1995 O1)

Icarus 141, 1–12 (1999) Article ID icar.1999.6159, available online at http://www.idealibrary.com on

Post-Perihelion HST Observations of Comet Hale–Bopp (C/1995 O1) H. A. Weaver and P. D. Feldman Center for Astrophysical Sciences, Department of Physics and Astronomy, Johns Hopkins University, Baltimore, Maryland 21218-2695 E-mail: [email protected]

M. F. A’Hearn Department of Astronomy, University of Maryland, College Park, Maryland 20742

C. Arpigny Institut d’Astrophysique et de G´eophysique, Universit´e de Li`ege, B-4000 Li`ege, Belgium

J. C. Brandt Laboratory for Atmospheric and Space Physics, University of Colorado, Boulder, Colorado 80309

and S. A. Stern Southwest Research Institute, Space Studies Department, Boulder Extension Office, 1050 Walnut Street, Suite 426, Boulder, Colorado 80302 Received October 8, 1998; revised February 25, 1999

by asymmetric ejection of H2 O molecules from the nucleus, but we have not attempted to model this effect. Except for this offset, conventional models for the spatial distribution of OH provide a good match to the data from November 1997. The observed spatial profile for the August 1997 observation is well matched by our model at large cometocentric distances but is considerably flatter near the nucleus (within ∼200 of the continuum peak) than predicted. The available evidence strongly suggests that an optical depth effect, rather than the production of OH from a population of icy grains in the coma, is responsible for the observed flattening of the OH spatial brightness profile. °c 1999 Academic Press Key Words: comets; comets, composition; spectroscopy, comets.

Post-perihelion observations of Comet Hale–Bopp (C/1995 O1) were made with the Hubble Space Telescope (HST) on 27 August 1997 (r == 2.48 AU; ∆ == 2.99 AU), 11 November 1997 (r == 3.38 AU; ∆ == 3.31 AU), and 19 February 1998 (r == 4.44 AU; ∆ == 4.40 AU) using the newly installed Space Telescope Imaging Spectrograph (STIS). The STIS CCD was used to image the comet, and use˚ using the ful spectroscopy was obtained between 2000 and 3190 A G230LB, G230MB, and G230L gratings. The morphology of the images closely resembled that obtained at similar heliocentric distances preperihelion, but the dust production rates may have been slightly lower post-perihelion. We find no evidence for any companions to the nucleus in the STIS images, but fairly bright objects (up to ∼20% of the brightness of the main nucleus) could be easily “hidden” within the strong coma jets. Emissions from OH and CS were observed in the spectra and were used to derive production rates for H2 O and CS2 . As with the dust, the gas production rates appear to be somewhat smaller post-perihelion than preperihelion. The two-dimensional STIS data allowed us to map the spatial distribution of the OH emission with a spatial resolution of ∼0.00 1, and the ∼6-A˚ spectral resolution of the G230MB grating permitted a detailed examination of the OH excitation. The relative intensities of the rotational lines in the OH(0,0) band are fairly well matched by a standard fluorescence excitation model. The OH spatial brightness profile was slightly asymmetric (∼20% brightness differences when comparing two different directions), and the intensity peak was offset by ∼3200 km from the continuum peak for both the August and November 1997 observations. Perhaps this offset may be explained

1. INTRODUCTION

We began systematic monitoring of Comet Hale–Bopp (C/1995 O1) with the Hubble Space Telescope (HST) in September 1995, shortly after its discovery the preceding July. The preperihelion observations ended in October 1996 and are described elsewhere (Weaver et al. 1997). Unfortunately, Hale– Bopp violated the HST solar elongation angle constraint (elongation angles must be ≥50◦ ) between early-November 1996 and late-August 1997, and this precluded making any HST observations during the near-perihelion period (the perihelion date was 1 April 1997) when the cometary activity was highest. The post-perihelion HST program began on 27 August 1997, and 1 0019-1035/99 $30.00 c 1999 by Academic Press Copyright ° All rights of reproduction in any form reserved.

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additional observations were made on 11 November 1997 and 19 February 1998. Table I is a log of all the HST observations of Hale–Bopp, but this paper is primarily devoted to the analysis of the post-perihelion data. Although our program was primarily spectroscopic in nature, we begin with a discussion of the post-perihelion HST images of Hale–Bopp. The high spatial resolution of HST permitted a detailed view of the morphology of the inner coma, which we compared to that observed preperihelion. Photometric analyses of the images were made to estimate the dust production rate and its variation with heliocentric distance (r ). We also examined the images for any evidence of companions to the nucleus. Most of this paper is devoted to our spectroscopic results. We determined molecular production rates from analysis of OH and CS emissions in the ultraviolet (UV) spectra of a Hale–Bopp. We obtained moderate spectral resolution (∼6 A ) data of the OH(0,0) band that allowed us to examine the molecular excitation and compare it to models. We also mapped the OH emission at unprecedented spatial resolution, which allowed us to measure the excitation as a function of cometocentric distance and to search for any evidence of an extended source of H2 O molecules in the coma. 2. OBSERVATIONS

The post-perihelion observations of Comet Hale–Bopp employed the Space Telescope Imaging Spectrograph (STIS), which was installed during the HST second servicing mission in February 1997. STIS uses CCDs to perform imaging and spectroscopy from near-UV through near-infrared wavelengths and multianode-microchannel-array (MAMA) detectors to perform UV imaging and spectroscopy. A variety of gratings and aperture sizes are available for undertaking low-, medium-, and high-resolution spectroscopy (see Kimble et al. 1998 for further details). The August 1997 and November 1997 observations were made before the MAMAs were available, so the CCD was used for both imaging and spectroscopy. We used the G230LB grat˚ ing with the 0.500 × 5200 slit (the spectral slit width was 13 A −1 ˚ and the dispersion was 1.3 A pixel ) to obtain an overview ˚ alspectrum covering the spectral range from 1685 to 3065 A, though the rapidly falling CCD sensitivity at the shortest wavelengths limited the useful range to wavelengths longward of ˚ We used the G230MB grating with the 200 × 5200 ∼2000 A. ˚ and the dispersion was slit (the spectral slit width was 6.0 A −1 ˚ 0.15 A pixel ) to cover the strong OH(0,0) band at sufficient spectral resolution to resolve, at least partially, the individual rotational lines. The February 1998 spectroscopic observations were made using the MAMA with the G230L grating and the 200 × 2500 a slit (the spectral slit width was 132 A and the dispersion was a 1.58 Aa pixel−1 ) to provide a broad spectral range (1570 to 3180 A) with high throughput. The CCD was used only for imaging in February.

