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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
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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
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270
240
210
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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
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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
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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
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1
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0.90
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WAVELENGTH (pm)
l
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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
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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
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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
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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+
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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.