European Space Agency European Mars Science and Exploration Conference: Mars Express & ExoMars ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007 MARS SOLAR WIND INTERACTION : FORMATION OF THE MARTIAN CORONA AND ATMOSPHERIC LOSS TO SPACE J-Y. Chaufray 1 , R. Modolo 2 , F. Leblanc 3 , G.M. Chanteur 4 1 Service d’Aéronomie du CNRS/IPSL, Reduit de Verrieres BP3 Route des Gatines 91371 Verrieres-le-Buisson, FRANCE. 2 Department of Physics and Astronomy University of Iowa, 203 Van Allen Hall Iowa City IA 52242-1479, USA 3 Osservatorio astronomico di Trieste, Via Tepolo 11 34131 Trieste, ITALY 4 CETP/IPSL, 10-12, Avenue de l'Europe 78140 Velizy Villacoublay, France [email protected]Section 1: A three dimensional (3-D) atomic oxygen corona of Mars is computed for periods of low and high solar activities. The thermal atomic oxygen corona is derived from a collisionless Chamberlain approach whereas the nonthermal atomic oxygen corona is derived from Monte Carlo simulations. The two main sources of hot exospheric oxygen atoms at Mars are the dissociative recombination of O 2 + between 120 and 300 km, and the sputtering of the Martian atmosphere by incident O + pick-up ions. The reimpacting and escaping fluxes of pick-up ions are derived from a 3D hybrid model describing the interaction of the solar wind with our computed Martian oxygen exosphere. In this work, it is shown that the role of the sputtering crucially depends on an accurate description of the Martian corona as well as of its interaction with the solar wind. The sputtering contribution to the total oxygen escape is smaller by one order of magnitude than the contribution due to the dissociative recombination. The neutral escape is dominant at both solar activities (1x10 25 s -1 for low solar activity and 4x10 25 s -1 for high solar activity) and the ion escape flux is estimated to be equal to 2x10 23 s -1 at low solar activity and to 3.4x10 24 s -1 at high solar activity. This work illustrates one more time the strong dependency of these loss rates on solar conditions. It underlines the difficulty to extrapolate the present measured loss rates to the past solar conditions without a better theoretical and observational knowledge of this dependency. .
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European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
MARS SOLAR WIND INTERACTION : FORMATION OF THE MARTIAN CORONA AND
ATMOSPHERIC LOSS TO SPACE J-Y. Chaufray1, R. Modolo
2, F. Leblanc
3, G.M. Chanteur
4
1Service
d’Aéronomie du CNRS/IPSL, Reduit de Verrieres BP3 Route des Gatines 91371 Verrieres-le-Buisson,
FRANCE. 2Department of Physics and Astronomy University of Iowa, 203 Van Allen Hall Iowa City IA
52242-1479, USA 3Osservatorio astronomico di Trieste, Via Tepolo 11 34131 Trieste, ITALY
4CETP/IPSL,
10-12, Avenue de l'Europe 78140 Velizy Villacoublay, France [email protected]
Section 1: A three dimensional (3-D) atomic
oxygen corona of Mars is computed for periods of
low and high solar activities. The thermal atomic
oxygen corona is derived from a collisionless
Chamberlain approach whereas the nonthermal
atomic oxygen corona is derived from Monte Carlo
simulations. The two main sources of hot exospheric
oxygen atoms at Mars are the dissociative
recombination of O2+ between 120 and 300 km, and
the sputtering of the Martian atmosphere by incident
O+
pick-up ions. The reimpacting and escaping
fluxes of pick-up ions are derived from a 3D hybrid
model describing the interaction of the solar wind
with our computed Martian oxygen exosphere. In
this work, it is shown that the role of the sputtering
crucially depends on an accurate description of the
Martian corona as well as of its interaction with the
solar wind. The sputtering contribution to the total
oxygen escape is smaller by one order of magnitude
than the contribution due to the dissociative
recombination. The neutral escape is dominant at
both solar activities (1x1025
s-1
for low solar activity
and 4x1025
s-1
for high solar activity) and the ion
escape flux is estimated to be equal to 2x1023
s-1
at
low solar activity and to 3.4x1024
s-1
at high solar
activity. This work illustrates one more time the
strong dependency of these loss rates on solar
conditions. It underlines the difficulty to extrapolate
the present measured loss rates to the past solar
conditions without a better theoretical and
observational knowledge of this dependency.
.
European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
MARS SURFACE MAGNETIC OBSERVATORY: A GEOPHYSICAL AND ENVIRONMENT (GEP)
EXPERIMENT FOR EXOMARS, S.Vennerstrom1, M. Menvielle
2, J.M. Merayo
1, S. Schwartz
3, P. Brauer
1,
C. Carr3, G. Chanteur
2, P.A. Jensen
1, B. Langlais
4, M.B. Madsen
5, M. Mandea
6, H. O’Brien
3, N. Olsen
1, S.M.
Pedersen1, F. Primdahl
1, P. Tarits
7, K. Whaler
8,
1Danish National Space Center, Technical University of
An approach to the inversion of the data available
from the MARSIS (Mars Advanced Radar for
Subsurface and Ionosphere Sounding) instrument on
Mars Express is described. The data inversion gives
an estimation of the materials composing the
different detected interfaces, including the impurity
(inclusion) of the first layer, if any, and its
percentage, by the evaluation of the values of the
permittivity that would generate the observed radio
echoes.
The data inversion method is based on the analysis
of the surface to subsurface power ratio and the
relative time delay as measured by MARSIS. The
MARSIS resolution permits us to identify layered
structures present in the subsurface with a depth
resolution of 150 m. A volume scattering and a
multilayer analysis has been performed in order to
analyze the influence of these scattering process on
the obtained results. The data inversion has been
performed at several frequencies to estimate the
frequency dependent parameters affecting the
behavior of the radar echoes.
A preliminary relative calibration has been
performed to determine the capability of MARSIS
to resolve different surface dielectric constants. In
this calibration, based on the estimate of surface
backscattering, the influence of the ionosphere has
also been taken into account. The constraints, due to
the known geological history of the surface, the
local temperature and the thermal condition of the
observed zones and the results of other instruments
on Mars Express and other missions to Mars, have
to be considered to improve the validity of the
utilized models.
The interpretation of radar data require
discrimination between signals arising from
subsurface interfaces and those coming from the
surface topographic features not immediately below
the radar so that the time delay between
transmission and reception is the same (surface
‘‘clutter’’). The main complexity, pertaining to the
data inversion, is related to the accuracy needed on
the values of the dielectric constant on the surface
( ’m(0)), as well as on the accuracy in the radar data
influenced by various causes as, for instance, the
ionosphere residual distortion.
Taking into account that along the orbits the echo
frames exhibit a non stationary behavior, due to the
shape of the surface and subsurface, in order to
obtain a proper inversion, the frames have been
selected only in regions of MARS that are
moderately flat as can be determined, a priori from
MOLA data and by the echoes’ behavior. In this
case, where the surface backscattering is frequency
independent, the echoes should have a shape as
narrow as possible according with the pulse
bandwidth and the weighting network.