For all imaging observations, we used the CCD with the F28X50LP filter, which is a “long-pass” filter having a sharp ˚ and a long-wavelength reshort-wavelength cutoff near 5500 A sponse that follows the intrinsic quantum efficiency of the CCD itself to its limit at ∼1.1 µm. The CCDs have 1024 × 1024 pixels and a plate scale of 0.00 0508 pixel−1 , but the effective field-ofview is 2800 × 5100 as the filter occults part of the field. Our total time allocation for the post-perihelion Hale–Bopp program consisted of eight orbits, two each in August 1997 and November 1997 and four in February 1998. For the August and November observations, we used one orbit to obtain acquisition images (two 10-s exposures in August and two 20-s exposures in November) and G230MB CCD spectra (total exposure time of 28 min in August and 27.4 min in November) and one orbit for confirming images (one 2-s and one 20-s exposure in August; one 4-s and one 40-s exposure in November) and G230LB CCD spectroscopy (total exposure time of 30 min in August and 30.4 min in November). During the February 1998 observations, we obtained two acquisition images (50 s each) and two confirming images (20- and 50-s exposures). Most of the time in February was used to obtain spectra with the G230L grating, and the total exposure time was 103.7 min. The acquisition images were not calibrated and were used solely to verify that the comet was found and that the nucleus was centered in the spectrograph aperture. In all cases, the autonomous acquisition procedure worked perfectly, the nucleus was centered in the aperture to an accuracy of ∼0.00 05, and the tracking also appeared to be accurate to ∼0.00 05. During the February observations, we lost two confirmation images and ∼16 min of spectroscopy due to temporary loss-of-locks of the guide stars that were used to stabilize the tracking. All of the confirmation images and all spectrograph data were calibrated using the standard pipeline processing at the Space Telescope Science Institute (STScI). The spectrograph data were recalibrated in July 1998 using the best available on-orbit calibration files. The absolute sensitivity is thought to be accurate to ∼10–15% for both the imaging and spectrograph data presented here (see the discussions on the STIS absolute calibration at http://www.stsci.edu/instruments/stis/). 3. IMAGING RESULTS

The post-perihelion image morphology was very similar to that observed preperihelion at the same heliocentric distance. The August 1997 STIS images displayed several jet-like structures, giving the coma a “porcupine” appearance. As the comet moved farther from the sun, the number of jets progressively decreased (see Fig. 1). By performing absolute photometry on the HST images, we were able to calculate Afρ, the product of the dust albedo, the dust filling factor, and the radius of the effective circular aperture used during the observations (see A’Hearn et al. 1984). This quantity is proportional to the observed continuum flux and was defined to provide an aperture-independent measure

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FIG. 1. Three post-perihelion STIS images of Comet Hale–Bopp are displayed on a logarithmic intensity scale along the top row. The intensities are normalized to the same peak value and are displayed using identical stretches in each case. Division of each image by its azimuthally averaged version enhances asymmetric coma structures; the images produced in this way are displayed along the bottom row, again normalized to the same peak intensity and using a common stretch. Each image is 200 × 200 pixels (1000 × 1000 ) in size, and the horizontal bars correspond to 5000 km at the comet. Celestial north is straight up and east is to the left. See the text for further discussion.

of the dust production rate. For dust having an average radius a (µm), density d (g cm−3 ), geometric albedo A p , and flowing outward from the nucleus with velocity vdust (km s−1 ), the dust mass production rate, Q dust (kg s−1 ), is given by Q dust =

(0.67)advdust (Afρ) Ap

when Afρ is in meters. We calculated Q dust using vdust = 0.13rh−0.5 (Weaver et al. 1997), a = 10, d = 1, and A p = 0.04. Given the large uncertainties in the dust size distribution, density, albedo, and outflow speed, Q dust derived here may be different from the actual dust mass loss rate in Hale–Bopp. As shown in Table I (and displayed in Fig. 10, which is discussed in more detail later), the values of Afρ and Q dust indicate that the dust activity in Hale–Bopp was slightly lower postperihelion (∼30% or so). However, this apparent systematic trend in the HST data could be an artifact because our temporal sampling was very sparse, and short-term temporal variability could have easily disguised any long-term trends. We further note that we quote Afρ and Q dust at ρ = 100 (where ρ is the projected distance from the nucleus), and the values can vary by up to 50%, depending on the specific value chosen for ρ.

The dust-to-gas ratio for Hale–Bopp became systematically smaller with decreasing heliocentric distance, as is common for comets (see A’Hearn et al. 1995). At r = 4.79 AU, log(Afρ/ Q OH ) was −23.2 cm s molec−1 , while the value decreased to −24.2 cm s molec−1 at r = 2.5 AU. Hale–Bopp was clearly an exceptionally dusty comet, as a “typical” value for log(Afρ/ Q OH ) for r ∼ 3 AU is approximately −25.1 cm s molec−1 . As was the case for the preperihelion data, we found that the optical Afρ values (derived from the STIS images; we calculated Landolt R-band magnitudes, which have an effective ˚ are typically 1.7–2.0 times larger than wavelength of ∼6500 A) a the UV values near 2900 A (derived from the STIS spectra). Assuming that the dust albedo increases linearly between the UV and optical regions, a ratio of 1.7–2.0 corresponds to an average a reflectivity gradient of 14–19% per 1000 A. Using spectraa of nine comets covering a wavelength range of 3500 to 6500 A, Jewitt and Meech (1986) a obtained an average reflectivity gradient of ∼13% pera 1000 A, with individual values ranging from 5 to 18% per 1000 A. Thus, the UV to optical reflectivity gradient in Hale– Bopp falls within the range of values observed in the Jewitt and Meech survey, but is slightly redder than their average value. In our analysis of the STIS images, we could not find any evidence for companion nuclei. Deconvolved images taken in

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FIG. 2. Detailed view of the inner coma (25 × 25 pixels, or 1.00 3 × 1.00 3) of Comet Hale–Bopp during the 19 February 1998 observations. The frame to the far left is the raw image displayed using a logarithmic intensity stretch. The middle image was produced by dividing the raw image by the azimuthally averaged raw image, in order to enhance any asymmetric structures and to improve the detectability of any faint companions. The frame to the far right was produced by subtracting the azimuthally averaged image from the raw image. Both the middle and right-hand frames show features just north of the nucleus (whose approximate position is marked by the ×), but we ascribe this to temporal variation rather than to the presence of a companion. The compass shows the orientation for all images, and the horizontal bar is 500 km at the comet. See the text for further discussion.

