Spectral emissivity measurements of Mercury`s surface indicate Mg

Transcription

ARTICLE IN PRESS
Planetary and Space Science 57 (2009) 364–383
Contents lists available at ScienceDirect
Planetary and Space Science
journal homepage: www.elsevier.com/locate/pss
Spectral emissivity measurements of Mercury’s surface indicate Mg- and
Ca-rich mineralogy, K-spar, Na-rich plagioclase, rutile, with possible
perovskite, and garnet
A.L. Sprague a,, K.L. Donaldson Hanna b, R.W.H. Kozlowski c, J. Helbert d, A. Maturilli d,
J.B. Warell e, J.L. Hora f
a
Lunar and Planetary Laboratory, University of Arizona, 1629 E. University Blvd., Tucson, AZ 85721-0092, USA
Brown University, Providence, RI 02912, USA
c
Susquehanna University, Selinsgrove, PA 17870, USA
d
Institute for Planetary Research, DLR, Rutherfordstrasse 2, 12489 Berlin, Germany
e
Institutionen for Astronomi och Rymfysik, Uppsala Universitet, Uppsala, Sweden
f
Harvard-Smithsonian Center for Astrophysics, Cambridge, MA 02138, USA
b
a r t i c l e in f o
a b s t r a c t
Article history:
Received 8 October 2008
Received in revised form
14 January 2009
Accepted 15 January 2009
Available online 29 January 2009
Mid-infrared 2-D spectroscopic measurements from 8.0 to 12.7 mm of Mercury were taken using Boston
University’s Mid-Infrared Spectrometer and Imager (MIRSI) mounted on the NASA Infrared Telescope
Facility (IRTF) on Mauna Kea, Hawaii, 7–11 April 2006. Measurements reported here cover radar bright
region C, a dark plains region west of Caloris Basin, and the interior of Caloris Basin. By use of spectral
deconvolution with a large spectral library composed of many mineral compositions and grain size
separates, we fitted, or ‘‘unmixed’’, the Mercury spectra. We find mineral suites composed of
magnesium-rich orthopyroxene and olivine, Ca-, Mg-, Na-rich clinopyroxene, potassium feldspar, and
Na-bearing plagioclase feldspar. Both Ca- and Mg-rich garnet (pyrope and grossular, respectively) are
apparently present in small amounts. Opaque minerals are required for spectral matching, with rutile
(TiO2) repeatedly providing the ‘‘best fit’’. However, in the case of the radar bright region C, perovskite
also contributed to a very good fit. Caloris Basin infill is rich in both potassium feldspar and Na-rich
plagioclase. There is little or no olivine in the Caloris interior smooth plains. Together with the high
alkali content, this indicates that resurfacing magmas were low to intermediate in SiO2. Data suggest
the dark plains exterior to Caloris are highly differentiated low-iron basaltic magmas resulting in
material that might be classified as oligoclase basalts.
& 2009 Elsevier Ltd. All rights reserved.
Keywords:
Mercury
Mercury’s surface composition
Iron-poor mineralogy
Spectroscopy of mercury’s surface
Mercury’s formation
1. Introduction
Here we present results of Mercury surface observations using
IRTF on Mauna Kea Hawaii. We used MIRSI (http://www.cfa.harvard.
edu/mirsi/) to obtain mid-infrared spectra from three regions on
Mercury’s surface. It is important to know the composition of
Mercury’s surface to build an understanding of the volcanic and
thermal history of the planet, infer the composition of its crust
and model the composition of the mantle on the assumption that
the lavas were extracted from it by partial melting. To this end, we
discuss in this manuscript, measurements from radar bright
region C, dark plains west of Caloris Basin, and the interior of
Caloris Basin. The data presented here were obtained in 2006.
Data reduction and analysis have been time consuming because
Corresponding author. Tel.: 520 621 2282; fax: 520 621 4933.
E-mail address: [email protected] (A.L. Sprague).
0032-0633/$ - see front matter & 2009 Elsevier Ltd. All rights reserved.
doi:10.1016/j.pss.2009.01.006
we have been pioneering the data analysis procedures which are
described here along with some details of the instrumentation
used and the observations. First results from analysis of some
Mercury and lunar data obtained during the same time period
have been presented previously (e.g. Hanna et al., 2006; Sprague
et al., 2008), but this is the first time that a detailed description of
observations, methods, results, and implications of these data
from these three specially chosen regions have been presented.
1.1. Brief historical context
Many ground-based observations of Mercury and two NASA
spacecraft have made measurements of Mercury. The surface
superficially resembles that of the Earth’s Moon, covered with
silicates, heavily cratered over much of the surface, with bright
crater ejecta rays against a darker background (cf. Strom et al.,
1975; Davies et al., 1978). Imaging by the Mercury Dual Imaging
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System (MDIS) on MESSENGER spacecraft has revealed volcanic
vents, pyroclastics, and other evidence of extrusive volcanism
(Head et al., 2008; Murchie et al., 2008). Briefly stated, Mercury’s
surface materials are low in FeO which is common in silicates on
the Moon, asteroids, and Earth (cf. Vilas, 1988; Warell et al., 2006.
McClintock et al., 2008). Microwave imaging and modeling has
observed Mercury’s regolith to be more transparent (lower in
Ti and Fe) than that of the Moon (cf. Mitchell and de Pater, 1994).
Previous mid-infrared observations of Mercury’s surface at other
locations (Sprague et al., 1994, 1997a, 2002; Sprague and Roush,
1998; Emery et al., 1998, Cooper et al., 2001), have discerned a
heterogeneous surface chemistry based on emissivity maxima,
transparency minima and variations in spectral activity in Reststrählen bands between 7.7 and 13.5 mm. These results are
complementary to those from the Mercury Atmospheric and
Surface Spectrometer (MASCS) (McClintock et al., 2008), MDIS
(Robinson et al., 2008) and other data from instruments obtained
during the first MESSENGER flyby of Mercury because the same
regions on Mercury’s surface were traversed (Solomon et al.,
2008). Neither the MDIS nor the MASCS discerned the FeO
absorption band from Mercury’s surface materials at the locations
measured. The synergy between the MESSENGER observations
and our ground-based measurements will be discussed in detail in
the sections to follow.
2. Geologic context of three regions measured
The footprint of each location (Fig. 1A–C) has been outlined in
white on a grey scale on a portion of the departure mosaic
obtained by the MDIS during the first Mercury flyby in January
2008. The size and location of the three regions is estimated from
imagery obtained at the time of data acquisition. Our long slit
spectra permit the acquisition of spatially resolved regions on the
surface no smaller than 400–600 km either perpendicular to or
along the spectrograph slit. Thus our measurements blend the
signal from many smaller regions into our signal and the resulting
spectra are a mixture of spectra from all rock, mineral, and soil
types in the footprint. We provide compositions of extended
regions which are the estimated abundances of mineral phases
contributing to the signal in our data. Thus, we do not discern
smaller compositional units which may exist within the regions,
but the blend of all mineral phases in the region. Using the
location of the slit on the Earth-facing disk of Mercury and the
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ephemerides for the day and time of observation, we computed
the locations along the slit in latitude and longitude.
2.1. Radar bright region C (RBC)
Goldstone-VLA X-band imaging of Mercury (Slade et al., 1992,
Butler et al., 1993) observed a bright radar return signal over a
large irregular region centered N of the equator at about 1201E.
This high radar backscattering region included and surrounds an
area that corresponds approximately with one of the ‘‘Goldstein
features’’ seen several years before in pioneering radar imaging of
Mercury (Goldstein, 1970, 1971). The high radar backscatter
properties indicate a rough and blocky surface possibly composed
of excavated material generated during impact, or some other
process that caused roughness on the scale of several decimeters.
The region has acquired the familiar name of ‘‘radar bright spot C’’
(Harmon, 1997). The region appears as an irregular radar bright
region in the radar albedo map adapted from Butler et al. (1993)
(section of Fig. 2 with grid lines) and as a rough, heavily impacted,
and irregularly surfaced region in Arecibo S band (12.6 cm
wavelength) images in dual polarization, delayed-Doppler imaging which found the area to correspond to heavily cratered
terrain with a relatively fresh 125 km diameter impact crater
centered at 1141E, 111N latitude surrounded by prominent lobate
ejecta also within the region (Harmon et al., 2007) (bottom
section of Fig. 2). The extended system of bright impact ejecta rays
which appears bright against the dark background east of radar
bright spot C may contribute slightly to, but does not explain, the
high radar backscatter over the entire region. The region was
chosen for discussion here because ground-based imaging by
Ksanfomality (2004) suggested the presence of a large basin to the
west. Ejecta from the basin-forming impact may be in this region
and thus have spectral signatures from the apparently heavily
cratered and highly radar backscattering characteristics. So far
there is scant evidence for the basin suggested by Ksanfomality
(2004) in the MESSENGER imaging but additional images and a
variety of illumination geometries will be necessary to determine
the presence or extent of an ancient, highly modified feature.
MDIS imaging Wide Angle Camera departure color sequences
show the radar bright region C and surrounding terrain to have
many circular bright albedo areas, corresponding to heavily
cratered terrain with inter-fingering of material 10% lower in
reflectance than the planetary average (Robinson et al., 2008).
2.2. Dark plains west of Caloris Basin (DPWCB)
Fig. 1. MDIS departure image from the first fly by of Mercury by the MESSENGER
spacecraft illustrates the three regions of this study. Region A—radar bright region
C (RBC), rough on the scale of decimeter wavelength and longer. The prominent
impact crater centered at 1141E, 111N, its lobate ejecta and the heavily cratered
regions surrounding this impact contribute to the high radar backscatter. Region
B—dark plains west of Caloris Basin (DPWCB). Region C—Caloris Basin (CB). The
white outlines indicate the approximate footprint of mid-infrared spectral imaging
for each region. Mercury mosaic is courtesy of Applied Physics Laboratory and
Johns Hopkins University, www.jhu.messenger.misison.
The smooth dark plains surrounding and west of Caloris Basin
exhibit a crater density 40% less than on the plains interior to
Caloris Basin and are therefore thought to be volcanic plains
younger than those in the Caloris Basin interior (Strom et al., 2008).
Similar dark plains material, classified by Robinson et al. (2008) as
generally low-reflectance material (LRM) extends to the south of
Caloris Basin and covers much of the side of the southern
hemisphere imaged during the flyby. Analysis of our spectra
presents direct evidence for mineral phases and relative abundances
for the dark plains region (white outlines B, Fig. 1). Spectral
measurements of the dark plains surface materials presented in this
paper were obtained on 8 and 9 April 2006 (frames #063 and #098,
respectively). The data and ‘‘best fit’’ modeling of data obtained on
those two separate days enhance our confidence in the results.
2.3. Caloris Basin (CB)
One exciting discovery of the first flyby was evidence for
intrusive and extrusive volcanic activity in and around Caloris
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enhanced signal of neutral atmospheric potassium emission
(Sprague et al., 1990).
Our observations within Caloris Basin identify mineral phases
consistent with MDIS imaging that shows a relatively high albedo
with a redder slope than the surrounding dark plains. We discuss
measurements of materials in Caloris Basin obtained in two
consecutive days (8 and 9 April 2006; frames #071 and #080,
respectively) and outlined in white Fig. 1C.
3. Observations
Spectral measurements of Mercury and the Moon’s surfaces
were taken using Boston University’s MIRSI mounted on the NASA
IRTF on Mauna Kea, Hawaii. Observations were obtained during
daytime observations on 7–11 April 2006 covering longitudes
between 1001 and 1841E. Mid-infrared standard stars, b Pegasus
(b Peg) and a Bootis (a Boo), were measured in the same
observing mode as Mercury. At the time of the observations,
Mercury was between 7.4 and 7.9 arcsec in diameter, 0.46 AU
from the Sun, with the sub-Earth longitude varying from 1881 to
1621E longitude.
