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SURFACE COMPOSITION OF MERCURY FROM REFLECTANCE SPECTROPHOTOMETRY FAITH VILAS NASA Johnson Space Center Reflectance spectra of Mercury have been obtained periodically from 1963 through 1984. Since 1969, these observations were made in an effort to learn about the surface mineralogical composition of Mercury. Using the phases of the planet around maximum elongations, Mercury's 6Y3851day rotational rate, and the theory of bidirectional rejectance spectroscopy, some spatial resolution across the planet has been obtained. A very shallow absorptionfeature, which has been attributed to Fe* + in orthopyroxenes, is evident in two recent, highquality spectra but is noticeably absent in a third. This dzfference cannot be explained by rejected light from dzfferent terrain. No noticeable spectral differences exist between the portion of Mercury photographed by Mariner 10 and the unimuged portion of the planet. All of the rejectance spectra display the same slope seen in lunar rejectance spectra, attributed to Fe- and Ti-bearing agglutinates in the lunar regolith. Gravitationalfocusing of meteoric material at Mercury's heliocentric distance and the extreme surface temperatures experienced by Mercury's sunward hemisphere both suggest that the regolith formation rate would be higher than the Moon's, and that agglutinates of unknown composition cause the spectral slope. Electromagnetic radiation received from Mercury at the Earth emanates from two sources: (1) diffuse reflected sunlight in the near-ultraviolet, visible and near-infrared spectral range; and (2) thermally emitted radiation beginning at 1.6 pm resulting from incident sunlight heating the planet's surface and reradiating at longer wavelengths. The method of determining the mineralogy of rock samples by observing how the crystal structure of certain minerals changes the spectrum of diffuse reflected light was pioneered during the late - 60 F. VILAS 1960s (see, e.g., Burns 1970; Adams and Filice 1967) and continues to be studied today (see, e.g., Pieters 1983). Concurrently, photoelectric photometry was developed during the 1960s and applied directly to telescopic observations of the planets. In the late 1960s, groundbased telescopes with narrowband filter photometers were turned toward the Moon and other solar system objects and used to measure how reflected sunlight is affected by the surface regolith. Samples returned from the various Apollo missions in the late 1960s and early 1970s confirmed the basaltic surface mineralogy suggested by the reflectance spectra, thus affirming the validity of this approach to determining the surface composition of a planetary body (see, e.g., Adams and McCord 1970). Today, reflectance spectrophotometry remains the dominant method of remotely sensing the surface mineralogical composition of solar system objects. Reflectance spectrophotometrycan be used as a probe of up to 100 pm depths of surface regolith materials, depending upon the composition (and therefore crystal structure) and particle size of the material (Pieters 1983; Morris and Mendell 1984). Spectral reflectance data can potentially constrain our knowledge about Mercury's surface composition. Indeed, it could even constrain the primordial bulk chemistry of the planet; the assumption made here is that Mercury's surface was not altered or eradicated by a major event such as the volatilization of the outer layer(s) of the planet (Chapter by Cameron et al.), or catastrophic collision (Chapter by Wetherill). In this case, the presence, amount and form of FeO on the planet's surface would serve as an indication of the oxidation state and amount of volatiles in the outer portion of the planet, which is expected to contain silicates. The primary indicator sought in these spectra is the presence and characteristics of an absorption feature centered near 1.0 pm caused by interelectronic transitions of Fe2+ in an octahedral site in olivine [(Mg,Fe),SiO,] or pyroxene [(Mg,Fe)SiO,] or both. Transition energy differences cause the absorption features to be centered at different spectral locations. If a regolith contains both olivine and pyroxene, the spectral width of the absorption band will be broader. These absorption features are common in reflectance spectra of the Moon and asteroids (see, e.g., McFadden et al. 1984; Adams and McCord 1970). In the context of several models for the condensation of Mercury, the presence and amount of volatiles could show where, and under what circumstances, planetary formation occurred. Lewis (see his Chapter) discusses why various condensation and accretion models are inadequate alone to explain Mercury's