mid-January 1998 with the 3.6-m ESO telescope in La Silla, Chile using the ADONIS adaptive optics system showed evidence for a companion located roughly 0.00 3 east of the main nucleus (Boehnhardt, personal communication). Given that this separation should be easily resolved in the STIS images (∼6 pixel separation), we performed image processing on the STIS images in an attempt to detect any companions near the nucleus (see Fig. 2). A strong coma jet is seen nearly due north of the nucleus in all images, but there are no obvious companions. By adding model PSFs to the observed February 1998 image, we determined that we could have easily detected a companion 6 pixels east of the nucleus as long as it was brighter than ∼5% of the brightness of the main nucleus. If the companion was located within the strong coma jet, then any companion fainter than ∼20% of the brightness of the nucleus would have escaped detection during our inspection. If a companion was in a bound orbit around the main nucleus, perhaps their relative separations changed in such a way that they could be resolved during the ESO observations but not during the HST observations one month later. However, we also do not see any obvious companions in the August 1997 and November 1997 HST images. We note that Sekanina (1999) claims to have detected companion nuclei during his analysis of the preperihelion HST images of Hale–Bopp, but those results have been questioned (Weaver and Lamy 1999). In any case, the techniques we applied to the post-perihelion images are not capable of detecting a companion as faint and as close to the main nucleus as the companion(s) reported by Sekanina. 4. SPECTRAL IMAGING RESULTS

During the post-perihelion observations, the only molecular emissions detected in the spectra were from OH and CS. The

OH(0,0), (1,1), and (1,0) bands were detected during the August 1997 observations, as was the CS(0,0) band (Figs. 3 and 4). The observed brightness ratio of the OH(0,0) and (1,1) bands was 13.5 ± 3.5, which is consistent with the theoretical value of 17 (Schleicher and A’Hearn 1988; the OH(1,0) band was observed through a different-sized aperture, but its relative brightness also appears to be consistent with the theoretical value). During the November 1997 observations, the OH(0,0) band was definitely detected, the OH(1,0) band was marginally detected, and the CS(0,0) band was not detected (Fig. 5). No molecular emissions were detected during the February 1998 observations (Fig. 6). A strong continuum, due to scattering from dust in the coma, was detected during all spectroscopic observations and could be a accurately measured at wavelengths longward of ∼2000 A. The color of the dust derived from the STIS post-perihelion observations is very similar to that observed during UV observations of other comets (Feldman and A’Hearn 1985) with a reddening of a ˚ between 2600 and 3050 A. The preperihelion HST ∼5%/100 A spectra extended to slightly longer wavelengths and generally indicated that there is a flattening of athe color (toward neutral) for wavelengths longward of ∼3100 A (the data are displayed in Fig. 3 of Feldman (1999)). Neither the pre- nor post-perihelion measurements of the cometary continuum showed any evidence for strong spectral features over the entire wavelength range from ˚ (e.g., there is no graphite feature near 2200 A). ˚ 2000 to 3190 A Following standard convention, we assumed that OH and CS are photolysis products of H2 O and CS2 , respectively. The observed OH(0,0) band brightnesses were converted to apertureaveraged OH column densities assuming optically thin conditions and using “quenched” fluorescence efficiency factors, or g-factors (see further discussion below). The OH column densities were then converted to H2 O production rates (Q H2 O ) at the nucleus using a spherically symmetric “vectorial” density

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struction produced OH, the OH was ejected isotropically in the frame of the dissociating H2 O molecule with a velocity of 1.05 km s−1 , and that the OH lifetime was 1.6 × 105r 2 s. Column densities for CS were calculated assuming that the CS emission was optically thin and using a g-factor of 7 × 10−4 photons s−1 molec−1 (Jackson et al. 1982). We then used a standard Haser model (Haser 1957) to calculate the CS2 production rates (Q CS2 ), assuming that a CS molecule was produced for

FIG. 3. The top panel is a STIS spectral image taken on 27 August 1997 using the CCD with the G230MB grating. The slit dimensions are 200 × 5200 (spectral by spatial dimensions). The two dark horizontal streaks at ∼1200 above and below the continuum emission are due to support structures that stabilize the slits. The bottom plot is a spectrum extracted from the top image, using a 200 high effective aperture centered on the continuum emission. We fit a solar spectrum to the continuum and subtracted it from the data to produce the emission spectrum shown here. Both the OH(0,0) and OH(1,1) bands were detected.

model (Festou 1981, Budzien et al. 1994). We assumed that the H2 O outflow velocity was 0.8r −0.5 km s−1 (for the analysis of International Ultraviolet Explorer (IUE) data, we assumed a slightly larger H2 O outflow velocity of 0.85r −0.5 km s−1 because the IUE aperture was significantly larger than the apertures used for the HST observations and the effective outflow velocity is thought to increase with increasing aperture size), the total H2 O lifetime was 8.2 × 104r 2 s, 84% of the H2 O de-

FIG. 4. The top panel is a STIS spectral image taken on 27 August 1997 using the CCD with the G230LB grating. The slit dimensions are 0.00 5 × 5200 (spectral by spatial dimensions). The bottom plot is a spectrum extracted from the top image, using a 200 high effective aperture centered on the continuum emission. We fit a solar spectrum to the continuum and subtracted it from the data to produce the emission spectrum shown here. Weak emissions from the CS(0,0) and OH(1,0) bands were detected.

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(O P12 [5/2]) have observed intensities that are significantly different from the model prediction. We extracted spectra at various positions along the slit (i.e., at various spatial positions) and find no evidence for any excitation changes with cometocentric distance. Unfortunately, the quenched and unquenched models (“quenched” means that the populations of the two 3-doubled levels of the lowest rotational level are equally populated due to collisions) are nearly identical at the heliocentric radial velocities of the post-perihelion observations, so we cannot verify that the quenched model is the appropriate one to use in

FIG. 5. The top plot shows the G230MB spectrum from 11 November 1997 using an effective aperture of 200 × 200 centered on the continuum peak. The continuum was subtracted to create this emission spectrum, which shows the OH(0,0) band. The bottom plot shows the G230LB spectrum from the same date using an effective aperture of 0.00 5 × 200 centered on the continuum peak. The continuum was subtracted to create this emission spectrum, but no significant spectral features are detected.

