM can still be pursued. The 
first HST spectroscopic results are found in Astrophysical Journal Letters (1991 August 20). 

A discussion of the UV spectroscopic capabilities of the HST for studies of the ISM is the 
emphasis of this short review which is organized as follows: A brief overview of the 
diagnostic power of UV spectroscopy is found in § 2 while the spectroscopic capabilities of 
the HST are discussed in §3. Examples of recent scientific results are found in §4. Because 
of the author's affiliation with the GHRS science team and his access to GHRS data much of 
the emphasis will be on results from the GHRS. 

2. THE DIAGNOSTIC POWER OF ULTRAVIOLET SPECTROSCOPY FOR 
STUDIES OF THE ISM 

The spectrographic instrumentation on HST provides access to the rich UV region of the 
spectrum for absorption and emission line spectroscopy. The UV spectrographs aboard HST 
have high efficiency from 11 50 to 3200 A and limited capability from 11 00 to 1 150 A. Table 1 
lists some of the atomic and molecular species having their resonance lines or lines from low 
lying excited levels in the 1 100 to 3200 A region of the spectrum. Access to the UV makes 
possible the direct detection of absorption by such abundant atoms as C, N , O, Mg, Si, etc. in 
a number of ionization states including those found in cool neutral gas ( C I, C II, N I, O I, etc) 
and those found in the hot interstellar medium ( C IV and N V). In a number of cases, ions 
from adjacent ionization states are available which can be used to probe physical conditions in 

33 



interstellar clouds (e.g C I, C II, Mg I, Mg H, S I, S II, S III, P I, P II, P III, etc.). 

The UV provides important tracers of cold neutral interstellar gas and permits studies of 
elemental abundances in that gas. With high spectral resolution and high signal to noise ratios 
it will be possible to begin to get information on many of the rarer elements such as those listed 
near the bottom of Table 1. In addition many molecules that play important roles in interstellar 
chemical reaction networks become available for direct study. The extension of some 
sensitivity to 11 00 A due to the selection of a LiF window for the short wavelength detector 
for the GHRS makes possible studies of the (0,0) and (1,0) Lyman band systems of H2 and 
HD toward bright stars. 

The UV provides important tracers of the hot ISM with the doublet of N V near 1240 A 
being the most important among those accessible to HST. N V peaks in abundance in 
collisionally ionized gas with a temperature near 2x10^ K. Unfortunately the important O VI 
doublet near 1032 and 1038 A is not observable with HST except in the spectra of QSO's with 
adequate redshift to shift the doublet to wavelengths longward of 1 150 A. 



TABLE 1 

ATOMS AND MOLECULES WITH UV RESONANCE LINES 

AND ATOMS WITH LINES FROM LOW LEVEL FINE 

STRUCTURE STATES WITH 1100 < X < 3200 A 



HI.DI 


H2,HD 


B I, B II , B III 




CI, CI*, C I**, C II. C II*, CIV 


CO,C2.CO+,CH2 


NI,NV 


N2, CN+. NO. NO+ 


01, 01* ,01** 


H2O 


Nal 




Mg I, Mg II 


MgH+ 


Al I, Al n, Al III 




SiI,SiII,SiII*,Sini,SiIV 


SiO 


pi.pn.pin.pv 




s I, s n, s m 


cs 


ClI 


HCl 


Sc II. Sc III 




Ti II, Ti III 




V I, V II, V III 




Cr I, Cr n 




Mn I, Mn II 




Fe I, Fen. Fell*. Felll 




Co I. Co II 




Nil.NiU 




CuI.CuII 




Zn I, Zn II 




Gal, Gall 




Gel.Gell 




As I. As II 




Se I, Se n 




Krl 





In addition to the direct detection of a wide range of atomic and molecular species for 
abundance studies, UV spectroscopy also permits measures of the physical conditions in 
interstellar clouds. Temperature information is available from the measured Doppler 
broadening of spectral lines or from the presence of certain species (i.e. N V suggests gas near 
IxlO-'K). Interstellar gas density can be inferted by studying the excitation of fine structure 
levels or from measures of abundances of gas in adjacent states of ionization (i.e. Mg I/Mg II, 

34 



C I, C II, S I/S II, etc). In several cases, atomic isotope shifts are large enough to allow the 
direct study of important isotope ratios (i.e. D/H). In other cases isotope ratios can be inferred 
for atoms found in molecules (i.e. C^^O/ C^^o, and HD/ H2) . 

Figure 1 provides an example of the spectroscopic richness of the UV. The spectrum is of 
the bright rapidly rotating star Zeta Ophiuchi and was obtained with the GHRS Echelle A mode 
using the small science aperture. The spectral resolution is approximately 3.5 km s"^. In the 6 
A region of the spectrum illustrated, interstellar absorption lines are found for Kr I 
X1235.83, Ge II ?il237.059, Mg II A.1239.925 and Mg II X1240.395. Kr I and Ge II 
have not been previously seen in the interstellar gas. 



20 



15 _ 



10 



0.0 




1235 



1236 



1237 



1238 



1239 



1240 



Figure 1. GHRS Echelle mode measurements of the far- UV spectrum of Zeta Ophiuchi 
from 1235 A to 1241A revealing interstellar absorption by Kr I, Ge II, and Mg II. The 
spectrum, obtained with the GHRS and the small (0.25"x0.25") entrance aperture, has a 
resolution of 3.5 km s'^ . The Mg II doublet lines reveal a two component structure with the 
strongest component at \(helio) = -75 km s'^ and a weaker component at -27 km s'K The 
signal to noise ratio in the spectrum is approximately 40/1. 

3. HST SPECTROSCOPIC CAPABILITIES 

Three HST instruments have spectroscopic modes that can be used for UV interstellar 
studies. The characteristics of these modes are listed in Table 2. The modes in the FOG and 
FOS can be used for low resolution (FWHM > 200 km s'^) measurements of interstellar 
emission and or absorption lines and for the study of the continuous extinction due to 
interstellar dust. The FOS also has a spectropolarimetric mode that will permit important 
studies of the UV polarizing properties of interstellar and circumstellar dust . 

It is expected that the primary instrument for studies of the interstellar gas will be the 
Goddard High Resolution Spectrograph (GHRS) which has modes providing low, medium 
and high spectral resolution operating in the UV from 1 100 to 32(X) A. 

The spherical aberration of the HST has adversely affected the spectrographs in several 
ways. To preserve spectroscopic resolution the small entrance apertures must be used which 
results in a substantial reduction in throughput. For example with the GHRS the estimated 
light loss when using the 0.25"x0.25" small science aperture (SSA) is approximately a factor 
of 3 for X >1700A and a factor of 4 for X. < 1700A. The light loss is greater if the object is 
poorly centered on the aperture. The 2"x2" large science aperture (LSA) accepts more 
radiation but the resultant spectrum is degraded in resolution by about a factor of 2 to 3, 
depending on the mode. 



35 



TABLE 2 
HST ULTRAVIOLET SPECTROSCOPIC CAPABILITIES 



mode 



wavelength 
range (A) 



resolution ^ 
-FWHM(kms-l) 



comment 



Goddard High ResoluUon SpectropraphfGHRS") 


LSA 


SSA 










f2"x2"> 


(,25"> 


^25") 




G140L 


1100-1800 


250kms-l 


130 kms-l 


Side lb 


G140M 


1100-1800 


25 


12 




Side l'' 


Echelle A 


1150-1800 


8 


3.5 




Side l'' 


G160M 


1150-2100 


30 


15 




Side 2 


G200M 


1600-2300 


25 


13 




Side 2 


G270M 


1900-3200 


25 


12 




Side 2 


EcheUe B 


1700-3200 


8 


3.5 




Side 2 


Faint Obiect Camera fFOC> 




0,1" X 20" slit 






long slit spectrographic facility 


3600-5400 


260 km s-1 






first order 




1800-2700 


260 






second order 




1200-1800 


260 






third order 




1150-1350 


260 






fourth order 


Faint Obiect Spectroffranh (POS^ ' 


c 


0,25" X 2,0" 


^d 






G130H 


1150-1608 


210 km s-1 




G130H has low sensitivity 


G190H 


1576-2332 


250 








G270H 


2227-3306 


250 








G160L 


1150-2523 


1200 









* The approximate resolution (FWHM [km s"^]) is given for the various spectroscopic modes. The resolution 

depends on the aperture choice. When using the larger apertures a single number is a poor measure of the 

resolution. The spectroscopic spread function has a sharp core and very broad wings with a significant fraction of 

the energy going into the wings. 

^ At the time of writing (6 September 1991) side 1 of the GHRS was experiencing electrical problems with a 

low voltage power supply. 

'' The FOS has many other modes which operate at visual wavelengths. It also has a spectropolarimetric mode 

that can be used to study interstellar polarization in the ultraviolet. 

^ The FOS resolution is only listed for the 0.25" x 2.0" slit. A large selection of other aperture choices is 

available. However, the resolution will be degraded when using larger apertures. 



When using the SSA, the GHRS intermediate resolution modes have a resolution 
comparable to the high resolution mode of the Copernicus satellite which is about a factor of 2 
higher than the high resolution mode of the lUE observatory. The GHRS high resolution 
Echelle modes have a resolution about 3 times higher than the Copernicus satellite and 6 times 
higher than lUE. Figure 2 Illustrates absorption line measurements of the 2803.531 A Mg II 
line toward Mu Columbae obtained with the GHRS operating in several modes. Although the 
LSA Echelle-B spectrum appears to have somewhat higher resolution than the G270M -SSA 
measurement, the G270M measurement provides more wavelength coverage and a stable 
spectroscopic spread function. The appearance of the Echelle-B LSA measurement will vary 
with telescope pointing accuracy and with telescope focus. 

Figure 3 compares LSA and SSA measurements of interstellar Fe II absorption obtained 
with the G270M grating. The degradation in resolution when using the LSA is substantial. 
Broad wings are apparent on the interstellar lines for the LSA measurements. 

36 



Figure 2. GHRS observations of 
theMgll X2803.531 interstellar 
absorption line toward fj. Columbae 
in the G270M mode (with the 
SSA), in the Echelle-B mode (with 
the LSA), and in the Echelle-B 
mode (with the SSA). The 
multicomponent nature of the 
absorption which is readily apparent 
in the Echelle-B-SSA data is less 
apparent in the G270M-SSA 
measurement. The effective 
resolutions (FWHM) from top to 
bottom are approximately 12, 8 and 
3.5 km s'^ . 



fx Col Mg II 2803 



•o 

01 



o 

2 




ECH-B LSA 



iAX^V^WHA. 



0_ 




_L 



-100 



-50 



50 

Velocity (km/s) 



100 



150 




2371 



2375 



2379 



2383.0 



2387.0 



2391.0 




2375.0 



2379.0 



2383.0 



2387.0 



)1.0 



Figure 3. GHRS observations of interstellar Fe II X 2374.461 and X 2382.765 absorption 
toward HD 93521 as observed in the G270M mode with the large science aperture, LSA (top ) 
and small science aperture, SSA (bottom). The degradation of spectral resolution when using 
the LSA at intermediate resolution is readily apparent. The effective resolution (FWHM) for the 
upper spectrum is 25 km s'^ compared to 12 km s'^ for the lower spectrum. However, the 
very broad wings on the spectroscopic spread function for the LSA measurement creates the 
broad wings on the resulting interstellar line profiles. 



37 



The principal advantages the GHRS offers over previous instruments for UV interstellar 
studies are: 1. Higher spectral resolution. The 3.5 km s'^ resolution will permit the study of 
conditions in individual interstellar clouds. 2. Low noise photon counting detector . The 
Digicon detectors are capable of very high S/N spectroscopy (S/N = 160/1 has been 
demonstrated). The detectors can be used to observe objects well beyond the reach of lUE 
which is severely background noise limited. The combination of high resolution and high 
signal to noise permits the study of very weak interstellar features which are important for 
accurate abundance measurements. 

At the time of writing, side 1 of the GHRS was experiencing intermittent electrical 
problems with a low voltage power supply. Side 1 contains the G140L, G140M and EcheUe A 
modes while side 2 contains G160M, G200M, G270M, and Echelle B modes. If side 1 is lost 
for future use, the GHRS will still have wavelength coverage down to 1150 A with the 
G160M mode. However, the resolution is substantially reduced compared to that provided by 
Echelle A (15 versus 3.5 km s'^). Echelle B on side 2 would provide a high resolution 
capability in the middle UV from 1700 to 3200 A . 

4. SCIENTIFIC RESULTS 

This section overviews the principal results presented in several papers on interstellar 
absorption found in the special Astrophysical Journal Letters issue reporting the first HST 
spectroscopic results. Those papers involving the ISM concern the interstellar gas toward t, 
Persei, the gaseous matter in the circumstellar environment of p Pictoris, and Milky Way 
absorption in the direction of the bright QSO 3C273. The results for (3 Pictoris are discussed 
elsewhere in this volume. 

Several more recent HST results also pertain to the ISM. GHRS measurements of 
interstellar absorption by D and H toward Capella are discussed by Linsky( this volume). New 
GHRS measurements of weak absorption lines in the spectrum of C, Ophiuchi from Cardelli, 
Savage & Ebbets (1992) are briefly described below. 

4.1 Interstellar gas toward ^ Persei 

The star t, Persei (HD 24912) is an 07.5 III star at an estimated spectroscopic distance of 
540 pc in the direction 1 = 160.4O and b = -13.1° The large vsini ( 216 km s'^) and 
substantial reddening E(B-V) = 0.32 make ^ Persei well suited for studies of narrow 
interstellar absorption lines. Ultraviolet observations of interstellar absorption toward t, Persei 
were obtained by the GHRS in October 1990 and January 1991. The measurements were 
obtained with the GHRS operating in the Echelle modes with the light of ^ Persei placed in the 
small (0.25"x0.25") entrance aperture. Observations were obtained at 26 different setup 
wavelengths with each observation providing approximately 5 to 10 A coverage of the 
ultraviolet spectrum. The total integration time for each spectral region observed was typically 
3 to 6 minutes. The signal to noise achieved in the reduced data ranges from 15 to 100 
depending on the wavelength and the accuracy of centering the star in the small aperture. The 
results of the analysis of the GHRS data for b, Persei are found in Cardelli et al. (1991), Savage 
et al. (1991) and Smith et al (1991). Cardelli et al. present gas phase abundance results for the 
diffuse clouds toward the star. Savage et al. discuss how elemental abundances vary with 
velocity, and Smith et al. consider physical conditions in the densest portions of the cloud 
toward t, Persei where molecules are found. The data were processed using the techniques 
described by Cardelli et al. (1991). Sample interstellar line profiles plotted on a velocity basis 
are illustrated in Figure 4. 

The GHRS setup wavelengths provided for the detection of interstellar absorption lines 
from the following ions: C I, C I*,C I**, C II, C II*, C IV, O I, O I*, Mg I, Mg II, Al III, Si 
II, Si IV, P II, S I, S II, S III, CI I, Cr II, Mn II, Fe II, Ni II, Cu II, Zn II, and the CO 
molecule. The velocity structure of the absorption toward t, Persei is as follows: 1. The 
diffuse clouds which are traced by the weaker absorption lines of neutral and once ionized 
atoms absorb between +5 and +15 km s'^. 2. Ionized gas absorption is traced by the GHRS 
measures of Al III, S IE, Si IV and C IV. The Al III and S III lines indicate broad absorption 

38 



extending from -20 to +20 km s"l. The lines of Si IV and C IV also absorb in this velocity 
range although their profile shapes are considerably different than those for Al III and S HI. 
N V >,1238 absorption was not detected in the low S/N spectrum obtained by the GHRS at 
1238 A. 3. The strong ultraviolet lines of species such as Fe II, Si II, O I, and C II reveal 
additional absorption at velocities different from those of the diffuse clouds. In particular, the 
lines of Fe 11 and Si 11 reveal components near -5 km s'^ and +25 km s"l. These strong UV 
lines permit the detection of low column density gas with gas phase abundances very different 
than those found in the diffuse clouds toward q Persei. 

Accurate measures of total column densities for a large number of atomic species have been 
determined from the ^ Persei absorption line measurements (Cardelli et al. 1991). A summary 
of the results is shown in Figure 5 which illustrates gas phase abundances relative to solar 
versus the condensation temperature for each element. The GHRS allows the measurement of 
very accurate column densities of not only abundant elements such as C, N, and O but also for 
many of the rarer elements such as Cr, Ni, Cu, and Zn. 




-60 -40 -20 0.0 20 « 60 -60 -40 -20 0.0 20 W 60 -60 -40 -20 0.0 20 40 60 
Velocity (km s'') Velocity (km s"') Velocity (km s"') 



Figure 4. Relative intensity versus heliocentric velocity for a sample of interstellar lines in 
the spectrum of t, Persei. The zero-level of each line is indicated on the vertical axis in each 
panel. For Si II and Fe II the availability of absorption lines with a wide range of oscillator 
strength permits an evaluation of the nature of the absorption over a wide range of velocity. 
The low column density absorption components near -5 and +25 km s'^ are readily apparent in 
the sequence of Fe II XX 2374.46, 2382.77, 2600.17 profiles illustrated. In several cases 
more than one absorbing species appears on the short portion of the ultraviolet spectrum 
illustrated. The measurements are from the GHRS Echelle mode with t, Persei in the small 
aperture and have a resolution of approximately 35 km s'K 



39 



o 



o> 
o 








1 '''' 1 
wsak ltd* 


-^^ — r 




''-' 


1 


1 


1 


1 


1 






! 


damping H damping 
wings _ wings 


















-1 




c 


■ u 

weak' 
lint 






Zn 




! 




waok 2 

lin«s - 

damping j=j 
wings U 


P 




_ 




- 










Cu 


■ 
Vn 


Mg 






- 


-2 


- 


0"(N/H)- j}Cr(N)onl» 

1 1 ^^ 1 


-VX— L 




1 ,'/' 


-1 ^^ 


1 


1 


Cr 

1 


Fe 

i 
1 
Ni 

1 


- 



50 100 200 650 700 1100 1200 1250 

Condensation Temperature (K) 



1300 1350 



Figure 5. Average depletion for dominant ions arising in the diffuse clouds toward ^ Persei 
plotted against the element's condensation temperature . The filled symbols illustrate 
depletions and their errors derived from weak absorption lines while the open symbols 
illustrate results from strong lines with Lorentzian damping wings. The extra error bars 
include the additional uncertainty associated with the errors in the reference hydrogen column 
density from Bohlin, Savage, & Drake (1978). The discrepancy between the open and filled 
symbols, especially for Mg II indicates potential errors in the adopted f values. This figure is 
from Cardelli et al. (1991) 



The 3.5 km s'^ resolution provided by the Echelle mode of the GHRS means that in many 
cases the absoqjtion line measurements obtained by the spectrograph will be either fully 
resolved or close to being fully resolved. When an absorption line is fully resolved the 
observed optical depth of absorption is given by 

T(v) = In [ Io(v) / lobs(v) ] = 2L£2 a N(v) = 2.654x10-15 f X N(v), 

mgC 

where, IqCv) and lobs(v) ^^ *^ continuum and observed intensity, respectively, N(v) is the 
column density per unit velocity in atoms cm'^ ( km s-l)-^, f is the oscillator strength of the 
line, and A. is the wavelength in A. . When an absorption line is not fully resolved we refer 
to the optical depth as derived above as the apparent optical depth, ta(v), and the 
corresponding column density as the apparent column density, N^Cv). 

The high resolution absorption line data can be converted into measures of N^Cv) 
extending over a large range of v by combining the absorption measurements for lines of 
different strength. The weak lines will determine Na(v) at velocities where the column density 
per unit velocity is large while strong lines will determine N^Cv) at velocities where it is small. 
If a given species has a number of lines, it is possible to construct a complete N^Cv) profile 
and use the empirical information in the region of velocity overlap from one line to the next to 
infer the presence or absence of unresolved saturated structures in the derived profiles. When 
Na(v) profiles derived firom lines differing in the value of Xf by more than a factor of 2 agree in 
a region where they overlap in velocity, the work of Savage & Sembach (1991) has 
demonstrated that it is reasonable to assume that unresolved saturated structures are not 
influencing the values of N^Cv). Comparisons of curves of Na(v) for different elements 

40 



permits a study of how elemental abundances change with velocity due to the effects of 
differential depletion. Such a comparison is found in Figure 6 from Savage et al. (1991) 
which shows log N^Cv) curves for O I , C II, Mg II, Si II, Fe II, S II, Mn II, and Zn II 
adjusted for solar abundance differences. At velocities where the curves for different elements 
coincide, the elements are found in the gas phase with the solar abundance ratio. At velocities 
where the curves differ one is directly viewing the effects of differential depletion. For 
example, in the velocity range from v = +5 to + 15 km s*^ the curve for Fe EI in Figure 6 lies 
below that for O I by about 1.8 dex due to the incorporation of Fe into interstellar dust. In the 
cloud near v = +25 km s'^, the various curves nearly overlap suggesting nearly solar 
abundance ratios for O I, Si II, Fe II, S II, Mg U. Evidently elements have been returned to 
the gas phase through some process like shock processing of interstellar dust in the medium 
absorbing near +25 km s'^. 

The interstellar Une data for ^ Persei also provide significant diagnostic information about 
the physical conditions in the denser portions of the diffuse clouds along the sight line. Smith 
et al. (1991) have analyzed the data for C I, C I*, C I** , and CO to estimate the temperature, 
density, and pressure in the clouds. The highest density component near 6 km s'^ is estimated 
to have a very high pressure [ log (P/k) > 4.3] while the component at 10 km s"l has a 
pressure a factor of 10 lower. Such estimates are of great importance for detailed modeling of 
atomic and molecular conditions in diffuse clouds. 



Figure 6. Curves of apparent 
column density versus velocity 
adjusted for solar abundances for the 
sight line to ^ Persei for O I,C II, Mg 
II, Si II, Fe II, S II, Mn II , and Zn 
II. The value of [OIXJq for a given 
element was taken to be the 
logarithmic abundance difference for 
that element compared to oxygen and 
was obtained from the solar system 
abundances of Morton (1991). Thus, 
all the curves are corrected for solar 
abundance differences and are 
referenced to oxygen, a species that is 
only slightly depleted in diffuse 
interstellar clouds. At velocities 
where the curves for two species 
overlap, the two species therefore 
have solar abundance ratios. The 
large vertical separation of some of 
the curves at velocities corresponding 
to those of the diffuse clouds ( v ~ +5 
to +15 km s'^) is caused by the 
depletion of the various elements into 
interstellar dust. Note that Lorentzian 
wings appear on the profiles for 
species with very strong lines (i.e. C 
II, O I and Mg II) . This figure is 
from Savage et al. (1991). 



X 

\ 

o 

I — I 

+ 

> 







16 



15 



13 



12 



-^ \ \ 

«»«». I (ref) 
00000 C II (0.37) 
^,,,+ Mg II (1.34) 
«e«, Sll (1.66) 
Doooo Mn II (3.40) 
^,,^, Zn II (4.28) 
.... Si II (1.38) 
Fe II (1.42) 



I — r 



1 — r 



11 




-40 -50 -20 -10 0.0 10 20 30 40 50 50 
Velocify (km s"') 



41 



Figure 7. Examples of weak 
absorption lines seen in Echelle mode 
observations of C, Ophiuchi. The 
species and wavelengths are listed on 
the left while^ the observed equivalent 
widths in mA are listed on the right. 
The absorption for Zn 11 illustrates 
the two component structure seen in 
strong lines. The principal 
component is at -15 km s'^ with a 
weaker component near -27 km s'^ . 
The weak lines are only reliably 
detected in the -15 km s' ^ 
component. C 11] X2325 is not 
detected with a 2a upper limit of 
0.8mA in a spectrum with SIN = 
15011. The detection of Kr 1 
absorption is particularly significant 
since Kr is not expected to form 
bonds and will probably not be 
incorporated into the interstellar dust. 
Hence, measures of Kr provide direct 
abundance information about an 
important element created by a 
combination of r- and s-process 
nucleosynthesis. This figure is from 
Cardelli, Savage, & Ebbets (1992). 



1 1 r- 



Cll X2325 



-No 1 X2854 



-1 1 i 1 — 

Wx(mA)- 

<0.8 



7.49 




-40 -20 20 

Vo (kms-') 



4.2 Weak Absorption Lines Toward C, Ophiuchi 

In May 1991 the bright O star C, Ophiuchi was observed with the Echelle modes in order to 
evaluate the scattered light in the spectrograph. Those science verification data are also 
providing important information about the ability of GHRS to detect weak absorption lines. 
The first scientific results from the analysis of the ^ Ophiuchi data is found in Cardelli, 
Savage & Ebbets (1992). Figure 7 provides a sample of the C, Ophiuchi weak line 
measurements for lines of C II], Na I, O I, Cu II, Kr I, Ga II, and Ge II. In addition a strong 
line of Zn II is illustrated. The Zn II absorption reveals the two component absorption toward 
C, Ophiuchi known from optical measurements with the principal absorption occurring near -15 
km s"^ and weaker absorption near -27 km s'^. The various weak lines only record 
absorption in the higher column density component near -15 km s'^ except for C II] which is 



42 



not detected with a small (2o) upper limit of 0.8 mA. Ga, Ge and Kr represent the heaviest 
elements so far detected in the ISM. The derived column densities show that Ga is depleted by 
1.2 dex from the gas phase while Ge is overabundant by 0.2 dex. The detection of Kr is 
particularly important since it is not expected to form chemical or mechanical bonds and should 
reside primarily in the gas phase. The 14.0 eV ionization potential of Kr I means that the 
neutral atom is the dominant ionization state in the neutral ISM. Assuming Kr to be 
undepleted, the measurements imply a logarithmic cosmic abundance of 2.95 on the scale 
where [H] = 12.00. Since Kr is produced by a mixture of s- and r-process nucleosynthesis, 
further observations in other sites will provide significant insights about nucleosynthesis and 
interstellar enrichment processes without the difficulty of understanding the complications 
produced by the presence of interstellar dust. 



4.3 Absorption by the Galactic Corona Toward 3C273 

Spectra of the bright QSO 3C273 were obtained by FOS (Bahcall et al. 1991) and by 
GHRS (Morris et al. 1991). Those spectra provided the important cosmological result that 
the Lyman alpha forest persists to zero redshift and that there is a significant decrease in the rate 
which the Lyman alpha forest thins out, occurring at some redshift less than about 2. The HST 
UV spectra of 3C273 also contain significant information about the gaseous halo or corona of 
the Milky Way. Lines associated with gas in the Milky Way disk and halo from H I, O I, C H, 
C II*, C IV, N I, N V, Si II, Si III, Si IV, S II, Mg II, Mn II, Fe II and Ni II are seen. 
Figure 8 shows the far-UV portion of the FOS spectrum obtained with the G130H grating 
(resolution ~ 250 km s"^) and a small oortion of the GHRS G160M spectrum obtained with the 
large aperture (resolution ~ 25 km s"^ but with wings extending to ± 70 km s"^). The Milky 
Way lines seen in this 35 A portion of the spectrum include: N V X^ 1238.82 and 1242.80, 
S II U1250.58, 1253.81, and 1259.52 and Si II ^1260.42. The other lines, attributed to 
the Lyman alpha forest, appear at ?l?i1242.17, 1247.54, 1251.46, and 1255.70 A. The 
Lyman alpha line at 1251.46 is particularly strong with an equivalent width of 120 mA. A 
comparison between the FOS spectrum and the GHRS spectrum shows the high sensitivity of 
the GHRS to the detection of weak absorption lines. 

The lines from abundant elements associated with the neutral ISM of the Milky Way are 
very strong and have widths (FWHM) extending from 1 10 km s'^ (O I) to 145 km s'^ (C II 
and Si II). These lines are tracing low column density high velocity dispersion neutral and 
weakly ionized gas toward 3C273 ( 1 = 290° and b = 65°). The lines from the highly ionized 
gas (Si IV , C IV, and N V) are also very strong and broad. The direct detection of N V is 
quite significant because it traces collisionally ionized gas near 2x10^ K. Gas at this 
temperature cools very rapidly. Its origin may be associated with the cooling gas of a Galactic 
fountain (Shapiro & Field 1976; Edgar & Chevalier 1986) or with thermal condensations in 
cosmic ray driven fountains (Hartquist & MorfiU 1986). 

The very strong Milky Way absorption toward 3C273 reveal a problem all extragalactic UV 
spectroscopic observers will face— that of allowing for the presence of Milky Way disk and 
halo absorption lines in the spectra of extragalactic objects. 



43 



1200 



1300 



1400 



1500 




T 1 r 



pttj|*%%(tll% 



1250 
Wavelength (A 



1270 



Figure 8. FOS spectrum of 3C273 from Bahcall et al. (1991) and GHRS spectrum from 
Morris et al. (1991). The FOS spectrum (top) was obtained with the G130H grating using the 
0.25"x2.0" slit and has a resolution of about 210 km s'^ . The GHRS spectrum (bottom) was 
obtained with grating G160M using the 2.0"x2.0" aperture and has a resolution of 30 km s'^ 
(FWHM) but with broad wings extending ±70 km s'^ . The lower curve in the bottom panel is 
the S/N in the spectrum shown as the middle curve. The upper curve in the lower panel 
illustrates a deconvolution of the original spectrum as described in Morris et al. Of particular 
significance for the Milky Way is presence of strong N V doublet absorption at XX1238.80 
and 1242.80 A. This absorption may be produced by cooling 2x10^ K gas in the Milky Way 



corona. 



Acknowledgements The efforts of many people have contributed to the success of the 
early HST spectroscopic program to study the interstellar medium with the FOS and GHRS. 
The author appreciates support from NASA through contract NAS5-29638. 



44 



References 

Bahcall, J.N., Jannuzi.B-T., Schneider, D.P., Hartig, G.F. , Bohlin, R., & Junkkarinen, 

V. 1991, ApJ, 377, L5 
Bohlin, R.C., Savage, B.D., & Drake, J.F. 1978, ApJ, 224, 132 
Cardelli, J.A., Savage, B.D., Bruhweiler, F.C., Smith, A.M., Ebbets, D.C. Sembach, K.R., 

& Sofia, U.J. 1991, ApJ, 377,L57 
Cardelli, J., Savage, B.D., & Ebbets, D.E. 1992, ApJ, (submitted) 
Edgar, R.J, & Chevalier, R.A. 1986, ApJ, 310,L27 
Hartquist, T.W., & Morfill, G.E. 1986, ApJ, 311, 518 
Hester, J.J. et al. 1991, ApJ, 369, L75 
Jakobsen , P. et al.l991, ApJ, 369, L63 
Kondo et al. (eds), 1987, Exploring the Universe with the lUE Satellite, (D.Reidel Pub. 

Co.: Dordrecht) 
Morris, S.L., Weymann, R.J., Savage, B.D., & Gilliland, R.L. 1991, ApJ, 377, L21 
Morton, D.C. 1991, ApJS, (in press) 
Savage, B.D., Cardelli, J.A., Bruhweiler, F.C. , Smith, A.M., Ebbets, D.C. & Sembach, 

K.R. 1991, ApJ, 377, L53 
Savage, B.D., & Sembach, K.R. 1991, ApJ, ( in press) 
Shapiro, P.R., & Field, G.B. 1976, ApJ, 205,762 
Smith, A. M., Bruhweiler, F. C, Lambert, D. L., Savage, B. D., Cardelli, J. A., 

Ebbets, D. C, Lyu, C.-H., & Sheffer, Y. 1991, submitted to ApJL 
Spitzer, L. 1988, PASP, 100, 518 
Spitzer, L. & Jenkins, E.B. 1975, ARAA, 13, 133 



45 



FOS Observations of the Absorption Spectrum of 3C 273 ^ 

J.N. Bahcall^, B.T. Jannuzi^, D.P. Schneider, 

Institute for Advanced Study 

School of Natural Sciences, Princeton, NJ 08540 

G.F. Hartig, R. Bohlin 

Space Telescope Science Institute 

3700 San Martin Drive, Baltimore, MD 21218 

V. Junkkarinen 

University of California, San Diego 

La JoUa, CA 92093 

Abstract 

We describe the FOS observations of the absorption line spectrum of 3C 273 and 
compare the results with the GHRS observations of the same quasar. Three Ly-a lines 
appear to be produced by gas in the Virgo cluster or by the halos of galaxies associated 
with the Virgo cluster. We identify a total of seven Ly-a absorption systems with equiv- 
alent widths greater than 0.2 A. The evolution of the number density of Ly-a clouds 
carmot be determined with confidence by comparing only 3C 273 observations with 
those of large redshift quasars. The inferred redshift- dependence of the number den- 
sity depends critically upon whether or not the Virgo-cluster lines cire included or upon 
an uncertain extrapolation of the equivalent width distribution for strong Unes found 
at large redshifts. 

1 INTRODUCTION 

We report here on ultraviolet observations of 3C 273 made with the high resolution 
gratings of the Faint Object Spectrograph (FOS, see Ford 1985) as part of the science 
verification program of the Hubble Space Telescope. In §11, we summarize the observations 
and in §111 we describe the measurement of the Hnes. In §IV, we discuss the identification of 
Ly-a lines. In §V we compare our results with those obtained by Morris et al. (1991) with 
the GHRS and in §VI we discuss the redshift evolution of the Ly-a systems. We summarize 
our main scientific results in §VII. 



^ Based on observations with the NASA/ESA Hubble Space Telescope, obtained at the Space Telescope Science 
Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA 
contract NAS5-26555. 

^Guest Observer with the International Ultraviolet Explorer sa.ie\lite, which is sponsored and operated by the 
National Aeronautics and Space Administration, by the Science Research Council of the United Kingdom, 
and by the European Space Agency. 



46 



2 OBSERVATIONS 

We observed 3C 273 on 14-16 January, 1991 using the three high resolution {R = 1300) 
gratings, G130H, G190H, and G270H, centered, respectively, on 1300 A, 1900 A, and 2700 A. 
The spectra cover the region between 1150 A and 3300 A with a gap from 1600 to 1650 A. 
Five apertures were used with each grating. They are the 0.25" x 2.0" slit, the three circular 
apertures with diameters of 1.0", 0.5", and 0.3", and the 4.3" by 4.3" square aperture. A 
paper currently in preparation (Jannuzi, Hartig, Bahcall, and Schneider 1992) wiU compare 
and analyze the data from all the apertures. In this talk, we use the data from different 
apertures to verify results obtained with the sUt or to reject spurious features. The spectral 
resolutions (FWHMs) of the data obtained through the sUt are 1.1, 1.5, and 2.0 A for the 
three gratings. We also observed 3C 273 with the International Ultraviolet jEip/orer satellite 
on 7, 13, 15-17, and 23 January 1991 in order to set the zero point of the flux calibrations 
and to check upon possible flux variability. 

Figure 1 of Bahcall et al. 1991 (hereafter, Paper I) shows the reduced HST slit data from 
all three gratings. The typical signal-to-noise ratio (SNR) is between 40 and 50 per diode 
in the G130H data (2000 s exposure), « 60 per diode in the G190H data (1400 s exposure), 
and ss 100 per diode in the best-studied regions of the G270H data (1400 s exposure). The 
signal-to-noise ratios for the data are better than that expected to be obtained for the Key 
Project Quasar Absorption Line Survey, for which the signal-to-noise ratio will typically be 
30 per diode in each of the gratings. Further details of the observations are given in Paper I. 

We were fortunate in having observations through the five different apertures, which 
enabled us to reject a number of spurious features. For example, we rejected narrow features 
that had similar strengths and shapes when observed through the sht and the 4.3" aperture, 
since real lines have increased hne widths (due to the degraded HST PSF) in the larger 
aperture. 

3 MEASUREMENT OF ABSORPTION LINES 

Details of the measurement of the Hues are given in Paper I. For inclusion in the complete 
sample, we required that all lines be 3(7 detections with observed EWs greater than 0.250 A 
in the G130H or G190H data or greater than 0.150 A in the G270H data and that the 
measured equivalent widths be consistent in the the slit and the 0.3" data. 

Table 1 hsts the complete sample of 36 absorption lines. The columns contain: the 
measured line center, equivalent width, error in equivalent width, identification, vacuum 
wavelength of identification, difference between observed and laboratory wavelength, and 
comments. 



47 



TABLE 1: Ultraviolet Absorption Lines 



-^obo 


EW 


cew 


Identification 


AA 


Comment 


(A) 


(A) 


(A) 


Ion 


Ao 


(A) 




1190.44 


0.463 


0.111 


Sill 


1190.42 


0.02 




1193.35 


0.414 


0.102 


Sill 


1193.28 


0.07 




1200.22 


0.982 


0.085 


N I 


1199.90 


0.32 




1206.44 


0.553 


0.078 


Si III 


1206.51 


-0.07 




1215.7: 


7.0: 


1.5: 


Ly Q 


1215.67 


0.06 


EW uncertain: 


1219.80 


0.371 


0.039 


Ly a 






cz = 1020 km s-i 


1222.12 


0.414 


0.092 


Ly a 






cz = 1590 km s"^ 


1224.52 


0.240 


0.081 


Ly a 






cz = 2185 km s~^; not in cpl. sp. 


1238.60 


0.183 


0.076 


N V 


1238.81 


-0.21 


Not in cpl. sp. 


1243.04 


0.178 


0.101 


N V 


1242.80 


0.24 


Not in cpl. sp. 


1260.08 


0.789 


0.057 


Sill 


1260.42 


-0.34 


+ SII (1259.53) ? 


1275.23 


0.251 


0.059 


Ly a 






cz = 14,700 kms-i 


1296.52 


0.287 


0.057 


Ly Q 






cz = 19,950 kms-^ 


1302.08 


0.372 


0.050 


01 


1302.17 


-0.09 




1304.40 


0.395 


0.060 


Sill 


1304.37 


0.03 




1317.08 


0.292 


0.096 


Ly a 






cz = 25,030 km 8-^ 


1325.10 


0.238 


0.057 


Ly a 






cz = 27, 000 km s~^; not in cpl. sp. 


1334.57 


0.586 


0.056 


CII 


1334.53 


0.04 




1335.75 


0.168 


0.058 


CII* 


1335.70 


0.05 


Not in cpl. sp. 


1361.53 


0.146 


0.072 


Ly Q 






cz = 35, 990 km s"^; not in cpl. sp. 


1393.86 


0.479 


0.036 


Si IV 


1393.76 


0.10 




1402.69 


0.261 


0.042 


Si IV 


1402.77 


-0.08 




1526.77 


0.477 


0.047 


Sill 


1526.72 


0.05 




1548.26 


0.561 


0.049 


CIV 


1548.20 


0.06 




1550.75 


0.402 


0.050 


CIV 


1550.77 


-0.02 




1670.92 


0.534 


0.036 


AlII 


1670.81 


0.11 




1855.63 


0.281 


0.053 


Aim 


1854.72 


0.91 




1862.95 


0.182 


0.049 


Aim 


1862.78 


0.17 


Not in cpl. sp. 


1878.03 


0.259 


0.068 








Broad; Flat Field feature? 


2026.55 


0.266 


0.035 


Mgl 


2026.47 


0.08 




2065.09 


0.440 


0.061 










2344.16 


0.727 


0.031 


Fell 


2344.21 


-0.05 




2352.30 


0.192 


0.044 










2374.57 


0.612 


0.032 


Fell 


2374.46 


0.11 




2382.77 


0.828 


0.031 


Fell 


2382.76 


0.01 




2577.87 


0.344 


0.081 


Mnll 


2576.89 


0.98 




2586.67 


0.844 


0.050 


Fe II 


2586.64 


0.03 




2594.36 


0.283 


0.066 


Mnll 


2594.50 


-0.14 




2600.25 


0.980 


0.051 


Fell 


2600.18 


0.07 




2606.71 


0.155 


0.026 


Mnll 


2606.47 


0.24 




2796.27 


1.098 


0.025 


Mgll 


2796.35 


-0.08 




2803.51 


0.993 


0.024 


Mgll 


2803.53 


-0.02 




2852.85 


0.392 


0.027 


Mgl 


2852.97 


-0.12 





48 



4 LINE IDENTIFICATIONS AND IMPLICATIONS 

For identifications, we used a standard set of ultraviolet absorption lines (Bahcall 1979) 
that correspond to the strongest allowed, one-electron, dipole transitions from ground or 
excited fine-structure states of abundant elements. Of the 36 observed lines in the complete 
sample, 28, or 78%, are identified at zero redshift with strong interstellar lines from the 
standard Ust. The rms wavelength difference between the measured and the 28 standard 
lines is 0.26 A. 

The characteristics of the interstellar lines towards 3C 273 that are inferred from our 
observations are discussed in Paper I (see also Ulrich et al. 1980). 
Ly-a 

absorption at small redshifts apparently produces 5 of the 8 lines in the complete 
sample that are not identified with Galactic absorption. The total range over which we 
have good observations exceeds 1770 A. If the 8 unidentified lines were uniformly sprinkled 
over this entire range, then the Poisson probability is about 0.2% that 5 of 8 lines would be 
confined, as observed, into the 192 A between the rest wavelength of Ly-a and the redshifted 
Ly-a emission line. One of the remaining three lines may be identified as a possible flat 
field artifact (at 1878.03 A). In addition, there are three other real lines that are apparently 
Ly-a absorption systems but which have EWs too small to be in the complete sample. These 
lines, at 1224.52 A (EW = 0.24 A), at 1325.10 A (EW = 0.24 A), and 1361.53 A (EW = 
0.15 A), are included in Table 1 and correspond to values of cz of 2,185 km s~^,27,005 km s~\ 
and 35,995 km s~^. 

The Ly-a systems have H I column densities of 14 ^ log iV ^ 16 for 6 w 35 km s~^. The 
maximum EW that any line could have and not be in the complete sample (0.25 A) limits 
the amount of metals present in the clouds to log N ^ 15. The absence in our spectra of 
metal lines from the Ly-a systems (hereafter, called clouds) is not surprising; we would not 
expect to detect the metal lines unless the hydrogen is exceedingly highly ionized. 

We emphasized in Paper I that the Ly-a lines at 1219.8 A and 1222.1 A are most 
likely caused by gas in the Virgo cluster or in halos of galaxies associated with the Virgo 
cluster. The heliocentric velocities of the two Ly-a systems are 1020 ± 17 km s~^ and and 
1591 ± 45 km s~^. The Virgo cluster has an average velocity of 1158 km s~^ (Huchra 1985) 
and contains galaxies with a broad range of velocities that bracket the two Virgo Ly-a clouds. 

Figure 1 shows the spectrum in the region between 1203 A and 1236 A. The absorption 
line at 1224.52 A (which was not included in the expanded part of the spectrum shown in 
Figure 2 of Paper I) has a measured equivalent width of 0.24 A and is a real feature (ap- 
proximately 3a significance), although it is weaker than the lines at 1219.8 A and 1222.1 A. 
We believe that the line at 1224.52 A is also caused by absorbing gas associated with the 



49 



4 


— Y6 


' 1 ' ' 

- I 

rJV| 
1 ^ 


>% BB 


3 


- 






n V 


2 


_ L 




^ 


— 


1 


- 


v 




V : 




; ^ 


1 1 1 



1210 



1220 



1230 



Figure 1. The region between 1203 A and 1236 A in the slit spectrum. The 
resolution is w 1.1 A. The most prominent feature is the strong Galactic Ly a 
absorption line. The geocoronal Ly a emission line is seen in the center of the 
Galactic absorption hne. The absorption line shortward of Ly-a is produced by 
Si III in the Geilactic interstellar mediimi. The three marked absorption features 
located on the red shoulder of the Galactic Ly a Une axe probably produced by gas 
in the Virgo Cluster or by halos of galaxies associated with the Virgo Cluster. The 
flux units are 10~^^erg s~^ cm~^ A"^. 



50 



Virgo cluster. Huchra (1985) notes that most of the galaxies in the direction of the Virgo 
cluster with velocities less than 3000 km s~^ are cluster members. 

The large gas cloud HI 1225+01 discovered by Giovanelli and Haynes (1989) has a 
systemic velocity of 1275 km s~^ and may well be associated with the Virgo cluster. The 
center of HI 1225+01 is located about 250 kpc x (f(Hl 1225 + 01)/20 Mpc from the Une of 
sight to 3C 273 and could be the source of one or more of the Virgo Ly-a absorption lines. 

5 COMPARISON OF FOS AND GHRS OBSERVATIONS 

Morris et al. (1991) identified absorption features in observations of 3C 273 that were 
taken through the Large Science Aperture of the Goddard High Resolution Spectrograph 
(GHRS). The results of the FOS and the GHRS observations were reduced independently. 

A direct comparison of the FOS and the GHRS line lists is of great interest since it 
provides an objective estimate of the possible systematic errors in the measurements, in the 
data analysis, and in the interpretations. 

Table 2 compares the wavelengths and equivalent widths measured by the group using 
the GHRS (Morris et al. 1991) and by the group using the FOS (Bahcallei al. 1991 ) 
for lines at zero redshift (hues arising from the galactic interstellar medium, ISM). Some 
regions of the spectrum were not observed with the GHRS and some weak lines reported by 
the GHRS group were not listed by the FOS group. There are 24 lines in common in the 
two Usts. The agreement between the measurements for these lines is generally good. The 
root-mean-squared wavelength agreement is 

< (Afos - Aghrs)' >'/' = 0.33A , (1) 

and the root-mean-squared fractional difference between the equivalent widths is 

< {EWros - EWonKsf I EWl„,^, >"^ - 0.34 , (2) 

where ^^H^average IS the average of the FOS and GHRS measured equivalent widths. 

There is no evidence of any systematic displacement of measured line centers between 
the FOS and the GHRS data as a function of observed wavelength. In fact, almost half of 
the root-mean-squared wavelength discrepancy is contribute by only two lines. 

Figure 2 shows that there are seven ISM Unes with GHRS measured equivalent widths 
less than 0.200 A. For all seven of these lines, the equivalent width measured with the 
FOS is larger than the equivalent width measured with the GHRS. Several different factors 
contribute to this skewed distribution. The most obvious explanation is a selection bias: 
GHRS Unes with a measured equivalent width of less than 0.2 A can only appear in the 



51 



TABLE 2: Comparison ISM Absorption Lines 



GHRS A 


FOS A 


GHRS EW 


FOS EW 


Ident. 


1190.5 


1190.44 


0.426 


0.463 


Si II 


1193.5 


1193.35 


0.389 


0.414 


Sill 


1200.4 


1200.22 


0.713 


0.982 


N I 


1206.8 


1206.44 


0.445 


0.553 


Si III 


1215.0 


1215.7 


7.8 


7.0 


H I 


1238.75 


1238.60 


0.145 


0.183 


N V 


1242.72 


1243.04 


0.052 


0.178 


N V 


1250.53 


.... 


0.164 


.... 


SII 


1253.78 





0.183 





SII 


1259.50 


.... 


0.274 




SII 


1260.40 


1260.08 


0.590 


0.789 


Si II 


1302.07 


1302.08 


0.415 


0.372 


01 


1304.29 


1304.40 


0.385 


0.395 


Sill 


1334.49 


1334.57 


0.622 


0.586 


CII 


1335.72 


1335.75 


0.116 


0.168 


CII* 


1370.09 




0.079 




Ni II 


1393.64 


1393.86 


0.367 


0.479 


Si IV 


1402.65 


1402.69 


0.217 


0.261 


Si IV 


1423.04 




0.043 






1526.66 


1526.77 


0.551 


0.477 


Sill 


1534.71 




0.078 






1548.16 


1548.26 


0.530 


0.561 


CIV 


1550.72 


1550.75 


0.388 


0.402 


CIV 


.... 


1670.92 




0.534 


AlII 


1807.98 


.... 


0.118 




Sill 


.... 


1855.63 




0.281 


Aim 


.... 


1862.95 




0.182 


Al III 





2026.55 


.... 


0.266 


Mgl 


.... 


2344.16 




0.727 


Fell 





2374.57 




0.612 


Fell 




2382.77 




0.828 


Fell 


2576.71 


2577.87 


0.197 


0.344 


Mn II 


2586.46 


2586.67 


0.747 


0.844 


Fe II 


2589.98 




0.033 






2594.27 


2594.36 


0.142 


0.283 


Mn II 


2600.01 


2600.25 


0.942 


0.980 


Fell 


2606.28 


2606.71 


0.090 


0.155 


Mn II 


2796.31 


2796.27 


1.167 


1.098 


Mgll 


2803.49 


2803.51 


1.078 


0.993 


Mgll 


.... 


2852.85 


.... 


0.392 


Mgl 



52 



FOS list if the equivalent width measured with the FOS data is greater than the GHRS 
measured equivalent width. This is because the minimum equivalent width for inclusion in 
the complete FOS sample was 0.250 A. In addition, the higher resolution available in the 
deconvolved GHRS spectra permits the resolution of the FOS line at 1260.08 A into two 
lines at line at 1259.50 A and 1260.40 A. Finally, it is possible that the two groups assumed 
different levels for the continuum. 

ISM Equivalent Widths (in Angstroms) 



CO 

Exq 

in 
O 



.c 


1 


1 1 


1 


1 1 


1 


1 1 


1 


1 1 


1 1 1 


/- 




- 


















y *- 


1 


- 














■ 


m/ 


■ 


.8 


- 












■ 


■ y 




- 


.6 


- 






1 


■ 


■ 


/ 






- 


.4 


- 


■ 


■ 


/ 


V 

■ 










- 


.2 
n 


■ 


1 1 


1 


1 1 


1 


1 1 


1 


1 1 


1 1 1 


1 1 r 



.2 .4 .6 .8 

GHRS EWs 



1.2 



Figure 2. The measiured equivalent widths (EWs) of ISM absoprtion Unes from 
the FOS data (BahcaJl ei al. 1991) plotted against the results from the GHRS data 
(Morris et al. 1991). 



Table 3 compares the Ly-a lines that were reported by the GHRS and the FOS groups. 
The agreement is generally good. No GHRS equivalent widths were reported for the three 
Virgo lines (two were detected). The only striking disagreement is the absence in the GHRS 
list of the line in the FOS list at 1317.08 A. We reexamined two separate FOS flat-fields and 
verified that the feature in question is not created by our flat fielding. We also compared the 



53 



TABLE 3: Lyman Alpha Absorption Lines 



GHRS A 


FOS A 


GHRS EW 


FOS EW 


1220.0 


1219.80 




0.371 


1222.2 


1222.12 


.... 


0.414 


.... 


1224.52 




0.240 


1242.17 




0.027 


.... 


1247.54 




0.032 


.... 


1251.46 




0.120 


.... 


1255.70 




0.074 




1275.19 


1275.23 


0.144 


0.251 


1276.54 




0.068 


.... 


1289.79 




0.052 


.... 


1292.84 


.... 


0.063 


.... 


1296.57 


1296.52 


0.302 


0.287 




1317.08 




0.292 


1322.16 




0.075 


.... 


1324.96 


.... 


0.027 


.... 


1325.22 


1325.10 


0.057 


0.238 


1361.63 


1361.53 


0.126 


0.146 


1393.86 




0.333 





strength of the feature in the observations through the five different apertures and verified 
that the 1317 A hne is, as expected if it is a real feature, strongest in the narrow slit and 
broader in the large 4.3" aperture. We conclude that the 1317.08 A line is an intrinsic feature 
of the spectrum. 

The GHRS team analyzed lines with equivalent widths much less than the lower limit 
we adopted for inclusion in our complete sample. The lower hmits adopted by the GHRS 
team were, respectively, 0.025 A and 0.050 A, for "weak" and "strong" hues. The equivalent 
width limit we adopted was 0.250 A. This diflference between the equivalent width limits is 
reflected in Table 3 by the presence of a number of weak hnes reported by Morris et al. but 
not in Paper I. There are difficulties (see, e.g., the discussion of Morris et al. 1991 ) associated 
with including the small equivalent width lines. The signal to noise level decreases as one 
goes to smaller equivalent widths so that it becomes more difficult to be be sure that the 
weaker lines are not caused by flat-fielding or other sources of noise. 

6 EVOLUTION OF Ly-a SYSTEMS 

The number density and equivalent width distribution of Ly-a clouds at large redshifts 
has been determined by major surveys of many different sources. We compare with our 
observations of 3C 273 the number of Ly-a systems predicted for small redshifts by the 



54 



formulae used to fit the observations at large redshifts. Using the conventional parameter- 
ization (Sargent, Young, Boksenberg, and Tytler 1980; Murdoch, Hunstead, Pettini, and 
Blades 1986; Lu, Wolfe, and Turnshek 1991), the expected number, N, of Ly-a clouds in the 
spectrum of a nearby quasar may be written 



iV(^en„ Vr > W^cutoff) = ^0(1 + l)-' [{I + Z^r.)'^' - l] exp [- (-^) 



(3) 



where VKcutoff is the minimum considered rest equivalent width. The values of Aq that 
are determined by extrapolating large redshift observations are uncertain, but have been 
estimated recently by Murdoch et al. (1986) and by Lu et al. (1991). 

The most direct comparison between expectation and observation is obtained by tak- 
ing the ratio of predicted to observed number of systems assuming the same parameters, 
including minimum equivalent width, Wcutoff, as for the large redshift surveys. Thus for the 
Murdoch et al. survey, with an equivalent width limit of 0.32 A, one expects 0.76 systems 
for 3C 273 . For the Lu et al. analysis, one expects 0.506 systems with equivalent widths at 
least equal to 0.36 A. We see that from Table 1 that there are exactly two hues that sat- 
isfy the equivalent width limits of the large redshift surveys, the Virgo lines at 1219.80 and 
1222.12 A. Including the Virgo Hnes, the ratio of observed to expected number of systems is 
2.6 for the Murdoch et al. survey and 4.0 for the Lu et al. analysis. If we represent in the 
conventional way the evolution of the number density between small and large redshifts by 
a power law, dN/dz oc (1 + z^, then we obtain 

7 = 1.0 Murdoch et al. parameters; 7 = 1.2 Lu e< al. parameters. (4) 

We stress that for direct comparison with the analyses of large redshift samples only the 
two strong Ly-a systems (presumably) produced in the Virgo cluster can be used. Obviously, 
there are large systematic and statistical uncertainties in the value of 7 that are not included 
in Eq. (4). A number of independent lines of sight must be studied before the evolution of 
the Ly-a clouds can be reliably determined. 

We can obtain some additional information by considering Ly-a systems with EWs 
above 0.200 A, since Murdoch et al. (1986) (see their Figure 3) have presented evidence that 
the equivalent width distribution has the canonical exponential form above this limit but 
changes shape for weaker lines. According to the FOS observations summarized in Table 1, 
the line of sight to 3C 273 contains seven Ly-a systems above 0.200 A. The observed number 
of lines with equivalent widths greater than this value of M^cutoff is 6 times larger than 
predicted by the parameters determined by Murdoch et al. and is 12 times larger than 
predicted by the parameters of Lu et al. . Even ignoring the lines that apparently arise from 
the Virgo cluster, the number of observed Ly-a systems is larger than expected at the 90% 
confidence level. 



55 



The two strong Virgo Ly-a lines are not included in the GHRS analysis since they did 
not report the equivalent widths of these lines. As discussed by Morris et al. (1991), the 
determination of the evolution of the Ly-a clouds with the published GHRS observations 
requires additional information about the weak hues at large redshifts. Therefore, Morris 
et al. used the form and parameters derived by Murdoch et al. (1986) for the equivalent width 
distribution. However, Morris et al. also noted that the conclusions of Murdoch et al. referred 
to a domain of much larger EWs and therefore the GHRS group investigated other studies 
of weaker lines at large redshifts. The uncertainties and controversies involved in using weak 
lines include the difficulty of separating real lines from spurious features and the possibihty 
that the shape of the equivalent width distribution changes with redshift for weak lines. 
The listener is referred to the papers of Morris et al. (1991), Murdoch et al. (1986), and Lu 
et al. (1991) for a discussion of these points. 

7 SUMMARY AND DISCUSSION 

Approximately a quarter of a century ago, before the discovery of any absorption features 
in the spectra of quasi-stellar sources, it was predicted (Bahcall and Salpeter 1966 ) that 
intergalactic gas in clusters of galaxies would produce ultraviolet quasar absorption lines. In 
the first direct test of this idea, we find in the direction of 30 273 three Ly-a absorption 
systems that are probably associated with the Virgo cluster. It is not clear whether or not 
our observations confirm this hoary prediction since the lines could arise from individual 
clouds in large halos of galaxies within the cluster (see Bahcall 1975). One possible way of 
distinguishing between gas in halos of galaxies and gas between galaxies is in the width of 
the hues. Intergalactic gas might well produce broad absorption lines reflecting the velocity 
dispersion within the cluster potential, but there may also be locaUzed clumps of gas between 
galaxies in the cluster that produce relatively narrow lines. 

The observation of a larger-than-expected number of Ly-a absorption lines in the spec- 
trum of 30 273 is an encouraging sign for future HST studies of quasar absorption line 
systems. However, a reliable determination of the cosmic evolution of the Ly-a systems will 
require observations of a number of independent lines of sight. 

We express our gratitude to the FOS instrument team for building a superb instrument. 
This work was supported in part by NASA contract NAS5-29225 and STSOI grant #2424. 



56 



REFERENCES 

Bahcall, J. N. and Salpeter, E. E. 1966, Ap. J. Letters 144, 847. 

Bahcall, J. N. 1975, Ap. J. Letters 200, LI. 

Bahcall, J. N. 1979, in Scientific Research with the Space Telescope, lAU Colloquium 54, 
ed. M.S. Longair and J. W. Warner (superintendent of documents, U.S. Government 
Printing Office, Washington, DC, 20402. 

Bahcall, J. N., Jannuzi, B. T., Schneider, D. P., Hartig, G. F., Bohlin, R., and Junkkarinen, 
V. 1991, Ap. J. Letters 377, L5, (Paper I). 

Ford, H.C. 1985, Faint Object Spectrograph Instrument Handbook, (Space Telescope Science 
Institute: Baltimore). 

Grady, C.A., and Taylor, M.A. 1989, lUE Data Analysis Guide, lUE Newsletter, No. 39. 

Giovanelli, R. and Haynes, M.P. 1989, Ap. J. L., 346, L5. 

Huchra, J. P. 1985, in The Virgo Cluster of Galaxies, eds. O.-G. Richter and B. Binggeli, 
(European Southern Observatory: Garching), p. 181. 

Jannuzi, B.T. , Hartig, G. F., Bahcall, J. N., and D. P. Schneider 1992, in preparation. 

Lu, L., Wolfe, A.M., and Turnshek, D.A. 1991, Ap.J., 367,19. 

Lynds, C. R. 1971, Ap.J. (Letters) 164, L73. 

Morris, S. L., Weymann, R. J., Savage, B. D., and GiUiland, R. L. 1991, Ap. J. Letters 377, 
L21. 

Murdoch, H.S., Hunstead, R.W., Pettini, M., and Blades, J.C. 1986, Ap.J. 309, 19. 

Sargent, W.L.W., Young, P.J., Boksenberg, A., and Tytler, D. 1980, Ap.J.Suppl. 42, 41. 

Sargent, W.L.W., Boksenberg, A., Steidel, C. C. 1988, Ap.J.Suppl. 68, 539. 

Ulrich, M. H., et al. 1980, M.N.R.A.S. 192, 561. 



57 



RESULTS AND SOME IMPLICATIONS OF THE GHRS OBSERVATIONS 

OF THE LYMAN a FOREST IN 3C273 



Ray J. Weymann 

Observatories of the Carnegie Institution of Washington 

813 Santa Barbara Street 

Pasadena, CA 91101 

USA 



Abstract. 

The results of the first GHRS exposures of 3C273 and related observations are sum- 
marized, and their implications for our understanding of the Ly a forest assessed. The 
potential for further understanding of the Ly a forest and intergalactic medium through 
additional HST and ground-based programs is briefly discussed. The number of low 
redshift Ly a lines found is substantially higher than expected based upon extrapola- 
tion from ground-based data using current estimates of the slope of the log (dN/dz) 
vs. log (l-|-z) relationship. However, we consider the more significant results to be the 
small value of the slope (~ 0.80) between redshift and 2 and the establishing of the 
very existence of low redshift Ly a clouds in numbers sufficient for their properties to 
be investigated in some detail. 



1. INTRODUCTION 

The discovery of numerous Ly a clouds in the line of sight between ourselves and 
3C273 (Morris et al. 1991, hereafter MWSG; see also Bahcall et al. 1991, hereafter 
BJSHBJ) presents us with the opportunity of studying the evolution of the properties 
of the Ly q forest clouds (and possibly the intergalactic medium as well) for over about 
90 % of the age of the Universe. Since an account of the observations and reductions 
has already been given in MWSG, in §2 we simply summarize the results of these 
observations and touch upon a few highlights. In §3 we discuss what can be inferred 
about the evolution of the Ly a forest based only upon some extremely scanty current 
data. Finally, in §4 we discuss some of the open questions concerning the Ly a forest, 
and some possible future observations, both space and ground-based, which may shed 
some light on these questions. 



2. SUMMARY OF THE 3C273 OBSERVATIONS 



58 



In Table 1, collated from Table 1 of MWSG and Table 1 of BJSHBJ, we list all 
the possible Ly a absorption lines seen in either the GHRS or FOS observations. As 
described in MWSG, the quantity -log P is a measure of the probability, P, that the 
absorption feature arises by chance from photon statistics. Although an analysis of 
the noise properties of these exposures by Gilliland et al. (1991) suggests that the 
noise characteristics are well-approximated by Poisson statistics, the weaker absorption 
features should be viewed with some caution since some very weak features due to 
variations in photocathode response could still be present even though the FP-SPLIT 
mode was used. From inspection of Table 1 it is seen that the agreement in wavelength 
between the GHRS and FOS observations is quite satisfactory, though the agreement 
in equivalent widths between the GHRS and FOS observations is only fair and becomes 
poorer for the weaker lines. The following points deserve comment: 

i) A line at A1317.08A was identified as Ly a in BJSHBJ but is almost certainly 
galactic Ni II. Curiously, this line was not seen in the GHRS G160M exposure, 
but is clearly visible in the deconvolved GHRS 140L exposure, with a measured 
wavelength differing by only 7 km s~ from the FOS value. The corresponding Ni II 
line at A1370 is present in the GHRS 160M exposure at the expected wavelength 
to within a few km s~ . We therefore do not list the A1317A line in Table 1. 

ii) There has been considerable discussion lately over the value of the minimum doppler 
parameter in the Lyman a forest clouds, (c/. Webb and Carswell 1991; Hunstead 
and Pettini 1991 and references therein.) It is thus of interest to examine the 
minimum doppler parameter found in the GHRS 3C273 data. The narrowest well- 
determined doppler parameter is for the line at Al36lA. It is only "well-determined" 
in the sense that this line appears on two separate GHRS G160M exposures, and 
two separate Voigt profile fits yield column densities and doppler parameters in 
good agreement. After these two exposures are coadded we find the best value of 
b to be 22.0 km s"" with a formal \-a uncertainty of 3.2 km s~ and a value for 
X per degree of freedom of 1.035 with 50 degrees of freedom. In fact, the doppler 
parameter is not too well constrained. Fixing the doppler parameter at various 
values and allowing the fitting routine to fit the column density and redshift, we find 
that the x probability drops below 5% {i.e., the fit starts to become discernably 
poor) only for b outside the range 12 < b < 37 km s~ . (In addition to the 
uncertainties due to photon counting statistics, since these observations were taken 
through the LSA and the PSF for this mode is somewhat uncertain (Gilliland et al. 
1991), there is some additional uncertainty in the doppler parameter and the value 
for the doppler parameter just quoted differs from that given in MWSG because a 
slightly different PSF was assumed.) 

iii) As noted in MWSG, the Si IV doublet arising in the galactic halo is anomalous 
in that we measure an equivalent width ratio of A1394 to A1403 of greater than 
2.0. Moreover, the centroid of the A 1394 component is significantly discrepant in 
velocity compared to other high ionization galactic halo lines. This led MWSG to 
suggest that another line was blended with the Si IV A1394 line, presumably another 
intervening Ly a line. As described in MWSG, one can infer the optical depth of the 
extra contributor by using the A 1403 line as a template for the Si IV column density 
vs. velocity and applying that to the observed A1394 feature. Figure 1 shows the 
result of this procedure. The profile with the large equivalent width is the observed 



59 



Table 1: Lyman o Lines in 3C273 



^GHRS 


^FOS 


EWgh/js 


EWfos 


cz 


b 


logN 


-log(P) 


Note 


(A) 


(A) 


(mA) 


(mA) 


(kms-i) 




(cm^) 






1220.0 


1219.80 




371 


1018 






L 


a 


1222.2 


1222.12 




414 


1591 






L 


a 


1242.17 




27 




6535 


23 


12.87 


6.0 


b 


1247.54 




32 




7859 


49 


13.01 


5.2 




1251.46 




120 




8826 


36 


13.48 


>7.5 




1255.70 




74 




9872 


111 


13.30 


>7.5 




1275.19 


1275.23 


144 


251 


14678 


24 


13.55 


>7.5 




1276.54 




68 




15011 


17 


13.04 


>7.5 




1289.79 




52 




18278 


63 


13.06 


5.9 




1292.84 




63 




19031 


34 


13.12 


>7.5 




1296.57 


1296.52 


302 


287 


19950 


28 


14.13 


>7.5 




1322.16 




75 




26261 


56 


13.23 


>7.5 




1324.96 




27 




26952 


45 


13.05 


5.9 


c 


1325.22 


1325.10 


57 


238 


27016 


18 


13.57 


>7.5 


c 


1361.63 


1361.53 


126 


146 


35995 


20 


13.52 


>7.5 




1393.86 




331 




43943 


37 


13.97 


D 


d 



Probability code: L=GHRS observation with G140L only; D=not resolved; see text for 

explanation of -log(P) values. 

cz values from GHRS wavelengths unless noted. 

Notes: 

a: on wing of galactic Lyman a, cz from FOS 

b: line did not improve x^ of fit to NV doublet 

c: parameters uncertain, dependent upon deblending; these two lines unresolved by FOS 

d: Unresolved from Si IV A1394 galactic halo line; see text for details 



60 




1393 



1393.5 1394 

lambda (Angstroms) 



1394.5 



Fig. 1 — The result of deblending the A1394 complex into the Si IV 
contribution and the inferred Lya contribution. The profile with the large 
equivalent width is the observed profile (after 5— smoothing) and the profile 
with the smaller equivalent width is the inferred Si IV. The large amount 
of residual absorption is attributed to a Lya line. 