high density, and how the amount of FeO might be an indicator of the formation process. Wasson (see his Chapter) proposes that the area near Mercury is the formation region for the enstatite chondrites, and would have little or no FeO. Evidence of FeO in the planet's surface material would suggest that the accretion phase of Mercury's formation sampled material from a greater range of heliocentric distances than covered by Mercury's narrow feeding zone. Thus, spectral reflectance studies of Mercury have emphasized - SPECTROPHOTOMETRY OF MERCURY 61 the search for a shallow absorption feature centered near 0.9 pm, seen in some spectra but not in others. This chapter discusses the basic procedure for reducing spectrophotometric data of Mercury; the controversies surrounding (and the implications of) the existing spectra of Mercury, as well as a methodology for defining the portion of Mercury's surface contributing the greatest amount of light to an individual spectrum, including its application to these spectra. DATA ACQUISITION AND REDUCTION As viewed from the Earth, Mercury has a maximum angular elongation of approximately 28" from the Sun. Due to the close visible proximity of the two objects, observations of Mercury must be conducted either during daylight or during twilight while the Sun is below the horizon but Mercury is still visible. This restriction introduces some unavoidable observing problems: careful attention must be paid to avoid scattered sunlight in the telescope tube and instrumentation. The limited dynamic range (operating range within which the detector produces a measureable 1:1 output signal for each input irradiance level) of earlier detectors caused saturation of the received planetary signal during daylight hours. As a result, all spectral reflectance observations reported here were made during twilight except for the 1984 CCD spectrograph measurements. Spectrophotometric observations must also be corrected for the attenuation of light by the Earth's atmosphere along the line of sight from the object to the telescope. The atmospheric thickness will vary from a minimum at the zenith to a maximum along the horizon. The general extinction formula describes the signal attenuation corresponding to the extinction T, caused by the atmosphere at a given wavelength h, and the atmospheric thickness at that location. The airmass X is a logarithmic definition of the atmospheric thickness defined as 1 at the zenith and increasing to infinity at a zenith angle of 90" (the horizon). The signal S, is the expected counts for that spectral interval at the airmass X, and So, represents the counts at an airmass equal to zero. Differential refraction (the dispersion of light at different wavelengths by the thick Earth atmosphere) becomes a problem at high airmass, especially for data acquired at the near-ultraviolet and blue wavelengths obtained with narrowband filter photometers. Active telescope guiding by the observer can compensate where necessary for the images positioned differently on the instrumentation due to refraction. Water present in the Earth's atmosphere (telluric H20) has absorption features centered near 0.73,0.82 and 0.93 pm. The water content of the atmosphere can change throughout an observing session, and observations of Mercury at high airmass can aggravate the H 2 0 absorp- F. VILAS 62 tion. Removal of the effects of water absorption in the Earth's atmosphere from Mercurian reflectance spectra is probably the most difficult task in the data acquisition and reduction procedure. The presence of water absorption affects part of the spectrum covered by the Fe2+ silicate absorption feature. Spectrophotometry of Mercury has been obtained in the visible and nearinfrared spectral region with a variety of instruments, however, the data reduction procedure has generally remained the same. [Some exceptions for specific characteristics of certain instruments have also been included in the data reduction procedures of individual observers. For example, Chapman and Gaffey (1979) describe corrections for coincidence and misalignment of optics in the dual-beam photoelectric photometers which were used to take Mercury data (McCord and Adams 1972a,b; Vilas and McCord 1976). Vilas and Smith (1985) describe the uniformly illuminated "flat field7'observations necessary to correct for variations in the spectral responsivity of individual pixels across a 2-dimensional CCD camera.] For a given spectral interval AX and integration time, the signal received from a solar system object during one observation S, is calculated as where C, is the raw photon count of the object plus the background sky, B, is the background sky