every CS2 molecule that was destroyed, the CS2 and CS outflow velocities were collisionally equilibrated at 0.8r −0.5 km s−1 , the CS2 lifetime was 500r 2 s (Huebner et al. (1992) calculates that the lifetime of CS2 lies between 315 and 345 s at 1 AU, depending on the solar activity, but we have adopted a slightly larger value for consistency with our group’s previous work on other comets observed by HST and IUE), and the CS lifetime was 1 × 105r 2 s (Jackson et al. 1982). The OH rotational excitation is fairly well matched by a standard fluorescence equilibrium model that takes into account the strong heliocentric radial velocity dependence of the excitation (Schleicher and A’Hearn 1988; see Fig. 7). Only the two weak, ˚ ˚ (R Q21 [3/2]) and 3106.95 A satellite transitions at 3072.95 A

FIG. 6. The top plot shows the G230L spectrum from 19 February 1998 using an effective aperture of 200 × 200 centered on the continuum peak. Unlike the other spectra shown in this paper, the continuum has not been removed in this case. Also plotted is the scaled solar spectrum that we used to estimate the continuum contribution (Woods et al. 1996). The latter was subtracted from the data to produce the emission spectrum plotted in the lower panel. The apparent a features near 2800 A are probably due to a mismatch between the actual and modeled solar Mg II line. Faint emission is detected in the vicinity of the CS(0,0) a band centered near 2576 A (the region between the two vertical lines) but this is probably due to an incorrect continuum subtraction. If present, OH(0,0) band emission would appear between the two dashed vertical lines.

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and gaseous emissions have been observed for other cometary radicals (e.g., CN, C2 , and NH2 ), and that an explanation was proposed based on the different dynamics of dust and gas (see Delsemme and Combi 1983). However, this explanation rested on the premise that the peak in the gaseous emissions, rather than the continuum peak, defined the position of the nucleus, which

FIG. 7. This plot compares the observed OH(0,0) band spectrum to that predicted by the fluorescence excitation model model of Schleicher and A’Hearn (1988) for the “unquenched” case at r˙ = +24 km s−1 , which is close to the velocity of the comet. (We do not have access to a model for the exact velocity of Hale–Bopp.) The unquenched and quenched models are essentially identical at this r˙ . Most of the OH emission is well explained by the model.

interpreting the HST observations of Hale–Bopp. (As discussed in Weaver et al. (1997), the unquenched g-factors were about 1.5 times smaller than the quenched values for the preperihelion observations.) The spatial brightness profile of OH has some interesting properties (see Figs. 8 and 9). The OH brightness distribution is clearly asymmetric, but the magnitude of the asymmetry is only ∼20% in the two directions sampled by the HST observations. The peak of the OH column brightness is offset by ∼3200 km from the peak of the continuum brightness. As the solar phase angle was small during the HST observations (e.g., φ = 18◦ on 27 August 1997 and 17◦ on 11 November 1997), the offset cannot be clearly associated with either the sunward or tailward directions. The direction of the offset in the OH brightness peak is in the same hemisphere as the strongest dust jet. (The slit axis was oriented at ∼45◦ to the direction of the strongest dust jet.) Perhaps preferential ejection of H2 O in a particular direction could explain the observed offset, but we have not attemped to model such a scenario. Since the peak in the OH brightness was only ∼1.00 3–1.00 5 from the nucleus (assuming that the peak in the continuum brightness profile is coincident with the position of the nucleus), this offset would be difficult to detect during ground-based observations. The magnitude of the observed offset presumably depends both on the degree of asymmetric outflow from the nucleus and on the spatial resolution employed during the observations. For a fixed degree of asymmetry, we would expect lower resolution observations to yield larger offsets. Indeed, we observed Comet 103P/Hartley 2 with STIS at a geocentric distance of 0.82 AU on 2 January 1998 and found that the peak in the OH spatial profile was offset (tailward this time) by only ∼230 km from the continuum peak (Weaver et al., unpublished). We note that similar offsets between the continuum

FIG. 8. The top plot shows spatial brightness profiles derived from the G230MB spectral image taken on 27 August 1997. The bottom profile is from a region that includes only continuum emission, and the top profile is the spatial profile of the OH(0,0) band after continuum subtraction. The OH profile is flatter than that of the continuum, as is expected for a dissociation product. The peak of the OH emission is offset by ∼3200 km from the peak in the continuum emission, which is at the position of the dashed vertical line. The bottom plot compares the OH spatial brightness profile to that predicted by a model. The peak brightness in the model is centered on the estimated position of the observed OH peak. The model has been convolved with the instrumental point spread function, but still rises toward the nucleus faster than the observed profile. Otherwise the model matches the observed profile fairly well, and we derive that 2.2 × 1029 ≤ Q H2 O ≤ 2.6 × 1029 .The apparent data dropouts near ±1200 from the continuum peak are due to the presence of support structures that stabilize the slits.

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FIG. 9. The top plot shows the OH(0,0) band spatial brightness profile derived from the G230MB spectral image taken on 11 November 1997. The peak of the OH emission is offset by ∼3200 km from the peak in the continuum emission (which is at the position of the dashed vertical line). The bottom plot compares the OH spatial brightness profile to that predicted by a model. The peak brightness in the model is centered on the estimated position of the observed OH peak. The model has been convolved with the instrumental point spread function and matches the observed profile fairly well. We derive that 6.0 × 1028 ≤ Q H2 O ≤ 7.0 × 1028 .The apparent data dropouts near ±1200 from the continuum peak are due to the presence of support structures that stabilize the slits.

we feel is unlikely. Perhaps further modeling will illuminate the cause of the observed offset in the peak of the OH distribution, but such modeling is beyond the scope of the current paper. For the August 1997 data, the OH brightness distribution in the core (i.e., near the nucleus) is much flatter than predicted by the standard models (e.g., vectorial, Haser, or hydrodynamic, all of which predict essentially the same profile shape within ∼104 km of the nucleus), even after convolution with the STIS point spread function. Two possible explanations for the flat-