The data inversion, taking into account models of
inclusion distribution in the first layer, and data
from SHARAD/MRO that show multilayer structure
of the first layer with higher depth resolution,
provides a solution, in terms of determination of the
dielectric constant of the subsurface, compliant with
the knowledge accuracy of the surface scattering.
The obtained dielectric constants are higher than
those pertaining to the material confined by the
extreme models considered possible by geologists
and their values show an unexpected compatibility
with a presence of liquid water mixed with solid
material.
ANNUAL CHANGE OF MARTIAN DDS-SEEPAGES. Horváth, A. (1, 2, 3), Kereszturi, Á. (1, 5), Bérczi, Sz. (1, 4), Sik, A. (1, 5), Pócs, T. (1, 6), Gesztesi, A. (3), Gánti, T. (1), Szathmáry, E. (1,7)
(1) Collegium Budapest (Institute for Advanced Study), 2 Szentháromság, H-1014 Budapest, ([email protected]); (2) Konkoly Observatory, H-1525 Budapest Pf. 67; (3) Budapest Planetarium of Society for Dissemination of Scientific Knowledge, H-1476 Budapest Pf. 47, ([email protected]); (4) Eötvös University, Dept. G. Physics, Cosmic Mat. Sp. Res. Gr. H-1117 Budapest, Pázmány 1/a.; (5) Eötvös University, Dept. of Physical Geography, H-1117 Budapest, Pázmány 1/c; (6) Eszterházy Károly College, Dept. of Botany, H-3301 Eger Pf 43, ([email protected]); (7) Eötvös University, Dept. of Plant Taxonomy and Ecology, H-1117 Budapest, Pázmány 1/c; Hungary.
Introduction: The signs of surface water found by MGS
(on MOC images [1]), Mars Odyssey (neutron data [2]) and Mars Express (spectral data, [3]) play important role in understanding surface processes – especially probable life forms – on Mars. There are signs of recent liquid water on Mars like the gullies formed probably during high obliquity [1, 4, 5] and dark slope streaks which could be formed by gravitational mass movements or water seepage [6, 7, 8].
We discovered and analysed a possible third group of phenomena presumably produced by liquid water on the surface, called DDS-seepage. These are originated at dark dune spots (DDS). (Dark dune spots appear in the defrosting surface in late winter–early spring in the polar regions of Mars [9, 10]).
Most of the DDS-seepages can be found at the steep slopes of the dark dunes in craters and the intercrater areas and we could study not only great number of these seepages [11, 12] but also could observe their changes from one Martian year to the other.
Fig. 1 The crater where we studied the dunefield and DDS-seepages. The frame refer the belt of Fig. 2b (MGS MOC image)
Data and methods: The DDSs and the DDS-seepage
structures were identified visually on images from the MGS MOC and measured manually with Surfer software, the topograthic data were from MGS MOLA measurements. The maximal error of the morphometric results is 30%.
The surface studied is about 41 square kilometers where there we found 750 dark dune spots and 440 DDS-seepage formations.
We analyzed a crater (coordinates: 150.8°W, 69.2°S and diameter ca. 70 km, Fig. 1) based on two images of the same region in spring, but with one martian year difference (E07-
00808 and R07-00938; Fig. 2), almost in the same phase of the seasonal cycle of the DDS-phenomenon.
Fig. 2 MGS images of the same locality from 2001 and 2003 with DDS-seepages on the slopes. Enlarged view of the frames with more details about the seepage-flows are given in Fig. 4
Morphological charateristics and annual change of
DDS-seepages:. The dark and grey streaks from these DDS’s suggest that the frosted layer has been partly or totally defrosted (Fig. 3a, b, 4a, b).
The main characteristics of the DDS-seepages are: • the dark streaks originate from DDS (Fig. 3a-d, 4a, b), • based on MOLA data they point downslope away from
DDSs (Fig. 3a-d -see arrow, 4a, b), • slope having angles between 18– 31 degrees (Fig. 3a-d), • most streaks become narrower at the foot of the hill (Fig.
3a, c, 4a, b),
Lunar and Planetary Science XXXVI (2005) 1128.pdf
• at their lower end a spot indicates that the downflown material has accumulated there (ponds, Fig. 3b, d),
• the darkness of the streaks is variable (Fig. 3a-d, 4a, b), • the phenomenon annually appears on Mars (Fig. 4).
Fig. 3 Enlarged view of Fig. 4 frames where we can observe the main characteristics of the DDS-seepages. Arrow shows flow direction to all imeges (a-d)
Concluding model of the DDS-seepage: According to our
earlier model the DDS forming defrosting process contains possible biological components [11, 12]. For these biological components (the Martian Surface Organisms MSOs) the defrosting process cycle begins in spring when the MSOs begin their activity and help enhance the melting of water. The molten water seepage starts flowing downwards between the ice cover and the frozen soil. First the grey color exhibits the thinner frosted layer, later the final dark color of the DDS exhibits the naked surface of the dark dunefield [13].
Summary: The morphology and annual occurence of the
DDS-seepages on slopes were studied. Our results suggest the temporal presence of liquid water on polar dune surfaces bellow the CO2 frost cover. This could be one of the few current examples of liquid water on Mars.
The water-related model of the DDS-seepage phenomena gives better interpretation of the observed slope features than the dust avalanche model [6] because of 1) the presence of ponds, and 2) the overwhelming majority of DDS-seepages narrows towards the lower end of the streak.
Our result are consonant with the water ice detected by Mars Express next to the CO2-frost, and agree partly with the suggestions of other authors on the possible presence of liquid water on Mars today [10].
Acknowledgements: Authors thank for the use of MGS MOC
images of NASA and Malin Space Science Systems [14]. The ESA ECS-project No. 98004 is highly acknowledged.