count, and D, is the dark count (background count generated by the detector). Since photon counts comprise the data, the calculated signal-to-noise is based upon Poisson statistics. The spectral interval is physically defined by the attributes of the instrumentation (e.g., the passband of a narrowband filter, the grating resolution of a spectrograph). Observations of a calibrated standard star are obtained across a large airmass interval which includes ideally the airmass interval covered by Mercury when observations are made of the planet. For a given spectral interval, the extinction coefficients are calculated using a least-squares fit to the logarithm of the counts vs airmass for all of the observations of the standard star during one observing session. The standard star observation closest in airmass to a specific Mercury observation is identified, and the extinction coefficients are used to adjust the standard star counts to those counts expected at the same airmass as the planet's observation as In S,, = In S, + 7, (X - XC) (3) where S, is the count for the wavelength interval at the airmass X, 7, is the calculated extinction coefficient and S, is the count of the standard star at the planet's airmass X,. The inverse logarithm of this expression produces the corrected standard star value. The corrected standard-star spectrum and the planet's spectrum are independently scaled to 1.OO at a specified wavelength. The reflectance value for SPECTROPHOTOMETRY OF MERCURY 63 Mercury, corrected for the spectral signature of the reflected sunlight, is calculated as ( ) Mercury standard star (standard star) --Mercury Sun Sun The selection of appropriate standard stars is difficult. To date, the selection and calibration of standard stars for planetary spectral reflectance data has been accomplished in two different ways. One method has produced a net of bright standard stars located around the ecliptic (P. Owensby personal communication) observed using the same instrument used for the planetary observations. These stars have been calibrated using observations of Apollo lunar landing sites from which lunar soil samples have been returned to the Earth. Laboratory reflectance spectra of the lunar soil samples have then been used to correct the standard startApollo landing site spectra to become a standard startsun spectrum. This is the only method known of calibrating standard stars which closes the loop between observations of star and planetary object, and calculation of objectlsun, although there remain questions of how representative the returned samples are of the large site area measured telescopically. Alternately, Hardorp (1980) has conducted a careful search for stars whose spectra are solar analogues, documenting in his publications when variations in absorption features exist. If, for a given spectral interval, the ratio of the scaled solar analogue star to the theoretical Sun is approximately 1.O, then the second step in Eq. (4) can be eliminated, i.e., Mercurytsolar analogue star sufficiently represents MercurytSun. A reduction in the calculation steps should eliminate one source of calculated uncertainty. Various observers have tested the resulting reflectance spectra obtained using a solar analogue star. McFadden et al. (1984) compared the reflectance spectra of asteroid 2100 Ra-Shalom ratioed to solar analogue 61 Cygni B and to standard star a Equulei. They find that the spectra are consistent within the calculated errors, except for the near-ultraviolet wavelengths, and they endorse the continued use of both of these standard star sets. For the Eight-Color Asteroid Survey (ECAS), Tedesco et al. (1980) calculated the differences between ECAS photometry of stars that Hardorp designated as spectrally identical to the Sun, and stars he determined were close to solar analogues. Interpolating between the values for these filters to obtain higher spectral resolution data for these stars, Vilas and Smith (1985) note that the differences between most of the Hardorp solar analogue stars and the assumed true Sun are less than the scatter in their asteroid reflectance spectra. They conclude that object-to-solar analogue star ratios represent the object-to-Sun reflectance spectra adequately. The advent of the use of spectrograph and detector combinations, such as CCD spectrographs, which provide comparatively high resolution from 0.95 to 1.OO pm, introduces a further problem. The spectra of some early-type stars (e.g., E Aquarii) contain the hydrogen Paschen lines. These lines, which are smeared in the wider passband filters used in astronomical photometers, are resolved in the higher-resolution photometry. If the values of