tened brightness profile are: (1) an “extended source” for the H2 O molecules that are the ultimate source of the observed OH and (2) optical depth effects. As H2 O ice grains were detected during infrared spectroscopy of Hale–Bopp at heliocentric distances of 6.8 AU (Davies et al. 1997) and 2.8 AU (Lellouch et al. 1998), it seems natural to suggest that sublimation of these grains in the coma could provide the extended source necessary to explain the flattened OH spatial brightness profiles. Furthermore, asymmetric ejection of the icy grains might also explain the observed offset from the nucleus of the peak OH emission. However, several properties of the HST Hale–Bopp data are inconsistent with these hypotheses. If icy grains were responsible for the flattening of the August spatial profile, then the effect should be even more pronounced for the November spatial profile because the lifetime of icy grains at 3.4 AU should be almost 1000 times longer than the lifetime of icy grains at 2.5 AU (see Hanner (1981) for a general discussion of the lifetime of icy grains as a function of grain size, composition, and heliocentric distance). Yet, the OH spatial brightness profile for the November data is matched very well by the standard model (Fig. 9). Similarly, if icy grains were responsible for the offset of the OH peak brightness, then this offset would be much larger in November, contrary to what is observed (see above). So although there appears to be credible evidence that icy grains were present in the coma at the heliocentric distances sampled by the HST observations, we cannot find any manifestation of their presence in the OH spatial brightness profiles. For lines-of-sight passing near the nucleus, the OH column density was ∼1–2 × 1014 cm−2 for the August 1997 observations and about five times smaller for the November 1997 observations. If the optical depth (τ ) of the OH lines was ∼0.3 for the August observations, then that could explain both the saturation observed in the August spatial brightness profile and the lack of saturation in the November data. We now examine the optical depth of the OH lines. Under cometary excitation conditions, there are six absorption lines responsible for pumping the UV OH emission bands, and all of these lines originate from one of the two 3-doubled levels of the J = 3/2, v = 0, 2 53/2 state of OH. The oscillator strengths of these transitions range from 0.48 to 5.8 × 10−4 , with an average value of 3.3 × 10−4 . The optical depth at line center for a Doppler-broadened line is given by π e2 2 f τ= m e c νD

r

ln 2 Nl , π

where e and m e are the charge and mass of the electron, respectively, c is the speed of light, f is the absorption oscillator strength for the transition being considered, νD is the full width at half-maximum (FWHM) of the Doppler-broadened line (in Hz), and Nl is the column density of the lower level for the transition. (The factor πe2 /m e c is numerically equal to 2.64 × 10−2 cm2 Hz.) The line center optical depth is inversely proportional to the linewidth, and our main uncertainty in calculating

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the optical depth relates to the choice of the linewidth. Usually one assumes that the linewidth is equal to twice the outflow velocity of the OH molecules, or ∼2–2.5 km s−1 , which is also close to the observed width of OH radio lines (Bockel´ee-Morvan and G´erard 1984). Assuming that the FWHM (in Hz) of the OH lines is given by νD =

2v ν, c

where v is 1 km s−1 and ν is the frequency of the transition (in Hz), and that half of all the OH molecules are in the absorbing level (since about half of the molecules are in each of the two 3-doubled levels), we estimate that the line center optical depth of the strongest OH line was ∼0.17 for the August 1997 observation, which is about a factor of 2 smaller than needed to explain the observed saturation in the core of the OH spatial brightness profile. However, the linewidth used for this calculation is probably much broader than is appropriate for a pencil beam passing near the nucleus. For this latter case, let us consider the possibility that the absorption is governed by a lineshape characteristic of a gas in thermal equilibrium at a temperature T , for which the FWHM is given by r −7

νD = 7.162 × 10 ν

T , M

where M is the molecular weight of OH in atomic mass units (=17). In order for the strongest OH line to have an optical depth of 0.3 for the August 1997 observation, the effective temperature would have to be ∼460 K, which is well above the expected kinetic temperature of the collisionally dominated region of the coma. (The formula derived by Biver et al. (1997) from a fit to observed molecular rotational temperatures gives a nominal temperature of 36 K at a heliocentric distance of 2.5 AU. Note that a linewidth of 2 km s−1 corresponds to T ∼ 1470 K.) In other words, one does not have to assume an unusually narrow line profile in order to demonstrate that significant optical depths for the OH lines are possible for the Hale–Bopp observing conditions. Therefore, we conclude that optical depths of a few tenths or so could easily be achieved for lines-of-sight passing within a few thousand kilometers of the nucleus of Hale–Bopp in August 1997, and that line saturation is mainly responsible for the observed flattening of the OH spatial brightness profile in the inner coma. Both the offset from the nucleus of the peak in the OH spatial brightness distribution and the saturation of the OH spatial brightness profile in the inner coma can affect the calculated H2 O production rates. For the August 1997 observation, the OH spatial brightness far from the nucleus (see Fig. 8) indicates that Q H2 O is ∼2.4 × 1029 s−1 , which can be compared to a value of ∼2.0 × 1029 s−1 derived from a 200 × 200 aperture centered on the nucleus. Similarly, an analysis of the OH spatial brightness profile for the November 1997 data (see Fig. 9) yields a water production rate of ∼6.5 × 1028 , whereas a value of ∼5.7 × 1028

FIG. 10. Heliocentric variation in the production rates of H2 O, CS2 , and dust derived from HST and IUE observations. Filled symbols are derived from the post-perihelion STIS data discussed in this paper. The statistical errors are generally smaller than the symbol sizes. Downward-pointing arrows are 3σ upper limits.Different species have different heliocentric variations, indicating that no single physical mechanism controls their behavior. H2 O sublimation curves were computed for three cases: an isothermal spherical nucleus (dashed curve), a nucleus whose rotation axis is pointing directly at the Sun (dot-dashed curve), and the local rate at the subsolar point (solid curve). The comet appears to be somewhat less active post-perihelion than preperihelion.

would be derived for a 200 × 200 aperture centered on the nucleus. The preperihelion H2 O production rates were generally derived from small-aperture observations centered on the nucleus, which probably means that they are underestimated by ∼10–15% due to optical saturation and/or brightness offset effects. However, the values of Q H2 O given in Table I and plotted in Fig. 10 do not include any such correction, primarily because we do not have the data required to make an accurate correction. The heliocentric variation in the gas and dust production rates is plotted in Fig. 10. The post-perihelion production rates appear to be somewhat smaller than the preperihelion values for both the gas and dust, with the H2 O production rates showing the largest asymmetry (up to a factor of ∼2). The asymmetry for CS2 and the dust may not be real, as Hale–Bopp often exhibited short-term temporal variability of a magnitude comparable to the observed pre- to post-perihelion differences. However, temporal variability probably cannot explain the observed asymmetry in the H2 O production rate. We have compared our molecular production rates to those derived by others (Fig. 11). CS was observed extensively at radio wavelengths (Biver et al. 1997, 1999), and production rates of CS2 derived from those data are in excellent agreement (within ∼20%) with the UV results when there were nearly simultaneous observations. Generally, the radio and UV values for Q CS2 seem to be consistent with each other, and any apparent discrepancies could be reasonably ascribed to temporal variability. (Owing to CS2 ’s short lifetime of ∼500 s, temporal variability is much more easily observed in the CS emissions than in the OH emissions.) The radio observations made at heliocentric distances smaller than ∼2.5 AU, where HST and IUE could not observe, show

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WEAVER ET AL.