Fig. 4 Annual changes of DDS-seepages. (a) 2001-08-13, (b) 2003-07-13, (c) combination of 2001 positive and 2003 negative images
References: [1] Malin, M. C. and Edgett, K. S. (2000) Evidence for recent groundwater seepage and surface runoff on Mars, Science 288, 2330-2335. [2] Boynton, W. V. et al (2002), Distribution of Hydrogen in the Near-Surface of Mars: Evidence for Subsurface Ice Deposits, Science 297, 81-85. [3] Bibring, J.-P. et al (2004) Perennial water ice identified in the south polar cap of Mars, Nature 428, 627-630. [4] Costard, F., Forget, F., Mangold, N., Peulvast, J. P. (2002) Formation of Recent Martian Debris Flows by Melting of Near-Surface Ground Ice at High Obliquity, Science, 295, 110-113. [5] Christensen, P. R. (2003) Formation of recent martian gullies through melting of extensive water-rich snow deposits, Nature 422, 45-48. [6] Treiman, A. H. (2004) Martian slope streaks and gullies: origin as dry granular flows, Lunar Planet. Sci. XXXIV, #1323. [7] Miyamoto, H., Dohm, J. M., Beyer, R. A., Baker, V. R. (2004) Fluid dynamical implications of anastomosing slope streaks on Mars, Journal of Geophysical Research, 109, E6, CiteID E06008. [8] Motazedian, T. (2003) Currently Flowing Water on Mars, Lunar Planet. Sci. XXXIV, #1840. [9] Edgett, K.S., Supulver, K. D. and Malin, M. C. (2000), Spring defrosting of Martian polar regions: Mars Global Surveyor MOC and TES monitoring of the Richardson Crater dune field, 1999-2000, Mars Polar Sci. and Explor. II, #4041. [10] Bridges, N. T.., Herkenhoff, K. E., Titus, T. N., and Kieffer H. H. (2001) Ephemeral dark spots associated with Martian guillies. Lunar Planet. Sci. XXXII, #2126. [11] Horváth, A., Gánti, T., Gesztesi, A., Bérczi, Sz., Szathmáry, E. (2001) Probable evidences of recent biological activity on Mars: appearance and growing of dark dune spots in the South Polar Region. Lunar Planet. Sci. XXXII, #1543, LPI, Houston. [12] Gánti, T., Horváth, A., Bérczi, Sz., Gesztesi, A., Szathmáry E. (2003) DARK DUNE SPOTS: POSSIBLE BIOMARKERS ON MARS? Origins of Life and Evolution of the Biosphere 33: 515-557, Kluwer Academic Publishers, Netherlands. [13] Horváth, A., Bérczi, Sz., Kereszturi, Á., Pócs, T., Gesztesi, A., Gánti, T., Szathmáry, E. (2004) Annual change of outflows from Dark Dune Spots in the Southern Polar Region of the Mars, IV. European Workshop on Exo-Astrobiology (EANA), Great Britain, 22-25 November 2004, Abstract book, p. 91. [14 ] http://www.msss.com/mo_gallery/
Lunar and Planetary Science XXXVI (2005) 1128.pdf
European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
“ODD” MARTIAN DICHOTOMY AND ITS HARMONIC INTERPRETATION. G.G. Kochemasov.
necessary to the origin of life as we know it, and
others are specifically produced by living
organisms. However, these molecules have never
been detected on Mars, either from observations or
in situ space probes. Therefore, relevant questions
related to organics are: are organic molecules
actually present at the surface of Mars; where are
they; what is their concentration; under which form
can we find them.
Indeed, even if endogenous organic molecules
were never synthesized, at least those brought by
exogenous sources, like interplanetary dust
particles, should be present in detectable amount.
Moreover, the track endogenous organic molecules
should not be dropped out because some terrestrial
molecules are known to be able to resist over
periods of several billion years without being
degraded.
It thus appears that organic molecules could be
present at the surface of Mars, even if they have
significant chances to undergo a partial or total
chemical evolution. Within the framework for the
search for organic molecules by present or future
space experiments, we are developing the MOMIE
laboratory experiment (Martian Organic Material
Irradiation and Evolution) in order to determine how
the organic species can evolve at the Martian
surface. We thus propose to implement this type of
research with the assistance of an experimental
setup designed for the study of the behaviour of
organic molecules under conditions mimicking, as
close as possible, the environmental conditions of
Mars surface (e.g. UV radiation, temperature…).
We focused the first part of our study on the
influence of UV radiation on organic molecules
relevant to Mars. We showed that if globally
molecules are destroyed by UV radiations which
should be present at the Martian surface, the
destruction rates differ from a molecule to another
[1]. Moreover, it appears that some species could be
converted into molecules resistant to solar UV. We
present here the results of this study and the
potential influence it could have on the investigation
of the surface of Mars, seeking for organics.
References: [1] Stalport F. et al. (in press), Adv. Space
Res.
European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
Martian Organic Molecules Analyzer (MOMA) Gas Chromatograph (GC) : objectives, principle and
preliminary design C. Szopa1, F. Raulin
2, P. Coll
2, M. Cabane
2, F. Goesmann
3 and the MOMA team.
1Service
d’Aéronomie (SA), University Paris VI, Paris France, 2Laboratoire Interuniversitaire des Systèmes Atmosphériques
(LISA), University Paris XII, Créteil France; 3MPS, Lindau Germany. Contact : [email protected]
The seek for organic molecules (not yet found)
on Mars is of primary importance for Martian
astrobiology because these molecules are closely
connected to life as we know it. With this aim, the
Martian Organic Molecules Analyzer experiment
was pre-selected to be part of the scientific payload
of the Exomars mission. Its goal is to point out the
presence of organic molecules in the soil samples
collected at the surface and sub-surface of Mars by
the rover, and to identify their nature.
MOMA consists of three complementary
analytical sub-systems devoted to detect and
identify a wide range of organic compounds which
could be present in the atmosphere or which could
evolve from soil samples heated and pyrolyzed. The
analysis relies on gas chromatographic and mass
spectrometric measurements.
We present here what could be, at this early
stage of development, the gas chromatograph part
devoted to separate and to bring information for the
identification of organic and inorganic gases. From
this preliminary design, we can estimate the
performances it could reach and give clues on the
scientific return we can expect from its analysis.
DO MEGA IMPACTS LEAVE CRATERS? CHARACTERIZING MEGA IMPACTS AND THEIR RELATION TO THE MARS HEMISPHERIC DICHOTOMY. Margarita M. Marinova1, Oded Aharonson1, and Erik Asphaug2, 1Caltech, 150-21, Pasadena, CA 91125, [email protected], [email protected], 2University of California, Santa Cruz, Earth Sciences Dept., Santa Cruz, CA 95064
Introduction: The most clearly visible feature on
Mars is the hemispheric dichotomy: the difference in elevation (~4 km), crustal thickness (~30 km), rough-ness, and impact crater density between the Northern and Southern hemispheres [1,2]. The depression in the northern hemisphere encompasses ~35% of the planet's surface, equivalent to an average diameter of 7700 km [2]. The dichotomy boundary is expressed both as steep scarps and gentle slopes [2,3,4].
Despite the crustal dichotomy's prominent nature, its formation mechanism remains unknown. The pos-sible formation mechanisms fall in the categories of endogenic and exogenic. For endogenic processes, degree-1 mantle convection is often invoked [e.g. 5]. Exogenic scenarios call for a single mega impact [2] or multiple smaller impacts [6]. If the crustal dichotomy is formed by a mega impact, the impact must not shat-ter the planet or produce sufficient melt to obliterate all surface and crustal evidence of the impact.
We investigate whether the Mars crustal dichotomy may have formed by a single mega impact. This first requires characterizing planetary-scale impacts, which have not been extensively studied; these impacts differ from the thoroughly studied smaller impacts due, in part, to the importance of surface curvature in plane-tary-scale impacts. Due to surface curvature it is ex-pected that material redistribution, and thus melt dis-tribution, would differ from that resulting from small impacts, and the change in crater properties with im-pact angle may be more prominent. We focus on the effect of planetary-scale impacts on early Mars. We compare the results of these simulations to observa-tions to evaluate whether a single mega impact may have formed the dichotomy. Particularly, we investi-gate the depth of penetration of the projectile, the amount of melt produced, and the redistribution of excavated material.