early-type standard starISun ratios used in narrowband filter photometry are interpolated for the purpose of correcting the higher-resolution spectra for reflected sunlight, the interpolated values will not have sufficient spectral resolution to correct for the structure introduced into the spectrum by the Paschen lines. Even if subsequent stellar calibrations account for the Paschen lines, ratios of solar system objects to early-type standard stars can introduce noise in the spectra. Changes in seeing, instrument flexure, or exact positioning of the target in a spectral slit have required subpixel interpolation of CCD reflectance spectra (see, e.g., Vilas 1985; Vilas and Smith 1985; Buie 1984). The 0.95 to 1.00-pm spectral range covers part of the spectral reflectance absorption feature caused by Fe2+ in olivine or pyroxene or a combination of both, crucial to studies of Mercury. Using a solar analogue star would prevent the introduction of unnecessary structure in the 0.95- to 1.0-pm spectral range, caused by the subpixel misalignment of two spectra, from degrading a high-resolution reflectance spectrum. In this chapter, the spectral reflectance curves are all scaled to the value of 1.00 at 0.7 pm, a spectral range common to all of the data sets discussed, allowing the data to be intercompared. Other scientific literature on Mercury present data scaled differently, but the significance of the data is unaffected by the choice of scaling. One other factor must be considered in reduction of Mercury reflectance data: the high surface temperatures of the illuminated side (590 - 725 K) cause the thermal component of the radiation to become greater than the reflected component beyond 1.6 pm. Corrections must be made to remove the thermal radiation component, in order to consider the spectral reflectance characteristics of Mercury at longer wavelengths where some silicate absorption features are prominent. Clark (1979) described one method of removing this component from circular-variable-filter (CVF) data of Mercury. These analyses are left to future observers and are not discussed further in this chapter. CONSTRAINING THE SPATIAL RESOLUTION OF MERCURY'S SURFACE The reflectance spectra obtained of Mercury to date spatially cover the integral illuminated portion of the planet presented to the Earth at the time and date of the observations. Some correlation of the integral planet spectrum with surface terrain is desired. The planet's orbital and rotational periods, the conditions of illumination, and the observational geometry can be used to derive some spatial resolution of the planet's surface. When Mercury is at maximum I 1 SPECTROPHOTOMETRY OF MERCURY 65 elongation, the phase (portion of the planet illuminated) dictates the longitudinal interval of the planet illuminated during one observing session (longitudinal dimension covered is equal to the difference between 180" and the phase angle). The rotational period of Mercury is taken to be equal to two-thirds of its orbital period or 58.6462 days. Table I contains information about the locations and physical ephemerides of Mercury on all of the dates for which available spectral reflectance observations were made. An eastern elongation as seen from the Earth shows TABLE I Observational Data of Mercury Observation Date (UT) 16 Jun 63 14 Aug 63 15 Aug 63 18 Aug 63 19 Aug 63 22 Aug 63 25 Aug 63 26 Aug 63 19 May 64 20 May 64 22 May 64 15 May 65 18 May 65 17 Jun 69 18 Jun 69 26 Dec 69 13 Mar 72 29 Sep 74 5 Oct 74 6 Oct 74 7 Oct 74 8 Mar 75 9 Mar 75 10 Mar 75 11 Mar 75 21 Apr 76 22 Apr 76 23 Apr 76 24 Nov 84 Phase Anglea (deg) Bright Limb Terminator. Longitude Longitude Hemisphere Telescope Instrumentb ((leg) (deg) (m) 328.4 46.2 50.8 65.5 70.3 85.4 100.7 105.9 269.6 275.2 286.1 331.1 345.3 109.6 115.2 156.3 358.5 107.1 138.2 143.6 149.0 83.6 88.6 93.6 98.6 167.9 172.4 177.2 146.9 246.8 156.7 160.0 170.0 173.4 183.5 193.7 197.0 205.6 208.9 215.4 234.1 242.4 51.1 54.0 265.6 89.4 213.3 231.8 234.7 237.5 341.8 345.2 348.6 352.0 270.5 271.1 271.9 255.1 S S S S S S S S S S S S S unknown unknown S N S S S S S S S S N N N S 0.4 0.4 0.4 0.4 0.4 0.4 0.4 0.4 0.4 0.4 0.4 0.4 0.4 - 1.5 0.9 0.9 0.9 0.9 0,9 0.9 0.9 0.9 0.9 2.2 2.2 2.2 1.O NFP NFP NFP NFP NFP NFP NFP NFP NFP NFP NFP NFP NFP NFF' NFP NFP NFP NFP NFP NFP NFP NFP NFP NFP NFP CVF CVF CVF CCD aPositive phase angle represents eastern elongation; negative phase angle represents western elongation. bNFP = Narrowband filter photometer; CVF = Circular variable filter photometer; CCD = CCD spectrograph. 