MIRSI is an imaging spectrograph with the 10-mm grism
covering 8–14 mm with resolving power (l/Dl) equal to 200 and a
slit width of 0.600 . The MIRSI detector is a 16-channel 320 240
Si:As IBC array developed by Raytheon; each channel measures
20 240 pixels. Daytime observations are kept short to keep the
detector response within the linear regime in all 16 channels
characterized by the same gain. More details about MIRSI can be
found in Deutsch et al. (2003) and Kassis et al. (2006, 2008).
3.1. Mercury and standard star observations
Fig. 2. (top) Goldstone-VLA X-band imaging showing RBC as the dark irregular
patch between the equator and 301N and 100–1401E longitude (adapted from
Butler et al., 1993). (bottom) Arecibo S-band dual polarization, delay-Doppler
imaging showing heavily cratered terrain in the same region (adapted from Fig. 7,
Harmon et al., 2007).
Basin. An extensive radial fracture system (Pantheon Fossae) with
over 100 graben likely results from stress faulting following
surface uplift, caused by pressure from upwelling magma into a
system of dikes near the basin center or volcanic vents observed in
and around Caloris Basin (cf. Head et al., 2008, Murchie et al.,
2008). While fracturing and faulting within Caloris Basin was
known from Mariner 10 imaging and previously related to a
period of uplift of the crater floor, possibly caused by either
exterior loading lateral crustal flow (cf. Thomas et al., 1988,
Watters et al., 2005), the MDIS imaging from the flyby has
permitted a much better understanding of the extensive volcanic
activity that has taken place within the Caloris rim material
(Murchie et al., 2008). It has been proposed that the extensive
fracture system within the basin, along with the high temperatures at the longitude of Caloris Basin at perihelion, facilitate the
diffusion of potassium (K) into Mercury’s exosphere and cause the
Mercury was observed near maximum western elongation in
the early morning as it transited. b Peg (M2.5 II–III spectral type
star) was chosen because of its proximity to Mercury in the sky
and its relatively large flux in the mid-infrared. Long slit spectra
were obtained for b Peg prior to and following Mercury observations. Mercury and b Peg spectra were taken in chop/nod mode
as is standard technique for mid-infrared observations. We
obtained narrow band filter images of the target (both Mercury
and standard star) at several wavelengths between 7.7 and 13 mm
(spectral data extend from 8.0 to 12.7 mm) for flux calibration. In
addition, an image shows the extent to which atmospheric
turbulence smears the image of Mercury and, by inference, the
view of Mercury’s surface in the long slit (see Fig. 3). Note that the
regions along the slits in Fig. 3 are reprojected onto an MDIS
image with a different geometry in Fig. 1.
The location of the slit on Mercury’s surface can be determined
using ephemerides for Mercury for the exact time of spectral
integration. A good estimate of the seeing smear can be made by
comparing the observed image of Mercury to a model computed
with no seeing smear. Using both model and image, the location of
the slit on the disk and the number of distinguishable sectors
Fig. 3. Images, taken in MIRSI’s imaging mode, just prior to the spectral image
integration for (a) RBC, (b) DPWCB, and (c) CB. The location and width of the MIRSI
spectrograph slit at the time of integration is indicated by the parallel black lines.
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along the slit in latitude and longitude can be determined. The
signal along the portion of the slit that falls over the disk of the
planet, which is typically 6–8 arcsec in diameter, can be divided
into latitude regions for spatially resolved spectral study. The
spatial extent of the spectrograph slit width and CCD array pixel
are 0.6 and 0.3 arcsec. Mercury was 7.7 arcsec in diameter at the
time of measurement. With perfect ‘‘seeing’’ the spectrum would
be 26 lines (pixels) from north pole to south pole. With
atmospheric seeing, the spectrum contains more lines and, when
plotted, the flux approximates a Gaussian which is typical for slit
spectroscopy of Mercury’s surface (see Sprague et al., 1997b) for a
detailed study of this phenomenon and seeing deconvolution
method. Typically we use a simple Gaussian fit to the continuum
from limb to limb along the slit which may fall anywhere on the
illuminated disk and is not necessarily subtending a full diameter.
For the sake of simplicity we divide the spectrum into four equal
regions with the cut off at 2 sigma from the peak. Each frame has a
slightly different seeing smear but on average each spatially
discrete sector is 600 km mapped along the slit. We name each
spectrum by date of collection, identify the spatial sectors as north
(N), north-mid (NM), south-mid (SM), and south (S), respectively,
and archive spectra for future analysis. In this paper we present
the north mid-latitude sector over radar bright C and the dark
plains west of Caloris Basin. In the case of Caloris Basin it was
necessary to extract several lines from our spectral images that
did not fall into one of the standard four divisions.
4. Data reduction
4.1. Sky removal, dispersion correction, and image rectification
Along with emitted light from Mercury or the standard star,
MIRSI spectral images contain signal from the telescope, instrument, day sky, and dark current. Signals from the sky and dark
current are easily removed through a subtraction of chop and nod
frames. Signal from the sky and dark current must be removed
from a spectral image to increase its signal-to-noise ratio. While
images are collected in chop/nod mode, the secondary mirror
chops on target storing a spectral image in frames 0 and 1 of a FITS
file; the telescope then nods to the sky and the secondary mirror
chops on the sky storing a spectral image in frames 2 and 3 of the
same FITS file. Spectral images are dark noise and sky corrected by
subtracting the differenced sky frames from the differenced target
frames ((target 1sky 1)(target 2sky 2)) in the usual fashion
for mid-infrared telescopic observations.
MIRSI images exhibit spectral and spatial distortions, curvature
in both x (wavelength dispersion) and y (spatial dimension)
coordinates of the spectrum due to refraction of light as it passes
through the optics. The correction of spectral and spatial
distortions is done using tools in Image Reduction and Analysis
Facility (IRAF) (Tody, 1986, 1993). A lunar spectrum obtained
during the same observing period was chosen to fit dispersion
functions across both lines and columns (320 240 pixels). The
chosen lunar spectrum had several distinct low albedo regions
(dark shadows, maria), several rows tall that were used to fit the
dispersion function in the y-coordinate. Ozone and water absorption bands were chosen in the lunar image to fit the dispersion
function in the x-coordinate. IRAF tools were used to fit dispersion
functions curvatures with polynomials and then pixels were rebinned for linearization in both dimensions. The dispersion
functions were then applied to all spectra.
Telluric absorptions measured with the MIRSI grism and
telluric absorptions in atmospheric spectra collected with the Jet
Propulsion Laboratory’s (JPL) MkIV Interferometer were used for
calculating the wavelength dispersion of MIRSI’s 10 mm grism. The
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JPL MkIV Interferometer is a high-resolution Fourier Transform
Infra-Red (FTIR) Spectrometer designed to remotely sense the
composition of the Earth’s atmosphere by the technique of solar
absorption spectrometry (Toon, 1991). Spectral observations range
from low air mass (solar zenith angle ¼ 201) to high air mass
(solar zenith angle ¼ 851) under warm and cold conditions (from
plus 101 to negative 15 1C). MkIV spectra were convolved to the slit
function of the MIRSI instrument to smooth the high-resolution
MkIV spectra to the lower resolution MIRSI spectra. The slit
function for the 10 mm grism is a Gaussian function with a sigma
of 1.6 pixels (or 0.0306 mm assuming a linear spectral dispersion
of 0.0191 mm). A running Gaussian smooth was performed for a
2-sigma Gaussian width of 1.0, 1.2, and 1.6 pixels full-width halfmaxima. Five telluric features observed in both the MkIV and
MIRSI spectra were chosen: an H2O feature at 8.111 mm, three O3
features at 9.486, 9.573, and 9.651 mm and a CO2 at 12.709 mm.
Pixel positions in the x-coordinate for the five telluric features
were found in the MIRSI calibration spectrum and were then fitted
with a polynomial function to map the wavelength scale.
4.2. Correction for telluric absorptions, the stellar spectral shape and
thermal continuum
After the spectral images had been corrected for instrumental
effects it was necessary to remove telluric absorptions, stellar
spectral shape, and the thermal slope of Mercury’s rough surface
to prepare Mercury spectra for spectral ‘‘unmixing’’ (hereafter
called UM). We used the standard method for removing telluric
absorptions in telescopic spectral data (taking a ratio of the target
spectrum to a spectrum of a well characterized mid-infrared
standard star collected at the same air mass as the target
spectrum). Stellar spectral images were collected as close as
possible on the sky and in time to Mercury spectra in order to
approximate the same atmospheric depth and opacity. In the
cases when the identical air mass was not achieved, we
‘‘corrected’’ stellar spectra to the air mass of Mercury. This was
made possible by computing extinction coefficients from sequential stellar spectra for each MIRSI wavelength and interpolating
the actual spectrum to the airmass at which the Mercury
spectrum had been obtained. After division of Mercury spectra
by standard star spectra to complete telluric corrections, the
quotient was corrected for the stellar spectral slope and features.
Infrared standard stars’ spectral shapes have been determined by
Cohen et al. (1995, 1996). Stellar spectra contain absorption bands
due to SiO and CO fundamentals and a general continuum whose
shape is controlled mainly by effective temperature. After
correction of the Mercury spectra for both stellar shape and
telluric absorptions we divide by a rough surface thermal model
(Emery et al., 1998) computed for the exact ephemerides of the
target during the observation. This division removes the spectral
slope introduced by Mercury’s hot rough surface and prepares it
for deconvolution processing (spectral unmixing).
5. Spectral analysis
Our Mercury spectra have been analyzed by spectral UM using
an established spectral deconvolution algorithm based on the
principle that the emitted or reflected energy from a multimineralic surface is a linear combination of the energy radiated
from each component in proportion to its areal percentage
(Ramsey, 1996; Ramsey and Christensen, 1998). Using this
assumption, computation of the percentage of spectral endmember minerals with known particle size and density approximates the abundance of each end-member present. Ramsey and
Christensen (1998) generated laboratory mixes with varying
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abundances of hornblende, microcline, oligoclase, and quartz and
performed a blind retrieval test by selecting all end-members.
Results for the blind retrieval test showed that for each mineral in
the mixture differences between the UM abundance and the actual
abundance could be as great as 12%, but the average difference was
4%. For some of the laboratory mixtures, the UM fitted minerals that
are not present in the laboratory mixture, but these are fitted with
abundances of less than 5%. This indicates that the linear spectral
deconvolution method can be used to predict mineral abundances
to within 5% in the cases where the spectral end-members are
spectra from the actual sample mineral mixes. An iterative process
of comparison of linear combinations of spectra end-members to
the target data is made. A root mean square (RMS) error to the fit to
the target data over the entire spectral range and a residual model
spectrum are generated for each ‘‘best fit model’’ (BFM) spectrum. A
blackbody can also be added to the spectral end-member UM to
accommodate variations in the spectral contrast between the input
spectrum and the best model fit (Hamilton et al., 1997).
For our Mercury spectral fitting, the best model fits are
dependent our spectral libraries, built with many mineral compositions and different grain size fractions. They are unlikely to have the
exact composition and grain size mixture of the regolith on Mercury
(which of course we do not know). Thus best fit models are
approximations of what might be present on Mercury’s surface or
representative of mineral phases and grain sizes present given the
assemblage of spectra used in the spectral library during the UM.
Hamilton and Christensen (2000) found that end-members chosen
by the deconvolution process for the best model fit will be
overestimated if some spectral end-members of the exact mineral
composition are absent from the spectral library. Hamilton et al.
(1997) demonstrated this with the olivine solid solution (fayalite to
forsterite) and cautions that the compositions of the spectral endmembers chosen will only approximate those of the actual target
spectrum. Minor and trace constituents usually are not accurately
chosen because their contributions to the bulk spectrum may in
some cases be slight and the algorithm will judge their contribution
as null. Exceptions to this case occur when the minor constituent
has a distinctive spectral feature that fits a spectral feature in the
target spectrum. An additional complication that may lead flawed
abundance computations is that a spectrum of the correct mineral
phase may be present in the spectral library but not measured at
the correct grain size to perfectly match that of the regolith for the
region measured. To our knowledge no good quantitative study of
these effects has yet been published.