TABLE I Log of HST and IUE Hale–Bopp Observations Q H2 O Q CS2 Q CO2 (1028 ) (1026 ) (10 28 ) (molecules s−1 )

A fρ (m)

Q dust (103 kg s−1 )

— — — ≤0.29 1.2 — 1.0 — 0.98 1.2 2.7 2.2 4.6 3.0

— — — ≤21 — — — — 4.4 3.0 — 8.8

20 ± 5

— — — ≤1.7 2.0 — 2.5 — 6.2 10 14 21 26 27

3.4

1410 1190 1430 640 1060 1410 1170 1280 1050 1210 1520 2120 7290 1950

12.0 10.1 12.4 5.60 10.8 14.5 12.3 13.6 11.7 14.1 18.5 27.3 93.9 26.4

— — —

24 6.5 ≤1.5

2.3 ≤1.4 ≤0.45

— — —

1240 990 890

17.2 11.7 9.2

1

r˙ (km s−1 )

φ (deg)

BOH

6.82 6.80 6.59 6.36 4.79 4.65 4.43 4.35 4.00 3.65 3.31 2.97

6.31 6.33 6.51 6.72 4.80 4.42 3.85 3.68 3.01 2.74 2.79 2.95

−14.9 −14.9 −15.1 −15.4 −17.2 −17.4 −17.7 −17.9 −18.4 −19.0 −19.6 −20.2

7.6 7.9 8.7 8.2 12 12 11 11 4.4 8.2 16 19

2.69

3.04

−20.8

19

— — — ≤6.3 10 ± 2 — 20 ± 5 — 90 ± 14 140 ± 25 110 ± 20 390 ± 30 830 ± 100 740 ± 100

— — — ≤14 — — — — 13 ± 4 11 ± 4 — 41 ± 6

PC, FOS

— — — ≤18 35 ± 5 — 60 ± 15 — 345 ± 40 770 ± 70 1000 ± 100 3600 ± 200 4500 ± 800 7000 ± 700

STIS STIS STIS

2.48 3.38 4.44

2.99 3.31 4.40

+21.2 +19.5 +17.7

18 17 13

14,000 ± 1000 1300 ± 250 ≤130

1000 ± 300 ≤250 ≤30

Date (UT)

r Instrument

31 Aug. 1995 3–4 Sep. 1995 26–27 Sep. 1995 23–25 Oct. 1995 7 April 1996 21 April 1996 13 May 1996 20 May 1996 22–23 June 1996 25–26 July 1996 25 Aug. 1996 23 Sep. 1996

IUE IUE WFC PC, FOS PC, FOS IUE IUE PC PC, FOS PC, FOS IUE PC, FOS

17–18 Oct. 1996 27 Aug. 1997 11 Nov. 1997 19 Feb. 1998

(AU)

BCS (rayleighs)

BCO

Note. WFC and PC refer to the wide-field and planetary modes, respectively, of WFPC2. r˙ is the heliocentric radial velocity of the comet. φ is the phase angle (Sun–comet–Earth angle). B is the observed average column brightness in the spectrograph aperturea (1 rayleigh = 106 photons cm−2 s−1 ) for the OH(0,0) band a a near 3090 A, the CS 1v = 0 band sequence near 2576 A, and the CO(1,0) Cameron band near 1995 A. We used a rectangular FOS aperture whose size was 3.00 66 by 1.00 29, except that a circular aperture of diameter 0.00 86 was used for some observations on 18 October 1996. The FOS aperture was usually nearly centered on the nucleus, but the aperture was offset ∼5–600 east and ∼0.00 5 east for the April 1996 and October 1996 FOS observations, respectively. The IUE observations were made with a rectangular aperture(9.00 3 × 1500 ) that was centered on the nucleus.The STIS OH(0,0) band brightnesses are average values in a 200 × 200 aperture centered on the nucleus, but the H2 O production rates were derived from model fits to the observed OH spatial brightness profile. The STIS CS(0,0) band brightness is the average value in a 0.00 5 × 200 aperture a that was centered on the nucleus. Q is the calculated production rate at the surface of the nucleus. Afρ is proportional to the observed continuum flux at ∼6500 A. Dual values are given for some of the September 1996 entries, and these refer to pre- and post-outburst measurements. The quoted errors are 1σ values, except that 3σ values are quoted for entries in which only upper limits are available.

that Q CS2 was much larger than predicted by the power law derived from the UV data. Thus, either there was a change in the power law for CS2 production for r ≤ 2.5 AU or the power law fit to the UV CS2 production rates is not valid. The radio data do not show any clear evidence for a change in the power law dependence of Q CS2 near r ∼ 2.5 AU. Perhaps it is more likely that the fit to the UV observations is not valid, either because of the limited heliocentric distance covered, or because temporal variability skewed the results. There are several independent sources of data on the H2 O production rate in Hale–Bopp. The Infrared Space Observatory (ISO) made direct, spectroscopic observations of H2 O on 26 September 1996 (r = 2.93 AU), 6 October 1996 (r = 2.82 AU), and 29 December 1997 (r = 3.89 AU) (Crovisier et al. 1997, 1999). After ensuring that consistent model parameters are used, we find that the ISO preperihelion values for Q H2 O appear to be in excellent agreement with those of HST (i.e., the ISO points fall nearly along a curve connecting the HST points; there were no simultaneous observations). The ISO data also seem to confirm the conclusion reached above from the HST results that the preperihelion activity for H2 O was somewhat higher than postperihelion activity levels. The post-perihelion H2 O production rate reported by Stern et al. (1999; 2.6 × 1029 at r = 2.33 AU)

based on UV OH imaging observations made by the SouthWest Ultraviolet Imaging System (SWUIS) also appears to be in excellent agreement with the HST results. Comparison of the HST/IUE H2 O production rates with those derived from radio OH observations (Biver et al. 1999, Colom et al. 1999) and ground-based UV OH photometry (Schleicher et al. 1997; Schleicher has provided us with production rates that have been converted from Haser model to vectorial model values using the empirical relation derived in Cochran and Schleicher (1994)) is more complicated. Both the radio and ground-based UV production rates can sometimes be significantly larger (up to factors of 2 to 3) than the HST/IUE values, while at other times they can be in nearly perfect agreement. The large apparent discrepancies during September and October of 1996 may be related to the large outburst in activity detected by HST on 23 September 1996 and the different response times of the different instruments to temporal variability (i.e., the much larger apertures used by the radio and ground-based observers imply that their production rates are effectively averages over timescales much longer than those for the HST observations). The large apparent discrepancies at the largest heliocentric distances (i.e., for r ≥ 3.5 AU) may also be related to aperture size, if the sublimation by icy grains contributed significantly to the H2 O content of