61 



blend (after 5-smoothing), while the profile with the smaller equivalent width is 
the predicted Si IV A1394 profile. Evidently there is a very large residual amount 
of absorption. Alternatively, we may find the best fit to the observed profiles using 
Voigt profiles for the Si IV doublet and an additional feature at A1394. The former 
method has the advantage that it does not require the assumption of a Voigt profile, 
but has the disadvantage that it does not take account of the finite PSF (at least as 
we have applied this method.) The Ly /9 line should be observable with the GHRS 
140M grating as and when side 1 is operable. 

3. THE EVOLUTION OF THE LY a FOREST 

The justifiable excitement and gratification over the discovery that Ly a forest lines 
appear to exist at very low redshifts has led, in our view, to overemphasis of the degree 
to which the density of lines exceeded "expectations". There were really only two 
bases upon which predictions of the line density at very low redshift could be made: 1) 
Extrapolation of fits to the line density from ground-based data. 2) Models. Needless 
to say, both of these methods carried with them enormous uncertainties. 

With respect to the extrapolation, not only must the extrapolation of the line 
density per se be carried out over a huge range in log (1+z), but one must tacitly 
assume that the distribution of equivalent widths also does not vary unless one is 
comparing samples whose limiting equivalent widths are the same. In addition, the 
data base for ground-based observations is weighted rather heavily to redshifts above 
2, so that despite the very large number of lines, 7, (the slope of the log(dN/dz) vs. log 
(1+z) relation) is not that well determined. For example, in an analysis using over 900 
lines and 38 QSOs, Lu et al. (1991) obtained a value of 7 = 2.75 ± 0.29 (when lines in 
the sample likely to be affected by the proximity effect are removed), but found some 
evidence that evolution steepened below a redshift of ~ 2.3, with a broken power law 
fit having a value of 7 = 4.21 below this value. On the other hand Rauch et al. (1991), 
using a smaller number of lines, but based upon higher-resolution data (and heavily 
weighted to redshifts above 2) found 7 = 1.68 ± 0.80 for a sample having the same 
cutoff in line strength. It is possible that a global fit for 7 over all redshifts might yield 
values not very different from the Rauch et al. value. In this sense the number of lines 
at low redshift is not "unexpected". On the other hand, it is clear that the predicted 
number of lines at low redshift is incompatible with the fit obtained by Lu et al. , for 
both the single power law and especially for the broken power law. More significantly, 
when a value of 7 is determined from the HST data at z ~0.05 and ground-based data 
at z~2.0, it is much flatter than either the Lu et al. or Rauch et al. value. 

These results hold when objects other than 3C273 are considered. While it borders 
on the foolish to quote results based upon only 4 objects, especially when vastly more 
extensive data will be soon forthcoming (and indeed are already in hand), some results 
based upon the 4 objects available to the author at the time of writing are as follows: 

i) For weak lines — irst equivalent width (REW) > 50mA — we believe the 3C273 sam- 
ple is probably complete. From Table 1 there are 10 such lines over the range from 
z=0.016 to z=0.151. (This includes the inferred line at A1394 but not the two 
Virgo cluster lines, since the minimum z sampled by the G160M exposures does not 
include these lines. The upper limit is adopted, as in MWSG, to take account of 
the proximity effect from 3C273.) We compare this with a sample having the same 
REW cutoff taken from the data published by Carswell et al. (1991) for QllOl-264. 



62 



A simple linear fit to the two locally determined values of dN/dz gives a value of 7 
= 0.82, close to the maximum likelihood value quoted in MWSG. 

ii) For lines of intermediate strength (REW > 200mA — cj. BJSHBJ), we can adjoin to 
the G160M 3C273 list of 4 lines (we use the GHRS values of the equivalent widths) 
the two additional Virgo cluster lines seen in the FOS and G140L GHRS spectra. In 
addition, we judge the PKS2155-33 observations of Boggess et al. (1991) to also be 
complete to this limit. In this object, Maraschi et al. (1988) reported the presence of 
two absorption lines which they interpreted as intervening Lycv. However the HST 
observations show^ that the feature measured at A1237 and tentatively identified 
by Maraschi et al. as N V is at A 1236 which is sufficiently far from the galactic 
halo N V A123S line that we think it unlikely to be attributable wholly to N V. 
In addition, they show that the feature at A=1285 consists of two well-separated 
(AV~600 km s~^ ) features. The local value of dN/dz at z~2.0 for lines with REW 
> 200mA is again taken from the Carswell et al. data for Ql 101-264. In computing 
the local value of dN/dz for z~0 for this range of line strengths a matter of principle 
arises: It will be noted that the galactic Lya line has not been included in the list, 
since it is clear that our preferred position in a spiral galaxy makes all lines-of-sight 
atypical at zero redshift. To what extent should our preferred position in the local 
supercluster and rather near the Virgo cluster cause us to give less than full weight 
to the Virgo cluster clouds? A large number of lines-of-sight not passing through 
the Virgo cluster will soon make this question moot, so we shall not dwell on it, but 
simply quote two values of 7 for this range of line strength: With the two Virgo 
clouds we find 7 = 0.54 , and without the two Virgo clouds, 7 = 0.83. 

iii) Finally, for lines with REW > 360mA — the threshold adopted by Lu et al. — we 
add lines from the two objects CS0251 and PG121H-143 (Burbidge et al. 1991). 
In these two objects there are an additional two certain and two probable Lya lines 
with REW > 360mA. ^ As in case ii) above, we considered both the cases in which 
the Virgo lines were and were not counted, and cases in which the two probable 
lines just mentioned were and were not counted. For the high redshift end of the 
fit we used the line density found by Lu et al. at z=2.3. The corresponding values 
of 7 range between .73 and 1.24. 

The actual values of 7 thus determined are, of course, very uncertain due to the 
small number of lines in all three cases considered. Neverless, if one assumes the validity 
of the Lu et al. fit, the number of lines actually observed at low redshift would be 
extremely unlikely to occur by chance. In other words, the Lu et al. fit does not extend 
to low redshifts, but becomes much flatter. 

With respect to whether the number of zero redshift lines was "unexpected" from 
the point of view of models, a number of authors have, after the fact, shown that such 
behavior can arise quite naturally. Indeed, the possibility of a flattening of 7 or even 
of a turnup in the number of lines per unit redshift at low redshifts was explicitly 
noted several years ago by Bechtold et al. (1987; see their Figure 8) using a simple 



^ I thank Drs. Boggess and Bruhweiler for permission to quote this result in advance 
of publication. 

^ I thank Dr. V. Junkkarinen for communicating these two spectra and Drs. Burbidge 
and Junkkarinen for permission to quote these results in advance of publication. 

63 



pressure-confined quasi-static model with the confining pressure undergoing adiabatic 
expansion, together with their best estimate of the evolution of the intergalactic ionizing 
radiation field. 



4. CURRENT PROBLEMS AND FUTURE OBSERVATIONS 

A detailed review of models of the Lya clouds and their evolution is beyond the 
scope of this paper. However, the following comments are relevant in connection with 
possible future observations, both ground-based and with HST. 

It goes almost without saying that an important constraint on models will be pro- 
vided by the detailed picture of the evolution of the Lyo forest line density over the 
entire range from 0.0 to ~5.0 and the range from 0.0 to ~1.6 will undoubtedly be 
provided by the Quasar Absorption Line Key Project provided only that the HST and 
spectrographs do not lose further capability. As noted in §3 however, there is still room 
for substantial improvement in delineating the line density as a function of both line 
strength and redshift, especially in the regime from ~1.6 to ~2.3. It would also be use- 
ful to devise a measure for the line density which is not sensitive to the decomposition 
into multiple components, perhaps along the lines pioneered by Webb et al. (1991) and 
Jenkins and Ostriker (1991) in the context of the Gunn-Peterson trough analysis. 

In §2 we noted the interest in the possible existence of lines with very small values 
of the doppler parameter. However the question of the origin of doppler parameters 
with large formal values is also of interest [cf. Ranch et al. 1991). Values of up 
to 50 and even 60 km s~ are derived even when based upon high resolution data. 
Such values surely cannot represent thermal widths, and the question is whether they 
are to be attributed to bulk motion or discrete components. If the latter, then as 
noted by Ranch et al. , the two-point correlation function will have a strong peak at 
velocities of order 20-50 km s~ , though the interpretation presumably has nothing 
to do with gravitational clustering, but would more likely reflect fragmentation in a 
hydrodynamical process. Simulations show that it is generally possible to distinguish 
between these two origins for the super-thermal doppler parameters at resolutions now 
being employed for ground-based observations, but better signal-to-noise is required. 

At least for 3C273 it is feasible to carry out such observations as well, and this 
should be done to see if the origin of the large formal doppler parameters is the same 
at low redshift as it is at high redshifts. 

The question of the characteristic size of the Lya clouds is still an important open 
question. The results of Foltz et al. (1984) on the pair of QSOs 2345+007A,B have been 
widely quoted as providing a characteristic size of the clouds of a few kpc. Subsequently, 
Steidel and Sargent (1990) examined this pair at lower resolution but higher S/N and 
concluded that it was unlikely to be a gravitational lens, but rather a true double 
QSO. This change in the geometry would imply that the characteristic size was as 
much as an order of magnitude larger than the value which Foltz et al. suggested. 
Most recently however, these same authors (Sargent and Steidel 1991), on the basis of 
further data, have concluded that the weight of the evidence favors the gravitational lens 
interpretation after all. However, they have discovered additional absorption systems 
exhibiting metals, and at least 3 of the 4 Lyo lines found by Foltz et al. to be common 
to the two lines of sight belong to these systems. Thus, only at most one of the lines 
in common to the two lines of sight is a Lya-only system. The consequence of this is 
that the constraint on the sizes of the clouds is decidedly relaxed and becomes more in 
the nature of a lower limit. Intensive spectroscopic study of Q2345A,B should clearly 

64 



be a high priority project for the new generation of large telescopes, as also stressed by 
Sargent and Steidel. 

In the meantime, at least one ambiguity in connection with the interpretation of the 
common lines in double or gravitationally lensed QSOs appears to have been cleared 
up. Until now, there has been no compelling reason to believe that lines in common 
represent the same physical cloud, as opposed to a swarm of small cloudlets. Recently 
however, Smette et al. (1991) have shown that the equivalent widths of the lines in 
common between the pair UM671A,B are strongly correlated, suggesting that single 
large clouds are involved and setting a firm lower limit of order 20 kpcs to the cloud 
sizes. This pair too should be studied at higher resolution and high S/N. 

The characteristics sizes above refer to redshifts ~2, and it would evidently be of 
considerable interest to place similar constraints on the characteristic sizes at very low 
redshifts. At least one pair. Ton 155/156 offers the possibility of examining this question 
with HST, although an lUE exposure (Malkan 1991) of Ton 155 suggests that a strong 
Ly limit system may be present at a redshift of about 1.2-1.4 which has absorbed most 
of the far UV flux. 

A final question of considerable interest which can be readily investigated for the low 
redshift Lya clouds towards 3C273 (and of course towards all other QSOs) involves the 
question of the association of the clouds with optical images (if any) and the statistical 
correlation of the clouds with galaxies and clusters of galaxies (if any). 

As pointed out by MWSG, direct imaging of the clouds themselves by means of 
their recombination radiation does not appear feasible except for extremely high col- 
umn densities, unless the clouds are exposed to UV radiation fields several orders of 
magnitude higher than the estimated value of the intergalactic radiation field at zero 
redshift. Direct "imaging" in 21 cm radiation is also not feasible. However, it could well 
be that the Lyo clouds are shreds of H I on the outer edges of, e.g., dwarf galaxies which 
might have some Hq emission. Such dwarf galaxies may have extremely low surface 
brightnesses and would be very difficult to detect in continuum radiation, especially very 
near a bright object like 3C273. In a collaboration involving S. Morris, R. Schommer, 
R. VVeymann and T. Williams, a very preliminary set of observations were obtained 
by Schommer and Williams at the CTIO 4m telescope using the Rutgers Fabry-Perot 
Interferometer, and scanning over the 1600 km s~ and 1000 km s~ regions with a 
FWHM of 2.5A. No obvious features appear which are clearly above the noise (of a few 
electrons), although some faint emission filaments about 30-40 arcsec from the QSO 
may possibly be present at about the noise level (Schommer 1991). Before any decision 
can be made about the reality of such possible features, much more extensive data will 
need to be obtained. In a separate collaboration, custom narrow-band filters have been 
obtained to search for emission features at the two Virgo cloud redshifts as well as the 
redshift space around the feature at 125lA, using the Smith-Terrile coronograph at the 
Las Campanas duPont telescope. 

In addition to a search for emission features closely associated with the Lya clouds, 
an obvious related program is to check for statistical correlations (or anti-correlations) 
between individual galaxies and groups and clusters of galaxies and the Lya clouds. 
Indeed, preliminary investigations along these lines have already been carried out by 
Salzer (1991) and by Jaaniste (1991). In particular, Salzer has examined currently 
available redshifts in the neighborhood of 3C273 to identify those galaxies closest to 
the Virgo Lya clouds and to look for structures which might be associated with some 
of the more distant clouds. Interestingly, one of the closest associations appears to 
be between one of the clouds and the large H I cloud 1225-f-Ol, a possible association 

65 



already noted by BJSIIBJ. The data is still too scanty for any definitive conclusions to 
be reached concerning the extra- Virgo clouds. Collaborative programs are underway 
at LCO and CFHT to provide extensive redshift and color data on galaxies near the 
3C273 line-of-sight. 

As a final quasi-philosophical remark, it is certainly not clear that the Lyo: systems 
found at very low redshift represent the same phenomenon seen at high redshift. There 
has been extensive debate over whether the low- and high-column density Lya systems 
at high redshift are best regarded as members of a "single population" or should be 
considered two separate populations. Suppose the result should emerge that the low 
redshift Lya clouds are indeed associated with galaxies and clusters of galaxies (we 
already know that two such clouds are.) There may then be a tendency to conclude that 
the low redshift clouds belong to a different population than the high redshift clouds, 
since the high redshift clouds "don't cluster like galaxies". (A more accurate statement 
is that the two point correlation function for the Lya-only systems shows very little 
structure whereas that for the higher column density systems exhibiting C IV does.) 
This conclusion would be premature. In general, the existence of a property {e.g., the 
correlation function) spanning a wide range as another parameter {e.g., column density 
or redshift) varies is not an argument in and of itself for two populations. To make a 
convincing case for two populations, it will be necessary to trace the strength of the 
association between the Lya clouds and galaxies (if such an association is found at low 
redshift) back to higher redshifts and then find other properties {e.g., metal abundance) 
which discriminate between Lya clouds which do and do not associate with galaxies. 
Only then could one consider the two-population case to have been made. 



REFERENCES 

Bahcall, J.N., Jannuzi, B.T., Schneider, D.P., Hartig, G.F., Bohlin, R., and Junkkarinen, 

V. 1991, Ap. J. (Letters), 377, L21. (BJSHBJ) 
Bechtold, J., Wcymann, R.J., Lin, Z., and Malkan, M.A. 1987, Ap. J., 315, 180. 
Boggess, A., Bruhweiler, F., Kondo, Y., Urry, M., Grady, C, and Norman, D. 1991, in 

preparation. 
Burbidge, E.M., Cohen, R., Junkkarinen, V. and several other FOS team members. 

1991, in preparation. 
Carswell, R.F., Lanzetta, K.M., Parnell, H.C., and Webb, J.K. 1991, Ap. J., 371, 36. 
Foltz, C.B., Weymann, R.J., Roser, H.-J., and Chaffee, F.H. 1984, Ap. J. (Letters), 

281, LI. 
Gilliland, R.L., Morris, S.L., Weymann, R.J., Ebbets, D., and Lindler, D. 1991, in 

preparation. 
Hunstead, R.W., and Pettini, M. 1991 in Proc. of the ESO Mini-Workshop on Quasar 

Absorption Lines ESO Scientific Report No. 9 Feb. 1991, p. 11. 
Jaaniste, J. 1991, (preprint submitted to Baltic Astronomy.) 
Jenkins, E.B. and Ostriker, J. P. 1991, Ap. J., 376, 33. 
Lu, L., Wolfe, A.M. and Turnshek, D.A. 1991, Ap. J., 367, 19. 
Malkan, M. 1991, private communication. 
Maraschi, L., Blades, J.C., Calanchi,C., Tanzi, E.G., and Treves, A. 1988, Ap. J., 333, 

660. 
Morris, S.L., Weymann, R.J., Savage, B.D., and Gilliland, R.L. 1991, Ap. J. (Letters), 

377, L21. (MWSG) 



66 



Rauch, M., Carswell, R.F., Chaffee, F.H., Foltz, C.B., Webb, J.K., Weymann, R.J., 

Bechtold, J., and Green, R.F. 1991, submitted to Ap.J. 
Salzer, J.J. 1991, submitted to Astron. J.. 
Sargent, W.L.W. and Steidel, C.C. 1991, preprint 
Schommer, R. 1991, private communication. 
Smette, A., Surdcj, J., Shaver, P.A., Foltz, C.B., Chaffee, F.H., Weymann, R.J., 

Williams, R.E., and Magain, P. 1991, submitted to Ap.J. 
Steidel, C.C. and Sargent, W.L.W. 1990, A. J., 99, 1693. 
Webb, J.K., and Carswell, R.F. 1991 in Proc. of the ESO Mini-Workshop on Quasar 

Absorption Lines ESO Scientific Report No. 9 Feb. 1991, p. 3. 
Webb, J.K., Barcons, X., Carswell, R.F., and Parnell, H.C. 1991, preprint. 



67 



Hot Stars and the HST 



R.P. Kudritzki 

Institute fur Astronomie und Astrophysik der Universitat Miinchen 

Scheinerstr. 1 

8000 Miinchen 80 

Germany 



Abstract. 

The HST is ideally suited to the observation of hot massive stars which emit much of 
their prodigious energy output as radiation in the ultra-violet region of the spectrum. 
These most luminous of stars can be identified directly in galaxies as distant as the 
Virgo cluster or indirectly through their illumination of giant HII regions, such as 
the 30 Doradus complex in the Large Magellanic Cloud. They are therefore ideal 
standard candles and tracers of young populations providing important information 
about abundances, star formation, energetics of the ISM (radiation, stellar winds) and 
nucleosynthesis. 

On the other hand, thanks to dramatic advances in NLTE model atmosphere tech- 
niques the methods of quantitative spectroscopy of hot stars have experienced great 
progress. Model atmospheres are now available that include the opacity of thousands 
to tens of thousands of lines fully in NLTE and take into account the radiation hy- 
drodynamics of stellar winds. This has opened the door to determine precisely the 
stellar parameters of luminosity, effective temperature, gravity, mass, radius, distance, 
chemical composition as well as the stellar wind parameters, mass loss rate and velocity 
structure not only in our own galaxy but also in local group galaxies and somewhat 
beyond. In addition, new ionizing model atmosphere fluxes are becoming available that 
will allow a more realistic interpretation of nebular recombination spectra. 

The crucial parameter determining the properties of hot stars is metallicity. It 
affects the ionizing energy distribution, the spectral appearance, the stellar wind prop- 
erties and the formation and evolution of hot stars. With HST it will be for the first 
time possible to study quantitatively the physics of massive stars in galaxies of different 
metallicity, in particular by obtaining high quality ultraviolet spectra of hot stars in 
the Magellanic Clouds. 

In this connection, we report first HST observations obtained with the GHRS of 
the 03f star Melnick 42 in the 30 Doradus complex of the LMC. A first analysis of the 
excellent spectra reveals that with a luminosity of 2.3 x 10 L(7, and a present mass of 
100 A/0, the object is one of the most massive stars known. An estimate of abundances 
indicates that iron and oxygen are very likely reduced by a factor of four relative 
to the sun, whereas carbon is more strongly depleted and nitrogen is approximately 
solar. The terminal velocity of the stellar wind is 3000 km/sec. The mass-loss rate is 

68 



4 X 10~ MQ/year, with a large uncertainty. 

Most of the content of this paper has been discussed in recent reviews by Kudritzki 
and Hummer (1990) and Kudritzki et al. (1991) and the very recent publication on 
first results obtained with the GHRS in the Ap.J. Letters by Heap et al. (1991). In 
consequence, to avoid simple duplication of paper, only this summary is published here. 

REFERENCES 

Heap, S.R., Altner, B., Ebbets, D., Hubeny, I., Hutchings, J.S., Kudritzki, R.P., Voels, 
S.A., Haser, S., Pauldrach, A., Puis, J., Butler, K., 1991, Ap. J. (Letters), 337, L29. 

Kudritzki, R.P., Hummer, D.G., 1990, Ann. Rev. Astr. Ap., 28, 303. 

Kudritzki, R.P., Gabler, R., Kunze, D., Pauldrach, A., Puis, J., 1991, in Massive Stars 
in Starbursts, STScI Symp. Series No. 5, eds. C. Leitherer et al. p. 59 



69 



GHRS FAR-ULTRAVIOLET SPECTRA OF CORONAL AND 
NONCORONAL STARS: CAPELLA AND 7 DRACONIS 



Jeffrey L. Linsky^ and Alexander Brown 

Joint Institute for Laboratory Astrophysics 

University of Colorado 

Campus Box 440 

Boulder, CO 80309-0440 

USA 



Kenneth G. Carpenter 

NASA Goddard Space FHght Center 

Code 681 

Greenbelt MD 20771 

USA 



Abstract. We report on the first GHRS spectra of two very different late-type giant 
stars - Capella and 7 Dra. CapeUa is a 104 day period binary system consisting of two 
stars (G9 III and GO III) each of which shows bright emission lines formed in solar-like 
transition regions and coronae. By contrast, 7 Dra is a hybrid-chromosphere star with 
very weak emission lines from high-temperature plasma. Low-dispersion spectra of these 
stars covering the 1160 to 1717 A spectral range show unresolved emission lines from 
neutral species through N V. The very different surface fluxes detected in the spectra of 
these stars suggest different types of heating mechanisms. Moderate dispersion spectra 
of Capella show intersystem lines of C III, N III, III, IV, Si III, and S IV, which 
are sensitive to electron density. Echelle spectra of hydrogen and deuterium Lyman-a, 
Fe II, and Mg II permit measurements of the cosmologically interesting D/H ratio and 
the properties of the interstellar medium on the 13 pc line of sight to Capella. 

1. INTRODUCTION 

The Goddard High Resolution Spectrograph on the HST places new observational 
capabilities in the hands of astronomers studjang the atmospheres of stars and the 
interstellar medium. Both lUE and Copernicus have obtained ultraviolet spectra of 
bright sources in the 1170-3200 A spectral region, but the GHRS wiU expand our 
observational capabilities enormously in at least four ways: 

'Staff Member, Quantum Physics Division, National Institute of Standards and 
Technology 



70 



• The higher throughput of the GHRS and the low background of its Digicon detec- 
tors support photon-limited observations of much fainter sources than heretofore 
feasible and the measurement of weaker emission and absorption lines which are 
buried in the noise of existing ultraviolet spectra. 

• The GHRS can obtain spectra with signal/noise well in excess of 100:1 (Carpenter 
et al. 1991), a major improvement over lUE. This is critical for measuring lines 
profile shapes, Doppler-imaging experiments, and for studying individual velocity 
components in interstellar absorption lines. 

• Small science aperture (SSA) spectra are not noticeably degraded in spectral res- 
olution by the spherical aberration. Wahlgren et al. (1991) determined that at 
1940 A moderate dispersion G160M spectra have a resolution of 28,000 and the 
echeUe spectra have a resolution of 87,000. Large science aperture (LSA) spectra 
are degraded in resolution by a factor of 2 compared to prelaunch expectations, 
but spectral deconvolution techniques can recover most of the lost resolution for 
point sources when the signal/noise is sufficiently large. Except for echelle spectra 
of a few very bright sources obtained with a rocket instrument, the GHRS is the 
highest resolution and most sensitive ultraviolet spectrograph in operation. 

• The very low scattered light level of the GHRS gratings and the solar-blind detec- 
tors make observations of the ultraviolet spectra of very red stars possible. These 
properties are essential for studies of interstellar deuterium, for example. 

In this paper we provide examples of these new capabilities, which wiU yield major 
scientific benefits in the study of cool stars and the interstellar medium. At the same 
time, we should recognize that lUE beautifully complements the strengths of the GHRS 
by its broad spectral coverage in single exposures, its capability to monitor sources over 
many time scales, and its continuing success in observing targets of opportunity. 



2. LOW DISPERSION SPECTRA OF COOL GIANTS 

On 15 April 1991 we obtained GTO low dispersion G140L spectra of Capella, a 
104 day spectroscopic binary system consisting of a slowly rotating G9 III primary 
(Capella Aa) and a more rapidly rotating GO III secondary star (Capella Ab). (See 
Battan, Hill, and Lu 1991 for a discussion of the system parameters.) On the basis 
of lUE spectra, Ayres and Linsky (1980) showed that the GO III star dominates the 
ultraviolet emission line spectrum of the system. During SV, we obtained on 6 April 
1991 low dispersion spectra of the K5 III star 7 Draconis, a member of the class of 
hybrid- chromosphere stars. The ultraviolet spectra of these stars are characterized by 
high-velocity blue-shifted absorption features due to a cool wind (75-200 km s~ ) and 
fcdnt emission lines formed at temperatures up to 150,000K (Hartmann, Dupree, and 
Raymond 1980; Drake, Brown, and Linsky 1984). 

Low dispersion spectra with the G140L grating provide a means of rapidly observ- 
ing broad spectral regions (288 A at one time) with enough spectral resolution (2,000 
with the SSA and roughly 1,000 with the LSA) to measure the fluxes of most important 
emission lines, although higher resolution spectra are needed to separate close blends. 
We first inspect the low dispersion spectra to identify the major differences between an 
"active" star like Capella Ab and a very inactive star like 7 Dra. 



71 



Figures 1 and 3, which display the 1170-1710 A region of Capella, should be com- 
pared with Figures 2 and 4, which display the 1260-1740 A region of 7 Dra. We note 
immediately that the spectrum of Capella is dominated by bright emissions, including 
the resonance lines of C II, Si IV, C IV, and N V formed at temperatures of 20,000- 
150,000K (see Table 1). The brightest feature is Lyman-Q, despite strong interstellar 
hydrogen absorption in its core (see below). In the Sun this line is formed at 40,000K as 
a result of ambipolar diffusion (Fontenla, Avrett, and Loeser 1991), and we suspect that 
the line is formed at similar temperatures in Capella. Emission lines of other neutral 
species are formed in the chromosphere at T < 8,000K and are very weak, except for 
the C I multiplet near 1657 A. 

Table 1. Comparison of Emission Line Surface Fluxes (log units) 



Multiplet 


logT 


Sat. 


V711 Tau 


Capella 


Sun 


7Dra 


aBoo 






Limit 


(RS CVn) 


(GOIII) 


(G2V) 


(K5III) 


(K2III) 


C IV 1549 A 


5.0 


6.0 


5.62 


5.43 


3.73 


2.08 


<2.00 


Si IV 1400 A 


4.8 




5.05 


4.88 


3.37 


1.60 


<2.00 


Si III 1892 A 


4.6 




4.68 


5.26 






1.82 


C II 1334 A 


4.3 




5.46 


5.20 


3.67 


1.93 


<2.00 


I 1304 A 


3.9 




5.24 


4.85 


3.62 


3.70 


3.75 


C I 1657 A 


3.8 




5.07 


5.01 




2.51 


2.84 


Mg II 2800 A 


3.8 


7.2 


6.82 


6.92 


6.07 


4.73 


5.25 



The low dispersion ultraviolet spectrum of 7 Dra shows a very different appearance 
as both Lyman-a and the I 1304 A multiplet dominate over the high-temperature 
Unas. The temperature at which the Lyman-a line forms in inactive K giants is not 
known, but the I Lines are formed in the chromosphere as a result of pumping by 
Lyman-^ (Haisch et al. 1977). The other prominent emission lines are from neutral 
species and the fourth positive bands of CO, all formed at temperatures below 8,000K. 
While the CO bands could be identified in lUE spectra of Arcturus (Ayres, Moos, and 
Linsky 1981; Ayres 1986), the GHRS spectra are of much higher signal/noise and will 
permit more detailed analysis. The high temperature resonance lines of C II to N V 
are all present but very weak compared with the low-temperature emission lines. 

This quaditative difference in the spectra of CapeUa and 7 Dra can be made quan- 
titative by measuring the observed emission line fluxes and converting them to surface 
fluxes by dividing by the square of the stellar angular radii. The latter may be inferred 
simply from the stellar visual magnitudes and colors (Linsky et al. 1979). The surface 
fluxes are given in Table 1, together with corresponding values for the more active 
RS CVn system, V711 Tau (Byrne et al. 1987), the quiet Sun (Ayres, Marstad, and 
Linsky 1981), and the slowly rotating K giant Arcturus (a Boo; Ayres, Simon, and Lin- 
sky 1982). Mg II and other chromospheric line fluxes were obtained from Ayres et al. 
(1982) and Simon, Linsky, and Stencel (1982). We also list the maximum observed 
C IV and Mg II surface fluxes for the youngest and most rapidly rotating stars without 
obvious circumstellar disks. Vilhu (1987) calls these fluxes "saturated" in the sense 
that they represent the maximum radiative emission from a star completely covered 
with "active regions", which are undoubtedly locations of strong magnetic fields. 



72 



E 

u 

\ 

en 



3 



T3 
9) 
> 

a> 
en 

O 



2 








Cape 


HIa 






E-11 


1 




1991 April 15 
Phase = 0.28 


1.5 




— H Ly-a 


25.6 seconds 
GUOL 


E-11 


Si III 




1 C 11 


1 
E-11 


C 111 

1 




r 


1 1 


SI IV 








N V 










5 

E-12 


- 




n 








- 





*w 


^ 


Uu 


J 


u 


LA ii 





1150 1210 1270 1330 

Wavelength 



1390 



1450 



Figure 1.: The GHRS low dispersion spectrum of Capella obtained with the G140L 
grating. The 1170-1450 A spectrum contains emission Hnes formed at 20,000-150,000K. 



< 



CM 

* 
* 

E 
o 

~\ 

0) 

X 

D 



T3 

> 

u 
<U 

in 
i3 

O 



1.5 
E-13 



1.125 
E-13 



7.5 
E-14 



3.75 
E-14 



7 Draconis 



S 1 - 


r-- 1 1 1 




- 


;■" 


CO 


- 


1 1" 


- 


C It; 


:; Si IV 


C IV 

1 


w 


UJ 


VAImww 


WJ 



1260 1320 



1380 1440 

Wavelength 



1500 



1560 



Figure 2.: The GHRS low dispersion spectrum of 7 Draconis obtained with the G140L 
grating. The 1260-1550 A spectrum contains the bright I resonance hne multiplet 
(1302, 1304, and 1306 A) blended with S I Hnes. The high-temperature transition region 
lines are very weak compared to those in the Capella spectrum. The fourth positive 
bands of CO are indicated. There are many weak emission lines that are not identified 
here, but no evidence for photospheric or other continua. 



73 



CN 
* 
« 

E 

o 

\ 

0) 
X 

3 



T3 

0) 
X) 

O 



2 




Capella 






E-n 

1.5 


1 

c 

f 


1 1 

1991 April 15 

Phase = 0.26 

25.6 seconds 

G140L 


E-11 




C 1 


1 








E-11 


- 




- 






He II 
1 


III 


5 










E-12 


- 






. Id" 






C 1 


,1 .m\ 





.A.«-,.^V-A»^tr»MV^ 


1 1 



1425 



1500 



1575 
Wavelength 



1650 



1725 



Figure 3.: Same as Fig. 1 except for the 1425-1710 A region. This spectrum is domi- 
nated by the C IV resonance lines formed at 100,000K. Note the Hell 1640 A Hne and 
the intersystem III] line at 1666 A, which is a part of a density-sensitive multiplet. 
The underlying continuum is from the GO III star in the system. 



E 
o 



0) 



X 



0) 
t 

<n 
XI 
O 



1.5 

E-13 



1.125 
E-13 



7.5 
E-14 



3.75 
E-14 







y Draconis 



1 




-l I 1 

C 1 








1 + 






- 


c 


He II 
1 




- 


CO 


C IV 




III 


CO 


\kX 


w^ 


u*Wl 


wm 


A 



1450 1510 1570 1630 

Wavelength 



1690 



1750 



Figure 4.: Same as Fig. 2 except for the 1450-1740 A region. Unlike the Capella spectra, 
low- temperature emission hues dominate over the high-temperature C IV lines. There 
is no evidence for the photospheric continuum in this spectrum or out to 1840 A. 



74 



The data in Table 1 indicate that the surface fluxes of the high-temperature lines 
for Capella Ab lie about a factor of 3 below the saturated limit, and those for the 
shorter period V711 Tau system lie even closer to this limit. A natural explanation for 
this behavior is that a large fraction of the surface area of these active stars is covered 
by "plages" where the magnetic fields are strong and the heating rate is at or near 
its maximum possible value. Indeed, large plages have been identified on the surface 
of AR Lac by Doppler imaging techniques (NefF et al. 1989), and Linsky (1990) has 
shown that the surface fluxes in the plages of the RS CVn systems AR Lac, II Peg, and 
V711 Tau are near the "saturated" limit. 

On the other hand, the surface fluxes for the high temperatures lines for 7 Dra lie 
nearly a factor of 10,000 below the "saturated" limit. They are, in fact, the smallest 
surface fluxes ever measured on a cool star. Previously, the smallest values were the 
uncertain upper limits for a Boo listed in Table 1 obtained with lUE. One could in- 
terpret the very smaJl surface fluxes of high-temperature lines on 7 Dra as indicating 
that the fraction of the surface of such slowly rotating inactive stars covered by active 
regions is ~ 10""*, less than 10% of the plage coverage of the quiet Sun. This hypothe- 
sis is possible but is not easily tested observationally. More likely heating mechanisms 
are acoustic waves generated by the known convective motions in the photosphere, or 
perhaps magnetoacoustic waves if weak magnetic fields are present. Cuntz (1987) and 
Cuntz and Luttermoser (1990) have computed models of the K giant star a Boo in 
which a stochastic distribution of acoustic wave periods leads to the occasional coa- 
lescence of individual shocks into very strong shocks that produce high-temperature 
plasma. Our observations of 7 Dra and our proposed observations of a Boo and other 
stars wiU extend the measurement of surface fluxes to even smaller values to test these 
and other competing theories. 



3. MODERATE DISPERSION SPECTRA OF COOL GIANTS 

We now inspect the moderate dispersion spectra of Capella obtained through the 
LSA during our GTO program. These spectra have a nominal dispersion of 10,000 or 
30 km s~^. Figure 5 shows a spectrum containing the C IV resonance lines obtained 
with the G160M grating. These line profiles appear to be smooth Gaussians with no 
identifiable structure or splitting. Since the components of the Capella system have a 
radial velocity separation of 53.5 km s~ at phase 0.28, the absence of splitting or line 
asymmetry confirms that one star, the GO III star as determined in previous studies, 
contributes most of the flux. The absence of line structure indicates that no single 
bright plage was on the surface of the GO III star. This also is consistent with earlier 
studies, but more active RS CVn systems like AR Lac show enhanced discrete features, 
superimposed on otherwise smooth line proflles, which are thought to be produced by 
bright plages that are Doppler-shifted by stellar rotation. We will observe the C IV 
lines in AR Lac at many phases to map the location of bright plages regions using the 
Doppler imaging technique. 

The FWHM of the C IV 1548 A line is 217 km s'^ , while for the 1550 A line 
it is 186 km s~^ . These widths are much larger than the predicted thermal width, 
AX J) = 14.4 km s~^ , and the instrumental width of 30 km s~ , but are consistent with 
lUE observations at quadrature (Ayres 1984). The line flux ratio /1548//1550 = l"^"^ 
is significantly smaller than the ratio of gf values which is 2.0. These data indicate 



75 



« 

E 
o 

^^ 

en 

0) 

X 

3 



T3 
0) 

O 



3 


C 


apella 


E-11 

2.4 
E-11 


1 

c 


1 1 

IV UV1 

1991 April 15 
Phase = 0.28 
154 seconds 
G160M 


1.8 
E-11 


- 


1 


1.2 

E-11 


- 


- 


6 
E-12 




- Si II 

UV2 / 


V \ m 



1530 



1540 



1550 
Wavelength 



1560 



1570 



Figure 5.: A GHRS moderate dispersion spectrum of Capella obtained with the G160M 
grating. The C IV resonance lines are well- resolved in this spectrum. The profiles are 
smooth with no evidence for isolated plage regions on the surface of the GO III star. 



< 

in 
~\ 

CM 

« 
* 

E 

o 

~\ 

en 

i_ 

X 

3 



O 



1875 



5 




Capella 




E-11 

4 
E-11 


1 


^Si III ] 


1 
1991 April 15 
Phase = 0.28 
256 seconds 
G200M 






C III 


] 


3 

E-11 


r- 


S 1 UVl 


- 


2 
E-11 


- 






1 


1 
k 


1 
E-11 


- ^WM 


W 





1 1 1 


1 



1885 



1895 1905 

Wavelength 



1915 



1925 



Figure 6.: A GHRS moderate dispersion spectrum of Capella obtained with the G200M 
grating. The intersystem Hnes of Si III and C III are in emission superimposed on the 
photospheric absorption line spectrum. 



76 



that both turbulence and opacity broaden these lines, and that the more opaque line is 
optically thick. Such data will provide new constraints on acceptable model atmospheres 
for Capella and other stars. 



3.1 Density-sensitive Line Ratios 

Figure 6 shows the presence of the Si III] and C III] intersystem lines in a moderate 
dispersion G200M spectrum. The line fluxes can be measured without too much confu- 
sion above the photospheric absorption line spectrum of the star. Intersystem lines of 
III] at 1660 and 1660 A are shown in Figure 7, and the intersystem lines of IV] and 
S IV] are shown in Figure 8. To our knowledge, the S IV] have never been detected pre- 
viously in a stellar spectrum, except for the Sun, while the other intersystem lines have 
been detected by lUE in the spectra of several stars but with poor signal/noise. Clearly 
the GHRS can measure accurate fluxes for these faint lines. We note that FWHM = 
124 km s~^ for the Si III] line, while the predicted thermal width, AXj) = 6.0 km s~\ 
and instrumental width is 30 km s~ . The narrower width of this Line compared to the 
C IV 1550 A line is consistent with the Si III] line being turbulently broadened but with 
no opacity broadening, as is expected for intersystem lines that should be optically thin. 

Intersystem lines are important, because they provide independent measures of the 
electron density at the plasma temperatures where the ions are abundant. This can 
be seen by considering a simple three-level atom in which level 1 is the ground state, 
transition 1-3 is aJlowed, and transition 1-2 is an intersystem transition. For example, 
in Si III the 1-3 transition would be the 3s^ ^S — 3s3p ^P resonance Line at 1206 A, 
and the 1-2 transition would be the 3^^ ^S — 3s3p P intersystem line at 1892 A. The 
statistical equilibrium equations for this three-level atom are: 

ni [ne(7i2 + -Bl2-^12j = '"'2 [^eC'21 + ^21] 

ni [neCi3 -|- 513J13J = 713 [ueC^i + A31] , 

where tij is the population of level i, Cij is the colHsional rate for the i-j transition, J^j 
is the mean radiation field in the i — j transition line, and A^j and Bij are the Einstein 
A and B rates. Since the observed flux, fij oc rijAji, the flux ratio of the permitted to 
the intersystem line is, 

/31 _ C'i3/Ci2 



/21 [^ + 1] 



When the first term in the denominator becomes appreciable, i.e. when Ue > 
OAA21/C21, then collisional de-excitation of the upper state of the intersystem line 
is important and the flux ratio depends exphcitly on the electron density. At higher 
densities colUsional de-excitation from the upper state of the allowed transition (not 
included in the above equation) also becomes important, and the flux ratio is no longer 
sensitive to density. Table 2 summarizes the density range over which the prominent 
ions with ultraviolet intersystem lines are density sensitive. Figures 6-9 demonstrate 
that the GHRS can provide beautiful spectra containing these lines that can form 
the observational basis for accurate numerical models of stellar chromospheres and 
transition regions for late-type stars. An example is the model of (3 Dra computed by 
Brown et al. (1984) on the basis of earlier lUE spectra. 



77 



■\ 

CM 

* 

* 

E 

u 



en 

0) 



X 

3 



t 

en 
O 



1 5 






Cape 


la 






E-11 

1.2 
E-11 


He II 
1 


1 

UV12 


1 

C 1 
UV2 


1 

III ] 


1991 April 15 
Phase = 0.28 
102 seconds 
G160M 


9 


m 






E-12 








\\l 


1 ■ 


6 

E-12 


- 






1 




. 


3 
E-12 


m 


wk 


iiW 


iiUkll 


U- 





1 


iii]iYHi[iir " ■■ nr'i 1 
1 1 . 



1635 



1645 



1655 
Wavelength 



1665 



1675 



Figure 7.: A GHRS moderate dispersion spectrum of Capella obtained with the G160M 
grating. Note the intersystem lines of O III], which are density sensitive. 



Capella 




1380 



1390 



1400 
Wavelength 



1410 



1420 



Figure 8.: A GHRS moderate dispersion spectrum of Capella obtained with the G160M 
grating. Note the intersystem lines of IV] and S IV], which are density sensitive. 



78 



Table 2. Density-sensitive Line Ratios in the 1170-2350 A Region 



Ions 


logT 


Wavelengths (A) 


Range of log(Ne) 


CII 


4.0 


2323.5, 2324.7, 2325.4, 2326.9, 2328.1 


7-9 


Si III 


4.6 


1294.5-1303.3 (6 lines), 1892.0 


9-12 


III 


4.6 


1660.8, 1666.2 


9-13 


NIII 


4.8 


1746.8, 1748.6, 1749.7, 1752.2, 1754.0 


8-10 


SIV 


4.9 


1404.8, 1406.0, 1416.9 


10-13 


OIV 


5.1 


1397.2, 1399.8, 1401.2, 1404.8, 1407.4 


8-12 


V 


5.4 


1218.4 


10-13 



4. HIGH DISPERSION SPECTRA: THE INTERSTELLAR MEDIUM 
AND D/H RATIO FOR THE LINE OF SIGHT TOWARDS CAPELLA 

We discuss finally our beautiful echelle spectra of CapeUa obtained through the 
SSA, which have a measured spectral resolution (Wahlgren et al. 1991) of 87,000, cor- 
responding to 3.4 km s~ . Our objective in obtaining these spectra was to determine the 
D/H ratio and the physical properties of the interstellar medium along the 13 pc line of 
sight towards CapeUa. For this purpose Capella is a bright emission Line source against 
which we measure the opacity of resonance lines formed in the interstellar medium. 
These data are also useful for other purposes as we shall see. 

Figure 10 shows the spectrum of the Mg II h (2803 A) and k (2796 A) resonance 
lines obtained with the Ech-B grating. These spectra show the narrow interstellar ab- 
sorption lines, which are spectrally resolved and do not go to zero flux after correction 
for scattered light. The analysis of the line profiles provides information on both the 
line opacity and broadening. To the right of the interstellar lines one can see the self- 
reversal of the emission line from the G9 III star, and to the left one can see a portion 
of the self- reversal of the emission line from the GO III star. These features are barely 
present in lUE spectra at this phase. The shape of the composite emission line will be 
useful in testing chromospheric models of these stars. 

We show the Ech-A spectrum of the Lyman-a region in Figure 11. The broad 
stellar Lyman-a emission line is mutilated by the interstellar hydrogen Lyman-a ab- 
sorption feature and a narrow interstellar feature due to deuterium Lyman-a centered 
at -0.32 A relative to the hydrogen absorption line. The deuterium line has been seen 
in Copernicus and lUE spectra of Capella (e.g. Murthy et al. 1990) and other stars, 
but this spectrum is the first in which the line has been spectrally resolved. The small 
amount of instrumental scattered light can be measured from the minimum fiux seen 
in the saturated interstellar core of the hydrogen absorption line. The central depth of 
the deuterium feature is a measure of its optical depth, and the shape of the hydrogen 
absorption feature can be used to measure its opacity. A detailed analysis of this spec- 
trum, which is now under way, will provide a very accurate measurement of the D/H 
ratio along this line of sight. This wiU be important for inferring the primordial D/H 
ratio, which is a major constraint on models of the very early universe. 



79 



2.7 
E-11 



Capella 



in 

^ 1.8 
I E-11 

"\ 

I. 

u. 9 

-D E-12 
> 
a> 

M 

O 



1 

1991 April 15 

Phose = 0.28 

102 seconds 

G140M 


1 1 

1 
UV2 


1 


1 


- 


S 1 UV9 
II 


II ■" '■ 


1 


.,— ,^AK4u.«f.«** 


1 Si III UV4 


Si II UV3 





1285 1290 1295 1300 1305 1310 1315 

Wavelength 

Figure 9.: A GHRS moderate dispersion spectrum of Capella obtained with the G140M 
grating. The Si III lines are density sensitive. 



< 
~\ 

» 
« 

E 
o 

\ 



X 

3 



X» 

t 

o 



1.2 




Co 


pella 




E-10 


1991 April 15 

Phose = 0.28 

666 seconds 

ECH-B 


1 1 




9 


\ . Mg 11 UVl 


E-11 


" 




~ 


6 
E-11 


- /J 


Mg 


1 

1 UV3 1 


\ - 


3 

E-11 


■"-J 


V. 


/^-^^-Ny 


u 


0.0 


1 


1 1 





2791 2794 2797 2800 2803 2806 

Wavelength 

Figure 10.: A GHRS high-dispersion spectrum of Capella obtained with the Ech-B grat- 
ing. Each of the Mg II resonance lines (2796 and 2803 A) shows the narrow interstellar 
absorption line and self-reversed emission from the G5 III star (to the right) and the 
GO III star (to the left). 



80 



< 

en 

« 
« 

E 

o 

\ 

en 

(U 



1) 

t 

<U 

O 



4 
E-11 



3 
■11 



2 

E-11 



Capella 



1 
11 



0.0 



1 — 

1991 April 15 

Phose = 0.28 

3686 seconds 

ECH-A 



1212 1213 




1214 1215 1216 

Wavelength 



1217 



Figure 11.: A GHRS high-dispersion spectrum of Capella obtained with the Ech-A 
grating. Superimposed on the stellar Lyman-a emission line is interstellar absorption 
due to hydrogen and deuterium. 



X 

3 



T3 

O 

o 

(/I 



1.5 



USM Profiles 



0.5 




-45 



45 

Velocity (km/s) 



90 



Figure 12.: Comparison of the interstellar absorption lines of deuterium, Mg II, and 
Fe II on a common velocity scale. 



81 



Figure 12 compares the interstellar absorption lines of Mg II h, Fe II 2600 A, and 
deuterium on a common wavelength scale. Since the D and Fe ions differ by a factor 
of 28 in mass, the different line widths provide a means for separating thermal from 
turbulent broadening. There appears to be only one velocity component in this line of 
sight, but we are investigating whether the dip at the center of the deuterium line may 
indicate a second cooler velocity component. When the instrumental properties of the 
GHRS are better understood, we will publish what we hope will be a definitive value 
for the D/H ratio and interstellar properties for this line of sight. 

This work is supported by NASA Grant S-56500-D to the National Institute of 
Standards and Technology. We wish to thank Tom Ayres for his suggestions. 



REFERENCES 

Ayres, T. R. 1984, Ap. J., 284, 784. 

Ayres, T. R. 1986, Ap. J., 308, 246. 

Ayres, T. R. 1988, Ap. J., 331, 467. 

Ayres, T. R., Marstad, N. C, and Linsky, J. L. 1981, Ap. J., 247, 545. 

Ayres, T. R., and Linsky, J. L. 1980, Ap. J., 241, 279. 

Ayres, T. R., Moos, H. W., and Linsky, J. L. 1981, Ap. J. (Letters), 248, L137. 

Ayres, T. R., Simon, T., and Linsky, J. L. 1982, Ap. J., 263, 791. 

Batten, A. H., Hill, G., and Lu, W. 1991, Pub. A.S.P., 103, 613. 

Brown, A., Jordan, C., Stencel, R. E., Linsky, J. L., and Ayres, T. R. 1984, Ap. J., 

283, 731. 
Byrne, P. B., Doyle, J. G., Brown, A., Linsky, J. L., and Rodono, M. 1987 Astr. Ap., 

180, 172. 
Carpenter, K. G., Robinson, R. D., Wahlgren, G. M., Ake, T. B., Ebbets, D. C., Linsky, 

J. L., Brown, A., and Walter, F. M. 1991, Ap. J. (Letters), 377, L45. 
Cuntz, M. 1987, Astr. Ap. (Letters), 188, L5. 

Cuntz, M. and Luttermoser, D. G. 1990, Ap. J. (Letters), 353, L39. 
Drake, S. A., Brown, A., and Linsky, J. L. 1984, Ap. J., 284, 774. 
Fontenla, J. M., Avrett, E. H., and Loeser, R. 1991, Ap. J., 377, 712. 
Haisch, B. M., Linsky, J. L., Weinstein, A., and Shine, R. A., 1977, Ap. J., 214, 785. 
Hartmann, L., Dupree, A. K., and Raymond, J. C. 1980 Ap. J. (Letters), 236, L143. 
Linsky, J.L. 1990, in Active Close Binaries, ed. C. Ibanoglu (Dordrecht: Kluwer Academic), 

p. 747. 
Linsky, J. L., Worden, S. P., McClintock, W., and Robertson, R. M. 1979, Ap. J. Suppi, 

41, 47. 
Murthy, J., Henry, R. C., Moos, H. W., Vidal-Madjar, A., Linsky, J. L., and Gry, C. 

1990, Ap. J., 315, 675. 
Neff, J. E., Walter, F. M., Rodono, M., and Linsky, J. L. 1989, Astr. Ap., 215, 79. 
Simon, T., Linsky, J. L., and Stencel, R. E. 1982, Ap. J., 257, 225. 
Vilhu, O. 1987, in Cool Stars, Stellar Systems, and the Sun, ed. J.L. Linsky and R.E. 

Stencel (Berlin: Springer- Verlag), p. 110. 
Wahlgren, G. M., Leckrone, D. S., Shore, S. N., Lindler, D. J., Gilliland, R. L., and 

Ebbets, D. C. 1991, Ap. J. (Letters), 377, L41. 



82 



HIGH RESOLUTION UV SPECTROSCOPY OF THE CHEMICALLY PECULIAR B-STAR, CHI LUPI 



David S. Leckrone 
NASA, Goddard Space Flight Center 

Sveneric G. Johansson 
Department of Physics, University of Lund 

Glenn M. Wahlgren 
Astronomy Programs, Computer Sciences Corporation 

Abstract. Science assessment observations of the bright, ultra-sharp-lined B- 
peculiar star, chi Lupi, with the GHRS have provided an ultraviolet spectrum 
of unprecedented detail and photometric accuracy. The observed profile of the 
resonance line of Hg II at 1942 A confirms the reality and extreme nature of 
the Hg isotope anomaly in this star. In the surrounding 10 A spectral interval 
we observe for the first time lines of Ru II, As I, Ge II and Zr III. The data 
provide an ample demonstration of the inadequacies of the currently available 
atomic data base for the quantitative interpretation of high resolution 
ultraviolet spectra. 



1. INTRODUCTION 

In the preceeding papers Rolf Kudritzski and Jeff Linsky have treated, 
respectively, the hot, massive OB stars with their powerful winds and 
turbulent atmospheres, and the cool, late-type stars with their convective 
envelopes and dynamic chromospheres. In this discussion we are concerned with 
the "lukewarm" stars in the intermediate effective temperature range between 
about 8500 K and 15,000 K, whose stable photospheres we observe for the 
express purpose of deriving accurate elemental abundances. 

This interval of B and A spectral types constitutes an important "cut" through 
the HR diagram. It includes "normal" B and A dwarfs, whose main-sequence 
lifetimes are only a few hundred million years. Abundances derived for such 
stars presumably represent the composition of the interstellar medium, from 
which the stars formed, at a much more recent epoch than do solar abundances. 
Thus, they provide more suitable reference values, for example for studies of 
abundance depletion in the present interstellar gas, than do solar abundances. 
This temperature interval also includes a small number of relatively bright, 
highly evolved Population II field stars, which closely resemble blue 
horizontal branch stars in globular clusters. Elemental abundances determined 
for some species in these field horizontal branch (FHB) stars reflect the 
results of CNO processing and dredge-up on the red giant branch. For most 
elements, however, abundances provide a direct measurement of the composition 
of the interstellar medium at a very early epoch of galactic evolution. 

Of primary interest for this paper are the 10 to 20 % of B and A main-sequence 
stars that are classified as "chemically peculiar" (CP) stars. There are two, 
apparently unrelated, sequences of such stars - those which possess magnetic 
fields, and those which do not. The nature and origin of the spectroscopic 

83 



anomalies in these stars has been an enigma for nearly a century (e.g. Lockyer 
and Baxandall 1906). The high resolving power and photometric integrity of the 
HST/GHRS offers the opportunity for a major advancement in our understanding 
of these bizarre objects. 

The normal B and A stars, CP stars and FHB stars are apparently "well-behaved" 
subjects for analyses utilizing classical, LTE, plane-parallel model 
atmospheres and the associated spectral synthesis techniques. They are too hot 
to have convective photospheres and too cool to have significant winds, 
turbulence or mass loss. They emit sufficient ultraviolet flux to allow 
efficient UV spectroscopy down to Lyman-alpha. However, the derivation of 
meaningful elemental abundances for such stars may be complicated by processes 
of radiatively-driven diffusion, leading to chemical fractionation and an 
inhomogeneous radial distribution of the various ions present in their highly 
stable photospheres. It is this physical process, involving a competition 
between gravity and radiation pressure, that is currently the most widely 
accepted explanation for the peculiar abundances measured for the CP stars. It 
is a primary objective of our GTO program with the GHRS to critically test the 
quantitative predictions of diffusion models, as well as to evaluate 
alternative possibilities. 

Access to ultraviolet wavelengths is critical to this work. Their is a paucity 
of lines in the visible spectra of B and A dwarfs, and only a small sample of 
the periodic table is represented in studies which rely on ground-based 
spectroscopy alone. At UV wavelengths one can observe numerous intrinsically 
strong lines of low-abundance elements. Combining UV and optical-wavelength 
spectra gives one access to transitions from multiple ionization states and to 
resonance or low-excitation lines, which minimizes the possibility of large 
systematic errors due to departures from LTE. Finally, the rich UV absorption 
line spectra of these stars provides both a challenge and an opportunity for 
atomic physics. Absorption-line spectra of singly or doubly ionized elements 
are essentially impossible to observe in the laboratory, but are easily 
observed in the stars. We demonstrate in this paper that in trying to 
quantitatively interpret the strengths of the UV transitions we have observed, 
we are pushing contemporary knowledge of atomic structure to its limits. 

The following sections present the salient properties of our target, chi Lupi, 
and describe two scientific investigations - 1 . an attempt to independently 
confirm the reality and magnitude of the isotope anomaly in Hg, first detected 
in ground-based observations of a single line, Hg II X3984, and 2. a search 
for lines of elements whose abundances are unknown from optical-wavelength 
spectra. The former represents an important step prior to our extensive GTO 
program to thoroughly investigate the abundance and isotope anomalies in Hg, 
involving observations of strong resonance or low excitation UV transitions of 
Hg I, Hg II and Hg III in several stars. The latter begins the process of 
"filling in" the periodic table in order to systematically study the patterns 
of abundance anomalies from element to element as well as from star to star. 



84 



2 . PROPERTIES OF chi Lupi 

The target star for our observations, chi Lupi, is ideally suited to the 
purposes of the Science Assessment Program, that is to demonstrate the 
capabilities of the HST and GHRS in the presence of a severely aberrated point 
spread function. In addition to being bright (V = 3.9), so that integration 
times could be kept relatively short, chi Lupi possesses an exceedingly sharp- 
lined absorption spectrum containing a complex assortment of lines ranging 
from very weak to strong. Thus, it provides a good vehicle with which to 
assess the resolving power and effective S/N ratio of the GHRS. That chi Lupi 
is among the most sharp-lined of early-type stars results from its low 
projected rotational velocity, v sin i < 1.2 km/sec (Dworetsky and Vaughan 
1973) and a "classical" microturbulent velocity parameter = 0.0 km/sec 
(Adelman, et al. 1991, in preparation). Chi Lupi is a double-lined 
spectroscopic binary. The primary has Tgff = 10,650 K, log g = 3.8, while for 
the secondary Tg£f = 9,200 K, log g = 4.2. Near 1940 A, the wavelength region 
of interest here, the primary-to-secondary light ratio in the continuum is 
about 6.6. We see lines of the secondary in our observation, but they are 
generally very weak. 

Chi Lupi is one of the more extreme members of the non-magnetic sequence of 
chemically peculiar stars of the "HgMn" class. Both ground-based and lUE 
spectra indicate that Mercury is approximately 100,000 times overabundant in 
chi Lupi's photosphere, with respect to the solar-system value, although it 
must be noted that the abundance of Hg in the solar system is itself poorly 
known (e.g. Leckrone 1984). Platinum appears to be about 10,000 times 
overabundant (Dworetsky, et al. 1984). 

As mentioned previously, the shape and position of Hg II X3984 suggests that 
Hg in chi Lupi's photosphere is dominantly in the form of ^'-''^Hg, the heaviest 
isotope of Hg (White, et al. 1976). For comparison, only 7% of the Hg in the 
normal terrestrial isotope blend one finds in a thermometer is 204Hg_ 
Similarly, it appears that the heaviest isotopes of Pt are also overabundant 
in chi Lupi. The mercury anomalies are particularly important. Any physical 
model which seeks to explain the origin of the abundance anomalies in CP stars 
must be able simultaneously to reproduce a huge absolute overabundance of Hg 
and an extreme abundance distribution of the Hg isotopes in the line forming 
region of chi Lupi's atmosphere. Early attempts to create such theoretical 
models, based on radiatively-driven diffusion, may be found in Michaud, et al. 
(1974). To test these models, we must compare the strengths of lines from 
three ionization states, Hg I , II and III, in chi Lupi and in other Hg-rich 
stars of various effective temperatures. And to do that requires ultraviolet 
spectra of high quality. The present Science Assessment observations serve the 
additional purpose of providing an essential check on these seminal results, 
obtained from Hg II X3984, prior to the commitment of further HST time to the 
extensive study of the Hg anomaly. 



3. THE OBSERVED SPECTRUM 

Figure 1 illustrates a GHRS Echelle spectrum of chi Lupi, centered on the Hg 
II resonance line at 1942.3 A. These Science Assessment data were obtained on 

85 



February 11, 1991. The 0.25 arcsec small science aperture (SSA) of the 
instrument was used for the observation. This transmitted the central peak of 
the OTA point spread function into the spectrograph, while rejecting the broad 
PSF "skirt", which results from spherical aberration. Consequently, the 
resolving power of the GHRS anticipated prior to launch (X/6X == 87,000) is 
achieved in this mode, but with efficiency reduced by about a factor of four 
compared to pre-launch expectations. The total integration time was 2278 sec. 
To achieve proper sampling, the spectrum was quarter-stepped across the 
detector's diode array, so that one quarter of the total integration time was 
devoted to each of 2000 sample points in the spectrum. The S/N ratio per 
sample point near the continuum is approximately 100. The 1942.3 A Hg II line 
is plainly visible near the center of the displayed spectrum. Perhaps of even 
greater interest is the complex and remarkably detailed array of absorption 
lines seen in the surrounding 10.4 A interval. 



o 



3 
O 

u 

2 



20 
19| 

18 

17 

16 

15 

14 

13 

12 

1 1 

10 

9 

8 

7 

6 

5 

4 

3 

2 

1 









_L 



1936 1938 1940 1942 1944 

Wavelength (A) 



1946 



1948 



Figure 1. GHRS Small-Science-Aperture Echelle spectrum of chi Lupi, centered 
on the resonance line of Hg II at 1942.3 A. 



4. THE MERCURY ANOMALY 

In Figure 2 we have "zoomed in" on a one Angstrom segment of the spectrum, 
containing the Hg II line. The observed line profile has a well defined 
shape and appears to be relatively free of distortions due to blends. Figure 3 



86 



1.0 - 



^^. 



0.8 : 



CZ 0.6 

0) 

Ni 

"o 
E 
o 0.4 



0.2 



0.0 



T — I — I — I — I — I — r 



T — I — 1 — I — I — 1 — r — I — I — I — r- 




Mn II 
Fe II 



-J I I 1 I I L_ 



Hg II 



1941.8 1942.0 1942.2 1942.4 

Wavelength (A) 



1942.6 



1942.8 



Figure 2. Profile of the Hg II X1942.3 resonance line in chi Lupi, observed 
with the GHRS. 



0.0 



T — I — I — I — I — I — r — 1 — I — I — r- 




_J I — I — I — 1 — 1 — I — 1 1 — I I — I — L 



) Hg II Isotopic Lines 
. I ......... I 



1941.8 1942.0 1942.2 1942.4 

Wavelength (A) 



J- 



1942.6 



1942.8 



Figure 3. Theoretical profiles of Hg II X1942.3, calculated for various 
mixtures of Hg isotopes. Solid - q = 0.0 (solar mix), long dashes - q = 1.0, 
medium dashes - q = 2.0, short dashes - q = 3.0 (value estimated from Hg II 
X3984) . 



87 



shows our theoretical spectrum calculations for this same interval. The Hg II 
line is in fact a composite of eleven individual isotopic and hyperfine 
components of diverse strength, with central wavelengths ranging from 
1942.2240 to 1942.2994 A. For convenience we have chosen to calculate various 
isotope blends using a one-parameter model defined by White, et al. (1976). 
The logarithmic isotope mix parameter, q = 0.0 for the terrestrial blend of Hg 
isotopes. This corresponds to the solid curve in Figure 3. White, et al. 
estimated q = 3.0 for chi Lupi, based on their observations of Hg II X3984. 
This case, plotted with small dashes in Figure 3, corresponds to a mixture 
made up of about 99% ^'^^Hg, about 1% ^^^Hg, and tiny traces of the other 
isotopes. Two intermediate cases (q = 1.0 and 2.0) are also plotted in the 
figure. Our purpose in Figure 3 is to demonstrate that the shape, width and 
central wavelength of the Hg II X1942 profile are sensitive to variations in 
the relative abundances of the Hg isotopes. 

We can now simply superpose the observed and theoretical Hg II profiles, as 
shown in Figure 4. It is clear from this comparison that the mixture of 
isotopes in the line-forming region of chi Lupi ' s photosphere deviates 



-1 — I — I — I — I — 1 — I — I — r- 



—I — I — 1 — I — I — I — I — r 




O.U L_l I I I I 1 I I I I I I I I I— 



_l 



JLu. 



solid - observed 

dashes — q=0.0, solar mix 

dots - q = 3.0, 99% Hg(204) 

-I — I — I I I I — I I I I I — t — I — I — I — L 



1941.8 1942.0 1942.2 1942.4 

Wavelength (A) 



1942.6 



1942.8 



Figure 4. Comparison of observed and theoretical profiles of Hg II X1942.3. 



strongly from the solar-system isotope blend. The isotope anomaly is both real 
and extreme. The observed profile is reasonably well matched by the 
theoretical model with q = 3.0, but in fact is indistinguishable from the case 
of pure ^^'*Hg, to within the observational uncertainties. The one-parameter 
model used to describe the isotope mixture is somewhat arbitrary (although it 
has an empirical basis described by White, et al.). It must be emphasized. 



88 



however, that we are unable to define any alternative mixture of Hg isotopes 
that would lead to an equally good fit to the observed profile. 

The fit to the observed Hg II profile is not perfect. In particular, we are 
unable to compute a profile which is quite as narrow as the observed one. 
Also, there are significant departures from a good fit at the deepest part of 
the line core. It is possible that we are observing subtle evidence that 
mercury is not homogeneously distributed, but is concentrated in higher, 
cooler atmospheric layers. A more sophisticated, non-LTE model atmosphere and 
rigorous treatment of the radiative transfer problem in the high photosphere 
are also called for. Details of the analysis of the Hg II X1942 feature can be 
found in Leckrone, Wahlgren and Johansson (1991). 



5. "FILLING IN" THE PERIODIC TABLE 

The abundances of approximately eighteen chemical elements have been 
"reliably" determined from optical-wavelength spectra of chi Lupi ' s 
photosphere. For the present discussion we loosely define a "reliable" 
abundance as one that is based on at least two spectral lines that give more 
or less the same answer. In only six cases are abundances from ground-based 
data derived from lines of more than one ionization state of a particular 
element. Thus, in most cases it is difficult to assess the magnitude of 
systematic errors due to departures from LTE in the ionization equilibria. So 
we know relatively little about the patterns of elemental abundances in chi 
Lupi. We only know, on the basis of the study of a few elements, that its 
abundances are extremely anamalous in some cases. 

It has been a long-standing objective of Space Astrophysics to remedy this 
kind of problem by extending the observations to ultraviolet wavelengths where 
one can find many intrinsically strong lines of trace elements and of 
ionization states which are not well represented in the visible. Of course the 
instruments flown on Copernicus and on the lUE have allowed considerable 
progress to be made. But as one can see in Figure 5, with the HST and GHRS we 
have stepped into a new, largely unexplored spectroscopic universe. The 
comparison shown here between an lUE high resolution spectrum of chi Lupi and 
the GHRS Echelle observation is not intended to belittle the capabilities of 
the lUE. The latter observatory has been and will continue to be an immensely 
important tool for astrophysics. Instead, the comparison illustrates a new 
capability, not available before from any instrument. The S/N ratio ( = 15) of 
this single lUE observation could be improved perhaps to 40 or 50 by coadding 
multiple lUE images, obtained with the star properly offset in the lUE large 
aperture. However, the resolving power in the lUE observation cannot be 
improved beyond what is shown here. The GHRS observation of chi Lupi is, we 
believe, the most detailed ultraviolet spectrum of any star obtained to date, 
except perhaps for the Sun. 

In the 2 A interval shown in Figure 5 are several examples of elements or 
ionization states seen for the first time. These include Zr III, As I, Ru II, 
and Ge II. Lines of Zr II are observed in ground-based spectra of chi Lupi and 
other CP stars, but Zr II is the minority ionization state, sensitive to 
departures from LTE. We have identified three well-resolved and unblended 



89 



lines of Zr III, the majority ionization state, in the observed 10.4 A 
interval . 