5.1. Previous use of deconvolution algorithm in mid-infrared
analyses
Previous successes interpreting target spectra with this
spectral deconvolution algorithm include: Hamilton et al. (1997)
who fitted whole Martian meteorite samples representing the SNC
meteorites, Feely and Christensen (1999) who modeled whole
rock samples that had previously been studied using petrographic
techniques, Wyatt et al. (2001) who determined the phase
abundances of whole igneous rocks and comparing the results
with phase abundances determined by electron microprobe
mapping, and Milam et al. (2007) who modeled complex mixtures
of plagioclase sands and other mineral phases. In addition,
Donaldson Hanna and Sprague (2009) have identified compositions of HED meteorites and Vesta from deconvolution of
laboratory and telescopic mid-infrared spectra, respectively.
5.2. Spectral libraries and spectral library end-members
For our Mercury spectral studies it was necessary to build a
large spectral library of many mineral compositions and for many
grain size separates (a good discussion of the effect of small and
hyperfine grain sizes on rock, mineral, and soil spectra can be
found in Mustard and Hays, 1997). Mercury’s surface is well
comminuted from meteoroid impact and its thermal infrared
spectrum has been shown to be dominated by small grain size
particles from observations in the VISNIR (Dollfus and Auriere,
1974; Warell and Blewett, 2004) and in the mid-infrared
(cf. Henderson and Jakosky, 1997; Emery et al., 1998). Even
though evidence suggests there is little or no FeO absorption in
near-infrared spectra from Mercury (see Warell et al., 2006) we do
not assume any single starting set of rock or mineral spectra that
might be a priori ‘‘suitable’’ like lunar or martian suites of
minerals, or only minerals low in FeO. Our spectral libraries have
spectral end-members (a term used to denote the spectra in the
library that may be used to fit the target spectrum during the
deconvolution; from Christensen et al., 2000) with a broad range
of composition representing rock-forming minerals on Earth, the
Moon, asteroids, and meteorites. We use spectra from several
large spectral libraries including reflectance spectra of plagioclase
glasses and typical lunar samples (Nash, 1991; Nash and Salisbury,
1991, respectively) and spectra from the ASTER collection (Hook,
1998). The ASTER collection includes laboratory reflectance
spectra from two sources: Johns Hopkins University (JHU)
(Salisbury et al., 1987, 1988, 1991) and JPL (Grove et al., 1992).
JHU end-member samples were sieved into two grain size
fractions (0–74 and 74–250 mm) and reflectance measurements
were obtained using a Nicolet* 5 DxB interferometer spectrometer. JPL end-member samples were sieved into three grain size
fractions (o45, 45–125, and 125–500 mm) and reflectance
measurements were obtained over the 2.2–25.0 mm spectral range
using a Nicolet interferometer spectrometer. Many more mineral
species and phases were available to be easily downloaded from
the Brown University RELAB spectral library (Pieters and Hiroi,
2004) where samples are of varying grain size and are measured
over the 2.0–25.0 mm spectral range using a Thermo Nicolet Nexus
870 spectrometer. We also obtained end-member samples of
varying grain size fractions from the USGS spectral library (Clark
et al., 2007) measured using a Nicolet Fourier Transform Infra-Red
(FTIR) interferometer spectrometer covering 1.3–150 mm.
We used selected emittance spectra from two sources. (1) The
Arizona State University (ASU) spectral library (Ruff et al., 1997,
Christensen et al., 2000) where end-member samples available
are mostly of a single grain size (710–1000 mm). Thermal emission
spectra are collected under a nitrogen purge at over the
2.0–25.0 mm spectral range using a Mattson Cygnus 100 interferometric spectrometer. (2) The Planetary Emission Laboratory
(PEL) established to support future mid-infrared spectral measurements of Mercury’s surface by the Mercury Emission Radiometer and Spectrometer (MERTIS) (Helbert et al., 2007; Helbert
and Maturilli, 2008; Hiesinger et al., 2008; Maturilli et al., 2006,
2008) as part of the BepiColombo mission (Benkhoff et al., 2009)
to Mercury that is scheduled for a 2014 launch. Laboratory
measurements are made of many mineral types and of four grain
size fractions: some of them in the o25, 25–63, 63–90, and
90–125 mm range (Maturilli et al., 2006), others in the o25,
25–63, 63–125, and 125–250 mm range (Maturilli et al., 2008).
Thermal emission spectra are collected over 6.3–22.0 mm using a
Fourier transform infrared spectrometer (Maturilli et al., 2006).
Systematic measurements are made at incremental temperatures
appropriate for Mercury’s regolith; 100–700 K (Helbert and
Maturilli, 2008). These spectra are accumulated and categorized
in the Berlin Emissivity Database (BED). The chamber for the
vacuum measurements has been purchased but the laboratory
setup and the expected spectral library are not yet available.
Our Mercury spectra are emittance spectra emanating from
surface materials at a wide range of temperatures, mostly hotter
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369
than room temperature, in a vacuum environment. Laboratory
spectra obtained in such conditions are not available. Thus large
reflectance and emissivity spectral libraries generated at room
temperature and one bar pressure are used in our deconvolution
process. We inverted the reflectance spectra to emittance using
Kirchoff’s law (Emissivity ¼ 1Reflectance), the well-known
relationship that is regularly used in remote sensing applications
(Salisbury et al., 1994). All of the measurements of the spectral
libraries listed above were made in reflectance at room temperature and at one bar atmosphere. Small differences between
biconical reflectance and hemispherical reflectance spectra of
the same laboratory sample have been thoroughly discussed by
Salisbury et al. (1994) and there are slight departures from an
exact application of Kirchoff’s Law to the former. The additional
uncertainty in the deconvolution is likely to be commensurate
with the small differences in conversion of the two types of
reflectance spectra. We combined the spectra measured in
emittance with the inverted reflectance spectra to form our
working spectral end-member library. Diagnostic absorption
bands may not be quantitatively reproduced among these
different types of spectra for the same exact mineral phase and
grain size but their use is the best we could do at this time.
6. Previous results and spectral library comparisons
We have found it instructive to use the same spectral libraries
used to successfully model Howardite, Eucrite, and Diogenite
(HED) meteorites, aubrite meteorites, Vesta, and the Moon to UM
spectra from Mercury. If we restrict ourselves to the same mineral
phase end-members for deconvolution of Mercury spectra the
results are not satisfactory. This is an important point because it
may not be obvious that some spectral end-member libraries have
mineral suites that are suited for the HED meteorites or aubrites,
for example, but not for Mercury. It was not obvious until these
comparisons were made to the HED library, the Vesta library, and
the lunar library that no good matches over the entire wavelength
range were possible. This important result is now examined in
some detail. Figs. 4–6 illustrate this point. Donaldson Hanna and
Sprague (2009) found good spectral deconvolution matches to
spectra of differentiated meteorites Johnstown (diogenite) and
Bholghati (howardite) with excellent correspondence to the
correct mineral suites known to be present in them. The spectral
deconvolution algorithm produces a best fit using the primary
minerals for the Howardite meteorite (low- and high-Ca pyroxene
and plagioclase feldspar) with minor amounts of chromite and
olivine (Fo89). The major phase end-members (greater than
10 vol%) are plotted in Fig. 4. Major minerals chosen in the best
model fit included low-Ca pyroxene, high-Ca pyroxene, and
plagioclase. Two low-Ca pyroxenes were chosen: a BED pigeonite
with a grain size fraction of 25–63 mm and USGS hypersthene
(PYX02) with a grain size fraction of 4250 mm. Compositions of
both the pigeonite and hypersthene end-member samples fit
within the large phase range measured by Furhman and Papike
(1981). The deconvolution algorithm overestimated the modal
abundance of low-Ca pyroxene by 6% over the measured range.
The high-Ca pyroxene chosen for the best fit model is a RELAB
Zagami pyroxene with a grain size fraction of 0–50 mm. The modal
abundance of high-Ca pyroxene (24%) determined by the
deconvolution algorithm falls within the measured range and
other derived mineral phases and abundances for Bholghati fit
within the definition of a Howardite meteorite.
Also shown in Fig. 4 is the deconvolution UM for Johnstown,
the largest representative by mass of the diogenite class of
meteorites (Floran et al., 1981). Diogenites are dominated
(approximately 84–100 vol%) by orthopyroxene with minor
Fig. 4. Spectral unmixing of HED meteorites. Figure adapted from Donaldson
Hanna and Sprague (2009). For discussion see Section 6, page 14.
Fig. 5. Mercury spectrum #071 from Caloris Basin is compared to best fit models
computed using three different spectral libraries: (top) minerals found in aubrite
meteorites, (middle) minerals found in HED meteorites, and (bottom) lunar sample
soils. For discussion see Section 6, page 15.
minerals (approximately 0–5 vol%) chromite and olivine and
accessory minerals (approximately 0–2 vol%) diopside, troilite,
metal, a silica phase, and rare phosphates (Mittlefehldt et al.,
1998). The spectral deconvolution algorithm produces a best fit
using the primary mineral for the diogenite meteorite (low-Ca
pyroxene) with minor amounts of plagioclase and olivine. The
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370
spectral end-members that gave good fits to those targets, did not
give good fits to Mercury data supports our previous work that
indicates Mercury has a distinctly different surface composition.
We will demonstrate, in the next section, that much better fits can
be obtained if we expand the contents of the spectral library to
include more rock-forming mineral phases and minor mineral
phases of more grain sizes than were required for our HED,
aubrite, and lunar spectral fitting discussed in this paper.
7. Results
Fig. 6. Lunar (top) and Mercury (bottom) spectra are fit using the same set of
minerals from the spectral library used to fit several lunar spectra (Donaldson
Hanna et al., 2007). We were not able to achieve a good best fit with the lunar
mineral spectral library. For discussion see Section 6, page 15.
best fit modeled mineral percentage for low-Ca pyroxene and
plagioclase fit within approximately the published data for
Johnstown; however, the olivine abundance is over estimated by
14%. The spectrally derived mineral modes and abundances for
Johnstown fit within the definition of a diogenite meteorite. The
interested reader can find the details of the Bholghati and
Johnstown study in Donaldson Hanna and Sprague (2009).
Our first examples with Mercury are shown in Fig. 5. Spectrum
#071 from Caloris Basin has been ‘‘unmixed’’ using the aubrite
mineral spectral library (top), the HED-Vesta mineral library
(middle), and the lunar soil spectral library (bottom). After many
trials of deconvolution with many different subsets of the spectral
end-member library we finally converge on the BFM. The BFM is
defined in two ways: first by examining the residual spectrum
from the deconvolution UM and using the eye for comparing key
spectral features and second, by minimizing RMS error of the UM
across the spectral region used in the deconvolution for a given set
of spectra permitted for use in the spectral unmixing.
Sometimes the BMF is excellent (determined by visual study of
target spectrum and BFM) and matches most of the spectral
features in the Mercury spectrum. However, sometimes the BFM
results only match a portion or none of the Mercury spectrum (or
other target). In cases of excellent visual match the RMS deviation
over the entire spectrum is lower than in the cases where the BFM
exhibits visible departures from the Mercury spectrum. There are
some occasions when our spectral library does not contain either
the required mineral phases or grain size samples of the mineral
phases or both and we cannot match the target spectrum. In the
cases shown in Fig. 5, no BFM reproduces all of the Mercury
spectral features, although some are reproduced. Most notable is
the match of two emissivity maxima at 8.2 and 8.5 mm in the top
and middle panels. These peaks are those of enstatite. No
combination of lunar soil spectra matches these peaks. Fig. 6
shows a good fit to a spectrum from Grimaldi Basin, also obtained
with MIRSI during the same observing period. We used a library of
terrestrial mineral laboratory spectra with mineral phases known
to be in lunar rocks and soils. Obviously this spectral library,
which is adequate to model lunar spectra, does not yield a
satisfactory fit to the Mercury spectrum, or to other Mercury
spectra in our data set.