66 F. VILAS the western side of Mercury to be illuminated; however, the eastern longitudes of the Mercurian hemisphere facing the Earth are illuminated. The reverse is true for a western elongation. Mercury is locked into a quasi-commensurability in which 54 Mercurian sidereal periods are equal to 13.00600 terrestrial tropical years. Thus, the Mercurian physical ephemeris repeats itself every 13 years. During one 13 year interval, northern hemisphere observations favor Mercurian central meridian longitudes of 90 and 270", while the longitudes of 0 and 180" are poorly viewed. However, southern hemisphere observations cover equally well all central meridian longitudes except for 90 and 270". This unusual observing circumstance may have contributed to the late discovery of Mercury's rotational period. Smith and Reese (1968) describe the implications of this phenomenon: ". . . had Schiaparelli or Antoniadi made their observations from Africa, Australia or South America, the true rotation period of Mercury might have been established long ago." Recent observations of the spectral reflectance of Mercury have been conducted at both northern and southern hemisphere observatories, providing relatively complete longitudinal coverage of the planet. Photometric theory can be used to define a subsection of the illuminated portion of Mercury's disk which would contribute the greatest amount of reflected sunlight to a reflectance spectrum. Using Hapke's (1981) bidirectional reflectance formula for a planetary surface composed of closely packed particles of arbitrary shape (Hapke's Eq. 16), Vilas et al. (1984) calculated the distribution of bidirectional reflectance along the illuminated Mercurian surface for different phase angles, assuming a constant albedo planet. The brightest portion of the planet lies at the luminance equator on the planet's limb. As more of the planet's disk is illuminated, the distribution of brightness across the surface becomes more uniform. Figure 1 shows the longitude interval vs the date of observation for the reflectance spectra of Mercury. Insufficient spatial resolution exists for the Earth-based observations to distinguish among the surface terrain types seen in Mariner 10 imagery. In this discussion, the assumption is made that spectral differences on the planet's surface could at best be distinguished between the two dominant geologic terrain types, the intercrater plains and the smooth plains encompassing Caloris Basin (Trask and Guest 1975). Figure 2 shows a sketch of the distribution of these two terrain types as inferred for spectral reflectance studies from Mariner 10 imagery, groundbased radar and albedo data (Chapter by Clark et al.; Zohar and Goldstein 1972; Murray et al. 1972). REFLECTANCE SPECTRA Harris (1961) reported the earliest broadband UBVRI photoelectric measurements of the planet Mercury (see the Chapter by Veverka et al.). Both spectral range and spectral resolution were extended by Irvine et al. (1968) with observations made of Mercury taken around four maximum elongations of the -I I I I I I- I __3111(1 - I I I 4 1 MAR 75 : - I I I -7 _ I I C I I I I DEC69I26- 2 0 JUN~~{~?I 0 MAY~~-[/~I 22 - 192625 22- I I I 181514JUN 6 3 1 16- 360 (0) - 1 I 1-1 I I 3 - I I $----.111 - I I I I I i I I I -1 -1 I- I II 1 I1 330 300 270 240 210 - 180 I 1 150 120 90 MERCURlAN LONGITUDE ( " ) 60 30 0 (360) Fig. 1. Longitude intervals for observation dates of Mercurian reflectance spectra. The broad side is the bright limb tapering to a point for the terminator (indicating the greater contribution of light to the spectrum from the bright limb area). The hatched lines mark the 10" to 190° interval photographed by Mariner 10. 360 (0) 270 180 90 Mercurian Longitude, 0 (360) (O) Fig. 2. The distribution of the intercrater plains (lined portion) and the smooth plains encompassing Caloris Basin (blank portion) on Mercury's surface as inferred for spectral reflectance studies. ~ 68 F. VILAS planet during the years 1963 through 1965. These observations were made using 9 or 10 different narrowband interference filters having passbands centered from 3147 A to 10635 A. In this review, the corresponding relative intensity for each filter was calculated from the data tabulated by Irvine et al. (1968) using the relationship I I I and scaling the intensities to 1.0 for a value interpolated for 0.7 pm from between the values of the 6264 A and 7297 A filters. Vilas et al. (1984) grouped these data around the maximum elongations of Mercury which occurred on 13 June 1963 (23"W), 24 August 1963 (27"E), 24 May 1964 (25"W), and 6 May 1965 (27OW). Figure 3 shows plots of the unweighted mean and standard deviations of the spectral reflectance values taken on dates around these elongations. Irvine et al. do not list errors for the daily Mercury data, and general errors for