We attribute the weak feature observed near 1937.6 A to As I X1937.594, on the 
basis of close wavelength coincidence and the lack of any other candidates at 
that wavelength. Our wavelength scale registration is accurate to 1-2 mA (see 
discussion in Leckrone, et al. 1991). Moreover, only a small number of 
features in the 10.4 A interval do not have solid identifications. We will 
have to search for lines of Arsenic at other UV wavelengths to be certain of 
this identification. However, if verified, we believe this is the first 
detection of Arsenic in any star, including the sun. 




Fe III 



n n r 



. I 



1937.1 



1937.6 



1938.1 
Wavelength (A) 



1938.6 



1939.1 



Figure 5. Comparison of lUE and GHRS Echelle observations of chi Lupi 



The observed feature near 1938.0 A is, we believe, about a 50-50 blend of Ge 
II X1938.007 and Fe III X1937.990. A companion Ge II line in this resonance 
multiplet, at 1938.890 A, is unresolved from the blend with Ni and Fe lines 
seen in Figure 5. Thus, we must also look elsewhere in the UV to confirm the 
identification of Ge II. A few Ge I lines have been identified in the solar 
spectrum. But the present observations, if confirmed, constitute the first 
detection of Ge II, and hold out the promise that we will be able to derive Ge 
abundances in the sharp-lined early-type stars. 

The detection of five well-resolved lines of Ru II in our 10.4 A interval is 
particularly exciting. Two of these lines are seen in Figure 5. Ru I is seen 
in the optical-wavelength spectrum of the sun and other late-type stars. 
Ruthenium is found to be about two orders of magnitude overabundant in the 
late-type, "heavy metal" or S stars, where it is a component of the s-process 



90 



neutron capture chain that leads to formation of the unstable element 
Technetium (see e.g. Wallerstein 1984). With the GHRS we now have the ability 
to measure the abundance of Ru in CP and other early-type stars. In chi Lupi 
Ru also is about two orders of magnitude overabundant, as we shall show in the 
following section. Given that models of the production of the abundance 
anomalies in CP stars based on nucleosynthesis are now very much out of favor, 
we are reluctant to suggest that the overabundance of Ru in chi Lupi results 
from recent s-processing, as it does in the S stars. However, prudence 
dictates that one should check the UV spectrum of chi Lupi for the lines of Tc 
II, and we plan to do so. 

We also see in Figure 5 a moderately strong line of Pt II near 1937.4 A. Our 
spectrum of chi Lupi contains many line of Pt I and II. This should not have 
come as a surprise, since the star's photosphere is overabundant in Pt by a 
factor of 10"*, and given the well-known richness of the Pt spectrum in the 
ultraviolet. After all, we use Pt lamps as wavelength calibration standards on 
GHRS, FOS, lUE and other space instruments. But it is nevertheless interesting 
to see in the absorption line spectrum of chi Lupi a mirror image of the 
emission line spectrum produced by the GHRS Wavecal lamp. 



6. THE ATOMIC DATA PROBLEM 

The line density in the spectrum illustrated here is obviously high. The 
dominant contributors are transitions from the second spectra of the iron 
group elements, V, Cr, Mn, Fe and Ni. To accurately synthesize the spectrum, 
using codes such as Kurucz's SYNTHE routine, one needs comprehensive and 
accurate atomic data - wavelengths, transition probabilities, and line 
broadening parameters for all transitions which make a noticeable contribution 
to the line opacity in the observed wavelength interval. Even the spectra of 
ions which are not of direct astrophysical interest are important, because 
their lines may be blended with those of other species of primary scientific 
interest. Their inclusion in the calculated synthetic spectra facilitates the 
accurate estimation of the level of the line-free continuum. Moreover, a large 
number of unidentified lines, produced for example by iron group ions, would 
add great confusion to the process of identifying lines produced by rarer and 
more interesting elements. 

The massive library of atomic data, calculated by Kurucz (1991) using the 
Cowan Code, provides the only reasonably comprehensive database with which one 
can begin to quantitatively interpret complex UV spectra. Although Kurucz has 
calculated data for over 50 million iron-group transitions, fewer than 2% of 
these involve atomic energy levels which have been accurately measured and 
classified, using laboratory spectra. It is this relatively small subset of 
transitions that have accurate enough wavelengths in the Kurucz database to be 
useful for computing synthetic spectra. Kurucz's library also includes 
compilations of transitions, calculated or measured, for elements both lighter 
and heavier than the iron group. 

We now face a dilemma, an extreme example of which is illustrated in Figure 6. 
In this plot of 1.2 A of chi Lupi's spectrum there is virtually no agreement 
between the observed spectrum and the theoretical spectrum, calculated with 

91 



the Kurucz atomic data base. This is a problem both with the completeness of 
the atomic database and with our knowledge of atomic structure, level mixing 
and configuration interactions. 






O 

E 



1.2 : 



1.0 7^ 



0.8 r 



0.5 



0.4 - 



0.2 



0.0 

1938.8 1939.0 1939.2 1939.4 1939.6 1939.8 1940.0 

Wavelength (A) 

Figure 6. First attempt to theoretically synthesize a 1.2 A spectral interval 
of chi Lupi, using atomic data from most recent Kurucz calculations. 



7*" 


''■''''' 1 1 


'*''••■ 


'1 1 


I 1 I I 1 I ' 


, . 1 . i 





1 1 


; 


solid - GHRS observotion 












; 


- 


dashes - theory 












: 


T 




/Cr 


^~\r 


*>vv 


V'-^ 




-; 


' 


1 1 N 


\i 


\ 


r 


\r\f 


1 


\ " 


}[ 


- 


I 


ll n 


\ 




Cr II 


\ 1 








'] 


V : 


- 


u / / 




Al 


II 












U' 


1 - 


- 


1 ""^ '" W 














I 


|- 


- 






Fe II I 




1 




V 


t 


~ 


/ 






yreii 


1 

1 




Ij 


\ 


1 1 


e III " 










[i 








; 


'- 


1 I 


Cr II 














Cr II 


- 


\ 


1 1 

1 1 Ri. 
1 1 (^^ 


S) 




Ru 
(uv 


h 










1 


~__ 


1 • 






Fe 11 


\/ 




_^ 


\ 








v 




; 


- 


1 

Ge II 








pt 11 




j 


"i 


'■'''''■ 1 1 


1 1 1 1 t 1 1 


' 1 ' • ' ' 1 


t 1 1 1 ■ 1 





1 1 1 1 1 


1 1~ 



Kurucz has not yet included data for Ru or Pt, so that no calculated features 
appear at the wavelengths of these lines. He did include a "guess" for the 
transition probability of the Ge II line, apparently based on relative 
laboratory line intensities. This estimated gf-value is obviously much too 
large. 



We see an extremely anomalous calculated Fe II feature near 1939.7 A. The 
transition is b ^Pi/2 ~ w ^^212' Whenever we encounter Fe II transitions which 
involve this upper level, whether in analysing GHRS or lUE data, the Kurucz 
gf-values seem to be badly in error. We believe this results from the 
difficulties in accurately treating the mixing of closely coincident atomic 
levels in the Cowan Code calculations. In this case the Kurucz calculations 
produce a short-range perturbation between w ^^212 ^"^ ^ ^^3/2 which is not 
verified by laboratory line intensities. Less extreme examples of the same 
problem are seen in the two Cr II lines, XX 1939.149, 1939.902, one of which 
is calculated much too weak and the other somewhat too strong. 



There are also "simple" problems of wavelength accuracy. For example, the 
wavelength of the calculated Fe III line at 1939.105 A should now be increased 
by 10 mA, based on new measurements of the Fe III spectrum by J. Ekberg at the 



92 



1.2 



1.0 



X) 



o 
E 



0.8 r 



0.6 r 



0.4 



— T—I 


1 r I 1 


T-r- 


' 1 


T T y 


' I 


r-T-TT' 


I I I I 


. , . . . 


1 I I I 1 


' 1 ' 




1 T ..... . 


• ■_ 


- 


solid 


- 


GHRS observat 


on 
















: 


- 


dosh 


es 


- theory 




















: 


-_ 


























~ 


- 




















yV 


<%« 




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f^ 


\j 


i 


n^ 




/y 




/^ 


K^ 


V / 


A^ 


^ 




1 


L 




1 


1 

1 


\' 


/ 


V 






\ 






\ 


/^ 


\1 


: 






1 
1 












1 






\ 


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V 


i 








1 
1 


















\ 


IV 1 

1 1 


I: 


r 


u 






1 
1 


v 










v' 




Fe II ', 




\- 


: 






w' 












"d 






1 


: 


I 








% 






















; 


r 


Fe III 




























- 


: 


Ni II 






























: 




















Ru 


II 










: 




1 


Ru II 
ob wavelen 


gth 








lob 


wavelength 




1 


: 


: 
























pt II 


: 






























-, 


' ' ■ ■ ■ 




1 1 


i 1 1 


1 1 






1 n . . 





1 1 1 


1 1 1 1 1 1 


1 1 1 1 1 1 1 1 





0.2 



0.0 

1938.8 1939.0 1939.2 1939.4 1939.6 1939.8 1940.0 

Wavelength (A) 

Figure 7. Improved spectrum synthesis, using published wavelengths for lines 
of Ru II. 






O 

E 



1.2 


solid 


1 < . , 1 . 

- GHRS 


' ' 1 

Dbservotion 


' ' 


f T ' T 1 r- 


' ■ 1 


1 ' ' ' 


7 T' f I I I 


I ' ' 


T T n~ 


r T » r 




: dosh 


es - theory 
















l 


1.0 


n 


A 


^^,f 




y^ 


^r 


f^ 


-^ 






1 


0.8 


" I 


n 


\ 






/ 




\ 
\ 
1 




IM 


' v 


0.6 


~ 1 / 










II 




\ 
\ 
1 




1 V 


\ 


0.4 


j- 






















- 




" 


Ru 
shifted - 


II 

-0,016 A 






Ru II 
shifted -0.015 


A 








1 


0.2 


r 
















V 




— 




: [ 


N(Ru)/N(^ 


H)] = +1,93 


dex 


obove 


solar obundonce 










E 


0.0 











1 ,1 1..J— L 


. . 1 


llll 




1 1 


LX l.-L- 


rrr 



1938.8 1939.0 1939.2 1939.4 1939.6 1939.8 

Wavelength (A) 



1940.0 



Figure 8. "Final" spectrum synthesis with calculated Ru II lines arbitrarily 
shifted by -0.015 A. Ruthenium is about 2 dex more abundant in chi Lupi's 
atmosphere than in the sun. 



93 



University of Lund. Such an apparently minor adjustment makes the difference 
between a clear discrepancy, and reasonable agreement, between theory and 
observations at the position of the Fe III line. 

In Figure 7 we illustrate the first effort to resolve some of these problems. 
We obtained improved estimates for the transition probability of the Ge II 
X1938.890 line, using calculated values for Si II, which is homologous with Ge 
II and checking these with measured values for other lines in the same 
multiplet in Ga I, which is isoelectronic with Ge II. These two approaches 
yield consistent values of log gf = -3.6, instead of the value, 0.0, "guessed" 
in the Kurucz database. We added Pt lines to the database, with transition 
probabilities from Cowan-Code calculations for the classified transitions. As 
shown in Figure 7, the result is slightly too strong for the Pt II line near 
1939.8 A. We simply eliminated the super-anomalous Fe II line at 1939.698 A 
from the calculation. And we shifted Fe III X1939.1 to its refined laboratory 
wavelength. 

Of special note in Figure 7 are the two lines of Ru II (UV multiplet 5). We 
estimated transition probabilities for these lines using the Cowan Code, and 
in fact these calculations converged to the measured log gf-values for the 
corresponding transitions of Fe II, which is homologous with Ru II. As shown 
in Figure 7, the calculated relative line strengths are approximately correct 
and the relative spacing of the two lines in wavelength corresponds to that 
observed in the stellar spectrum. However, the absolute laboratory wavelengths 
(which come from Shenstone and Meggers 1958) differ from those observed in the 
stellar spectrum by 16 mA. In fact all five well-resolved Ru II lines observed 
in our 10.4 A interval are shifted from their published laboratory wavelengths 
by -0.016 A. We do not know whether this discrepancy reflects a calibration 
problem in the laboratory measurements or an astrophysical phenomenon. New 
measurements of the Ru II spectrum will soon be made at Lund, using the new 
Fourier Transform Spectrometer there. For purposes of estimating the Ruthenium 
abundance in chi Lupi, however, we have simply shifted all of the calculated 
Ru II lines by -0.016 A. The result is illustrated in Figure 8. On this basis 
we conclude that Ru is overabundant by about 2 dex in chi Lupi relative to the 
solar value. 



7. CONCLUDING REMARK 

We cannot say at this moment whether the discovery of a particular, previously 
unobserved species, such as Ruthenium, Arsenic or Germanium, will ultimately 
provide the decisive clue about the origin of the truly bizarre abundance 
anomalies observed in stars such as chi Lupi. These new GHRS data do 
illustrate vividly, however, that with superb resolution and S/N ratios, we 
are beginning to realize the full potential of UV space spectroscopy as 
originally envisioned by its pioneers decades ago. 



94 



REFERENCES 

Dworetsky, M.M. , & Vaughan, A.H., Jr. 1973, ApJ, 181, 811. 

Dworetsky, M.M., Storey, P.J., & Jacobs, J.M. 1984, Phys. Scripta, T8, 39. 

Leckrone, D.S. 1984, ApJ, 286, 725. 

Leckrone, D.S., Wahlgren, G.M., & Johansson, Se. G. 1991, ApJ, 377, L37. 

Lockyer, N., & Baxandall, F.E. 1906, Proc. Roy. Soc. (London), 77, 550. 

Michaud, G., Reeves, H., & Charland, Y. 1974, A & A, 37, 313. 

Shenstone, A.G., & Meggers, W.F. 1958, J. Research Nat. Bur. Std. , 61, 373. 

Wallerstein, G. 1984, J. Opt. Soc. Am. B, 1, 307. 

White, R.E., Vaughan, A.H., Jr., Preston, G.W. , & Swings, J. P. 1976, ApJ, 204, 

131. 



95 



Hubble Space Telescope Optical Performance 

Christopher J. Burrows' 

Space Telescope Science Institute, 3700 San Martin Drive 
Baltimore, Maryland 21218, USA 

ABSTRACT 

The Hubble Space Telescope suffers from spherical aberration. Although much that is 
scientifically valuable can be done with the telescope in its present condition, we must 
install corrective optics on-orbit in order to enable many key programs. An analysis of 
imaging data obtained on-orbit gives the same results as measurements on the null lens 
used to fabricate and test the primary mirror, so such optics can be designed with 
confidence. The assumed conic constant on the primary mirror for all the corrective 
optics, -1.0139(5), is consistent with measurements by four major independent 
methods. Aligning the new optics will be very demanding, because of the large slope of 
the wavefront to be corrected. If the images are to be diffraction limited, the pupil at 
the corrective element must be aligned to better than one percent of its diameter. Some 
other residual effects of the spherical aberration will remain after installation of the 
corrective optics, primarily in the pointing and collimation of the telescope. We 
summarize the present imaging performance of the observatory, and compare it with 
the expected performance when corrective optics (COSTAR and WFPC 2) are installed 
on-orbit. 

2, EVTRODUCTION 

This is a general review of the status of and prospects for the Hubble Space Telescope 
(HST) optics, and the current and projected imaging performance of the observatory, 
written one year into the mission. Despite the spherical aberration, HST is now 
routinely producing high resolution images and spectra. Nevertheless its performance is 
seriously degraded relative to expectations. It has almost completed an extensive 
commissioning period, made more difficult by the spherical aberration, so now is a 
good time to review what we have learned, and where we are going. We concentrate 
here on the optical performance of the various observatory subsystems. The scientific 
results have been discussed extensively elsewhere, '■- and in these proceedings. The 
present optical imaging performance, and pointing control system performance are 
generally well understood. ^ We choose not to repeat information given in (3). Instead, 
we concentrate on giving the reader some additional background on the imaging 
performance and how it can be modelled. We concentrate here on understanding 



'Affiliated with the Astrophysics Division, Space Science Department of the European Space Agency. 

96 



physically the current imaging performance and compare it with the expected 
performance which should be substantially attained after the planned servicing mission. 

The HST will be serviced on orbit with a shuttle visit presently scheduled for 
November 1993. At that time, the bulk of the problems caused by the spherical 
aberration will be fixed. That mission will involve the most ambitious on orbit satellite 
repair ever undertaken. It will probably include replacement of the solar panels, gyros, 
the High Speed Photometer and the Wide Field/ Planetary camera (WFPC). The solar 
panels are affecting the pointing performance primarily during terminator transits 
because thermal gradients across the bistem array supports are causing them to rapidly 
bend and then oscillate. There are also impulsive torques on the spacecraft at other 
times due to stiction in the tensioning mechanisms. One of six rate gyro assemblies 
failed after nine months on orbit. A second gyro has failed, and a third has shown 
evidence of possible failure since the workshop. The Faint Object Camera (FOC), 
Faint Object Spectrograph and High Resolution Spectrograph will probably be largely 
corrected by a replacement for the High Speed Photometer, COSTAR. COSTAR is 
designed to deploy corrective optics in front of the scientific instrument apertures.'* 
Finally, but most critically, the Wide Field Planetary Camera (WFPC) which is the 
main scientific instrument of the observatory will be replaced by a similar instrument 
with internal correction. 

3. MEASUREMENT OF SPHERICAL ABERRATION 

In order for the corrective optical schemes to work, the spherical aberration must be 
precisely characterized. None of the schemes presently contemplated allows for on- 
orbit adjustment of the amount of compensating spherical aberration that they 
introduce. Measuring the amount of spherical aberration to the precision required has 
been a major challenge. It has been determined from historical optical test data, from 
recent measurements on the test equipment and from on orbit measurements with both 
the main cameras. Remarkably, these methods all agree to within about 1/45 wave, 
which is sufficient for the corrective optics. The loss of encircled energy in a 0. 1 
arcsecond radius resulting from such an error is only about 2%. 

The on-orbit determination of the spherical aberration has evolved from low order 
phase aberration least squares fitting procedures (which were used to originally 
diagnose the spherical aberration), to involved procedures involving simultaneous 
solutions to the pupil phase map and pupil obscurations at three different focal settings 
as reported by Roddier in these proceedings. Many groups have developed such 
techniques but we only discuss our results here. The wavefront fitting procedure was 
first developed and applied within days of the WFPC first light images in May 1990. 
Since then it has evolved considerably by including the secondary mirror support 
spiders, primary mirror support pads and field dependent WFPC internal obscurations 
(which are misaligned in an unexpected but determined manner). The algorithm works 
best on well sampled out of focus images at long wavelengths with good pointing 
stability. Such images have been obtained in both the PC and FOC. Given the pupil 

97 



obscurations^, the image is determined in the Fraunhofer approximation once the 
distribution of phase errors on the wavefront in the exit pupil is fixed. These errors can 
be expanded in a series of orthogonal polynomials, and the coefficients deduced by a 
least squares fit to the observed data at various focus positions. The rms wavefront 
error expressed in waves at 632.8 nm is converted to a change in the conic constant on 
the primary mirror by dividing by 35.3. (This factor is correct rather than the 36.03 
one naively deduces by doubling the sag on the mirror primarily because the reflected 
ray does not exactly retrace its path). Our result is that the conic constant on the 
primary mirror is 1.01410(45), from fifteen measurements in the planetary camera 
(PC6), after subtracting a correction of 0.0010 due to a manufacturing error in the 
camera. The FOC gives 1.01394(85) from four measurements and no significant 
camera aberrations are believed to contribute. Figure 1 shows a comparison of an 
observed star image in PC6 with filter F547M at a focus setting close to the telescope 
paraxial focus, with the results from phase retrieval. These same techniques are now 
being applied to the telescope collimation, and should be applied to generate theoretical 
point spread functions at the adopted focus position for use in decon volution. 




Figure 1 Comparison of fits to data for defocussed images 



The on-ground determination at Hughes, Danbury (formerly Perkin Elmer) has 
proceeded by recertifying the null lens used in their manufacture of the primary 
mirror. The combination of null lens and primary mirror was designed to reflect a 
spherical wave exactly back on itself as illustrated in Figure 2. The autocollimated 
wavefront from the assembly as built was measured with an interferometer, and 
changes were made to the figure of the primary mirror in order to get straight fringes. 
As a result errors in the construction of the null caused a corresponding error in the 
primary, and recertifying the null now comes close to a direct measurement of the 
primary mirror figure. The null consists of two spherical mirrors, which reimage the 
interferometer focus onto the center of curvature of the HST primary, and introduce 



98 



most of the required optical path changes, together with a refractive field lens that 
images the interferometer pupil onto the primary mirror 




To HST primary 
mirror under test 



1^ 



Field Lens 



Interferometer 
focus 



Upper null mirror 



Lower null mirror 



Figure 2. The null lens used in the manufacture of the HST primary mirror 

The field lens positioning error primarily responsible for the spherical aberration in 
HST has now been measured as 1.305 mm and the cause of the error is believed to be 
understood^. In addition the two mirrors in the null lens have been remeasured as 
separated by an amount that differs from the design by 79 microns. This is also a 
significant error, and has recently been explained by measurements of the null mirror 
radii, which are in error by about the correct amount to explain the discrepancy, given 
the way they were positioned. The original spacing was set from the centers of 
curvature, while the recent measurement was from the surface vertices. After the 
remeasurement of the mirror radii, the reflective null lens data indicate that the conic 
constant on the primary is 1.01378(31). 

Finally, a less precise (but more accurate!) refractive null lens was made that was used 
in place of the reflective null early in the testing. Archival interferograms from the 
period when the mirror was polished with this null indicate a conic constant of - 
1.01314(60). There is reason to believe that the secondary mirror was made correctly. 

Figure 3 summarizes the status of all the above primary mirror conic constant 
estimates. All the measurements agree quite closely with the adopted value of 
1.0139(5) which is being used in the design of the corrective optics. The nominal conic 
constant for the primary is -1.0022985. The difference between these values 
corresponds to an RMS wavefront error of -0.410 waves at 633 nm, and the edge of 
the primary is 2.2 microns too low. (which can be compared to its nominal sag relative 
to a sphere of 0.2mm) 



99 



1.015 



1.0145 



1.014 



1.0135 



1.013 



1.0125 

Primary Mirror 
Conic Constant 



[Error in 
camera 



N ull 
M e tro lo g V 



R e f ra c tiv e 
N ull 



P la n e ta ry 
Camera 



Fa in t O bje c t 
C a m e ra 



Adopted 
Value 



z(ll) = -0.256(7) 
K=-1.01378(31) 



zdl) = -0.242(16) 
K = -1.01314(60) 



z(ll) = -0.264 (10) 
K = -1.01410(4S) 
dz(ll) =-0.022 



z(ll) = -0.260(19) 
K = -1.01394(85) 



z(ll) = -0.255(14) 
K = -1.0139(5) 



Figure 3. Estimates of the HST primary mirror conic constant. 



4. CORRECTIVE OPTICS ERRORS 

Because of the accuracy of our knowledge of the spherical aberration, on-orbit 
compensation for errors in the prescription is not planned. On the other hand, there are 
extremely tight positional tolerances for the corrective optics which will necessitate on- 
orbit centering adjustments. The reason for this is as follows. All the corrective 
schemes being pursued involve reimaging the primary mirror on to a corrective element 
which has exactly the opposite deformation to the primary mirror error. Thus the 
optical path delays at the edge of the pupil introduced by the primary mirror that 
presently cause marginal rays to focus 40 mm too far back are cancelled by 
compensating advances introduced by the corrective optic. The compensating term 
varies as the fourth power of the pupil radial position, and therefore changes rapidly 
near the edges. Slight misalignments then lead to large amounts of wavefront 
perturbation proportional to the derivative of the aberration (coma), and to the amount 
of misalignment. If the system is misaligned by as much as one percent, it will cease 
to be diffraction limited. A seven percent misalignment (which is about equal to the 
misalignments in the existing WFPC) would lead to as much RMS coma as we 
presently have in spherical aberration. Of course, the misalignments in the existing 
camera would not have significantly damaged the image quality, if the OTA wavefront 
had been nominal. Similarly, the scale of the pupil image must be right. If it is in error 
by 2 percent, we again fail to be diffraction limited. These issues are illustrated in 
Figure 4. 



100 




Error* «hoMfn xlO 




Errors ihown x^0 



One percent diameter 
misalignment of the image 
of the primary mirror on 
the corrective optic gives 
7/100 waves rms of coma. 



Two percent diameter 
error of the image of the 
primary mirror on the 
corrective optic gives 7/100 
waves rms of spherical. 



Figure 4. Sensitivity of corrective optics to misalignments and scale errors. 



5. GUIDING PERFORMANCE 

The servicing mission will probably not include changeout of the three Fine Guidance 
Sensors (FGS). These detectors are affected by the spherical aberration both in their 
coarse track and fine lock modes but for different reasons. Each FGS consists of a 
large pickoff mirror that takes a quadrant of the field between 10 and 14 arcminutes off 
axis and directs the beam onto an aspheric collimating element. The beam is then 
steered by two sets of moveable prisms onto a beamsplitter and thence to one of two 
orthogonal Koesters prism interferometers. The faces of the Koesters prisms are 
reimaged onto photomultiplier tubes with square 5 arcsecond field stops. 

In coarse track the field of the four PMT in a given FGS is nutated around the star 
image, so that the system works somewhat like a quadrant detector. Performance is 
degraded because the star image is aberrated, so the image on the field stops is not 
sharp. The result is that the pointing performance is degraded by about 2 magnitudes. It 
still scales at the square root of the number of photons, but the rms jitter expected for 
14.5 magnitude guidestars of 20 milliarcseconds is only realized for 12.5 magnitude 
guidestars. 

In fine lock the performance is limited by the transfer function of the interferometers. 
Because of the presence of misalignments in the FGS, the beam is not exactly centered 
on the Koesters prism at interferometer null. This would not affect performance 
appreciably in the absence of spherical aberration from the OTA, but the fringe 
visibility is drastically reduced by the position dependent phase errors that result. 

Operational changes to planning, scheduling and pointing control system control loop 
gains are being made to mitigate the degraded performance of the FGS, and solar array 



101 



induced jitter. For bright stars, we have excellent pointing stability in fine lock away 
from terminator transitions. The changes are making such performance possible on 
fainter (and therefore more common) guide stars, and avoiding loss of lock at the 
terminator. However, it is not expected that acceptable fine lock performance will be 
achieved on stars with visual magnitudes fainter than about 13.5, so coarse track (with 
its associated larger pointing errors) will continue to be the only available guiding mode 
over much of the sky. 

6. TELESCOPE COLLIMATION 

The Optical Control System (OCS), contains a set of three radial shearing 
interferometers (WFS), one in each FGS. They were designed to collimate the 
telescope (ensure that the optical axes of the hyperbolic primary and secondary mirrors 
coincide). They do not work because of the spherical aberration. There is always some 
zone on the primary mirror with such a rapid rate of change of the wavefront that the 
fringe separation in the instrument is smaller than the instantaneous field of view of the 
image dissector tubes, so the fringe visibility drops. In principle, they therefore can 
only accurately measure aberrations with significant angular dependence (such as 
astigmatism), but are poor when measuring aberrations that only have radial 
dependences (such as focus and spherical aberration), or that have angular dependences 
that look like wavefront tilt over any narrow annulus (such as coma). As a result the 
present collimation of the telescope was set by tilting the secondary mirror until the 
camera images were symmetrical and decentering it (with compensating tilts) until the 
OCS indicated zero astigmatism. 

Hughes Danbury and Pierre Bely at STScI have independently modelled the FGS in 
fine lock, and are able to predict the effects of differing telescope collimation on the 
FGS transfer functions. The results to date indicate that the performance of the FGS 
can be significantly improved at least for FGS 1 and 3, but at the cost of the images in 
the WFPC and FOC. Now that the effects of collimation on FGS performance are 
better understood, and the effort to numerically understand the spherical aberration is 
completed, the collimation will be revisited with a view to improving the FGS 
performance while compromising the camera images as little as possible. To this end, a 
systematic study of the phase retrieval results at numerous secondary positions is 
underway, together with a sequence of dedicated HST observations to perform a 'coma 
sweep'. So far, results from both cameras seem consistent, and it is believed that the 
telescope has been operating with about 1/15 wave of coma. Improved knowledge of 
the collimation can be used to generate better theoretical point spread functions. 

6. IMAGING PERFORMANCE 

The Science Working Group (SWG) has defined the focus setting of the secondary 
mirror to be used for HST observations by requiring that it gives the maximum 
encircled energy in O.I arcseconds radius for the FOC at 486 nm. Figure 5 shows a 
series of theoretical encircled energy curves as a function of focus setting, assuming the 
nominal conic constant. In practice the encircled energy is lower than these curves 

102 



indicate primarily because of microroughness scatter from the mirror surfaces, but the 
shape of the curves seems to agree with observations. 



Enerrclad Energy 4^6.000 0.4B7043 




Peak value of 
0*t75 hereon 
lencircled aiergy 
in 0.1 arcsec radius 



Paraxial focus Adopted focus @ 13.5mm 



Marginal focus @41.2mm 



Figure 5. Predicted Encircled energy for the FOC as a function of focus position. 

The peaks in the curves shift towards the paraxial focus as the wavelength is 
decreased. In the limit of short wavelengths, (neglecting microroughness), we get the 
ray trace limit which can be written in closed form, and is an adequate approximation 
for some purposes. The transverse aberration of a ray at fractional pupil coordinate 

;c(with 0.33 <x< 1) is then T^(x^-xd), where T^ is the transverse spherical aberration, 

of a marginal ray in the paraxial focal plane (about 3.07 arcseconds), and d is the 
defocus as a proportion of the longitudinal spherical aberration - d is 0.327 at the SWG 
focus. At this focus, the geometric image radius is obtained by setting x=\ and comes 
out as 2.07 arcseconds. Further it is easy to see that the zone of the mirror that 
geometrically is in focus is at jc= 0.572. From the fact that the effective aperture 
contributing to the core is roughly 2/3 the expected size, one might predict that the core 
would be about 50% broader than expected, and detailed diffraction calculations 
confirm this intuition. Specifically, the FWHM at 633 nm is 75 milliarcseconds, 
instead of 53 milliarcseconds, and it is proportional to wavelength to a reasonable 
approximation. (At longer wavelengths it oscillates somewhat about the linear trend due 
to interference phenomena, remaining constant between 850 and 950 nm for example.) 

The WFPC differs slightly from the FOC in that the reimaging cameras contain 
secondary mirrors that obscure part of the center of the beam, and also contain about 
1/30 wave of extra spherical aberration. The result is that the encircled energy curve 
for a 0. 1 arcseconds radius aperture is shifted by 3.7 mm further from the paraxial 
focus, and the encircled energy at the SWG focus is about 12% instead of a possible 
peak value of 17%. The central obscuration in the cameras moves in a predictable way 
from ray trace calculations, although it is not centered at the center of each chip as one 



103 



might expect from the optical design. The result is that star images change in a 
predictable way as one changes field position. An example is shown in Figure 6 which 
compares predicted and observed PSFs at several field positions in PC6 through filter 
F547M. The theoretical predictions are parameter free being based entirely on the 
known pupil function of the camera, the measured spherical aberration of the OTA and 
the SWG focus definition. Obviously one could improve the agreement further by 
introducing known misalignments in the telescope (primarily coma), and modelling the 
measured phase errors on the primary and secondary mirrors. This work is ongoing, 
and hopefully will lead to better simulated PSFs for deconvolution purposes. In the 
meantime, a library of PSFs with the above nominal parameters has been produced by 
at the STScI, using the TIM software. Both the library and the software are available to 
interested researchers on request. 




Figure 6 Comparison of predicted and observed PSFs at several field positions in PC6 

104 



7. REFERENCES 

1. Astrophysical Journal (Letters) 369 (1991) (all papers) 

2. Bulletin of the American Astronomical Society Vol 22 pp 1275-1284 (1990) 
(Abstracts to session 49 of the January 1991 A AS meeting) 

3. Burrows, C.J. et al. "The imaging performance of the Hubble Space Telescope" 
Astrophysical Journal (Letters) 369, 21 (1991) 

4 Brown, R. A. and Ford, H.C (Editors) " Report of the HST strategy panel" , 
Space Telescope Science Institute special publication (1991) 

5. Allen, L. et al. "The Hubble Space Telescope optical systems failure report" 
NASA publication (November 1990) 



105 



INTRODUCTION TO THE GODDARD HIGH RESOLUTION SPECTROGRAPH (GHRS) 

John C. Brandt 

Laboratory for Atmospheric and Space Physics 
University of Colorado 
Boulder, CO 80309-0392 
USA 

The principal presentation on the Goddard High Resolution Spectrograph (GHRS, nee HRS) and its status will 
be made by Dr. Dennis Ebbets in the following paper. This introduction presents the GHRS from the perspective of 
the Preliminary Design Review (PDR), which began on December 12, 1978. Some archival research has uncovered 
my notes and viewgraphs^ from the "kickofF' talk for that meeting. Here, I review the materials presented at the 
PDR and comment where appropriate. 

First, GHRS is very much a team effort requiring the cooperation of scientists, engineers, technicians, 
programmers, and support personnel. Because the scientific investigations to be carried out are from a variety of 
astronomical disciplines, the Science Team is large. It originally consisted of twelve members and now numbers 
sixteen; see Table 1. 



Table 1 
The GHRS Investigation Definition Team 

John Brandt (PI) - University of Colorado, Boulder 

Sara Heap (Co-PI) - Goddard Space Flight Center 

Edward Beaver - University of California, San Diego 

Albert Boggess - Goddard Space Flight Center 

Kenneth Carpenter - Goddard Space Flight Center 

Dennis Ebbets - Ball Aerospace 

John Hutchings - Dominion Astrophysical Observatory 

Michael Jura - University of California, Los Angeles 

David Leckrone - Goddard Space Flight Center 

Jeffrey Linsky - Joint Institute for Laboratory Astrophysics 

Stephen Maran - Goddard Space Flight Center 

Blair Savage - University of Wisconsin, Madison 

Andrew Smith - Goddard Space Flight Center 

Laurence Traflon - University of Texas, Austin 

Frederick Walter - State University of New York, Stony Brook 

Ray Weymann - Observatories of the Carnegie Institution of Washington 



The areas of intended scientific investigation were summarized by the PDR viewgraphs shown in Figure 1. 
These objectives dictated an ultraviolet instrument (wavelength range = 1100-3200A) with the instrumental 
capabilities given by PDR viewgraph shown in Figure 2. We wanted to attempt a simple approach and design; see 
Figure 3. The GHRS has one major moving device, the "Carrousel," which is used to position the various gratings 
and acquisition mirrors. 



106 



MAJOR SCIENTIFIC OBJECTIVES 

1. THE INTERSTELLAR MEDIUH 

Very Local Gas in the Interstellar Medium 

Molecule Formation and Selective Depletion of Heavy Elements in Dense Clouds 
By Studying Distant Stars' Spectra, Determine Composition & Distribution of 
THE Gas in Adjacent Spiral Arms, Galactic Halo, and Magellanic Clouds 
Search for as Yet Undetected Simple and Very Complex Molecules in Space 

2. MASS LOSS BY STELLAR WINDS AND THE EVOLUTION OF THE OUTER ATMOSPHERES OF STARS 

OB Supergiants in the Magellanic .Clouds 

Coronal Winds in late-Type Stars 

Mass Loss, Chromospheres, Circumstellar Shells in Red Giants 

Mass Transfer in X-Ray (& Other) Binary Stars 

3. ABUNDANCES OF THE ELEMENTS AND STELLAR EVOLUTION 

Abundances in Stars with Wide Age Range to Determine Chemical Evolution of 
the Galaxy 

^. EXTRAGALACTIC SOURCES 

Limited But Important Quasar Studies 

Physical Investigation of Nuclear Regions of Seyfert Galaxies 

5. THE SOLAR SYSTEM 

Atmospheric Structure in Jovian Planets and Their Satellites 
Auroral Activity on the Planets 
Abundance of Deuterium in Comets 

Fig. 1 - Summary of the scientific objectives for the GHRS. 

DESIGN DRIVERS 



Ultraviolet Response - llOOA - 3200A 

Spectral Resolution - R = 2 x 10^ (15 km/s) 

R = 1.2 X 10^ (2.5 KM/s) 

High Sensitivity 

High Photometric Precision 

Angular Resolution - none within field of view 

Fig. '2 - Summary of the desired instrumental capabihties for the GHRS. 

The detectors chosen were the 512-channel Digicons with (Da LiF/CsI faceplate/photocathode combination for 
short-wavelength response and long-wavelength rejection; and (2) a MgF2/CsTe combination for long-wavelength 
response. The Digicons have a very high dynamic range (- 10^). 



107 



CAMERA MIRRORS 
CROSS DISPERSERS ■ 



CARROUSEL 



DIGICON 
CsTe / MgF, 




COLLIMATOR 



DIGICON 
CsI/LiF 



SLIT PLATE 



Fig. 3 - Cartoon of the GHRS. 



T 



T 



T 



— I r 

INTERSTELLAR SOOIUM 02 LINE IN < ORi 



6x10* 5 
I2»I0* 2 5 
22x10* 135 



-5 



\N-\-f 



-y 



10 15 20 25 30 

V, [Km S"*] 



35 



Fig. 4 - The interstellar D2 line at different spectral resolutions; the obse.-vations at R = 6 x 105 are from 
rlobbs (1969). 



The R = 2 X 104 mode uses four first-order gratings. The highest-resolution mode uses an echelle grating and 
cross dispersers; it became the R = 1.0 x 10^ mode (instead of the R = 1.2 x IQS mode) when the desired echelle could 
not be fabricated. A replica of a grating was substituted. Finally, note that our R = 2 x 103 mode is not mentioned 
At the time of PDR, it was in the design stage but not yet approved. We envisioned it as a reconnaissance mode and 
as a useful backup to the FOS in the wavelength range 1100-1750A. 

At the time of the PDR, we wished to stress that the inclusion of an R = 1.0 x 105 mode was not arbitrary To 
estabhsh this point, we showed (see Figure 4) the spectrum of the interstellar sodium D2 line at various values of R. 
This value of this inclusion has been completely confirmed. 



108 



The PDR presentation discussed some further details of the instrument, made comparisons with previous 
spaceborne ultraviolet instruments, and concluded with a summary of an instrument with major scientific 
capabilities. (Recall that the scientific instruments were originally selected provisionally and were subject to later 
confirmation). Some of these capabilities have survived the HST spherical-aberration problem intact. 

Specifically, my PDR notes contain the assertions of a "... powerful instrument above the atmosphere" that 
"should produce a huge variety of astrophysical results and discoveries!" Even allowing for the "puffing" of a selling 
environment, things have turned out surprisingly well, as evidenced by the GHRS early results papers already in 
press. "^ 



Notes 

1. Some of these materials were from our original proposal, " A High Resolution Spectrograph for the Space 
Telescope," HRS-680-77-01, July 1977. 

2. Ap J. Letters . 377, No. 1, 10 August 1991 issue. 



109 



STATUS OF THE GODDARD HIGH RESOLUTION SPECTROGRAPH IN MAY 1991 

Dennis Ebbets 

Ball Aerospace Systems Group 

PO Box 1062 ARl 

Boulder, CO 80306 

John Brandt 

LASP University of Colorado 

Campus Box 391 

Boulder, CO 80309 

Sara Heap 

NASA Goddard Space Flight Center 

Code 680 

Greenbelt, MD 20771 

Abstract. At the time of this workshop the Orbital Verification of 
the GHRS had been completed, and the Science Verification was well 
under way. This presentation summarized the state of our knowledge 
about HST pointing accuracy, target acquisition procedures, 
sensitivity, spectral resolution, stray and scattered light, 
wavelength calibration, photometric precision and time resolution. 



1. ACCURACY OF INITIAL HST POINTING 

In January, and again in February, 1991, a GHRS to FGS "Fine 
Alignment Test" was performed, whose goal was to measure the 
location of the GHRS science apertures in the coordinate system 
defined by the FGSs. The test executed properly, the expected data 
were obtained and analyzed, and the positions of the apertures were 
updated. Since that update there have been 22 successful target 
acquisitions, all of which found the star within +/-3 arc seconds 
of the initial pointing. Thirteen targets were found within +/~ 1 
arc second, including all 9 for which the celestial coordinates had 
been measured with the GASP system at STScI. It now appears that 
the geometrical alignments between the science instruments and the 
pointing control system are accurate enough to support routine 
target acquisitions. We recommend that observers use GASP 
coordinates whenever possible, and SAO positions for brighter 
targets. It is important to be careful about such details as 
Equinox (1950 or 2000) , proper motions, and the epoch of the proper 
motion. When specifying a GHRS target acquisition, we recommend the 
use of "search-size=3" with GASP coordinates, and "search-size=5" 
otherwise. Figure 1 shows a histogram of initial pointing errors 
compiled between February and April 1991. 

2. TARGET ACQUISITION OPTIONS 

Three types of acquisitions have been exercised and found to work 
well. We have used "interactive acquisitions" with great success. 
The GHRS commands the HST to execute an outwardly growing spiral 
search pattern, and generates a "field map" of the 2x2 arc sec 



110 




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Figure 1. Accuracy of initial telescope pointing 



14 



NOAO/IRAF V2.9EXP0RT simonSleo Tue 09:45:50 08-Jan-91 
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Figure 2. Sample calibration for grating G140L 



111 



aperture at each point. Software on the ground assembles the maps 
into a mosaic, from which the observer identifies the target. A 
slew request then moves the telescope to place the target at the 
center of the aperture. We have verified that all flight and ground 
software procedures are working properly, making interactive 
acquisitions a reliable option for many targets. 

If the count rate can be accurately estimated for the proposal, an 
"onboard acquisition" will produce excellent results in a shorter 
time. Our standard procedure has either four or five steps. We 
first "search" until the target is detected, followed by a "locate" 
in the 2x2 arc sec aperture. We follow this by a second locate, to 
"peakup" in that aperture. A "map" then shows the final centering. 
If the smaller 1/4x1/4 arc sec aperture is desired, a blind offset 
is made, followed by a map of that aperture. Good coordinates, a 
proper "faint" limit, and use of the "double locate" ensures a 
successful onboard acquisition. We recommend a faint limit of 
approximately 40% of the best estimate of the count rate, and a 
"step time" chosen to produce at least 200 counts. 

A third option is an initial acquisition of a nearby star followed 
by a blind offset to the real target. We tested this with a 
separation of approximately twenty arc seconds, and found the 
target within a few tenths of an arc second from the center after 
the offset. A "peakup" then improved the position. 

We have found that the current algorithm for "peakup" in the small 
science aperture does not work well, and we do not recommend its 
use. The problems are a combination of the broad image structure, 
jitter, and inherent imprecision in the method. An improved 
algorithm has been designed for a future update to the flight 
software. 

We had not attempted any acquisitions of moving targets, nor any 
WFC assisted acquisitions at the time of this workshop. 

3. SENSITIVITY AND SPECTRAL RANGE 

The photometric sensitivity of the GHRS has been calibrated using 
measurements of three UV standard stars - BD+7 5d325, HD93521 and Mu 
Columbae. The throughput for the 2x2 arc sec aperture has been well 
determined, and is 0.4 to 1.1 times the prelaunch estimates. Most 
of the discrepancies result from the 30% loss of light in the 
aberrated image, and errors in the ground based calibration at the 
shortest wavelengths. There is no evidence for deterioration of the 
sensitivity since 1984. The 1/4x1/4 arc sec aperture transmits 1/5 
to 1/3 of the light captured by the 2x2 arc sec aperture. The 
recommended wavelength intervals for the first order gratings are: 



112 



GRATING 


RECOMMENI 


)ED RA 


G140L 


1050 - 


- 1900 


G140M 


1050 - 


- 1900 


G160M 


1150 - 


- 2300 


G200M 


1600 - 


- 2300 


G270M 


2200 - 


- 3300 



COMMENT 



2nd order overlap for lambda > 2300 
2nd order overlap for lambda > 2300 
2nd order overlap for lambda > 3300 

Figures 2 and 3 show the large aperture calibration for grating 
G140L, and the ratio of small to large aperture throughput. 
Complete sensitivity information will be available in the Proposal 
Instructions, the GHRS Instrument Handbook, and in the IDT's End of 
SV Report. Figure 4 shows a spectrum of a starburst knot in a 
spiral arm of the Seyfert galaxy NGC 1068, obtained with G140L and 
calibrated with the data shown in Figure 2. 

4. SPECTRAL RESOLUTION 

The image of the small science aperture maps onto one diode in the 
GHRS detectors. For observations obtained with a star in this 
aperture, the spectral resolving power is unaffected by the OTA 
image structure. The two onboard spectral lamps have apertures and 
optical paths identical to the science aperture. The illumination 
is similar enough that the profiles of their emission lines serve 
as reliable proxies of the line spread function and resolution. We 
have measured the profiles of hundreds of calibration lines, and 
have verified that the GHRS internal optical focus is essentially 
perfect, and the full resolution planned for the instrument is 
available. Figure 5 shows a histogram of line widths for Echelle A, 
with a peak at 1.05 diode widths. We have verified that the line 
width is constant, showing no variation with wavelength, location 
on the detector, or echelle order. Figure 6 shows the line widths 
converted to resolving power for the Echelle. 

Observations made with the target centered in the large aperture 
suffer a loss of resolution of approximately a factor 2. The line 
spread function has a sharp core and significant but truncated 
wings. Experiments by members of the GHRS science team and the 
STScI have demonstrated that deconvolution can be performed if the 
S/N of the raw data is adequate. Three techniques have been 
explored so far, a "block-iterative" restoration, the "Richardson- 
Lucy" algorithm, and a "Fourier Quotient" approach. All three 
preserve the location of features in the spectrum and greatly 
improve the contrast between blended features. Equivalent widths 
may not be preserved, so quantitative measurements may be better 
served by small aperture data. Figure 7 shows one example of a 
small aperture spectrum, and a deconvolved large aperture exposure 
for comparison. 

We recommend the following observing strategies if spectral 
resolution is an important goal. Use the small science aperture 
with increased exposure time to compensate for the lower 
throughput. Use a "step-pattern" with four samples per diode width. 
Use the default "comb-addition" of four. Keep the duration of 



113 



NOAO/IRAF V2.9EXP0RT simonSleo Wed 08= 15; 13 09-Jan-91 
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Figure 3. Small/large aperture throughput 




Figure 4. G14 0L spectrum of a starburst in NGC 1068 



114 



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115 




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Figure 8. Removal of echelle scattered light 



116 



individual subexposures to no longer than ten minutes. Merge 
individual "fp-split" segments together carefully by coaligning 
spectral features. Figure 9 shows a short spectral interval of the 
star Chi Lupi observed with various combinations of apertures and 
gratings. 

5. STRAY AND SCATTERED LIGHT 

Many of the scientific programs of the GHRS require very precise 
measurements of line profiles. Accurate profiles in turn require 
that all background components be removed from the raw data. We 
investigated four possible types of scattered light. Near angle 
grating scatter is essentially the far wings of the instrumental 
line spread profile. For the low and medium resolution gratings 
this appears to be negligible beyond two diode widths or so from 
line center. In the high resolution modes both the Echelle and 
cross disperser gratings produce faint but broad wings which 
scatter light for many angstroms. Light is scattered both along the 
direction of dispersion, and into the interorder regions, and 
requires careful removal. We recommend using a "step-pattern" 
which samples the interorder light with the science diodes, 
especially for Echelle exposures. Figure 8 shows Echelle 
observations of saturated interstellar absorption lines in the 
spectrum of Xi Per. After removal of the background the cores of 
the lines show essentially zero residual intensity, as they should. 
A third effect is "red leak", in which longer wavelength light 
could be scattered into the field of view and superimposed on the 
true ultraviolet signal. We have observed uv emission lines from 
very cool stars, and detect no spurious "continuum" light between 
the lines. Side 1 modes in particular are extremely "solar blind." 
Figure 10 shows our observations of chromospheric emission lines 
from Gamma Dra at low, medium and high resolution. The fourth 
effect is telescope scattering, in which light from nearby bright 
stars is scattered into the aperture during observations of nearby 
fainter targets. We measured the signal as a bright standard star 
was stepped away from the GHRS aperture, and found a residual 
fraction of 2E-5 at 16 arc seconds off center. 

6. PHOTOMETRIC PRECISION AND SIGNAL TO NOISE RATIO 

We have quantified four effects which influence the photometric 
quality of GHRS data. Photon statistics dominate the signal to 
noise ratio, and exposure times should be based primarily on this. 
The dark noise in the detectors is less than 0.01 counts per diode 
per second. Dark noise is significant only if the source count rate 
is less than roughly five times this rate. Scattering in the 
Echelle creates another statistical noise source, which can be 
minimized with proper smoothing and removal of the background. 
Finally, small scale irregularities in the photocathode 
sensitivity, gradients, blemishes, granularity etc. contribute 
noise on spatial scales from the full detector width to pixel to 
pixel variations. These can be accommodated somewhat by using the 
"fp-split" procedure. Figure 11 shows signal to noise ratio 
achieved in some prelaunch calibration spectra. The data follow the 



117 



OBSERVATIONS OF chi. Lupi 



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I I I I I I I I I I I I I I I I I I I 



1937.0 1937.5 1938.0 1938.5 1939.0 
WAVELENGTH 



Figure 9. Comparison of resolving power with GHRS large and small 
science apertures, medium and high resolution gratings 



118 



T 0.0 t 




1301 1302 1303 1304 1305 1306 1307 



to 15.0 

CM 
I 

a 



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EchA I 

15.0 h Resolution = 
93,000 

10.0 h 



5.0 
0.0 



1301 1302 1303 1304 1305 1306 1307 
Wavelength (A) 



Figure 10. 



Emission lines from the chromosphere of Gamma Dra, a 
cool star with no ultraviolet continuum emission. The 
GHRS is "solar blind" to scattered long wavelength 
light, and produces no spurious "red leak" signal. 



119 



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Figure 11. The signal to noise ratio is dominated by photon 
statistics at low signal levels, and by detector 
irregularities at high levels. 



5000 



Interstellar and Circumstellar Fe II Features 




2596 2599 2602 2604 

Wovelength (A) 



2607 



2610 



Figure 12 . Variable and nonvariable features in the Beta Pic 
spectrum are easily distinguished with accurate 
wavelength scales and good photometric precision. 



120 



expected "root N" behavior up to S/N = 50. The departures at higher 
signal levels show that photocathode irregularities at the 1-3% 
RMS level dominate if they are not calibrated out of the data. 

7. WAVELENGTH CALIBRATION 

The GHRS contains two onboard Pt-Ne hollow cathode lamps, whose 
spectrum was calibrated for precise wavelengths at NIST. 
Observations of these reference spectra have been used to establish 
the relationships between wavelength and geometrical location for 
all grating and Echelle modes, at a wide range of carrousel 
positions. We formulate the problem is such a way that a polynomial 
describes the relationship at any given carrousel position, and the 
coefficients of the fit, the "dispersion constants" can be smoothly 
interpolated to carrousel positions not explicitly calibrated. We 
have settled on a cubic representation. Our data reduction process 
evaluates this basic step, and makes adjustments for offsets 
between the lamp and science apertures and a small systematic 
thermal drift. There are smaller errors associated with geomagnetic 
effects, miscentering of the star in the aperture, and carrousel 
repeatability that are not modelled at this time. Experiments with 
a large data base of lamp spectra, and a growing set of 
observations of interstellar absorption lines indicates that our 
wavelength calibrations are internally consistent to approximately 
+/- 1 km/sec, with an absolute zero point uncertainty of about 3 
km/sec. Figure 12 shows two of our medium resolution spectra of 
Beta Pictor is, taken approximately three weeks apart. The excellent 
registration of the non-variable interstellar components 
illustrates the accuracy of the wavelength scale, and our 
confidence in measurements of the time variable circumstellar 
components. 

8. TIME RESOLUTION 

The two basic spectroscopic operating modes ACCUM and RAPID, can 
both produce data with higher time resolution than has been 
previously possible. ACCUM allows use of the full suite of flight 
software features, and can achieve approximately thirty seconds 
time resolution for the duration of target visibility. RAPID mode 
bypasses the flight software, so no substepping or data quality 
checking can be performed, but can achieve time resolution between 
0.05 and 12.75 seconds. If the "sample time" is longer than 0.35 
seconds very long observations can be accommodated by the tape 
recorder. Shorter sample times limit the duration to approximately 
20 minutes. Our science team has used RAPID with 0.4 second sample 
time for observations lasting seven orbits. 



121 



9. SUMMARY - WHAT IS RIGHT, WHAT IS NOT? 

At the time of this workshop, the initial pointing accuracy and 
target acquisition procedures were working very well. We had no 
serious difficulties getting stars into the large science aperture. 
Sensitivities had been calibrated and found to support use of the 
entire 1050 - 3300 A spectral range. The planned spectral 
resolution is achieved by placing the target in the small science 
aperture. Observing procedures and data reduction algorithms allow 
accurate removal of scattered light in the Echelle modes. The 
signal to noise ratio is dominated by photon statistics up to about 
S/N = 50. Routine wavelength calibrations are accurate to better 
than one pixel. All operations, commanding and flight software 
aspects appear to be working well. 

Precise and reliable centering of targets in the small science 
aperture is still being worked on. Improvements to data base 
parameters and flight software algorithms will hopefully improve 
the utility of this aperture. The throughput of the small aperture 
has suffered significantly from the aberrated telescope image, but 
hopefully will be improved by the proposed COSTAR instrument. 
Unanticipated sensitivity to thermal and magnetic environments 
require some special care. The photocathode irregularities have not 
been fully calibrated, and their removal is not yet automatic. 

Acknowledgement: The results discussed in this paper represent the 
work of dozens of individual scientists, engineers and managers. 
We gratefully acknowledge the contributions of the members of the 
Investigation Definition Team and their many colleagues. The 
tireless efforts of many individuals at HSTPG and the STScI were 
required to bring the OV and SV programs to fruition. The 
engineers from Ball Aerospace, and the HST operations staff 
contributed immensely to the development of the GHRS and the 
implementation of its scientific program. 



122 



Early Operations with the High Speed Photometer 



J. W. Percival, R. C. Bless, and M. J. Nelson 

Space Astronomy Laboratory 

1150 University Avenue 

Madison, WI 53706 

USA 



Abstract. The performance of the High Speed Photometer (HSP) during the Orbital 
Verification (OV) and the Science Verification (SV) programs of the Hubble Space 
Telescope (HST) is described. The HSP is operating as designed, and all hardware is 
fully operational. The HSP has been seriously affected by the degraded point spread 
function (PSF) of the telescope system, the telescope pointing calibration, and the jitter 
in the spacecraft pointing. 

1. INTRODUCTION 

The design and basic operation of the HSP has been described elsewhere (Nelson 
et al. 1991 and White, 1990). It can produce high speed (12 /zS) photometry in 27 
narrow, medium, and broad filters from 1200 to 7500 Angstroms. The HSP has no 
moving parts, and selects targets and filters in a two step operation. First, to select a 
target, the HST is maneuvered so that the target image falls at a particular point in 
the focal plane, where it passes through the desired filter and a 1 arcsecond observing 
aperture. The filtered target image then falls on the face of an Image Dissector Tube 
(IDT). Next, the HSP acquires the target by magnetically steering the IDT's read beam 
to the point on the tube face on which the filtered image is faUing. In practice, there is 
in intermediate operational step. The star is first acquired in a 10 arcsecond aperture, 
and an on-board target acquisition is performed that centers the star in the aperture 
of choice. 

The large number of HSP apertures, and the need for precise, repeatable, and stable 
HST pointing to produce high quahty photometry conspired to comphcate and lengthen 
the HSP OV and SV activities. The HSP is nearing the end of these activities, and 
has begun to observe scientifically interesting targets, including a rapidly oscillating Ap 
star, and a stellar occultation by Saturn's rings. 



2. INTERNAL CALIBRATIONS 

The primary internal calibration for the HSP is to measure the magnetic deflection 



123 



currents that move the IDT read beam to the part of the tube face on which the desired 
filter/aperture combination is imaged. Initially, this was done using the bright earth as 
a flat field source, backlighting the aperture plate. We make a crude image by stepping 
the read beam through a grid, making a photometric measurement at each point. These 
images are analyzed to yield a deflection coordinate pair for any given aperture. 

We discovered that the earth is not a very flat field. It is spikey on small (1 km) 
spatial scales, probably due to cloud tops and strips, and scattered sunlight from water. 
The spikes were 5-50 times the average expected brightness of the earth. Operational 
changes were made for the solar sensitive detectors to lower the tube gain when exposed 
to the bright earth, and Orion was selected as a flat field for some filters on these detec- 
tors. The operational changes and much longer exposure times delayed the calibration 
of these detectors by several months. The deflection calibrations for each aperture are 
now known to within about 0.05 arcseconds. 



3. ALIGNMENT CALIBRATIONS 

The other alignment critical to normal HSP operations is the calibration of the focal 
plane positions of the photometry apertures. This calibration is performed by scanning 
the HST in a grid pattern on the sky, while doing time-series photometry through the 
selected aperture. An analysis of the time- varying signal, combined with a post-test 
knowledge of where the spacecraft was pointing, yields a focal plane coordinate pair for 
each aperture. 

Early in the mission, the fine guidance sensor (FGS) calibration was so poorly 
known that the alignment stars could not be positioned reliably within the diameter 
of the 10 arcsecond finding aperture. ReUable, repeatable FGS calibrations only began 
appearing in January 1991. Focal plane calibrations have proceeded smoothly since 
that time, and the HSP aperture positions in the HST focal plane are now known to 
within 0.02 arcseconds. Figure 1 shows contour plots of the same star in two different 
HSP photometry apertures. 



4. OPERATIONS 



4.1. Bright Earth 

The bright earth problem discussed above resulted in operational changes for the 
two solar sensitive IDT's and for the photomultiplier tube (PMT). Whenever the cur- 
rent pointing is about to be occulted by the bright earth, the IDT gain is lowered by 
decreasing the high voltage. After the occultation, the voltage is restored to its nor- 
mal observing value. The PMT has a less fragile reflective, rather than a transmissive, 
photocathode, so the operational change for it was simply to turn off the bright object 
protection software for this tube during the bright earth events. The bright object 
protection is reenabled after the occultation. 



4.2. Target Acquisition 

The HSP has an on-board target acquisition mode. A crude image, described in 



124 



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125 



Section 2 above, is made of the 10 arcsecond finding aperture after moving the HST to 
place the target star within it. The image is sent to an on-board computer program that 
performs a simple centroid calculation on the target, and issues a small angle maneuver 
request to center up the target. The relative positions of the HSP apertures are known 
to great accuracy, so after the centering operation the target can be reliably moved to 
any photometry aperture with subsequent small angle maneuvers. 

The image is oversampled with the 1 arcsecond read beam, and the original al- 
gorithm was designed to operate on the expected extremely narrow PSF. The much 
broader actual PSF presents a larger than expected image to the simple nearest-neighbor 
centroid algorithm, which results in a slightly less accurate centroid. We have found 
that doing the acquisition twice in a row causes an improvement in the net centering. 
We find that the target acquisitions converge rapidly (two is enough) and is repeatable 
to within 0.05 arcseconds. This is now done automatically. The proposer need not 
request the iterated acquisition. 



4.3. Length of Exposures 

The HST Science Data Formatter (SDF) protocol was designed with the imaging 
instruments in mind. The SDF operates in a three dimensional data space whose 
axes are words per hne (WPL), lines per frame (LPF), and frames per observation 
(FPO). The idea is that WPL and LPF are determined by the image size, the image is 
integrated before the science data transfer is started, and that once begun, the transfer 
should proceed at a high rate. The SDF expects hues of data to follow each other 
in the LPF dimension within a strictly observed timeout period of 10 ms. The third, 
FPO, dimension represents successive frames so the SDF allows the science instrument 
to insert an arbitrary amount of time between frames. 

The HSP does not collect a whole frame before sending. It has no buffers that large. 
It begins the transfer to the SDF as soon as the first line of data has been collected. 
Because of the 10 ms SDF timeout in the LPF dimension, the HSP must be prepared 
to transmit the next line of data within 10 ms of the previous one, which it can do only 
in the high-speed (MHz) photometry regime. For lower speed photometry, the HSP is 
forced to coUapse the LPF dimension to unity, sending each hne as a new frame. The 
final piece in this puzzle is that the HSP FPO counter is only an 8-bit quantity, which 
Umits the HSP to 256 frames. This places an upper Umit on the number of samples per 
science observation of 2 * WPL * FPO, or about a half miUion 8-bit samples (half as 
many 16-bit samples, of course). This constrains the sample time for long observations. 
As you lengthen the observation but keep the total number of samples fixed, then each 
sample must last longer to span the time. A detailed analysis can be found in White 
(1984). 

One interesting observation has already been subjected to this constraint. In our 
5.6 hour time series of a rapidly oscillating Ap star, the quarter million sample limit 
forced each sample to be no shorter than 82 milliseconds. We have proposed a simple 
commanding fijc to eliminate this constraint. This fix, when implemented, will allow 
arbitrarily long data sets. This commanding fix is now being reviewed by the STScI. 



5. PERFORMANCE 



126 



5.1. Throughput 

The throughput of the HSP filters has been measured with stars during the FGS 
alignment activities. The results agree, with one exception, with HSP model predictions 
modified only for the broadened PSF. We find the expected 50% reduction for the 1 
arcsecond photometry apertures, and a 75% reduction for the 0.65 arcsecond forward 
facing polarimetry IDT. 

The exception is the 5500 Angstrom filter (F551W) on the solar sensitive IDT. We 
find an unexplained 60-70% loss of fight in each of the four apertures on that filter. We 
are investigating this curious result. 



5.2. Linearity 

In our Hnearity test, we looked at four stars between 5 and 11*'' magnitudes. 
Exposure times were chosen to produce a signal to noise (S/N) ratio of 50. Table 1 
shows the known and measured V magnitudes. The measured value was derived from a 
magnitude-count relation with a linear color term. The agreement is satisfactory given 
the S/N of the HSP data. 



Catalog Mv 


Measured My 


5.111 


5.119 


7.247 


7.267 


8.060 


8.035 


11.070 


11.067 



Table 1: HSP linearity. 



5.3. Photometric Performance 

The photometric performance of the HSP is aff^ected by the HST spherical aberra- 
tion because the broadened PSF has significant energy at the edge of the photometry 
aperture, a situation that the system was not designed to encounter. The presence of 
energy at the edge of the aperture increases our sensitivity to two effects, the ability to 
point the HST in a repeatable way, and the stability of the pointing during an exposure. 

The repeatability of the pointing places a Hmit on our photometric accuracy. If a 
star is measured two different times, and the position of the PSF changes slightly from 
one time to the next, then we will measure a slightly different count rate as more or 
less fight passes through the aperture. Figure 2 shows the predictions of a model that 
numerically integrated the energy under the flight PSF as a function of miscentering in 
the aperture. The target acquisition repeatability of about 0.05 arcseconds implies a 
lower Hmit of about a miUimagnitude in absolute photometry. 

The HST exhibits some pointing instabiUty that is detectable by the HSP. The 
day/night terminator reaction can cause the stellar image to move nearly completely 
out of the photometry aperture, and the induced jitter can inject power into a time 
series at frequencies ranging from 0.1 Hz to 10 Hz. Figure 3 shows the 5.6 hour time 
series of the Ap star. Note the loss of fight at the transition from night to day (the 
low part of the sinusoid). The sinusoidal variation occurs at the orbital period, and is 
a topic under investigation. It may have to due with thermal effects in the HSP or in 



127 




0.05 0,1 

Miscentering of Star in ArcSecoads 



Figure 2: Photometric error vs. centering error for a 1" aperture. 



5.6 hour time series showing night/day pointing instability 



700 



600 




Figure 3: Data dropouts at night/day boundary. 



128 



FFT showing improvement in HST jitter 



V 

o 

a. 




2 4 

Frequency (Hz) 



Figure 4: Improvement in HST jitter seen by the HSP. 



129 



the HST guidance system. 

Figure 4 shows part of the Fourier transform of this time series. Note the lack of 
detectable power at 0.6 and 2 Hz, where power was earlier seen in HSP data. This 
improvement in the Fourier domain is apparently a result of the improved Solar Array 
Gain Augmentation (SAGA) software running in HST. This SAGA fix has improved the 
performance of the guidance system at the terminator crossings. The data gUtches in 
Figure 3, while large in magnitude, are short in duration (seconds rather than minutes) 
and appear to die out quickly. A more extensive jitter test is being designed by the 
STScI. 



6. SUMMARY 

The HSP is performing as designed, except for the effects of the HST spherical 
aberration. The instrument is linear to within the 2% accuracy of the data. The 
on-board target acquisition is repeatable to within 0.05 arcseconds, which places a 
theoretical lower limit on the photometric accuracy at about 1 miUimagnitude. The 
HSP has nearly completed its Science Verification activities, and is now beginning its 
GTO science program. 



REFERENCES 

Nelson, M. J., Bless, R. C, and Percival, J. W. 1991, Photometry From Space, ASP 

Meeting, June 1991. 
White, R. L. 1990, Hubble Space Telescope High Speed Photometer Handbook, Space 

Telescope Science Institute. 
White, R. L. 1984, Timing Considerations for HSP Data Collection, STScI Instrument 

Science Report HSP-001. 



130 



Early Commissioning Astrometry Performance of the 
Fine Guidance Sensors 

G. F. Benedict, W. H. Jefferys, Q. Wang, A. Whipple, E. Nelan, 

D. Story, R.L. Duncombe, P. Hemenway, P. J. Shelus, B. McArthur, 

and J. McCartney 

University of Texas 

Austin, TX 78712 

O. G. Franz, L. Wasserman, and T. Kreidl 
Lowell Observatory 
Flagstaff, AZ 86001 

Wm. F. van Altena and T. Girard 

Yale University 

New Haven, CT 06511 

L. W. Fredrick 
University of Virginia 
Charlottesville, VA 22903 

Abstract. We discuss astrometry with the Fine Guidance Sensors and 
explore various factors Umiting their performance. These results were 
obtained before starting either the Orbital Verification or Science 
Verification programs. 