That we obtained good fits to telescopic data from the Moon,
and differentiated meteorites, gives us confidence in the spectral
deconvolution method. That the same spectral libraries with
We report the approximate compositions for radar bright C
(RBC), the dark plains to the west of Caloris Basin (DPWCB), and
Caloris Basin (CB). Intermediate, mafic, and rock compositions are
indicated by Mg-rich chemistry, K-spars, either orthoclase or
sanidine, Na-bearing plagioclase, and high Ca-pyroxene as
dominant phases. Minor Ca- and Mg-garnet, and opaque phases
are also strongly indicated. Opaque phases indicated are rutile
(TiO2), and possibly, perovskite (CaTiO3) in radar bright C region,
and or troilite (FeS) in materials of Caloris Basin and surrounds.
The end-member spectra for ilmenite in our spectral library were
never chosen in any UM.
Obviously, the mineral phase identifications from UM outcomes
can only contain mineral phases for which there are spectra in the
spectral libraries. By repeating deconvolution trials with spectral
libraries containing a wide range of end-members, a pattern of
minerals of a range of compositions and grain sizes becomes
apparent. For example after hundreds of trials the only olivine
spectral end-members chosen were Mg-rich tending toward the
forsterite endmember (Mg2SiO4) of the olivine solid solution
which is represented (Mg, Fe)2SiO4. We present alternative BFM for
the three locations measured. We also provide model ‘‘bad fits’’ as
a basis for discussion of the results of omitting or including certain
mineral phases. Comparison of results permits a visual method of
assessing the goodness-of-fit and the degree to which the results
may be considered definitive or suggestive.
7.1. Radar bright region C
7.1.1. Exploration of orthopyroxenes
Best fit models to RBC repeatedly require Mg-rich orthopyroxene as well as other mineral phases that will be discussed in
Section 7.1.2. We decided to take a detailed look at our spectral
library and determine if we had any end-member available that
would exactly match the spectral features of orthopyroxene in our
Mercury spectrum from this region. We included orthopyroxene
spectral end-members containing a broad range in Mg, Fe, and Ti.
Hamilton (2003) gives an excellent illustrated discussion of
changes in mid-infrared spectral bands as a function of chemical
elemental abundances of major and minor elements. We found
that no one end-member matched all the features in the Mercury
spectrum but several matched some of the features. This is
illustrated in Fig. 7 where the closest matching spectra are plotted
along with the Mercury data from RBC.
Close examination of Fig. 7 shows that the Mercury spectrum
has many features of Mg-rich orthopyroxene, particularly enstatite. This is illustrated by the close match between 10.8 and
11.3 mm in the double peak feature. The grain size has a strong
effect on the 11.7–12.7 mm region where the Mercury spectrum
exhibits the characteristics of small grain size (USGS (usgs) and
RELAB (rel) spectra with grain size o180 and 0–74 mm, respectively). This grain size range is consistent with previous
measurements on Mercury’s surface detailed in the earlier Section
1.1. Our spectral library does not contain the end-member
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371
Fig. 7. Mercury data from RBC are plotted along with a selection of orthopyroxene
spectra from our working spectral library. It is apparent that the double peaked
feature between 10.8 and 11.3 mm is that of Mg-rich orthopyroxene, probably
enstatite (compared to enstatite_710_1000_Hamilton_ASU) but we do not have a
perfect match in both grain size and composition.
orthopyroxene spectra that can match all of the spectral features
likely caused by orthopyroxene mineral phase.
7.1.2. Other mineral phases
Early in the deconvolution ‘‘unmixing’’ we found that the best
matches were always from the small grain size mineral separates
(0–25 mm; 25–63 mm; 0–74 mm; and so on) even though the
coarse grain size separates of the same compositions were
included as end-members. Several mineral phases were repeatedly chosen, including Mg-rich olivine, Mg- and Ca-rich pyroxenes, and Mg- and Ca-rich garnets and potassium feldspar
(K-spar). However, for K-spar, both orthoclase and sanidine were
chosen for different regions. We also found that an opaque phase
was required for the best fits but no one opaque phase was chosen
exclusively. We found that rutile was the significant contributing
opaque phases to the BFM. Perovskite and troilite were also
chosen. Ilmenite (FeTiO3) was chosen in two cases when spectra
from no other opaque were available in the spectral unmixing.
Plagioclase end-members chosen in the deconvolution are
repeatedly Na-rich trending between bytownite and albite.
In an effort to determine our uncertainties and to explore
cosmochemical differences that might be exhibited by the spectra
and the good fits, we ran many models with restricted spectral
end-member options for comparison. For example, we made
many deconvolutions where we permitted only orthoclase or
sanidine, but not both and found the BFM for each. There are only
minor visual spectral differences between the orthoclase and
sanidine spectra in our library (Fig. 8, top panel). For RBC, the best
UM included orthoclase or sanidine 100% of the time, but the
choice did not seem to have an obvious effect on the best UM
spectrum. Orthoclase and sanidine were chosen equally often.
This presented serious ambiguity in our ability to identify which
phase was present and prevents discussing the petrologic
significance of discrimination between the two phases. Also
difficult to visually determine was the effect on goodness-of-fit
caused by opaque phases that might be present at Mercury and
present in our library (Fig. 8, bottom three panels). Every
Fig. 8. (a) JHU orthoclase and sanidine, 0–74 mm grain size fractions; (b) rutile
(TiO2), spectra most often providing best fit for opaque phase; (c) perovskite
(CaTiO3) chosen as opaque in radar bright region C; (d) troilite was chosen when
rutile was not permitted, (e) ilmenite (Fe2TiO3), never chosen, when rutile or other
opaque was permitted.
deconvolution that had a satisfactory fit contained minor amounts
of garnet. The garnets chosen were always the Mg- (grossular) or
Ca- (pyrope) phases (Fig. 9). Almandine, the Fe-rich garnet, was
never chosen although several spectra of different grain size
fractions were in the library.
To explore subtle differences between end-member choices, we
ran five cases of ‘‘unmixing’’ the Mercury spectrum from RBC with
orthoclase as the only permitted K-spar choice. We then repeated
the same five cases with sanidine as the only K-spar choice. The
five cases altered which opaque phase was permitted as a possible
end-member in the solution. We include the fits from work done
for a Lunar and Planetary Science Conference poster (Sprague et al.,
2008). The spectral deconvolutions for that study were obtained
using an 89 member spectral library. These BFM are shown in
cases for each of the sanidine and orthoclase trials. All other UM
were made with a much larger spectral library with 327 members.
The choices for opaques were: (a) rutile included, 89 member
spectral library, (b) rutile (TiO2), (c) ilmenite (FeTiO3) (permitted
but not chosen), (d) perovskite (CaTiO3), and (e) troilite (FeS).
ARTICLE IN PRESS
372
Fig. 10 shows the same Mercury spectrum fit with the options
described for the two cases where the K-spar is restricted to be
orthoclase (left panel) and restricted to sanidine (right panel).
A presentation of the best composition results are detailed in
Table 1 for the cases where orthoclase was required as the K-spar
chosen, if any. Table 2 gives the results for the cases where
sanidine was required if a K-spar was chosen. The restriction on
the K-spar mineral phase does result in slight differences for other
miner phases chosen. The restriction of the opaque phase also
causes slight changes in the amount of the other mineral phases.
However, the relative order of abundance (%) does not change.
This indicates that the fits are quite robust. The names of the
spectral library members can be read from Tables 1 and 2. The
abbreviations of the spectral library members contain enough
information to identify the mineral phase, the grain size fraction
of the end-member spectrum, and the laboratory providing the
spectrum.
7.2. Dark plains west of Caloris Basin
Fig. 9. Magnesium- and calcium-rich garnet spectra from two different grain size
separates are shown to illustrate that both the spectral structure and the grain size
are important to spectral fitting. Note especially the difference in emissivity from
10 to 12.7 mm with small grain size having the highest emissivity.
Spectra and spectral deconvolution UM models from two
independent observations of the DPWCB are shown in Fig. 11 (left
panel, 8 April, #063; right panel, 9 April, #098). Spectra were
obtained on different days and used different standard star
spectra for telluric corrections. This independence provides an
excellent test of our data acquisition, analysis, and modeling. We
ran many deconvolutions of both spectra with a wide variety of
spectral libraries (including lunar and meteorite spectra, including
hydrated minerals and sulfates, restricting spectra end-members
to anhydrous minerals, etc.). All of the BFM had similar endmember choices with some variations in grain size and some
cases mineral phases. As with RBC, we found it instructive to
Fig. 10. (left panel) The Mercury spectrum from over RBC shown (black points) along with unmixing models (grey points) for which K-spar was orthoclase: (a) rutile
permitted, small library; (b) rutile permitted, more end-members added; (c) ilmenite permitted, rutile not permitted, ilmenite not chosen; (d) perovskite and ilmenite
permitted, perovskite chosen; (e) troilite and ilmenite permitted, troilite chosen. (right panel) The Mercury spectrum from over RBC is shown (black points) along with
unmixing models (grey points) for which K-spar was sanidine: (a) rutile permitted, small library; (b) TiO2 rutile permitted as opaque phase; (c) FeTiO3 ilmenite permitted,
rutile not permitted, ilmenite not chosen; (d) perovskite and ilmenite permitted, perovskite chosen; (e) troilite and ilmenite permitted, troilite chosen.
ARTICLE IN PRESS
Table 1
#046 Radar bright C (RBC) best-fit model (BFM) test results for orthoclase (O-a) to
(O-e) case study corresponding to the left panel of Fig. 10.
(O-a) #046 Radar bright C RMS for BFM
a
LPSC BFM+orthopyroxene to fit 10.6–11.6 mm
Rutile_0–45_15sm—JPL
b
ort_3f_gs0–74—JHU
enstatite_nmnh128288.7447_gs25—USGS
ort_90c_gs0–63—BED
olv_7p_fo66_gs0–74—JHU
hyp2_180c_gs125–250—BED
pyrope_180c_gs0–25—BED
olivine_gds71.16579_fo91—USGS
olivine_ki3189.16976_fo60_gs25—USGS
grossular_garnet_125–500_15sm—JPL
grossular_ws485.8753_gs171—USGS
0.0024
oAbundance (%)
37
19
10
10
9
7
3
2
1
1
1
Total
100
(O-b) #046 Radar bright C RMS for BFM
Rutile+more end-members added
orthoclase.3f—JHU
enstatite 0–45—JPL
rutile 45–125—JPL
oligoclase 90c 63–90—BED
enstatite_0_45_nmnh128288.7447—USGS
hyperstene180c 25–63—BED
orthoclase 90c 0–25—BED
olivine_0_60_ki3054.16862 (Fo66)—USGS
bytownite green 180c 63–125—BED
pyrope 180c 63–125—BED
grossular_0_65_nmnh155371.8586—USGS
0.0022
oAbundance (%)
28
15
11
10
10
9
7
5
4
0.7
0.3
Total
100%
(O-c) #046 Radar bright C RMS for BFM
Ilmenite permitted, not chosen
hedenbergite.1f—JHU
orthoclase.3f—JHU
bytownite green 180c 25–63—BED
andesine hs142.2101—USGS
pyrope.1f—JHU
enstatite_0_45_nmnh128288.7447—USGS
grossular_0_210_ws484.8698—USGS
0.0033
oAbundance (%)
38
26
14
8
7
3
4
0.4
Total
100
(O-d) #046 Radar bright C RMS for BFM
Perovskite permitted, chosen
olivin7f_emiss_lib—JHU
diopside_0–45_15sm_emiss_lib—JPL
enstat1c_emiss_lib—JHU
orthoclase_nmnh142137.17360_gs0–74—USGS
PV-EAC-002_MagnetCove_perovskite_gs0–45—RELAB
PV-EAC-020_Tapira_perovskite_gs0–45—RELAB
emiss_hyperstene180c0_25—BED
orthoclase_125–500_15sm_emiss_lib—JPL
PV-EAC-024_Afrikanda_perovskite_gs0–45—RELAB
emiss_avg_orthoclase90c0_63—BED
0.0023
oAbundance (%)
25
16
15
12
11
7
4
4
4
2
Total
100
(O-e) #046 Radar bright C RMS for BFM
Troilite permitted, chosen
enstat1f_emiss_lib—JHU
EA-EAC-001C_HeatedCanyonDiablo_troilite—RELAB
orthoc3f_emiss_lib—JHU
enstatite_0–45_15sm_emiss_lib—JPL
olivin9p_emiss_lib—JHU
hypers1c_emiss_lib—JHU
orthoclase_hs13.17296_gs74–250—USGS
pyrope1f_emiss_lib—JHU
0.0018
o Abundance (%)
34
23
18
8
4
4
3
3
373
Table 1. (continued )
olivine_ki3054.16862_Fo66_gs25—USGS
emiss_grossular180c125_250—BED
2
1
1
Total
100
The computed statistical error between the best-fit model (BFM) and the Mercury
spectrum over the entire spectral interval is defined as the root mean square (RMS)
and is provided for each case.