Modeling: We use a fully 3 dimensional Smoothed Particle Hydrodynamics (SPH) model to simulate the impacts. SPH is a Lagrangian model in which an object is represented by particles, where each particle’s mass remains constant, but its size, pressure, internal energy, and density change in response to ex-ternal forces. SPH has been extensively used for simu-lating the Moon-forming impact [7]. The 3 dimen-sional nature of the code allows the simulation of im-pacts at any impact angle. In our simulations we nomi-nally use 200,000 particles, giving a resolution (parti-cle diameter) of about 115 km. The semi-empirical
Tillotson Equation of State (EOS) is employed [8]. Figure 1 shows a snapshot of a simulation of a 60 deg impact (measured from the horizontal).
Figure 1. Snapshot of an impact simulation: t = 25 min after impact. Impact parameters: v = 6 km/s, Dimpactor = 860 km, Eimpact = 1.45x1029 J, Dcrater ~ 8000 km, impact angle = 60 deg.
Planet Initial Conditions. Mars’ initial pressure
profile in the simulation is set to hydrostatic. In order to be able to calculate melt production, we require a realistic initial internal energy profile. We assume the surface and core-mantle boundary temperatures from parameterized convection models [9], and impose an adiabatic compression heating profile in the planet to obtain the mantle and core internal energies. Early Mars is likely to have had a convecting mantle and core, resulting in an adiabatic profile. The bulk materi-als for the mantle and core are taken to be olivine and iron, respectively.
Equation of State: The proper implementation of initial conditions requires using the appropriate mate-rials for the mantle and core. The Tillotson EOS li-brary does not include an olivine-like material, so to match mantle density we create our own olivine EOS. We use the same parameterization and formulation as the Tillotson EOS. Density [10], bulk modulus [11], and heat capacity [12] values were obtained from the literature; all other values were set to the average of available representative materials (basalt, granite, an-orthosite lpp & hpp, andesite). Our model of Mars
matches the known planet radius and mass, and the pressure profile (Pcentral,model = 50 GPa, similar to ref [13]) and core size (rcore,model = 1600 km, within range of ref [14]) are within the expected range.
Depth of Penetration: We calculate the depth of penetration of the deepest 10% of the impactor, which is effectively the depth of the transient impact crater cavity. We consider this depth as it relates to the mag-nitude of gravity waves that are sent through the planet as a result of the impact. That is, a deep depth of pene-tration implies large amplitude waves and significant disruption of the planet’s surface by these waves.
Melting Criteria: We calculate melt production us-ing two criteria: a high pressure criterion and a low-pressure (energy) criterion. In the high pressure melt-ing criterion, material shocked above its threshold pressure melts upon decompression. This is commonly referred to as pressure melting. For olivine and basalt, the pressure melting threshold is ~75 GPa [15]. The low pressure melting criterion is effectively an energy melting criterion. For particles at low pressure (no more than one particle depth into the planet) we as-sume that melting occurs when the internal energy exceeds TmeltCp + Hfusion, where Tmelt is the melting temperature, Cp is the heat capacity, and Hfusion is the heat of fusion for the material. We do not take into account energy melting at depth in the planet and thus we underestimate melt production. We do, however, expect that we take into account all melting occurring close to the planet’s surface, thus we can evaluate the extent of preservation of surface and impact features.
Figure 2. Penetration of deepest 10% of impactor. Colours represent impact angle; rcore = 1600 km, RMars = 3400 km; Eimpact = 1.45x1029 J.
Impacts Parameter Space: We simulate impacts with velocities of 6 to 50 km/s (where 5 km/s is Mars' escape velocity and 50 km/s is twice Mars' orbital velocity), impact angles of 90 (perpendicular to the planet surface), 75, 60, 45, 30, and 15 degrees, and impact energies sufficient to create 4000 to 12,000 km craters (following the gravity regime scaling in [2]).
Results: Figures 2 & 3 show some results on the depth of penetration of the impactor and mass of melt produced for an 8000 km crater impacts, for different impact velocities and angles.
Results indicate that slower and vertical impacts penetrate the deepest, thus producing significant grav-ity waves and the strongest disturbance of the plane-tary surface. At constant impact energy, the impator’s momentum is inversely proportional to the impact ve-locity, thus slower (higher momentum) impacts pene-trate deeper. The smaller depth of penetration of oblique impacts is due to their grazing nature. Thus, faster and lower angle impacts result in less disruption of the planet.
The maximum melt production is for about 15 km/s impacts. In the case of high pressure melting, the low impact velocities (6 km/s) do not generate a shock
wave upon impact since the sound speed in olivine is comparable to these impact velocities, and therefore the high pressure melting production is negligible. As the impact velocity increases, the shock wave is stronger and therefore produces more high pressure melting. At high velocities we interpret the decrease in melt production as due to the decrease in the impactor size and therefore a smaller volume is exposed to the strong shock. In the low pressure (energy) melting criterion, the melt production is generally constant for all impact velocities, as expected for constant impact
Figure 3. Total melt produced (dashed) and melt re-tained on the planet (solid) in terms of a global equiva-lent layer depth on Mars; a few percent of the melt is ejected into space. 10 km depth = 5.1x1021 kg. Eimpact = 1.45x1029 J.
Seventh International Conference on Mars 3354.pdf
energy. However, there is a trend of higher melt pro-duction at low impact velocities, which is due to lar-ger, slower impactors depositing energy over a larger volume of the planet. Since the material is already close to its melting point, this increase in internal en-ergy results in larger melt production.
We can visualize the volume of melt produced as a global equivalent layer (GEL) of a given thickness over the surface of Mars (fig. 3). In these units, a verti-cal impact produces the equivalent of 30-40 km deep melt over the entire planet and oblique impacts (15-30 deg), produce a GEL layer of 5-10 km. While the GEL depths are useful to visualize the total melt volume, they fail to represent the spatial distribution that ulti-mately determines whether surface features are pre-served.
The distribution of melt is a key factor in determin-ing whether a mega impact erases all the evidence of its occurrence. Figure 4 shows snapshots of the distri-bution of melted and non-melted material at 12 min
(a,c) and 2.1 hrs (b,d) after the impact of a 15 km/s impactor at 90 deg (a,b) and 15 deg (c,d). The figures represent slices through the planet. It is seen that in the case of the head-on (90 deg) impact, the depth of pene-tration is down to the core-mantle boundary, there is significant excavation, and the resulting area with a surface melt pool is extensive. In addition, the exca-vated material re-impacts the planet, thus covering much of the surface with melt. A melt pool is formed at the antipode of the impact. In the case of the oblique impact, the depth of penetration is smaller, and the resulting melt pool is more restricted. The simulations shown in figure 4 highlight the difference in resulting crater structure, and melt production and distribution due to the change in impact angle. The simulations show that the resulting melt pool in the vertical impact case covers ~85 deg of the planet’s circumference while for the oblique, 15 deg impact it spans ~35 deg of the planet’s circumference in the downrange direc-tion. Because much of the excavated material reaches
(b) (a)
(d) (c)
Figure 4. Production and redistribution of melt (red and orange); 15 km/s impact at 90 deg (a,b) and 15 deg (c,d); 12 min (a,c) and 2.1 hrs (b,d) after impact. The head on impact produces more melt and a more extensive melt pool than the oblique impact. Eimpact = 1.45x1029 J.