1 . Astrometric Use of FGS 

The Hubble Space Telescope contains three fine guidance sensors 
(FGS). While two are used for pointing control, the third is available for 
astrometric measurements. These measurements fall into two broad 
categories, position mode (POS) and transfer scan (TRANS) mode. A 
detailed discussion of these modes and the required post-observation 
processing can be found in Bradley et al. (1991). A principal goal of 
Astrometry Orbital Verification is the selection of the astrometer: which 
FGS of the three available will give the best performance in both position 
and transfer mode. Currently, this critical choice will be made in 
November 1991, after the final mirror moves to minimize coma and 
astigmatism are made. 



131 



1.1 Transfer scan mode 

Since it affords us an introduction to the response function of the 
FGS, we shall first discuss TRANS mode. Fig.l presents an example of the 
characteristic response of the FGS 1 Y-axis for Upgren 69 in NGC 188, a 
star thought to be without a companion. This curve (often referred to as an 
S-curve) is generated by scanning the 5 arcsec square FGS entrance 
aperture over the star. A similar curve will be produced for the x-axis. 
Positive attributes of any transfer function include large modulation, which 
is the left peak to right valley amplitude, and the detailed shape of the 
curve. A transfer function should ideally have only one positive peak and 
one negative valley. Comparing the ideal with the actual, we see three 
peaks and three valleys in Fig. 1, Minimizing these secondary peaks and 
valleys is a prime consideration. 

A double star will produce a transfer function which is the sum of 
two overlapping single star transfer functions. It is obvious that the detailed 
shape of the single star transfer function will affect binary detection. The 
sample curve, for a V=9.58 star, will become very much noisier for a 
fainter star. This noise, too, will affect the detection hmits for duplicity. 

The pre-launch TRANS mode performance goal for binaries was 10 
mas separation detectability for two stars differing in magnitude by less 
than 0.75 magnitudes. Unfortunately, the spherical aberration and lack of 
critical collimation of the telescope along with the intemal mis-alignments 
within each FGS also affect the shape of the transfer function. The ultimate 
capability of the FGS for double star astrometry will not be known until 
the telescope is properly collimated. 

1.2 Position mode 

In POS mode, the fine guidance electronics searches for the first 
zero crossing in the FGS response curve after the first deep minimum, 
traveling right to left along the curve in Fig. 1. The positions of the star 
selector coordinates are averaged for some period of time while the fine 
guidance electronics keeps the star at this position, called the null point. 

Fig. 2 illustrates the planning and the mechanics of a typical (but, as 
yet, unrealized) POS mode observation. The star selectors move the FGS 
entrance aperture to any location within the pickle-shaped region. To 
measure the position of a target relative to a field of reference stars, we 



132 



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star was Upgren - 69 in 
NGC 188. 



File : Prox Plan 1992 ep2000 decdeg 

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HST Astrometry 
POS mode Performance Goals 

Measure 10 stars in 20 minutes 

Formal uncertainly in positions: 0.0027 arcsec 

V = 17 limiting magnitude 




Figure 2 - An astrometry observation planning 
chart for Proxima Centauri. Plotted are reference 
stars and positions for Proxima Centauri at various 
times between 1950 and 2000 A.D. 



133 




command the star selectors to the predicted position of the target, obtain 
star selector readings averaged for some specified period of time 
depending on the star magnitude, then proceed to do the same for each 
reference star. The limiting magnitude will be determined primarily by the 
height of the peak-to-peak modulation of the transfer function (S-curve) 
because as the noise increases with fainter stars, the S-curve itself will 
become buried within the peaks and valleys of the background noise. The 
original pre-launch POS mode performance goal, 2.7 mas per 
measurement on a V=17 magnitude star, will require that many of the 
performance issues discussed in the next section be resolved. Aside from 
the scientific loss inherent in degraded astrometric performance, these 
issues, if unresolved, impact guiding. 

2. Performance Issues 

Both the environment of the FGS and the internal conditions of the 
FGS directly determine the interferometric response and hence the ability 
of the FGS to make astrometric measurements. This section describes these 
difficulties and in some cases identifies possible solutions. With a clear 
understanding of these extrinsic and intrinsic conditions, we can estimate 
our potential on-orbit performance. 

2.1 Problems Extrinsic to FGS 

2.1.1 Spherical Aberration 

This major blow to HST performance has most often been discussed 
in the context of camera science (e.g., Hester, et al., 1991). For many 
months the FGS's were thought to be immune to spherical aberration. A 
series of 'N Points of Light' tests proposed by the Astrometry Science 
Team, in which the same star is observed in TRANS mode in 'N' locations 
within each FGS field of view, have demonstrated otherwise. The results 
for one such test (a Five Points of Light) for all FGSs are shown in Fig. 3, 
and will be discussed in greater detail below, 

2.1.2 Collimation 

Once it was determined that the Optical Control System Wavefront 
Sensors were not usable in the presence of spherical aberration, Hughes 
Danbury Optical Division and the Astrometry Science Team proposed to 
explore the secondary mirror tilt and decenter (collimation space) using the 
'N Points of Light' tests. Coma, the result of misalignment of the HST 



134 




Figure 3 - Summary of a "5 Points of Light" test done day 066, 1991. 
Plotted are x-axis S-curves for FGS 1 and 3, y-axis S-curves for FGS 2. 
Positions within the FGS field of view are blackened within each FGS 
"pickle". Note the variation of modulation with position within the "pickle" 
for FGS 3. 



135 



secondary mirror relative to the primary, produces a characteristic 
deformation of the intrinsic transfer function of each FGS. This 
deformation can consist of some combination of modulation reduction and 
the introduction of additional and spurious peaks and/or valleys. 

This very productive series of observations discovered as much 
about each FGS as about the state of coUimation of HST. The unhappy 
conclusion, corroborated by ground testing of a flight-spare FGS, is that 
the large spherical aberration of the as-built primary mirror, in the 
presence of internal FGS misahgnments, produces a signature in the 
transfer functions which mimics coma. The pair of transfer scans of the 
same star in the same FGS shown in Fig. 4 illustrates the problem. At one 
secondary mirror position the S-curve has far more modulation than at the 
other position. 

The ultimate result of the 'N Points of Light' tests was to prove the 
existence of FGS intemal misalignments, a not inconsequential reward, 
since the detailed shape of the transfer function influences binary detection. 
For FGS 3 (Fig. 3) the intemal misalignments cause the shape of the 
transfer function to vary with location within the FGS. 

Another conclusion from these tests was that the secondary mirror 
position itself can perturb the transfer function shape. It appears to be 
impossible to obtain high quality S-curves from all three FGS units 
simultaneously. It is hoped that the secondary mirror will eventually end 
up in a position which will provide very high quality S-curves from one 
FGS, lesser quality, but still usable S-curves from another, and guiding- 
quality S-curves in the third FGS. 

The principal effect of lack of collimation for FGS astrometry at this 
point has been to delay our choice of astrometer. Once the WFPC and FOC 
have determined the best coma and astigmatism secondary mirror position, 
a final 'N Points of Light' test will be carried out to select the astrometry 
FGS. From transfer function variations seen as the secondary changes 
positions, it has become clear to us that a stationary secondary mirror is of 
prime importance. 
2.1.3 Jitter 

A major contributor to the FGS error budget, jitter, has its strongest 
source in the response of the Solar Arrays to terminator crossing. This 
stimulus occurs roughly eighteen times each 24 hours. The strongest modes 



136 



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two different secondary mirror positions, day 116 (left) and day 122 
(right) in 1991. 



137 



excited have frequencies near 0.1 Hz and 0.6Hz. The effects are strongest 
for about 6 minutes after each terminator passage, although there are 
random 'quakes' throughout orbit night or day. 

Fig. 5 displays an example of the jitter problem. We present 
TRANS mode data for a star observed as part of a preliminary thermal 
test. The five consecutive scans, each taking one minute, clearly indicate 
motion along the scan axis, especially in the third panel. 

Fig. 6 shows what the guiding FGSs "see" at the terminator. We plot 
the sum of the two PMT channels on each axis against time. FGS 3 and 
FGS 2 are installed in HST such that the Y axis of FGS 3 is parallel to the 
X axis in FGS 2, The disturbances seen by the FGS seem highly correlated, 
which raises the possibility that they could be removed. Note the presence 
of both the O.lHz and 0.6Hz oscillations. These data were acquired before a 
partial fix via on-board software (SAGA) was installed. In an optimistic 
sense, jitter is a temporary problem, since NASA is committed to 
exchanging the flawed Solar Arrays with a redesign much less sensitive to 
temperature variations. This upgrade is scheduled to occur during the first 
refurbishment mission. 

2.2 Problems Intrinsic to FGS 
2.1.1 Internal misalignments 

As discussed above, the 'N Points of Light' tests demonstrated 
conclusively that spherical aberration, coupled with intemal FGS 
misalignments, can change the shape of the transfer function. For FGS 1 
and FGS 2, the misalignments perturb the transfer function in a similar 
fashion throughout the FGS field of view. For FGS 3 (Fig. 3) the intemal 
misalignments cause the shape of the transfer function to vary with location 
within the FGS. 

Each FGS has a filter wheel installed near an aperture stop. One 
position in this wheel contains a 2/3 aperture stop. A partial, but for some 
projects most unsatisfactory, fix consists of observing with the FGS 2/3 
aperture. This greatly reduces the consequences of FGS misalignment, 
often restoring the transfer function to a near normal shape and 
modulation. Unfortunately, the use of the 2/3 aperture reduces our limiting 
magnitude. Rare is the scientifically interesting target surrounded by bright 
reference stars! Relying on the 2/3 aperture to 'fix' misalignments also 



138 



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duration was 100 seconds. Jitter affects all but scan 2. Multiple scans in 
TRANS mode are now standard procedure. 



139 



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Figure 6 - Top, FGS 3 y-axis, bottom, FGS 2 
X-axis. The sum of the two PMT channels on 
each axis is plotted against time. Easily seen are 
the 0.1 Hz and 0.6 Hz components to the HST 
jitter. Note the correlation between FGS 3 and 
FGS 2. These are parallel axes, due to the 
placement of FGS 3 and FGS 2 within HST. 



140 



prohibits the use of the scientifically useful bandpass limiting filters 
installed in each FGS. 

2.2.2 Optical Field Angle Distortion 

Achieving the specified 2.7 mas relative positional measurement 
accuracy will require that we map and remove distortions known to be 
present in the FGS optical train. These distortions are caused by, for 
example, non-flamess in the pickoff mirror and figure imperfections in the 
FGS asphere. We lump all effects together and call them optical field angle 
distortion (OFAD). Ideally, one would like to simply observe a field of 
stars whose relative positions are known to 0.5 mas. Analyzing the 
residuals of a fit of FGS data to these known positions would give 
information on the intrinsic FGS distortions. Unfortunately, no such field 
exists. The FGS itself is the only device accurate enough to measure the 
effects of optical field angle distortion. 

How then will we determine the distortion? We shall observe a rich 
star field (such as NGC 188 or NGC 5617) as follows. We obtain POS 
mode data for about 25 stars in each FGS. We do this fifteen times, varying 
the pointing by some fraction of the width of an FGS, to achieve significant 
overlap. Next, the data are corrected for known star selector encoder 
errors and HST orbit-induced velocity aberration. We finally subject these 
data to overlapping plate techniques using GaussFit (Jefferys et al., 1991), 
solving for star positions and distortion coefficients simultaneously. The 
constraint that the relative star positions are unchanged for each 
observation set allows us to determine the optical effects intrinsic to the 
FGS. 

To demonstrate the overall correctness of the approach, we provide 
some results from a preliminary OFAD. These results (to be discussed at 
greater length in Wang et al., 1991) are not a sufficiently accurate mapping 
of positions and distortions to achieve our goal of 2.7 mas relative 
positions. The observations were done in coarse track rather than fine lock. 
We obtained only five pointings, rather than fifteen. Nonetheless, they do 
show that our methodology works. 

For this sparse data set we restricted ourselves to a model with order 
five or fewer terms for the distortion coefficients. We also included 
ground-based astrometry as a check, knowing that if the approach was 



141 



viable, these would have far larger errors in position than did the HST 
data. 

The results are encouraging, especially considering the non-optimal 
observing conditions and degree of overlap. Fig. 7 displays the x and y 
residuals (0-C) plotted against the x coordinate. The residuals plotted 
against the y coordinate have a similar distribution. Both residual plots are 
for the HST data. As expected, the ground-based residuals (not shown) 
were considerably larger. Fig. 8 presents the x and y residuals as a function 
of position within FGS 1 , showing that we have mapped the distortions 
with 3 to 5 mas rms residuals. Note that the smaller residuals in y are 
probably due to the smaller jitter along this axis for this FGS unit. The y 
axis is perpendicular to the direction of the dominant flapping mode of the 
solar arrays. To obtain the full field 1-2 mas will require many more 
pointings, fine lock, and the scheduling legerdemain required to avoid 
terminator crossings. 

Finally, each time the secondary mirror is moved, the 
milliarcsecond calibrations of the field distortions are obliterated. This 
stands as a compelling astrometric argument for a stable secondary mirror. 

3. Transfer Scan Mode Scientific Results 

While waiting for the telescope to be optimally collimated we have 
accomplished some early science using the transfer scan mode. By 
observing a presumed single star close in time to the observation of a 
suspected double, and by observing the presumed single star at the same 
location within an FGS as the target star, we have inspected a set of stars 
for duplicity and have measured the relative positions of the components of 
a known double. 

3.1 Hyades Binary Search 

As part of the Early Release Observation (ERO) program, S-curves 
for about a dozen stars in the Hyades were obtained and examined in an 
effort to detect previously unknown binaries. One star has a suspicious 
looking S-curve and may be double. The detailed results of this survey 
(Franz et al., 1991b) are to appear in The Astronomical Journal. 



142 



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143 



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— -60. IS 



- -60.2 



— -60.25 




I I f I I I I I I I 1 I 



■60,35 



-60.45 



6 4 2 6-2-4-6 

X coordinates In FGS frame (arc-min) 



Figure 8 - X and y residuals as a function of position within FGS 1 for a 
preliminary OFAD study. 



144 



3.2 The Binary ADS 11300 

For another contribution to the ERO program, we chose a known 
double, very close to periastron. One accurate measurement near 
periastron may serve to define its orbit. The details of this successful 
observation can be found in Franz et al. (1991a). The Astrometry Science 
Team will continue to monitor this star to better define the orbit and use it 
as a test of the FGS, since the separation is predicted to go below 0.01 
arcsec in 1992. 

4. Summary 

HST will likely have (after final collimation) one fully capable 
astrometer, probably one fully capable guider in addition to the astrometer, 
and one guider capable of 7 mas guiding using the 2/3 aperture. 

Jitter is ultimately a soluble problem. The interim solution of control 
law modification will eventually render some parts of all orbits quiet 
enough to reach our original relative position accuracy goal. The 
replacement of the existing solar arrays will provide a permanent fix, and a 
continuously stable platform. 

The preliminary OFAD results are encouraging. We have 
demonstrated the logic and utility of our basic approach. 

Finally, transfer scan mode double star science looks very 
promising, although we may have to restrict our observations to a 
particular location within one FGS. 



145 



References 

Bradley, A., Abramowicz-Reed, L., Story, D., Benedict, G., and Jefferys, 
W. 1991, PASP, 103, 317. 

Franz, O.G., Kreidl, T.J.N., Wasserman, L.W., Bradley, A.J., Benedict, 
G.F., Hemenway, P.D., Jefferys, W.H., McArthur, B., McCartney, J.E., 
Nelan, E., Shelus, P.J., Story, D., Whipple, A.L., Duncombe, R.L., 
Fredrick, L.W., and van Altena, Wm. F. 1991a, ApJ, 377, L17. 

Franz, O.G., Wasserman, L.H., Nelan, E., Lattanzi, M.G., Bucciarelli, B., 
andTaff, L.G. 1991b, to appear in AJ 

Hester, J., Light, R., Westphal, J., Currie, D., Groth, E., Holtzman, J., 
Lauer, T., and O'Neil, E. 1991, AJ, 102, 654 

Jefferys, W., McArthur, B., and McCartney, J.E. 1991, BAAS, 23, 997. 

Wang, Q., Jefferys, W. 1991, "Bootstrap statistical analysis of the Hubble 
Space Telescope Optical Field Angle Distortion", in preparation 



146 



A REVIEW OF PLANETARY OPPORTUNITIES AND 
OBSERVATIONS WITH THE HUBBLE SPACE TELESCOPE 



Reta Beebe 

Department of Astronomy, New Mexico State University 

P.O. Box 30001/ Department 4500 

Las Cruces, New Mexico 88003-0001 

U.S.A. 



Abstract 

This review discusses the anticipated capabilities of the Hubble Space Telescope lui 
observations of solar system bodies, reviews the results achieved from preliminary ob- 
servations, suggests future observations and addresses improvements that will simplify 
and enhance planetary observations. 

1. INTRODUCTION 

The possibility of obtaining near simultaneous high resolution, multispectral im- 
ages and ultraviolet spectra of selected regions of solar system bodies provides us with 
opportunities to obtain both spatial and spectral information concerning links between 
chemistry and dynamics of solar system bodies. But, obtaining desired observations is 
complicated by the differential motion of solar system objects relative to background 
stars. Although full implementation of software to track moving targets has not yet 
been completed, progress is being made. So far we have obtained data that expand our 
knowledge concerning temporal variability of planetary atmospheres and supplements 
earth-based observations of comets. This review: 

1) addresses the effectiveness of deconvolution of the 
Wide Field/ Planetary Camera (WF/PC) images, 

2) considers the anticipated capabihties of Hubble 
Space Telescope (HST) for a variety of solar system 
objects, 

3) reviews early observations of Pluto, Comet Levy, 
Mars, and Jupiter, 

4) presents a preliminary analysis of the Saturn 
data obtained in November 1990, 

5) previews future observations and 

6) concludes with a summary of problems yet to be 
solved and goals we expect to attain. 



147 



2. DECONVOLUTION OF WIDE FIELD/PLANETARY CAMERA IM- 
AGES 

The WF/PC team obtained multispectral images of Saturn in the wide-held mode 
on August 26, 1990. At this time, less than 6 weeks after opposition, Saturn was easily 
accessible and the globe showed typical east-west banding and color-dependent limb- 
darkening similar to that observed by the Voyager spacecraft. Because the structure of 
the rings was known, these data could be used to evaluate artifacts introduced by the 
deconvolution. The resolution of Encke's division, which is at the limit of earth-based 
resolution, was convincing evidence that HST could provide useful imaging data. 

In late September a major disturbance developed in Saturn's equatorial region. In 
mid-November a series of multispectral images were obtained with the WF/PC. Figure 
1 uses one of these images to illustrate the deconvolution problem. The image on the 
left is the original raw image from the planetary mode (chip P6) of the WF/PC and 
the image on the right is the reduced product, generated by the Lucy method (Lucy, 
1974), assuming a constant theoretical point spread function over the image. Careful 
inspection of the images reveals that deconvolution is limited by uncertainties in the 
fiat-field and noise in the data and that the assumption of a constant point spread 
function is not the limiting factor at this time {i.e., there is no discernible distortion 
of Saturn's rings). 




Figure 1. Deconvolution of WF/PC Data. This image was obtained with the F588N 
filter at 1:39:53 UT on Nov. 17, 1990. The raw image is on the left and the deconvolved 
image is on the right. (North is at the top and east to the right in all images in this 
article.) 



3. ANTICIPATED CAPABILITIES OF THE HUBBLE SPACE TELE- 
SCOPE 

The initial specifications for the HST and its associated camera and spectrographs 
promised the opportunity to obtain observations that would enhance our knowledge of 
both short term and longer term variable phenomena associated with solar system ob- 
jects. The opportunity to sample at 1.5 hour intervals, to obtain high spatial resolution 
through relatively broad-band filters and to acquire high spectral resolution of areas on 



148 



solar system bodies would provide enhanced capabilities. We will cohskIci ^uiiu u\ i h 
opportunities for some important targets. 



3.1 Pluto 

The fact that Pluto and Charon have recently completed a series of mutual occulta- 
tions (which occur every 124 years) has enabled us to determine the radii and masses of 
the two bodies and has revealed that the albedo and spectral response of the two bodies 
differ. This new information, coupled with the fact that the two bodies have recently 
passed through perihelion and developed thin atmospheres due to maximum insolation, 
motivates us to seek evidence of atmospheric absorption on the two bodies. Although 
the expected lifetime of HST is short (15-17 yrs) relative to Pluto's seasons, the pos- 
sibility of estabhshing surface conditions on the two bodies at a time near maximum 
solar heating is challenging. 

The radii of Pluto and Charon are 1,150 km and 595 km, respectively, and the sep- 
aration of their centers-of-mass is 19,640 km (see rehable tables in Beatty and Chaikin, 
1990). At their current distance from earth (see the Astronomical Almanac, U.S. Gov. 
Printing Office), the diameter of Pluto would subtend a little more than 0.1" and Charon 
about half that. Near perihelion these two bodies have a maximum separaliou ul abuui 
0.9". The angular size of a pixel in the planetary mode (PC) of the WP'/PC is 0.043". 
therefore the two bodies would be separated by about 20 pixels and the dianioters ol 
Pluto and Charon would span about 2.5 pixels and 1.3 pixels, respectively. Thus, the 
system would be a desirable target for multispectral imaging, as well as astrometric 
measurements to refine our knowledge of the system. But, even though WF/PC im- 
ages have been successfully deconvolved and maximum separation of the two bodies, 
0.9", occurs every 3.2 days, the current point-spread function presents an obstacle in 
the way of obtaining individual spectra. 



3.2 Comets 

Observations of comets from ground-based sites are fraught with problems. When 
the comet is bright enough to obtain many of the desired observations its proximity 
to the sun severely hmits the length of the observing window and requires daylight 
observations or dealing with a large zenith angle. These observing aspects reduce 
spatial resolution and severely hamper photometric observations. 

There is increasing evidence that the rotation rate of cometary nuclei is on the order 
of hours. When the patterns of outgasing of Comet Halley are considered, it is apparent 
that the possibility of observing a comet every orbit of HST, at 1.5 hour intervals. 
for multiple orbits several times during a period of weeks would be highly desirable. 
These observations, obtained under constant viewing conditions would greatly enhance 
our understanding of the composition and dynamics of a selected set of representative 
comets. H. Weaver will present WF/PC observations of Comet Levy elsewhere in this 
conference proceedings. 



3.3 Mercury and Venus 

Although, in 1974-75, Mariner 10 encountered Mercury three times, the nature of 

149 



the orbit was such that half of the surface area of the planet was not observed. When 
Mercury is at maximum elongation a pixel of the PC mode of the WF/PC would span 
only 30 km on the surface of Mercury, but the angle between Mercury and the sun is 
17 to 28 degrees. This geometry places Mercury well inside of the HST safety limit 
for near-sun observations. Venus is also within the safety limit at about 46 degrees 
maximum elongation. Therefore, even though the PC mode would yield 22 km/pixel, 
observations of Venus are not possible without relaxing the HST sun safety limit. 



3.4 Mars 

The Earth-Mars distance of 0.666 to 0.381 AU at opposition makes Mars a likely 
target for HST. Because the synodic period is 2.135 years, favorable observing conditions 
will occur biannually. 

Again, the ability to observe the planet at 1.5 hour intervals is useful for monitoring 
events associated with the onset of global dust storm which occur on time scales of hours 
and days. Because Mars and Earth rotate in the same direction with periods of 24.6229 
and 23.9345 hr, respectively, the planet appears to rotate less than 10 degrees per day, 
providing poor coverage of planet-wide events from earth-based stations. 

The last Martian opposition occurred in late November 1990, with a minimum 
Earth-Mars distance of about 0.52 AU. At this distance, an image of Mars spans 420 
pixels in the PC mode of the WF/PC, and each pixel corresponds to about 16 km 
on the Martian surface. The fact that a second favorable opportunity would not arise 
for more than two years led Philip James and his co-investigators to ask for special 
consideration. They were granted the first General Observer time on HST. James will 
report on early results from the WF/PC and spectroscopic observations in this section 
of the conference proceedings. 



3.5 The Outer Planets 

Even though Jupiter, Saturn, Uranus and Neptune are more remote, the large scale 
of their atmospheric features recommends them as desirable HST targets. Table 1 
summarizes the anticipated resolution of features within these cloud decks. 

Table 1. The Resolution of the Giant Planets with the Planetary Mode of the Wide 
Field Planetary Camera 



Planet 



Apparent(l) 

Equatorial 

Diameter 



Pixels 
Subtending 



km/pixel 



Jupiter 

Saturn 

Uranus 

Neptune 



31.5 to 46.8" 

15.2 to 18.7 

3.5 to 3.9 

2.1 to 2.3 



735 - 1090 

353 - 435 

81 - 90 

50- 53 



196 - 132 

342 - 277 
646 - 581 
990 - 934 



(1) The equatorial diameters represent an annual range given in the Astronomical 

Almanac 

With temperatures less than 150 K , the thermal response times of the visible cloud 
decks of these planets are on the order of years and it seems that occasional sampling 



150 



would probably suffice to define their atmospheric state. Although lliib appeals lu In- 
true for Uranus, it is not the case for Jupiter, Saturn or Neptune. These planeLs have 
sizable internal heat sources, with Jupiter, Saturn and Neptune emitting 1.65, 1.82 
and 2.70 times more energy than the absorbed solar radiation. As a result, convective 
processes play a large role in determining the state of the visible cloud deck. The 
resulting cloud features are controlled by the strong zonal (east-west) wind patterns 
and planet-encircling phenomena can develop in days. Wave phenomena, which change 
on the scale of minutes and hours, are also present. Therefore, HST opportunities, 
when combined with the historical ground-based data sets, the short term coverage by 
Pioneer and Voyager flybys and the hoped for Galileo and Cassini observations, will 
enhance our understanding of these giant planets. 

Table 2. Observational Constraints for the Giant Planets 





Jupiter 


Saturn 


Neptune 


Revolutions of 

HST per Rotation 

of the Planet 


6- 7 


6- 7 


10 - 11 


Exposure Time(l) 
per Pixel of 
Smear (sec) 


10 


30 


450 


Dilution Factor(2) 


1.00 


0.30 


0.03 


Maximum Measurement(3) 
Accuracy (m/sec) 


4 


8 


14 



(1) based on observed rotation rates at low latitudes. 

(2) assumes distance from sun = semimajor axis of orbit. 

(3) assumes it is possible to reach one pixel accuracy in location 

of a cloud feature on 2 consecutive rotations of the planet. 

An unanswered fundamental question involves the degree of variability of the zonal 
winds as a function of latitude (Beebe and Youngblood, 1979; Ingersoll and Cuzzi, 1969; 
Ingersoll, et al. , 1981)). Conflicting models of the mechanisms that maintain the global 
circulation include a thin shell model, essentially a scaled model of the earth's atmo- 
sphere, and a deeper model that includes the convective envelope and assumes the 
winds are generated by the tendency for the rapid rotation to force rising convective 
cells into a cylindrical flow. The WF/PC in PC mode will allow short enough exposures 
to avoid smearing due to rotation of the planet while providing adequate spatial resolu- 
tion to identify small cloud features moving with the winds. Unhke the Voyagci- \ idei; 
cameras that were insensitive to red light, the WF/PC can provide observations with 
an 889 nm narrow-band filter that spans a wavelength interval dominated by methane 
absorption in the upper atmosphere. These data, when combined with other filters and 
selected spectroscopic observations, will supply information about the vertical structure 
of cloud systems. Table 2 summarizes the observing constraints for resolving longitudi- 
nally varying cloud structure on Jupiter, Saturn and Uranus. The period of revolution 
of Jupiter and Saturn are 11.86 and 29.458 years and the incHnation of the equators to 
their orbits are 3.12° and 26.73°, respectively. When the time scales associated with 
the seasonal aspects of the atmospheres of these two planets are considered, it is easy 



151 



to see that the data that can be acquired with HST will contribute a valuable archive 
for atmospheric studies. 

4. SELECTION OF TARGETS FOR EARLY DATA ACQUISITION 

The arrival of Comet Levy prompted early observations of a cometary target. These 
observations (see H. Weaver, et al. this proceedings) have shown that multispectral 
imaging in time steps on the order of hours can be used to determine the periods of 
rotation of comets. The brightening of the inner coma varies as active regions rotate in 
and out of our hne-of-sight and periodic behavior can be derived from time sequences. 

Sequences of observation which include multispectral imaging and ultraviolet spec- 
tra of Mars (see P. James, et al. this proceedings) and imaging of Titan (see J. Caldwell 
this proceedings) have been acquired and problems of deconvolution and calibration of 
the P6 chip in the WF/PC have been investigated. 

Early multispectral imaging of Jupiter, which utihzes all four chips of the PC mode 
of the WF/PC, has been acquired by the WF/PC team. Problems of flat-fieldiiig. 
deconvolving and mosaicking the data are being addressed. Tlu' initial icsmIin will In- 
discussed in the next section. 

In late September 1990, a major disturbance occurred in Saturn's equatorial region. 
Only two other major equatorial disturbances had ever been observed. They occurred 
in 1876 and 1933. The fact that the 1990 disturbance grew rapidly and the three events 
were separated by intervals of 57 years (two Saturnian years = 58.92 years) suggested 
that they might be seasonally induced convective disturbances. STScI responded to 
this event by granting a group of us, J. Westphal, W. Baum, R. Beebe, J. Caldwell, 
E. Danielson and A. IngersoU, a target of opportunity which allowed us to acquire 6- 
color imaging for two rotations of the planet, spaced to allow 20 to 30 hours between 
observations of the same portion of the cloud deck. Early results from these observations 
will be discussed in the next section. 

No observations of Neptune have been acquired. Ground-based imaging in red 
(619 nm) and near infrared (890 nm) by H. Hammel (1989) reveals the white clouds 
associated with the Great Dark Spot. In 1989, Voyager measurements established a 
rotational period of 18.33 h for the Great Dark Spot near 20°S latitude while the Voyager 
rotation period at 42°S was 16.76 hours (Hanomel, et al. , 1989). In comparison, Belton 
et al. (1981) derived a dominant period of 17.73 hours and secondary periods of 18.56 
and 18.29 hours from whole disk photometry. Whether this indicates that there was 
an additional large feature at mid-latitudes in 1981 or whether the Great Dark Spoi 
decelerated before 1989 by moving northward or by some other mechanism i.s noi known 
It does indicate, however, that changes occur, not only on a time scale ol minutes, hours 
and days (Smith, et al. , 1989), but also years. The resolution that HST tan obtain on 
this 2.2" disk would be adequate to monitor the Great Dark Spot, especially with the 
889 nm filter where the high white clouds would have a maximum brightness. 

No observations of Uranus have been attempted. The bland, near featureless cloud 
deck that was observed by Voyager 2 (Smith, et al. ,1986) relegates it to low priority. 
However, it should be noted that Uranus and Neptune cannot be assumed to be similar 
to Jupiter and Saturn. With masses more than five times smaller than Saturn and 
average densities 1.8 to 2.4 times greater, this pair of planets represents an intermediate 
type planet when compared to the terrestrial or jovian planets. In contrast with Jupiter 
and Saturn, Uranus and Neptune have westward equatorial winds. These differences 
make the more accessible Neptune an interesting target for HST. 



152 



5. INTERPRETATION OF THE DATA 



5.1 Jupiter 

When Voyager 1 and 2 spacecraft were at a range of 20 million km from the planet, 
about 20 days before nearest encounter, the narrow-angle camera images had a resolu- 
tion that was equal to the original specification of the PC mode of WF/PC. Figure 2 
shows a Voyager 1 view of Jupiter on the left that has a viewing aspect similar to the 
HST image on the right. The Voyager image was obtained on Jan 26, 1979, 38 days 
before closest encounter when the resolution was 350 km/pixel, half the optimum PC 




Figure 2. A Comparison of Voyager and Hubble Space Telescope Images. The 
Voyager 1 image was obtained on Jan 26, 1979 with a resolution of 350 km/pixel and 
the HST image was obtained with the F547M filter on Mar 11, 1991. 

resolution. The Red Spot is rotating off the visible disk and one of three white ovaLs. 
located near 30°S latitude, is west of the Red Spot. (These three ovals formed in 1938. 
from an overall increase in the reflectivity of this latitudinal region. As the white clouds 
evolved, they separated into three storm centers where the dark intervening regions were 
designated A-B, C-D, and E-F. The white storm centers retained this designation as 
they contracted and stiU bear the unhkely names of FA, BC and DE.) The oval in the 
Voyager image is BC, with FA located to the east and DE to the west (Beebe, et al. , 
1989). These ovals drift eastward relative to System HI ( 870.536°/day, Riddle and 
Warwick, 1976) at 2.6 to 5.5 m/sec (Beebe and Youngblood, 1979). Not only do the 
ovals catch up and pass the Red Spot in about two and a half years, but the spacing 
between them varies. The image on the right in Fig. 2 is a deconvolved, green image 
(F547M) that was obtained with HST on March 11, 1991, almost one Jovian year after 
the Voyager image. Again the Red Spot is at the limb, however, this time the white oval 



153 



is FA, 50° to the west of the Red Spot, while DE and BC are 66° and 87° east of the Red 
Spot, respectively. Although the turbulence to the west of the Red Spot appears very 
similar in the Voyager and HST images, this region experienced a general brightening, 
where the entire dark belt became white and featureless in June 1989, similar to the 
aspect during the 1973-74 Pioneer 10 and 11 encounters. When Jupiter emerged from 
conjunction in late summer 1990, the belt had returned to the Voyager- like appraraiuf 
seen in Fig. 2. The HST image serves as a test to determine how niaiiy small tddu"- 
are available as markers of atmospheric motion. Based on our experience with the 
Voyager images, the resolution is adequate to obtain a new map of zonal (east-west) 
wind velocities. Images separated by about 20 hours, will be used to derive zonal winds 
as a function of latitude. The results can be compared with a similar set of Voyager 
data, obtained one Jovian year earher (Ingersoll, et al. , 1981) during the same season, 
to determine the extent to which the zonal winds have changed. The WF/PC team has 
already obtained a preliminary data set which will allow this comparison. 




Figure 3. Ground-based Images of the Onset of the Saturnian Storm. The image on 
the left was obtained at 1:49:03 UT on Oct 4 and the image on the right was recorded at 
1:31:41 UT on Oct 7, 1990. Both images were obtained at New Mexico State University 
with a broadband blue filter and 0.2" sampling. 

5.2 Saturn 

On September 25, 1990, amateur observers reported a bright white spot in Sat- 
urn's northern equatorial region. Subsequent ground-based observations obtained at 
New Mexico State University (see Fig. 3) revealed that, although the storm expanded 
rapidly, it retained an identifiable nucleus that translated eastward at a rate of 400.25 
-|-/- 0.82 m/sec (Beebe, tt al. , 1991) relative to System III longitude (System III is 
based on the modulation of the radio signal and presumably the planet core. See the 
The Astronomical Almanac, Section E79, Gov. Printing Office for definition of this 
system.) This translation rate results in an eastward displacement of the storm svstem 
relative to the core of the planet of -32.884 +/- 0.067 deg/day. At li Ottubti U). 
1990, the nucleus was located at 212.1 -|-/- 1.6 deg. west longitude. This information 
was used to predict the location of the storm nucleus on November 9-18, when it was 
possible to develop an observing sequence for the planet as a moving target. 



154 



5.2.1. The HST Data Set 

The multispectral data set, obtained during the November 1990 Saturn Target of 
Opportunity, consisted of exposures with the P6 chip of the PC mode of the WF/PC. 
The planet was nearly centered on the chip and the plate scale was such that the ansae 
of the rings were not imaged (see Fig. 1). The sequence included images for two 
consecutive orbits on November 9 and 11, 1990. They were followed on Nov. 17 and 




Figure 4. Multicolor HST Images. These images were obtained with the F439W, 
F547M, F889N and F718M filters, clockwise from upper left. The original storm nucleus 
is visible. 

18 with a more extensive set of observations that included filters F439W, F547M, and 
F718M. Not only would these broad-band filters allow color composites to be gener- 
ated, but the short exposures ranging from 0.8 to 4.0 sec would also minimize rotational 
smearing and allow determination of wind speeds and wave motion in the atmosphere. 
Two narrow-band filters, F889N, centered within a strong methane absorption band, 
and F588N, a continuum reference, were selected to form a discriminator of vertical 
structure. A final filter, F336W, was included to provide insight into the aerosol prop- 
erties of the upper atmosphere. Data were obtained over 7 orbits of HST on November 
17 from Ih 36m to 13h 13m UT. On November 18 observations from Oh l'2\u Lu 131i 
27m, and an additional set at 22h 48m to 22h 58m, were obtained to iiiaxiiin/.c i lie luifil 



155 



time interval and better define drift rates of the clouds. Figure 4 illustrates the color 
dependence of clouds associated with the Saturn disturbance. This figure contains im- 
ages obtained with broad-band blue, green and red filters that show large differences in 
structure. A fourth image illustrates the extent of absorption by atmospheric methane 
at 889 nm. Because the rings contain no methane, there is a large difference in surface 
brightness between the globe and rings. This creates problems during the deconvolution 
process and reduces the use of this filter as an indicator of vertical structure. 

5.2.2 The Equatorial Wind Field 

Pairs of broad-band blue and green images separated by approximately 20 hours, 
containing the same large-scale cloud features, were registered and map-projected. This 
allowed controlled measurements of observed translations of cloud features which were 
converted to zonal and meridional wind speeds. Preliminary analysis of the data re- 
vealed that the rates of translations of clouds in the equatorial region were less than 
those at the time of the Voyager encounters (Smith, et al. , 1981; Smith d al. .1982). 




10 15 

ZONAL WIND (m/sec) 



HST 



VGR 



Figure 5. A Comparison of the Average HST and Voyager 2 Winds. These data 
are averaged zonal winds where all points within -I-/-0.5 degrees latitude are assigned 
equal weight and averaged. 

Figure 5 illustrates the differences. At the time of the Voyager flybys there was little 
longitudinally dependent structure in the equatorial region. The few visible features 
were quickly sheared apart by the latitudinal variation in the zonal wind. The data 
are plotted in Fig. 5 as a function of planetographic latitude (the angle formed when 
a normal to the local surface on an ellipsoidal planet intersects the equatorial plane of 



156 



the planet). 

Many of the features in the HST data are obviously associated with a planet- 
encircling wave pattern that has been generated as a consequence of the initial distur- 
bance. Due to the fact that small wisps, not well-formed cloud systems, were measured 
in the Voyager images, an obvious interpretation of the discrepancy between the two 
data sets in Fig. 5 is that the Voyager 2 data represent the unperturbed zonal wind, 
and the difference between the HST and Voyager translation rates represents the phase 
velocity of the equatorial waves. If this is the case, the storm would be bounded on the 
poleward side by an eastward wind with a relative speed of about 75 m/sec. The tact 
that ground-based observations yield no evidence of additional convectivc distui haiues 
similar to the initial disturbance (Beebe, et al. , 1991) supports llns inui |iui ,ii lun 
Thus, we assume that our analysis of the cloud structure is justified. 

5.2.3. Wave Analysis of the Cloud Structure 

Broad-band green images from November 17, 1990, selected to span all longitudes, 
were processed to remove the limb darkening and then map-projected. The data were 
obtained over a period of 9.85 hours during seven HST orbits as the planet rotated in 




Figure 6. The Amplitudes of Components of Fourier Series. The amplitudes were 
obtained by fitting the variation in surface brightness in bins spanning 1 degree latitude 
and 360 degrees longitude. Both phase and amplitude were free parameters. 

front of the camera. In order to reduce these images to some instant in time, the HST 
wind profile in Fig. 5 was utilized to correct each image for the latitudinal distortion 
that would occur over the intervening time interval. All seven images were remapped 
to reduce them to the time when the fourth image was obtained. This resulted in 
shifts as large as 6.7 degrees in the equatorial regions and half that at 20°N in images 



157 



1 and 7. The mosaic of the map-projected images, corrected to a standard time, was 
achieved. In regions where the maps overlapped, we selected the image that minimized 
the incident and emergent angles, reducing residuals from removal of limb darkening 
with the Minnaert function. 

The resulting mosaic spans 360 degrees with a resolution of 0.5 degrees in longitude 
and latitude. Each latitudinal bin, 0.5 degrees wide, was fit with a Fourier series that 
contained up to 36 terms, where the phase and amplitude of each term was calculated. 
The 3-D plot in Fig. 6 illustrates the results for wavenumbers less than 27 (representing 
standing wave structure with wavelengths greater than 13 deg. longitude). Here the am- 
plitude, which corresponds to surface brightness is plotted as a function of waven umber 
on the X-axis and latitude on the y-axis. 

Figure 7 is the corresponding mosaic. The dominance of wavenumber 2 is caused 
by the two bright regions that are bracketed along the top of Liie niusaic in Fig. 7. 
while the reality of the enhanced amphtudes for n = 6 is illustrated by the arrows that 
show a repeating pattern in the cloud structure. Although there are definite wave-hke 
structures near 15''N latitude, there is no evidence from this analysis that they form a 
standing wave. Instead, inspection of these latitudes on map-projections separated in 
time by about 20 hours reveals that these features are behaving as expected. There is a 
strong latitudinal gradient in the zonal winds and they are responding to this shear by 
twinning and recombining. Thus, although they tend to span about 15 degrees, they 
do not constitute a standing wave. 




Figure 7. An F547M Map-Projection. These data were obtained on Nov 11, 1990, 
and show the cloud structures that yield the enhanced n = 2 and 6 amplitudes. The 
original storm nucleus is indicated by the first arrow on the left. 

This preliminary analysis indicates that wavenumber 6 characterizes the gross wave 
structure. Surprisingly n = 6 is present in the polar hexagon at 76°N latitude (Alli- 
son, Godfrey and Beebe, 1990). The significance of this value is not obvious and there 
is much to be done to understand the November 1990 data set. To enhance our un- 
derstanding of the equatorial region's response to a convective disturbance, additional 
observations of its dying stages and the subsequent recovery of the atmosphere are 
needed. 



6. PREVIEW OF OBSERVATIONS TO COME 



158 



Already the WF/PC team has acquired multispectral imaging of Jupiter that will 
provide complete longitudinal coverage of the planet as well as a preliminary check on 
the stability of the zonal winds. In June 1991, they acquired 3 multispectral data sets 
of Saturn that spanned 6 HST orbits, providing a post-conjunction map of the wave 
structure of the abating disturbance. These data are now being reduced and analyzed 
and indicate that no additional disturbances have occurred. 

In addition to these preUminary observations, systematic sets of observations that 
combine imaging and observations with the FOS or HRS to define temporal variations 
within the Martian, Jovian, Saturnian and Neptunian atmospheres are desirable. Be- 
cause these bodies span 15" to 48" and have no extreme contrasts across the visible 
disks they can be deconvolved into useful data. In addition, prevailing zonal winds tend 
to cause molecular variations to be latitudinally dependent and lo extend u\(i laigf 
enough areas of the planet that spectra obtained with arcsecond apeitures will ie\t^al 
differences in molecular concentrations. When these observations are coiiibiued willi 
multispectral imaging, they will provide insight into couphng between the troposphere 
and stratosphere of these planets. Limited observations of this sort are planned for 
Cycle 1 and later. 

The usefulness of HST observations to investigate the dynamical properties of 
comets has been demonstrated. Unfortunately, the reluctance of comets to announce 
their arrival dates wiU tend to require that we be granted Targets of Opportunity to 
observe them. In a Uke manner, the abrupt development of Saturn's equatorial storm 
and the speed with which Jupiter's belts can mix and change dramatically will cause the 
planetary community to be strong contenders, along with novae watchers, for director's 
discretionary time. 



7. CONCLUSIONS 

The aberration associated with the mirrors has made it more difficult to acquire 
spatially resolved multispectral imaging and has increased exposures with the FOS 
and HRS. However, we have shown that HST will contribute vitally needed long- term 
data sets that will provide insight into the structure, dynamics and energy balance of 
planetary atmospheres and comets. There are still problems associated with reducing 
the effort necessary to obtain the data and to optimize their quality. The urgent problem 
of tracking moving targets, will be solved by completing planned software enhancements 
at the Institute. Other factors, such as improving our capability to flat-field the WF/PC 
data are being worked on. We, the planetary community, look forward to utilizing this 
facihty for acquiring unique observations, as well as systematic acquisition of data sets 
that will enhance our understanding of the temporal variabihty within the Solar System. 

ACKNOWLEDGEMENTS 

I thank J. Westphal and R. Light for their efforts to deconvolve the data, A.S. 
Murrell for his dedication to obtaining ground-based observations to define the storm 
for targeting HST and C. Barnet, L. Huber and P. Sada for their support in reducing 
and analyzing the Saturn data. Ground-based observations used for targeting and 
interpreting this data were supported by NASA grant NAGW-1802. 



159 



REFERENCES 

Allison, M., D.A. Godfrey, and R.F. Beebe 1990, Science, 247, 1061. 

Beatty, J.K. and A. Chaikin 1990, The New Solar System (Sky Publishing Corporation, 

Cambridge, Mass.) 298-291. 
Beebe, R.F. and L.A. Youngblood 1979, Nature, 280, 771. 
Beebe, R.F., G.S. Orton and R.A. West 1989, in Time Variable Phenomena of the 

Jovian System, NASA SP-494, ed. M.S. Belton, R.A. West and J. Rahe, (U.S. 

Government Publication) p. 245. 
Beebe, R.F., C. Barnet, P.V. Sada and A.S. Murrell 1991, Submitted to Icarus. 
Belton, M.J., L. Wallace and S. Howard 1981, Icarus, 46, 263. 
Hammel, H.B. 1989, Science, 244, 1165. 
Hammel, H.B., R.F. Beebe, E. M. De Jong, C.J.Hansen, C. D. Howell, A. P. Ingersoll, 

T.V. Johnson, S.S. Limaye, J. A. Magalhaes, J.B. Pollack, L.A. Sromovsky, V.E. 

Suomi and C.E. Swift 1989, Science, 245, 1367. 
Ingersoll, A. P. 1990, in The New Solar System, ed. J.K. Beatty and A. Chaikin (Sky 

Publishing Corporation, Cambridge, Mass) p. 139. 
Ingersoll, A.P. and J.N. Cuzzi 1969, J. Atmos. Sc, 26, 981. 
Ingersoll, A.P., R.F. Beebe, J.L. Mitchell, G.W. Garneau, G.M. \agi and J. Mullci 

1981, J. Geophys. Res., 86, 8733. 
Lucy, L.B. 1974, A. J., 79, 745. 

Riddle, A.C. and J.W. Warwick 1976, Icarus, 27, 457. 
Smith, B.A., et al. 1981, Science, 212, 163. 
Smith, B.A., et al. 1982, Science, 215, 504. 
Smith, B.A., et al. 1986, Science, 233, 43. 
Smith, B.A., et al. 1989, Science, 246, 1422. 



160 



OBSERVATIONS OF MARS USING HUBBLE SPACE TELESCOPE OBSERVATORY 

Philip B. James, Univ. Toledo, R. Todd Clancy and Steven W. Lee, 
Univ. Colorado, Ralph Kahn and Richard Zurek, Jet Propulsion 
Laboratory, Leonard Martin, Lowell Observatory, and Robert 
Singer, Univ. Arizona. 

1. INTRODUCTION 

The lack of a continuous record of martian meteorology or of 
volatile cycles on Mars for extended periods of several martian 
years seriously hinders efforts to understand the physics of the 
martian atmosphere - surface system. Despite the fact that Mars 
has received relatively intense scrutiny from spacecraft, these 
observations are limited to only a few isolated time periods; 
and, inasmuch as these missions were primarily interested in high 
resolution geology, more than a small fraction of the planet's 
surface was rarely covered during a particular time period. The 
earth based record is limited by the relatively short periods 
surrounding oppositions when telescopic observations can yield 
useful data. Because of the incommensurability of the orbital 
periods of Earth and Mars, the martian season seen during one 
opposition will not be observable again for eight martian years. 

The lack of a continuous synoptic record of the planet is 
a serious impediment to understanding the martian weather and 
climate (for an up to date review of martian phenomena consult 
Mars, U. Arizona Press) . For example, the martian global 
duststorms are a meteorological phenomenon which seems to be 
unique to Mars. Large storms were observed from Earth in 1956, 
from Earth and Mariner 9 in 1971, from Earth in 197 3, and by 
Viking twice in 1977 and in 1982. There is no documentation for 
any other global duststorm event, leading to the facetious 
hypothesis that "spacecraft cause duststorms." A more reliable 
record is needed in order to determine the nature of a longer 
term cycle as well as to establish mechanisms which can lead to 
such events in some years but not in others. Our incomplete 
knowledge of the temporal distribution of major dust storm events 
on Mars is the best known consequence of the lack of such a 
record, but the situation is much the same for interannual 
variability in the behavior of surface condensates in the polar 
regions and for the behaviors of clouds in different years. 

To test one possible technique for remedying this situation, 
we have embarked on a three year program of Mars observations 
using the Hubble Space Telescope. Although the solar pointing 
constraint eliminates 45% of the 780 day synodic cycle from 
possible observation, this is still a great improvement over the 
two to five months of each cycle (depending on the orbital 
geometry of the two planets at opposition) that can be profitably 
be used for earth based observations. During the initial phase 
of the project we have imaged Mars on five dates in a variety of 
spectral bands; the observations completed during Cycle are 
listed in the following table: 



161 



01-02-91 


359 


13.5" 


21.7 


kiti/px 


300 


02-07-91 


16 


9.4" 


31.1 


km/px 


300 


03-20-91 


35 


6.6" 


44.3 


km/px 


300 


05-15-91 


60 


4.8" 


60.9 


km/px 


190 

300 

70 



Date Ls Size Scale LCM Filters 
(deg) (arcsec) (km/px) (deg) 

12-13-90 349 16.5" 17.7 km/px 190 890N, 673N, 588N, 502N 

300 439W,336W,230W 
70 

673N,413M,FOS 

673, 413, 336, 230, EOS 

673N,413M 

673N,413M 
502N,336W,230W 

Due to the fact that the Solar System Target software was 
not yet active, these observations were entered as fixed targets; 
the co-ordinates of Mars at the exact time of the observations 
were required for this. Much of the credit for the successful 
scheduling of this program can be traced to Marc Buie who 
performed these calculations. Most of the integration times were 
very short, less than 1 second. Therefore, scans were generally 
not needed despite the large apparent motion of Mars. In fact, 
three exposures were often made on a single target, producing a 
drift of -100 pixels during the time interval between exposures. 
The 2 minute F2 30W exposures did require a scan, which was 
successful in all cases. 

All of the Mars exposures were successful except for the 
first three targets which were to have imaged the 190 central 
meridian. These images were located at the outer edge of PC 6 so 
that only 50% - 33% of each exposure was recovered. Subsequent 
December exposures were recovered by upl inking a command to HST 
in real time to adjust its pointing by 15"; the success of this 
remedy, without which the December data would have been 
worthless, illustrates the advantage of having investigators at 
the Space Telescope Science Institute during observations. The 
error was ultimately traced to an aberration correction which had 
been inserted at two different steps in the preparation of 
instructions for HST. 

Unfortunately, attempts to acquire images of a solar type 
star for photometric calibration and for point spread function 
proved to be much less successful. Images of HD23169, a G2V 
star, were acquired in December but were underexposed and 
therefore not useful for PSF's. Analysis of the failure was 
inconclusive; the images seemed consistent with a star which was 
one magnitude fainter than catalog values. The second attempt 
failed in February when a cosmic ray event in the South Atlantic 
Anomaly caused a safing. The final attempt on HD23169 failed in 
March when a properly exposed image of the star was found on PC 5 
rather than PC 6. 

162 



2. SCIENTIFIC OBJECTIVES 

The scientific interests of the observing team focus on the 
atmosphere of Mars and the interactions between the atmosphere 
and surface of the planet. The primary objective of the 
investigation is to monitor seasonal changes on the planet 
through as much of its annual cycle as possible. As noted above, 
Hubble Space Telescope is a potentially valuable tool for 
monitoring Mars. Near oppositions, the expected scale of HST 
images was comparable to that of the Viking approach images which 
provided resolution of the martian surface similar to a 
terrestrial weather satellite (Figure 1) . Even more important, 
the resolution expected for images acquired when Mars is near 
solar conjunction was comparable to Planetary Patrol images 
acquired near oppositions (Figure 2) . Therefore, except for the 
solar elongation constraint, HST could reliably monitor 
conditions on Mars throughout the seasonal cycle. 

More focused scientific goals for the project include: 

1. Observing the surface albedos of various units on the planet 
and seasonal variations in these albedos. The shapes and albedos 
of various surface units of Mars change as a function of time. 
Part of the variation is seasonal; originally ascribed to polar 
melting and to possible vegetation, these changes are now thought 
to be a result of shifting dust cover on the planet. Part of the 
variation is not seasonal but responds to changes over a longer 
time scale, perhaps related to a longer dust storm cycle. We 
have intentionally chosen the Syrtis Major area, where some of 
the most prominent albedo features occur, for frequent 
monitoring with HST in order to attempt to quantify the albedo 
changes which occur and to separate seasonal and longer term 
variations. 

2. Observations of dust storms on Mars. As noted above, the dust 
cycle on the planet is still relatively unknown. Dust phenomena 
on the planet span the range of scales from dust devils to huge 
storms that completely envelop the planet. Many intermediate 
scale storms, 100-1000 km in size, were observed by Viking 
orbiters. Only the larger storms are generally picked up by 
terrestrial observers, and the statistics on these are modulated 
by the 15 year opposition cycle acting on the dust storm season, 
which spans the spring-early summer seasons in the south. 
Frequent monitoring of regions with the potential to initiate 
dust storms by HST during the next dust storm season, in HST 
Cycle 2, is proposed as a continuation of this work. 

3. Measurement of atmospheric opacities. The background opacity 
of the atmosphere is affected by both condensate hazes and by 
aerosols. The opacity observed by Viking varied substantially 
during the mission being affected mainly by the dust cycle. 
Observed surface reflectances will be analyzed using a radiative 
transfer model in order to derive optical depths and properties 
of scatterers in the atmosphere. 

4. Measurements of seasonal and interannual variations in the 
distribution of ozone on Mars. Ozone is a trace ingredient in 

163 




Figure 1: This approach image was taken by Viking Orbiter 1 
shortly before its orbit insertion in 1976. The scale of the 
approach images is similar to that which was expected for HST 
observations of Mars near favorable oppositions. These images, 
similar in scale to terrestrial weather satellite views, are 
ideal for synoptic monitoring of the planet. Photo credit: NASA. 



164 




Figure 2 : Th 
Kea in 1986 
the angular 
when Mars is 
its synodic 
this image, 
sees in May, 
despite the 
and National 



is photographic image of Mars was acquired from Mauna 
during the excellent opposition of that year, when 
size was about 24 arc sees. WFPC images acquired 

only 3-5 arc seconds, which is true during most of 
cycle, are expected to have better resolution than 

Deconvolved WFPC images obtained when Mars was 4.8 

1991, have proven that this is still the case 
optical problems. Photo credit: Lowell Observatory 

Geographic Society. 



165 



Mars' atmosphere which displays substantial geographic and 
temporal variability. Ozone is destroyed through chemical 
reactions involving the OH radical, so ozone concentration is 
(anti) correlated with water vapor abundance. The driest regions 
of the planet, e.g. the winter polar regions, are therefore 
places where maximum ozone is expected. Ozone can be detected 
through absorptions near 230 nm. FOS scans of the planet were 
performed to provide spectra in this region which can be used to 
determine the concentration of ozone molecules in various regions 
of the planet. In addition, comparison of WFPC images using the 
230W and 336W filters is being used to attempt to map ozone on 
the planet. 

5. Multispectral mapping of surface units on Mars. The 
reflectance spectra of the surface of Mars contain features which 
are diagnostic of the minerals which are present on the surface. 
Comparisons of the spectra of various surface units contributes 
to understanding the geological history of the planet. Though 
the set of filters used does not provide the spectral resolution 
available using other techniques, the surface resolution possible 
is generally greater than in other experiments. The HST results 
will be used in concert with other data sets to map the 
compositions of surface units on the planet. 

6. Seasonal changes in polar caps and polar hoods. Numerous 
polar regressions have been observed by earth based astronomers 
and by spacecraft. It has been shown that the polar regressions 
in different martian years are different, and there has been some 
speculation about possible relationships between these variations 
and those associated with the dust cycle. HST will enable 
determinations of polar cap boundaries during years in which 
these data cannot otherwise be obtained. In addition, monitoring 
in red and blue filters will permit separation of the atmospheric 
hood from the surface cap; though the hood is possibly the most 
dynamic phenomenon on the planet, it is one of the least 
understood. 

7. Observation of diurnal and seasonal development of clouds. 
Condensate clouds occur in various regions of the planet. These 
are often diffuse hazes which appear near the limbs due to 
condensation in a cold atmosphere. There are regions which 
display discrete, optically thick clouds: Tharsis, Elysium, and 
Hellas are examples. The distribution of these clouds seems to 
be determined by topography, but there are substantial variations 
from year to year. HST will provide ample resolution to document 
these clouds at all times in Mars' synodic cycle. Limited 
diurnal data can be obtained by imaging on consecutive orbits. 



166 



3. PROCESSING WFPC IMAGES OF MARS 

The scientific benefit of scheduling our GO program so early 
in Cycle was only slightly negated by various practical 
difficulties associated with being the first GO program. The 
images were to receive initial processing which included flat 
fielding. The images which were emitted from the end of the 
processing "pipeline" in some cases had as many apparent defects 
as the raw images, and in all cases the images contained 
blemishes due to dust particles on components of the optical 
system. Post launch flat fields taken on PC 6 with the various N 
series filters which were used to image Mars in the visible and 
near infrared portions of the spectrum did not exist, and in some 
cases there were not even any pre launch flats available. The 
default for the processing software in the absence of any flat 
was to divide the raw image by a unit image and to proceed as if 
the image had been properly flatted. Valid flat field images 
were acquired for several of the filters which we used through 
the generosity of members of the WFPC team. A raw, unflatted 
image of the Syrtis Major region acquired using F588N is shown in 
Figure 3a. 

Even images which had been divided by the appropriate flat 
field had residual blemishes due to dust specks on the pyramid 
and filters. These appeared as bimodal light-dark, nearly 
circular blotches roughly 5-10 pixels in diameter. They occurred 
in regions where there were large intensity gradients in the 
image, especially near the limbs of the planet. Particular care 
was needed to remove these blemishes since their scale is similar 
to the features on the martian surface which are of interest and 
since the deconvolution routines will further extend their 
influence. An example using the 588N filter of an image which 
has been flat fielded is shown in Figure 3b; the image has been 
greatly stretched to reveal the blemishes near the limb. 

The method used to remove the blemishes is as follows: 

1. The image is carefully compared to the flat field image to 
make certain that the blemish to be removed is indeed a residual 
of the optical system rather than a feature on the planet. 

2. A square image generally 20-25 pixels on a side centered on 
the blemish is extracted from the main image. 

3 . The IRAF routine IMSURFIT is used to create an image from a 
two variable polynomial fit to the unblemished border of the 
extracted image. Generally a third or fourth order polynomial 
has proved to be sufficient to give a consistent image. 

4. The polynomial fit is not adequate to replace the undesired 
portion of the image because it is mathematically smooth and 
stands out if used to replace the blemish. The IRAF routine 
MKNOISE is used to make a suitably noisy image out of the 
polynomial fit. 

5. The "patch" is reinserted into the image to replace the 
blemish. 

Inasmuch as most of the blemishes occurr in the relatively 
featureless limbs of the planet, we do not believe that this 

167 




Figure 3a. This WFPC image of Mars through the 588N filter has 
not been flat fielded. In addition to a blocked column, numerous 
blemishes caused by dust on the pyramid and filter are present; 
inasmuch as their scale is similar to the surface features on 
Mars that are of interest, their careful removal is essential. 



168 




Figure 3b. After flat fielding, blemishes still remain in regions 
of large intensity gradient; this can be seen in the 588N picture 
of Figure 3a which has been stretched to bring out these features 
near the limb. This image also reveals that the point spread 
function has greatly extended the limbs of the planet; this 
effect makes determinations of limb profiles suspect. 



169 



cosmetic procedure has any detrimental effect on the validity of 
the scientific data contained in the images. The result of this 
process for the 588N Syrtis image is shown in Figure 3c. 

Because the three attempts to acquire images of a solar type 
star on PC 6 to use as a photometric calibration and as a point 
spread function failed to provide the necessary data, computer 
generated point spread functions supplied by James Westphal were 
used to deconvolve the images. The implementation of the 
Richardson Lucy deconvolution procedure which is contained in the 
X version of the STSDAS package was used to deconvolve the 
images. This routine assumes that the point spread function is 
constant across the image. This is clearly not true for the 
early Mars images which are 400 pixels in diameter. However, 
except for the first set, the images are centered in the chip; 
and star field images suggest that large distortions appear 
mainly near the edge of the chip. 

Forty to sixty iterations of Lucy were performed on the 
images. Less iterations led to reductions in resolution while 
more iterations produced little perceptible increase in 
resolution but made the images noticably more noisy. Estimation 
of resolution is somewhat subjective since the visibility of a 
surface feature depends on albedo contrast as well as on size; 
even a large crater is invisible if there is little contrast 
between it and the surrounding terrain. Surface features on the 
planet suggest that the sub earth resolution is at least 50-75 
km, i.e slightly better than 0.2" The Richardson Lucy routine 
therfore restores the resolution to within a factor of two of the 
actual surface resolution which might have been expected in the 
absence of spherical aberration. The final, deconvolved 588N 
image which has been used as an example herein is shown in Figure 
3d. 

Color composits are useful scientifically in identifying 
yellow dust clouds which have only small contrast in brightness 
at individual wavelengths. However, the public appeal of color 
composits is probably their greatest asset. One of the first 
tasks facing the team was production of such a color image from 
the available data. 

There are two potential problems in color compositing: 
registration and color balance. Both of these proved to be 
present in producing the color image which appeared in Life, 
Astronomy, Sky & Telescope, etc. Registration was a problem 
because Mars rotated perceptably between the exposures. The 
exposures which were used for blue and green were obtained on the 
same fixed target and were only slightly displaced from each 
other in time; the 889N image, used for red, was exposed on a 
different fixed target and was displaced by a greater amount from 
the other two. Registration at the limbs of the identically 
sized images would lead to fuzziness and color halos at the 
boundaries of major albedo features, such as Syrtis. 
Registration on the albedo features, which was finally used in 
the published composite, leaves a color halo around the limb of 

170 




Figure 3c. The blemishes have been cosmeticly removed from the 
image of Figure 3b using techniques described in the text. 
Without further processing, the resolution of this image, 
acquired when Mars was 16.5 arc seconds, is similar to Planetary 
Patrol photographs. 



171 




Figure 3d. The same image in as in Figure 3c is shown after 4 
iterations of the Richardson-Lucy algorithm. The surface 
resolution has been restored to within a factor of two of what 
was originally anticipated. 



172 



the planet. 

Availability of flat fields for the various filters forced 
us to use 889N as red and 588N as green. Because the surface 
reflectance of Mars is quite steep between 500 and 600 nm, the 
latter situation caused considerable problems in attaining 
correct color balance. Mars is considerably brighter at 588nin 
than at 502nin, and the contrast between light and dark areas is 
much greater in the "yellow" filter than in green. The result is 
that the image is somewhat greener than most color images of the 
planet. Subsequently, we have composited images from May using 
673N for red and 502N for green, and the color in the resulting 
image is much closer to what is expected for Mars. 



173 



4. PRELIMINARY RESULTS 

Preliminary analysis of the images obtained by HST during 
the five observation sequences has rewarded our optimism 
concerning the potential scientific value of HST for monitoring 
martian phenomena. Even the 4.8 arc second images acquired in 
May have sufficient resolution to reveal details of the albedo 
boundaries on the surface. The images taken in May, when Lg 
equaled 60 , clearly reveal the multicomponent "W" clouds in the 
Tharsis - Valles Marineris region as well as clouds associated 
with Elysium Mons. They also show the north polar surface cap, 
which is tilted earthward at this season; the coverage from three 
different central meridians will permit a detailed comparison of 
thee boundary of the cap with excellent spacecraft data at the 
same L_ from Mariner 9 and Viking for three martian years as well 
as with earth based regression data. The potential scientific 
value of these 4.8 arc second images augurs well for the 
investigation of the 1992 "classic dust storm" season which will 
use the similar scale images planned when Mars again emerges from 
the 50 solar interdict. 

Processed images from December, January, February, and March 
are shown together in Figure 4; all of these were taken throught 
the 673N filter and were processed using a flat field for the 
"nearby" 656N filter. The use of the flat for the spectrally 
adjacent filter with a similar bandwidth does a reasonable job, 
at least superficially. This is probably because the size and 
structure of blemishes produced by dust in the optics is mainly 
wavelength dependent. The 656N flat did have an artifact not 
apparent in our 673N images which may be due to a pinhole leak in 
the former filter. The prominent dark feature which looks like 
an inverted map of Africa is Syrtis Major; this area is thought 
to slope from the bright Arabia region on the west to the Isidis 
impact basin on the east. The region is probably dark because of 
slope induced winds which scour dust from the surface. Syrtis is 
a region which exhibits many wind induced streaks which support 
this hypothesis. The bright region to the south of Syrtis is the 
huge Hellas impact basin. This is often the site of dust or 
condensate clouds, and during southern winter the basin is 
covered with bright carbon dioxide frost. The dark, east to west 
arc which bisects Hellas is a relatively new albedo feature which 
has not appeared on most past albedo maps; likewise, the dark 
knob to the west of Hellas is darker than it usually appears. 

Figure 5 displays the corresponding blue images of the Syrtis 
region. Surface contrasts are reduced in blue, and Syrtis Major 
is barely visible. The filter 413M, which had been selected as 
our "blue" filter was replaced in December only by 439W because 
the CCD chips had not yet had a UV flood; the sensitivity of 413M 
was more suspect than that of the W series selection. The early 
pictures show the north polar hood prominently as well as 
extensive clouds to the north of the south polar region. This 
season, which is near the equinox, is historically one of the 
least active times on the planet from a meterological point of 

174 




Figure 4: Four images of the Syrtis Major face of Mars acquired 
in December, January, February, and March when the angular 
diameters were 16.5, 13.5, 9.4, and 6.6 arc sees respectively. 
The images, which used the 673N filter, are shown at the correct 
relative scale. Even the smallest size images reveal detail 
comparable to very good terrestrial photographs, making them 
scientifically useful. Credit: STSCI/NASA and Univ. Colorado. 