a
LPSC model for (a) refers to Lunar and Planetary Science Institute Conference
Abstracts: Sprague et al. (2007).
b
Spectra are named to provide identification of the mineral phase, grain size,
and laboratory. For example, the second entry, orthoclase model (O-a),
ort_3f_gs0–74—JHU can be interpreted as: orthoclase sample 3fine with grain
size separates of 0–45 mm obtained from the JHU spectral library. The constituents
are listed in relative order of upper limit percent abundance that resulted from the
BFM of a particular case, such as O-a, O-b, O-c, O-d, etc.
explore the chemistry of the dark plains by restricting the endmember choices to elucidate which mineral phases were essential
to good fits and which were either non-essential or clearly
not present on the surface. The BFM cases are shown as case (a) in
Fig. 11 both left and right panels. A clear trend of Na-rich
plagioclase (albite, oligoclase, bytownite), Ca-, Na, and Mg-rich
clinopyroxene (augite), Mg-rich orthopyroxene, and pyrope, the
Mg-rich garnet phase is present.
It is obvious that the only good fits are those of case (a) for both
spectra from the dark plains. The names of the spectral endmembers can be read from Tables 3 and 4. As with Tables 2 and 3,
the abbreviations of the spectra end-members contain enough
information to identify the mineral phase, the grain size fraction
of the end-member spectrum, and the laboratory providing the
spectrum.
We learn something about Mercury’s composition by studying
the poor fits (b–d). An opaque phase is required for a good fit to
the data. Only rutile gives a good fit. When ilmenite is chosen at
28% (see Table 3, case (c)), the fit is poor. Also, pyrope
Mg3Al2(SiO4)3 is required for a good fit to the spectra (compare
panels (a) and (b) in both DP #063 and DP #098). The presence of
these magnesium aluminum silicates in the dark plains units is
surprising in that there is no related evidence for deep excavation
similar to the region in RBC. Never the less our spectral
deconvolution of the Mercury spectra are BFM with pyrope
included as a spectral end-member. We have determined that
both the spectral peaks and grain size in the garnet spectrum
chosen by the deconvolution BFM is important. Fig. 11 case (a) in
both left- and right-hand panels are best fit with the inclusion of
pyrope which provides the spectral peak at 10.5 and 11.4 mm
without providing too much emissivity in the 11.5–12.7 mm region
where the spectral emissivity is declining in the Mercury spectra
for both #063 and #098 (compared to the same spectral region in
Fig. 9). Examination of Tables 3 and 4 show that the 75–250 mm
grain size pyrope spectrum is chosen in the BFM rather than the
0–74 mm grain size pyrope spectrum because the continuum
height is too great in the smallest grain size separate to mix with
the other constituents of the BFM and retain the continuum
height of the Mercury spectrum. The larger grain size separate
provides the peak contributions without raising the spectrum
excessively in the 11.5–12.7 mm spectral region.
7.3. Caloris Basin
The mineralogy of CB is unlike that of RBC and the DPWCB.
While there are some common mineral phases chosen, the basin
infill is distinct with a significant abundance of potassium
ARTICLE IN PRESS
374
Table 2
#046 Radar bright C BFM and test results for sanidine (S-a) to (S-e) case study
corresponding to the right panel of Fig. 10.
(S-a) #046 Radar Bright C RMS for BFM
LPSCa post+only orthopyx to fit 10.6–11.6
labradorite_0–45A_15sm—JPL
b
olv_7f_fo66_gs0–74—JHU
enstatite_nmnh128288.7447_gs25—USGS
rutile_hs126.19708_15sm—USGS
hyp2_180c_gs125–250—BED
san_180c_gs0–25—BED
andesine_hs142.2101_gs74–250—USGS
lab_1c_gs75–250—JHU
olivine_ki3189.16976_fo60_gs25—USGS
grossular_garnet_45–125_15sm—JPL
olivine_gds71.16534_fo91—USGS
0.0030
oAbundance (%)
21
15
15
13
11
10
8
2
2
1
1
Total
100
(S-b) #046 Radar Bright C RMS for BFM
Rutile+more end-members added
sanidine.1f—JHU
hypersthene_0_45_nmnhc2368.10362—USGS
rutile hs126.19708—USGS
bytownite 90c 0–63—BED
enstatite nmnh128288.7447—USGS
labradorite 45–125A—JPL
olivine ki3054.16862 (Fo66)—USGS
grossular hs113.8529—USGS
0.0023
oAbundance (%)
31
14
13
12
11
6
6
4
3
Total
100
(S-c) #046 Radar Bright C RMS for BFM
Ilmenite permitted, not chosen
hedenbergite.1f—JHU
sanidine.1f—JHU
grossular hs113.8529—USGS
enstatite nmnh128288.7447—USGS
olivine gds71.16638 (Fo91)—USGS
labradorite 90c 63–90—BED
pyrope.1f—JHU
0.0028
oAbundance (%)
34
26
16
10
5
4
2
0.9
0.5
Total
100
(S-d) #046 Radar Bright C RMS for BFM
Perovskite permitted, chosen
sanidi1f_emiss_lib—JHU
olivi10f_emiss_lib—JHU
olivin7f_emiss_lib—JHU
diopside_0–45_15sm_emiss_lib—JPL
enstatite_0–45_15sm_emiss_lib—JPL
grossular2f_emiss_lib—JHU
enstat1c_emiss_lib—JHU
PV-EAC-022_IceRiverComplex_perovskite—RELAB
pyrope1c_emiss_lib—JHU
grossular_garnet_45–125_15sm_emiss_lib—JPL
0.0022
oAbundance (%)
24
19
18
10
6
6
5
5
4
3
1
Total
100
(S-e) #046 Radar Bright C RMS for BFM
Troilite permitted, chosen
enstat1f_emiss_lib—JHU
sanidi1f_emiss_lib—JHU
TB-TJM-050_ParagouldChondrite_troilite—RELAB
emiss_sanidine180c125_250—BED
olivi12c_emiss_lib—JHU
grossular_garnet_0-45_15sm_emiss_lib—JPL
0.0019
oAbundance (%)
47
17
16
9
6
2
2
2
grossular_ws485.8753_gs171—USGSS
1
Total
100
a
LPSC model for (a) refers to Lunar and Planetary Science Institute Conference
Abstracts: Sprague et al. (2007).
b
Another example of how to read the spectral nomenclature: olv_7f_fo66_gs0–74—JPL may be interpreted as olivine number 7f of chemistry
corresponding to forsterite 66%, fayalite 54%, grains size 0–74 mm.
feldspar (best fit with orthoclase). Fig. 12 illustrates seven UM for
spectrum #071 and 5 UM for spectrum #080. The sequences (a–g)
in the left panel explore the best fit minerals including garnet,
while restricting the K-spar to either orthoclase or sanidine, and
permitting amphibole. The sequence (a–e) in the right-hand panel
of Fig. 12 tests the BFM for inclusion or exclusion of garnet,
orthoclase, sanidine, rutile, and ilmenite. Inclusion of orthoclase,
garnet, and rutile result in the lowest RMS and BFM for both #071
and #080. In addition, for #071, the presence of hornblende
results in the BFM of Fig. 12a. The identification of orthoclase
supports the ground-based observations of Sprague et al. (1990)
who observed enhanced column abundances of atomic potassium
in Mercury’s exosphere over Caloris Basin. Also indicated in the
UM is the presence of amphibole, a mineral that has water of
hydration as part of its lattice structure. Two examples of
amphibole whose spectra were indicated in UM are hornblende,
(Na, Ca)2(Mg, Fe, Al)5(SiAl)8O22(OH)2, and glaucophane, the
Na-rich end-member of the glaucophane-riebeckite solid solution
Na2(Mg, Fe, Al)5Si8O22(OH)2. Hornblende is an important and
widespread rock-forming mineral common in igneous plutonic
rocks on Earth. It can vary greatly in composition and is found in
basic igneous rocks often associated with significant amounts of
titanium (up to 10% TiO2 in analyses). That the TiO2 occurs
without the association of Fe is consistent with our results for all
of the UM from CB where rutile is the only TiO2 phase indicated.
Hornblende occurs especially in diorites and syenites. This is
particularly interesting because Emery et al. (1998) found spectral
similarities from measurements of Mercury’s surface made from
the Kuiper Airborne Observatory over this region that were
suggestive of alkali syenites or diorite. For the glaucophaneriebeckite solid solution series, substitution of ferrous and ferric
iron for magnesium and aluminum occurs depending on available
elements in the melt. These two mineral phases are logical
members of mineral suites that are expected in Caloris Basin
which has a highly fractured floor from extensional tectonics
described in detail by Strom et al. (1975), Melosh and McKinnon
(1988), Watters et al. (2005) and the presence of more than 230
linear troughs part of the extensive system of radial graben
discovered in high-resolution MDIS imaging (cf. Murchie et al.,
2008, Watters et al., 2008). In addition, multiple episodes of
magmatic instrusion and resurfacing are indicated by wrinkle
ridges concentric to the basin walls and within the interior
(Watters et al., 2009). The radial graben, cracks, and fissures
within the CB may be locations of secondary magmatic intrusion
where, on Earth, such sodium-rich amphiboles have been found.
Alternatively, a second (or more) episode of magmatic intrusion
may have included a hydrous magmatic composition.
We found early in the UM process that Na- and Ca-rich
hornblende and or Na-rich amphibole was included in chosen UM
mineral phases. At first we were worried that this was an
unrealistic mineral to be present in such a hot environment on a
presumably dry planet but hornblende and amphibole are stable
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375
Fig. 11. Two independent spectra from the DPWCB (#063 and #098, left panel and right panel, respectively) are shown along with four spectral unmixes (panels a–d) to
illustrate best-fit choices and chemical composition at this location are shown above. Remarkably, for both spectra, best unmixed minerals include about 10% pyrope, the
Mg-rich garnet phase. Left and right panels: (a) best-fit case—best unmixed minerals; (b) best unmixed minerals, no pyrope; (c) ilmenite permitted, no rutile; (d) no
opaque.
against out gassing or conversion to non-hydrated species at
Mercury’s highest temperature (Berry and Mason, 1959). In
support of this identification are the results of the fast ion particle
(FIPS) measurements (Zurbuchen et al., 2008) which, by modeling
the ratio of the charge to mass of many ions found strong evidence
for OH+ and other water group ions during the first flyby of
Mercury in the vicinity of Caloris Basin. Because no improvement
was shown for Mercury spectrum #080 from Caloris Basin with
either hornblende or spessartine, unlike the case for #071, there
are no corresponding panels (f) and (g) for #080. While
suggestive, the evidence for hydrated mineral species is not
conclusive. The water group ions observed by the FIPS instrument
could be explained by surface chemistry in the regolith induced
by solar wind impact on the surface of Mercury (Gibson, 1977).