Seventh International Conference on Mars 3354.pdf
escape velocity, no significant amounts of material, including melt, re-impact the planet. Thus, at constant energy and for a given impact velocity, the more oblique impacts produce much smaller melt pools and do not distribute molten material over the planet. A similar trend is also apparent as the impact velocity is increased.
Crustal redistribution is another important con-straint, since current Mars crustal thickness estimates [1] show no crustal thickening at the highlands - low-lands boundary. Figure 5 shows crustal distribution for a 15 km/s impact at 90 deg and 15 deg impact angle at 2.1 hrs after the impact event (same impacts as shown in figure 4). For the vertical impact there is apparent crustal thickening around the crater, while in the oblique impact there is crustal thickening only on the
downrange side of the impact crater and the thickening appears to be less than in the vertical impact case. This limited example shows the significant difference in crustal redistribution as a function of impact angle. Further work is needed to determine the crustal thick-ening from various impacts.
Conclusions: Our simulations provide insight into planetary scale redistribution and melting of crust fol-lowing mega impacts. As a first order observation, at constant impact energy, we note the large discrepancy between vertical and oblique impacts, where the change in impact angle has a more exaggerated effect than seen in smaller (flat surface) impact events. We see that head-on and intermediate velocity impacts produce the largest amounts of melt and disrupt the planet significantly, while the slowest and head-on impacts distribute melt over much of the surface. Oblique and fast impacts produce less melt and disrupt the planet to a lesser extent, thus allowing a signature of the impact to remain. Our results show that mega impacts need not obliterate the evidence of their occur-rence and the possibility of forming the Mars hemi-spheric dichotomy by an impact should be further ex-amined.
(a)
References:
[1] Solomon S.C. et al. (2005) Science 307, 1214-1220. [2] Wilhelms D.E. and S.W. Squyres (1984) Nature 309, 138-140. [3] Smith D.E. et al. (1999) Science 284, 1495-1503. [4] Aharonson O., Zuber M.T. and Rothman D.H. (2001) JGR 106, 23,723-23,735, 2001. [5] Zhong S. and Zuber M.T. (2001) EPSL 189, 75-84. [6] Frey H.V. and Schultz R.A. (1988) GRL 15, 229-232. [7] Canup R.M. and Asphaug E. (2001) Nature 412, 708–712. [8] Tillotson, J. H. (1962) General Atomic, San Diego, California, Report No. GA-3216, July 18. [9] Hauck S.A. and Phillips R.J. (2002) JGR 107, 10.1029/2001JE001801. [10] Klein, Mineral Science, pg 493. [11] T.J. Ahrens (Ed.), Mineral Physics and Crystal-lography: A Handbook of Physical Constants. Am. Geophys. Union, AGU Ref. Shelf 2, 45–63. [12] Hashimoto A. (1983) Geochem J. 17, 111-145. [13] Bertka C.M. and Fei Y. (1998) EPSL 157, 79-88. [14] Yoder C.F. et al. (2003) Sci-ence 300, 299-303. [15] Melosh. H.J., Impact Cratering: A Geological Process. Oxford University Press, 1989.
(b)
Figure 5. Crustal redistribution from a 15 km/s impact at 90 deg (a) and 15 deg (b). The crust (red), impactor (dark blue), mantle (light blue), and core (green) are shown. The excavation of crust (red) and its thickening around the impact crater are apparent. Eimpact = 1.45x1029 J.
Seventh International Conference on Mars 3354.pdf
European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
METNET ATMOSPHERIC SCIENCE NETWORK FOR MARS A.-M. Harri
1, R. Pellinen
1, M. Uspensky
1,
T. Siili1, V. Linkin
2, A. Lipatov
2, H. Savijarvi
3, V. Vorontsov
4, A. Ivankov
4 1Finnish Meteorological Institute,
Helsinki, Finland. 2Russian Space Research Institute, Moscow, Russia.
3University of Helsinki, Finland
4Babakin Space center, Moscow, Russia. Ari-Matti.Harri @fmi.fi / Phone +358 50 337 5623
A new kind of planetary exploration vehicle for
Mars is being developed. The MetNet mission to
Mars is based on a new semi-hard landing vehicle
called Mars Meteorological Lander (MML). The
scope of the MetNet Mission is eventually to deploy
several tens of MMLs on the Martian surface using
inflateable descent system structures. The MML
will have a versatile science payload focused on the
atmospheric science of Mars. Detailed
characterisation of the Martian circulation patterns,
[3] Combes et al. (2006), LPSC XXXVII, Abs. #2010. [4]
Bell J.F. et al. (2004), Science, vol.305, p. 800-806. [5]
Soderblom et al. (2004), Science, vol.306, p. 1723-1726.
European Space Agency European Mars Science and Exploration Conference: Mars Express & ExoMars ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
THE MINIATURIZED MÖSSBAUER SPECTROMETER MIMOS II: FUTURE DEVELOPMENTS FOR EXOMARS AND PHOBOS-GRUNT G. Klingelhöfer
1, D. Rodionov
1,2, M. Blumers
1, L. Strüder
3, B.
Bernhardt4, I. Fleischer
1, C. Schröder
1,5, R.V. Morris
5, J. Girones Lopez
1, G. Studlek
1.
1Johannes Gutenberg
Universität Mainz, Institut Anorganische und Analytische Chemie, Staudinger Weg 9, D-55099 Mainz,
Germany. 2Space Research Institute IKI, Moscow, Russia.
Distributions as Indicators of Magnetic Field Topology
near Mars, J. Geophys. Res., Submitted, 2007.
European Space AgencyEuropean Mars Science and Exploration Conference: Mars Express & ExoMarsESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
A NEW MESOSCALE MODEL FOR THE MARTIAN ATMOSPHERE A. Spiga1 and F. Forget1. 1Laboratoire de Météorologie Dynamique, Université Pierre et Marie Curie, Paris [email protected]
Figure 1. Wind velocity field ~20 m above the surface in the Tharsis region after 2.5 elapsed simulation sols. Ls is 190° (sol 389), local time is ~4h (12h UTC). Note that the downslope winds amplitudes in the vicinity of Olympus Mons (~30-35 m/s) are consistent with e.g. Rafkin et al. results [8]. The model is run with hydrostatic option and includes 50 vertical levels from the surface to ~15 Pa. Horizontal grid is 50x50 with resolution 60 km. Dynamical timestep is 74 seconds. Full physics are included, and computed each dynamical timestep.