175 




Figure 5: These are the blue filter (F413M) images corresponding 
to the red images in Figure 4. The shorter wavelengths show 
primarily atmospheric features on the planet, and surface 
contrast is much reduced. Credits: Same as Figure 4. 



176 



view. 

The ultraviolet imaging capabilities of HST provide a unique 
opportunity to study the surface and atmosphere of Mars in this 
relatively unexploited wavelength region. In particular, the 
strong ozone absorption near 230 nm makes it possible to map 
ozone concentration through differencing the 230 and 33 6 nm images. 
Preliminary use of this method reveals strong ozone absorption in 
the north polar region during late winter, as expected from the 
low water vapor content in the atmosphere at that time, and 
reveals other interesting correlations with various topographic 
and surface features. The differencing technique will be verified 
and calibrated using spectral scans of the planet in the ultra 
violet portion of the spectrum using the Faint Object 
Spectrograph (FOS) ; the latter data have been inspected to verify 
that signal to noise is as expected, but detailed analysis of 
those data has not yet been undertaken. Inasmuch as Mars is much 
brighter at wavelengths in excess of 600 nm than at ultra violet 
wavelengths, red leaks in the ultraviolet filters could easily 
jeopardize this part of the experiment. Inspection of the 230N 
images strongly suggests that the red leak is no worse than 
indicated in Figure 4.7.3.1 of Version 2 . 1 of the WFPC Instrument 
Handbook. There is no apparent residual of the albedo patterns 
which would be produced by exposure to the longer wavelengths. 

The two sequences of pictures in Figures 4 and 5 clearly 
illustrate that our expectations regarding the potential value of 
these images for monitoring Mars have been confirmed. Even the 
smallest images, acquired in May (not shown here) show detail 
comparable to Planetary Patrol images acquired at times of 
opposition. HST's Cycle 2 will encompass a major portion of the 
next dust storm season, and frequent monitoring of the planet has 
been proposed for that period. That Cycle will also afford the 
opportunity of revisiting the same seasons imaged during the last 
few months in order to search for existence and causes of 
interannual variability. Cycle 2 will provide monitoring of Mars 
leading into the Mars Observer Mission. We hope that, using the 
new instrumentation to be installed on HST in 1993, it will be 
possible to make observations which will complement the 
experiments to be conducted on that mission. 

We wish to express appreciation to the large number of people at 
Space Telescope Science Institute whose assistance has been 
invaluable. We especially thank Ed Smith, who has worked with us 
in our attempts to overcome data analysis problems since the 
first day of the project. We also appreciate the assistance of 
James Westphal, who supplied much needed flat fields and psfs 
early in the project. 

Reference 

B. Jakosky, H.H. Kieffer, M. Matthews, and C. Snyder editors, 
MARS, University of Arizona Press, Tucson (1991) . 



177 



DECONVOLUTION AND PHOTOMETRY ON HST-FOC IMAGES 



C. Barbieri, G. De Marchi, R. Ragazzoni 
Astronomical Observatory of Padova 
Vicolo dell'Osservatorio, 5 
35122 Padova, Italy 



Abstract. Due to the peculiar characteristics of the PSF of HST, a careful analysis with 
many deconvolution experiments must be performed on HST-FOC images in order to 
understand their property. We briefly present some aspects of our work on the subject. 

1. INTRODUCTION 

The strongly aberrated PSF of HST requires a large amount of effort in the field of 
image deconvolution, under conditions quite unusual if compared with those found in 
ground-based optical and radio data analysis. 

In fact, generally from the ground and before HST the equivalent PSF of an astro- 
nomical instrument (telescope plus atmosphere) was known only with a rough approx- 
imation, due to the stochastic behaviour of the atmosphere itself and of its perturba- 
tions. On the other hand, the reconstruction of images taken with radiotelescopes is 
fundamentally based on deconvolution techniques; in this latter case, though, the PSF 
is known with high accuracy (because it essentially coincides with the instrumental 
beam) and the collected raw images are characterized by a high signal-to-noise ratio 
(SNR). 

In the HST case, a rather complex PSF is known with a good degree of accuracy 
and it seems to be fairly stable (apart from human modifications, like those arising from 
the changes of the focus setting, and so on). However, most of the images (especially 
FOC's) are strongly photon-Umited, i.e. characterized by a high poissonian noise. 

2. FOC FRAMES DECONVOLUTION USING CLEAN 

Using IDL as a framework (because of its flexibility), two different implementations 
of the CLEAN algorithm were performed and extensively tested: a standard one, based 
on Hogbom (1974) and Segalovitz et al. (1978), and an enhanced version, capable also 
to treat extended objects, based on Steer et al. (1984) and Wakker et al. (1988). 

In order to test the efficiency of the procedure and to see what happens inside the 
procedure, a real time control of the behaviour both of the raw and of the cleaned image 
is allowed in our implementation. 



178 







"■ <■■■; ■ ■ ■ 


'•::>.\^-*fi.'--r.::- 


■■ '■^ ^■ 










.-^ 










; n--^. 


■ ■ y- 






l^ 


,' ■ 


■V ■ v' ^ / 




'l ■ 






' f*' 


: ^'•:.':f 






! s 











^ 












» 




• 




* 


* 

• 


•• 


* 








^ 


• « 


« 


•• 








* 




' 


♦ 


• 


8- 







Figure 1: Action of Clean on an HST-FOC frame of R136a. See the text for an expla- 
nation. North IS bottom-left to top-right. East on the right. 



179 



The operator, looking at the display, can interactively see the cursor moving from 
a point to another, indicating where objects are found and subtracted in the raw frame 
and simultaneously added in the cleaned one. 

In other words, the growth of the stars in the cleaned frame and their corresponding 
disappearence in the raw one can be inspected during the execution of the procedure. 

In our opinion, this step is of great importance, because programme execution 
can be thus interactively checked, avoiding the use of a black-box like procedure in 
deconvolution. 

The typical run of this algorithm takes few minutes to some hours, depending on 
the choices for the loop gain, on the threshold and on the complexity of the treated 
image. In this case, the interactive display procedure becomes unuseful and so the 
option can be conveniently switched off. 

Figure 1 shows one of the first results obtained on R136a. The raw image is treated 
with Clean, producing a set of locations where peaks arising over a certain threshold 
have been found. Each location is characterized by both positional and intensity infor- 
mation, since the latter is retrieved by subtracting the PSF from the raw frame and 
adding back the residual. As a matter of visualization, the final cleaned image is pro- 
duced convolving the location data with the nominal PSF characterized by a FWHM 
typically of the order of the expected HST diffraction limited performances. 

3. USING CLEAN TO SUBTRACT SINGLE UNRESOLVED SOURCES 

In the case of a frame containing few unresolved sources, such as bright foreground 
stars, not scientifically interesting, Clean can be forced to remove them and their dele- 
terious halo, even if the star peaks themselves are slightly saturated. 

Actually, FOC is a photon counting device and the occurrence of the saturation 
effect can become so heavy that one of the statements which Clean techniques are 
based on {i.e. the fact that the height of the peak is proportional to the brightness of 
the star as it is for its optical halo) can fail. 

In such cases (in practice easily detectable by simple inspection of the raw frame) 
Clean can be modified using, instead of the peak value, the average properly scaled 
value of a group of pixels, faUing within an annulus surrounding the star. 

This procedure, applicable to non crowded fields, allows one to perform subtractions 
of single stars in a very efficient way. Actually, the positional information is retrieved, as 
always, using the location of the peak, while the brightness information can be obtained 
fitting a zone of the PSF free of saturation. 

Moreover, the effects of a sUght blurring or of the oversamphng can become negli- 
gible, thanks to the fact that peaks locations are allowed to vary during the numerous 
iterated subtractions in order to match the shape and the sub-pixel position of the star 
to be subtracted. 

An example of application can be seen in Fig. 2 

4. BLURRING PSFs IN ORDER TO CLEAN BLURRED IMAGES 

Even if the great part of HST observations are taken in fine-lock mode, sometimes 
a loss of lock can happen. As a consequence finding a way to perform restoration and 
deconvolution even in such conditions is, in our opinion, of fundamental interest. 

In fact, even when the frames clearly look trailed it is not so easy to give a detailed 



180 




Figure 2: An example of image processing in IDL environment: a) 
The original SN 1987a image; b) the same with the two bright stars 
subtracted; c) stretched in order to circularize the bright ring and 
d) converted to polar coordinates. 



181 



description of the blurring. Such information can be retrieved, with high signal to noise 
ratio and using array manipulation based on FFT, comparing the Auto Correlation 
Function (ACF) of an unblurred (i.e. collected in Fine Lock mode) PSF with the 
ACF of the whole frame (supposed trailed). This is rigorously true if there are neither 
extended sources nor crowding effect in the image; anyway, due to the weighting nature 
of ACF, a small number of close stars in a rich frame does not essentially modify the 
whole ACF. 

Comparison could be possibly obtained following two methods. The first requires 
deconvolution of the ACF of the raw image using the ACF of the untrailed PSF as 
synthesized beam. It is very simple, for instance, to deconvolve via standard Clean: 
this allows one to immediately see the presence of trailing, since Clean associates to 
the main peak of the image ACF a set of displaced position instead of only one single 
well defined location (as it would have been if the image had not been trailed). On 
the other hand it is also possible to analyze the shape of the isolevels in the main peak 
of the trailed image ACF, comparing them with the shape of the isolevels of the PSF 
ACF. 

This latter method is very simple and we found its results (amount and direction 
of trailing) are in good agreement with respect to those given by the former. 

An important consequence follows: any strongly PSF-based deconvolution tech- 
nique can get advantage from the preparation of a PSF blurred in the same way as 
the whole raw frame. So, we have found it possible to produce a synthetically trailed 
PSF (knowing both the amount and the direction of image trailing) and to use that, 
instead of the normal one, to deconvolve the raw image, for instance by Clean or Lucy 
algorithm. 

Results are shown in Figure 3. 

More difficult is the treatment of images taken with HST in oscillating conditions, 
for the trailing due to the spacecraft oscillation translates into a space-invariant blurring 
component plus a rotation of the field. Such latter effect is space- variant and so the 
described technique is not able to take it into account. 



5. COMBINED ACTION OF LUCY AND CLEAN 

The Richardson-Lucy iterative technique is able, at least in principle, to perform 
the deconvolution of HST-FOC frames. Our experiments, that are similar to those 
obtained by others (Adorf, 1990), show that the effects of a restoration of this kind are 
essentially significative on the outer halo of the PSF. 

Only a very large number of iterations, with particular additional constraints, seems 
to produce point-hke sources if performed on an image of point-Uke sources convolved 
with a typical FOC PSF. 

In such a framework, we have made some simple attempts to merge the Lucy 
algorithm with Clean. It is worth noting that this approach is quite similar to Meier's 
(1990), who proposes a merging between MEM (Maximum Entropy Method) and Clean. 

Lucy's capabihty of enhancing SNR is very good, but the algorithm requires many 
iterations to reach a high degree of resolution. So we have thought to apply Lucy (for 
instance 20 iterations) to both the raw image and the PSF and then use the lucy-ed 
PSF to deconvolve via Clean the lucy-ed image. 

Due to the space- variant effect of the Lucy algorithm on the raw image, the relia- 
bility of the method (as far as the photometric precision is concerned) is quite low. As 
a consequence the method is recommended only for obtaining positional information 



182 




Figure 3: a): The trailed raw frame (M14), b) deconvolution using 
a blurred PSF, c) and d) the ACF of the raw frame and of an 
unblurred PSF; in the inset a magnification of the peak of the ACF 
is shown. 



183 



for single low SNR unresolved sources. 

An example of application is shown in Figure 4. 



6. A PHOTOMETRIC APPROACH 

As a main consequence of the spherical aberration, the large halo of the HST-PSF 
requires a special care when any sort of aperture photometry is attempted. 

A minimal mathematical description of the aperture photometry in presence of 
large PSF halo suffices to give insight into a possible way to overcome the problem and 
suggests useful photometric procedures. 

Aperture photometry on position x,y can be efficiently described introducing the 
window function W (see Figure 5). 

Saying R the raw frame, the sum of counts falling within the aperture defined by 
W centered in Xj , y ,■ can be expressed by 

Fj = Jw{x- Xj,y- yj) • R{x,y)dxdy (1) 

The true sky, assumed composed by N single unresolved sources of intensity !{, can 
be described by the function S: 

N 
S = Y^ h ■ ^{^ - '^hV ~ Vi) (2) 

while the collected distribution is given by the following convolution: 

N 
R = S® PSF =Y.h- PSF{x -xi,y- yi) (3) 

i=l 
Once locations Xj,yj of the found stars are known, the aperture photometry pro- 
cedure gives the following set of N measurements: 



Fj = J W{x-Xj,y-yj)-Rdxdy = ^ h J W{x -Xj,y -yj)- PSF{x -Xi,y -yi)dxdy 

So, defining the matrix Pij as: 

Pi,j = J ^{^ - ^j,y -Vj)- PSF{x -Xi,y~ yi)dxdy (5) 

while 

F = [Fi,F2,...,Fn]; I = [Ii,I2,---,In] (6) 

equation (4) becomes: 

F^ = P#I (7) 

where # indicates the row by column product between matrices. Note that no halo 
means P = 1. This suggests a way to perform aperture photometry, via inversion of 
the P matrix, since from equation (7): 

I=p-'^#F^ (8) 

184 




Figure 4: a): Raw frame (R136a), b): the same Lucy-ed, c): Clean- 
ing, and d) convolution with a gaussian beam. 



185 



y . 




a> 





V(x,y) 




^ 


■///. 




j^i 


^y/ 




'^ 


J/^ 


X 


'% 


w^ 





y . 


# 


V(x-x,y-y) 


b> 


y ■ 


M 


1 






^ 


^ 








X 


X 



Figure 5: a): Definition of the function W^(x,y), b): the function 
W{x - x,y -y) is able to define the aperture location in any point 

{x,y). 



186 



Extensive testing of this method is now in progress. The estimation of the reachable 
photometrical accuracy is the subject of a future work, where a comparison between this 
and other methods is performed, via some numerical simulations with various degrees 
of crowding and different shapes of the luminosity function. 

Initial tests show, as expected from a rough analytical estimation of the error (forth- 
coming), that satisfactory results can be obtained even with a high degree of crowding, 
provided the luminosity function is narrow enough. 

7. CONCLUSIONS 

Feasibility of a typical radioastronomical technique, like Clean, has been shown to 
be effective on the images of HST-FOC. However, some problems arise when the SNR 
is very low and, unless one is interested only in the brightest objects, the use of Clean 
becomes difficult, because it can produce false detections. This can be avoided following 
the suggestions given in section 5, although the non-linearity of Lucy's algorithm pro- 
duces some minor problems related to the positional dependance of the PSF. Anyway, 
this is really not a severe problem, as we are interested in deconvolution primarily as a 
mean for locating objects positions in the frame and not in the evaluation of fluxes in 
deconvolved images. Finally, a new way to perform photometry on HST-FOC images 
is here indicated. 



ACKNOWLEDGEMENT 

Thanks are due to prof. F.Bortoletto for his kind advices and help in data analysis. 
We are indebted to Dr. A.Nota and Dr. F.Rampazzi for useful suggestions and careful 
reading of the manuscript. 



REFERENCES 

Adorf, H.M. : 1990, ST-ECF Newsletter, 14, 8. 

Hogbom, J. A. : 1974, Astron. Astrophys. Suppi, 15, 417. 

Meier, D.L. : 1990, The restoration of HST images and spectra, STScI workshop, 20-21 

August 1990 eds. R.L.White and R.J.Allen, 113-120. 
Segalovitz, A., Frieden, B.R. : 1978, Astron. Astrophys., 70, 335. 
Steer, D.G., Dewdney, P.E., Ito, M.R. : 1984, Astron. Astrophys., 137, 159. 
Wakker, B.P., Schwartz, U.J. : 1988, Astron. Astrophys., 200, 312. 



187 



FOC Images of the Gravitational Lens System G2237+0305 



P. C^anel'^ R. Albrecht^-^-^, C. Barbieri^-'', J. C. Bladesl'^ A. Boksenberg^'^ 

J. M. Deharveng^'^, M. J. Disney^'^, P. Jakobsen^'^, T. M. Kamperman^'^° 

I. R. Kingl'^S F. Macchettol•3'^ C. D. Mackayl'12, F. Paresce^'^-^, G. Weigelt^-l^ 

D. Baxter^, P. Greenfield^, R. Jedrzejewski^, A. Nota^'"^, W. B. Sparks^ 



'Member FOC Investigation Definition Team 

'Space Telescope European Coordinating Facility 

'Astrophysics Division, Space Science Department of ESA 

^Observatorio Astronomico di Padova 

^Space Telescope Science Institute 

'Royal Greenwich Observatory 

'European Southern Observatory 

'Laboratoire d'Astronomie Spatiale du CNRS 

'Department of Physics, University College of Cardiff, Wales 

'"Laboratory for Space Research, Utrecht 

"Astronomy Department, University of California, Berkeley 

"Institute of Astronomy, Cambridge 

"Max Planck Institut fur Radioastronomie, Bonn 



1. Introduction 

The gravitational lens G22374-0305, discovered by Huchra et al. (1985), appears as a 
result of an extremely fortuitous alignment of a background QSO at z = 1.695 with the 
nucleus of a 14th magnitude foreground galaxy at z = 0.039. This lens produces four 
distinct QSO images (see Racine, 1991, for the best ground-based images) arranged in 
a roughly symmetrical cross, centered on the nucleus of the galaxy. Models of this lens 
presented by Schneider et al. (1988) and Kent and Falco (1988) imply the alignment is 
better than 0.1 arcseconds. 

Although ground-based images of this lens with excellent seeing (^ 0.48 arcsec 
FWHM) have resolved the four QSO images, clearly better resolution is required to (a) 
improve or confirm their positions and magnitudes, (b) better determine the galaxy's 
nuclear structure which has an important effect on the QSO images, and (c) search for 
the fifth image predicted near the nucleus by current lensing theory. 

2. Observations 

The observations reported here were obtained through the f/96 camera on 27 August 
1990 and 19 December 1990 and comprise three images. The first image was a 597s 



188 




Figure 1: (a) (left) The central 256x256 (~ 5.5"x5.5") region of the F502M image, (b) 
The residual image obtained after subtracting the lensed quasar images. [Note: The 
cores of the quasar images do not come away cleanly because of differing amounts of 
non-linearity compared to the PSF used.] 

acquisition exposure taken through the F430W filter (close to a Johnson B Filter). 
This image has 512x1024 pixels, where the pixels are rectangular and have a size of 
% 0.044x0.022", resulting in a field size of ^ 22"x22". The second image (see Figure 
la) is a 512x512, 1496s exposure taken through the F502M filter, (approximately a Gunn 
g filter). In this image the pixels are ^ 0.022" square, giving a field of ^ ll"xll". The 
third image was also an f/96 exposure through the F342W filter for 3 x 1200s. The 
results from the first two images were reported by Crane et a/., 1991. 

The brightest of the lensed images, B, has a peak count of 430 counts in the F502M 
image and the corresponding count rate is well within the FOC linear range for point 
sources (Paresce, 1990). The diffuse source seen between the QSO images is the nucleus 
of the lensing galaxy and has a peak count of ~45 counts per pixel. Figure lb shows 
the residual image of the galaxy with the quasars subtracted. 



3. Results 

These images allow us to determine the relative magnitudes of the individual images 
of the quasar, the galaxy, and to set an upper limit on the brightness of any fifth image. 
We also determine the positions of the individual images with very high precision. 
Except for the result on the fifth image, these results are reported in greater detail by 
Crane et a/.(1991). 



3.1. Positions 

The positions of the individual quasar images were measured using a simple cen- 
troiding algorithm in IRAF. The results are given in Table 1 below, and are compared 



189 



Object 


AX 


AY 


AE 


AN 


A 


0.000 


0.000 


0.000 


0.000 


B 


0.108 


1.796 


-0.672 


1.673 


B-Yee 






-0.68 


1.68 


B-Racine 






-0.671 


1.682 


C 


-0.976 


0.941 


0.626 


1.202 


C-Yee 






0.62 


1.20 


C-Racine 






0.617 


1.203 


D 


0.646 


0.761 


-0.854 


0.517 


D-Yee 






-0.85 


0.530 


D-Racine 






-0.853 


0.530 


Galaxy 


-0.209 


0.917 


-0.093 


0.936 


G-Yee 






-0.08 


0.94 


G-Racine 






-0.073 


0.938 



Table 1: Relative Positions (in arcsec). 

to the results of Yee(1988) and Racine(1991). The agreement with the results of Racine 
is quite good, except for the position of the galaxy nucleus which is off by about a FOC 
pixel. The center determined in the FOC image is closer to the D quasar image. 

3.2 Photometry 

The relative brightness of the individual quasar images was determined by summing 
the flux inside fixed apertures. The details of the procedures used are given in Crane et 
a/.(1991) The results are summarized in Table 2. We note that the reported brightening 
of image B reported by Pettersen(1990) is confirmed. 

Table 2: Relative Magnitudes. 



Object 


A</ 




AB 


A(/ 


AR 




A<? 


Ar 




(1) 




(1) 


(1) 


(2) 




(3) 


(4) 


A 


0.00 




0.00 


0.00 


0.00 




0.00 


0.00 


B 


-0.14 




-0.14 


-0.12 


0.53 




0.21 


-0.10 


C 


0.67 




0.70 


0.78 


1.14 




0.69 


0.33 


D 


0.89 




1.02 


0.80 


1.37 




0.92 


0.83 


Date 


27 Aug 90 


27 


Aug 90 


19 Dec 90 


18 Aug 


88 


25 Sep 87 


13 Oct 85 



Notes: (1) This paper. The relative magnitudes have an error bar of ±0.05 for component 
B and ±0.10 for components C and D. An estimate of the g and B magnitudes of the 
A component is 17.74 ±0.10 and 17.96 ±0.07 respectively, (2) Irwin et al. (1989), (3) Yee 
(1988), (4) Schneider et al. (1988). 



190 



3.3 Fifth Image 

The 3600s image in the UV was searched for a fifth image at or near the position of 
the galaxy nucleus. A careful subtraction of the quasar images resulted in no detectable 
image brighter than 250 times fainter than image B. This corresponds to 5.8 magnitudes 
fainter that image A. This result is uncorrected for the relative difference in extinction 
between the center of the galaxy and image A. Using the extinction law found of Nadeau 
et ai, 1991, the extinction might be as much as one magnitude greater at 3450A than 
at the wavelength where Racine(1991) claims to have found a fifth image. Thus our 
limit would be 4.8 magnitudes fainter than image A and is just consistent with Racine's 
claimed detection at 4.5 magnitudes fainter than image A. 



REFERENCES 

Crane, P., et ai, 1991, Ap. J. (Letters), 369, L59. 

Huchra, J., Gorenstein, M. Kent, S., Shapiro, I., Smith, G. Horine, E,. and Perley, R., 

1985, A. J., 90, 691. 
Irwin, M.J., Webster, R.L., Hewett, P.O., Corrigan, R.T., and Jedrzejewski, R.I., 1989, 

A. J., 98, 1989 
Kent, S.M., and Falco, E.E., 1988, Ap. J., 96, 1570. 
Paresce, F., 1990, The Faint Object Camera Handbook (Baltimore: Space Telescope 

Science Institute) 
Pettersen, B.R., 1990, I.A.U. Circ. No 5099. 
Nadeau, D., Yee, H.K.C., Forrest, W.J., Garnett, J.D., Ninkov, Z., Pipher, J.L., Ap. J., 

376, 430. 
Racine, R., 1991, A. J., , 102, 454. 
Schneider, D.P., Turner, E.L., Gunn, J.E., Hewitt, J.N., Schmidt, M., and Lawrence, 

C.R., 1988, A. J., 95, 1619. 
Yee, H.K.C., 1988,^. J., 95, 1331. 



191 



REDUCTION OF PG1115+080 IMAGES 



Edward J. Groth, Jerome A. Kristian, S. P. Ewald, J. JefF Hester, 
Jon A. Holtzman, Tod R. Lauer, Robert M. Light, Edward J. Shaya, 
and the rest of the WFPC Team: William A. Baum, Bel Campbell, 
Authur Code, Douglas G. Currie, G. Edward Danielson, S. M. Faber, 
John Hoessel, Deidre Hunter, T. Kelsall, Roger Lynds, Glen Mackie, 
David G. Monet, Earl J. O'Neil, Jr., Donald P. Schneider, 
P. Kenneth Seidelmann, Brad Smith, and James A. Westphal 



1. THE DATA 

The data are three exposures in PC6 through F785LP obtained on March 3, 1991. 
The exposure times are 120, 400, and 400 seconds. The data are reduced with the 
"standard" WFPC reduction scheme: A-to-D correction, DC bias subtraction, AC bias 
subtraction, dark current subtraction, preflash subtraction, and flat field normalization, 
using the best available calibration data. The exposures are combined into a weighted 
average normalized to 400 seconds exposure time, so one DN (data number) is about 
17.25 electrons. At this step, cosmic rays are removed by intercomparison of the three 
images. 

2. THE GOAL 

The lensing object can be seen in the processed image. One would like to subtract 
the four QSO images to leave behind a clear picture of the lens. 

3. THE PROBLEM 

Due to various glitches there is no high signal-to-noise PSF observation contempora- 
neous with the PG1115+080 observations. Since the data were obtained, the secondary 
mirror has been moved several times in an attempt to improve the performance of the 
FGSs, so it is unlikely that a PSF suitable for subtraction will ever be obtained. 



4. THINGS THAT DON'T WORK: A "THEORETICAL" PSF 

One of the things we tried was a PSF from the STScI library of PSFs. The library 

192 



PGl 115+080 Stretch = 1/20 Full Scale 



ii*iiT*«iiiiiiiiriii«imiiiiiiiininiiiMiiiiiiiiiiiiin 



O 

-^ 
OJ 



OJ 



o 
o H 

OJ 



I „ u I _slL. I 






Lens 



200 



220 
X 



240 




J— o 



contains PSFs calculated "from first principles." While the library PSFs are qualita- 
tively similar to the actual PSFs — one can make correspondences between the tendrils 
and rings, etc. — the library and actual PSFs differ in quantitative details which are 
important for the kind of subtractions required here. 

The subtractions of the library PSF yields an image in which the lens is obscured 
by incomplete removal of the outer parts of the PSF. We attempted to calculate more 
accurate PSFs but were not successful. 



5. THINGS THAT DON'T WORK: A LOW S/N PSF 

Observations of Q0957+561 were obtained the same day as those of PG1115+080. 
One of the QSO images in these exposures is sufficiently well separated from the lens and 
the other image that it can be used as a PSF. Unfortunately, one of the two exposures 
with F785LP was badly jittered, leaving only a single 350 second exposure to be used 
for the PSF. Although the core of the PSF is well exposed, the halo is not. Using this 
object for subtraction introduces so much noise in the resulting image that the lens is 
obliterated. 



6. SOMETHING THAT WORKS: AP LIB (BUT IT'S HARD) 

Observations of AP Lib were obtained the same day as those of PG1115+080. These 



193 



observations include three exposures through F785LP in PC6. The exposure times are 
30, 500, and 500 seconds. These data were processed through the standard reduction 
in the same way as the PG1115+080 observations. In this case, the combination of the 
images into a weighted average also takes account of the fact that the central four pixels 
of the 500 second exposures are saturated and uses only the data from the 30 second 
exposure for these pixels. 

An advantage of the AP Lib exposure is that it's very high signal-to-noise: a satu- 
rated core means that the halo is well exposed. Another advantage is that AP Lib is 
centered on PC6 only about 55 pixels from the center of the PG1115+080 images. 

A big disadvantage is that AP Lib is not a point source: there is a galaxy underneath 
that fills the entire detector! 

However, it appears that AP Lib can be well approximated as a point source plus a 
concentric, circularly symmetric galaxy. It should be possible to take advantage of this 
symmetry. 



7. ASSUMPTIONS AND PROCEDURES 

Assume that the AP Lib image is a circularly symmetric smooth galaxy concentric 
with a point source manifested as the PSF. Note that this assumption is probably not 
quite correct. The convolution of the PSF with a smooth function should give back a 
smooth function. But, at the center, the galaxy in AP Lib may have structure on scales 
comparable to the structure in the PSF. Thus, the validity of this assumption must be 
judged by how well the procedure works. 

In any case, with this assumption, the model is that everything in the PG1115+080 
image with the exception of the lens can be represented as: 



s{rj) = j:ai{A{vj-ri)-G{\vj-r,\)) 



where ^(rj) is the signal in pixel rj, Tj is the center of QSO image i,i = 1,2,3,4, Oj is 
the relative strength of QSO image i, A is the AP Lib image, and G is the circularly 
symmetric galaxy profile in the AP Lib image. 

This model is fit to the PGl 115+080 image using weighted least squares. The 
AP Lib image is translated to each QSO position with bi-cubic interpolation. The 
galaxy profile, G, is represented as 101 numbers giving the value of the profile at radii 
from to 100 pixels; linear interpolation is used to center G at each QSO image. 
Altogether there are 105 parameters estimated by the fit: four QSO amplitudes and 
101 numbers in the profile. Errors are determined by propagation of errors using the 
read and photon noise in the PGl 115+080 image. The fit is performed for three cases: 
In case 1, a patch of 12 pixel radius centered on the lens is excluded from the fit. In 
case 2, the lens is not excluded, case 3 is a fit to a simulation, whose description is 
omitted due to space considerations. The following table summarizes results from the 
fits: 

Case Pixels Degrees of x X 

in Fit Freedom Before After 

1. Lens Excluded 46180 46075 815738 48772 

2. Lens Included 46621 46516 861359 50097 

3. Simulation 46180 46075 832000 47856 

and the following shows the QSO parameters 



194 



QSO 


Oj (Case 1) 


aj (Case 3) 


Case 1 Rescaled 


Al 


0.1180 ±0.0017 


0.1159 ±0.0013 


0.1150 


A2 


0.0819 ±0.0012 


0.0795 ± 0.0009 


0.0798 


B 


0.0204 ± 0.0003 


0.0199 ± 0.0003 


0.0199 


C 


0.0315 ±0.0005 


0.0301 ± 0.0004 


0.0307 



8. RESULTS OF THE SUBTRACTION 

Once the QSO amplitudes, a{, and the galaxy profile, G, are determined, the QSOs 
can be subtracted, leaving a picture that contains only the lens (and possibly the fifth 
image!). The results shown are for case 1, the lens excluded fit. The results for case 
2, the lens included fit, are similar, except that the galaxy profile is a little higher at 
radii corresponding to the distance of the lens from the two brighter QSO images. The 
subtraction then leaves the lens slightly fainter and leaves a slight hole to the upper left 
of the two brighter QSO images at about the same distance as the lens. 

PGl 115+080 Lens, Hard Stretch 



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9. FUTURE WORK 

Future work will attempt to improve the subtraction, to deconvolve the lens, and 
then to improve the lens model based on these data. Additional observations will be 
proposed in order to obtain a higher signal-to-noise image of the lens. 



195 



OPTICAL AND UV POLARIZATION OBSERVATIONS OF THE M 87 JET 



P. E. Hodge, F. Macchetto, W. B. Sparks 
Space Telescope Science Institute 
3700 San Martin Dr 
Baltimore, MD 21218 
USA 



The f/96 relay of the Faint Object Camera (FOC) contains three linearly polarizing 
filters with nominal position angles of 0, 60, and 120 degrees. These filters are described 
in some detail by Paresce (1990). Observations of the knot A region of the jet of M 87 
were taken with the FOC in the blue and ultraviolet through the polarizing filters for the 
purpose of determining the orientations of the filters. Polarization maps were created 
from these data, and the results were compared with 2 cm VLA observations (Owen, 
Hardee, Cornwell, 1989). 

The observations were taken on 1991 April 3. Each of the six exposures for polar- 
ization was of 1500 seconds duration. Three exposures were taken through the F430W 
filter, one with each of the polarizing filters, and then three exposures were taken 
through the F220W filter, one with each polarizing filter. The approximate peak wave- 
length and bandwidth of the F430W filter are 396 nm and 87 nm respectively, while 
for F220W the values are 226 nm and 47 nm (Paresce 1990). 

The results are shown in Figures 1 and 2, both of which show polarization vectors, 
with the orientation indicating the magnetic field direction, and the length proportional 
to the polarized flux. The horizontal and vertical axes are labeled in image pixel 
coordinates. Figure 1 shows the data taken through the F430W filter, and Figure 2 
shows the data taken through the F220W filter. The image orientation and scale are 
shown to the right of the plot border. The scale of polarized flux is indicated by a line 
at the lower right within the plot border showing the length of a bar for a polarized 
flux of 100 counts detected by the FOC during the full exposure of 4500 seconds. 

The polarization map taken with the F430W filter is smaller than that with the 
F220W filter because the telescope was moved between the second and third of the 
F430W images in order to include knot C in the field of view. A common section 
within the overlap region was extracted. All three of the F220W images were taken at 
the same location. 

The basic calibration of these images was performed by the Post Observation Data 
Processing System at the Space Telescope Science Institute. For FOC images, this 
calibration includes dividing by a flat field and correcting for the geometric distortion 
of the optics and camera. For further details, see Greenfield et al. (1991). 

The background level was measured for each image by taking averages in regions 
away from the jet, and the values of about 15 counts per pixel for F430W and one 



196 




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198 



count per pixel for F220W were subtracted from the images. To improve the signal-to- 
noise, the F430W images were block averaged with a 10 x 10 pixel box, and the F220W 
images were block averaged with a 16 x 16 pixel box. The image scale of the FOC at 
f/96 is 0.02217 arcsecond per pixel (Greenfield ei al. , 1991). These images have not 
been deconvolved. 

The three polarizing filters differ somewhat in throughput. The 60-degree filter has a 
short-wavelength cutoff near 220 nm, while the and 120-degree filters extend below 150 
nm (Paresce 1990). When the polarizing filters are combined with F430W, the difference 
in throughput is less than one percent. With the F220W filter, on the other hand, the 
throughput of the 60-degree filter is only about 2/3 that of the other two polarizing 
filters. In order to have any confidence in the polarization measurement for F220W, 
this factor must be accurately determined. The reflectivities of the HST mirrors, the 
transmission curves of the various filters, and the sensitivities of the detectors were 
measured prior to launch. Home (1990) has written a program called XCAL to calculate 
the throughput of the HST with various instrument configurations and different spectral 
distributions of the incident light. We used XCAL to calculate the relative throughput 
of the three polarizing filters together with either the F430W or F220W filter, and then 
we used these values to normalize the images. When running XCAL, we specified that 
the input light was unpolarized and had a power-law spectral distribution. We estimated 
that the spectral index was -|-1 (F-lambda increases with decreasing wavelength), based 
on the F430W and F220W fluxes in knot A. 

The maximum flux was in the knot A region, with 334 counts through the F430W 
filter and 67 counts through the F220W filter. The maximum amplitude of polarization 
was 111 counts for F430W and 30 counts for F220W. Thus the degree of polarization in 
knot A was of order 30 to 40 percent. Note that the figures show polarized flux rather 
than percent polarization, and the direction shown is that of the magnetic field. 

The position angles of polarization in the strongly polarized regions in these images 
are predominantly either parallel or perpendicular to each other. In this situation 
there are two arrangements of the polarizing filters that give identical results. The 
two arrangements are mirror images of each other around the direction of polarization. 
This ambiguity would not be an issue during routine observations because the filter 
positions would already have been calibrated, but these data were taken for the purpose 
of verifying the orientations of the polarizing filters. One arrangement agrees with the 
engineering drawing for the filter wheel, and we assume it is correct, but the other 
arrangement is not ruled out by these observations. 

REFERENCES 

Greenfield, P., et al., In-Flight Performance of the Faint Object Camera of the Hubble 

Space Telescope, Proc. Soc. Photoopt. Instrum. Eng., in press. 
Home, K., 1990, XCAL Users Manual, Space Telescope Science Institute, Baltimore. 
Owen, F. N., Hardee, P. E., Cornwell, T. J., 1989, Ap. J., 340, 698. 
Paresce, F., 1990, Faint Object Camera Instrument Handbook, Space Telescope Science 
Institute, Baltimore. 



199 



THE NON-PROPRIETARY SNAPSHOT SURVEY: 

A Search for Gravitationally-Lensed Quasars 

Using the HST Planetary Camera 

D. Maoz\ J.N. Bahcall\ R. Doxsey^, D.P. Schneider\ 
N.A. Bahcall^ O. Lahav^ and B. Yanny^ 

1. Institute for Advanced Study 

2. Space Telescope Science Institute 

3. Princeton University Observatory 

4. Institute of Astronomy 

The Snapshot Survey is an imaging survey of bright quasars using HST's Plan- 
etary Camera (PC). Short exposures (2 or 4 minutes) are taken during gaps in the 
scheduled observing program, when the telescope would otherwise be idle. All images 
are obtained using only the gyroscopes for pointing and guiding, thus saving the time 
necessary to acquire guide stars, and also allowing us to monitor routinely the gyro 
performance. Targets are distributed throughout the sky, so only short slews (a few 
degrees) are required to move the telescope from any approved science target to a 
nearby Snapshot target. Snapshot targets are assigned only after all other programs 
have been scheduled. The resulting data are non-proprietary, and can be obtained 
from the STScI User Support Branch. Further details can be found in Bahcall et al. 
(1991). 

The scientific purpose of the currently operating Snapshot Survey is to search for 
evidence of gravitational lensing among known distant, intrinsically luminous quasars. 
Despite the spherical aberration of HST's primary mirror, the sharp core of the point- 
spread function, containing ~ 15% of the light, permits high spatial resolution studies 
of closely separated bright point sources. The existing point-spread function permits 
the detection of multiple images at subarcsecond separations, which cannot be easily 
probed from the ground. 

As of mid- April 1991, 89 short exposures, through two filters, of high luminosity 
quasars from a well-defined sample have been obtained. Useful high-resolution im- 
ages of 30 quasars have resulted. None show evidence of multiple images caused by 



200 



gravitational lensing. Simulations show that multiple images with brightness ratios 
of up to several magnitudes would have been detected, if present, down to image 
separations of ^ 0.1". This seems to be in conflict with several ground based surveys 
(e.^., Crampton et al. 1989, Surdej et al. 1989) who reported that a large fraction 
of quasars have subarcsecond multiple images. The paucity of lensed quasars found 
so far (0 out of 30) suggests lensing is a rare phenomenon, and argues against re- 
cently popular-again cosmologies involving a universe dominated by a cosmological 
constant. In such cosmologies (Fukugita and Turner 1991) multiple images would be 
detected in about 10% of our sample. As more improved-quality data are obtained, 
this study will allow a stronger confrontation with the ground-based surveys and the 
theoretical models. 

The Snapshot Survey has uncovered several engineering problems in the obser- 
vatory's performance, which have already been corrected. In particular, we have 
encountered large telescope pointing errors (typically 20") and drift rates (typically 
4.5 milli-arcseconds/sec, or 3 to 4 times that expected) when solely under gyro con- 
trol. We have determined that stellar aberration corrections are not applied in the 
current control system when HST is operating solely on gyros. The stellar aberration 
due to the motion of the earth around the sun is 

^= -sin <^ = 20.5" sin <?i , (1) 

c 

where v is the earth's velocity around the sun, c is the speed of light, and (f) is the angle 
between the earth's velocity vector and the direction of the object being observed. 
The orbital motion of HST about the earth contributes an additional aberration term 
with an amplitude of 5". The spacecraft's centripetal acceleration around the earth 
will cause the object's position to drift at a rate 

dd Idv . ^ . « -1 

-— = — ;-sm^ ?s 5.5smy mas s , (2) 

dt cdt ' ^ ' 

where 6 is the angle between the spacecraft's acceleration vector and the direction of 
the object. 

Comparing these relations to the observations above, we see that the lack of stel- 
lar aberration correction can account for much of the pointing error and of the drift 
rate. Since April 1991, the HST operation procedures have been modified to invoke 



201 



the on-board stellar aberration corrections for observations carried out under gyro 
control. The correction of this problem will probably shorten the time required to 
acquire guide stars in FGS-guided observations. The low intrinsic gyro drift rate (now 
that the stellar aberration correction has been implemented) may make gyro-guided 
observations attractive to other observers as well. We will continue to monitor the 
performance of the gyros throughout HST's Cycle 1 of observations. 

REFERENCES 

Bahcall, J.N., Maoz, D., Doxsey, R., Schneider, D.P., Bahcall, N., Lahav, O., and 

Yanny, B. 1991, ApJ, in press 

Crampton, D., McClure, R.D., Fletcher, J.M., and Hutchings, J.B. 1989, AJ, 98, 1188 

Fukugita, M. and Turner, E.L. 1991, MNRAS, in press. 

Surdej, J. 1989, in "Gravitational Lensing" eds. Y. Mellier, B. Fort, and G. Soucail, 

(Berlin: Springer- Verlag), p88 



FIGURE CAPTION 

Segments of typical Planetary Camera exposures of five Snapshot Survey quasars, 
with a variety of brightnesses and trail lengths. The image scale is 0.043" pixel" , 
and the field for each panel is 8.6" on a side. The orientation of the images is random, 
according to the HST roll angle at the time of the exposure. The gray scale is set 
individually for each image, such that the darkest hue corresponds to the number of 
counts pixel"^ in the brightest part of the quasar. The numerous dark specks in the 
images are charged-particle events. The filter used is indicated following the object 
name ("V" for F555W, and "I" for F785LP). The lower right-hand panel shows a 
simulated image of a quasar and a secondary image 2 magnitudes fainter, separated 
by 0.3". The primary image corresponds to a 17th magnitude object trailed at a rate 
of 4.6 mas s~^ in the 120 s exposures, or a 17.8 magnitude object trailed at a rate of 
2.4 mas s~^ in the 230 s exposures. 



202 



"^^■^■T ' ".'. "VJ 



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SIMULATION 



203 



FAINT OBJECT SPECTROGRAPH OBSERVATIONS OF CSO 251 



R.D. Cohen, E.A. Beaver, E.M. Burbidge,V.T. Junkkarinen, 

R.W. Lyons, and E. I. Rosenblatt 

Center for Astrophysics and Space Sciences, UCSD 

La JoUa, CA 92093-0111 



1. OBSERVATIONS AND REDUCTIONS 

The QSO CSO 251 (z=0.0788, F;^(5500A)= 2 x 10~^^) was observed as an early 
release observation (3065) with the FOS. Observations with the G130H grating, with 
lA per diode, and the "blue" Digicon detector were made on January 7, 1991. Sampling 
is improved by quarter-diode sub-pixel steps. In this mode the spectrum spans the fuU 
array sampUng from 1150A to I6OOA. 



2 — 



Ly a 



CSO 251 



Si IV 




<*«Ht^lt»*t^J^ 



1200 



1600 



1400 
WAVELENGTH (A) 

Fig. 1 - Observed spectrum of CSO 251 obtained with the FOS on the HST. Identified 
emission Hues at z = 0.0788 are marked above the spectrum and absorption lines in the 
Galaxy are marked below the spectrum. The trace at the bottom of the figure shows 
the 1(7 error. The emission Une at 12 15 A is geo-coronal Lya. 



204 



Four spectral exposures were made with the I'.'O circular aperture, selected as the 
best compromise between resolution and throughput efficiency. The exposure time per 
pixel is 1680 seconds. This UV spectrum is shown in figure 1. A fifth exposure was 
100 seconds using the 4'.'3 square aperture. The flux calibration of the large and small 
aperture data is consistent to better than 5%. Observations of Ha were taken at Lick 
Observatory within two weeks of the FOS observations, and observations of H/3 were 
made less than two months prior. While previous tests had shown a cyclic drift in 
position for the FOS red detector due to an interaction with the earth's magnetic field 
(Junkkarinen et al., 1991) no conclusive evidence has been found that such an effect 
exists in the blue detector. The background was scaled to match that predicted for the 
geomagnetic coordinates of our observations (Rosenblatt et al., 1991). 



2. EMISSION LINES 

The broad emission lines present in QSO spectra are probably produced by mate- 
rial, either in clouds or in an accretion disk, moving at velocities around 10000 km s~ 
at distances of order one to ten parsecs from the source of continuum radiation. The 
emission-line profiles produced by this gas depend on both the dynamics of the clouds 
and the emission properties of the individual clouds. Several models for the broad- 
emission-line region have been developed that predict shapes for line profiles which are 
in reasonable agreement with observations. These models include: radiatively accel- 
erated clouds (Blumenthal and Mathews, 1979), clouds accelerated by quasar winds 
(Weymann et al., 1982), and cloud motion along parabolic orbits (Kwan and Carroll, 
1982). Because the overall line parameters do not provide a definitive test of the pos- 
sible dynamical models, more subtle systematic properties of the emission-line profiles 
must be considered. Systematic QSO emission-line profile differences (Mathews and 
Wampler, 1985) and redshift differences (Gaskell, 1982 and Wilkes, 1984, 1986) have 
been observed. The interpretation of these data in terms of the dynamics, internal ob- 
scuration, and geometry of the broad-emission-line region is limited by the quality of 
the observations and the statistical nature of the problem. These data already provide 
direct evidence for an inhomogeneous, multi-component broad-emission-line region; 
with further observations, our understanding wiU improve. 

Wilkes found that, on average, the high-ionization lines were blue-shifted with 
respect to H/3, while the low-ionization lines were red-shifted. Observations of Ha in 
the IR by Espey et al. (1989) show a median redshift difference of 1000 km s~ between 
Ha and C IV A1549. Lya is usually near the C IV redshift. This difference might be 
the result of a broad-emission-line region that contains a low-ionization region that 
is optically thick and a separate optically thin high-ionization region (see for example 
Mathews 1986). The Lya emission-line profile compared to Ha can be used to estimate 
the Lya/Ha ratio in different regions. One prediction of this model is that at velocities 
where the optically thin component dominates, the Lya/Ha ratio should be near the 
recombination value i.e. large. Alternatively, some of the emission can come from the 
margins of the accretion disk (CoUin-Souffrin et al., 1988). 

Normalized emission line pairs can be divided to compare line flux ratios as a 
function of velocity (Shuder, 1982, 1984). With simple assumptions, Shuder showed 
that the Ha and H/? profiles indicate that velocity increases inwards in the broad-line 
region of Seyfert galaxies, although the effects were less pronounced in QSOs. Detailed 
comparisons of several lines, including those from both the fully and partially ionized 
zones, combined with predictions from photo-ionization models may allow us to reach 



205 



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0.4 

0.2 

0.0 

-0.2 

-10000 -5000 5000 

Velocity (km/sec) 



10000 



Fig. 2 - Broad emission lines in CSO 251 normalized and plotted on a velocity scale. 
The obvious smooth area in the wing of Lya indicates where N V has been removed. 
Other lines have been removed as described in the text. 



lor 



D 

X 

3 




-10000 -5000 5000 

Velocity (km/sec) 



10000 



Fig. 3 - The profile of Lya divided by the profile of Ha. The vertical scale shows true 
relative fluxes. 



206 



more detailed conclusions. Such comparisons require lines observed in the UV and 
optical or optical and IR. To date, the optical emission lines in such comparisons have 
been of higher resolution and signal-to-noise than the other lines. However, with the 
FOS and HST, emission lines from O VI out to Ha can be observed in low-redshift 
QSOs with comparable resolution and signal-to-noise. An additional advantage is that 
Lya in low-redshift QSOs is not severely cut up by Lya forest absorption lines. 

Our analysis for these lines is not complete. However, although the peaks of all 
lines occur at almost the same velocity, profile differences between the lines are evident. 
Normalized profiles of Lya, H/3, and Hq are shown in figure 2, and the division of the 
Lya profile by Ha is shown in figure 3. Broad emission lines at the positions of N V 
and Fe II were removed from the profiles of Lya and H/3 respectively, while a smooth 
continuum was fit through all UV absorption lines. Narrow lines of [N II] are not seen, 
nor were any narrow components of the permitted lines. Profile comparisons made after 
subtracting our best estimates of the narrow lines are similar to those shown here. 

While the division of U./3 by Ha (not shown) is reminiscent of the profile divisions 
shown by Shuder, the Lya division is surprising. This line ratio first falls with increasing 
velocity, and then rises again. This behavior is similar to that shown in Carroll and 
Kwan (1985), but we hesitate to draw broad conclusions based on observations of three 
lines in a single object. 



REFERENCES 

Blumenthal, G. R., and Mathews, W. G. 1979, Ap. J., 233, 479. 

Collin-Souffrin, S., Hameury, J.-M., and Joly, M. 1988, Astron. and Astroph., 205, 19. 

Espey, B. R., CarsweU, R. F., Bailey, J. A., Smith, M. G., and Ward, M. J. 1989, Ap. J., 

342, 666. 
Gaskell, C. M. 1982, Ap. J., 263, 79. 
Junkkarinen, V. T., et al. 1990, B.A.A.S., 22, 1282. 
Kwan, J., and Carroll, T. J. 1982, Ap. J., 261, 25. 
Mathews, W. G. 1986, Ap. J., 305, 187. 

Mathews, W. G., and Wampler, E. J. 1985, P.A.S.P., 97, 966. 
Rosenblatt, E. R., et al. 1990, B.A.A.S., 22, 1283. 
Shuder, J. M. 1982, Ap. J., 259, 48. 
Shuder, J. M. 1984, Ap. J., 280, 491. 
Weymann, R. J., Scott, J. S., Schiano, A. V. R., and Christiansen, W. A. 1982, Ap. J., 

262, 497. 
Wilkes, B. J. 1984, M.N.R.A.S., 207, 73. 
Wilkes, B. J. 1986, M.N.R.A.S., 218, 331. 



207 



FOC OBSERVATIONS OF R136a IN THE 30 DORADUS NEBULA^ 

G. Weigelt2'3, R. Albrecht2'4.5, c. Barbieri^'^, J. C. Blades^'^ A. Boksenberg^-S, P. Crane^-^, J. 
M. Deharveng2'i°, M. J. Disney^-ii, P. Jakobsen^'^, T. M. Kamperman^-i^ I. R. King^-^^, p. 
Macchetto^'S''^, C. D. Mackay^'^'*, F. Paresce^-^-^ D. Baxter^, P. Greenfield^, R. Jedrzejewski^, 
A. Nota^-'^, W. B. Sparks^ 

^ Based on observations with the NASA/ESA Hubble Space Telescope, obtained at the Space 
Telescope Science Institute, which is operated by AURA, Inc., under NASA contract NAS 5- 
26555, ^Member FOC Investigation Definition Team, ^MPI fiir Radioastronomie, "^Space Tele- 
scope European Coordinating Facility, ^Astrophysics Division, Space Science Department of 
ESA, ^Osservatorio Astronomico di Padova, ^Space Telescope Science Institute, ^Royal Green- 
wich Observatory, ^ESO, "^Laboratoire d'Astronomie Spatiale du CNRS, ^^ Department of 
Physics, University CoUege of Cardiff, ^^SRON - Space Research Utrecht, ^■'Astronomy De- 
partment, University of California, Berkeley, ^''institute of Astronomy, Cambridge. 

1. OBSERVATIONS AND DISCUSSION 

R136a is the central object in the 30 Doradus nebula (see Walborn 1973; 1986; 1990 and 
references therein). The physical nature of R136a has been the subject of controversy over the 
last few years. One suggestion was that R136a might be a single object with a mass of the 
order of 1000 solar masses. The other suggestion was that R136a is a compact star cluster and 
that it consists of several and WR stars. Observations of R136a by speckle techniques have 
resolved 8 stars within 0.7 arcsec diameter (Weigelt and Baier 1985; Neri and Grewing 1988). 
The HST observations described here (Weigelt et al. 1991) confirm that R136a is a compact 
star cluster. There is good agreement with the speckle observations. 

The raw image shown in Fig. 1 was taken on 1990 Aug. 23 (filter F346M and neutral 
density (ND) filter F8ND; FOC f/96 mode; coarse track mode; exposure time 600 s). The 
count number in the brightest pixel is only 136 since too many ND filters were used. The 
background light is caused by the wings of the psf (spherical aberration). Fig. 2 shows the 
same R136 image after application of the image restoration method CLEAN. 

Fig. 3 shows a high-resolution image of R136a reconstructed from a FOC f/288 exposure 
(1990 Aug. 23; f/288 mode; 1.7 arcsec region; filter F253M plus ND filter F4ND; exposure time 
900 s; count number in the brightest pixel is only 27; CLEAN reconstruction). The photometric 
accuracy of the reconstructed image is not very good since the raw image is very noisy. The 
separation of the bright, close double star al-a2 is ~ 0.11 arcsec. In the f/288 raw image 
the star R136a2 is ~ 0.4 (±0.2) magnitudes fainter than R136al, while the stars R136a3 and 
R136a6 are ~ 0.6 (±0.3) magnitudes fainter than R136al. AU 8 stars resolved by holographic 
speckle interferometry can be found in both the f/96 and the f/288 FOC images. 

Walborn (1986) has calculated the mass of the brightest component al on the assumption 
that the V-magnitudes of a2 and a3 are not more than ~ 0.3 mag fainter than al. In this case 
he finds an upper limit for the mass of al of ~ 250 solar masses. The speckle observations 
have shown that in the red the magnitude differences of al, a2, and a3 are ~ to 0.3. From 
the HST observations we now know in addition that at 2550 A the magnitude differences of 



208 




Figure 1. FOC f/96 raw image of R136 (filter F346M + F8ND). 



209 




57 

|l 

2 _ I* ^ - ®4 



1 



■^% 



3 — • 
6 - • 

Figure 2. Image of R136 reconstructed from the FOC f/96 image shown in Fig. 1. 



210 



al, a2, a3, a6 are only ~ 0.4 to 0.6. This means that the HST observations support Walborn's 
conclusion that the upper limit for the mass of R136al is ~ 250 solar masses. 

ACKNOWLEDGEMENTS. The FOC is the result of many years of hard work and important 
contributions by a number of highly dedicated individuals. In particular, we wish to thank ESA 
/r5T'Project Manager R. Laurance, the ESA/ F5T Project Team, and the European contractors 
for building an outstanding scientific instrument. The FOC IDT Support Team, D. B., P. G., 
R. J., and W. B. S., acknowledge support from ESA through contract 6500/85/NL/SK. P. C. 
and I. R. K. acknowledge support from NASA through contracts NAS5-27760 and NAS5-28086. 

REFERENCES 

Neri, R. & Grewing, M. 1988, A&A, 196, 338 

Walborn, N.R. 1973, Ap. J. (Letters), 182, L21 

Walborn, N.R. 1986, in lAU Symp. 116, p. 185, Walborn, N.R. 1990, in lAU Symp. 148, p. 145 

Weigelt, G. & Baier, G. 1985, A&A, 150, L18 

Weigelt, G., Albrecht, R., Barbieri, C, Blades, J.C., Boksenberg, A., Crane, P., Deharveng, 

J.M., Disney, M.J., Jakobsen, P., Kampemann, T.M., King, I.R., Macchetto, F., Mackay, CD., 

Paresce, F., Baxter, D., Greenfield, P., Jedrzejewski, R., Nota, A., & Sparks, W.B. 1991, Ap. 

J. (Letters), 378, L21 



2 - 

1 - 




3 - ^^ 





6- Jm 



0.5" ^. ^ 



Figure 3. Image of R136a reconstructed from a FOC f/288 exposure (filter F253M+F4ND). 



211 



GHRS CHROMOSPHERIC EMISSION LINE SPECTRA OF 
THE RED GIANT a TAU 

Kenneth G. Carpenter (NASA - Goddard Space Flight Center) 

Richard D. Robinson (Astronomy Programs - Computer Sciences Corporation) 

Dennis C. Ebbets (Ball Aerospace System Group) 

Alexander Brown and Jeffrey L. Linsky (JILA - Univ. of Colorado & NIST) 



1. INTRODUCTION 

The K5 III non-coronal giant a Tau was observed during the GHRS Science Assess- 
ment Observation (SAO) Program to assess capabihties of the spectrograph important 
to the study of narrow emission hne sources. A region near 2325 A was chosen since it 
contains intercombination Hnes of C II and Si II, which have very small intrinsic widths 
(no opacity broadening), as well as stronger Hnes of Fe II, Ni II, and Co II. The observa- 
tions were made through both the Large and Small Science Apertures (LSA and SSA) 
in both medium (G270M) and high (Echelle-B) resolution modes to allow a determina- 
tion of the relative instrument performance in the four observing configurations. The 
initial scientific results of the program are presented in Carpenter et al. (1991). In this 
paper, we discuss the instrument performance in more detail and present additional 
scientific results . 



2. INSTRUMENT PERFORMANCE 

2.1 Wavelength Calibration 

Figure 1 shows typical results from the default wavelength calibration procedure 
compared with those based on calibration lamp exposures taken at the same carousel 
position and close in time to the science observations. The top figure (a) shows the 
Pt wavelength cahbration spectrum for the G270M SSA exposure, where a default 
caUbration has been apphed. The dashed Unes indicate the expected locations of the 
calibration hnes. An offset of 2.83 diodes (36 km/s) was identified between the two. A 
careful reduction of the data using a near-simultaneous internal lamp exposures allows 
the computation of more accurate dispersion constants and wavelength offsets. The 
residuals from a second order polynomial fit to the measured hne positions are shown 
in the lower panel (b) and are seen to be on the order of 0.1 diode- widths (1.2 km/s). 
Table I summarizes the precision of various levels of wavelength cahbration, obtained 
from these and other SV/SAO observations. 

2.2. Effect of Spherical Abberation on Sensitivity 

Figure 2 shows the relative throughput of the Small and Large Science Apertures 
(SSA/LSA) versus wavelength. This curve is based on all SAO/SV/GTO observation 
sets where data were taken at the same wavelength in the same grating mode through 
both apertures and where we beheve the target was well-centered in the SSA. The 
point at 2325 A is based on the a Tau observations, while the targets used at the other 



212 



data points are indicated on the plot. Table II shows the counts rates seen through 
the four grating/aperture combinations used in the Alpha Tau program and compares 
them to each other and to pre-launch (pre-spherical aberration) expectations. 

2.3. Comparative Line Profiles 

Figure 3 illustrates the differences in the observed line profiles in four different ob- 
serving modes. The line profiles obtained in modes G270M/LSA, G270M/SSA, and 
Ech-B/LSA are compared to the 'true' profiles obtained with Ech-B/SSA (dashed- 
line). The wavelength shift of the G270M/LSA data can be attributed to a relatively 
poor calibration, since this observation is the only one of the four without a cal-lamp 
exposure at a nearby wavelength. The lines shown here are almost fully resolved in the 
G270M/SSA mode, although a few very minor differences can be still be seen between 
it and the Echelle SSA observation. 



3. SCIENTIFIC RESULTS 

3.1, Detection of Non-photospheric UV Continuum 

The ability to confidently detect the presence of a weak continuum in cool stars 
is one of the major advantages of GHRS over lUE. Figure 4 shows the G270M LSA 
observation (which has the highest photometric precision of the four observations), 
along with the expected photospheric flux (from a standard Kurucz line-blanketed LTE 
model with T^ff of 4000 K) from a Tau. Log(flux) is plotted to clearly display the 
weak and strong flux regions of the spectrum on a single plot. The factor used to scale 
the computed fluxes to flux-at-earth was derived by forcing agreement between the 
lUE and model fluxes at 3200 A. The observed stellar spectrum, which is weU-above 
background even between the strong emission features, is substantially above the 4000 
K photospheric flux. This excess flux is most likely the first detection of chromospheric 
continuum emission from a cool giant star. 



3.2. Line Profile Analysis 

The top panel in Figure 5 shows gaussian fits to a Co II line and a self-reversed Fe II 
line. The former is well-fit with a single gaussian, the latter by a combination of an 
emission plus absorption gaussians, where the absorption gaussian is shifted by about 
1.5 km/sec to the red of the emission gaussian. The lower panel shows gaussian fits to 
two of the C II] (UV 0.01) lines found in the Echelle data. The lines cannot be fit by 
single gaussians, but are well-represented by a two-gaussian fit, where the gaussians have 
the same central wavelength, but substantially different FWHM and maxima. However, 
the nearby Si II] (UV 0.01) lines (not shown) are well- represented by single gaussians. 
The profiles exhibited by the C II] lines are similar to those which can be generated, as 
discussed by Gray (1988), by the full-disk integration of an anisotropic velocity field. 
Alternatively, Harper (1991) has shown that such profiles can be generated using models 
of hybrid bright giants in which v^^^;, increases with Tg in the Une formation region. 
Similar physics may be responsible for the profiles seen in a Tau. The differences in 
the Si II] and C II] profiles suggests differences in the extent and/or location of the line 
formation regions for the two species. 

213 



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2355 2360 




2330 



2340 2350 
Wavelength (A) 

Figure 1 



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Figure 2 



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2325.5 2326 
Wavelength (A) 

Figure 3 



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T«ff-^003 K 



2340 2350 

Wovelength (A) 

Figure 4 




2330.5 2331.0 2331.5 2332.0 2332 5 




Figure 5 



214 



Table I 
Properties of the Wavelength Calibrations 



GrafiiiR 


Dispersion 


Wavclongt li 




Accuracy of Wavelenptli Scale* 




Default 


Calil,. 


SPYHAL Calil). 


Full Calih. 




(A/diode) 


Range (A) 


max offset 


max error 


max offset 


max error 


TTiax offset 


max error 


CHOI, 






(diodes) 


(km/s) 


(diodes) 


(km/s) 


(diodes) 


(km/s) 

15-10 


0.572-0.573 


1050-1800 


3 


470-285 


0.5 


78-48 


0.1 


GlIOM 


0.0.56-0.052 


1150-1700 


3 


46-28 


0.5 


7.6-4.5 


0.1 


1.5-0.9 


GIf.OM 


0.072-0.066 


1200-2000 


3 


54-30 


0.5 


9.0-5.0 


O.l 


1.8-1.0 


G200M 


0.081-0.075 


1600-2400 


3 


46-28 


0.5 


7.G-4.7 


0.1 


1.5 0.9 


G270M 


0.096-0.087 


2200-3200 


3 


39 24 


0.5 


6.5-4.0 


0.1 


1.3-0.8 


Ech A 


0.011-0.017 


10.50-1730 


3 


10. 


0.5 


1.6 


0.1 


0.33 


Ech B 


0.019-0.035 


1C80-3200 


3 


10. 


0.5 


1.6 


0.1 


0.33 



Maximum expected error in the absolute wavelenglh scale caused by thermal and magnetic drifts in the detector. 
Improvements can be expected if a known fiducial wavelength exists in the spectrum 



Table II 
Effects of the Spherical Aberration 

Throughput for the Large and Small Aperture 



grating aperture peak count total counts SSA/LSA degradation! 
(2325.8 A) per A throughput 



G270M 


LSA 


21.5 


76.0 




2.0 


G270M 


SSA 


11.3 


22.6 


0.30 


4.5 


Ech B 


LSA 


14.5 


117.3 


- 


2.0 


Ech B 


SSA 


5.0 


36.6 


0.31 


4.5 



I Degradation relative to the prc-launcli expectations 



REFERENCES 

Carpenter, K. G., Robinson, R. D., Wahlgren, G. M., Ake, T. B., Ebbets, D. C., and 

Walter, F. M. 1991, Ap. J., 377, L45. 
Gray, David F. 1988, chapter 1 in 'Lectures on Spectral Line Analysis: F, G, and K 

Stars' (The Publisher: Arva, Ontario). 
Harper, G. 1991, MNRAS, submitted. 

We acknowledge support of NASA to NIST through grant S-56460-D. J. Linsky is a 
Staff Member, Quantum Physics Div., NIST. 



215 



lUE FAR-ULTRAVIOLET SPECTRA OF CAPELLA AND 7 DRACONIS 
FOR COMPARISON TO HST/GHRS GTO OBSERVATIONS 



Thomas R. Ay res ^ 

Center for Astrophysics and Space Astronomy 

University of Colorado 

Campus Box 389 

Boulder, CO 80309-0389 

USA 



Abstract. I present reference spectra from the lUE Archives to compare with recent HST/GHRS 
observations of Capella and 7 Draconis. The comparison demonstrates graphically the enormous 
increase in sensitivity and spectral resolution afforded by the GHRS. At the same time, the HST 
tracings reveal that much of the faint structure in coadded lUE spectra is genuine: structure 
that seasoned lUE observers would tend to dismiss as noise. 

1. INTRODUCTION 

A previous paper (Linsky. Brown, eV" Carpenter 1991: this volume) reported the results 
of low-, moderate-, and high-dispersion spectroscopy of two bright late- type stars using the 
GHRS of HUBBLE during Science Verification and early GTO activities. The two targets - 
Capella (a Aurigae A [G9 III + GO III]) and 7 Draconis (K5 III) - present very different energy 
distributions in the vacuum ultraviolet. Capella - the archetype "active-chromosphere" giant 

- is dominated by bright, iiigh-e.xcitation emissions (like C IV AA1548,.'30). In contrast, 7 Dra 

- a typical "non-coronal" giant ~ is dominated by low-excitation species (like the I 1305 A 
multiplet), and its high-excitation spectrum is quite weak (in fact, thought to be entirely ab- 
sent prior to HST). The observational work described by Linsky and collaborators consisted of 
G140L low-resolution spectra of both stars, covering the range ll.'jO-lTSO A; medium-resolution 
(G140M, G160M, k G200M) spectra of selected intervals of Capella containing diagnostically- 
important emission lines; and ECH-A & ECH-B spectra of the bright chromospheric emissions 
of H I (A1215 Lyo) and Mg II (AA2795,2802 doublet) of Capella. The low-dispersion GHRS 
spectra have a factor of ss 5 higher spectral resolution than the comparable SWP-LO mode 
of the lUE; the medium gratings of the GHRS are comparable in spectral resolution (through 
the LSA) to the lUE echelle ("HI") mode; and the GHRS echelles have a factor of « 8 higher 
resolving power (through the SSA) than the lUE SWP-HI or LWP-HI modes. Furthermore, the 
HST/GHRS is considerably more sensitive than the lUE by virtue of its 25x larger collecting 
area, high-throughput spectrometers, and low-noise detectors. Nevertheless, the lUE has been 
accumulating spectrograms of a diverse set of cosmic targets for nearly a decade and a half. 
Thus, a comparison between HST/GHRS and lUE spectra is useful not only to demonstrate the 
extraordinary advance represented by the new Great Observatory, but also as an independent 
validation of the data quality of the aging but prolific Explorer. 