More observations in the mid-infrared and by instruments on
MESSENGER and BepiColombo will be required before we can
determine if hydrated minerals are present on Mercury’s surface.
If they are, it must be determined if they are present as original
mineral phases or as altered phases following copius and
perpetual solar wind H and O ion impact and implantation
chemistry (cf. Gibson, 1977) or exogenous material delivered from
meteoroids or comets.
8. Three region summary and brief comments
The results detailed in Section 6 in Figs. 10–12 and Tables 1–6
are now summarized in Table 7 for the orthoclase and sanidine
case study for radar bright region C and Table 8 for the two spectra
from each of the dark plains west of Caloris Basin and Caloris
Basin. Because it is unlikely that Mercury’s surface at any of the
locations measured has exactly the chemical composition and
grain size of the library spectra chosen for the BFM listed in Tables
1–6, we have summed all similar mineral phases into broad
categories. For example, for each Mercury spectrum and for each
UM case, we summed all plagioclase abundance estimates (all
contained Na) and entered the result in the category ‘‘Na-bearing
plagioclase’’. We have retained the UM abundances but point out
that these abundances are upper limits for the reasons discussed
in Section 4 and we expect that there are other mineral phases at
the locations measured for which we had no appropriate spectral
end-member and thus they are missing from the table. The RMS
for each BFM is provided for easy comparison.
The major similarities and differences among the three regions
are: (1) RBC and CB both have relatively high abundances of
K-spar (polymorph phases of KAlSi3O8). K-spar in the hightemperature crystalline structure of sanidine occurs in volcanic
rocks such as rapidly crystallized basalts. This is the preferred
phase of K-spar for radar bright C. The spectra of CB are well
matched with K-spar in the crystalline structure of orthoclase.
Orthoclase is found in a wide variety of igneous and metamorphic
rocks and has crystallized at intermediate to low temperatures.
We suggest both monoclinic K-spars are in the smooth plains of
Caloris Basin—sanidine in rhyolite or other volcanic rocks and
orthoclase in the radial dyke system observed by MESSENGER in
MDIS imaging (Murchie et al., 2008). There is no indication of
K-spar in the mineral suites of the DPWCB, while it is a major
phase in the smooth plains of the basin interior. The relatively
high abundance of K-spar is a likely one cause of the high
reflectance (I/F defined as the observed radiance divided by the
solar irradiance from a normally solar-illuminated Lambertian
disk). Caloris Basin smooth plains material are 15% higher than
Mercury’s hemispherical average (Murchie et al., 2008). (2) RBC
has very high abundances of orthopyroxene and low to no high-Ca
clinopyroxene. The smooth plains within Caloris Basin have
high abundances of high-Ca clinopyroxene and lesser amounts
of orthopyroxene. (3) The dark plains west of Caloris have more
Mg-rich olivine (10%) than the interior of the basin. The olivine
spectral end-members chosen, although Mg-rich, have small
amounts of iron present. The same is true for the orthopyroxene
spectral end-members chosen in the BFM cases (a) and (b) for the
dark plains (see Tables 3 and 4).
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376
Table 3
#063 Dark plains west of Caloris Basin (DPWCB) BFM and test results
corresponding to the left panel of Fig. 11.
Table 4
#098 Dark plains west of Caloris Basin (DPWCB) BFM and test results
corresponding to the right panel of Fig. 11.
(a) #063 Dark plains west of Caloris Basin RMS for BFM
Best unmixed minerals
aug_1f_gs0_74—JHU
rutile_0–45_15sm—JPL
bytownite_0–45_15sm—JPL
pyrope_1f_gs0_74—JHU
ALBITE_0–45_15sm—JPL
hyp2_180c_gs0_25—BED
albite_45–125_15sm—JPL
aug2_180c_gs125_250—BED
olv2_180c_gs0_25—BED
oligoclase_hs110.16239_gs74–250—USGS
0.0049
oAbundance (%)
23
14
11
10
9
8
8
6
5
4
1
Total
(a) #098 Dark plains west of Caloris Basin RMS for BFM
Best unmixed minerals
aug2_180c_gs0_25—BED
hyp_1c_gs75_250—JHU
pyrope_1c_gs75_250—JHU
olv_9f_fo88_gs0_74—JHU
lab_90c_gs0_25—BED
olv_12c_fo91_gs75_250—JHU
olv_180c_gs63_125—BED
0.0059
oAbundance (%)
22
15
11
10
10
10
10
4
3
3
2
1
0.04
100
(b) #063 Dark plains west of Caloris Basin RMS for BFM
Best unmixed minerals, no garnets
olv2_180c_gs0_25—BED
aug_1f_gs0_74—JHU
aug2_180c_gs125_250—BED
0.0086
oAbundance (%)
26
20
17
14
13
6
4
1
Total
100
Total
100
(b) #098 Dark plains west of Caloris Basin RMS for BFM
Best unmixed minerals, no garnets
hyp_1c_gs75_250—JHU
0.0121
oAbundance (%)
23
19
14
13
12
6
5
3
2
2
(c) #063 Dark plains west of Caloris Basin RMS for BFM
Ilmenite permitted, no rutile
aug_1f_gs0_74—JHU
ilmenite_hs231.11112_gs74–250—USGS
ilm_180c_gs25_63—BED
0.0141
oAbundance (%)
32
28
17
14
5
4
Total
100
Total
100
(c) #098 Dark plains west of Caloris Basin RMS for BFM
Ilmenite permitted, no rutile
0.0204
oAbundance (%)
43
29
13
11
2
1
(d) #063 Dark plains west of Caloris Basin RMS for BFM
No opaque
aug_1f_gs0_74—JHU
0.0133
oAbundance (%)
61
17
13
7
1
Total
100
Total
100
(d) #098 Dark plains west of Caloris Basin RMS for BFM
No opaque
0.0204
oAbundance (%)
43
29
11
2
1
Total
100
9. Discussion
9.1. Possible presence of microscopic iron metal blebs (npFe0)
Our results indicate more olivine in the dark plains region than
in Caloris Basin. This is consistent with the lower reflectance (I/F)
measured by MDIS on MESSENGER during the first flyby
(Robinson et al., 2008). Even though the olivine chosen in the
BFM is Mg-rich, some olivine sample spectra contained small
amounts of FeO. It is possible that the dark plains west of Caloris
are affected by the conversion of FeO in the olivine to npFe0. Noble
et al. (2007) found in detailed laboratory experiments that for the
visible and near-infrared, small amount of npFe0 embedded in
silica gels may cause changes in albedo and spectral slope. The
most important factors in predicting the result is the size of the
embedded npFe0 particles and the concentration. Some combinations of npFe0 size and concentration only darken the spectrum
with no change in spectral slope; other combinations both redden
and darken the spectrum. Comparison of asteroid, lunar, and
Mercury spectra showed that darkening and slope effects differed.
For Mercury, they found that npFe0 is larger and that Mercury
spectra in the visible and near-infrared are systematically lowered
in albedo across the entire visible and near-infrared range but
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377
Fig. 12. Spectra from CB obtained on two different days are shown along with several models to elucidate spectral similarities and differences between them and to
illustrate the changing spectra fits with altering the chemistry of the mineral phases by making restrictions to the spectral libraries. (left panel) CB #071: (a) best fit
minerals with hornblende; (b) best fit minerals, no hornblende; (c) K-spar is sanidine, best fit minerals but no garnet permitted; (d) K-spar is sanidine, ilmenite is
permitted, no rutile permitted; (e) sanidine, all best minerals but no opaque; (f) K-spar is orthoclase, best minerals; (g) best minerals with spessartine instead of pyrope.
(right panel) CB #080: (a) sanidine, best minerals; (b) sanidine, best minerals, no garnets; (c) sanidine, ilmenite, no rutile; (d) sanidine, no orthoclase or opaque; (e)
orthoclase permitted, no sanidine.
with little change in continuum slope. It is possible that the npFe0
in the dark plains is in the size range 410 nm and therefore only
darkens the spectrum without causing a change in the spectral
slope of the MDIS near-infrared multi-band spectrum. In contrast,
the smooth plains in Caloris have a higher (I/F), are rich in K-spar,
and have significantly less olivine, demonstrably less orthopyroxene, and more high-Ca clinopyroxene than the surrounding dark
plains. 5. All regions apparently have some Ca- and Mg-rich garnet
phases contributing to the spectral signatures. Alternatively, there
could be some other mineral phase missing from our spectral
library that exists on Mercury and exhibits the emissivity matched
by pyrope and/or grossular. Clearly, the discovery of ubiquitous
garnet in Mercury’s surface materials would be a major chemical
discovery. While the spectral evidence is highly suggestive, we
cannot definitively prove its presence because our mineral
spectral end-members are incomplete and spectra are not
obtained in a heated vacuum environment simulating the
conditions of Mercury’s daylight surface. A reanalysis of our data
when such spectra become available will likely clarify this issue.
9.2. General trends from these data and others
The general trends of compositions for the three regions
measured are very clear with Mg-, Ca-, Na-, Ti-rich mineralogy
dominating. K-spar is a major phase in two of the three regions
measured. The obvious lack of mineral phases containing larger
than a few %FeO is consistent with all prior observations (cf. Vilas,
1988 and references therein, Warell et al., 2006, Warell and
Blewett, 2004) and recent spectroscopic and multi-band spectral
imaging of Mercury’s surface during the first MESSENGER flyby
(McClintock et al., 2008; Robinson et al., 2008), respectively.
While the trend to Mg-rich orthopyroxenes, Ca-rich clinopyroxene, Na-bearing plagioclase, TiO2, and K-spar is notable for these
regions, at the same time there is clearly heterogeneity among the
mineral phases and relative abundances.
Strong heterogeneity from location to location on Mercury’s
surface has previously been demonstrated by other mid-infrared
spectral measurements. For the sake of brevity, results from
several previous observations and observing groups have been
ARTICLE IN PRESS
378
Table 5
#071 Caloris Basin (CB) BFM results corresponding to the case studies in the left
panel of Fig. 12.
(a) #071 Caloris Basin RMS for BFM
Sanidine, best minerals+hornblende
san_3f_gs0_74—JHU
hornblende_nmnh117329.10137—USGS
olivine_gds70.16357_Fo89_gs115—USGS
0.0080
o% Abundance
31
18
12
11
10
6
5
5
1
1
Total
100
(b) #071 Caloris Basin RMS for BFM
Sanidine, best minerals
aug_1f_gs0_74—JHU
hypersthene_0–45_15sm—JPL
san_3f_gs0_74—JHU
albite_gds30.359_gs74–250—USGS
0.0074
o% Abundance
20
16
14
10
7
7
7
7
6
5
1
1
Total
100
(c) #071 Caloris Basin RMS for BFM
Sanidine, best minerals, no garnets
san_3f_gs0_74—JHU
albite_gds30.359_gs74–250—USGS
aug_1f_gs0_74—JHU
rutile_0-45_15sm—JPL
0.0104
o% Abundance
35
13
11
10
9
8
7
7
1
Total
100
(d) #071 Caloris Basin RMS for BFM
Sanidine, best minerals, +ilmenite, no rutile
san_3f_gs0_74—JHU
aug_1f_gs0_74—JHU
oligoclase_hs110.16239_gs74–250—USGSS
0.0204
o% Abundance
37
30
24
8
1
Total
100
(e) #071 Caloris Basin RMS for BFM
Sanidine, best minerals, no opaque
san_3f_gs0_74—JHU
aug_1f_gs0_74—JHU
0.0208
o% Abundance
39
30
22
8
Total
100
(f) #071 Caloris Basin (CB) RMS for BFM
Orthoclase, best minerals, no sanidine
0.0154
o% Abundance
26
25
orthoclase_nmnh142137.17360_gs0–74—USGSS
albite_gds30.359_gs74–250—USGSS
grossular_ws484.8698_gs210—USGSS
19
18
7
3
2
1
Total
100
(g) #071 Caloris Basin RMS for BFM
No pyrope, or grossular permitted, but spessartine permitted
san_3f_gs0_74—JHU
aug_1f_gs0_74—JHU
spessartine_2f_gs0_74—JHU
glaucophane_0–45_15sm—JPL
olivine_gds70.16357_Fo89_gs115—USGSS
0.0115
o% Abundance
27
21
19
12
7
6
5
3
1
Total
100
The computed statistical error between the BFM and the Mercury spectrum over
the entire spectral interval is defined as the root mean square (RMS) and is
provided for each case.
organized according to longitude region, latitude, and suggested
composition in Table 9. References to the original observations
and analyses are noted. The results presented in this paper have
been added to Table 9 for completeness.