Introduction The new mesoscale model developed at Laboratoire de Météorologie Dynamique, aims to simulate Martian meteorology in realistic conditions at finer scales than regular GCMs: transition from large-scale to meso-scale, cyclogenesis and frontology (1000-100 km), mesoscale atmospheric circulation and waves (100-10 km), non-hydrostatic phenomena (10-1 km), and micro-scale circulation (<1 km).Dynamical core The dynamical core (i.e. the way atmospheric fluid dynamic equations are numerically solved) is adapted from the new generation WRF-ARW (Advanced Research Weather Research and Forecasting Model) terrestrial model [1]. Martian physical constants and time conventions are included.The WRF solver uses fully compressible nonhydrostatic Euler equations projected vertically on mass-based terrain-following coordinates [2], and horizontally on an Arakawa C-grid (with different possible map projections on the sphere). The temporal integration is computed with 3rd order Runge-Kutta split-explicit scheme [3], which integrates separately the meteorologically significant circulation and the acoustic modes. Compared to regular leapfrog time-integration schemes, the Runge-Kutta scheme leads to improved numerical stability and accuracy. The dynamical core includes a forward-in-time scheme for tracer dynamics.The model is designed to run idealized and real-case simulations in domains with horizontal resolution ranging from meter to kilometer scales. Several domains can be interactively nested to focus in a particular zone of interest. A gravity-wave absorbing layer at the top of the model is included. Lateral boundary conditions can be periodic, open, symmetric or specified. In the real-case mesoscale simulations, the 3D atmospheric starting state and the specified boundary conditions are interpolated from GCM fields or climatologies by the WRF Preprocessing System (WPS) adapted to Mars. In addition, the adapted WPS can handle any surface dataset at any resolution to initialize the static fields.
Martian physics The whole LMD/AOPP/IAA Martian physics, already used and validated in the LMD-Oxford GCM, are interfaced with the adapted WRF dynamical core. Thus, the resulting Martian mesoscale model features the entire “state of the art” Martian physical model from the LMD-GCM [4,5] : radiative transfer with CO2 gas absorption/emission and dust absorption, emission and diffusion; turbulent diffusion scheme; convective adjustment scheme; soil thermal conduction model; CO2 condensation processes; tracer (water ice, dust, chemical species) transport, dust sedimentation and lifting; microphysics; chemistry; NLTE processes in the thermosphere... The new mesoscale model benefits from the LMD/AOPP/IAA consistent and carefully validated physical representation of the Martian CO2, dust, water and aerosols cycles. In the future, minor adaptations will be required to include the upcoming enhancements of the LMD-GCM physics [6] derived from comparative studies with the recent Mars Express measurements. Adding external physical modules to the model, as well as turning on terrestrial schemes easily tunable to Mars (e.g. planetary boundary layer), will be very easy too. Applications The model can be applied e.g. to help interpreting surface pressure maps derived by OMEGA [7]. More generally, such a tool will enable the Martian community to get insights into a wide range of applications: gravity waves, dust devils studies, polar meteorology, atmospheric dynamics around craters and mountains, landing sites choice for future missions, convective processes, planetary boundary layer and turbulence (Large Eddy Simulations), tracer dynamics, aerosols and microphysics studies, paleo-climates local processes... References [1] Skamarock et al. (2005), NCAR Tech. Note [2] Laprise (1992), MWR 120. [3] Klemp et al. (2007), acc. MWR. [4] Forget et al. (1999), JGR 104. [5] Hourdin et al. (1993), JAS 50. [6] Forget et al., this issue. [7] Spiga et al. (2007), JGR 112 + this issue. [8] Rafkin et al. (2002), Nature 419.
mechanisms of terrestrial planets. Obtained directly
at the surface, in-situ thermal measurements in
planetary regolith allow the determination of the
near-surface heat flow, hence being an important
mean to characterize a planet’s thermal state.
Usually, the heat flow is obtained by combining
two separate measurements: thermal gradient in, and
thermal conductivity of the near-surface soil. In
order to obtain the thermal gradient, a depth
resolved measurement of the soil’s temperature is
needed. On unmanned missions, an instrumented
penetrator is well suited for such measurements.
Despite their importance, in-situ heat flow
measurements have so far only been performed on
the Moon. To estimate the Martian planetary heat
flow, scientists had to rely on indirect methods.
For ESA’s upcoming ExoMars mission, a ‘Heat
Flow and Physical Properties Probe’, the so-called
HP3 instrument, has been proposed as part of the
geophysics payload for the stationary lander
element. The HP instrument package consists of a
mole that carries a package of thermal and electrical
sensors to a depth of five meters.
During descent, sensors on the package will
measure the temperature, the thermal conductivity
and diffusivity, and the electrical conductivity and
relative permittivity of the soil as functions of depth.
After the mole has reached its final depth, the
package will go into a monitoring mode. Together
with the measurement of the thermo physical
properties of the soil, the long term monitoring of
the temperature-depth profile will for the first time
on Mars allow to determine the surface planetary
heat flow which is a key constraint for models of the
Martian volatile cycle as well as for planetary
thermal and habitability evolution models.
However, being an active system, the HP3
instrument inevitably dissipates heat into the soil
during penetration as well as monitoring phase,
thereby itself altering the soil thermal field around
it, first of all the temperature profile.
Hence, to be able to determine the soil’s
undisturbed temperature field, a detailed knowledge
of the instrument induced disturbances on the soil’s
thermal field is vital. The aim of this study is to
develop numerical methods that can be used to
predict these disturbances and to filter them out,
thereby improving the scientific return of the
instrument.
In order to simulate the operative phase from the
start of the penetration phase until the final depth we
use a 2D thermal mathematical model (developed in
ESATAN) including all hardware components and
the soil, with a complex dynamic connection to
simulate the relative motion of the probe in the soil.
To obtain the undisturbed status of the soil column
at a potential landing site (prior to the arrival of the
lander and start of the HP3 operative phase), a 1D
thermal mathematical model of the soil in thermal
equilibrium is used.
The overall goal of the thermal modeling in this
respect is to optimize the instrument’s operational
profile by determining the required duration
between conclusion of a hammering episode and
start of a meaningful thermal measurement. This
duration is essentially influenced by the need to
allow heat conducted into the regolith from mole
and front end electronics (FEE) dissipations to be
transported away, allowing to sense an essentially
undisturbed temperature field with the TEM sensor
suite.
Furthermore, the analysis can serve to introduce
dedicated design measures to minimize the
instrument induced disturbances on the thermal field
around the mole.
Future developments of this work will include
the development of dedicated models to be able to
simulate thermal vacuum tests currently being
carried out at DLR. References: Messina, G. et al. (2006), Thermal Analysis
of HP3, a penetrometer to measure the planetary surface
heat flow, IAC 2006 Conference proceedings.
European Space Agency European Mars Science and Exploration Conference: Mars Express & ExoMars ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007 OMEGA-CRISM CHARACTERIZATION OF MAFIC CRUSTAL COMPOSITION IN THE SYRTIS MAJOR REGION J. F. Mustard1, P. Thollot1, S. L. Murchie2, B. L. Ehlmann1, L. A. Roach1, F. Seelos2, F. Poulet3, J.-P. Bibring3, D. Baratoux4, P. Pinet4, Y. Langevin3, B. Gondet3. 1Dept. of Geological Sciences, Box 1846, Brown University, Providence, RI 02912 [email protected], 2JHU/Applied Physics Laboratory, Laurel, MD 20723, 3Institute d’Astrophysique Spatial, Université Paris 11, 91405 Orsay Cedex, France. 4UMR5562/DTP/OMP, 14, Av. E. Belin, Toulouse, 31400 France
Introduction: The mafic mineralogy of the martian crust records crust forming processes and the composition of melt source regions associated with volcanism [1]. Remotely sensed and landed measurements are dominated by the signatures of feldspar, pyroxene, and olivine and imply that, where exposed, the igneous crust is dominantly basaltic [2, 3]. Thermal infrared data (TIR) show two major divisions in crustal composition. Type I material, predominantly in the equatorial highlands, is basaltic, and Type II, found predominantly in the northern lowland plains, has been variously interpreted to be andesite or basaltic andesite [4], altered basalt with a significant component of hydrolytic weathering materials [5, 6], oxidized basalt [7] or silica-coated basalt [8].