'Guest Observer, International Ultraviolet Explorer. 

216 



i< 

T 

V) 

4 

'e 

o 

CD 



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I 



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- 1 


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a Aurigae Ab (GO III) 

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7 Draconis (K5 III) 



VA^Vs/\/'V^,VVhA^sAv,,,^^ 




J \ \ L 



J L 



1200 1300 1400 1500 1600 1700 
WAVELENGTH (A) 

Figure 1: - Coadded SWP-LO spectra of Capella and 7 Dra from the lUE Archives. 
Crosses flag a reseaux mark. Error bars indicate the statistical uncertainties of the coad- 
ded fluxes. Line identifications are provided in the Linsky et al. paper in this proceedings. 



2. LOW-DISPERSION lUE SPECTRA 

There are a large number of SWP-LO, SWP-HL and LVVP-HI spectra of Capella in the 
lUE Archives. I selected 4 representative SVVP-LO images for the present work, with exposure 
times of 0.5-2.5M (5M total) to increase the effective dynamic range. For 7 Dra, there are 
4 SWP-LOs in the Archives: the 3 reliable spectra total 500M. Fig. 1 illustrates the coadded 
SWP-LO spectra of the two stars. These tracings should be compared with Figs, f-4 from the 
Linsky et al. paper: the G140L exposure times were of order 0.5M for Capella and iOM for 
7 Dra. The difference in sensitivity of the two low-resolution modes is much greater than the 
simple ratio of the exposure times, because the S/N is higher in the HUBBLE spectra, and the 
noise refers to pixels that are 5x smaller in wavelength than those of the lUE. Encouragingly, 
the overall spectral structure in the 1150-1750 A is cjualitatively the same in the GHRS and 
lUE tracings. Nevertheless, a seasoned lUE observer would certainly hesitate to identify the 
3<T feature near 1550 A in the lUE coadded spectrum of 7 Dra as C IV; whereas the feature is 
highly significant in the GHRS spectrum, and the identification as C IV is made all the more 
secure because both components of the doublet are present, in the expected intensity ratio, at 
the higher dispersion of the G140L mode. 

3. HIGH-DISPERSION lUE SPECTRA 

I reduced a series of lUE SWP and LWP echelle spectra of Capella, taken near opposite 
radial- velocity extrema in the orbit. The original observations were conducted using a number of 
techniques (including pseudo-trailing, multiple Offset Reference Points, and graded exposures) 
to push the S/N and dynamic range of coadded spectra beyond the usual limits. 

Fig. 2 illustrates the Mg II h and k lines observed near the opposite orbital quadratures. 



217 




2790 



2795 2800 2805 

WAVELENGTH (A) 



Figure 2: - LWP-HI spectra of Capella in vicinity of Mg II h and A; resonance lines. 



The S/.\" in tiiese spectra is liigli (> 50:1): eacii I racing represents tlie sum uf at least three 
independent pseudo-trailed spectra, and iOM of total integration time. Also shown is a simple 
model of the relative contributions of the active GO III secondary ("F"), the less-active G9 
III primary ("G"), and the interstellar Mg II absorption components C'LISM"). The solid curve 
depicts the sum of the three model contributions: it is similar (at phase 0.29) to the HST/GHRS 
ECH-B spectra (see Fig. 10 in Linsky et al.) and schematically illustrates the origins of the 
distinct spectral structure in the high-resolution profiles. While the GHRS ECH-B spectrum 
also recjuired about IOM of integration, the S/N (and the noise characteristics) are better than 
the coadded lUE spectrum; again the noise refers to smaller wavelength steps; and the factor of 
sa 8 better resolution permits a whole new regime of scientific inquiry unavailable to the lUE. 

■ Fig. 3 illustrates several intervals in the sub-2000 A lUE spectrum of Capella coinciding 
with the medium-resolution (or ECH-A in the case of Lyo) GHRS spectroscopy reported by 
Linsky et al. Here, the lUE tracings represent the coaddition of two or more independent pseudo- 
trailed SWP-HFs with a total exposure time of 400M (A < 1800 A) or ^ 60M (Lycv & A > 1800 
A). The solid curves refer to phase 0.29 (similar to that of the GHRS work), while the dashed 
curves refer to the opposite orbital ciuadrature. The overall shift of the high-excitation emissions 
between the opposite velocity extrema is clear: they follow the fast-rotating chromospherically- 
active secondary star. Panel (a) should be compared with Figs. 11 and 9 of Linsky et al.; panel 
(b) with Figs. 8 and .5; panel (c) with Fig. 7; and panel (d) with Fig. 6. Aside from the stunning 
GHRS echelle spectrum of Lyo, the medium-resolution spectra of Capella are comparable in 
resolution to those of the lUE, although the S/N is clearly higher in most of the HUBBLE 
observations (in about l/.50-th the equivalent lUE exposure time!). This is particularly true of 
the fainter emissions in each interval: compare, for example, the diagnostically-critical O IV] 
lines in the 1400 A region. 



218 



T 

I 



o 




o 



X 

3 



Figure 3: 
vals near 




1885 



1890 



1895 



1900 



1905 



1910 



1915 



WAVELENGTH (A) 



- Coadded SWP high-dispersion spectra of Capella in selected wavelength inter- 
opposite quadratures in the binary orbit. Crosses flag reseau marks or saturation. 



4. CONCLUSIONS 

Space limitations prevent a more exhausive comparison than this. Nevertheless, one can 
safely conclude the following: (1) A large-aperture space telescope with a modern spectrograph 
and detectors produces beautiful (dare 1 say, solar-ciuality) vacuum-ultraviolet spectra; and (2) 
even so, the quality of the lUE spectrograms is surprisingly good, at least with respect to the 
preconceptions of this all-too-knowledgeable observer. Thus, the HST/GHRS presents not only 
a powerful new spectroscopic tool for the nineties, but also a critical validation of the quality of 
the spectral material in the extensive archives of the lUE, the workhorse UV space observatory 
of the eighties. 

This work was supported by NASA Grants NAG5-199 and NAG5-r2i5. 



219 



Faint Object Camera In-flight Performance 

Geometric Distortion, Stability and Plate Scale. 



Dave Baxter 

Space Telescope Science Institute 

3700 San Martin Drive 

Baltimore, MD. 21218. 



Abstract 



The geometric distortion characteristics of the Faint Object Camera have been analysed in 
great depth and it has been found that the distortion pattern is remarkably stable. The positional 
variations in the reseau pattern, over the central 512x512 region of the photocathode, from image 
to image, have an RMS value of ~1 pixel, (~40mas, 20mas and 7mas for the f/48, f/96 and f/288 
modes, respectively). Of this, 0.25 pixel represents the uncertainty in the individual reseau position 
caused by the effects rebinning and photon noise on the reseau itself. 

Low levels of saturation appear to have little or no effect on the stability of the detectors, 
however high flux rates across the full area of the photocathode, (particularly through the f/48 
relay), can cause a permenent change to the distortion pattern. 

The plate scales of the 3 imaging modes have been determined and are found to be very close 
to nominal. The values obtained are; f/48: 0.04526"/pixel, f/96: 0.02217"/pixel, and, f/288: 
0.007467pixel. 



1. Geometric Stability 

In order to carry out geometric correction of FOC data, i.e. to recover an image in 
vyhich the spatial relationships between objects are restored, a necessary requirement is 
that the geometric distortion field, shown in Figure 1, must be stable. By this we mean 
that there must be no significant change in the observed reseau positions with time. 

It has been noted that short term variation of the geometric distortion pattern occurs 
during the period immediately following FOC high voltage switch-on. During this time 
the observed reseau positions show an RMS deviation from the stable positions of approx- 
imately 3— 'l pixels. This period however, extends for only about 40 minutes, by which 
time the reseau positions have stabilised to within ~0. 25- 1.00 pixels. In order to avoid 
this period of instability, the scheduling software automatically inserts a time delay of 
40 minutes immediately following high voltage switch-on, which prevents exposures being 
taken during this time. 

Long term variation could possibly occur as a result of out-gassing in the instrument, 
however monitoring of the geometric distortion pattern over the last nine months has 
shown that, in general, both cameras are remarkably stable. As an example, in Figure 
2 we show the RMS deviation of the F/96 reseau positions from the mean positions. 
The mean positions are obtained by averaging the measured positions (from the central 



220 



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Figure 1: The total (i.e. optical+detector) geometric distortion pattern of both the F/48 
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on a grid interval of 60x60 pixels, with the vectors at a magnification of x2 for clarity. 









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300 320 340 

Dav Number ('from Ol-JAN-90) 



360 



380 



400 



Figure 2: The RMS deviation of the reseau positions, from the mean, over the period 1st 
September 1990, to present. 



221 



PATTERN ROTATION 



RAOWU. PATTERN SHIFT 




Ooy Humbw 
X-SCAUNC TACTOR 



y- SCALING f ACTOR 




Oav Number 



Dav Numbaf 



Figure 3: The effect of high saturation levels on the F/48 detector. This event took place 
on 22 July 1990 and caused a large change in the distortion field, involving ~ 1.5° of 
rotation, a 25 pixel pattern shift and a 1-2% change in the platescale. 

512x512 region) over the period from 1st September 1990, up to the present. It can be seen 
that the deviation is small, (% 0.85 pixels), and of this, ~ 0.25 is the intrinsic uncertainty 
in an individual reseau position due to rebinning and photon noise. 

One effect which has been noted however, is that the detector stability is rather susept- 
able to change if highly saturated. An event of this type occurred on 22nd July 1990 when 
both the F/48 and F/96 detectors were illuminated by the bright Earth to obtain a series 
of external flatfields. The resulting f/48 images were very highly saturated (>10 times 
the nominal saturation rate) and showed a sudden, and dramatic change in the distor- 
tion pattern amounting to a rotation of about 1.5 degrees and a shift of about 25 pixels 
(Figure 3). This change recovered gradually over the subsequent 5-6 weeks and the F/48 
geometric distortion pattern is now stable again, although not in the same position as 
prior to the saturation event. The current 'stable' position is offset by ~10 pixels from 
its pre- saturation position. The F/96 detector also showed a disruption of the geometric 
stability at that time, however since the F/96 pixel has a smaller angular area, (by a factor 
of 4) than the f/48 pixel, the incident count rate was smaller by the same factor and hence, 
the level of saturation was much lower. The F/96 detector returned almost immediately 
to the former stable position. 



222 



c 



o 
> 



o 
(J 

CO 

cu 



o 






-H 1 1 H 



CP 



Q 'U 

D 



H 1 1 H 



-1 



-0.5 0.5 

Platescale Variation in X (%) 



Figure 4: Variation of the F/96 platescale since 1st September 1990. 

2. Plate Scale 

The plate scale (i.e. the size of the pixels in arc seconds) has been determined for 
the two cameras in the FOC. This is done by taking a series of overlapping images of a 
crowded star field, moving the telescope between exposures by a known angidar offset. The 
measured distances (in pixels) between the same stars on adjacent exposures combined 
with the known offset (in arc seconds) then give us the plate scale. 

For the F/96 relay this was determined to be 0.02217 arcseconds pixel"^ (± 0.00010) 
and for the F/48, 0.04514 arcseconds pixel"^ (± 0.0005). These values are 'radial' plate 
scales and are within a few percent of the nominal values, vis. 0.022 arcsecond pixel" ^ 
for F/96, and 0.044 arcsecond pixel" ^ for F/48. Figure 4 shows that the platescale (at 
least for F/96), has remained very stable over the report period, staying within ±0.2% 
(~ 5xlO~^ "/pixel) of nominal. This is equivalent to an uncertainty in relative positions 
of s;0.025" over the width of the F/96 512x512 format. Although the analysis of the F/48 
data indicated a possible difference in the x- and y-platescales, subsequent examination of 
other data does not confirm this. Because of this discrepancy however, the F/48 platescale 
is assigned a somewhat higher uncertainty than F/96. A new proposal has been designed 
to determine unambiguously the absolute values for the platescales in the F/48 and F/96 
relays. This should be completed by late 1991. 



223 



The F/288 relay has a nominal value of 0.007 arcsecond pixel"^ and a derived value 
of 0.00746 arcsecond pixel"^ Because of the degraded optical performance of the HST, 
due to spherical aberration, an isolated pointsource scatters significant amounts of light 
out to a radius of ~ 2". This, combined with the small field size and the need for image 
restoration techniques, severely limits the usefulness of the imaging mode through the 
F/288 relay, and therefore we have no plans to investigate further, the platescale for this 
relay. 



224 



IN-FLIGHT PERFORMANCE OF THE FOG: 

EARLY ASSESSMENT OF THE ABSOLUTE SENSITIVITY. 



W.B. Sparks and the FOC IDT 
Space Telescope Science Institute, 
3700 San Martin Drive, 
Baltimore, MD 21218, 
USA. 



Abstract. Observations with the Faint Object Camera on the Hubble Space Telescope 
in the f/96 imaging mode indicate an absolute sensitivity consistent with nominal (as 
given in the Instrument Handbook) for wavelengths longward of about 2500 to 3000 A. 
Shortward of that, there is a smooth decline reaching approximately 60% relative to 
the baseline by 1200A. No secular changes have been identified at this stage. 



1. INTRODUGTION 

Knowledge of the absolute efficiency of the FOC in combination with the OTA 
is essential in estimating program feasibility. Here, observations of UV photometric 
standard stars observed during the OV phase (up to February 1991) are analysed to 
provide an initial assessment of the absolute sensitivity of the FOC as a function of 
wavelength and of time. The analysis procedure is simply to derive total count rates 
for each observation and to compare the result to a prediction using the best estimate 
of the input spectrum together with all component throughput curves and instrument 
DQE. Departures of 'observed/predicted' count rates from 1.0 indicate inconsistencies 
between the simulations and the observations. See Greenfield et al. 1991. 

The present analysis is based on data acquired for other purposes — focus moni- 
toring, UV first light, UV throughput monitoring (OLT) and FOC SAO PSFs (Science 
Assesment Observation Point Spread Functions). It is therefore incomplete in wave- 
length coverage and has sparse time coverage. More recent science verification observa- 
tions (proposal 1511 in particular) together with continued UV throughput and focus 
monitoring will enable a more thorough study of these issues to be undertaken in the 
future. 



225 



2. OBSERVATIONS AND ANALYSIS 



2.1. UV standard star observations 



Two UV photometric standard stars were observed during the period from launch 
to Day 51, 1991: BPM 16274 and GRW+70°5824, Bohlin et al. 1987, Turnshek et al. 
1989, Turnshek et al. 1990. Spectrophotometry from the visible through to the UV is 
available within CDBS. The stars were observed at a variety of wavelengths and for a 
variety of purposes. There are a total of 72 images. 



2.2. Analysis Procedure 

The principle observable is total count rate. To estimate this, the z'ra/ implementa- 
tion of daophot was used together with IDL. Counts interior to 3.-3 arcsec were used for 
the total, while the sky level was estimated from the outer region of the stellar intensity 
profile. Simulations of expected count rate assuming particular throughput curves were 
made using the iraj/stsdas package synphot, and the two were compared. 



3. RESULTS 

There is very strong Lyman a absorption situated within the peak transmission of 
the F120M filter for these UV photometric standards. The far-UV response estimate 
must therefore be considered more uncertain than the others, although both stars do 
give the same value for the sensitivity at Fr20M. Figure 1 shows the ratio of observed 
to predicted counts as a function of wavelength (the pivot wavelength, Koornneef et al. 
1986). A fourth order polynomial modification to the DQE curve using a least squares 
fit to these data (excluding the uppermost outliers) is also included, derived using the 
synphot program fitband. 

For wavelengths longer than about 2500 to 3000A there is no evidence for departures 
from the nominal DQE as given in the Instrument Handbook, and assumed for the 
'prediction'. There are some data points above a value of unity. The two reddest are 
also the observations with the narrowest filters. Because of this, the sinriulation becomes 
more uncertain and these two points should at this stage only be used as indicative that 
there are no serious problems in the red part of the spectrum. Shortward of 2500A there 
appears to be a fairly smooth decline with wavelength, reaching about 60% of baseline 
with the F120M (^ 1200A) filter. 



226 






CO 



SYNPHOT.PLRATIO 
2 I — I — I — I — I — I — I — I — I — I — I — 1 — I — I — r 



1.5 



T — I — r 



REFSPEC = none 
PHOT = fb_input.ddt 




J 1 1 1 I I I \ I \ I 1 I I I 





1000 2000 3000 4000 5000 6000 

WAVELENGTH (A) 



Figure 1. The ratio of observed to calculated count rate as a function of wavelength. 



227 



A potential source for apparent reduced UV sensitivity is large angle scattering from 
mirror micro-roughness. In order to assess whether this is a likely factor, the average 
of 15 F120M profiles from GRW+70°5824 was derived. Relative to the flux interior to 
4.2 arcsec, by 2.5 arcsec the profile has reached 97% and by 3.3 arcsec, it has reached 
99% of the total. There is no evidence therefore for much scattering into radii of order 
two to four arcsec. 

The values of 'observation/prediction' versus time were analysed for visible and far 
UV filters. The visible data have no significant trends at all, while in F120M there is 
marginal evidence for an early decrease, although again the statistical significance is 
poor. Continued UV throughput monitoring will be carried out. 

4. CONCLUSIONS 

Results of a preliminary investigation into the f/96 absolute sensitivity of the FOC 
have been presented along with a description of the analysis techniques used. The 
results indicate that the DQE is consistent with nominal (as given in the Instrument 
Handbook) for wavelengths longward of about 2500 to 3000 A. Short ward of there, there 
is a smooth decline reaching approximately 60% relative to the baseline by r200A. 

REFERENCES 

Bohlin, R.C., Blades, J.C., Holm, A., Savage, B.S., Turnshek, D.A. 1987, Standard 
Astronomical Sources for HST: 1. UV Spectrophoiometnc Standards. STScI Publication. 

Greenfield, P., Paresce, F., Baxter, D., Hodge, P., Hook, R., Jakobsen, P., Jedrzejewski, 
R., Nota, A., Sparks, W.B., Towers, N., Laurance, R., Macchetto, F. 1991, SPIE 
Conference on Space Astronomical Telescopes and Instruments, in press, STScI 
preprint 536. 

Koorneef J., Bohlin, R.C., Buser, R., Home, K.D., Turnshek, D.A. 1986, Synthetic 
Photometry and the Calibration of the Hubble Space Telescope, Highlights of 
Astronomy, 7, p. 833, ed. J. -P. Swings. 

Turnshek, D.A., Baum, W.A., Bohlin, R.C., Dolan, J.F., Home, K., Koornneef, J., Oke., 
J.B., Williamson, R.L. 1989, Standard Astronomical Sources for HST: 2. Optical 
Calibration Targets. STScI Publication. 

Turnshek, D.A., Bohlin, R.C., Williamson, R.L., Lupie, O.L., Koornneef, J., Morgan, 
D.H. 1990, Astron. J.,99,1243. 



228 



IN-FLIGHT PERFORMANCE OF THE FOC: FLAT FIELD RESPONSE. 



P. Greenfield and the FOC IDT. 
Space Telescope Science Institute 
3700 San Martin Drive 
Baltimore, Maryland 21218 



1. DETERMINING THE FLAT FIELD RESPONSE 

This paper addresses the spatial response of the FOC detectors. Details of the FOC's 
operation and other aspects of FOC performance may be found in Paresce (1990) and 
Greenfield et al. (1991) as well as other papers at this conference. 

The spatial response of the FOC's detectors, like most other photon counting detectors, 
is not constant over the field of view. Determining the flat field response is simple in 
principle, more difficult in practice. An ever present difficulty is the relatively limited 
linear count rate of the detectors. The fuU field video format (512z x 1024) becomes 
more than approximately 10% nonlinear at count rates above 0.05 counts pixel"^ s~^ 
using flat field illumination. This results in very long exposure times if counts of 
several hundreds or thousands are required which, in turn, means that data related 
to flat field response will be limited either in the number of counts, wavelengths, or 
formats obtained. 

In the visible, the onboard LEDs provide a convenient means of illumination, but, in 
the ultraviolet, it is more difficult to find a suitable source of illumination. The two 
candidates for UV flat field illumination have been the bright earth and a selected area 
of the inner Orion nebula. Although flat fields using the bright earth are attractive 
in the sense that illumination should be relatively flat because they are "streaked" 
over long paths on the ground and can be taken when the telescope would otherwise 
be idle, they have large variabiHty in brightness, possible visible light contamination 
through the filters, and require a large attenuation of the fight level. This is especially 
a problem for //48 where the filter selection is relatively Hmited and there are no 
neutral density filters, thus ruling out this source as a method of obtaining UV flat 
fields for this optical relay. Flat fields using the Orion nebula are compUcated by 



229 



the fact that the field illumination is not uniform. Fortunately, this region of the 
nebula looks relatively smooth in the UV and future observations should be able to 
determine the brightness distribution of the nebula, and therefore determine the UV 
spatial response. Nevertheless, so far virtually all information about the flat field 
response has been obtained from LED exposures or ground-based flat fields. 

2. FLAT FIELD PROPERTIES 

Figures 1 and 2 show //96 and //48 full format (512 x 1024 using "zoomed" pixels) 
flat fields. The //96 flat field was taken using the inner region of the Orion nebula 
and the F140W filter, and as a result the occulting fingers are visible. The //48 flat 
field was obtained using an LED. The nonuniform spatial response of the detectors is 
primarily due to the varying response of the bialkali photocathode which is evidenced 
by large and small scale variations (such as scratches), but there are other efi^ects also. 
Variations in the camera tube target response also results in localized nonuniformity 
of response. The most obvious of these is the apparent outline of the smaller 512 x 512 
video format within the larger video formats. This is apparently caused by degradation 
of the camera tube target by the extra dweU time of the electron read beam at the 
edges of the 512 x 512 format when it is In use. This format "burn-in" appears to 
worsen with time; however, it afl'ects a relatively small part of the total detector area. 
Also noticeable, but usually minor in its effect except for the smallest video formats, 
is the read-beam flyback after each frame scan. The photocathode response tends to 
vary the most in the far UV where areas and scratches may show up to 30% drops in 
sensitivity. 

Much of the apparent nonuniformity of the raw images is a result of detector geometric 
distortion where ±10 — 20% variations may result, mostly near the format's corners. 
Compare for example a //96 512 x 512 flat field shown in Figure 3 with the full 
format //96 flat field. There are clearly areas in the 512 x 512 image where there is 
noticible nonuniformity that do not appear in Figure 2 because the geometric distortion 
characteristics vary with the format being used. For that reason, flat fields obtained 
in one format cannot be applied to other formats without correction for geometric 
distortion effects. Except for areas near the edges where the distortion is severe, this 
component can be removed by applying a flux-conserving geometric correction to the 
images. Figure 4 is a plot of a row of Figure 3 to show more quantatively the size of 
nonuniformities in the flat field (note however, the relatively large photon noise because 
of the relatively low counts). Figure 5 shows the size of the effect of the geometric 
distortion on the //96 full format flat field response as a function of position. 

Both detectors exhibit a couple of forms of "pattern noise." One is believed to result 
from a moire fringe effect involving the camera tube's read beam interacting with 



230 






Figure 1 — An//96 full format exposure of the 
inner region of the Orion nebula taken with the 
F140W filter. The two coronographic fingers 
can be seen as well as the reseau marks, 
scratches, and the "bum-in" of the 512 x 512 
format in the center. The exposure has typically 
40 counts per zoomed pixel and is somewhat 
nonlinear. 



Figure 2 — An//48 full format exposure taken 
using an onboard LED. The dark comer in the 
upper left is a result of vignetting of the LED 
and does not appear in external exposures. 





Figure 3 — An//96 512 x 512 pixel format 
exposure taken using an onboard LED. 



100 200 300 400 500 600 

Figure 4 — A plot of row 300 of figure 3. 



231 




Figure 5 — This image shows the effect of the 
geometric distortion on the photometric re- 
sponse. Displayed is the result of geometrically 
correcting a perfectly flat full format image. 
Those areas that are brighter correspond to areas 
on the detector that had been stretched out and 
had larger than average pixels. The overlaid 
contours (the solid lines) are spaced every 5% in 
intensity relative to the center. 



L 



Figure 6 — An f/96 512 x 512 pixel format 
exposure taken using an onboard LED at a high 
intensity setting to highlight the diagonal pattern 
noise which is enhanced when the detector is 
driven into its nonlinear regime. 




Figure 7 — The amplitude of a two-dimensional 
FFT showing the presence ofa512x512 
format //96 image. The grid in the center is due 
to the reseaux. The diagonal pattern noise shows 
as extended peaks in the upper left and lower 
right while the 4-pixel pattern can be seen as 
sharp spikes in the center row. 




Figure 8 — The plot of the amplitude of the 
central row of the FFT showing the presence of 
the 4-pixel pattern. This pattern is believed to 
result from a digital clock in the electronics that 
has a 4-pixel period. 



232 



both a tube grid and the target diode array. The end result is a roughly sinusoidal 
modulation of the response with a spatial frequency of 3.3 pixels running in a nearly 
diagonal direction for the //96 detector with a ~ 5% RMS amplitude. The amplitude 
for the //48 detector is approximately half that. This pattern is generally enhanced 
when nonlinear count rates are encountered, and for that reason can easily be seen in 
Figure 6 which is an //96 LED exposure taken at relatively high count rates. The 
other pattern consists of vertical stripes appearing with a horizontal period of 4 pixels 
and originates from a digital clocking signal. The RMS amplitude of this pattern 
is 2 — 3% for both detectors. Figure 7 shows the amplitude of the FFT of a //96 
flat field where the presence of both types of patterns can be seen. Figure 8 is a 
plot of the center row of the FFT showing the characteristics of the vertical pattern. 
Finally, there is evidence that there is a significant fine scale pixel-to-pixel variation in 
response not characterized by any of the previous effects. It is most likely an intrinsic 
granularity of the photocathode, but that has not been clearly shown and has not been 
well characterized. 

Our estimates are that relative fluxes are currently accurate to only about the ~ 10% 
level if they are determined from an area of at least 10 pixels. Our long range goal is 
to improve the flat field calibration so that relative fluxes can be determined to the 3% 
level. 

3. ACKNOWLEDGEMENTS 

The Faint Object Camera is the result of many years of hard work and important 
contributions by a number of highly dedicated individuals far too numerous to list 
individually. In particular, we wish to thank the ESA HST Project Team and the 
European contractors for building an outstanding scientific instrument. All of the 
FOC images were taken using the NASA/ESA Hubble Space Telescope, obtained at 
the Space Telescope Science Institute, which is operated by AURA, Inc., under NASA 
contract NAS 5-26555. The FOC IDT Support Team: D. Baxter, P. Greenfield, and 
R. Jedrzejewski, acknowledge support from ESA through contract 6500/85/NL/SK. 

REFERENCES 



F. Paresce, Faint Object Camera Instrument Handbook, Space Telescope Science In- 
stitute, Baltimore, 1990. 

Greenfield, P., Paresce, F., Baxter, D., Hodge, P., Hook, R., Jakobsen, P, Jedrze- 
jewski, R., Nota, A., Sparks, W. B., Towers, N., Laurence, R., Macchetto, F., 
Proceedings of the SPIE Conference on Space Astronomical Telescopes Instru- 
ments, Orlando, Florida, 1 April 1991. 



233 



BACKGROUND NOISE REJECTION 

IN THE FAINT OBJECT SPECTROGRAPH 



Rosenblatt, E.I., Beaver, E.A., 
Linsky, J.B., and Lyons, R.W. 

Center for Astrophysics and Space Sciences 
University of CaJifornia, San Diego 
La JoUa, CA 92093-0111 



Abstract. We have modeled the background noise of the FOS "blue" MgF2 de- 
tector to investigate the optimal strategies for reducing the noise level in spectra. 
Background observations made with the FOS during the Orbital and Science Ver- 
ification periods (June 1990 to the present) have shown that the dominant source 
of noise is of a non-poisson burst character most hkely produced by Cerenkov radi- 
ation. This radiation will be emitted whenever a high energy particle traverses the 
detector faceplate, and can result in large portions of the diode array being flashed 
nearly simultaneously. We have modeled the effects of Cerenkov radiation in image 
tube faceplates by means of a Monte Carlo numerical simulation. This model pro- 
duces images and count statistics which are in good agreement with actual data. 
This simulated background data has allowed us to determine the rejection thresh- 
olds and frame times that yield the highest S/N ratio for a stellar source of any 
given flux level. 

1. MONTE CARLO SIMULATIONS OF FOS BACKGROUND NOISE 

A detailed analysis of FOS dark data has shown that the dominant source 
of background noise produced in the low altitude orbit of the HST results from 
Cerenkov radiation (Rosenblatt etal. 1991). This noise occurs whenever a high 
energy particle (E >, 300 MeV) such as a cosmic ray traverses the window of the 
Digicon detector and produces a burst of photons that can flash large portions of 
the diode array simultaneously. This source of noise generates at least 90% of the 
FOS background both inside and outside of the South Atlantic Anomaly. 

Without an accurate subtraction of the background, measurement of the shape 
and absolute level of the continuum of astronomical objects is compromised. Fig- 
ure 1 shows an uncalibrated spectrum of the quasar CSO 251 recently observed with 
the FOS. A third order polynomial fit to the diode response is also shown for the 
estimated background level. Note that even for this relatively bright quasar, the 
background is significant compared to the continuum flux level. For fainter objects, 
it is even more important to estimate and subtract the background accurately. 



234 



In order to better understand FOS dark background noise, we have developed 
a Monte Carlo model which closely simulates the FOS detector and the physical 
characteristics of Cerenkov radiation. This model allows a large number of parti- 
cle events to be simulated and analyzed. The code first generates random impact 
positions and angles (weighted by sohd angle) with respect to the faceplate. Since 
the number of Cerenkov photons produced depends on the atomic number of the 
incident particles, a cosmic ray abundance consisting of 91% protons and 9% alpha 
particles is input. Cones of Cerenkov hght are generated along the particle path in 
the faceplate. These cones are divided into narrower sub-cones for greater resolu- 
tion and accuracy in calcidating MgF2 absorption effects. The sub-cones are then 
projected down onto the photocathode and Poisson statistics are used to determine 
the number of photons incident on the photocathode within the sohd angle sub- 
tended by ciny given diode. Binomial statistics (based on photocathode Q.E.) are 
then used to determine whether a photoelectron is emitted from a specific area of 
the photocathode. Since the Digicon employs one-to-one imaging, when a photo- 
electron is emitted directly in fine with a diode, a count is registered in that diode. 
In this way, the differential and cumulative hits on each diode in the array can be 
traced throughout the simulation. 

One major advantage of this technique is that the counting statistics {i.e., the 
number of counts registered on the diode array per particle event) of the background 
noise can be analyzed. This information is difficult to acquire directly with the FOS 
due to the short time sampling required to separate one event from another. Our 
largest simulation to date includes over 16,600 particle events. Surprisingly, the 
vast majority of these events (~14,500) did not produce any counts in the array 
whatsoever. This result is due to a combination of factors including absorption in 
the faceplate, geometric dilution of photons, and poisson statistics. Another reason 
is simply that many of the projected conic sections did not intersect the diode 
array and therefore do not generate hits. Of the ~2000 events that did produce 
one or more counts, most generated 1-5 counts. However, a small number of events 
produced many hits at the array which quahtatively resemble the streaks observed in 
FOS dark data. The average count rate of this simulation (0.0056 cts s"^ diode" ^) 
agrees well with the observed rate (0.0062 cts s~^ diode"^) at low geomagnetic 
latitude where the background is at a minimum. 



FAINT OBJECT SPECTROGRAPH QUASAR SPECTRUM 



10000 



K300 



z 

o 
o 




250 375 

DIODE NUMBER 



Fig. 1 — An uncedibrated 
spectnun of the quasar CSO 
251 recently observed with 
the FOS is shown together 
with a third order polyno- 
mieilfit to the estimated back- 
grotmd level. Note that even 
for this relatively bright quasau', 
the background is signifi- 
cant compeired to the con- 
tinuum flux level. For fainter 
objects, it would be even more 
importeint to estimate and 
subtract the background ac- 
curately. 



235 



2. BACKGROUND NOISE REJECTION 



Due to the burst character of Cerenkov light, there is a potentially powerful 
tool available onboard the FOS that will ehminate a significant amount of the dark 
background noise from astronomical data. This rejection capabihty is suppUed by a 
burst rejection algorithm (known as REJLIM) that can be set to different thresholds 
such that any noise burst (summed over the entire diode array) registering at or 
above the threshold within some specified Uve time will be rejected from the data 
stream (all data within this time interval wLU be rejected). The FOS rejection 
software allows a minimal Uve time of 20 ms in which the total number of counts 
for the array are summed. Following each hve time is a deadtime of 10 ms in which 
the electronics are reinitialized. Simulated background data from our Monte Carlo 
model has allowed us to determine rejection thresholds and hve times that yield the 
highest S/N ratio for a stellar source of any given flux level. 

An example of our study is presented in Table 1. The background noise count 
rate was set to 0.01 cts s~^ diode" ^ in this simulation (corresponding to low geo- 
magnetic latitudes). A stellar source was modeled using poisson statistics with a 
mean of 0.004 cts s~^ diode" ^. For each stellar flux level, the hve time was varied 
from 20 to 500 ms and the threshold that resulted in the highest S/N ratio in each 
case was determined. Column 1 of Table 1 gives the Uve time in milliseconds, Col. 2 
is the threshold (in units of counts per hve time per array) that yielded the highest 
S/N ratio which is given in Col. 3, Col. 4 is the total time removed from the data 
stream due to rejected frames and deadtime, Col. 5 is the total deadtime. Col. 6 is 
the total hve time rejected. Col. 7 gives the actual exposure time after the reject 
time of Col. 6 is subtracted from the original integration time, Col. 8 is the fraction 
of stellar counts rejected, and Col. 9 is the fraction of dark counts rejected. 



Tabic 1 
Simulated Burst Noise Rejection Data 

Exp time = 1158a Stellai flu = 0.004 c/i/d Daik noiie = 0.01 c/i/d 

fiametime threshold S/N time lost deadtime rej time live time % star rejd % dark rejd 
(m») (<!"') ('ec) (sec) (sec) (sec) 

(1) (2) (3) (4) (5) (6) (7) (8) (9) 



20.0 


2 


1.31 


396.6 


386.0 


10.6 


761.4 


0.02 


0.72 


lOO.D 


2 


1.45 


184.3 


105.3 


79.0 


973.7 


0.12 


0.78 


300.0 


3 


1.33 


247.4 


37.4 


210.0 


910.6 


0.28 


0.77 


500.0 


4 


1.21 


322.7 


22.7 


300.0 


835.3 


0.36 


0.77 



236 



To illustrate the improvement in S/N that optimal burst rejection can provide 
we note that in the example above the S/N ratio found with a frame time of 100 ms 
and a rejection threshold of 2 is 25% higher than the S/N ratio that would be 
obtained without any noise rejection. Moreover, only 12% of the stellar flux was 
rejected, while 77% of the noise was eHminated. 

During the Spring 1991, the FOS obtained dark data with REJLIM enabled at 
threshold settings of 2, 3, 8, and 10. Figure 2 compares our Monte Carlo simulation 
(fiUed dots) with these observed data (asterisks). Although the Cerenkov model 
agrees fairly well with the observations, the addition of a 0.002 cts s~^ d~^ pois- 
son noise source (open dots) results in the best fit. The figure clearly shows that 
threshold settings of 8-10 or greater will have little effect on the observed count rate 
(roughly 90% of the counts registered relative to what would have been detected 
with REJLIM disabled). Thus, large thresholds will not reduce the background 
noise significantly. However, if the astronomical source is faint, the threshold can 
be set to small values, ehminating a substantial fraction of the taackground. 



Fig. 2— The Monte Carlo 
background noise simula- 
tion (filled dots) data is com- 
pared to actual FOS deirk 
data with REJLIM enabled 
(asterisks) at threshold set- 
tings of 2, 3, 8, and 10. Al- 
though the agreement is fcurly 
good, the best fit is achieved 
with the inclusion of a 0.002 
cts s~^ d~^ poisson noise 
source. Significant noise re- 
duction is achieved at smedl 
thresholds for faint astro- 
nomical sources. 



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. 


1 


—Jl 1 


1 


. . 1 . 


. i 1 







4 6 8 

THRESHOLD SETTING 



10 



12 



An important point to consider when applying REJLIM is zeroth order Hght, 
which in some observing configurations will register onto the diode array. Although 
the zeroth order hght falls on a different section of the array than the spectrum, 
it will still be summed within a frame time and compared to the threshold. Thus, 
with zeroth order fight included, an otherwise correct threshold setting might reject 
aU frames and no data wiU be acquired. For this potential problem to be avoided 
requires the specific diodes on which the zeroth order Hght fails to be turned off so 
that these counts are not included. 



The authors thank Rick Hier for useful discussions. This research was sup- 
ported by NASA NAS 5-24463/NAS 5-29293. 

REFERENCES 

Rosenblatt, E. I., Beaver, E. A., Cohen, R. D., Linsky, J. B., and Lyons, R. W. 
1991 SPIE Proceedings on Electron Image Tubes and Image Intensifiers II, ed. I. P. 
Csorba (Bellingham, WA: SPIE), 1449, p. 72. 



237 



DETECTION OF BINARIES WITH THE FGS: 
THE TRANSFER FUNCTION MODE DATA ANALYSIS. 

B. Bucciarelli, M. G. Lattanzi, and L. G. Taff 

Space Telescope Science Institute 

3700 San Martin Drive, Baltimore, MD 21218 

O. G. Franz and L. H. Wasserman 

Lowell Observatory 

Mars Hill Road 1400 W., Flagstaff, AZ 86001 

E. Nelan 

University of Texas 

Dept. of Astronomy, Austin, TX 78712-1083 



1. TRANSFER FUNCTION MODE GENERALITIES AND 
A DESCRIPTION OF THE DATA 

The Science Assessment Observations program (SAO) has given each sci- 
entific instrument on board the HST the first real opportunity to gather data 
usable for assessing, to a reasonable level of confidence, their scientific potential 
shortly after the problem in the telescope optics was found and well before the 
Science Verification (SV) had started. For what follows, it is therefore impor- 
tant to bear in mind that FGS3, the astrometer unit used to take our SAO 
observations was a totally uncalibrated device. 

The FGS SAO program was devoted to the detection of possible, yet undis- 
covered, binaries among the bright members of the Hyades cluster. Out of 
21 candidates, 16 targets were successfully observed in the so-called Transfer 
Mode, which consists of multiple scans through the target object executed in 
the following fashion. The FGS star selectors are driven in such a way that the 
aperture sweeps over the star at an angle of 45° to the X,Y FGS reference frame. 
The effect is to sample the visibility fringes (or Transfer Functions-TFs) of both 
the Koester prism interferometers in the FGSs (one per axis). The length of 
the scans is approximately 1.5" (on both axes) with an average sampling step 
of about 1 mas. Two representative scans, one taken with the neutral density 



238 



filter (Neut/ND5) and the other with the clear filter (Clr/583W), are shown in 
Figs, la and lb (dashed curve) for the X-axis. 

The TF scan raw data consist of a time series of photomultiplier counts 
and star selector encoder readings. These data must be transformed into the 
proper units and corrected for instrumental and other effects before the curves 
shown in Fig. 1 are obtained. The corrections include PMTs mismatch, sky 
background, spacecraft jitter, and velocity aberration. The imbalance between 
the two PMTs on the same axis is easily accounted for by averaging a set of 
the samples drawn from both tails of the TF scan. Sky background is not an 
issue since the targets are quite bright stars. Finally, given the early stage of 
the Science Verification observing, the refinements of jitter removal and velocity 
aberration were not considered (and only differential aberration really counts). 



2. REDUCTION PROCEDURES AND EVALUATION OF THE 
RESULTS 

Assume now we know the form of the Single Star (SS) TF. The hypothesis 
that the incoming light from two different sources, close by in the sky, is in- 
coherent and the application of the superposition principle yield the expected 
Double Star (DS) TF [^^(s)] in the form of a linear combination of two SS TFs, 
i.e., 

D{x) = A{Am)[S{x) -f B{Am)S(x + dx)] (1) 

(and its analogue for the Y-axis), where the second SS TF [S{x + dx)] is iden- 
tical to the first [5(x)] but displaced along the X-axis by dx, the DS projected 
separation. A{Am.) is a normalization factor, and B is the intensity ratio of the 
primary to the secondary star. Both quantities are, of course, functions of the 
magnitude difference (Am). The model just described is fitted to the observed 
TF curve and the parameters dx and Am derived. It is worth noticing here 
that two independent estimates of Am are available, one for each FGS axis. In 
practice, a grid of models is generated by varying dx and Am. Each model is 
cross-correlated wih the observed TF by computing the correlation integral 



/ 



D{t - u) TF{t)dt , (2) 



where the template function D is being cross-correlated with the actual visibility 
fringe TF; the sought value for u, which maximizes Eq. 2, represents the shift 
along the horizontal axis between the two functions. 

The best-fit model is chosen as the cross-correlation that minimizes the sum 
of the squares of the differences between the model and the observed one., viz. 



/' 



D(t - u„) - TF{t)]^dt = min , (3) 



239 



RIter: F5ND (5-mag neutr. dens.) 



0.4 



0.2 



X 

or 



-0.2 




7.5 
X-axis (arsec) 

Fig. 1(a) 



""> 1 r- 



0.2 



0.1 



-I 1 1 1 p 

Filter: F583W (Clear) 
mag: 9.99 V 
Guiding Mode: FL 



-0.1 - 




_L 



8.5 
X-axis (arsec) 

Fig. Kb) 



240 



where Un is the value maximizing Eq. 2. This approach has been preferred to 
direct applications of least-squares-like schemes for its robustness, in relation to 
the range of narrow separations (from 100 mas down to about 10 mas) where the 
astrometer FGS will make its most interesting detections, and the independent 
difficulty of giving sufficiently accurate initial guesses for dx and Am. How- 
ever, if felt necessary, the accuracy of the fit can be improved, now using the 
answers from the correlation technique, as initial guesses for the least-squares 
final adjustment. 

Before running the cross-correlation, the noise in the observed TF is smoothed 
via a piecewise low-order polynomial fit, where continuity is imposed at the bin 
boundaries up to the specified derivative order (continuous lines in Fig. 1). 
This polynomial smoothing increases the resolution of the subsequent cross- 
correlation, and makes it possible to compute the correlation integral analyti- 
cally. 

The current OTA and, possibly, residual aberrations and misalignments 
within the FGSs are producing field dependent aberrations across the FGS fields 
of view. In terms of a single star TF, this means that, at the moment, we are 
unable to successfully predict either its shape or its marked variation across the 
field of view. The nice properties of the theoretical, pre-launch, TF are gone. 
To find the actual signatures of a single star TF as a function of the position 
in the field of view, one must resort to in-flight calibrations, or, as for the SAO 
analysis, bootstrap one's way to find a single-star TF. To do so, we used data 
taken from the 9 POINTS OF LIGHTS experiment, a series of engineering ob- 
servations aimed to monitor the OTA-FGS optical characteristics as function of 
the secondary mirror position. Noticing that all the Hyades transfer scans were 
taken approximately in the same spot of the FGS field of view (FOV), we se- 
lected the single-star TFs on the basis of their resemblance with the 9 POINTS 
OF LIGHT TF taken in the nearest position of the FGS FOV. 

After having inspected all the scans of all the 16 targets (grouped per filter), 
we started our boostrap procedure by defining some scans to be single-star TFs. 
All these scans were co-added to produce an initial single-star TF template; 
then, each single scan was kept or rejected, on the basis of its resemblance to 
the template. Finally, the resulting single-star TF model was constructed by 
co-adding all the accepted scans. The templates obtained for the X-axis and 
the two filters used are shown in Figs. 2a and 2b. 

All the remaining stars, which observed TF had not contributed to the def- 
inition of the single-star TF, were tested for duplicity through the technique 
described above, and one binary was found. The complete results, together 
with other relevant information, are reported in Table 1. 

Based on the experience made with the SAO data analysis, we believe that 
FGS Transfer Function Mode astrometry can give binary star component sep- 
arations with a precision of about 5 mas, and derive magnitude differences to 
about 0.15 mag. After SV and Cycle 1 calibrations are carried out, we should 
be able to improve upon the present situation. 



241 



or 




-0.2 - 



X-axis (arcsec) 

Fig. 2{a) 




-0.1 - 



242 



TABLE 1. Summary of Observations 



Star Name" 


V 


B-V 


Status 


Filter 


Guiding Mode 


H115 


11'!'56 


+ 1':'38 


I* 


Clear'' 


G/ 


H198 


8.46 


0.72 


I 


Clear 


ct/ 


H230 


6.17 


0.46 


S" 


ND^ 


CT 


H246 


6.61 


0.41 


I 


ND 


CT 


H292 


9.11 


0.87 


s 


Clear 


CT 


H307 


7.15 


0.22 


s 


ND 


CT 


H312 


9.99 


1.06 


s 


Clear 


FL/ 


H316 


6.97 


0.44 


s 


ND 


CT 


H379 


7.49 


0.54 


s 


ND 


FL 


H388 


8.12 


0.66 


s 


ND 


FL 


H417 


9.52 


0.93 


s 


Clear 


CT 


H420 


9.03 


0.84 


s 


Clear 


FL 


H429 


5.90 


0.84 


s 


ND 


CT 


H507 


7.78 


0.54 


s 


ND 


FL 


H554 


8.66 


0.74 


s 


Clear 


CT 


H578 


8.51 


0.84 


D*^ 


Clear 


CT 



"From Hansen (1975). 
I = Indeterminate but probably single. 
•^S = Single. 
'^D = Double. 

''ND = F5ND neutral density filter (5 mag); Clear = F583W clear filter (Ag = 5830A 
FWHM = 2340A). 

■'Guidance Modes were FL = Fine Lock, CT = Coarse Track, or G = Gyroscopes. 



243 



REFERENCES 

Bradley, A., Abramowicz-Reed, L., Story, D., Benedict, G., and Jefferys, W. 

1991, PASP, 103, 317. 
Franz, O. G., Wasserman, L. H., Nelan, E., Lattanzi, M. G., Bucciarelli, B., 

and Taff, L. G. 1991, A. J., submitted. 



244 



Restoration of Images Degraded by Telescope Aberrations 



T. Reinheimer, D. Scherti and G. Weigelt 
Max-Planck-Institut fiir Radioastronomie 
Auf dem Hiigel 69 
D-5300 Bonn 1 
Germany 



1. Introduction 

The spherical aberration of the Hubble Space Telescope (HST) causes a point spread 
function (psf) which consists of a central core of about 0.1 arcsec diameter and a halo 
of several arcsec diameter. The core contains only about 12% of the total psf intensity. 
We have performed laboratory simulations of images degraded by telescope aberrations 
and photon noise (10 000 photon events per image). Spherical aberration was produced 
by suitable optics. The aberrated images were used to investigate the dependence of the 
reconstructed image on the applied image restoration method. The image reconstruc- 
tion methods Wiener filtering, Clean, Gerchberg method, Lucy-Richardson method and 
MEM were compared. 

2. Image Restoration Experiments vvfith Laboratory Raw Data 

The laboratory setup for the simulation of HST data is shown in Fig. 1. Spherical 
aberration is produced by using an achromatic telescope lens with the wrong orientation 
in the setup (plain surface on the side of the parallel beam). In front of the telescope 
lens a mask similar to the HST pupil function was inserted . The aberrated images in 
the focal plane of the telescope were recorded with a high-gain image intensifier (gain 
about 10 ) coupled to a CCD camera. The system was able to record individual photon 
events. Fig. 2 shows a diffraction-limited image of the laboratory object. The intensity 
ratios of the 4 stars are 1:0.61:0.53:0.23 . Fig. 3 shows the point spread function of 
the optical setup (spherical aberration). Fig. 4 is an aberrated raw image of the star 
cluster (Fig. 2) recorded with our optical setup. The image is degraded by spherical 
aberration and photon noise. The total number of photon events per image is ~ 10 000, 
the number of photon events in the brightest pixel is ~ 70. Figures 5-9 show the images 
reconstructed from the aberrated raw image (Fig. 4) by Wiener filtering (Helstrom 1967; 
Fig. 5), by the iterative image restoration method Clean (Hogbom 1974; Fig. 6), by the 
iterative Gerchberg method (Gerchberg 1974; Fig. 7), by the Lucy-Richardson method 



245 



(Richardson 1972, Lucy 1974; Fig. 8) and by MEM (MEMSYS-3 package, Gull and 
Skilling 1984; Fig. 9). 

The reconstruction of star 4 has the biggest error since it is faint and close to bright 
stars. A comparison of the reconstructed images, the aberrated raw image and the 
original object shows that most of the restoration methods were quite successful. In 
all reconstructed images all stars are clearly visible, whereas they are not visible in the 
aberrated image (see Reinheimer and Weigelt, 1992 for more quantitative details). The 
conclusion may be different if photon noise is more severe or if other object classes are 
observed. In future experiments we will study other object classes. 



REFERENCES 

Helstrom, C.W., 1967, J. Opt. Soc. Am. 57, 297 

Hogbom, J. A., 1974, Astron. Astrophys. Suppl. 15, 417 

Lucy, L.B., 1974, Astron. J. 79, 745 

Gerchberg, R.W., 1974, Opt. Acta 21, 709 

Reinheimer, T., Weigelt, G., "Deconvolution of Hubble Space Telescope Data: Computer 

Simulations and Laboratory Experiments", Conf. Proc. on "Restoration of HST 

Images and Spectra", 1990, ed. R. Allen (STScI), p. 88 
Reinheimer, T., Weigelt, G., "Restoration of Images Degraded by Telescope Aberrations" 

submitted to Pure and Applied Optics 
Richardson, W.H., 1972, J. Opt. Soc. Am. 62, 55 

Gull, S.F., Skilling, J., 1984, "The maximum entropy method" in Indirect Imaging, ed. 
J. A. Roberts, Cambridge Univ. Press 



laboratory star simulator 



HST simulator 




mercury 

vapor 

lamp 



condensor 

neutral 
density 
filters 



laboratory 
object 



aberration 
glass plate 



image 
intensifier 



<^^ 




<.t^ 





CCD 

camera 



coupling 
lens 



Fig. 1: Optical setup 



246 



Fig. 2: Diffraction-limited object, a 
star cluster. The stars are called 1, 2, 
3, and 4 from top to bottom. 



Fig. 3: Laboratory point spread fun- 
ction (simulated spherical aberration 
and HST FOC f/288 pupil function) 




Fig. 4: Laboratory image of the ob- 
ject degraded by spherical aberration 
and photon noise (~ 10000 photon 
events/frame or ~ 70 photon events in 
the brightest pixel) 



Fig. 5: Image reconstructed from Fi^ 
4 by Wiener filtering 



247 




Fig. 6: Image reconstructed from Fig. 
4 by Clean 



Fig. 7: Image reconstructed from Fig. 
4 by the Gerchberg method (30 iterati- 
ons) 




Fig. 8: Image reconstructed from Fig. 
4 by the Lucy- Richardson method (140 
iterations) 



Fig. 9: Image reconstructed from Fig. 
4 by MEMSYS-3 (140 iterations) 



248 



Coping with the Hubble Space Telescope's PSF: 
Crowded Field Stellar Photometry 

Eliot M. Malumuth 

Computer Sciences Corporation 

James D. Neill AND Donald J. Lindler 

Advanced Computer Concepts 

AND 

Sara R. Heap 

Goddard Space Flight Center 



1 Introduction 



The spherical aberration of the Hubble Space Telescope, HST, presents astronomers with a 
Point Spread Function, PSF, unlike any that they have had to deal with in the past. The PSF has 
a sharp core of approximately O'.'l and broad low surface brightness wings which have rings and 
tendrils that extend to over 2'/0 in diameter. Figure 1, a shaded surface plot of a PSF star taken 
from a Planetary Camera image, illustrates how much higher the surface brightness of the core is 
compared with the wings. Another complication is that the PSF varies with position in the field 
of view of the Wide Field Camera, WFC, and the Planetary Camera, PC. 

While the PSF is tantalizing 
to the astronomer who wishes to do 
photometry of crowded fields because 
of the sharp core, it is disappointing 
because the wings of nearby stars in- 
troduce a variable and unknown back- 
ground. 

2 Observations 

As part of the Science Assess- 
ment Program, PC images of the 30 
Doradus region of the Large Magel- 
lanic Cloud, LMC, were obtained with 
two different filters. Five 300 second 
exposures were taken with the F368M 
(hereafter U) filter and five 100 second 
exposures were taken with the F547M 
(hereafter V) filter. A further com- 
plication was that the V images all 
had saturated pixels in the cores of the 
brightest stars. We have repaired the 
V images as best we could by match- 
ing the unsaturated parts of the core with the same star in the U image. The compact cluster of 
stars at the center of 30 Doradus is known as R136. In these data R136 is located near the bottom 
center of the P6 CCD chip. In this work we only consider the data located on PC chip P6. 

Figure 2 shows the final U image displayed on a logarithmic scale. It is a median of the 5 
individual images. Inspection of the R136 region shows that there are ~ 150 stars in an area about 
the size of the PSF wings. The light between the stars in R136 is due to the overlapping of the 
wings of the PSFs of all of the stars. It is this background which must be correctly accounted for in 
order to use the cores to do accurate photometry. The average sky background has been subtracted 




Figure 1. A shaded surface plot of the PSF star taken from the PC 
image of 30 Doradus. It illustrates how sharp the core is, and how 
much higher surface brightness than the wings it is. 



249 



off of this image. 



3 Method 



In order to do stellar photometry in crowded fields with the HST, we have developed a simple 
approach that uses the known properties of the HST PSF. The following is a brief step by step 
description of this technique. 



1. Prepare a list of stars and 
their x, y positions. We used 
DAOPHOT (Stetson 1987) as modi- 
fied by Holtzman (1990) to find 402 
stars brighter than the local maxima 
in the PSF of the brightest stars. We 
found an additional 259 stars by in- 
specting the image. The positions of 
these additional stars were measured 
by fitting gaussians to the cores in 
both the X and y directions. 

2. Extract images of stars to use 
as PSFs. In the case of 30 Doradus, 
there is one bright, isolated star in the 
field. It is the star designated by the 
letter A, somewhat to the left of center 
in figure 2. The PSF derived from this 
stellar image has an area of 120x120 
pixels or 5'.'16x5'.'16. 




Figure 2. PC image of 30 Doradus. Tliis image is the median of 
five 300 second exposures using tiie F368M filter. 



3. Make an initial guess of the relative flux of each star, Fi. The ratio of the counts in the 
core of each star to the counts in the core of the PSF star is used for the initial guess of the relative 
flux. In practice we used the central 5x5 pixels (0'.'22x0'.'22) to give us the counts in the core. This 
will give an overestimate of the relative flux in the most crowded regions because the central 5x5 
pixel box will contain Ught due to the wings of the neighbouring stars. 

4. Produce a model of the field. We start with an image of the same size as the data image 
(800x800 for one PC chip) that has a data value of zero in each pixel. For each star, i, we register 
precisely the PSF image with the star's position using a bilinear interpolation, scale it by Fi, and 
add it to the model image pixel by pixel. When this has been done for each star we have a model 
of the field which can be compared with the data frame. 

5. Adjust the relative flux scale factor for each star. Once a model image is made the scale 
factors for each star is adjusted using the following equation. 



^x,-|-2 ^y,+2 



F' = F, X 



E._;tiErio(z^ 



Where F- is the new scale factor of the z*'' star, O is the observed image and M is the model image. 

The last two steps are repeated using the new scale factors until the convergence criterion is 
met. For these data we used a convergence criterion of 98% of the star scale factors change by less 



250 



than 3% between iterations. The behavior of the estimates of the scale factors is that stars which 
are isolated reach their final value quickly (1 or 2 iterations), while stars in crowded regions start 
out with an overestimate on the initial guess, drop below the final value after the first iteration, 
and then approach their final value asymptotically. For the crowded inner region of R136 it takes 
about 15 iterations to reach the final values of the relative flux scale factors. 

4 RESULTS 

The final model for the U image is shown in figure 3. This can be directly compared to figure 
2. On casual inspection it appears to be a very good representation of the observed image. A more 
quantitative way to evaluate the results is to produce a residual image. This is done by subtracting 
the final model image from the observed image. Figure 4 is the U residual image. This image 
reveals the problem of not having positional information for the PSF. For example the mismatch 
between the PSF and the stellar image is evident for the bright stars on the right side of the image. 
Otherwise, the residuals are fairly small. 

Another way of looking at the 
results is to compare cross-sectional 
plots of the observed image and the fi- 
nal model image. Figure .5 shows a row 
plot which crosses the stars R136c and 
Melnick 42 (Melnick 1985). The dot- 
ted line shows the data from the PC 
image and the dashed line is from the 
model image. Aside from the excel- 
lent overall agreement, notice how well 
the model matches the background be- 
tween the stars in the crowded region 
of R136. 

One of the advantages to this 
method is that it is a simple matter to 
use many PSFs. In theory, as many 
PSFs as there are stars may be used. 
To illustrate this, we have repeated 
the procedure using 4 PSF stars. The 
PSF stars used are shown in figure 2, 
and are labeled by letters A, B, C and 

D. The original photometry was used to clean stars from the vicinity of stars C and D. Figure 6 is 
the U residual image using the 4 PSF stars. Notice how the residuals of the stars on the right have 
decreased. 

Figure 7 is a comparison of the derived magnitudes using 1 PSF star and 4 PSF stars. In 
addition to the random errors at fainter magnitudes, small systematic differences between the 
photometry done with different PSF stars are evident. The differences for the fainter stars in 
the lower left quadrant, for which the PSF star was the same, are due to changes in the stars in 
neighboring quadrants. 




Figure 3. Final model image of the U image of 30 Doradus. This 
image can be compared with figure 2. 



251 




600 



400 



I Row 219 F368M 



o 
O 



200 



Figure 4. Residual image formed by subtraction figure 3 from figure 
2. The effect of the spacial variation of the PSF is seen as the wors- 
ening residuals in the upper right hand corner of the image. Notice 
the residuals of the rings and tendrils. 




-200 



PC Image 

Model 

Residual 



kAf 



V 



A 



■'^^MVi^^\f\yY*f1^r^'''■~- 



lA'Vww^ 



300 



400 



600 



Column 



Figure 5. A plot of row 219 crossing R136c and Melnick 42 (star D 
in fig. 2) in the PC image of 30 Doradus. The Observed image (fig. 
2) is shown as the dotted line. The model image (fig. 3) is shown 
as the dashed line. The residual is shown as the solid line offset by 
-100. Notice how well the background light in the region of R136 is 
matched. 



1.0 



0.5 



0,0 



'' I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I I j I I I [ I I I I I I I 1 1 1 
■ O Upper Left Quadrant ^ x ^ 

+ Upper Right Quadrant ♦ ^ o , 

'_ O Lower Left Quadrant * X " o 

_ X Lower Right Quadrant , t * * o 

' If "'it "* l" 

* «V«» SfOSM ° ♦• • o 



>- -0.5 - 



Figure 6. ResiduaJ image formed using the model made with 4 
PSF stars. The residuals in the lower left corner are identical to 
those in figure 4 because the same PSF stars (star A) was used for 
this quadrant. The residuals on the right hand side of the image are 
much better than in figure 4. This is especially true for the rings and 
tendrils of those stars. 



REFERENCES. 

Holtzman, J. 1990, PASP, 102, 806 
Melnick, J. 1985, A&A, 153, 235 
Stetson, P. 1987, PASP, 99, 191 



-1.0 




1 1 ■ ' ' ' I ' 







' ' I ' ' ' ' I ' ' ■ ■ I ■ ' ' ■ I ' ' ' ' ' ' 



12 



14 



16 



18 



20 



22 



U (Magnitudes) 



Figure 7. Comparison of the results for 1 PSF star and 4 PSF stars. In addition t 
the random errors at the fainter magnitudes there are small systematic differenc< 
between the photometry done with different PFS. 



252 



SOME ALGORITHMS AND PROCEDURES USEFUL TO ANALYSE 
HST-FOC IMAGES 



C. Barbieri, G. De Marchi, R. Ragazzoni 
Astronomical Observatory of Padova 
Vicolo dell'Osservatorio, 5 
35122 Padova, Italy 



Abstract. Four procedures are briefly described among those we have developed for 
the reduction of HST-FOC frames. 

Emphasis is given to those algorithms we think particularly useful for this kind of 
space-based images. 

1. INTRODUCTION 

Images collected by the Faint Object Camera need some particular procedures in 
order to be properly handled. Actually, most of the problems normally encountered 
analysing ground-based images are different from those arising while treating space- 
based images. 

In this poster we want to point out, with some examples, that a special care is re- 
quired even for simple operations, like recentering, smoothing, background subtraction, 
peak location and so on. 

2. FRAME RECENTERING VIA SHIFT AND ROTATION USING 
AUTOCORRELATION TECHNIQUES 

This section will discuss the problem of comparing and superimposing images taken 
from the ground and from the space. Ground-based images are characterized by ap- 
proximately the same degree of resolution, and the problem of comparing two or more 
of them is relaxed. 

Furthermore, the PSF is undersampled in space-based images while it is oversam- 
pled in ground-based, so the precise matching of HST with ground-based images be- 
comes a difficult task (even ignoring strong colour differences). 

Even the comparison of two FOC frames is not a trivial operation: actually, in 
Coarse Track mode we are never sure to be justified in ignoring field rotation. Due to 
the sharpness of the PSF core, even a shght rotation can destroy the precise alignment 
of stars in a crowded field. 

Precise and robust procedures capable of matching two generic images are therefore 



253 



needed. Recentering of frames via X-Y shifting, using auto correlation techniques, is 
more or less a common method. An extension of this method, allowing also for unknown 
relative rotation of frames, is here briefly described. 

The full procedure is shown in Figure 1 and is here summarized step by step: 

1. Each image is split into two areas: the inner one, with a radius of 1/6 of the image 
size, and the outer one, an annulus with internal and external radii respectively 
1/6 and 1/2 of the image size. The internal size of 1/6 has been somewhat arbi- 
trarily chosen in order to have enough area in the center and at the same time a 
not excessive rotation effect inside it. 

2. On the inner area the usual 2D auto correlation function (ACF) is performed. The 
position of the ACF's main peak gives the relative X-Y shifting; 

3. The outer areas are projected, using polar coordinates, on strips with height 1/3 
of the image size and width n times the size; 

4. Auto Correlation is performed on the couple of strips line-by-line; 

5. The sum of columns is performed, weighting each line by the value of the corre- 
sponding radius. 

G. Finally, \\\v amount of rrlativr rotation (radians) is obtained multiplying by 2 the 
position of the peak in the latter sum and dividing it by the size. 

It should be noted that performing auto correlation on the full image (instead of on 
the inner part only) in order to get the X-Y shift does not produce any improvement on 
the overall accuracy. Moreover, X-Y recentering is to be performed before projection 
and following rotation detection. 

Actually, each rotation around a point different from the centre of the frame trans- 
lates into an additive blurring of the final rotational Auto Correlation Function, i.e. in 
a loss of accuracy. 

The reached precision is of the order of some fraction of pixel, both for X-Y displace- 
ment and rotation. For the latter quantity, this linear error translates into a rotational 
error at the radius distance where the ACF is not negligible. In a typical crowded FOC 
frame 512 ^ 512 this means an error in the estimation of A6 « 5 . 



3. SOME SIMPLE ADAPTIVE FILTERS 

Smoothing frames in order to enhance the signal to noise ratio (SNR) is a common 
and useful operation. 

Working with photon Hniited images means sometimes dealing with abrupt SNR 
changes on the frame itself. In order to retain a SNR level approximately equal over 
the whole frame, some adaptive smoothing must be performed. 

Such adaptive filtering, while retaining the SNR constant on the entire image, leaves 
a varying resolution. In fact, a poor SNR calls for a strong smoothing, i.e. a loss in 
resolving power. We think that such a loss is due to physical and unavoidable reasons, 
and so no real information is lost. 

In Figure 2 some examples of such smoothing are shown. 

Given an estimation of the spatial resolution at a given SNR (for istance the FWHM 
of the typical PSF where SNR is greater than 10) the method can produce a map 



254 



IHII 




outer area 



outer area 



l.UO 




/\ 




0.72 


-A 


^JV 


A/v-A " 


OHI 


\^,/^ 


/ \j s^ 


'vv^ \,^,- 



ACF line by line 



Radius Weighted 



Ae 



Figure 1: Auto Correlation helps to solve the problem of recentering 
two frames shifted and rotated one with respect to the other. In this 
example the centre of a globular cluster taken with HST-FOC and 
convolved with a gaussian shaped beam is compared with a ground 
based frame in order to properly match the two observations. 



255 




Figure 2: a):Raw frame (G2237, the Einstein cross), b): the same 
filtered in an adaptive way, c): the adopted gaussian beam size for 
any point of the frame, d): a normal (space invariant) smoothing, 
for comparison. 



256 



showing the resolution for each point of the frame, in order to estimate the significance 
of faint, photon limited details in the raw frame. 

4. ABOUT THE SUBTRACTION OF UNDERLYING DIFFUSE 
OBJECTS 

Thanks to the sharp core of the HST PSF, a very simple technique can be used to 
subtract an underlying and diffuse object, like a nebulosity in a point-like field. 
The procedure can be summarized step by step as follows: 

1. Place a grid of assigned size on the image; 

2. Evaluate the lowest pixel value in each sub-image defined by the grid; 

3. The diffuse object is described as a smooth approximation to this set of values. 

The procedure is based on the assumption that there should be a prion no reason 
to have, on a sub-image defined by the grid, lower values than those given by the 
underlying object. 

In order to meet such requirements, care must be paid to the size of the gridding 
and to some initial smoothing in order to avoid exceedely low values due to Poisson 
fluctuations rather than to the real background. 

In Figure 3, as an example, the gridding is operated in a circular manner, in order 
to subtract the underllying galaxy in the assumption that its photometric behaviour is, 
at least approximately, only radial. 

From this example one can easily detect a typical drawback of the application to 
HST-FOC frames, i.e. the presence of reseau marks. At these locations the counts are 
lower than in the neighborhood. These extremely dark reseau marks are seen in the 
figure as circular dark rings. On the other hand the position of the reseau marks is well 
known in advance and it would be easy to remove them before applying the procedure. 