The range of compositions shown in Table 9 can be found in
extrusive igneous rocks, exposed intrusive dykes, sills, and domes,
plutons, impact melts, and material excavated from depth,
possibly as deep as the crust–mantle interface. In other words,
unlike the Moon, which is primarily anorthositic plagioclase,
Fe-, Mg-, Ca- and Na-bearing pyroxenes with lesser Mg- and Feolivine and up to 20% ilmenite (Taylor, 1992), Mercury displays a
wider and more complex chemical composition. Jeanloz et al.
(1995) concluded that Mercury did not experience widespread
basaltic volcanism whose magma originated deep within the
mantle. Instead, they suggested that the magma source for
exogenic resurfacing was shallow in origin, highly differentiated,
and likely to be alkali basalts and possibly, alkali-rich feldspathoids. Indeed, spectral observations of Mercury’s surface
found good spectral matches to alkali basalt (Sprague et al., 1994)
and some spectral features that were indicative of alkali syenites
in spectra obtained from the Kuiper Airborne Observatory (KAO)
while looking at 200–2601E longitude, near equatorial and midlatitude regions from east to west—Beethoven, Ts’ao Chan, Balzac,
Budh Planitia, to Tolstoj (Emery et al., 1998). However, in RBC, the
DPWCB, and within CB, no sodalite or nepheline end-member
spectra were chosen in any UM despite their presence in several
different grain size fractions in the spectral library. We therefore
conclude that the surface rock types in RBC, and DPWCB, are not
under-saturated in silica as is typical of feldspathoids. However,
the indication of low-intermedicate silica content (no or minor
olivine) and high alkali content in CB are suggestive of magma
that is close to under-saturated in silica. This striking regional
heterogeneity can be explained by different source magmas.
The modeled primary opaque phase on Mercury’s surface at
the three locations measured is rutile (TiO2), with possible
perovskite (CaTiO3) in the RBC. This is demonstrative of the
degree to which Mercury’s crustal material is iron-depleted
compared to the Moon and Earth where ilmenite is common in
basalts and other rocks. The abundance of rutile indicated in the
ARTICLE IN PRESS
Table 6
#080 Caloris Basin (CB) BFM results corresponding to the case studies in the right
panel of Fig. 12.
(a) #080 Caloris Basin RMS for BFM
Sanidine, best minerals
san_3f_gs0_74—JHU
aug_1f_gs0_74—JHU
enstatite_0–45_15sm—JPL
albite_hs143.524_gs74–250—USGSS
grossular_2f_gs0_74—JHU
sanidine_0–45_15sm—JPL
hypersthene_nmnhc2368.10362_gs0–74—USGS
grossular_2c_gs75_250—JHU
0.0104
o% Abundance
18
12
10
9
8
7
7
7
5
5
4
3
3
2
(b) #080 Caloris Basin RMS for BFM
Sanidine, best minerals, no garnets
san_3f_gs0_74—JHU
aug_1f_gs0_74—JHU
albite_hs143.524_gs74–250—USGS
0.0065
o% Abundance
29
14
10
9
9
8
7
7
6
Total
100%
(c) #080 Caloris Basin RMS for BFM
Sanidine, best minerals, ilmenite permitted, no rutile
aug_1f_gs0_74—JHU
ilm_180c_gs25_63—BED
0.0132
o% Abundance
29
22
22
18
6
2
2
Total
100%
(d) #080 Caloris Basin RMS for BFM
Sanidine, best minerals, no opaque
aug_1f_gs0_74—JHU
san_3f_gs0_74—JHU
enstatite_0–45_15sm—JPL
0.0137
o% Abundance
36
20
16
12
9
7
Total
100%
(e) #080 Caloris Basin RMS for BFM
Orthoclase permitted, no sanidine
rutile_0–45_15sm-JPL
orthoclase_0–45_15sm—JPL
grossular_2c_gs75_250—JHU
ort_90c_gs90_125—BED
0.0054
o% Abundance
23
16
11
11
9
8
8
7
6
1
Total
100%
The computed statistical error between the BFM and the Mercury spectrum over
the entire spectral interval is defined as the root mean square (RMS) and is
provided for each case.
379
UM (from 10% to 25%) is about the same as the abundance of
ilmenite in lunar basalts (Lewis, 1997). This abundance may be
overestimated because of the absence of spectra with the correct
composition or grain size in the spectral library. Laboratory
spectroscopy of rock chips demonstrates that Reststrählen
band structure is lost in albite and anorthosite following highpressure shock (details can be found in Johnson et al., 2002;
Johnson and Horz, 2003). Mercury’s regolith undoubtedly
contains some shocked material. We included the spectra of
shocked plagioclase feldspars, anorthosites, and pyroxenites but
they were not chosen. It may be because the spectra were from
chip surfaces and not finely sized particles. This puzzle may be
pursued in a future deconvolution when spectra of minerals
having undergone high-pressure shock are available at a range of
grain sizes.
The spectral shape of rutile is the important factor in its choice
for the BFM cases. As shown in Fig. 8 panels (b–e), the opaque
mineral phases have different spectral slopes in the mid-infrared
spectral region. The steep slope to lower emissivity of the rutile
spectra is a critical factor in matching the Mercury spectra. Could
the decreasing spectral emissivity in the Mercury spectra from
about 10 to 13 mm that is satisfied by the inclusion of rutile in the
BFM be an artifact of an improper thermal continuum removal?
We reexamined the way in which we correct the Mercury spectra
for the thermal continuum (using methods of Emery et al. (1998)
and Henderson and Jakosky (1997)) and determined that no such
decrease in emissivity would be caused by our thermal continuum
removal. It is notable that the ilmenite spectra were never chosen
for the BFM unmixes unless rutile was not permitted as a choice
(see Fig. 8). At the Moon, ilmenite (FeTiO3) is the most common
opaque. It is a less likely candidate for an opaque at Mercury
because of its Fe content and because spectra lack the drop in
emissivity from 10 to 12.7 mm.
Spectra from all regions are best fit with the inclusion of a
small amount (1–10%) of Mg- or Ca-rich garnet (pyrope or
grossular, respectively) with the exception of one case for CB
when neither pyrope nor grossular spectra were permitted as
choices (case g, #071) and the Mn-rich garnet spessartite
(Mn3Al2(SiO4)3) was chosen. Garnet is an important cosmochemical marker and may give important constraints for the formation
of Mercury and its evolution to the planet whose surface we
see today. It remains to be seen if the interpretation of small
amounts of ubiquitous garnet remains robust following more
observations and interpretation with spectral libraries prepared
for Mercury’s vacuum and temperature environment. It may be
that the MASCS spectral data will be able to elucidate this finding.
It is possible we will have to wait until the planned observations
by the thermal emission spectrograph MERTIS on BepiColombo to
verify this tentative discovery, or to discover some other mineral
phase that is providing the spectral shape required to fit the
Mercury data that is provided by Ca- and Mg- and Mn-rich garnets
in our spectral libraries. The same could be said of the spectral
shapes of plagioclase, pyroxene, and olivine. We think that the
identification of garnet is likely to be correct. The suggestion of
Ca- and Mg-bearing garnets in RBC is not be too surprising. The
heavily cratered, ancient material exhibits a roughness indicated
of deep excavation and overturn which very likely exposed
lithologies formed deep in the crust, or even at the mantle–crust
interface. On Earth, garnets are found in igneous rocks resulting
from episodes of extrusive volcanism that occurred on rapid
timescales as a result of some accompanying geophysical
disturbance such as impact or faulting. In these cases, there is
not time for dissolving garnets and subsequent equilibration with
the rising magma column. In the case of terrestrial potash
rhyolites, Mn-garnet is a minor phase (see more discussion below
in Section 9.3).
ARTICLE IN PRESS
380
Table 7
Upper limits for categories of mineral phases summed from BFM spectral choices listed in Tables 1 and 2 at radar bright region C (RBC) for orthoclase and sanidine cases
with different options for opaque phase. RMS 102 vales are given.
Radar Bright C (RBC) #046 orthoclase cases
Radar Bright C (RBC) #046 sanidine cases
a
b
c
d
e
a
b
c
d
e
102 RMS of BFM
Na-bearing Plagioclase
High Ca-clinopyroxene
high Mg-orthopyroxene
High Mg-olivine
K-spar
Garnet
Hornblende
0.24
0
0
17
12
29
5
0
0.22
14
0
34
5
35
1
0
0.33
15
38
18
0
26
3
0
0.23
0
16
19
25
18
0
0
0.18
0
0
46
6
21
5
0
0.30
31
0
26
18
10
1
0
0.23
24
0
25
4
31
3
0
0.28
11
34
20
2
26
6
0
0.22
0
10
16
37
24
10
0
0.19
0
0
56
4
23
3
0
Opaque phases below
Rutile
Perovskite
Troilite
37
0
0
11
0
0
0
22
0
0
0
23
13
0
0
13
0
0
0
0
0
0
0
0
0
4
0
0
0
16
Cases a, b, c, d, e, correspond to the cases in the left and right panels of Fig. 9.
Table 8
Upper limits for categories of mineral phases for a and b cases of BFM for dark
plains west of Caloris Basin (DPWCB) summed from Tables 3 and 4, and for Caloris
Basin (CB) summed from Tables 5 and 6.
Table 9
Comparison of results presented in this manuscript to historical mid-infrared
spectral observations.
Reference
Region
Dark plains (DPWCB)
Caloris Basin (CB)
Mercury file #
#063
#071
BFM cases
a
b
a
b
a
b
a
b
102 RMS of BFM
Na-bearing Plagioclase
High Ca-clinopyroxene
High Mg-orthopyroxene
High Mg-olivine
K-spar
Garnet
Hornblende
0.49
29
29
8
4
0
10
0
0.86
31
5
26
13
0
0
0
0.59
27
22
11
6
0
10
0
1.2
26
12
14
28
0
0
0
0.80
17
31
0
1
18
6
5
0.74
18
34
16
1
10
8
0
0.39
15
27
10
4
17
10
0
0.65
16
15
14
7
29
0
0
#098
#080
1
2
3
4
5a
6
7
Opaque phases below
Rutile
Perovskite
Troilite
19
0
0
26
0
0
25
0
0
19
0
0
22
0
0
14
0
0
17
0
0
18
0
0
Results from two spectra for each region are shown. RMS 102 values are given.
Cases correspond to top two cases from left and right panels of Figs. 10 and 11.
8
9b
10
9.3. Suggested rock types and formation conditions
Gillis-Davis et al. (2009) have given considerable evidence for
volcanic vents and rocks over large portions of Mercury’s surface.
Most notable are the vents and volcanic features within and
around CB. Thus, it is of interest to discuss examples of rock types
that result from differentiated primary Mg-rich olivine basalt
magmas. Some examples of those differentiated in volcanic
provinces are trachybasalts and oligoclase basalts. If the
primary low-Fe olivine basalt magma is differentiated and
undergoes plutonic crystallization, alkali gabbros and diorites
may result. Rocks expected from primary granidiorite form, by
differentiation and volcanic crystallization, dacites, latites andesites, and rhyolites.