Detailed analysis of OMEGA data in the Syrtis Major region show a diversity of compositions (Mustard et al., 2005; Pinet 2007) [9, 10] and indications of possible layering in the lavas and/or distinct alteration of the upper surface [11] (Baratoux 2007). Here we present the first results for the crustal composition of Mars derived from coordinated analysis of OMEGA (Observatoire pour la Mineralogie, l’Eau, les Glaces et l’Activité) [12] and CRISM (Compact Reconnaissance Imaging Spectrometer for Mars) [13] reflectance observations. [9, 10, 11]. For this initial analysis we focus on the pyroxene mineralogy. This work follows that of Baratoux [11] and Pinet [10]. With the higher spatial resolution of CRISM, we test the hypotheses presented by Baratoux [11] on the alteration of the crust and possible layering of compositions in the Syrtis Major volcanic region.
Datasets and Methods: CRISM is a visible-near infrared (VNIR) and infrared (IR) imaging spectrometer on the Mars Reconnaissance Orbiter (MRO) that can acquire high resolution targeted observations at 544 wavelengths from 0.36-3.92 µm at 15-19 m/pixel and multispectral mapping data with 72 wavelengths at 100-200 m/pixel [13]. We primarily focus on the multispectral observations. Data are processed to account for all instrumental effects and reduced to radiance. From these data, I/F is calculated and then corrected for solar incidence angle and the effects of atmospheric transmission absorptions using an approach similar to that used by the OMEGA experiment [9].
OMEGA is a VNIR and IR hyperspectral imager on the ESA/Mars Express mission [12]. It has a 1.2 mrad IFOV, a spatial sampling that varies from 300 m
(at pericenter) to 4.8 km (at 4000 km altitude), and a 7 to 20 nm spectral resolution in 352 spectral bands over 0.35-5.1 µm. Since entering orbit in January 2004, OMEGA has acquired global coverage between 1-2 km/pixel and high-resolution (<500 m/pixel) coverage for >5% of the planet.
Pyroxenes exhibit two distinct absorptions centered near 1 and 2 µm that result from electronic crystal field transitions of Fe in octahedral coordination [13, 14, 15]. To map the distribution of pyroxene, we use a method based on the Modified Gaussian Model [16]. For both instruments we use the 1.0-2.6wavelength range to avoid problems due to discrepancies in the spectra at the overlap between detectors.
Results The presence of HCP enrichment in the ejecta deposits of some of the craters in Syrtis Major was analyzed by [11]. They argue that this could be due to the presence of HCP-enriched lava flows at depth. Modeling suggests a depth of 300 m. The enrichment of HCP in some ejecta blankets is confirmed by CRISM. Full resolution CRISM observations reveal interesting details of the geology, including excavation of HCP-enriched rocks from beneath a cover of LCP-enriched materials and the complex nature of the Noachian Highland. Furthermore we see HCP enrichment in a number of craters <1 km in diameter. We will continue this analysis to refine the understanding of volcanic rocks in this important region. References: [1] McSween, H. Y. et al. (2003), JGR 108, 10.1029/2003JE002175. [2] Bandfield, J. L. et al. (2000), Science 287, 1626. [3] Mustard, J. F. et al. (1997), JGR 102, 25605-25616. [4] Hamilton, V. E. et al. (2001), JGR 106, 14733. [5] Wyatt, M. B., McSween, H. Y. (2002), Nature 417, 263. [6] Morris, R. V. et al. (2003), Sixth International Conference on Mars, LPI Contribution 3211. [7] Minitti, M. E. et al. (2002), JGR 107, E5, 10.1029. [8] Kraft, M. D., Michalski, J. R., Sharp, T. G. (2003), Geophys. Res. Let. 30, Art. No. 2288. [9] Mustard, J. F. et al. (2005), Science 307, 1594-1597. [10] Pinet et al. (this meeting). [11] Baratoux, D. et. al. (2007) JGR 112 E08S05. [12] Bibring, J-P. et al. (2005), Science 307, 1576-1581. [13] Murchie, S. et al., (2007) JGR, 112, E05S03. [12] Bibring, J-P. et al. (2004), ESA SP 1240, 37. [13] Burns, R. G., Mineralogic Applications of Crystal Field Theory, Cambridge University Press 1970. [14] Adams, J. B. (1974), JGR 79, 4829. [15] King, T. V. V., Ridley, I. (1987), JGR 92, 11457. [16] Sunshine et al. (1990), JGR 95, 6955-6966.
European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
OMEGA/MARS EXPRESS, WATER VAPOUR DAILY VARIABILITY OVER THE SOUTH POLE.
Riccardo Melchiorri1, T. Encrenaz
1, P. Drossart
1, T. Fouchet
1, D. Titov
2, L. Maltagliati
2, F. Forget
3, F. Altieri
4,
G. Bellucci4, Y. Langevin
5, J.P. Bibring
5
1LESIA/OBSPM, France,
2Max Planck, Germany,
3LMD, France,
4INAF - IFSI, Italy,
5IAS, France.
Introduction: The Martian Water cycle is one of
the main cycles that control the dynamic of the
Martian atmosphere. Recent observations has
shown a highly spatial and temporal variability. It is
not yet clear in which proportion these variabilities
are locally produced or if a dynamic of the
atmosphere redistribute them in the atmosphere,
specially concerning the Polar Regions.
The Polar Region is a peculiar and ideal place where
it is possible to observe a variability correlated with
the local time. We report on an daily variation of
water vapour on the south pole region (SPR),
observed by OMEGA/Mars Express during the
south spring-summer period (LS 250°-270°) outside
the CO2 ice cap.
Temperature and pressure taken from the EMCD [1]
model shows values close to the saturation point.
Being the morning temperatures lower than during
the day, it is possible that water vapour condenses
during the night and that it starts to sublimate in the
morning, expanding and redistributing in the
atmosphere.
We have developed a fast method to retrieve the
water vapour content of the OMEGA data, through
the analysis of the 2.6 m band, based on the
assumption that the Water vapour partial pressure is
proportional to the band depth [2, 3].
The totality of the OMEGA [4] orbits taken into
account starts with a lower value of water vapour
than at the end (10-20 pr- m of difference; Fig 1).
OMEGA has been designed to observe the day side
of the planet, which means that in nominal
conditions each orbit starts in the morning.