all observations made during monthly intervals (see Table IX in Irvine et al. 1968) are not included in the errors shown in Fig. 3. McCord first looked at Mercury's reflectance spectrum with sufficient spectral resolution to see mineralogical absorption features in the late 1960s and early 1970s (McCord and Adams 1972a,b; Vilas et al. 1984) using a photometer with 22 narrowband interference filters covering a spectral range of 3190 to 10530 A having passbands of -250 A. The calibration technique using areas on the Moon having known reflectance values was used in the reduction of spectral reflectance values taken near three maximum elongations of 23 June 1969 (23"W), 27 December 1969 (20°E), and 14 March 1972 (18"E). The data are scaled to 1.000 for the 6990 A filter. Figure 4 shows a spectrum composed of the unweighted mean and standard deviation of 4 observations made on the nights of 17, 18 June 1969 from an unknown location, and the unweighted mean and error of 6 observations taken on 13 March 1972 from KPNO. The spectrum formed from the unweighted mean and error of 6 observations taken on 26 December 1969 from CTIO is included in Fig. 5. During 1974 and 1975, Vilas and McCord (1976) obtained spectra on dates around the maximum elongations of 1 October 1974 (26"E) and 6 March 1975 (27OW) at CTIO using a photometer with 24 narrowband interference filters covering a spectral range of 3350 to 10640 A. Vilas et al. (1984) rereduced these spectra using improved standard star calibrations and grouped the spectra by phase angle. The weighted average and error of 37 observations obtained on 5 , 6 , 7 October 1974 formed one spectrum shown in Fig. 5. Eight observations of Mercury taken on 29 September 1974 and the weighted average and error of 56 observations obtained on 8, 9, 10, 11 March 1975, produced two additional spectra shown in Fig. 4. These data are scaled to 1.000 for the 6990 A filter. The spectral range was shifted to the near-infrared (0.65 - 2.5 pm), and spectral resolution was improved to 50 A when McCord and Clark (1979) SPECTROPHOTOMETRYOF MERCURY r 1.~1 . JUN 63 i I AUG 63 MAY 64 Fig. 3. Reflectance spectra obtained from 1963 to 1965. acquired reflectance spectra of Mercury near the maximum elongation of 28 April 1976 (21°E) at Mauna Kea using a circular-variable-filter photometer with 120 passbands. A composite spectrum of 51 observations made on the dates of 21, 22, 23 April 1976 was produced (Fig. 6). No further advances in instrumentation were applied to the problem of observing Mercury until 1984 when Vilas used a 2-dimensional CCD with a spectrograph to get 16 A resolution data across a more limited visible and near-infrared spectral range of 0.5 to 1.0 pm. The multiplexing advantage of Fig. 5. Reflectance spectra obtained with a narrowband filter photometer covering the Mercurian surface area containing both intercrater plains and smooth plains. W 0 z a 7 1 1 1 1 1 1 1 1 I0 1.04 - 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 1 - w MERCURY 1 LL W W > - 0.96 HIGHLANDS- I- a 1 W I I : J I I I 0.70 ~ I 1 l 0.80 I ~ l I 0.90 I I ~ l 1.OO WAVELENGTH (pm) l ~ l ~ 1.I0 Fig. 6. McCord and Clark's (1979) April 1979 reflectance spectrum of Mercury (individual points) compared with a laboratory spectrum of Apollo 16 (lunar highlands) soil showing the 0.89 pm Fez+ orthopyroxene absorption feature. The sloped continuum (reddening) seen in the other spectra has been removed. The spectrum covers the Mercurian surface area containing both smooth plains and intercrater plains. l l ~ l i - 72 F. VILAS this combination of instrument and detector allowed short exposures to be taken of the planet and sky concurrently, permitting observations to be made during the day when the planet was observed through a low airmass. A composite spectrum of 59 observations of Mercury taken on 24 November 1984 at CTIO near the maximum elongation of 25 November 1984 (22"E) covered an airmass range of 1.002 to 1.014 (Fig. 7). The greater dynamic range of the CCD accommodated the background sky brightness. The 2-dimensional CCD allowed the background sky to be mapped concurrently with the planet, thus accounting for temporal changes in sky conditions. The signal-to-noise ratio for this spectrum exceeds 100:1, and a better indicator of the quality of the spectrum is the scatter within the spectrum. Telluric H,O absorption is also very apparent in the spectrum, and is attributable to changes in atmospheric water content during the 6.5 hour delay between the observations of Mercury WAVELENGTH ( p m ) Fig. 7. Reflectance spectrum of the Mercurian surface area covering both intercrater plains and smooth plains, obtained in 1984 with a CCD