Given the possibilities described above, we examine the
specifics of Tables 7 and 8 and the summary of historical
observations in Table 9. Tyler et al. (1988) obtained mid-infrared
spectral measurements of Mercury and found a Christiansen
feature at 7.8 mm, a wavelength for an emissivity maximum that is
consistent with intermediate rock types such as rhyolite (Logan
10
East longitude
region
3201, 3141
338–3481
2501
240–2501
3301
295–3501
200–2601
240–2501
201S–201N
27517101
18017101,17017101
35017101, 21517101
201S–201N
275–3151
252–2921
0–301N
150–1601
125–1351
50–751N, 1101
50–751S, 166–2501
50–751N, 166–2501
Radar
0–301N
bright C
100–1401
Dark plains 0–301N
124–1551
10 Caloris
Basin
20–451N
152–1801
Reported result
Feldspathic, intermediate rock types
Alkali basalt
Labradorite
Labradorite, 52–61 wt% SiO2;
90% Feldspar and 10% Mg-pyroxene
Not modeled
Possible alkali syenite
90% Labradorite and 10% enstatite
44 wt% SiO2
45–57 wt% SiO2
50 and 42 wt% SiO2
52–61 wt% SiO2
Labradorite and clinopyroxene
Pyroxene (unspecified)
46–60 wt% SiO2
46–60 wt% SiO2
No FeO
Orthopyroxene 52–61 wt% SiO2
Clinopyroxene 52–61 wt% SiO2
Enstatite, forsterite, K-spar, labradorite
Ca- and Mg-garnet, rutile/perovskite
Na-plagioclase (albite, oligoclase), K-spar,
enstatite, Mg-rich hypersthene, Ca-, Na-, Alclinopyroxene (augite)
High Ca-, Na-, Mg-clinopyroxene, K-spar,
Na-plagioclase, minor Mg-rich
orthopyroxene, Mg-garnet, possible
amphibole
1
Tyler et al. (1988), 2Sprague et al. (1994), 3Sprague et al. (1997a, b), 4Emery et al.
(1998), 5Sprague and Roush (1998), 6Cooper et al. (2001), 7Sprague et al. (2002),
8
Sprague et al. (2007), 9Warell et al. (2006), 10This manuscript.
a
Improved modeling of Sprague et al. (1994, 1997a, b).
b
0.8–5 mm SpeX observations.
and Hunt, 1970). In fact potash rhyolites are good candidates for
Mercury’s surface based on the typical presence of either
orthoclase or sanidine and soda-rich amphibole and Mg-rich
pyroxene. Potash rhyolites also have green diopside augite and
manganese garnet in the linings of lithophysae (Williams et al.,
1955). Manganese is an effective darkening agent and may explain
the dark plains exterior to Caloris Basin and dark albedo layers in
ARTICLE IN PRESS
surface materials imaged by MESSENGER (Robinson et al., 2008).
While we cannot be certain of the identity of the exact rock types
on Mercury, we have narrowed the possible and likely rock types
to dacite, latite, trachybasalts, tonalites, possible rhyolitic pyroclastics and oligoclase basalts. We consider both extrusive and
intrusive rocks. This is especially justified by the recent discovery
of Pantheon Fossae within Caloris Basin and the implications of
formation by upwelling magma beneath the basin.
9.4. Cosmochemical implications of observations
Taken together, the body of mid-infrared and visible, near-infrared
spectral observations demonstrates that Mercury’s surface, and
probably its crust, is depleted in iron. The results described here
make it possible to answer the question ‘‘has space weathering been
the primary reason ground-based observations in the vis and near-ir
do not see the FeO absorption band?’’ That answer is ‘‘no, Mercury’s
surface is covered in very low-iron Mg- Ca- Na- K- silicates and
alumino silicates’’. The semi-volatile abundance is significant, as
demonstrated by the presence of 14–28% Na- and K-bearing feldspars
in RBC, the DPWCB, and CB. These results are consistent with Warell
and Blewett (2004) from Hapke modeling and Robinson et al. (2008)
who found most Fe to be bound in opaques and not in silicates.
Opaques such as FeS are chemically consistent with iron-free rutile.
Iron differentiation to the core must have been thoroughly
completed by the time of the freezing of the magma ocean if one
existed. If our observations included any primary crust produced
by crystallization from an early magma ocean (cf. Taylor, 1982), we
suggest it would be in the RBC region. If an ancient huge basin is
found to be as large as ground-based observations indicate
(Ksanfomality, 2004), it would lend support to giant impact
theory for an explanation of the large fraction of volume occupied
by Mercury’s presumed iron core (Cameron, 1985, Benz et al.,
1988). This catastrophic event is also consistent with the fact that
Mercury’s surface has retained its semi-volatiles Na, K, discovered
in Mercury’s exosphere by Potter and Morgan (1985, 1986),
respectively. In addition, Na+ and K+, S+ and water group ions were
discovered by Zurbuchen et al. (2008) during the first MESSENGER
fly by. Impact and recondensation models have shown much of all
but the most volatile material (H2O, FeO) will be retained and
recondensed following a large impact where debris is captured in
orbit around Mercury. These details are thoroughly described and
modeled by Benz et al. (2007) who examine the fractionation
induced by a giant impact on the proto Mercury having roughly
chondritic elemental abundances. Some ejected material will be
reaccreted. This accretion, which includes semi-volatile inventory,
can account for the anomalously high mean density of the planet.
It may be that we now have samples of Mercury in our
meteorite collections. Such meteorites will be characterized by
FeO content o1%, Mg-rich orthopyroxene, and olivine, if present.
Plagioclase will be 4Ab10 in modal chemistry. Mg- and/or
Ca- garnet may be present as will be potassium feldspar and
possibly sulfides. These are not entirely new ideas, Love and Keil
(1995) and Burbine et al. (2002) outlined somewhat different
expectations also including low-Fe chemistry, a reducing environment conducive to the formation of sulfides. We suggest that new
searches among our sample stores should look for meteorites
containing high Mg- and Ca-pyroxenes, Mg- and Ca- and possibly
Mn-garnets, Na-rich plagioclase, K-spar, and enstatite or very Mgrich olivine if olivine is present.
10. Caveats
Spectral end-members used in this analysis were primarily
reflectance spectra converted to emissivity by standard techni-
381
ques discussed in detail in Section 4.2 and some emissivity spectra
from the PEL and BED library as well as pyroxenes from ASU. All
spectra were obtained at Earth’s atmospheric pressure. Some
systems were under a N2 purge, or other purging system. Most
spectra were obtained in spectrometers operating with laboratory
air. Logan and Hunt (1970) demonstrated that the wavelength of
the Christainsen feature (emissivity maximum near 7–9 mm) was
changed in shape and possibly in wavelength if spectra were
obtained in a near vacuum rather than at 1 bar atmosphere.
Two of the authors of this paper, (Helbert and Maturilli, 2008)
have shown that some differences exist between spectra obtained
from the same sample at different temperatures, especially in the
region of the Reststrählen bands. Differences in spectral signature
at temperatures relevant for Mercury, 300–725 K need to be
quantified and characterized. Helbert and Maturilli are in the
process of doing so. In addition, a new vacuum chamber, suitable
for housing the sample for induction heating is under fabrication
and assembly. Already, they have obtained the first spectra in the
induction heating experiments and have observed important
changes spectral activity. For example, in stepwise heating of fine
grained quartz separates, an emissivity peak systematically
shifted from 14.4 to 14.5 mm as the temperature was increased
from 200 C to 500 C. For the peak near 12.3 mm, a region where
transparency features are very important in Mercury spectra, the
peak shifted less than 0.05 mm. The central wavelength of the
transparency minimum in a very low FeO labradorite sample
(FeOo0.08%) shifted from 12.1 to 12.2 mm over the same
temperature range. The central wavelength of the transparency
minimum is predictably related to overall bulk SiO2 content in a
rock or mineral powder. Thus, the SiO2% given in Table 9 are
considered approximate given as a range of values based on the
spectral data from Mercury and from the laboratory. However, the
central wavelengths of the transparency minima fall well within
the range of intermediate to mafic rocks. Spectral emissivity
maxima at the Christiansen frequency and in the Reststrählen
region did not shift in wavelength over the same change in
temperature. Therefore, we conclude that our identifications of
the Na-bearing plagioclase feldspars based on the maximum
wavelength and shape of the Christiansen feature is robust. We
expect that our identifications of orthopyroxene, clinopyroxene,
and olivine will also remain valid. There will probably be slight
changes in compositions and relative abundances. We cannot
predict if the differentiation we see between sanidine and
orthoclase will remain robust. Nor can we predict if each BFM
will require pyrope or grossular. However, despite these deficiencies in our present analysis, we believe we have done the best that
we could do in the present era. We look forward to the future
opportunity to test the results in this paper. Once the new PEL
facility is available with its new spectral library of spectra of many
grain sizes, measured while heated to Mercury temperatures and
obtained in a vacuum environment, we will reanalyze these
Mercury spectra and others called out in Table 9. We anticipate
that there will be some new discoveries, some of which may
demonstrate that what we have done here withstand the new
analysis and others do not.
11. Conclusions
1. Mercury’s surface is clearly heterogeneous in composition with
minerals formed from highly differentiated magma coming
from more than one extrusive event. The heterogeneity
strongly argues for more than one source magma. Spectral
unmixing indicates volcanic rocks ranging from trachytes (CB)
to Mg-rich oligoclase basalts (DPWCB) indicative of highly
evolved Mg-rich primary magmas. In both Caloris Basin and
ARTICLE IN PRESS
382
2.
3.
4.
5.
the dark plains west of Caloris Basin, the possibility of garnet
suggests rapid magma flow which carried garnets from
plutonic rocks in timescales that did not permit equilibration,
dissolving of the garnets before reaching the surface.
Evidence for intrusive igneous rocks is indicated by the mineral
suites Caloris Basin. K-spar in orthoclase crystalline structure,
and Ca-, Mg-, Na- pyroxene are characteristic of intrusive dykes
where slower cooling and retention of some water of hydration
may occur resulting in the presence of amphiboles (hornblende
and or glaucophane) as suggested in one case within our data.
At the three locations measured in this study there is little
evidence for feldspathoid-bearing rocks and thus magma
under-saturated in silica. However, the low abundance of
olivine and the high abundance of both orthoclase and Na-rich
plagioclase feldspar indicate magma at or close to the low
silica boundary.
The dark plains west of Caloris and generally forming an
annulus around the Caloris perimeter are younger than the
smooth plains interior to Caloris (Strom et al., 2008) but
systematically darker and their visible and near-infrared
spectra are bluer in slope than those of Caloris (based on
spectra composed of multi-band I/F measurements by MDIS
(Robinson et al., 2008). They provide a counter example to the
lunar space weathering style of ‘‘older is darker and redder’’.
Our measurements indicate that the albedo and slope
differences apparent in the visible and near-infrared are caused
by differences in the mineral content of the regolith. We
support the conclusions of Noble et al. (2007) who conclude
that not all solar system regoliths have the same space
weathering characteristics. They compare laboratory spectra
of npFe0 in silica gels to telescopic Mercury spectra at the same
wavelengths and conclude that npFe0 Mercury’s regolith is
larger in size and lower in concentration than in the lunar soil
and thus direct application of lunar space weathering analysis
is inappropriate for Mercury.
Reanalysis of the same and other data sets with the PEL
spectral library populated by spectra obtained at temperatures
appropriate for Mercury’s surface and cyclical heating and
cooling cycle in a vacuum environment may result in different
BFM. Our goal is to repeat the analysis exclusively using those
data when they become available.
Acknowledgements
The authors of this paper were Visiting Astronomers at the
NASA Infrared Telescope Facility which is operated by the
University of Hawaii under contract from the National Aeronautics
and Space Administration. We are especially grateful to Alan
Tokunaga and Eric Tollestrop for useful engineering time on the
telescope and Don Hunten for helpful discussions. Support for
Sprague, Kozlowski, and Donaldson Hanna was provided by NSF
Grant AST-0406796. Warell received travel support from the
Swedish Science foundation (VR).
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Spectral emissivity measurements of Mercury`s surface indicate Mg

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