5. 2D CENTERING VIA DERIVATIVE TECHNIQUES 

This simple procedure originates from the observation that the core of the PSF is 
undersampled in the FOC frames and some positional capability can be lost. 

In spite of the geometrical stability of the camera (which is very good, anyway) it 
can be useful to get precise positions of point-like sources on the frame as accurately 
as possible, even if this could not be related to an analogous position in the sky (for 
purposes of recentering, subtraction of stars, and so on). 

We have adopted a technique somewhat common in line centering, using derivatives 
of the third order (see, as an example. Figure 4). 

It is well known, in fact, that such an approach takes automatically into account 
any background described by a quadratic polynomial. 

This feature is particulary interesting in the case of crowded fields, where one needs 
an accurate centering of stars embedded in the halos of the other stars. 



257 




Figure 3: a):Raw frame (G2237), b): circular background estima- 
tion (see the text), c): results of the subtraction of the background 
from the raw frame. In the lower right plot a trace along the back- 
ground estimation is shown. Note the dark rings due to the reseau- 
marks. 



258 



First derivative Third derivative 




i\I14 field 



Figure 4: Centering of a star in a crowded field can take advantage 
from the use of the first and third order derivative of the image itself. 



259 



DECONVOLUTION OF AN FOC IMAGE USING A TIM-GENERATED PSF 



P. E. Hodge 

Space Telescope Science Institute 

3700 San Martin Dr 

Baltimore, MD 21218 

USA 



Science Data Analysts at the STScI have already computed a number of PSFs 
using the Telescope Image Modelling (TIM) software of Burrows and Hasan. These 
PSFs are the first of a catalog of PSFs that are to be prepared so that observers may 
deconvolve images taken with the Hubble Space Telescope. In order to get some feeling 
for the usefulness of these initial PSFs for deconvolving images taken with the f/96 relay 
of the Faint Object Camera (FOC), an image of a single star was deconvolved using 
the Lucy- Richardson algorithm with the appropriate TIM PSF. The central brightness 
increased by a factor of eight, but some structure in the wings was accentuated rather 
than suppressed. 

The image selected for deconvolution was a 900-second exposure of the star BPM 
16274 taken through the F210M filter plus a two magnitude neutral-density filter. The 
TIM PSFs are oversampled by a factor of two, so the IRAF blkrep task was run on 
the FOC image to match the pixel scales. The computed PSFs in this first set do 
not include aberrations other than spherical and focus offset. The PSF for F210M is 
one of the polychromatic PSFs, however, so it does include contributions from several 
wavelengths. We can expect substantially better agreement between the computed and 
observed PSFs as we develop a better understanding of the optical characteristics of 
the HST. 

The image was deconvolved using the lucy task in the stsdas playpen package in 
IRAF. Fewer than 30 iterations were required to bring chi-squared below one. The 
parameter values adu=l and noise=0 were used. 

Figures 1 and 2 show the original image and the deconvolved image using a grey 
scale that emphasizes the outer portions of the PSF. The same range of pixel values 
was used for both displays, even though the maximum value of the deconvolved image 
was much higher. Figure 3 shows a radial profile plot of the deconvolved image. The 
profiles of the original (not included here due to lack of space) and deconvolved images 
differ by a factor of eight in scale, and the original is a half pixel larger in radius, but 
otherwise both profiles are virtually identical in form. 



260 



[1] frame. 1.2: x - X[l/1] 




Figure 1. FOC f/96 image of BPM 16274 with F210M filter, scaled 
to show the outer portions of the PSF. 



261 



[2] frame. 2. 4: deconv - DEC0NV[1/1] 




Figure 2. Deconvolved image of BPM 16274, using the same display 
minimum and maximum and grey scale as Fig 1. 



262 



-I 1 1 1 1 r 



-i 1 1 1 1 1 1 1 r 



8000 



6000 



4000 



*•: X 



2000 










'C*' 



* ♦ + + t * 



J I u 



J L 



_I L 







6 
Pixels 



8 



10 



12 



Figure 3. Radial profile plot of the deconvolved image. 



263 



RAPID DECONVOLUTION OF HUBBLE SPACE TELESCOPE 
IMAGES ON THE NRL CONNECTION MACHINE 



Paul Hertz and Michael L. Cobb 

E. O. Hulburt Center for Space Research 

Naval Research Laboratory 

Washington, DC 20375-5000 

USA 



Abstract. We have developed a rapid, highly parallel image space based convolution 
algorithm for use on the NRL 16k processor Connection Machine. This supports an 
image reconstruction program which uses standard iterative algorithms, such as the 
Maximum Entropy Method or Richardson-Lucy Method; thus, when given a constant 
point spread function (PSF) it yields reconstructed images identical to those run on 
serial computers and workstations. Our parallel implementation offers two advantages. 
(1) The highly parallel Connection Machine allows us to use a PSF which varies across 
the field of view, more closely approximating the true HST PSF. Our current imple- 
mentation uses a 512x512 image and 256 PSFs, each of which is a 61x61 image array. 
We can handle up to 16k PSFs with no loss of throughput. (2) Image deconvolution is 
a highly parallel operation so our program runs very rapidly. A single MEM or RLM 
iteration requires less than 3 seconds of clock time and maintains a sustained perfor- 
mance of 1.1 Gflops. HST images can be deconvolved in a few minutes, rather than 
many hours as required on serial machines. This is an advantage when many different 
PSFs, background subtractions, etc., are being considered. 



1. DECONVOLUTION OF HUBBLE SPACE TELESCOPE IMAGES 

The spherical aberration errors associated with HST stiU produce diffraction limited 
information in the final images. The core, or Airy disk, of stellar images is the size of the 
diffraction hmit of the HST mirror but contains only 20% of the photons (Burrows et al. 
1991). The remaining 80% of the photons are distributed in a halo on the arcsecond 
scale size. For bright objects, deconvolution techniques can restore the images to the full 
diffraction limit creating images comparable to an unaberrated optical system (White 
and Allen 1990). Unfortunately the deconvolution techniques will not be able to push 
the faint end of the images to larger limiting magnitudes because the halo becomes lost 
in the noise of the images. 

Most iterative deconvolution methods take a current guess of the true image, con- 
volve with a PSF, and compare with the observed image. A correction term based on 
the residuals is determined and applied to the current guess; the process continues until 



264 



some convergence criteria are met. Common iterative techniques include the Maximum 
Entropy Method (MEM) (Cornwell and Evans 1985) and the Richardson-Lucy Method 
(RLM) (Richardson 1972; Lucy 1974). 

From a deconvolution point of view, HST images have two atypical characteristics. 
First, because of the reimaging optics in the WF/PC imaging system, the point spread 
function is space variant. A space variant PSF is one where the shape of the PSF 
depends on the location in the final image; thus no single PSF can be used to accurately 
characterize the image. The most computationally intense process in iterative image 
reconstriction techniques is the convolution of image and PSF. For space invariant 
PSFs, the convolution of two arrays is the product of their Fourier transforms, and the 
fast Fourier transform (FFT) is an integral part of most deconvolution efforts. FFTs 
can not be used for the convolution of WF/PC images since the PSF is space variant. 

The second atypical characteristic of WF/PC images is that, though PC images 
are sampled at the Nyquist frequency, WFC images are undersampled by a factor of 
two. Implicit in the use of FFTs is the assumption of Nyquist sampled images. If 
Fourier techiniques are used on undersampled images, aliasing becomes a problem and 
frequencies higher than the sampling frequency are aliased into lower spatial frequencies 
creating low frequency artifacts. In order to limit aliasing, the original image resolution 
must be degraded until the image becomes Nyquist sampled. 



2. A PARALLEL SOLUTION 

The Connection Machine is a massively parallel, single-instruction-multiple-data 
(SIMD) computer (HiUis 1987). The NRL CM-2 contains 16k processors, each with 
128 kbyte of memory and access to floating point coprocessors. The geometry of the 
processors is hardware configured as a hypercube, and is software configured to mimic 
the geometry of the problem. Additional hardware includes a data vault of striped disks 
connected by high speed parallel buses, video frame generation capabilities, 14 inch 
removable optical disks. Sun and VAX front-end machines, and a Tl link into the 
University of Maryland internet node. 

Our implemetation of an image space convolution algorithm on the CM-2 addresses 
both the issues of space variable PSFs and undersampled images in a robust, user 
friendly way. Our algorithm works in image space and does not use FFTs, thus the 
aliasing problem is minimized. We assume that there is a scale size over which the PSF 
can be considered space invariant. Each of these isoplanatic patches is assigned its own 
PSF. In the current implementation, isoplanatic patches range in size from 32x32 to 2x2 
pixels with the PSF in these patches being 61x61 pixels. In the case of a space invariant 
PSF, all isoplanatic patches are assigned the same PSF. The convolution subroutine 
is microcoded in CMIS, the CM instruction set, and relies on detailed knowledge of 
the CM geometry and communication hardware. The subroutine is C or FORTRAN 
callable, and the calling program, which executes the iterative image reconstruction 
algorithm, is currently written in C*, a parallel extension of C^"*^. 



3. RESULTS 

The code was benchmarked using both the MEM and RLM iterative techiniques. 
The MEM code was based on FORTRAN code provided by T. Cornwell of NRAO. WFC 
images of Saturn and the LMC open cluster NGC 1850 were provided by J. Westphal 



265 



and the WF/PC Instrument Development Team for testing of the algorithm. The lim- 
iting factor in our deconvolution efforts is a detailed knowledge of the PSF across the 
field of view. We have used PSF modeling software provided by P. Miller of Hugh- 
es Danbury. The TIM code developed at STScl is more accurate, but can not currently 
calculate the 256 PSFs required in a reasonable amount of time. 

A total of 8 runs were made with the CM deconvolution/reconstruction package. 
These runs include all combinations of two images (Saturn and NGC 1850), two PSFs 
(observed and modeled), and two iterative techniques (MEM and RLM). The observed 
PSF was obtained by chpping an isolated stellar image from near the center of a sparse 
WFC image. The clipped PSF is then replicated 256 times to create our observed 
PSF. The modeled PSF consists of 256 calculated PSFs evenly spaced throughout the 
512 x512 image. 

In the table we indicate the final values of x , which is calculated from the difference 
between the raw image and the deconvolved image convolved with the PSF, as well as 
the number of iterations and the run time for the calculation. I/O takes another 40- 
90 seconds depending on whether the data is stored on the frontend disk of the CM 
datavault and on whether video or graphics output is desired. For comparison purposes, 
comparable runs on serial computers would require between 1 and 16 hours. 

Note that the currently modeled space variant PSF gives results comparable to the 
observed space invariant PSF, but not significantly better. This is an indication of 
the lack of knowledge of the space variant properties of the PSF at the < 10% level. 
As understanding of the HST PSF improves, the results from the modeled PSF will 
be superior to those from the observed PSF. At that time, algorithms using space 
variant PSFs, such as the one described here, will yield results superior to those using 
a constant PSF. 



Parallel Image Reconstruction Test Runs 


Target 


PSF 


Method 


x' 


N^ter 


Run Time 


NGC 1850 


observed 


MEM 


1.391 


30 


89 sec 


NGC 1850 


observed 


RLM 


1.229 


20 




NGC 1850 


modeled 


MEM 


1.533 


30 




NGC 1850 


modeled 


RLM 


1.358 


20 




Saturn 


observed 


MEM 


1.009 


20 


90 sec 


Saturn 


observed 


RLM 


0.865 


20 




Saturn 


modeled 


MEM 


1.000 


20 




Saturn 


modeled 


RLM 


0.860 


20 





REFERENCES 

Burrows, C. J., et al. 1991, Ap. J. (Letters), 369, L21. 
Cornwell, T. J., and Evans, K. F. 1985, Astr. Ap., 143, 77. 
HiUis, W. D. 1987, Set. Am., 256, 108. 
Lucy, L. B. 1974, A. J., 79, 745. 
Richardson, W. H. 1972, J. Opt. Soc. Am., 62, 55. 

White, R. L., and AUen, R. J. 1990, The Restoration of HST Im,ages and 5pecira (STScI: 
Baltimore). 



266 



ON ORBIT MEASUREMENT OF HST BAFFLE REJECTION CAPABILITY 

by Pierre Y. Bely, Doris Daoii and Olivia Lupie 

Space Telescope Science Institute 

3700 San Martin Drive 

Baltimore, MD2r218 



>27 Mg 




Sourc* angle »27 d*gr««t 
Scitllt b» billlt IMn fcy primary minor 




IS to 77 aag 



Saum anglt 15 lo 27 oagrMi 

Scatlar bf oular batlla than by 

aacondary mirror 



1. HST BAFFLE DESIGN 

HST is extremely well baffled against the effect of 
off-axis bright sources such as the sun, moon and bright 
earth. Pointing restrictions and the aperture door fully 
protect against any effect from the sun. Light from the 
moon and bright earth is allowed to enter HST's tube, 
but baffles prevent direct illumination of the focal plane. 
Light can reach the focal plane only after deflection by 
several baffles or via scatter due to mirror dust. 

The mechanisms producing straylight in the focal 
plane fall into three regimes. 

At angles, larger than 27 degrees, light only reaches 
the focal plane after bouncing several times between the 
outer baffles or when scattered by the primary mirror 
dust. The effect is essentially proportional to the dust 
coverage on the primary mirror. 

For the middle angles, 15 to 27 degrees, light can 
reach the focal plane after bouncing from the rear of 
the outer baffle and secondary mirror baffle and subse- 
quent reflection by the secondary mirror. In this regime. 
the focal illumination is essentially independent of the 
mirror dust. 

For smaller angles, 15 degrees and below, light strike.* 
the primary mirror, is scattered by dust, and reaches 
the focal plane after reflection by the secondary mir- 
ror. Light scattered or diffracted by other surfaces (e.g 
secondary mirror spider) also contributes to focal plaiif 
straylight. 

The pre-launch determination of attenuation of off-axis light sources was made by Perkin Elmer and the 
Marshall Space Flight Center using computer modelling with the .APART jiackage and laboratory measurement 
of mirror dust scatter. The predicted attenuation factor is shown ni Figure 2. The APART detailed model 
was not run for angles smaller than 15 degrees. In this domain, light from the off-axis source hits the primary 
mirror directly and the resulting scatter by dust on the mirror becomes the dominant source of straylight. We 
have determined the attenuation at angles smaller than \y> degrees by using a simplified analytic model for the 
mirror scatter and extrapolating the APART model for the other scattering sources. 




Seurca angia < 15 Mgroat 
Oiract aoanar by pttmaiy ailn 



Figure 1 Scattering Regimes 



267 




Figure 2 Predicted BafQe attenuation factor. Tlie 
curves are for 0% (dotted), 2% (solid) and 5% (dashed) 
dust coverage on the primary mirror. The predicted pre- 
launch dust coverage was estimated at about 2% . 



;.0 -0 23 :o 40 50 60 70 80 90 ICO 
Ot(-o«is onoi* (♦rom ooint source) 



2. ON-ORBIT MEASUREMENTS OF STRAYLIGHT DUE TO OFF-AXIS SOURCES 



For angles less than 30 degrees the attenuation of the baffling system was measured on orbit using the 
moon as a source and the Wide Field Camera as an area photometer. The test consisted of measuring the focal 
plane illumination as a function of wavelength (F284W(VU\'). F336W(UV), F569W(V) and F675W(R)) at 4, 8 
20 and 30 degrees from the full moon. The sky background for the faintest exposure levels (20 and 30 degrees) 
•was measured at the subsequent new moon. 

The results are summarized in Figure 3. 

The on orbit data essentially confirms the validity 
of the detailed model in the 15 to 30 degree domain. 
At 30 degrees, strayiight from the moon is negligible 
compared to the zodiacal light level, as it was required 
by HST specifications. 

The 8 degree data confirms the prediction made 
with the simplified model but the 4 degree measurement 
is about 3 times brighter. This is likely explained by an 
underestimation of the complex scattering processes b\ 
surfaces other than the mirror (baffle, spider etc..) at 
very low angles. 

These results obtained at low angles where dust 
on the primary mirror is a primary source of strayiight 
suggest that the amount of dust on the primary mirroi 
is not substantiallv different from pre-launch estimate- 
(2%). 




3. CONCLUSION 



Off axis anaie 

Figure 3 Illumination of the focal plane by 
the full moon as a function of the off-axis angle. 
The on-orbil measured data is shown by point 
symbols for the various wavelengths and is to be 
compared to the predicted level shown as lines 
(solidiV, dottediVUV. dashed:UV, dot-dashed.R). 
The UV data (F284W) is affected by red leak in 
the WFPC and should not be relied upon. 



In conclusion, the results of this test indicate that tiie design requirement concerning strayiight from the 
moon has been met. The test confirms the validity of tiie model in tiie \o to 30 degree range, and hence suggests 
that the design requirement for the bright earth ha.s ai.'.o bet^n satisfied (strayiight less than the zodiacal light 
at 70 degrees from the bright earth limb). However, we intend to confirm the level of strayiight at large angles 
by measuring the background in selected WFPC frame-, taken over the bright earth. 



268 



APPENDIX 

Scheduling of Science Observations and 
Subsequent Data Processing 



TRANSFORMATION: 

THE LINK BETWEEN THE PROPOSAL 

AND THE 

HUBBLE SPACE TELESCOPE DATABASE 



ML. McCollough, H.H. Lanning, and K.E. Reinhard 
Computer Sciences Corporation/Space Telescope Science Institute 

Overview 

In order for a scientific program, specified in a proposal, to be executed by HST the information in the proposal 
must be translated into a set of parameters which can be interpreted and used by the Science Planning and 
Scheduhng System (SPSS), Science Commanding System (SCS), Observation Support System (OSS), and Post 
Observation Data Processing System (PODPS). The conversion of the proposal is performed by the 
"Transformation" software. Transformation is a rule based body of software, written in LISP, designed to 
convert the proposal into a series of relations which can be loaded into the Proposal Management Database 
(PMDB). In addition, Transformation provides products which are used by Science Planning Interactive 
Knowledge Environment (SPIKE) to do long term science planning. Figure 1 shows the flow of information 
from the proposal through Transformation into SPSS: 



Figure 1. This diagram 
shows how information 
from the proposal flows 
through Transformation 
to the various operational 
systems. 













PROPOSAL 














VALIDATION 














TRANSFORMATION 














SPIKE 




OSS • — 


SPSS 




1 


PODPS y — 
















(COMMANDING) 
SMS 










1 



The Proposal and Validation 

Observers initially enter their observing projects into the system through the proposal. An example of a 
proposal is shown in Figure 2. The format and outline of how to create a proposal are contained in the "Hubble 
Space Telescope Proposal Instructions". The major points of information from the proposal are the following: 

(A) Target Information : All the information necessary to observe the target of interest must be given 
(position, positional uncertainty, magnitude, etc.). It is from this information that the pointing of 
the spacecraft and the guide stars used are determined. 

(B) Exposures : These are the basic building blocks of the proposal and represent the observations which 
will be performed by the spacecraft. 

(1) Instrument: The scientific instrument used {WFPC. FOC, FOS. GHRS. HSP. or FGS). 

(2) Mode of Operation: The way in which the instrument is used (IMAGE, ACCUM, RAPID, 

etc.). 

(a) Optional Parameters: Adjustments to various instrument parameters for each mode of 
operation. 



270 



(3) Number of Observations: A single line can result in a multiple number of observations. 

(4) Exposure Time: This is the length of time that the instrument will collect photons. This can 

critically determine how an observation is performed and if the observation is possible. 

(5) Special Requirements: These determine how and when the observations are performed. 

(a) Structure: The order in which exposures are executed relative to one another is determined 

iSEQ, GROUP, etc.). 

(b) Timing: When exposures occur relative to each other and relative to an absolute time {AT, 

AFTER, etc.). 

(c) Real Time Contacts: The use of real time contacts (TDRSS. Tracking and Data Relay 

Satellite System) with the spacecraft are determined {INT ACQ, RT ANALYSIS, etc.). 

(6) Logsheet Comments: It is through exposure level comments that special scheduling and 

commanding requirements of exposures can be noted (not always completely describable by the 
special requirements). 

(C) Proposal Abstract and Description : In addition to the logsheet comments the proposal abstract and 
description relay much of the intent of the proposal to the people doing the scheduling and 
commanding of the spacecraft. 



EXPOSURE LOGSHEET 




















Id 
Pa 


251(P) 

qe: of 


1 


2 


3 


4 


5 


6 


? 


8 


9 


10 


11 


12 


13 


14 


15 


Ln 
Nm 


Seq 
Nam 


Target 
Name 


Instr 
Conf iq 


C^r. 
Mode 


Aper 
orFOV 


Spectral 
Element 


Centrl 
Waveln 


Optional 
Parameters 


Num 
Exp 


Time 


S/N 

Rel.T 
Ime 


Fix 
Ref 


Pr 


Special 
Requirements 


1 NGC224 


FOS/BL 


AC9 


4.3 


MIRROR 






1 


lOOS 


1 




INT ACQ FOR 2 


2 NGC224 


FOS/BL 


AC CUM 


0.5-PAI 


G270H 


2700 


STEP- 

PATT-STAR- 

SKY-BKG 


3 


200S 


1 






3 NGC224 


FOS/BL 


ACCUM 


0.5 


G190H 


1900 


STEP-PATT-DEF 


1 


lOOS 


1 






4 NGC224 


FOS/BL 


ACCUM 


0.25-PA 


GI30H 


1300 


STEP-TIME-1. 5 


1 


300S 


1 






5 NGC224 


FOS/BL 


ACCUM 


0.3 


G190H 


1900 


POLSCAN-SB 


1 


80S 


1 






6 NGC224 


FOS/RD 


ACQ/FIR 
rWARE 


4.3 


MIRROR 




MAP-BOTH, 

BRIGHT-4500.0 

FAINT-300.0, 

SKy-20 .0 


1 


300S 


1 




ONBOARD ACQ FOR 
7 


7 NGC224 


FOS/RD 


ACQ/PEA 
K 


0.3 


G570H 


5700 


TYPE-UP, 
SEARCH- 
SIZE-3, SCAN- 
STEP -0.2 


1 


15S 


1 




ONBOARD ACQ FOR 
8 


8 NGC224 


FOS/RD 


ACCUM 


0.3 


G570H 


5700 




1 


300S 


1 






9 NGC224 


FOS/RD 


ACCUM 


0.3 


.G.S.T.OH.... 


5700 


F.0.LS.CAN-.4A 


2 


lOOS 


1 






10 NGC224 


FOS/RD 


RAPID 


0.3 


G570H 


5700 


SUB-STEP-2, 
COMB-NO, READ- 

TIME-2.5 


1 


90S 


1 






11 HZHER 


FOS/RD 


ACQ/BIN 
ARY 


4.3 


MIRROR 




BRIGHT-4120.0 

FAINT-25.0,NT 

HSTAR-3 


1 


300S 


1 




ONBOARD ACQ FOR 
12 


12 HZHER 


FOS/RD 


ACQ/PEA 
K 


2.0-BAR 


G7e0H 




TYPE-DEF 


1 


20S 


1 




ONBOARD ACQ FOR 
13 


13 HZHER 


FOS/RD 


PERIOD 


2.0-BAR 


G7eOH 


7B00 


BINS-6, SUB- 

STEP- 

2,CYCLE- 

TIME-500,DATA 

-RATIO-5.0 


1 


1500 
S 


1 




AT21-AUG- 
89:13:08 


15 HZHER 


FOS/RD 


IMAGE 


4.3 


G7eOH 




Y-SIZE-18,Y- 

SPACE-18,X-5, 

Y— 3 


1 


600S 


1 







Figure 2. The table above is an example of a typical exposure logsheet. 
The exposure logsheet is how the observer expresses what observation needs 
to be performed. It is from this information that Transformation will create 
the observing structure and populate the PMDB. 



The proposal is submitted into the system through the Proposal Entry Processor (PEP) System. It is while the 
proposal is in PEP that exposures and defined sequences are expanded. Also, linkages between exposures are 
determined (both in ordering and timing of observations). Before the proposal reaches Transformation, it must be 
processed by "Validation" in the PEP system. Validation is software that checks the proposal syntax and 
populates the Internal Database (IDE). It is from the IDE that Transformation gets the files, containing the 
proposal information, with which it will populate the PMDE. 



271 



Observing Structure 

It is necessary for Transformation to create the observing structure which will be used by SPSS. The hierarchy 
created by Transformation is (from the smallest to largest structure): 



I. 


Exposure (Ex) 


II. 


Alignment (Al) 


III. 


Obset (Ob) 


IV. 


Scheduling Unit (SU) 



(I) Exposure : 

This is the basic building block of a proposal. Normally there is a one to one correspondence between 
an exposure and a logsheet entry on the proposal. Information on the observation such as. Scientific 
Instrument (SI) used, mode of operation, spectral element used, and aperture used are determined at this 
level. It is at this level most of the information necessary to command the Sis is contained. Also 
contained at this level is the information that PODPS finds necessary to do the post-processing of the 
observations. 



(II) Alignment: 

Exposures are grouped into alignments. An alignment deals primarily with the pointing of the VI axis 
of the spacecraft. Target position, roll of the spacecraft, and the timing of the observation which are 
used for scheduling are determined at this level. Also, the operational states of the detectors are fixed at 
the alignment level. 

(III) Ql?sct: 

Alignments are in turn grouped into larger structures called Observation sets (Obsets). Obsets are 
groups of alignments which use the same type of pointing control. In particular, groups of alignments 
which can use the same guide stars are often grouped together into the same Obsets. 

(IV) Schednlinf Unit: 

Finally, Obsets are built into large units called Scheduling Units (SUs). SUs are sets of Obsets which 
can be scheduled all at one time. SUs are the basic units which are used to build calendars in SPSS. It 
is between SUs that time Unkages are done. 



Transformation takes exposures created from the proposal and orders and merges them to form the observing 
structure. Once the order has been determined by Transformation, adjacent exposures are merged into alignments 
by a set of merging rules. In turn, Transformation will merge adjacent alignments into Obsets, and Obsets into 
SUs. The merging rules consist of reasons to merge and reasons not to merge. If there is any reason not to 
merge (no matter how many reasons there are to merge) the exposure (alignment or Obset) is not merged. Also, 
a lack of a reason to merge or not to merge is treated as a reason not to merge. For the test proposal given above 
a summary of the structure (and timing) determined by Transformation is shown in Figure 3. 

Alignment Times 

Another task of Transformation is to calculate the time it takes to perform the observation. This includes not 
only the exposure time but also all of the overhead necessary to operate the SI. The basic alignment time 
calculations for an observation are given by the algorithm shown in Figure 4. 



272 



TRANSFORMATION VERSION 12.0 
GENERATED 6-26-1991 14:34:46 
PROPOSAL 251 VERSION P 
TRANSFORMED USING FULL-TRANS 



PEPSI SEQU OBSET 


ALIGN 


EXP 


EXP 


AL SU 


ID PRIM SU 


TOLERANCE 


DELTA 


EXPOSURE LINE ID 


ID 


ID 


TIME 


TIME 








NUM 
















1.0000000 


01 


01 


01 


492 


492 0025101 002501 


000:00:00:00 


000:00:00:00 


INT-ACQ-DEC 


01 


02 


01 




960 


0025101 


000:00:00:00 


000:00:00:00 


INT-ACQ- 


01 


03 


01 




510 


0025101 


000:00:00:00 


000:00:00:00 


UPLINK 


















2.0000000#001 


01 


04 


01 


464 


464 


0025101 


000:00:00:00 


000:00:00:00 


2.0000000*002 


01 


05 


01 


407 


407 


0025101 


000:00:00:00 


000:00:00:00 


2.0000000*003 


01 


06 


01 


407 


407 


0025101 


000:00:00:00 


000:00:00:00 


3.0000000 


01 


07 


01 


454 


454 


0025101 


000:00:00:00 


000:00:00:00 


4.0000000 


01 


06 


01 


689 


689 


0025101 


000:00:00:00 


000:00:00:00 


5.0000000 


01 


09 


01 


1078 


1077 


0025101 


000:00:00:00 


000:00:00:00 


HOME 


01 


QA 


01 




292 


0025101 


000:00:00:00 


000:00:00:00 


6.0000000C 


02 


01 


01 


1041 


1040 0025102 0025102 


000:00:00:00 


000:00:00:00 


6.0000000F 


02 


02 


01 


1025 


1024 


0025102 


000:00:00:00 


000:00:00:00 


SETUP 


02 


03 


01 




948 


0025102 


000:00:00:00 


000:00:00:00 


7.0000000 


02 


03 


02 


948 


948 


0025102 


000:00:00:00 


000:00:00:00 


8.0000000 


02 


04 


01 


535 


534 


0025102 


000:00:00:00 


000:00:00:00 


9.0000000*001 


02 


05 


01 


568 


567 


0025102 


000:00:00:00 


000:00:00:00 


9.0000000*002 


02 


06 


01 


594 


593 


0025102 


000:00:00:00 


000:00:00:00 


10.0000000 


02 


07 


01 


276 


276 


0025102 


000:00:00:00 


000:00:00:00 


HOME 


02 


08 


01 




234 


0025102 


000:00:00:00 


000:00:00:00 


11.0000000 


03 


01 


01 


862 


862 0025103 0025103 


000:00:00:00 


000:00:00:00 


SETUP 


03 


02 


01 




1235 


0025103 


000:00:00:00 


000:00:00:00 


12.0000000 


03 


02 


02 


1235 


1235 


0025103 


000:00:00:00 


000:00:00:00 


HOME 


03 


03 


01 




151 


0025103 


000:00:00:00 


000:00:00:00 


13.0000000 


04 


01 


01 


1957 


1957 


0025103 


000:00:00:00 


000:00:00:00 


15.0000000 


04 


02 


01 


999 


4342 


0025103 


000:00:00:00 


000:00:00:00 


HOME 


04 


03 


01 







0025103 


000:00:00:00 


000:00:00:00 



Figure 3. The example above (Summary File) shows the structure and timing which 
resulted from the logsheet shown in figure 2. 



AL_TIME = 


AL_BEGIN + S ( EXP_TIME ) + AL_END 


AL_BEGIN = 


Alignment specific overheads which occur at the beginning of the 
alignment. 


EXP_TIME = 


Overheads and exposure time to complete a single exposure of the 
alignment (this quantity is summed over all the exposures in the 
alignment). 


AL_END = 


Alignment specific overhead which occurs at the end of the 
alignment. 


EXP_TIME = 


PRE_OVERHEAD + EXPTIME + POST_OVERHEAD 


PRE.OVERHEAD = 


Exposure level overheads necessary to prepare the SI for the 
observation and command the SI to perform the observation. 


EXPTIME = 


Time to perform the exposure and overheads incurred while taking 
the exposure. 


POST_OVERHEAD = 


Exposure level overheads necessary to read out the detector and 
return the detector to a state necessary to perform the next 
observation. 



Figure 4. Above is the basic algorithm used to calculate alignment times. 



273 



Populating the PMDB 

The final product from Transformalion which is used to load the PMDB is the assignment file. The assignment 
file consists of a set of relations (see list below) which describe the proposal, i.e. the way in which the 
observations will be done and how they will be scheduled. This file is in essence an IQL (Interactive Query 
Language) file [in the near future to be converted into an SQL (Standard Query Language) file] which is directly 
loaded into the PMDB. An example of the PMDB values is shown in Figure 5. 



Relation 


System 


Description 






U 


se 


d By 




Exposure 


Level : 










QEXPOSURE 




1, 


2 


3 


Exposure Level Information 


QELOGSHEET 






2 




Logsheet Information 


QESIPARM 






2 




SI Parameter Information 


QECOMMENTS 




1 






Exposure Level Comments 


QGACTINST 




1, 


2 




SPSS-SCS Interface 


Alignment 


Level : 










QALIGNMENT 




1 






Alignment Level Information 


QAPOSITION 




1 






Pointing Information 


QASI_STATES 




1, 


2 




Detector State Information 


QACOMMENTS 




1, 




4 


Alignment Level Comments 


Obset Level: 










QBS_OBSET 




1 






Obset Level Information 


QBWINDOWS 




1 






Scheduling Time Windows 


SO Level: 












QSCHEDOLING 




1 






SU Level Linkage Information* 


QSBRANCHING 




1 






SU Level Linkage Information* 


Target Le 


vel: 










QTARGETS 




1 






Target Information 


QTSYNONYMS 




1 






Target Related Information 


QTCOMMENTS 




1 






Target Level Comments 


Proposal 


Level: 










QPDESCRIP 




1 






Proposal Information 


QPABSTRACT 








3 


Proposal Information 


QPKEYWORDS 




1 






Proposal Information 


QPCOPROPSER 








3 


Proposer Information 


QPERSONNEL 




1 






Proposer Information 


QPPCOMMENTS 




1 






Proposal Level Comments 



1- SPSS 2- SCS 3- PODPS 4- OSS 
* In the near future the time linkage between SUs currently contained in these relations 
will be placed in the new relations QSLINK_INFO and QSLINK_SPEC. 

Test Scheduling and SMS Generation 

After the assignment file has been loaded into the PMDB, the "Proposal Preparation Group" of SPSS personnel 
must complete the analysis and processing of the proposal for scheduling and execution. This consists of three 
things: 

(1) Operational Problem Report (OPR) Fixups : Transformation has known problems and deficiencies. 
These issues must be addressed and fixed for proposals which are loaded into the PMDB. Also, as 
Transformation evolves and changes the new products must be examined for potential problems. 

(2) Test Scheduling : Once all fixups are completed on a proposal it is then test scheduled. The proposal 
is scheduled by itself on a test calendar of one week (the week it is supposed to first be observed if 
possible). This will allow scheduling concerns and problems to be addressed in advance of scheduling 
on the Flight Calendar. 

(3) Test Science Mission Specification (SMS) Generation : Once the test calendar has been made, it is 
then run through the SPSS software to produce a test SMS. The SMS is a detailed listing (in time) of 
the maneuvering and commanding of the spacecraft. It is during this stage that problems in the 
proposal which relate to the commanding of the Sis will be uncovered. If problems are found then 
fixes to remedy these problems are made to the proposal. 



274 



After all three of the above steps have been sucessfuUy completed the proposal is ready for scheduling for flight 
operations by SPSS. 

Figure 5. Below is an example of a portion of an assignment file that is used to load the PMDB 



append QBS_OBSET ( 
proposaljd = "00251 ", 
obsetjd = "01 ", 
saa_flag = "Y", 
target_opp = "N", 
parallel_can = "N", 
priority = 80, 
repeated = "N", 
order_spec = "N", 
recorder = "N", 
ac_ephemeris = "N", 
ac_clock = "N", 
ground_coord = "N", 
critic_type = "N", 
cntic_flag = "N", 
interleave = "N", 
interrupt = "Y", 
realtime_flg = "Y", 
service_type = "BOTH", 
linkjype = "BOTH", 
eng_32kbps = "N", 
facility = "SCI", 
reference = "N", 
software = "N", 
pcs_mode = "FGS", 
pcs_max_dur = 3600, 



append QALIGNMENT ( 
proposaljd = "00251", 
obsetjd = "01 ", 
alignmentjd = "04", 
alignjype = "DC", 
low_priority = "N", 
asfromefry = "N", 
excuteslew = "N", 
shadow = "N", 
interrupt = "N", 
interleaver = "N", 
calc_sam = "Y", 
time_require = 459, 
tape_recordr = "Y", 
primjarget = "251_1", 
targetjype = "P", 
calibrjype = "N", 
saa_avoid = "05", 
saaovr = "N", 
occ_ovr = "N", 
recovery_ovr = "N", 
targetacqsi = "03", 
camera_ast = "NONE", 
pointing_mde = "F", 
scan type = "N", 
scan coord = "C", 



scenario_acq = "COARSE2",fhstpar = "Y", 



saa_model = "02", 
saa_ovr = "N", 
recovery_ovr = "N", 
seq_target = "Y", 
fov_required = "N", 
brit_object = "Y", 
parall targ = "N", 
readyJor_gs = "N", 
fov_status = "N", 
readyjiag = "N", 
prev_acq_fl = "N", 
reacq_type = "N", 
reacq_sn = "COARSE1 
reacq_co_ovr = "N", 
acq_co_ovr = "N", 
reacq_tm_nsl = 320, 
max_slewint = 0, 
maxjnt_dur = 1 2600, 
min_slew = 20.0, 
fhstroll = "D", 
fhstfull = "D", 
fhstroin = "0", 
fhstroll2 = "0", 
fhstfuin = "0", 
pcsgap = "Y", 
reconf time = 60, 
slew_setjim = 30, 
si_motionJI = "Y", 
max_sep_dur = 86400, 
min_sep_dur = 0, 
version num = "01" 



dark_er_occ = 6.0, 
brit_er_occ = 1 5.0, 
si_parallel = "Y", 
gsss_request = "N", 
minsepdur = 0, 
max_sep_dur = 12600, 
saajlag = "Y", 
version_num = "01" 
append QAPOSITION ( 
proposaljd = "00251 ", 
obsetjd = "01", 
alignmentjd = "01", 
initial_pos = "B", 
alignjype = "AB", 
target_ref = "P", 
si_used = "FOS", 
def_aperjlg = "Y", 
coordjyp = "SICS", 
coord Jd = "YBL4_3", 
xoffset_aper = 0, 
yoffset_aper = 0, 
orientjype = "NM", 
ambiguity = "Y", 
aper1ure_viw = "ALL", 
target_view = "PNT", 
version num = "01" 



append OGACTINST( 
proposaljd •= "00251", 
obsetjd ■ "01", alignment_id = 
exposuretd - "01", 
activity_id = "SCIENCE", 
instrname » "CMAIN". 
instrver ■■ "01", 





APOSmON ( 


append QESIPARM ( 




Jd- "00251", 


proposal_id . -00251 -, 




- "01", allgnmem_id - "05" 


obset_id - -or. alignment_id - "05" 




)S."B-. 


exposure_id --01-. 




M . "NL". 


si_par_name - -INFTFLAG-. 




3l - "P-. 


si_par_ value - -NOCHANGE", 




. "FOS". 


versionnum - "Or 




r llg - -Y", 






ip . "SICS", 


append QESIPARM ( 




- "YBL0_5PRB-. 


proposaljd- -00251". 




aper - 0. 


obs6l_id - -or. alignm6nt_id - -05" 




aper - 0. 


exposure_id - -01". 




rpe - "NL". 


si_par_name - -PATTERN". 




Y - "Y". 


si_par_ value - "STAR-SKY-BKG". 




viw - "PNT". 


versionnum - "01" 




ew - "PNT", 






num - "01" 


append QESIPARM ( 
proposaljd - "00251 ", 




ASLSTATES ( 


obset_id - "01", alignment_id - "05" 




Ijd. "00251", 


exposure_id -"Or. 




- "01", alignment id- "05" 


sij>ar_name - "SCIHEADER", 




-OS". 


si_par_ value - "YES", 




- -BLUE". 


version_num - "01" 




te-"HVONB-, 






a - "HVONB", 


append QESIPARM ( 




num - "01" 


proposaljd - "00251". 

obset_id - "01". alignment_id - "05" 




*SI_STATES ( 


exposuro_id- "01". 




Ijd- -00251". 


si_par_nam6 - "TARGTYPE". 




."01", alignment id - "05" 


si_par_valu6 - "STAR", 




-OS", 


ver3ion_num - "01" 




to - "READY-, 


append QESIPARM ( 




e - -READY-, 


proposaljd - "00251". 




num - -or 


obset_id - "01". alignm8nt_id - "05" 
exposure_id- "01", 




EXPOSURE ( 


si_par_name - "COMRATE", 




Ijd --00251-. 


si_par_value - "4". 




- -or, alignment_id - "05" 


version_num -"01" 




id- -or. 






;cp. 


append QECOMMENTS ( 




Jg - "N-. 


proposaljd. -00251". 




lib . "N". 


obset_id - "01", 




ord."N", 


align_id - "05-, 




ume - 137. 


exposure_id = -or. 




itic - "N". 


page_num -"or. 




Id - -N-. 


comment type - "PC". 




--Y-, 


version^num - "01" 




type - -A", 






time = 200, 


append QALIGNMENT ( 




1--N-. 


proposal_id - "00251", 




ord - -C-. 


obset_id - "01", alignment_id - "06 




l--251_r. 


alignjype - "DC". 




igt - "N". 


low_pnority - "N". 




itjg . "N". 


astrometry - "N". 




- -FOS". 


excuie_slew - "N". 




p - "SICS", 


shadow - "N-. 




1 - "YBL0_5PHB-, 


interrupt - "N". 




iper - 0, 


interleaver - "N-, 




ipermO, 


calc_sam - -Y-, 




y--N-. 


time^require - 402. 




to - -0-, 


tape recordr -"Y". 




te - -0-, 


primjarget - -251_1-, 




SPECTROSCOPY-. 


targettype - -p-. 




lor - 4, 


calibr type • -N", 




i_pkt = 4, 


saa_avoid - "05-. 




d - -YPCHY-, 


saaovr - -N-, 




_ang =0,0. 


occovr - "N-, 




Jen =0 0. 


recovery_ovr » "N", 




-0. 


target_acqsi = "03", 




0. 


camera.ast - "NONE-. 
pointingLmde - 'F', 
scan Jype = "N-. 


samppi 


= 0. 
1 = 0. 


los_delector = "BLUE". 


scan_coord = "C". 


IchnI . 226, 


fhstpar - -Y-, 


nchnte ■ 0, 


dar1(_er_occ - 6 0, 


overscan - 5, 


bnt_©r_occ = 15 0, 


aperjd - "A-2", 


si_parallel = -Y-. 


polar_ic 


-■C-. 


gsss_request - -N" 



275 



PROPOSAL PREPARATION BY SPSS FOR SCHEDULING 

ON THE 

HUBBLE SPACE TELESCOPE 

K.E. REINHARD, H.H. LANNING, and W.M. WORKMAN, III 

Computer Sciences Corporation I Space Telescope Science Institute 

Overview 

Preparation of a proposal for execution on board the Hubble Space Telescope encompasses a great deal of 
manually intensive work by Science Planning and Scheduling System (SPSS) personnel. The preparation task 
includes tracking of the work status, detailed analysis of the structure and contents of the proposal, modification 
of the database values as required for proper execution onboard, generation of scheduling windows, test 
scheduling, and incorporation of the commanding and proposal changes necessary for execution. This 
preparation process is shown in the flow diagram on the next page. Throughout the process, the products are 
analyzed for potential errors in order to deliver a schedulable proposal. The test products are reviewed internally, 
and upon approval, delivered to flight preparation personnel within SPSS for guide star processing and final 
preparation for flight. 

PMDB Load 

Upon receipt of a proposal Delivery Notice from the Science Planning Branch (SPB), the Assignment File and 
Summary File which contain the proposal structure information, science and spacecraft activities, etc. are 
transferred to the Science Operations Ground System (SOGS) by Science Planning & Scheduling System (SPSS) 
personnel. The Assignment File, an IQL/SQL language file, is then loaded into the Proposal Management 
Database (PMDB). As part of the loading process. Scheduling Windows are set for the planned scheduUng time 
frame, link set information is established, and a number of standardized database values are input based upon 
current spacecraft/scheduling requirements. 

Proposal Structure Review 

Initial analysis of the proposal structure focuses on the evaluation of the correctness of the activities desired by 
the proposer. Such activities include looking for the proper arrangement of Interactive Target Acquisitions, the 
nature of Interruptions allowed, the combination of exposures in a given Observation Set (Obset), etc. If the 
structure is not consistent with the requirements of the proposal, or major Transformation problems are 
identified, the proposal will be returned to SPB along with the needed information to make the proposal 
schedulable. SPB then reworks the proposal and redelivers it to SPSS. 

TRANSFORMATION OPRs 

Transformation software problems or deficiencies identified with the products must be addressed prior to the final 
preparation and testing of the proposal to be scheduled. A list of Operations Problems Reports (OPRs) is 
reviewed in detail by examining the database values loaded into the various relations. If a problem is deemed to 
exist for a given case, the database is modified in accordance with the OPR. Following completion of the OPR 
analysis, the Obsets and Scheduling Units (SUs) may then be updated in order to prepare the data for test 
scheduling. The list of Transformation OPRs is updated as necessary as new OPRs are submitted and others 
fixed and installed in order to maintain as current and viable an operational system as possible. All stages of 
activities including the loading, review of OPRs, and subsequent preparation and testing are recorded in a History 
File in order to track all proposal work from input to final execution on board the spacecraft 

Updating 

After the initial proposal preparation has been completed the Obset data must be updated. The updating process 
involves generation of Space Telescope pointing data for single or multiple Obsets. The pointing data consists 
of calculation of target RA and DEC and aperture position in the Space Telescope V2-V3 coordinate plane, and 
orientation data for all underlying alignments and exposures. The updating process has to be executed any time 
changes are made to targets, alignments, or exposures of an observation set. The SU updating process is executed 
to prepare scheduling windows, observation times and user entered Science Instrument reconfiguration data. User 
entered values reflecting the targeted scheduling time are used during updating to determine the available windows. 
The windows calculated are 1) TV - Target dependent general visibility windows, which include the Sun and 



276 



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Cuudc Slkt Pur 6ma for mc id PCS pointxig 



ERROR 
ANALYSIS 



GSSS 

Softwnc iodmtifics 
firyfiH **^' GS piT» 
wfaxb i&icfy tbc 
SpKCcnfl pontiag 
•DdoncotMioo 
ro^uucniCDa givm 
lo ihcpropoMl 



■(response) 



I 



IGS ACOl 



(ERROR N 
ANALYSIS J 



SofTwan c<n^)uta V3 ocoinhDe roU dau for each 
GS pan For acb Obact U cakiilMca GS acj^namcn 
panmcun aod dcmmiDCi which pua piwxic tbc 
beat tappon ihraisbois tbc planned acba&iliag time 
frmc 



I 



Isultdating" 



(ERROR \ 
ANALYSIS ^ :; 



Soflwarc deterrmncj the mtcnccuon of the choaco 
GS pair viaibilxy wsidowi wuh tbc otfa^ Obaet 
wndowa to jxmwic the final DN /DO Scbctkiliiig 
wodowi 



[FLIGHT SQIEDULING^ 



I 



CCLIST 

BuiklcT addi the Tcqinrad SUi to tbc fbgbt Cand^laic List Hid 
fua tbc calendar Begin Hid Eod boLBxlary parvDcun 



CALENDAR BUILDING 

— "N. Builder acfaakiln lUpoanUcSUi in 
ANALYSIS } inonty okJct until aB mtaipa have 
. ^ been ohatuted 



OS SELECT 

Builder a dcds tbc mdivukial Guide Stsr 
Z' ERROR A p»n, froai ihe OS ACQ lut for e^li 
I ANALYSIS J 



Obaet acfaedulod tn the CalaxiMi 



I 



SMSGENT.RATIONi 



/^ ERROR \ 
i ANALYSIS t 



Builder CKOoMa SMS gajcnimg icAwaie 
lo read ibc flight c alcndar kkI prnducc the 
itniirDd r^rtTifngHng The builds* tbco 
irvicwa tbc wlput |>tk1uc(r for cxras ihM 
ndgbi TcquiTc calcTxlB' lebuilding . 



I 



ISMS REVIEWl 



Uaa'anuaiea a detailed exMnmaoori of tbc SMS to 
docmime if SMS a aaciafaciofy or oceda n 



AfUT successful SNC gaurautx], tbc ukt 
prcpBO a Noucc for ImiTn.i Review by 
Commaadaig ■>d It istnimog Tcxns. 




ISMS DKLlVERVl 



I 



,„ ^„ f ERROR \ 

^^^ \ ANALYSIS J 



U(cr cmfiilca DouficatioD a 
aodSCS review tcMna kkI I 
forfnal fbghip 



i^dudiiig comcDcau fnxn SPSS 
b it aloDg wub ibc SMS to PASS 



277 




Moon avoidance angle, 2) SN - Roll Normal windows and SO - Roll Off-Nominal windows which are based on 
the aperture position, target position, orientation type and V3 position angle, 3) RN - Restricted Normal 
windows and RO - Restricted Off-Nominal windows which are based on the guide star acquisition data set results. 
4) PC - Phase Critical windows, 5) SF - Surface Feature windows, 6) DN - Derived Normal windows and DO 
- Derived Off-Nominal windows which are the intersection of TV, SN, SO, RN, RO, PC, and SF windows, and 
7) SU - Scheduling Unit windows which are calculated to bound the DN and DO windows. The SU updating 
must be done before a candidate can be added to a Candidate List (CCLIST). SU updating must be repealed if any 
modifications are done to the underlying Obsets, alignments , and exposures such as a new Guide Star request, 
changes to science times, PCS Scenario usage, alignment parameters, etc. 

Test Scheduling 

Following the successful completion of Obset/SU updating, or resolution of updating problems, the candidate 
SUs contained within the proposal are placed onto a CCLIST. All candidates which can be scheduled usually will 
be, with the exception of large groups of identical SUs, in which case a small subset will be tested. Candidates 
whose target visibility windows are closed for the targeted timeframe must wait until their windows are open. 
The test calendar is reviewed and verified to be free of significant errors prior to generation of the command 
sequences. If scheduling problems are encountered or inconsistencies between scheduling requirements and the 
target time frame noted, problems must be resolved before proceeding. If severe enough, it may be necessary to 
return the proposal to SPB. 

Test SMS Generation 

A Science Mission Specification (SMS) is generated which contains the associated commanding required to 
execute the instrument operations and spacecraft maneuvers for the science observations. The SMS and all error 
output products are reviewed upon completion to verify that the timing of alignments is adequate, no errors exist 
in the target locations vs aperture locations, all planned exposures are present, and so on. If no significant errors 
are noted, the SMS may be sent out for Internal (STScI/Commanding) Review. On the other hand, if significant 
problems are encountered with timing, spatial scan parameters, etc., it may again be necessary to resolve the 
conflicts with SPB and/or the proposer. 

Change Requests 

SPB is notified of problems which have been identified in any of the error analysis processes of proposal 
preparation. These errors can result in one of two things; 1) the proposal being returned to SPB or 2) SPB 
sending a Proposal Change Request to be implemented that will correct the problem. Change Requests can also 
originate from the Proposer, and the Science Commanding System (SCS). These Change Requests can be 
implemented up to the point of Flight SMS Generation and have from minor to severe impact upon normal 
SPSS operations. 

Test SMS Delivery 

A Delivery Notice is prepared noting the proposal tested in the SMS to be reviewed. Special circumstances such 
as SUs which could not be tested due to closed windows or special scheduling requirements as provided in the 
proposal are described in the notice. At this point, specified SUs defined for the Flight SMS may be scheduled 
by the Flight Preparation crew. SPSS will be informed of any subsequent problems noted by the simultaneous 
Internal Review in progress. 

Candidate Pool 

Once a given proposal has passed proposal preparation, it is considered to be "flight ready"; that is, it is ready for 
flight SMS preparation activities as described below. Its associated SU's are now considered to be a part of the 
pool of scheduling candidates. At this stage nothing further is done with the SUs until SPSS is notified that 
they have been selected for scheduling on a specific flight SMS. SPB is responsible for providing the list of 
which usable candidate SU's from the existing pool are to be executed on each flight SMS. Currently, this 
notification is being done via the Flight SMS Order Form which is defined based partially on the Science 
Verification (SV) observing cycle requirements. The order form provides the list of SU's which are to be used 
from the candidate pool, as well as highlighting special proposal requirements such as scheduling priority, 
ordering of SU's relative to each other, pointing control requirements, etc. This form of notification has been in 
use from launch to the present As the mission progresses into the General Observer (GO) cycles and the schedule 



278 



requirements become less rigid, SPB intends to provide a software-generated list of SU's from the pool of 
candidates. 

Obset Updating 

The updating requirements for making an Obset ready for flight scheduling are the same as those covered in the 
Proposal Preparation phase with the addition of guide star processing. The updating starts with the generation of 
a list of all Obsets to be added to the CCLIST. This list can be used with the Updating Command Procedure to 
ready the Obsets for Guide Star Requests. This procedure generates all of the pointing data and sets the "ready for 
guide star request" flag which allows guide stars to be requested for the Obsets. Once the updating is completed 
and all errors have been corrected the Obsets are ready for the next step which is Guide Star Requesting. 

Guide Stars 

A major part in the flight SMS preparation activity involves the selection of guide stars for Fine Guidance 
System (FGS) pointing control. In fact, guide star processing takes up at least 15% of the time required during 
the flight SMS generation activities. This is a three stage process involving: 1) The identification of those 
Obsets which require guide stars and generation of the request, 2) the processing of the request by the Guide Star 
Selection System (GSSS), and 3) the processing of the guide star response data to apply acquisition specific 
selection criteria to the pool of candidate guide stars in order to achieve selection of guide star pairs which have 
the highest probabihty of success for the given acquisition. Of the total guide star processing time required. Steps 
2) and 3) involve the most in both the actual processing, results analysis, and troubleshooting. A majority of 
the Obsets can be processed automatically by the software from request through acquisition selection processing 
to provide satisfactory results the first time through. The remainder require the user to analyze processing results 
and interact with the software to produce the desired GS support. Problems due to physical constraints in the 
Field Of View (FOV) which limit the accessible guide star pair candidates using the default GSSS processing 
parameters are usually identified during step 2. These occur for sky regions which contain an extremely low 
density of field stars, extremely high density of field stars such as globular clusters, or fields which may be 
washed out in the GSSS catalog such as those near very bright stars, nebulosities, etc. Interactive processing 
may be used to modify the GSSS run time parameters in order to access other guide star candidates, or to generate 
diagnostics for analysis and documentation to show why a given observation cannot be supported due to real 
physical constraints. A recent quick survey of the PMDB and operations staff suggests that the ratio of the 
percent of Obsets per type of request to GS request/response processing time required breaks down approximately 
as follows: 



Type 


% Obsets 


% Time 


Automatic 


85 


60 


Interactive 


15 


40 



Once a pool of candidate GS's is returned to SPSS, the processing of step 3 computes sets of V3 roll ranges over 
which GS support is available. The candidate GS's which are used to make up these roll ranges are chosen based 
on a set of acquisition criteria. These ranges are then used by the SU updating function to compute the RN/RO 
GS support windows during the scheduling window computation. Following successful DN/DO window 
generation, the SU is ready for the next step; flight scheduling. 

Flight Scheduling 

This is the next manually intensive step in the SMS generation of activities. Creation of the CCLIST and the 
subsequent scheduling of activities are mechanically the same as described for the proposal preparation tasks. 
However, we are no longer working with SU's from just one proposal. In addition, there are special activities 
which must be prepared and scheduled to control Scientific Instrument (SI) states and Space Telescope (ST) 
pointing at the calendar boundaries. The scheduling scenario for the SU's hsted on the Flight Order Form occiu"S 
in three basic passes as follows. First , the time critical and other SU's with special scheduling requirements 
are put on the calendar manually. Second, automatic scheduhng software can be used to attempt scheduling of 
non-critical pointed SU's in priority order. This pass can be done in parallel with other SPSS activities (batch 
mode, etc.). The final pass is the most time consuming for the SPSS scheduler. It involves manual attempts to 
schedule the remaining pointed SU's. Problem SU's are analyzed to determine if the candidate can be scheduled 
on this calendar. Lower priority candidates may need to be removed at this time if that action would facilitate the 



279 



scheduling of higher priority "problem" SUs. The analysis results are reviewed and directed to the appropriate 
management levels for resolution. (For example, alignments whose science is too long to fit into an orbit are 
referred to SPB to see if a reduction in the time is possible.) Some candidates may be dropped from the Right 
Order Form and rescheduled at a later date due to these types of problems. After all the problems with pointed 
candidates are resolved, any internal calibration SU's are scheduled. Again, this third scheduling pass is the most 
manually intensive and can be iterated many times over. Once completed, the calendar is reviewed in detail to 
verify that all scheduling requirements have been met, and that no known problems (software-created, etc.) exist. 

GS Select 

The guide star processing stage above generates a pool of candidate pairs which support the calendar timeframe. 
The actual selection of the specific guide star pairs which will be used in flight is done after the calendar is 
complete and the exact schedule time is known. Problems at this stage are rare since the major guide star support 
problems have been worked out during the request and response processing stage. Since this is the last step prior 
to generating the SMS, the selection results are reviewed in detail at this time to verify that they satisfy the 
chosen FGS acquisition scenario for the Obset. This is a low manual impact part of the Flight SMS 
preparation activities. 

SMS Generation 

While this step in the processing is not manually labor intensive, it is one of the most time consuming with an 
average runtime of 3-4hrs for a one week SMS. The SMS generation software basically reads the calendar and 
queries the PMDB to determine what activities are being requested in the schedule. As it reads the calendar, it then 
extracts the instructions from the PMDB which are necessary to provide the science instrument or pointing 
control commanding for each activity, and then builds that commanding into the SMS. There are literally 
hundreds of activities which are expanded into thousands of commands for each one week SMS that is processed. 
The limited manual effort at this step comes after the SMS generation has completed. Then a brief error analysis 
is done to identify any problems in the SMS processing itself due to either proposal data or commanding 
deficiencies. Problems which are identified are resolved via PMDB fixups to the proposal data, commanding, etc., 
and by iterating through the previous SMS preparation steps to get the schedule back to its SMS readiness state. 

SMS Review 

A more detailed SMS Review is conducted by SPSS and Science Commanding System (SCS) on a SMS after it 
has been generated by SPSS. The initial review by SPSS consists of a series of checks on the selected Guide 
Stars, SMS file and the Database information used to generate the SMS. Upon completion of the SPSS 
review, the SMS is either sent for commanding review or returned to SPSS for fixes and regeneration. All errors 
found during the review process are catalogued in an error summary log. These error summary logs contain the 
status of the error and the solutions to be used in fixing them. If the SMS passes the Review process a detailed 
Delivery Notice which accompanies the SMS to PASS is generated. The SMS is then sent by electronic means 
to PASS at Goddard and will go through additional analysis and eventual uplink to the Space Telescope. 



Mailing address: 

K. E. Reinhard, H. H. Lanning, and W.M. Workman, HI 
Computer Sciences Corporation 
Space Telescope Science Institute 
3700 San Martin Drive 
Baltimore, Maryland 21218 



280 



THE SCHEDULING OF SCIENCE ACTIVITIES 

FOR THE 
HUBBLE SPACE TELESCOPE 

D.K. TAYLOR^. K.E. REINHARD^ , H.H. LANNING^ . D.R. CHANCE^ 

and 
E.V.BELL. 11^-'^ 

Overview 

The Science Planning and Scheduling System (SPSS) is the operational software portion of the Science 
Operations Ground System (SOGS) responsible for scheduling science activities onboard the Hubble Space 
Telescope (HST). In this presentation, we show a chronological order of the activities and features of SPSS that 
take an observing proposal from Transformation to execution. 

Once a proposal is entered into the relational Proposal Management Database (PMDB) by conversion software 
known as Transformation, a proposal consists of Scheduling Units (SU), Observation Sets (Obset), Alignments, 
and Exposures. These represent the observing structure used by SPSS in which the exposure is the basic 
building block containing the proposal logsheet information including Science Instrument (SI) used, mode of 
operation, spectral element, and aperture. Exposures are merged into alignments which manage the pointing of 
the spacecraft including target position, roll, and the timing of the observations. The alignments are merged 
into Obsets which control the type of pointing and acquisition of guide stars. The SU controls the execution of 
the Obsets, Alignments, and Exposures and is the major building block used in the construction of a detailed 
timeline of science activities known as a Calendar. The scheduling of each SU on a calendar consists of 
calculating guide star acquisitions, slew activities, target visibility, science instrument transitions, orbital 
characteristics, etc. Whenever observational requirements permit, the ordering of SUs is chosen to minimize 
slews and other time consuming activities onboard the spacecraft 

After an acceptable calendar has been built, a Science Mission Specification (SMS) is generated. A SMS is an 
ASCII file consisting of the expanded commands from the calendar, calculated alignment times, expanded 
exposure commands, and orbit relative commanding. The SMS is then sent from the Space Telescope Science 
Institute (STScI) to the Payload Operations Control Center Application Software Support (PASS) at Goddard 
Space Flight Center where it is merged with engineering commands and converted to binary for spacecraft upload. 

Shown below is a portion of the exposure logsheet from the proposal 03123, "Revised FOS Combined Mode II 
Target Acquisition", which was a Science Verification test proposal. 



Commen 
ACQUIS 



ts: LI 
ITION 



NE 1-4 DEFINE THE BINARY 
MODE TEST 



EXPOSURE LOGSHEET Id = 3123 (P) 


















Page: 1 


1 


2 


3 


4 


5 


e 


7 


8 


9 


10 


11 


12 


13 


14 




Ln 
Nm 


Seq 

Nam 


Target 
Name 


Instr 
Conf ig 


qper. 
Mode 


Aper 
orFOV 


Spectral 
Element 


Centrl 
Haveln 


Optional 
Parameters 


Num 
Exp 


Time 


S/N 

Rel.T 

ime 


Fix 
Re£ 


Pr 


Special 
Requirements 


1 


DEF 
BIN 


NGC- 
lBe-136 


« 


ACQ 


4.3 


MIRROR 






1 


33S 




1 


1 


INT ACQ FOR 

2:SEQ 1-4 NO 

GAP; CYCLE 0/1- 

213 


2 


- 


* 


* 


ACQ/BIN 
ARY 


' 


" 




BBIGHT-330000 
.,FAINT-275 


1 


5.6S 




1 


1 


ONBOARD ACQ FOR 
3; 


3 


- 


* 


- 


ACQ 


" 


* 






1 


33S 




1 


1 




3. 

50 


" 


TALED 


* 


ACQ/BIN 
.ARY 


0.3 


* 




BRIGHT-650000 
,FAINT-275 


1 


lis 




1 


1 


ONBOARD ACQ FOR 
4; 


4 


* 


NGC- 

188-136 


• 


ACQ 


4.3 


MIRROR 






1 


33S 




1 


1 





In the following section is a portion of the one week calendar pertinent to this test, a fraction of the SMS 
generated from this calendar, and various plots and charts showing the constraints and observing restrictions 
routinely encountered by SPSS. Certain sections of the calendar are circled and titled. In the following text, 
under the same headings, are descriptions of the calendar activities. Following the calendar descriptions are the 
same type of descriptions for the SMS portion of the example. 



1 with Computer Sciences Corporation / Space Telescope Science Institute 

2 with ST Systems Corporation / National Space Science Data Center - Goddard Space Right Center 



281 



CALENDAR 





00- 


??- 


43 


J Slews and 

014:00:27:53 


FHST Opdi 

MF Slew (AN= 


ites 1 










('on 


0,RR= 7,DE= 85,PA=250,OR= 


0, 


) 


) 


i^on 


00 


23 


53 


014:00:33:23 


MF 


FHST UDdt 


e <FULL ,MAN,E= 119,3, , ) 






) 


on 


no- 


4') 


05 


014:00: 49:06 


SI 


UP FOS 


READY 


03123-OG6 


08: 


01 


01 


on 


on 


4<t 


06 


014:00:58:36 


SI 


UP FOS 


RED LVONA 


03123-0G6 


08- 


01 


01 


014 


no- 


SB- 


36 


014:01:00:26 


ST 


UP FOS 


RED HVONA 


03123-0G6 


OB- 


01 


01 


014 


01 


09 


31 


014:02:01:51 




F/S AVD 


(EXT,L= 20.4) 


03123 


OS 


01 


01 


014 


01 


09 


43 





Main SU 0312308 *•*••■•**••*•*•••••**•*** 


... 


... 


•• 


014 


01 


09 


43 


014:01:30:26 


MF 


PCS AQ(FGS ,£• 69.COARSE2 ) 


03123-0G6 


08 


01 




014 


01 


13 


52 






NODE Crossing 3937 










014 


01 


29 


16 


014:02:10:13 




SLW FHST 


1 (EXT.I^ 37.01 


03123 


08 


01 


01 


014 


01 


29 


52 


014:01:33:02 


SI 


UP FOC 


DET STDBY96 


03028-OCl 


52 


01 


01 


014 


01 


29 


58 


014:02:05:17 




SHADOW 


(ENTRY) 










014 


01 


30 


26 


014:01:37:31 ■ 


HF 


Com 


<DN ,MA , ,E.2) 


03123-0G6 


08 


01 


01 


014 


01 


30 


26 


014:01:37:35 • 


HF 


Tar FOS 


YRD4 3 


03123-0G6 


08 


01 


01 


014 


01 


30 


26 


014:01:37:35 • 


MF 


Sci FOS 


YRD4 3 1 


03123-0G6 


OS 


01 


01 


014 


01 


32 


02 


014:02:01:28 




SLW FHST 


2 (EXT,L= 37.11 


03123 


08 


01 


01 


014 


01 


33 


02 


014:01:33:03 


SI 


UP FOC 


POWER 


03028-001 


52 


01 


01 


014 


01 


37 


35 


014:01:53:35 • 


MF 


Sci FOS 


2 


03123-0G6 


08 


02 


01 


014 


02 
02 
02 


01 
05 
14 


51 
17 
11 


014:02:46:16 




F/S AVD 


(ENT,T^ 20.4) 


03123 

1 03123 
02168-0D6 


08 

08 
01 


01 

01 
01 


01 


014 
014 


Science Instrument 

014:02:23:15 SI UP HRS 


Transition 

STDBY 


01 


f 014 


0? 


23 


14 


01' 


014 


0? 


23 


15 


014:02:23:16 


SI 


UP HRS 


DET2 STDBY2 


02168-006 


01 


01 


01 


i^ 014 


0? 


23 


16 


014:02:51:46 


SI 


UP HBS 


DET2 HV0N2 


02168-0D6 


01 


01 


oi; 


014 


0? 


46 


16 


014:02:51:36 


MF 


R£AQ/N(FGS , E= 50,COARSE2 ) 


03123-0G6 


08 


01 




014 
, 511 
^014 


02 
02 


46 

ML 
51 


16 
36 


J Slews and 

014:02:51:46 • 


FHST Dpda 

MF SAM (ANG 


tos I '"■"' 


03123 


08 


01 


01 


,= 0.0000,ROLL= 


07) 






;l,^ 


■irr? 


n? 


■^l 


IS 


014:05:00;1« 




f/S AVB 


(ESiT.L- 30.41 










014 


0? 