11

HST OBSERVATIONS OF HALE–BOPP

FIG. 11. (Top) We plot H2 O production rates derived from the HST and IUE UV OH spectroscopic observations (open boxes for preperihelion, filled boxes for post-perihelion), from preperihelion ground-based UV OH photometric observations (open circles; Schleicher, private communication), from radio OH observations (open diamonds for preperihelion, Biver et al. 1997; filled diamonds for post-perihelion, Colom et al. 1999), from the SWUIS UV OH observation (filled circle; Stern et al. 1999), and from ISO infrared H2 O observations (open triangles for preperihelion, Crovisier et al. 1997; filled triangle for post-perihelion, Crovisier et al. 1999). The post-perihelion radio values were derived using the inversion model of Despois et al. (1981) and can be up to three times smaller than values derived using the inversion model of Schleicher and A’Hearn (1988). For all the radio data, we divided the quoted OH production rates by 0.84 to convert them to H2 O production rates. The ISO H2 O production rates were scaled to the same H2 O outflow velocity used for the analysis of the HST and IUE data. The same H2 O sublimation curves that appear in Fig. 10 have been reproduced here as well. (Bottom) We plot CS2 production rates derived from the HST and IUE UV CS spectroscopic observations (open boxes for preperihelion, filled boxes for post-perihelion) and from radio observations of four CS rotational lines (open diamonds for preperihelion, Biver et al. 1997; filled diamonds for post-perihelion, Biver et al. 1999). The heliocentric variation in Q CS2 is ∼r −2.3 as derived from the HST/IUE observations and ∼r −3.5 as derived from the radio observations, which cover a much larger range in heliocentric distance. The statistical errors in the various production rates are generally comparable to or smaller than the symbol sizes.Downward-pointing arrows refer to 3σ upper limits.

the coma. Finally, we point out that different choices for model parameters (in either the density or excitation model) could be the source of some of the apparent discrepancies among the various determinations of Q H2 O from the OH observations, although such differences would tend to change only the absolute

production rates and not their heliocentric variation. (For example, changing vH2 O from 0.8r −0.5 to 1.0r −0.5 km s−1 would increase all values of Q H2 O derived from the HST observations by ∼25%.) Simple H2 O vaporization models (Weaver et al. 1997) were calculated in order to make comparisons with the observed sublimation behavior. Neither the pre- nor post-perihelion data are consistent with sublimation from an isothermal object (see Fig. 10). However, sublimation from a pole-on nucleus (in which the rotation axis points along the direction to the Sun and the entire surface is covered with ice), or a flat surface whose normal is parallel to the solar vector, is in reasonable agreement with the trends displayed by the observations. For the “flat-surface” case, the area of subliming ice is ∼330 km2 , which corresponds to a square region 18 km on a side. For the pole-on case the total area of subliming ice is ∼1330 km2 and the diameter of the nucleus, assuming that all of the observed H2 O is produced from sublimation at the nucleus, must be at least ∼20 km. If the diameter of the nucleus is 70 km (see Weaver and Lamy 1999), then the fractional active area is ∼9%, which is consistent with the range of values observed in other comets (see A’Hearn et al. 1995). For a 40-km-diameter nucleus, the fractional area increases to ∼26%, which is much larger than is typically observed but may be reasonable for a comet that has not made many close passages to the Sun. Alternatively, if the diameter of the nucleus is ≤40 km, the large H2 O production rates could be interpreted as evidence that a significant amount of the H2 O observed in the coma must then be coming from an extended source (e.g., sublimation from icy grains). However, as discussed earlier in this section, the spatial profiles of the OH emission do not show any convincing evidence for an extended source of H2 O, at least for heliocentric distances smaller than 3.4 AU. Notwithstanding the earlier discussions on the power law fit to the Q CS2 values and the possible effects of temporal variability on Q CS2 and Q dust , we still feel that the UV data support the conclusion (Weaver et al. 1997) that H2 O vaporization did not control either the CS2 or dust production in Hale–Bopp for heliocentric distances larger than ∼2.5 AU. The differences in behavior among the species are most apparent at the largest heliocentric distances, where both the CS2 and dust production rates seem not to follow the steep decline observed in Q H2 O between 3 and 4.4 AU. Many other species, in addition to OH and CS, were observed over a large range in heliocentric distance during radio observations of Hale–Bopp (Biver et al. 1997, 1999; see also the review by Bockel´ee-Morvan and Rickman 1999), and those results demonstrate even more dramatically that no single species seems to control the activity level of Hale–Bopp at large heliocentric distances. 5. SUMMARY

Post-perihelion HST observations of Hale–Bopp show activity levels that are somewhat below the preperihelion values. Although inadequate temporal sampling coupled with short-term

12

WEAVER ET AL.

temporal variability could explain this asymmetry for the CS2 and dust production rates, the asymmetry in the H2 O production rate is apparently a real pre- to post-perihelion difference. The excitation of the OH molecule is fairly well described by a standard fluorescence model. The peak in the spatial brightness profile of OH is offset by ∼3200 km from the continuum peak for both the August 1997 and November 1997 data, but this does not seem to be related to the presence of icy grains. The observed depression in the OH spatial brightness profile near the nucleus for the August 1997 data is most likely due to optical depth effects. ACKNOWLEDGMENTS We thank Andy Lubenow, Melissa McGrath, Stefi Baum, and the HST operations staff for their excellent technical support of these observations. We also thank John Debes for recalibrating the STIS spectroscopic data. Comments and suggestions from the referees, M. Combes and H. Rickman, and from D. Bockel´ee-Morvan, were very helpful and motivated us to perform further analysis that resulted in significant improvements in the paper. Financial support for this work was provided by NASA through Grants GO-6663.01-95A and GO7313.01-96A from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA Contract NAS5-26555.