spectrograph. SPECTROPHOTOMETRY OF MERCURY 73 and the observations of the Hardorp solar analogue standard star, SAO 147237, used to make the reflectance spectrum. A total of twelve composite reflectance spectra of Mercury have been produced from observations made around maximum elongations over the years of 1963 through 1984. Here, these spectra have been divided into three groups: early spectra, spectra covering solely intercrater plains and spectra covering terrain split roughly between intercrater plains and the smooth plains encompassing Caloris Basin. Early Spectra The spectra taken during 1963 to 1965 (Irvine et al. 1968) lack the spectral resolution to show detailed mineralogical information (Fig. 3). Vilas et al. (1984) drew some general conclusions about the surface compositional properties from these spectra. The most important feature of these spectra is that three of the spectra cover terrain unimaged by Mariner 10, while one spectrum covers terrain known to consist of primarily intercrater plains. The similarity between these four spectra suggests that there is no radical difference in mineralogical composition between the unimaged portion of Mercury and the area imaged by Mariner 10. At wavelengths >0.7 pm, these data exhibit increasing slope (reddening) with increasing phase angle, as seen for the Moon. Irvine et al. (1968) noted this same reddening effect when they grouped the spectra by phase angle, not by elongation date. Intercrater Plains Four of the Mercurian spectra (June 1969, March 1972, September 1974, March 1975) having adequate spectral resolution for surface compositional studies (Fig. 4) cover terrain in the 0-90" quadrant of the Mercurian surface, consisting structurally of predominantly intercrater plains. Most absorption features present in these spectra, however, can bc attributed to the incomplete removal of telluric H20. McCord and Adams (1972a,b) first interpreted the dip at 0.95 p m in the March 1972 spectrum as a possible Fe2+ pyroxene absorption. The dip at 0.82 pm in this spectrum indicates, however, that the incomplete removal of telluric H,O is the probable origin of this feature. Caloris Basin and Smooth Plains The smooth plains are considered to be younger than the intercrater plains, emplaced after the impact which formed Caloris Basin. A debate continues over whether the smooth plains are volcanic or impact-related in origin (see Chapter by Spudis and Guest). The spectra in Fig. 5 are those in the visible range covering roughly 50% intercrater plains and 50% smooth plains. A portion of the composite near-infrared spectrum of McCord and Clark (1979) covering the same area is shown in Fig. 6. McCord and Adams (1972a,b) interpreted the dip at 0.95 pm in the December 1969 spectrum as 74 F. VILAS due to pyroxene, but the spectrum shows absorptions near 0.73 and 0.82 pm which again suggests the incomplete removal of telluric H20. The three spectra with sufficient signal-to-noiseto merit serious consideration all cover approximately the same surface region (see Table I, Fig. 1). The October 1974 spectrum shows a shallow absorption feature beginning at 0.8 pm, centered near 0.9 pm, extending to 1.0 pm. The incomplete removal of telluric H,O is also noticeable in the spectrum. Vilas et al. (1984) compared this spectrum with the April 1976 composite spectrum which shows a very weak absorption feature beginning at 0.78 pm and extending to 0.95 pm, centered at 0.89 pm. McCord and Clark (1979) noted an absorption band depth of 4%, and compared the spectrum with the spectrum of Apollo 16 lunar highlands soil containing a minor amount of pyroxenes contributing -5.5% FeO to the composition of the soil. In contrast to these two spectra, the composite spectrum obtained on 24 November 1984 shows only the incomplete removal of the 02A band and telluric H20 absorption features (Vilas 1985). The absence of the absorption feature in the 1984 spectrum reopens the question of the existence of the weak feature seen rarely in the other Mercury spectra. There are two current opinions on the subject. One states that the proposed absorption feature does not exist, and that inaccuracies introduced into the formation of the October 1974 and April 1976 composite spectra produced the weak absorption features. All three spectra could be affected by the incomplete removal of telluric H20 due to the delays between the times that Mercury was observed and the times that the standard stars used for calibration (aLyrae in 1974, P Geminorum in 1976, SAO 147237 in 1984) were observed. (The thinner atmosphere at Mauna Kea's higher altitude reduces the amount of telluric H,O which would have affected the 1976 spectrum.) Vilas (1985) has also suggested that averaging spectra obtained on more than one observation date could have introduced additional problems in the 1974 and 1976 spectra. The phase angle difference among these three spectra is within 20". Using the results of laboratory studies conducted by Pieters (1983) on the effects of phase angle on the reflectance spectrum of the enstatite component of websterite (EN,,), Vilas has shown that the phase differences among these spectra could reduce the observed band depth by at most 6%. The complete disappearance of this feature could not be caused by differences in observational conditions alone. An alternate opinion is that, for reasons still not understood, the absorption feature appears under some conditions affecting Mercury and disappears with changes in these conditions. Earth-based observing conditions do not mask a feature, so physical reasons on the planet's surface must be examined. The surface temperature of the sunward portion of Mercury varies between 590 and 725 K depending upon location on the planet's surface and the planet's orbital position. Laboratory studies by Singer and Roush (1985) show SPECTROPHOTOMETRY OF MERCURY I I I I I I I I I 1 ! I I 75 that for temperatures up to 448 K the orthopyroxene absorption band located near 0.9 pm is enhanced in width and strength; the short-wavelength edge of the feature is constant in spectral location, the long-wavelength edge increases in location, and although the spectral position of the band minimum does not change, the depth increases up to 15%. At higher temperatures, the spectral position of the band minimum increases (Sung et al. 1977). For these Mercury reflectance spectra, the short-wavelength edge of the 0.9 pm absorption feature should not be obscured by the telluric H,O absorption. The issue of the tenuous existence of this absorption feature remains unresolved. What do these spectra imply for the surface composition of Mercury? The slope (reddening) present in all of the Mercurian spectra, and its similar behavior with changes in phase angle to the slopes seen in lunar spectra, suggest that similar material (either compositionally or in physical state) to the lunar regolith exists on Mercury. The slope in the spectra of different lunar terrains (maria, highlands) is caused by Fe- or Ti-bearing agglutinitic glasses, produced by micrometeoroid bombardment, in the surface regolith (Adams and McCord 1973). Metallic iron particles were also found with the lunar agglutinate samples. The environment of Mercury suggests that the regolith could contain agglutinates formed in a similar manner. The composition of these glasses remains unknown at this time. Cintala (1981) has shown that the melt production rate at the Mercurian surface is almost three times greater than at the hottest lunar point, due to the increased surface temperatures at Mercury and gravitational focusing of meteoric material by the Sun at Mercury's heliocentric distance. The flux of meteoric particles is at least eight times greater at Mercury than at the Moon (Leinert et al. 1981), suggesting that a much greater regolith formation rate is operating on the Mercurian surface, although that does not necessarily increase the fraction of agglutinates expected. However, the higher impact velocities might increase the fraction of agglutinates. The composition of possible Mercurian agglutinates remains unknown. The weak absorption feature has been attributed to Fe2+ in orthopyroxenes (McCord and Adams 1972a,b; McCord and Clark 1979) in the surface material. If the 0.9 pm pyroxene absorption feature is not present, as the most recent composite spectrum suggests, then information about the surface composition can be inferred only from its conspicuous absence. The lack of an absorption feature suggests that the surface of Mercury is highly reduced. Since condensation sequences proposed for the origin of the solar system generally have increasing FeO with decreasing temperature, and increasing volatile content with decreasing condensation temperature, a lack of volatiles on Mercury's surface could be inferred from the lack of FeO. Future spectral reflectance observations of Mercury in the visible and near-infrared spectral range could resolve the question of the existence of the Fe2 absorption feature more firmly. Valuable data on the mineralogical com+ 76 F. VILAS position can also be obtained in the thermal infrared and near-ultraviolet spectral range from Earth-based studies. The difficulties with groundbased observations remain, although improved instrumentation and observations of Mercury above the Earth's atmosphere from space will enhance future observing efforts.