51 


46 


014:03:00:16 « 


MF 


SCI FOS 


YRD4 3 3 


03123-0G6 


08 


03 


01 


014 


07 


51 


46 


014:03:00:16 • 


MF 


Com 


(UP ,SSA, ,W,1) 


03123-OG6 


08 


03 


01 


014 


03 


00 


16 


014:03:00:26 • 


MF 


SAM (ANG= 0.0000,ROLL= 


081 








014 


03 


00 


16 


014:03:08:38 




F/S AVD 


(EXT,T^ 45.0) 


03123 


08 


04 


01 


014 


03 


on 


26 


014:03:08:38 * 


MF 


Tar FOS 


YRD4 3 


03123-0G6 


08 


04 


01 


014 


03 


00 


26 


014:03:08:38 • 


MF 


Sci FOS 


YRD4 3 4 


03123-0G6 


08 


04 


01 


014 


03 


03 


58 


014:03:04:38 


SI 


UP HRS 


WARM 


02168-0D6 


01 


01 


01 


014 


03 


06 


49 


014:03: 42:07 




SHADOW 


(ENTRY) 










014 


03 


08 


38 


014:03:13:48 




F/S AVD 


(EXT,L= 53.5) 


03123 


08 


05 


01 


014 


03 


OS 


38 


014:03:08:48 ■ 


HF 


SAM (ANG= 0.0000,ROLL= 


09) 








014 


03 


08 


48 


014:03:13:48 • 


MF 


Sci FOS 


YRD4 3 5 


03123-0G6 


08 


05 


01 


014 


03 


13 


48 


014:03:13:58 • 


MF 


SAM (ANG= 0.0001,ROLL= 


.091 








014 


03 


13 


48 


014:03:38:35 




F/S AVD 


(EXT,L= 54.4) 


03123 


08 


06 


01 


014 
014 


03 

03 


13 
13 


58 
53 


J Target Vis 


ibility 

F/S AVD 


— rRD2 OBAR 

rRD2 OBAR 6 
(ENT.L= 20.4) 


03123-OG6 

03123-0G6 

03123 


08 
08 

"oF 


06 
06 
St 


01 
01 


C 014 


03 


38 


35 


014:04:22:59 


01\ 


014 


03 


42 


07 


014:04: 43:39 




SHADOW 


(EXIT) 










1^ 014 


03 


50 


54 


014:04: 10:42 




F/S OCC 


(BRIGHT EARTH) 


03123 


08 


06 


oj 


014 


04 


04 


38 


014:04:07:08 


SI 


UP HRS 


DET2 READY2 


02168-OD6 


01 


01 


01 


014 


04 


03 


08 


014:04:07:18 


SI 


UP HRS 


DET2 0PER2 


02168-OD6 


01 


01 


01 


014 


04 


?? 


59 


014:04:28:19 


MF 


REAO/N(FGS , E= 42,COARSE2 ) 


03123-0G6 


08 


01 




014 


04 


22 


59 


014:04:28:19 




F/S AVD 


(EXT,L= 20.3) 


03123 


08 


06 


01 


014 


04 


27 


18 






NODE Crossing 3939 










014 


04 


28 


19 


014:04:28:29 • 


HF 


SAM (ANG= 0.0001,BOLL= C 


.14) 








014 


04 


28 


19 


014:04:54:12 




F/S AVD 


(EXT.L= 30.4) 


03123 


08 


07 


01 


014 


04 


28 


29 


014:04: 36:00 ■ 


HF 


Sci FOS 


YRD4 3 7 


03123-0G6 


08 


0/ 


01 


014 


04 


33 


03 


014:04:37:53 


SI 


UP FOC 


DET OPER96 


03028-OCl 


52 


01 


01 


014 


04 


36 


00 


014:04: 37:36 • 


MF 


Sci FOS 


8 


03123-0G6 


08 


08 


01 


014 


:04 


37 


36 




End SU 0312308 •••• 




... 


... 


... 


014 


:04 


37 


18 


014:04:37:48 


ST 


UP HRS 


DET2 0BS2 


02168-0D6 


01 


01 


01 


014 


:04 


37 


36 


014:04:38:16 


SI 


DOWN FOS 


RED LVONA 


03123-0G6 


08 


08 


01 


014 


:04 


38 


16 


014:04:41:56 


ST 


DOWN FOS 


RED HOLD 


03123-OG6 


08 


08 


01 


014 
\/\ 


:04 


41 


56 


014:04:41:57 


SI 


DOWN FOS 


HOLD 


03123-OG6 


08 


.08 





\/\ 
016 


;01 


11 


08 


016:01:11:09 


ST 


UP FOS 


READY 


03123-0G6 


:05 


:01 


01 


016 


:01 


11 


09 


016:01:20:39 


SI 


UP FOS 


RED LVONA 


03123-OG6 


:05 


:01 


01 


21£ 


-01 


-?fl 


■19 


016:01:22:29 


fiT 


UP FOS 


RED HVONA 


03123-0C6 


;5S 


;51 


01 


016 


:01 
-01 


31 


.35 


016:02:15:28 




F/S AVD (EXT, L= 20. 4) 


03123 


:05 


:01 


01 



1)16.01.31-46 ni6-ni-S2-79 MT PC.9 XnirCI R^ 69 COAB.9R2 ) 031 23-QC.t : 05 : 01 



016:01: 35: 33 

ni6.ni.S2.29 01601.5934 » MF O^ 



NODE Crossing 3967 



016:01: 
016-01- 



016:01: 
016:01: 



52:29 
52-29 



016:01 
"16-01 



59:38 



016:02: 
016:02: 



59:38 
S9-3R 



30:31 
15:38 



MF Tar FOS 



SHADOW 
MF SCI FOS 



YRD4_3 
VR"4 3 



.B.81 031Z3-066;05;01 01- 



03123-006:05:01 01 
03123-0r.6:05;01 01 - 



01 



016:02:15-2(1 016:D2-1S:3B « MF SAM (AN(^ 0000 ROI.I.= 



016:02; 
016:02: 
016:02: 
016:02: 
016:02 
016:02 



15:28 
15:38 
15: 38 
26:43 
30:31 
35:41 



016:03:09:19 



016:03: 
016:03: 
016:03: 
016:03: 
016:03: 
016:03; 
016:03: 
016:03; 
016:03: 
016:03: 
016:03: 
016:03: 
016:03: 
016:03: 
016:03: 
016:03: 



08:19 
12:16 
13:39 
13: 39 
13:49 
13:49 
22:01 
22:01 
22:11 
27:11 
32:04 
36:53 
36:53 
37:03 
40: 31 
43:49 



016:02 
016:02: 
016:02: 
016:03: 
016:03; 
016:02; 
016:03; 



26:43 
24:08 
24:08 
08:19 
32:04 
56:28 
13-39 



F/S_AVD 
MF Set FOS 
MF Com 

F/S_AVD 

SHADOW 

F/S OCC 



(EXT,L= 36.21 
YRD4_3 3 

(UP ,SSA, ,£,2) 
(ENT,L= 14.8) 
(EXIT) 
(BRIGHT EARTH) 



016:03:13:39 



MF RPAn/N(Fn.'! P.= 40 rnARSR2 ) 



016:03 
016:03 
016:03 
016:03 
016:03 
016:03 
016:03 
016:03 
016:04 
016:03 
016:04 
016:03 
016:04 
016:03 



:13:49 
: 22:01 
:22:01 
:22:01 
: 22:11 
: 36:53 
:27:11 
:36:53 
:07:22 
:37:03 
:03;33 
: 43:49 
: 25:37 
:45:25 



F/S_AVD (EXT,L= 20.4) 

NODE Crossing 3968 

MF SAM (ANG= 0.0000,ROLL= 
F/S_AVD (EXT.L= 30.7) 

MF Tar FOS YRD4_3 

MF Sci FOS YRD4 3 4 



03123-006:05:02 

■ 031 

03123:05:03 01 

03123-006:05:03 01 

03123-006:05:03 01 

03123:05:03 01 

03123:05:03 01 

03123-0g6;05;01 



0.0000,ROLL= 
(EXT.I^ 45.2) 
YRD4_3 5 

YRD0_3 6 

(ENTRY) 
0. 0000, ROLL= 



MF SAM (ANG= 

F/S_AVD 
MF Sci FOS 
MF Sci FOS 

SHADOW 
MF SAM (ANG= 

F/S_AVD (EXT,L= 55.0) 

MF Sci FOS YRD4_3 7 

SLW_FHST 3 (EXT.L= 21.3) 
MF Sci FOS 8 



03123:05:03 01 

07) 

03123:05:04 01 
03123-006:05:04 01 
03123-006:05:04 01 
08) 

03123:05:05 01 
03123-006:05:05 01 
03123-OG6:OS:06 01 

.091 

03123:05:07 01 
03123-OG6:05:07 01 

03123:05:07 01 
03123-066:05:08 01 



016:03:45:25 



End SU 0312305 



016:04:34:44 016:05:24:13 

016:04:34:44 016:05:24:13 

016:04:36:36 016:05:22:17 

016:04:36:36 016:05:22:17 



TDRS 
TDRS 
TORS 
TDRS 



(WEST, MA , RET, VIS ) 
(WEST, HA ,FWD,VIS ) 
(WEST, SSA, RET, VIS ) 
(WEST, SSA,FWD, VIS ) 



016:04: 49:42 



]PCS Acquistion \_ 

(016:04:49:42 016:05:10:25 



MF PCS AQ(FGS , E= 52.COARSE2 ) 03123-006:06:01 

016:05:08:55 016:05:44:12 SHADOW (ENTRY) 

016:05:10:25 016:05:17:30 * MF Com (DN , HA . .W.l) 03123-066:06:01 01 



SMS 



OOMV /Ttext, TiKE^ (ORB, 3966, EAscNCR, oiH2i^-r Transition 

/ BEGINTEXT: t- 

BED Low Voltage To High Voltage 



FOS 



BEGINTEXT; 

RECON- 
ENDTEXT 

SMSTIME=1 991. 016: 01: 20: 39.000 

BEGINNING AtV COMMAND BLOCK YSHVC« 

RTSCTRL.FUNC(ACT) , RTSID (YHVCWO 94 ) , TIME= (OHB, 3966, EASCNCR; ; 
01H2 1M1 9. OOOS) 

SMSTIME=1991.016:01: 20: 39.000 

GROUP, pyHV_20,TI«E= (ORB, 3966, EASCNCR. 01H21M51 . OOOS) 

SMSTIHE=1991.016:01:20: 41 .00 

BEGINNING AfiV COMMAND BLOCK YSHVSET 

GROUP,PYFOCUS,FOCUS(2.eB) .TIME=(ORB. 3966, EASCNCR 

01H2 3M31.000S) 

SHSTIME=1991.016:01:22:21.000 

GROUP, PYHVDAC.KVOLTS (2.22E+01) , TIME=(CRB, 3966, EASCNCR 

01H2 3M32.000S) 

SMSTIME=1991.016:01:22:22.000 

BEGINNING AiV COMMAND BLOCK YSREFD 

GROUP, PYREFDAC, REFDAC(4 .IBSE+Ol) ,T1ME=(0RB, 3966, EASCNCR 

01H2 3M33.000S) 

SMSTIME=1991.016:01:22:23.000 



Guide Star 

Acquisition 



:GSACQ,ASTID(1) , CENTER (BOTH) ,CPNAME(PQ ( 

,GSllDEC(8.534364 004a3n77E+01) ,GS11FG 

,GS11ID(0^61901205) , GS11MAG<1 .27621EtO J 

,GS1 IRA (9. 96634 31 4 8 65571 6) .GSllRAD (6. s ilft JtiJiJ i;f^3 J ML+Ul 

,GS12DEC(8.53960564 4 93810 4E*01) ,GS12FGS(2) 

,GS 12 ID (04 61 900995) ,GS12MAG(1 .29619E+01) 

.GS12RA( 9. 66187 0114189522) ,GS12RAD<6. 68513454799537 3E+01) ; 

,GSlDOM(2) ,GS21DEC(8.5133354 46334 971E+01),GS21FGS(3) 

,GS21ID(04 61 900434) , GS21MAG ( 1 . 05773E+ 01) 

,GS21RA(7. 962712834845506) , GS21RAD ( 1 . 5E+01 ) 

,GS22DEC (6. 511 9835259691 07E+01) ,GS22FGS(3) 

, GS22IDt 04 61 900520) ,GS22MAG(1 . 13659E+01) 

,GS22RA {6. 98771570997109) ,GS22RAD(1 .5E+01) ,GS2DOM{2) 

,NOSLEW,NUM_PAIR(2) ,ACQTYPE(2) . FHSTBIAS(l) , GSllFT {F583W) ; 

.GSIIKIX ( 3. 2227 4 6 80 67 502 57E- 02) 

.GSIIKIY (4.042746903905828E-02) 

.GS11K3X(4.707391076601762E-01) 

,GS11K3Y (4. 3152 454594 9281 4E-01) ,GS11ML(1 14) , GS12FT (F5B3W) ; ; 

,GS12K1X(3.2547 031905507 03E-02) 

,GS12Kiy ( 4. 07 47 0328 77 062 74 E- 02) 

,GS12K3X (4. 660931 45459664E-01) 

.GS12K3Y (4.2S2305822068399E-01) ,GS12ML(97) , GS21FT (F58 3W) 

,GS21K1X (2. 4 006589450907 92E-02> 

.GS21KlY(2.570658e4 87 07524E-02) 

,GS21K3X (6. 1789242692 698 39E- 01) 

,GS21K3Y (5. 613451 652511787E-01) ,GS21ML(1207) 

.GS22FT(F583W) , GS22K1X (2 . 4 13789333687 464E-02) 

,GS2 2K1Y (2. 5 837 89 3595 90 67 6E- 02) 

,GS22K3X (6.1 375636552 48157E-01) 

,GS22K3Y (5.58259515 37 09073E-01) ,GS22ML(5B9) ,PLNTPRLX<0) 

, RCHVM<0 .0) , TARGETAO( PRIMARY) , WHICHACQ (BASELINE) , END= (ORB; ; 

, 3966, EASCNCR, 01H53M39. OOOS) . START=(ORB, 3966, EASCNCR 

, 01H32M56.000S) 

;SMSTIME=1991.016:01:31: 4 6.000 



Science 

Activities 

A 



'rTEXT,TIME=(ORB, 3966. EASCNCR, 01H5 3M39 

BEGINTEXT; 

START PROP=03123 , PROG=0G6 ,OBSET=05 

NGC-108-13 , FOS/RD , ACQIMAGE , 37.0 

NOCHANGE , A4_3 , MIRROR . G780H , 
Begin Ctoservation 

END TEXT 

;SMSTIME=1991.016:01:52:29.000 

; BEGINNING AiV COMMAND BLOCK YSFGWP 

:RTSCTRL. FUNC(ACT) , RTSID (YMOTR094) , TIME= (ORB, 3966. EASCNCR; 

,01H53M39.000S) 

; SMSTIME=1 991 . 1 6 : 01 : 52 : 29 . 000 

: GROUP, PYFMTCO. FMTCOOE ( ' 73 'X) . FORMAT (2) , TIME=(ORB. 3966 

, EASCNCR, 01H53M39.000S) 

;SMSTIME=1991.016:01:52:29.000 

:GROUP,PYFILTER,DIR(FWD) , FILTER (MIRRORA) ,TIME=(ORB, 3966 ; 

.EASCNCR. 01H53M40.0 00S) 

;SMSTIME=1 991. 016: 01:52: 30.00 

: SCIHDR. INSTID (FOS) .OBS_ID(01) .OBS_SET(05) , PROG_ID(0G6) ; 

.WORDll (32) ,T1ME=(0RB, 3966, EASCNCR, 01H53M4 1 . OOOS) 
i^MSTIME=l 991. 016: 01: 52: 31.000 / 

'^OMCON, INST_ID(FOS) ,OBS_ID(01) ,OBS_SET(05) .4cC^4CX>N PX 

.RATE (4. 0) .SERVICE (MAR) , TAPE_OPT (BACKUP) , END -i t» . «^, ^ ^«« 1 ;; 

, EASCNCR, 02H0OM12. OOOS) , START=(ORB, 3966, EASCNCR 

, 01H53M4 9. OOOS) , TAPE_BEG= (ORB, 3 966, EASCNCR, 01H56M49 . OOOS) ; ; 

, TAPE_END= (ORB, 3966. EASCNCR. 02H00M31- OOOS) 

VjSMSTIME-1991.016:01:52: 39.00 / 



Science 
Activities 

B 



: GROUP, PYMCSTEP,DIR(CW) , TIME=(ORB, 39 

, 01H53M51.000S) 

;SMSTIME=1 991. 016: 01: 52: 41.000 
: TABLE, YOCKHI.OVERLITE (3000000) .TIME 
,01H53M55.000S) 
;SHSTIME=1 991. 016: 01: 52: 45.000 

:TABLE, YOCKLO,OVERLITE(300 00 00) ,TIME= (ORB, 3966, EASCNCR 
, 01H53H56. OOOS) 

;SMSTIME=1991.016:01:52: 4 6.000 
BEGINNING A4V COMMAND BLOCK YSPTRNS 
.EGINNING A&V COMMAND BLOCK YSDEFL 







GROUP, PYDEFLEC,XOFFSET(0) , XBASE (A£ 

YOFFSET(-496) , YBASE (ASCMIRRO) ,Y_ 
, Y RANGE (1024) ,TIME=(ORB, 3966,EASCNC 
;SMSTIME=1991.016:01:52:47.000 
; BEGINNING AfiV COMMAND BLOCK YSDEFP 
: GROUP. PYPATTRN, INTS (1) .OVERSCAN (5) S> 
, Y_STEPS (64) ,TIME=(ORB, 3966, EASCNCR, 01 H54M02. 00 OS) 
;SMSTIME=1991.016:01:52:52.000 
; BEGINNING AfiV COMMAND BLOCK YSDAQP 

:GROUP,PYACOPAR.CHNNL1 (256) , CHNNLS (20) ,DEADTIME(1 .OEtOl) 
, INIT HYS (HYSTER) . INITSLI (SLICE) , INITXDF (XDEF) 
, INIT_YDF(YDEF) , LIVETIME (2 . 062 5Et01 ) , RE JLIMIT (NOREJLIM) 
. TIME=(ORB, 3966, EASCNCR, 01H54M0 7 . OOOS) 
; SMSTI ME=1 991 . 016 : 01 : 52 : 57 . 000 
; BEGINNING AfiV COMMAND BLOCK YSDRP 

: GROUP, PYPATTS, PATTERNS (1) ,TIME=(ORB, 3966, EASCNCR 
, 01H54M16.000S) 

;SMSTIME=1991. 016:01:53: 06.000 

: GROUP. PYREADS. READOUTS (1) ,TIME=(ORB, 3966, EASCNCR 
,01H54M17.000S) 

;SMSTIME=1 991. 016 101:53:07.000 

:GR0UP,PYCLEARS,CLEARS(1) ,TIME=(ORB, 3966, EASCNCR ; :/ 

\(01H54M18.000S) -.J 



282 



/'016:05:H:15 


016:06:03:07 


TDRS 


(EAST, MA , RET, VIS ) 


016:05:14:15 


016:06:03:07 


TDRS 


<EAST,HA ,FWD,VIS J 


016:05:16:02 


016:06:01:19 


TDRS 


(EAST, SSA, RET, VIS ) 


^. 016:05:16:02 


016:06:01:19 


TDRS 


(EAST, SSA.FWD, VIS ) 



016 
016 
016 
016 
016 
016 
016 
016 
016 
016 
016 
016 
016 
016 
016 
, 016 
( 016 
016 
016 
016 
016 
016 
( 016 



05:10:25 
05:10:25 



|TDRS Contact ^ 



FOS 
FOS 



YRD4_ 



03123-0G6 
03123-OG6 



06:01 01 
06:01 01 



05:17: 34 
05:37:22 
05:37:22 
05:40:14 
05:40:23 
05:44:12 
05:49:07 
05:51:04 
05:51:18 
06: 03:04 



016:05:33:34 ' MF Sci FOS 



O3123-0G6:O6:02 01 



016:06:03:04 
016:05:51:18 
016:05:51:04 
016:06:21:47 
016:06: 45:45 
016:06:09:58 
016:07:17:34 



SAA 05 
SAA 07 
SAA 03 
F/S_AVD 
SHADOW 
F/S_OCC 
SAA 03 



PCS Acqulstion 



(ENTRY) 

(ENTRY) 

(ENTRY) 

(ENT.L= 

(EXIT) 

(DARK 

(EXIT) 

(EXIT) 

(F.XTT) 



03123: 
03123: 



06:01 01 
06:01 01 



06:21:47 016:06:27:07 MF REAQ/N(FGS , E= 48,COARSE2 ) 03123-0G6 : 06: 01 



06:21:47 
06:25:43 
06:27:07 
06:27:07 
06:27:17 



016:06:27:07 



016:06:27:17 



TSTT 
016 
016 
016 
016 
016 
Jlli. 
r016 
^016 
016 
016 
016 
016 
016 
016 
016 
016 
016 
016 
016 
016 
016 
016 
016 
016 
016 
016 
016 
016 
016 
016 
016 
016 
016 



06:27:17 



OS: 55: 41 
06:35:47 
06: 35:57 
06: 35:57 
06:45: 45 
06:48:23 
06:48:23 



JTDRS Contact ^ 

016:06: 35:47 • MF Com 



F/S_AVD (EXT,L= 20.4) 

NODE Crossing 3970 
MF SAM (ANG= 0.0000,ROLL= 
AVD (EXT,L= 30.7) 

FOS YRD4 



03123:06:01 01 



.07) 

03123; 
03123-OG6: 



06:03 01 
06:03 01 



(UP 



0U:0«:?5:57 ' Mf sAM — Ia!J5S — 0.(1I1D0,R4U,= 



.SSA. ,W.2) 03123-0G6:06:03 01) 



016:06:48:23 F/SAVD (EXT,L= 

016:06:48:23 • MF Tar FOS YRD4_3 

016:06:48:23 • MF Sci FOS YRD4_3 

016:07:21 : 02 SHADOW (ENTRY) 



45.7) 



03123: 
03123-OG6: 
03123-OG6: 



Science Alignment 



06:48:33 
06:48:33 



016:07:00:34 
016:07:00:59 



MF Tar 
MF Sci 



FOS 
FOS 



.0000,ROLL= 
I IFJIT.L. 55.4) 
YRD4_3 
YRD4 3 5 



03I23-0G6 
03I23-0G6 



06:04 01 
06:04 01 
06:04 01 



06:05 01' 
06:05 01 



07:00:59 
07:00:59 
07:01:09 
07:07: 40 



016:07:01:09 
016:07:17:14 



SAM (ANG= 0.0000,ROLL= 0.10) 

F/SAVD (EXT,L= 44.2) 03123:06:06 01 

Cna Daooa/vAc 1^ ^^^ YRD4_3 6 03123-OG6:06: 06 01 

;>/ui massages | hrs warm oi408-oen:13:oi oi 



/'016:07:16:28 


016:07: 45:23 


SAA 05 


(ENTRY) 






016:07:16:28 


016:07:34:18 


SAA 07 


(ENTRY) 






016:07:17:14 


016:07:58:31 


F/S AVD 


(ENT,L= 


14.2) 


03123:06:06 01 


016:07:17:34 


016:07:32:39 


SAA 3 


(ENTRY) 






016:07:21:01 


016:07:30:12 


SAA 02 


(ENTRY) 






016:07:21:02 


016:08:22:36 


SHADOW 


(EXIT) 






016:07:25:50 


016:07:46:42 


F/S OCC 


(DARK 


EARTH) 


03123:06: 06 01 


016:07:30:12 


016:09:00:13 


SAA 02 


(EXIT) 






016:07:32:39 


016:08:58:48 


SAA 3 


(EXIT) 






016:07:34:18 


016:08:56:31 


SAA 07 


(EXIT) 






.016:07:45:23 


016:08:56:31 


SAA 05 


(EXIT) 







07:45:23 
07:58:31 
07:58: 47 
08:02:26 
08:04:07 
06:04:07 
06:04:17 
06:08:20 
08:10:50 
08:11:03 



016:08:12:39 



016:08:12:39 
016:08:13:19 



016:07:56:47 
016:08:04:07 
016:08:04:07 

016:06:04:17 
016:08:54:04 
016:08:11:03 
016:06:10:50 
016:06:11:00 
016:06:12:39 



016:08: 13:19 
016:08:16:59 



Sci FOS YRD0_3 7 O3123-0G6: 

F/S_AVD (EXT,L= 20.4) 03123: 

REAQ/N(FGS , E= 43,COARSE2 ) 03123-0G6: 
NODE Crossing 3971 
SAM (ANG= 0.0000,ROLL= 0.14) 

F/S_AVD (EXT.L= 31.2) 03123: 

Sci FOS YRD4_3 8 03123-OG6: 

UP HRS DET2 READY2 01408-OEN: 

UP HRS DET2 0PER2 01408-OEN: 

Sci FOS 9 03123-OG6: 



06:07 01 
06:06 01 
06:01 



13:01 01 
13:01 01 
06:09 01 



End SU 0312306 



SI DCfUN FOS 
SI DOWN FOS 



RED 
RED 



LVONA 
HOLD 



03123-OG6:06:09 01 
03123-006:06:09 01 



. >3MJTim-1991.0l(i.01.D9.0Q.eOO 



0G60500M 
0G60500Q 
OG60500R 

0G60500S 
OG60500T 

OG60500U 
0G60500V 

OG60500W 

0G60500Z 

0G605010 
0G6050II 
OG605012 
OG605013 

0G605014 
0G605015 
00605016 



Science 

Activities 

D 



^ 



/ . o: 



GROUP. PYUDLOAD.YSTEP^l (STAR) , YSTEP_ 
TIME=(ORB, 3966,EASCNCR, 01 HSflMl 9. 000 
SMSTIME-=1991 .016:01:53:09. 00 
BEGINNING AfiV COMMAND BLOCK YSRDY 
GROUP, PYACQMOO,ACQH_ADD (DOUBLE) . ACQt 

ACQM_SYN(NSlfNCSRT) , ACQHTA (SCI ) , ACQH lHlj 1 Wl lwe.lH i j> 
ACOMTRY (REJECT) ,TIME=(ORfl, 3966, EASCNCR, 01H54M23 . OOOS) 
SMSTIME=1991.016:01:53: 13.000 

GROUP, PYOVRLIT,OVR_LITE(PROTEC) . TIME= (ORB. 3966, EASCNCR 
01H54M24.000S) 

SMSriME=1991 .016:01:53: 14. 000 

GROUP, PYEFILL,TIME=(ORfl, 3966. EASCNCR, 01 HS4M27 . OOOS J 
SMSTIME=1991.016:01:53:17.000 
BEGINNING AtV COMMAND BLOCK YSENTRP 
RTSCTRL.FI;NC(ACT) ,RTSID(YSEP0093) , TIME= (ORB, 3966, EASCNCR; ; 
01H54M28.000S) 

SHSTIME=1991.016:01:53:18.000 
BEGINNING A4V COMMAND BLOCK YSUDL 

GROUP, PYIFUP.TIME= (ORB, 3966. EASCNCR, 01H57M09. OOOS) 
SMSTIME=1991.016:01:55;59.000 

GROUP, PYUSEFM2,TIME= (ORB, 3966, EASCNCR, 01H57M43 . OOOS) 
SHSTIHE=1 991. 016: 01:56: 33.00 

GROUP.PYSD EN, DMP TYPE (AUTO) ,TIME=(ORB, 3966, EASCNCR 
01H57M4'1.000S) 

SMSTIME=1991.016:01:56: 34.000 
BEGINNING AAV COMMAND BLOCK YSCOL 

GBOUP,PYEFILL, TIME= (ORB, 3966, EASCNCR, 01 H57M47. OOOS) 
SMSTIME=1991.016:01:56: 37.000 

GROUP, PYTDFLCK.YTDFLCK (RESET) , TIME=(ORB, 3966, EASCNCR 
01H57M48.000S) 

SHSTIME=1991. 016: 01:56: 38.00 

GROUP, PYIFUP,TIME= (ORB, 3966, EASCNCR, 01H57M49 . OOOS) 
MSTIME^1991 .016:01:56: 39. 000 



Science 

Activities 

E 



ITSCTRL, FUNG (ACT) , RTSID (YTDF093) ,TI 
01H57M50.000S) 

SMSTIME=1991.016:01:56: 4 0.00 
GROUP, PYTDFLCK.YTDFLCK (SET) ,TIKE=(C 
01H57M55.000S) 

SMSTIME=1991.016:01:56: 45.000 
GROUP, PYIFDOWN, TIME= (ORB, 3966, EASCNCR, 01H59M5e . OOOS) 
SMSTIME=1991.016:01:58; 4 8.000 

GROUP, PYSTPDMP,TIME= (ORB, 3966, EASCNCR, 01H59M59 . OOOS) 
SMSTIME=1991.016:01:58: 4 9.00 

GROUP, PYIFUP,TIME= (ORB, 3966, EASCNCR. 02HOOM02 . OOOS) 
SMSTIME=1991.016:01:58:52.000 

GROUP, PYDMPSHP,TIHE= (ORB, 3966, EASCNCR, 02HOOM03 . OOOS) 
SMSTIME=1991.016:01:58:53.000 

GROUP. PYSDEN, DMPTYPE (AUTO) ,TIME=(ORB, 3966. EASCNCR 
02H00M04. OOOS) 

SMSTIHE=1991.016:01:5B:54.000 
BEGINNING A4V COMMAND BLOCK YSENTRP 

GROUP, PYENPORT, SHUTTER (CLOSE) , TIHE=(ORB, 3966. EASCNCR 
02H0OMO5. OOOS) 

SMSTIME=1 991 . 01 6 : 01 : SB : 55 . 000 

GROUP, PCPDSLOL. TIME= (ORB. 3966. EASCNCR, 02H16M28 .OOOS) 
SMSTIME=1991.016:02: 15: 18.000 



i 



I sieving [ ^ 



/fSLEW,APER_EID(YRD4_3) , APEB_SID (YRD4_3) 
f ,END_DEC(8.5291623718ie03E+01) 

.ENDPA (1.121456604 00 3906E+02) ,END_RA(6. 953172341622021) 

.STRT_DEC(8.529162 3718ie03E+01) 

.STRT_PA(1 .1214 56604003906E+02) 

, STBT_RA (6. 953172341622021). TYPE (3) ,STABT= (ORB, 3966 
. , EASCNCR, 02H16M3e. OOOS) 
\SMSTIME=1991.016:02:15:28.000 



:GSACQ,CPNAME(PCPREACQ) ,NOSLEW, END=(ORB, GS RS-ACQ 
, 01H38M17. OOOS) .START- (ORB. 3967 , EASCNCR. L . ■ ■ »■■ ■ w . ■, ^-^-^^ ^ 
;SMSTIME=1991.016:0 3:08: 30.00 

: GROUP, PCPDSLOL, TIME= (ORB, 3967, EASCNCR, 01H3eM07 . OOOS) 
;SMSTIME=1991. 016: 03: 13: 40.000 



CALENDAR INFORMATION 
Slews and FHST Updates 

A basic attribute of HST is its ability to slew to any position in a reasonable amount of time (approximately 6 
degrees of arc per minute of time). Given the target coordinates, an eigenslew is calculated, allowing the 
spacecraft to maneuver in all three axes (pitch, yaw, and roll) simultaneously. In the example shown here, the 
slew angle is zero, meaning that the previous pointing was essentially identical to the new pointing. The 
position angle for nominal roll (defined as the sun lying in the half-plane given by the +V3 axis) is determined 
along with the solar avoidance angle. 

The coarsest pointing control mechanism is the Fixed Head Star Trackers (FHST). These are wide field imaging 
devices placed off the main optical axis of the telescope. The FHSTs are used to update the HST position 
uncertainty after maneuvers. This updating provides sufficient pointing accuracy such that the guide star 
acquisition to follow will have a high probability of succeeding. 

A very small slew which can occur within an Obset but between alignments is the Small Angle Manuever 
(SAM). These usually involve a small offset to place the target in a particular aperture of the SI or at a specific 
position within an aperture. 

Science Instrument Transitions 

The calendar building software determines the appropriate time for the science instruments (SI) to transition to a 
higher state before the exposures occur and down to a lower state after observation completion. Frequently, 
however, the Sis will not transition down completely if succeeding observations using the same instrument 



283 



follow closely in time. This can result in groups of SUs scheduling differently than each one separately. Also, 
other Sis for other proposals can transition during this SU if the observation permits. 

Target Visibility 

All target and geometrical information is incorporated at calendar building time. For a given target, the bright 
and dark limb avoidance angles are calculated, thereby determining when the Fine Guidance Sensors (FGS) can 
begin a guide star acquisition. The duration of target visibility (a function of the day on which the observation is 
scheduled), can be determined by when the target enters occultation. The calendar also shows Earth shadow 
crossings and node crossings. 

PCS Acquisition 

For all observations, a Pointing Control System acquisition (PCS ACQ) occurs before data taking begins. The 
acquisition can be done in GYRO mode (as is done for internals) or with the FGS (as is usually done for external 
observations). The calendar allots sufficient time for the PCS ACQ depending on the scenario: coarse track or 
fine lock and whether it is a one or two pair acquisition (cf. Figures 3-5). The expected pointing error in 
arcseconds is calculated knowing the previous pointing , the length of the slew to the new position, estimated 
drift, and time since the last update. 

After an observation has been interrupted due to Earth occultation or South Atlantic Anomaly (SAA) passage a 
re-acquisition (REACQ) occurs. The REACQ acquires the same set of guide stars used in the previous 
acquisition. Since the pointing has not changed, the probability of a successful acquisition is quite high and the 
time allocated for the REACQ is much less than that for the initial PCS ACQ. 

TDRS Contact 

From ephemeris data processed in SPSS, the visibility of the Tracking and Data Relay Satellite (TDRS) is 
calculated. Important information regarding the TDRS include: which satelUte (East or West), the type of service 
(multiple access [MA] or single service access [SSA]), and whether it is a forward or return link. When 
determined by the proposal, communication contacts (COMCONs) are established with the TDRS for either 
uplinks or downlinks. COMCONs are established at the alignment level, therefore requiring the alignment time 
to be sufficient for the COMCON (and any other activities). 

Science Alignment 

The designation "MF Sci" on the calendar refers to the Main Fixed Science alignment within an obset. Within 
this time span all exposures under this alignment must occur (unless the alignment is interruptible, in which 
case it may stop and resume at a later time). In the example shown, a target acquisition is also being executed. 
Obviously, if the alignment is pointed, the target must be visible for the duration of the alignment 

SAA Passages 

A major constraint in scheduling science instrument activities is the passage of HST through the SAA. Because 
the different Sis respond differenUy to the varying radiation intensities within the SAA, several models of the 
SAA are used (cf. Figure 1). A passage through one of the larger models (e.g.. Model 05) can last as much as 30 
minutes (cf. Figure 2). During the SAA passage, instruments such as the WFPC should not take data due to the 
higher background noise. Another major scheduUng constraint depends on the geometrical interaction of the 
SAA with the target visibility windows. If the two occur at the same time the desired observations may not be 
possible for many hours until HST's orbit does not intersect the SAA at all. This non-intersection happens 
about once a day and can last as long as 9 to 10 hours. Many observations that are SAA sensitive are targeted for 
these times. 



SMS TNFORMATTON 

SI Transition 

The statement at time 016:01:20:39 indicating the FOS transitioning from low to high voltage (HV) is expanded 
into the SMS block shown in the figure. Included are various command groups (essentially command 
subroutines with their arguments), comment blocks, and absolute and orbit relative times. The commanding 
following the Text Block starts with the HV turn on by ramping up the HV to 20 KV in 2 KV steps. The next 
step is the setting of the focus followed by the final adjustment of the HV value. The last commanding in this 
block is the calibration of the Digital to Analog converter. 



284 




Figure 1 

Plot of four South Atlantic Anomaly (SAA) models and the orbit of HST 

over an eight hour time span. Model 2 is relevant to the FGS, Model 3 for the FOC, 

Model 5 for the FOS, HSP, WFPC, and FGS for astrometry, and Model 7 for the 

GHRS. 



1991.014 



02 



1991.014:12 



1991.015 



03 



07 



05 





1 


1 1 


1 1 1 








1 1 


1 1 


1 1 1 


1 






1 1 


1 1 


1 1 1 


1 






1 1 1 


1 1 


1 1 1 


1 1 

















Figure 2 

Duration of HST passage through four SAA models over the 

the coiu'se of one day. 



285 




SPACE TELESCOPE 
FIELD OF VIEW 



NON TRACKING VWPA 

30 RADIUS 

7. 13 RIGHT ASC 

85 3SDfaiNATION 

0312300 01 111 AUGN 



CURSOfl 

00 RIGHT ASC 

00 DECLINATION 

00 V2 COORD 
000 V3 COORD 



EPOCH 
lOfll 01SD4 4B U 

BEGIN TIME 
1W101S04 4B4? 

END TIME 
ItNi 01704 40 42 



■ TARGETS 

□ TARGET RE F FONTS 

• GUIDE STARS 

. BRIGHT OBJECTS 



Figure 3 

Display of the stars from the 
Guide Star Catalog for the 
pointing of the Obset 03123:06. 




SPACE TELESCOPE 

FIELD OF VIEW 



NON TRACKING Wff A 

30 RADIUS 

7 13 RIGHT ASC 

85.35 DECLWATOH 

03123 00 01 Di ALIGN 



CURSOR 
00 RIGHT ASC 
00 DECUNATON 
00 V? COORD 
0.00 V3 COORD 



EPOCH 
1001 0101)4 40 43 

BEGIN TIME 
1M1 0150*40 42 

END TIME 
1001 01704 40 42 



• TARGETS 

a TARGET REFPOMTS 

• Gmoe STARS 

. BRIGHT OaiECTS__ 



Figure 4 

Guide Stars returned from GSSS 
for the Guide Star Reques t sent 
for Obset 03123:06. GSSS returns 
stars which have a high probability 
of acquisition (e.g.. no close binaries, 
stars of a limited magnitude range, 
etc.). However, only a small fraction 
(- 20 %) pass further criteria imposed 
by the scheduling software. 




SPACE TELESCOPE 
FIELD OF VIEW 



NON TRACKING W/PA 

30 RADIUS 

7 13 RIGHT ASC 

a5 35DECLINAIION 

03123 00 01 01 ALK^ 



CURSOR 
00 RGHT ASC 
00 DECLINATION 
000 V2 COORD 
00 V3 COORD 



EPOCH 
100101004 40 42 

BEGIN TIME 
1001 01504 40 42 

END TIME 
ie«l 017O4 49 42 



■ TARGETS 

a TARGETREFPONTS 

• GUIDE STARS 

. BRIOm OBJECTS 



Figure 5 

FOV display showing the Guide 

Star pair selected for the use in 

the acquisition for the Obset 

03123:06. 



286 



Guide Star Acquisition 

The guide star acquisition occuring at 016:01:31:46 becomes this rather lengthy block on the SMS. All 
necessary information regarding the acquisition is pulled from the PMDB and displayed as parameters and 
arguments. Included are the guide star pair identifications taken from the Guide Star Selection System (GSSS) 
catalog, coordinates of the stars in right ascension and declination, magnitudes, which star is the dominant (star 
that controls pitch and yaw), which is the sub- dominant (controls the roll), and other information controlling the 
actual guide star acquistion. Also included is the acquistion type (baseline, two step, and list), the Guide Star 
search radius values (upper search radius limit), the fillers to be used by each of the FGS and the guide control 
scenario (coarse or fine lock). 

Science Activities - A 

The science activity for the alignment 03123:05:01 begins thirty minutes after the SI transition block. This 
warmup time is standard for the FOS. The Text Block describes the values that will be set in the following 
commands and is equivalent to line 1 of the Exposure Log Sheet. The next activity is the setting of the Filter 
Grating Wheel (FGW) which is followed by a format code setup used by Post Observation Data Processing 
System (PODPS) for picture processing. 

COMCON 

The COMCON on the calendar at 016:01:52:39 now becomes the block shown on the SMS. The data rate is 
specified (here, 4 kb per second), the type of service (multiple access, return link), whether a science tape recorder 
backup will be done, and start and end times for the tape use. 

Science Activities - B 

The science activities continue after the COMCON with the final setting of the FGW. The following commands 
then sets the OVERLIGHT LIMIT value which is used to safe the instrument if the count reaches this value. 

Science Activities - C 

The next activity is setting up the magnetic deflection. This process determines where the spectrum is located on 
the photocathode and also sets the scale of the spectrum on the photocathode. The remaining part of the block 
sets up the pattern that is desired on the photocathode. OVERSCAN shifts the spectrum 1 diode to compensate 
for dead diodes. CHNNLl and CHNNLS determine the first diode and the number of diodes to use for imaging 
the spectrum. READOUT sets the number of readouts of the pattern but does not clear the diodes. CLEARS is 
the same as READOUT except it clears the diodes after the readout. 

Science Activities - D 

The beginning of this block is a setup for PODPS. It tells them what to expect in the frame (sky, star, 
background). This is followed by a formating of the FOS and the turn on of the OVERLIGHT protection which 
was set earUer. After this has completed a command is given to read out the engineering data. The last group of 
this block is the commanding to open the entrance port followed by another readout of the engineering 
data. 

Science Activities - E 

This block begins with the command PYTDFLCK which sets up the Take Data Flag (TDF) management 
governing when taking data is allowed. This is then followed by the SET parameter which is where data take 
begins. The remainder of the block is a cleanup after the observation. For the FOS this is known as the "Fire 
Break" which insures that the observation stops when it is supposed to. 

Slewing 

The type 3 slew block, a small angle maneuver at 016:03:13:50, is shown. The start and end aperture ids are 
given, along with starting and ending coordinates and position angle. Orbit relative time is given in addition to 
absolute time. 

GS Re-ACQ 

The guide star re-acquisition occuring at 016:03: 13:50 on the calendar is expanded on the SMS. Much less 
information is needed for a re-acquisition than for the initial acquisition due to the fact that the same guide stars 
are being used. 



287 



THE SCHEDULING EFFICIENCY FOR THE HUBBLE SPACE TELESCOPE 
DURING THE FIRST YEAR OF OPERATION 

E.V. BELL. II ^■^, K.E. REINHARD^ , and H.H. LANNING^ 
Introduction 

Prior to the launch of the Hubble Space Telescope (HST), estimates were made as to the ability of the Science 
Operations Ground System (SOGS) to schedule observations efficiently. These estimates ranged from the 
extremely pessimistic (0%), for those who thought SOGS incapable of the task, to optimistic values around 
35%. These latter estimates were based on several factors including the ability of HST to see the Tracking and 
Data Relay Satellites (TDRS), the penetration of HST's orbit into the South Atlantic Anomaly (SAA), target 
visibility, etc. HST completed the Orbital Verification (OV) phase of the mission in November 1990 and is 
currently in the Science Verification (SV) portion. Although the observations made during these early phases are 
not, in general, representative of the majority of the mission, they are indicative of the scheduling software's 
ability to cope with many of the extreme cases likely to be seen during the mission. This paper presents the 
results of the first year of scheduling observations on HST. 

General Overview of Scheduling Efficiency Since Launch 

Shown in Figure 1 is the total scheduling efficiency since launch (shown as the (percent) fraction of time during 
which the spacecraft performed some activity relative to the total time span of a given calendar). These activities 
include the total alignment time for each separate observation as well as numerous overheads (e.g., guide star 
acquisitions, target acquisitions, FHST updates, etc.). The efficiencies of the calendars generated by SPSS in this 
period range from a low value of 13% to a high of 62% with the majority of calendars falling in the 37-47% 
range. There are several features about these efficiency plots which need to be pointed out at this time. First, 
very early scheduling operations for HST were different than has been true more recently, both in terms of the 
type of proposals being scheduled as well as the time spans of the calendars. Characteristic of the first 20 days or 
so of scheduling are short interruptions between calendars created at STScl. These calendars contained planned 
gaps during which teams of engineers and scientists would analyze data and then upload new information to 
thespacecrafL During these gaps, spacecraft attitude was maintained by "Health and Safety" SMSs (Science 
Mission Specifications) generated by the Payload Operations Control Center (POCC) at the Goddard Space Flight 
Center. Second, there are visible gaps in between calendars which are of a longer duration (on the order of 1-3 
days). These are spacecraft safing events, times during which the spacecraft placed itself in a mode wherein it 
could not be damaged. Note that most of the safing events which have so far occurred happened within the first 
four months of operations, although two other events have happened fairly recently. Also, note that these do not 
include individual science insuiiment (SI) safmg events. SI safing events do not affect the overall functioning of 
the spacecraft, just the ability to perform observations with that instrument. Because the instrument is only 
recovered once the safing event has been fully analyzed and at such a time that the recovery can be performed 
without disrupting the operation of all the other instruments, these SI safing events are not visible on these 
plots. Third, early calendars generated by SPSS were of short duration (-18 hours to several days), whereas the 
current schedule (one expected to last for the duration of the mission) is to produce calendars covering seven days 
and running from Sunday midnight to Sunday midnight. This is driven not due to limitations in the ground 
support software, but because of the scheduling time periods for TDRSS. Finally, we wish to address how the 
actual observing timeline has reflected the timeline planned prior to launch. 

Originally, the OV period consisted of two equal portions, covering a four week period, during which the 
Marshall Space Flight Center was to have control of the spacecraft in the first half and Goddard Space Flight 
Center was to have control in the latter half. This orbital verification phase was intended to be used to check out 
the general health of all onboard support systems (the batteries, solar panels, on-board attitude control, etc.) as 
well as an initial checkout of the general health of the Sis. This period was to be followed by an eight month 
period of science verification (SV) during which the various operating modes of the Sis would be checked out. 
Shown above Figure 1 (and all subsequent plots) is the planned duration of OV and SV as well as the original 
planned start of the GO/GTO (General Observer/Guaranteed Time Observer) program. Also shown is the actual 
duration of OV and the beginning of SV. Although OV officially ended 202 days after launch (on Nov. 12, 
1990), some portions of OV were still being executed until very recently. Likewise, several SV proposals began 
to be executed some 60 days or more prior to the official beginning of SV (as the initial OV checkouts of some 
instruments were completed before others). The current timeline plans for the end of SV to be sometime late in 
1991. There are several reasons for the extension of both the OV and SV phases of the mission. First, several 

1 with ST Systems Corporation / National Space Science Data Center - Goddard Space Right Center 

2 with Computer Sciences Corporation / Space Telescope Science Institute 

288 



weeks of activity were involved in attempting to focus the telescope. Once it was discovered that the aberration 
of the mirror was to blame for the inability to find a single optimum focus, it was necessary to continue to 



AOual |<- 
Flumod I OV |<- 

100 r 



Orbital VeriEcilicxi 



ScienoB Verification 



SdctiCB Verificatioo ^^ 
•+« OTO/GO ■ 




100 200 

Days Since Launch 



Figure 1. 

The total efficiency of calendar 
scheduling performed by the 
Science Planning and Scheduling 
System portion of SOGS. The 
efficiency is plotted as the 
(percent) fraction of time during 
which the spacecraft performed 
any activity relative to the total 
time span of the calendar. 



perform observations with the WF/PC and FOC to determine the amount of aberration so that a point spread 
function could be determined. This greatly expanded the duration of OV. In addition, it was decided to perform 
several observations for early release to the astronomical community and the media in order to provide evidence 
of the capabilities of the telescope in spite of the spherical aberration. 

Time Spent in Various Activities 

Included in Figtire 1 are several different activities which are overhead activities. These are such things as target 
and guide star acquisitions, FHST updates, and slews and settling time. Several of these activities are broken out 
in the following figures. There are three main types of observing activities into which a calendar can be broken. 
These are main fixed observations (these include not only targeted observations but internal and earth limb 
observations), interleavers (activities which do not alter the attitude of the spacecraft and so can be scheduled 
diuing large gaps within a scheduling unit of another observation), and parallel observations. So far, interleavers 
have only been used on a few calendars (fewer than 20) and have accounted for less than 5% of the total time 
span of any given calendar. In addition, the current version of the ground-support software does not support the 
use of parallel observations. Therefore, these are not presented herein. Figure 2 shows the amount of time spent 
on each calendar in main fixed observations. Main fixed observations account for the majority of the total time 
presented in Figure 1. The amount of time since launch spent in main fixed activities (essentially the total of all 
the individual alignment times) has ranged from a low of 12% of the calendar span to a high of 52% with the 
bulk of the calendars ranging from 25- 30%. Little difference is seen in the efficiency between the OV period 



Acuul | < 

Plumed I OV \^ 
100 



i 

o 



O 



Ortrilal Verificatkn 



-*+< 



Sdeixe VerificatioD 



Sdentx Venficaticn «» 
*+« GTO/CX) ■ 



.3 




2C0 

Days Since Launch 



Figure 2. 

The efficiency of main fixed 
observations as scheduled by the 
SciencePlanning and Scheduling 
System portion of SOGS. The 
efficiency is plotted as the 
(percent) fraction of time during 
which the spacecraft jjerformed 
targeted, internal, and earth-limb 
observations relative to the total 
time span of the calendar. 



289 



(where much of the observing timeline was determined by committees of individuals) and the S V period to date 
(where the timeline is "optimized" by the use of the artificial intelligence program SPIKE). Shown in Figures 
3-5 are the percent time spent per calendar in the (re)acquisition of guide stars, the (re)acquisition of targets, and 
in slews and settUng, respectively. Important to note in Figures 3 and 5 are the relatively consistent amount of 



Adu^ L < OrbilBl Verificatian 
Raimd I CW t« Sa 



Science Venficatian — 
>f« GTO/GO ' 




100 200 

Days Since Launch 



Figure 3. 

The amount of time spent in 
guide star (re)acquisition for 
calendar spanning the first year 
of observations by HST. The 
time is plotted as the (percent) 
fraction of time which the 
pointing control system (PCS) 
spent acquiring or reacquiring 
guide stars relative to the total 
time span of the calendar. 



Aoual t ^ Orbilal Vcarificatwo i 

Planned | (fj f < Science Verificaboii 

loor — ' — ' — ' — ' — ' — ' — ' — ' ■" 



GTO/tXD . 



i;P 


80 






a 




o 








^ 


«) 


.a 




3 




ST 




< 


40 






0) 




ai) 




fa 




1- 


20 



200 

Days Since Launch 



Figure 4. 

The amount of time spent in 
target (re)acquisition for 
calendars spanning the first year 
of observations by HST. The 
time is plotted as the (percent) 
fraction of time which the 
pointing control system (PCS) 
spent acquiring or reacquiring 
targets relative to the total time 
span of the calendar. 



Aoual | < Orbital Venficalicn 



■*h 



Planned | OV |<- 
100° 



Science Verification 



Science Veiificaticn •»» 

>|< oro/cx) ' 




100 200 

Days Since Launch 



Figure 5. 

The amount of lime spent in 
slewing the spacecraft for 
calendars spanning the first year 
of observations by HST. The 
time is plotted as the (percent) 
fraction of time which the 
spacecraft spent slewing relative 
to the total time span of the 
calendar. 



290 



time spent in performing guide star acquisitions and in slewing the telescope (-5-10% and -5%, respectively). 
Most of the variation of these three activities can be attributed to the amount of time spent pointed at a particular 
target, as well as the number of internal observations scheduled on a given calendar. Figure 4, however, shows 
one of the primary differences between OV and SV activities, that is that much of the early OV observations did 
not require much time for target (re)acquisition. Much of the early timeline involved remaining at a particular 
attitude for extended periods of time, but few targets were actually involved. This was necessary to perform many 
of the checkouts of the pointing control system (PCS) which is involved not only in slewing and guide star 
acquisition, but in maintaining attitude. Target acquisitions have become a larger fraction of the calendar 
activities as the spacecraft spends more time in instrument checkouts and can probably be expected to account for 
5-10% of the calendar time as a norm. Not shown is the amount of time spent performing updates of the fixed- 
head star trackers (FHSTs) which accounts for less than 3% of the total calendar time. 

The Impact of TDRS Contact Time on Efficiency 

Another quantity which has a potential effect on scheduling efficiency is the amount of time required by a 
calendar to be in contact with one of the two Tracking and Data Relay Satellites (TDRS). The Tracking and Data 
Relay Satellite System (TDRSS) is the sole means of communicating with the spacecraft. The amount of time 
required for an observation and the availability of time on a given TDRS, as well as the ability of HST to see the 
TDRS, can have a significant impact on it's schedulability. SPSS can take all of this into account, but all final 
resolutions of confiicting requests for TDRSS time must be resolved between POCC, STScI, and the NCC 
(Network Control Center). The NCC is responsible for scheduling time on TDRSS. Shown in Figure 6 is the 
percent TDRSS request time for HST to date. Note that this is not the amount of time given to HST during 
final conflict negotiations nor does it lake into account the additional time for monitoring the spacecraft which is 
requested by POCC nor the time allocated on an emergency basis during spacecraft safemode events. This time 
does include the total amount of time required by the calendar for upUnks requested by the institute (for spacecraft 
and SI commanding), downlinks requested for real time activities (e.g., interactive acquisitions), and decision time 
needed by the observer. Only a minor fraction of the total calendar time («1%) is taken by decision time. Note 
that during early OV, a large fraction of time was needed to perform the necessary spacecraft commanding and to 
obtain data. During the latter portion of OV, however, and into early S V, requested TDRS time has settled down 
to a nearly constant rate of -5%. There is very little correlation, however, between the overall efficiency of a 
given calendar with how much TDRS time is requested (cf. Figures 1 and 6). Shown in Figure 7 is the amount 
of time during which either TDRS east was not available (generally during shuttle missions) or that TDRSS was 



Actual [ ^ Orbilal Venficatim i » [< S SdcDce Veiificaticn — — 
PliiiiEd I OV K Sdrai Verification >+« OTO/OO • 




100 200 

Days Since Launch 



Figure 6. 

Tlie amount of TDRS contact time 
requested by the Science 
Planning and Scheduling system 
to support observations. The 
time is plotted as the (percent) 
fraction of total TDRS contact 
time (for real-time acquisitions, 
decision time, SI commanding, 
and data transmission) relative to 
the total time span of the 
calendar. 



down (for upgrade or maintenance). TDRS east has the largest potential impact because during shuttle missions 
it is used exclusively for communications between the shuttle and mission control. These dead zones were 
calculated as a percentage of the total available TDRS east visibihty during the span of the calendar. Note that 
during each time that a shuttle mission was launched or expected to take place, TDRS east was completely 
unusable for spacecraft like HST. Smaller periods of TDRS unavailability (varying from around 2 to 6 hours) 
are usually the result of maintenance or upgrade of equipment or software at White Sands. These have little 
effect, however, on the overall efficiency of calendars (as can be seen by comparing Figures I and 7). This is 
because many observations requiring TDRS can be satisfied by simply requesting time on the remaining TDRS 



291 



(for which there are no extensive periods of dead time) or by scheduling the observation during an earlier or later 
time for which the TDRS down time is not a problem. This is a quite convenient feature of the manner in which 
the telescope is operated, since launch delays in the shuttle manifest can affect several weeks of TDRS east 
availability. 









Planned 


1 ov h 




Figure 7. 


100 












^ 80 


- 










The amount of TDRS East dead 


1 












time as experienced by HST 


H 60 












through day 220 of operations. 












The dead time is plotted as the 


^ 












(percent) fraction of TDRS East 




- 






p 


'-i 


dead time relative to the total 












TDRS East visibility time during 


t/5 
OS 
Q 20 

H 












the calendar. TDRS East 


- 










unavailability after day 220 are 












not shown in this figure. 




" , 


,h . n, .r" 


, 


1 


1 , , , 







100 200 300 

Days Since Launch 





Data Volumes 

The last topic to be examined is the amount of data so far generated by HST. This has a direct bearing on the 
efficiency of scheduling observations since one can very efficiently schedule activities on a spacecraft, but if no 
useful science activities are being performed or no data are being transmitted to the ground, it isn't a very efficient 
system. Shown in Figure 8 is the expected average daily data volume (in Gbits/day) for each calendar as 
calculated by the SPSS scheduling software. Note that this and the subsequent figure do not represent the actual 
data return of HST, but are expected returns. Very little data were being generated early in the mission. Figure 8 



Orbilal Vorifiabon 



Science VehfkCAticn 



Scioice Vehficatioo >— 
*+« OTOICO ■ 




Figure 8. 

The average daily data volume (in 
Gbits/day) as estimated by the 
Science Planning and Scheduling 
portion of SOGS. 



100 200 

Days Since Launch 



has denoted on it four special observations, the first a "movie" of Saturn's white spot which occupied a major 
portion of one calendar, the second a series of Mars observations, the third several scheduled observations of 
Jupiter, of which all but the first acquisition failed, and lastly some exposures of lo. Most of the early larger 
features were either the result of repeated data takes with the WF/PC and/or FOC to characterize the mirror or to 
support the early release observations (EROs). Shown in Figtire 9 is the anticipated maximum daily data volume 
(to the same scale as Figure 8) for each calendar. Note the large peak for the Saturn movie. Even with this large 
value, however, HST has still not generated the amount of data expected on a daily basis once the observatory is 
fully operational (somewhere in the neighborhood of 6 Gbits/day). 



292 



Actual |<- 



Orbilal VcrtficatiaD 



i| ov K- 



SdcncE VeriScadco 



ScaexKC Vcriflcaticin — — 
>+« GTO/00 ■ 




Figure 9. 

The maximum data volume during 
each calendar (in Gbits/day) as 
estimated by the Science 
Planning and Scheduling portion 
ofSOGS. 



100 200 

Days Since Launch 



Summary 

The scheduling efficiency of SOGS has so far supported the most optimistic estimates made prior to the launch 
of HST, around 30-40%. The overhead for each calendar amounts to some 15-20% necessary for supporting the 
science (guide star and target acquisition), a figure which is unlikely to change much during the course of the 
mission. It is expected that the overall efficiency of these calendars will improve from these values as more 
interleaver activities are available for inclusion and as SOGS is modified to support parallel observations. 
TDRSS availability, although not a major impact on the efficiency of a given calendar, can affect whether or not 
a given proposal will schedule during a particular period, and the TDRS time required by a given proposal can, 
of course, make the difference between an observation which is easy to schedule and one which is impossible. 
In addition, although no evidence currently exists to support the claim that artificial intelligence pre-scheduling 
of observations can improve the efficiency (this may be due to the largely manual effort still required to schedule 
many of these early observations) it may be that this will change as the nature of the proposals being scheduled 
have more to do with the more "normal" GO/GTO observations. 



Mailing address: 

E. V. BeU, II 

Code 933.9 

National Space Science Data Center 

NASA-Goddard Space Flight Center 

Greenbelt, Maryland 20771 

K. E. Reinhard and H. H. Lanning 
Computer Sciences Corporation 
Space Telescope Science Institute 
3700 San Martin Drive 
Baltimore, Maryland 21218 



293 



ROUTINE SCIENCE DATA PROCESSING OF HST OBSERVATIONS 

Daryl A. Swade^ , Sidney B. Parsons^ , Phil Van West^ , Sylvia Baggett^ , 
Mark Kochte^ , Daryl Macomb^ , Al Schultz^ , and Ian Wilson^ 

Computer Sciences Corporation 
3700 San Martin Drive 
Baltimore, MD 21218 

Abstract. All science observations performed by the Hubble Space 
Telescope (HST) will be automatically processed by the Routine Science 
Data Processing (RSDP) pipeline at the Space Telescope Science Institute 
(STScI). Monitoring and maintenance of pipeline activity is the 
responsibility of the Post Observation Data Processing System (PODPS) 
branch. 

1. HST TO PODPS DATA FLOW 

Data from the Hubble Space Telescope are transmitted from the 
spacecraft to White Sands [TDRSS (Tracking and Data Relay Satellite 
System)] ground station by telemetry through a Tracking and Data Relay 
Satellite. From there the data are transmitted to NASA communication 
(NASCOM) at GSFC (Goddard Space Flight Center) by domestic 
communications satellite and to the DCF (Data Capture Facility) at GSFC. 
DCF transmits the packetized data to the PODPS RSDP pipeline at STScI 
via ground links which are maintained at STScI by the Computer 
Operations Branch (COB). 

2. RSDP PIPELINE PROCESSING 

In the absence of any errors, RSDP reception and processing of 
data through Calibration will proceed automatically once data receipt 
has been initiated. Therefore all processing described in this section 
below requires no operator intervention. RSDP pipeline processing of 
HST science observations is supported by a Science Support Schedule from 
the Science Planning and Scheduling System (SPSS) and real-time activity 
and observer comment files from the Observation Support System (OSS). 
If a problem occurs at any step in the pipeline processing the 
observation is sent to "trouble" where the problem can be investigated 
and hopefully repaired by PODPS personnel before the observation is 
reinserted into the pipeline. A schematic representation of PODPS is 
shown in Figure 1. 

2.1 Data Partitioning 

For every observation received the data are partitioned into 
packetized information sets with one packet equal to one VAX record. 
These records are sorted by Packet Format Code into files to form an EDT 
data set with a Standard Header Packet, Unique Data Log, and science 
data. Other informational and trailer files are created by PODPS. 
These files are assigned a rootname which is derived from programmatic 
information in the SHP and follows the convention ipppssoot where i is 
the science instrument (v=HSP, w=WFPC, x=FOC, y=FOS, and z=GHRS), ppp is 
the program id, ss is the obset id, oo is the observation number within 
the obset, and t is the version (T=tape-recorded, R=real-time. . . ) . 

Data Partitioning then performs a time ordered sort of the science 
packets, checks for the correct number of packets received to detect 

^Staff member of the Space Telescope Science Institute 



294 




<1) 



295 



missing packets, and processes the DCF Quality Accounting Capsule to 
compute a total quality weighted error sum for each observation. If 
this error sum exceeds a database specified threshold value all files 
pertaining to the observation are removed from the RSDP pipeline and 
transferred to the appropriate trouble directory. 

2.2 Data Editing 

The output products of Data Partitioning are analyzed by the Data 
Editing process. Among the activities carried out here are insertion of 
PODPS fill data to serve as place holders of missing packet segments and 
detection of DCF fill data and Reed-Solomon corrections. If Data 
Editing is successfully completed an Edited Information Set is created 
and the observation is queued for Generic Conversion. The Edited 
Information Set is retained and archived to the EDT class. 

2.3 Generic Conversion 

The first step in Generic Conversion is Data Evaluation. The 
flags and indicators (F&I) required for the unique specification of the 
parameters which control the reformatting of the packetized data into a 
waivered FITS structure are dredged from the telemetry and compared to 
those specified in the Project Management Data Base (PMDB). If a flags 
and indicators mismatch between the telemetry (actual) and the PMDB 
(predicted) occurs, the telemetry values are used. Invalid F&I as well 
as F&I mismatches contribute towards the total weighted error sum. Once 
again, if the sum exceeds a database specified threshold the observation 
is removed from the RSDP pipeline and the observation's files are 
transferred to trouble. 

If the error sum is less than the threshold value, the packetized 
information set is reformatted into a waivered FITS structure, i.e., 
data and header files, called a Generic Edited Information Set (GEIS). 
This conversion uses bit locations specified in the Project Data Base. 
If any problems are encountered during this process the observation is 
removed from the RSDP pipeline and transferred to trouble. 

Upon successful completion of Generic Conversion any observation 
flagged as requiring calibration is sent through the calibration 
process. 

2.4 Calibration 

The individual HST Science Instrument Teams will supply through 
the Telescope Instrument Branch (TIB) all the information for 
calibration performed by the RSDP pipeline. Pipeline calibration 
consists of instrument specific algorithms such as wavelength 
calibration, flat-fielding, absolute flux, etc. As calibration 
standards change based on knowledge gained through observational 
experience or evolution of instrument performance, PODPS will update 
calibration reference files and tables. In addition, the observer will 
have the capability of recalibrating an observation from the GEIS files 
with Space Telescope Science Data Analysis System (STSDAS) tools. 

At the end of calibration all files produced in Generic Conversion 
and Calibration (uncalibrated GEIS files and calibrated data) are queued 
for archiving to the CAL class. 

3. STANDARD OUTPUT PRODUCTS 

At the end of RSDP pipeline processing PODPS produces either film 
files from which a print is made or a laser plot of the first group in a 
spectral observation. Generated for each science instrument are - 

FOC & WFPC: calibrated images on film 

HRS & FOS: uncalibrated plot of counts vs. channel number 



296 



and calibrated plot of flux vs. wavelength 
HSP: plot of raw counts vs. tine 

In addition to film products and plots PODPS creates a data 
quality report which is distributed along with the hardcopy output and 
is archived as ancillary data (class ASA). The PDQ (PODPS Data Quality) 
file contains the predicted as well as the actual observation 
parameters. Three keyword fields within the PDQ file contain 
information about the usefulness of the observation: 

QUALITY - one word or one phase that describes the overall quality 
QUALCOMl - comments about the usefulness of the observation 
QUALC0M2 - summarized significant OSS comments 

The choice of QUALITY kesrword is based upon the intrinsic merits 
of the observation and geared for the archival user. The standard PODPS 
quality keywords are: 

OK No apparent problems 

NOISY High background, low S/N 

WEAK-SIGNAL No target seen with decent S/N 

(if a targetted observation) 
DATA-DROPOUTS More than ca. 2% missing, or affecting 

probable area of interest 
SATURATED Majority of pixels "overexposed" 
NO-COUNTS Zero-level data 
BLANK-IMAGE No features visible 
POOR Other problems affecting probable 

scientific usefulness 
UNKNOWN Unable to judge usefulness 
NOT-DISPLAYED Undisplayable with current software, not enough 

time during shift, or a calibration exposure. 

4. DATA DISPERSION 

EDT and CAL data sets are archived to the Data Management Facility 
(DMF) optical disk and this data will normally be available to the 
General Observer (GO) from the HST archives within two days after the 
observation is performed. FITS tapes will be made for the GO and these 
tapes along with hardcopy output will be available to the GO from DSOB 
(Data Systems Operations Branch) within five days of the observation. 
Data loses proprietary status after one year at which time it is 
available for use by the astronomical community. 



297 





DATE DUE 


















































































































































CAVLORD 






PRINTEOINU S A. 



WELLESLEY COLLEGE LIBRARY 



3 5002 03 



10 1004 



u 



Astro qQB 500.268 . F57 I991 



■ "^^^ First year of HST 
observations 

Astro qQB 5O0. 268 . F57 1991 



The First year o± HST 
observations 



1/