REFERENCES A’Hearn, M. F., R. L. Millis, D. G. Schleicher, D. J. Osip, and P. V. Birch 1995. The ensemble properties of comets: Results from narrowband photometry of 85 comets, 1976–1992. Icarus 118, 223–270. A’Hearn, M. F., D. G. Schleicher, P. D. Feldman, R. L. Millis, and D. T. Thompson 1984. Comet Bowell 1980b. Astron. J. 89, 579–591. Biver, N., and 11 colleagues 1997. Evolution of the outgassing of Comet Hale– Bopp (C/1995 O1) from radio observations. Science 275, 1915–1918. Biver, N., and 22 colleagues 1999. Long term evolution of the outgassing of Comet Hale–Bopp from radio observations. Earth Moon Planets, in press. Bockel´ee-Morvan, D., and E. G´erard 1984. Radio observations of the hydroxyl radical in comets with high spectral resolution. Kinematics and asymmetries of the OH coma in c/Meier (1978 XXI), c/Bradfield (1979 X), and c/Austin (1982 g). Astron. Astrophys. 131, 111–122. Bockel´ee-Morvan, D., and H. Rickman 1999. C/1995 O1 (Hale–Bopp): Gas production curves and their interpretation. Earth Moon Planets, in press. Bockel´ee-Morvan, D., J. Crovisier, and E. G´erard 1990. Retrieving the coma gas expansion velocity in P/Halley, Wilson (1987 VII) and several other comets from the 18-cm OH line shapes. Astron. Astrophys. 238, 382–400. Budzien, S. A., M. C. Festou, and P. D. Feldman 1994. Solar flux variability and the lifetimes of cometary H2 O and OH. Icarus 107, 164–188. Cochran, A. L., and D. G. Schleicher 1993. Observational constraints on the lifetime of cometary H2 O. Icarus 105, 235–253. Colom, P., E. G´erard, J. Crovisier, D. Bockel´ee-Morvan, N. Biver, and H. Rauer 1999. Observations of the OH radical in Comet C/1995 O1 (Hale–Bopp) with the Nan¸cay radio telescope. Earth Moon Planets, in press. Crovisier, J., K. Leech, D. Bockel´ee-Morvan, T. Y. Brooke, M. S. Hanner, B. Altieri, H. U. Keller, and E. Lellouch 1997. The spectrum of Comet

Hale–Bopp (C/1995 O1) observed with the Infrared Space Observatory at 2.9 astronomical units from the Sun. Science 275, 1904–1907. Crovisier, J., K. Leech, D. Bockel´ee-Morvan, E. Lellouch, T. Y. Brooke, M. S. Hanner, B. Altieri, H. U. Keller, and T. Lim 1999. The spectrum of Comet Hale–Bopp as seen by ISO. In The Universe as Seen by ISO, ESA SP-427. European Space Agency, in press. Davies, J. K., T. L. Roush, D. P. Cruikshank, M. J. Bartholomew, T. R. Geballe, T. Owen, and C. de Bergh 1997. The detection of water ice in Comet Hale– Bopp. Icarus 127, 238–245. Delsemme, A. H., and M. R. Combi 1983. Neutral cometary atmospheres. IV. Brightness profiles in the inner coma of Comet Kohoutek 1973 XII. Astrophys. J. 271, 388–397. Despois, D., E. G´erard, J. Crovisier, and I. Kazes 1981. The OH radical in Comets: Observation and analysis of the hyperfine microwave transitions at 1667 MHz and 1665 MHz. Astron. Astrophys. 99, 320–340. Feldman, P. D. 1999. Ultraviolet observations of Comet Hale–Bopp. Earth Moon Planets, in press. Feldman, P. D., and M. F. A’Hearn 1985. Ultraviolet albedo of cometary grains. In Ices in the Solar System (J. Klinger, D. Benest, A. Dollfus, and R. Smoluchowski, Eds.), pp. 453–461. D. Reidel, Dordrecht. Festou, M. C. 1981. The density distribution of neutral compounds in cometary atmospheres. I. Models and equations. Astron. Astrophys. 96, 69–79. Hanner, M. 1981. On the detectability of icy grains in the comae of comets. Icarus 47, 342–350. Haser, L. 1957. Intensity distribution in the head of a comet. Bull. Acad. R. Sci. Li`ege 43, 740–750. Huebner, W. F., J. J. Keady, and S. P. Lyon 1992. Solar photo rates for planetary atmospheres and atmospheric pollutants. Astrophys. Space Sci. 195, 1–294. Jackson, W. M., J. B. Halpern, P. D. Feldman, and J. Rahe 1982. Production of CS and S in Comet Bradfield (1979 X). Astron. Astrophys. 107, 385–389. Kimble, R. A., and 88 colleagues 1998. The on-orbit performance of the Space Telescope Imaging Spectrograph. Astrophys. J. 492, L83–93. Lellouch, E., J. Crovisier, T. Lim, D. Bockel´ee-Morvan, K. Leech, M. S. Hanner, B. Altieri, B. Schmitt, F. Trotta, and H. U. Keller 1998. Evidence for water ice and estimate of dust production rate in Comet Hale–Bopp at 2.9 AU from the Sun. Astron. Astrophys. 339, L9–L12. Schleicher, D. G., and M. F. A’Hearn 1988. The fluorescence of cometary OH. Astrophys. J. 331, 1058–1077. Schleicher, D. G., S. M. Lederer, R. M. Millis, and T. L. Farnham 1997. Photometric behavior of Comet Hale–Bopp (C/1995 O1) before perihelion. Science 275, 1913–1915. Sekanina, Z. 1999. Detection of a satellite orbiting the nucleus of Comet Hale– Bopp (C/1995 O1). Earth Moon Planets, in press. Stern, S. A., W. C. Colwell, M. C. Festou, P. M. Tamblyn, and J. Wm. Parker 1999. Hale–Bopp (C/1995 O1) near 2.3 post-perihelion: SWUIS measurements of the H2 O and dust production. Astron. J., submitted. Weaver, H. A., and P. L. Lamy 1999. Estimating the size of Hale–Bopp’s nucleus. Earth Moon Planets, in press. Weaver, H. A., P. D. Feldman, M. F. A’Hearn, C. Arpigny, J. C. Brandt, M. C. Festou, S. A. Stern, and G. P. Tozzi 1997. The activity and size of the nucleus of Comet Hale–Bopp (C/1995 O1). Science 275, 1900–1904. Woods, T. N., D. K. Prinz, J. London, G. J. Rottman, P. C. Crane, R. P. Cebula, E. Hilsenrath, G. E. Brueckner, M. D. Andrews, O. R. White, M. E. VanHoosier, L. E. Floyd, L. C. Herring, B. G. Knapp, C. K. Pankratz, and P. A. Reiser 1996. Validation of the UARS solar ultraviolet irradiances: Comparison with the ATLAS-1, -2 measurements. J. Geophys. Res. 101, 9541–9569.