This phenomenon gives us the possibility to study in
detail the growth of water vapour in the atmosphere
during the day for this period .
Data analysis: This period is characterized by a
maximum of water vapour in the air (reaching 15
ppt- m) and a ground temperature close to the water
saturation. No water ice is spectrally detected on the
ground by OMEGA.
We estimate a quasi constant production of water
vapour of 0.5 ppt- m/hour; 8 ppt- m at 3 AM (local
time) to 18 ppt- m at 6 PM (Fig. 2). Our
observations do not cover the whole day, which
makes impossible to understand if during the
“night” the water vapour locally condenses on the
ground, if it is driven away outside the SPR or if it
condenses again on the CO2 ice cap. However if the
water locally condenses, it should happen in
between 7 PM and 2 AM and should be detectable
by OMEGA. References: [1] Forget F. et al., 1999, J. Geophys. Res.
104, 24155-24176. [2] Melchiorri R. et al 2007, Planetary
and Space Science 55 333–342. [3] Encrenaz, T. et al.,
Figure 1. Arsia (left) and Pavonis (right) Montes.
European Space Agency European Mars Science and Exploration Conference: Mars Express & ExoMars ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007 PRESSURE AND TEMPERATURE CHARACTERISTICS OF POSSIBLE EXOMARS LANDING SITES A. Kuti1, A. Kereszturi2,3. 1Eotvos Lorand University of Sciences Department of Astronomy, H-1518 Budapest, Pf. 32., Hungary, 2Collegium Budapest Institute for Advanced Study, 3Hungarian Astronomical Association e-mail:[email protected]
Introduction: The aim of our work is to develop a method, which is able to approach some macroscopic (in the free atmosphere) and microscopic (inside voids of near surface granular materials) environmental parameters (pressure, temperature, vapor content, possibility of H2O condensation etc.). Such parameters are useful in planning the work and observations of future surface probes on Mars, especially ExoMars. In this abstract only some analysis of TES based p/T conditions are summarized.
Working methods: In the analysis we have
chosen three regions (Amazonis-, Isidis-, Chryse Planitia) for the possible landing sites of ExoMars, which show scientific interest and fit to the engineering constrains too, i.e. they are between 10S 45N latitude and height below 0 m level. Temperature and pressure data were derived from Mars Global Surveyor (MGS) Thermal Emission Spectrometer (TES) measurements [1], using “vanilla” software. Our search has been restricted only to surface observations. We have retrieved data for solar longitudes of 105˚-107˚ (northern hemisphere summer) and 285˚-286˚ (northern hemisphere winter) in the three studied regions. Daytime and night-time data were taken around 2 pm and 2 am, local true solar time.
Results: example curves of the analysis are
visible below.
Figure 1. Temperature (top) and pressure
(bottom) curves for one landing site in summer (left) and winter (right)
As seen in Figure 1, there is a slight increase in
temperature values from south to north, which corresponds well with expectations. Variations of surface pressure feature this area (300˚E-330˚E, 8˚S-
45˚N), although a definite ascent can also be seen towards the northern latitudes. The two distinct curves on the top right panel illustrate the difference between daytime and night-time temperature values at Ls 285˚. These temperature variations are most significant in a ~10˚ interval around the equator, and can be as high as 100 K on the very same latitude. Mid-summer temperatures in the studied southern regions are higher than northern area temperatures. Surface pressure values though are higher in the northern winter, and do not drop below 4 mbars. Surface temperature and pressure parameters of the three potential landing sites are summarized in the table.
lon: 300E-330E lon: 195E-225E lon: 75E-105E
Ls=105-
107 Ls=285-
286 Ls=105-
107 Ls=285-
286 Ls=105-
107 Ls=285-
286
min T [K] 249.43 184.72 237.69 190.8 241.15 203.29
max T [K] 277.95 311.11 278.86 305.06 277.8 304.93 min p [mbar] 3.82 4.32 4.44 4.65 3.09 3.19
day- time
max p [mbar] 6.77 8.1 7.82 8.78 5.61 8.35
min T [K] - 109.71 - 147.58 - 141.51
max T [K] - 209.85 - 207.74 - 216.33 min p [mbar] - 4.36 - 5.01 - 3.02
night-time
max p [mbar] - 8.62 - 9.51 - 8.99
Table Representative p, T values for the possible landing-sites
Conclusion: The predicted and previously
observed p/T parameters are useful for the planning of observations with GEP [2] and of cloud, aerosol and water vapor content with Pancam [3] on ExoMars, as well for detectors on other future probes like the proposed MiniHUM on MSL too [4]. In the next step we are to implement water vapor related parameters and estimate condensation processes, including microphysical predictions inside pore spaces.
References: [1] Christensen, P.R. et al. (2001) J. Geophys.
Res. 106, 23823-23872. PDS Geoscience Node [2] GEP-ExoMars: a Geophysics and Environment Observatory on Mars. J. Biele, S. Ulamec, T. Spohn, D. Mimoun, P. Lognonné, and the GEP team, Lunar and Planetary Science XXXVIII (2007) 1587. [3] Context for the ESA ExoMars Rover: the Panoramic Camera (PanCam) Instrument Andrew D. Griffiths1, Andrew J. Coates, Ralf Jaumann, Harald Michaelis, Gerhard Paar, David Barnes, Jean-Luc Josset and the PanCam team. [5] MiniHUM – a miniaturized device to measure trace-humidity on Mars, D. Möhlmann, First Landing Site Workshop for the 2009 Mars Science Laboratory #45033 (Times New Roman, 9pt.) Smith, J. and J. Doe (2002), JGR 107, DOI:10.1029/2001JE123456. [3] Smith, J. et al. (2003), LPSC XXXV, Abs. #1234.
European Space Agency European Mars Science and Exploration Conference: Mars Express & ExoMars ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007 PSA/PDS DELIVERY OF DIGITAL TERRAIN MODELS AND ORTHOIMAGES DERIVED FROM HRSC DATA Th. Roatsch1, K.-D. Matz1, R. Jaumann1,2, G. Neukum2, D. Heather3. 1Institute of Planetary Research, German Aerospace Center (DLR), Rutherfordstrasse 2, 12489 Berlin, Germany. 2Remote Sensing of the Earth and Planets, Freie Universitaet Berlin. 3ESTEC, Noordwijk, The Netherlands. [email protected] The High Resolution Stereo Camera (HRSC) onboard Mars Express has been operating successfully in Martian orbit for more than 3.5 years (4700 orbits). Images taken during this time period became available to the public through the archives at the Planetary Science Archive (PSA) at ESA and the Planetary Data System (PDS) at NASA. So far, only radiometrically and geometrically calibrated data have been delivered. We now also began delivery of high precision Digital Terrain Models (DTMs) and orthoimages derived from the HRSC
stereo images (Gwinner et al., this conference). The DTM data from the first 6 months of the mission are to be delivered to the PSA and PDS by the end of 2007 and will be made available on both archives as soon as they are validated. Further data will then be delivered on a regular basis to both the PSA (http://www.rssd.esa.int/PSA) and PDS (http://pds-geosciences.wustl.edu/missions/mars_express/).
European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
PULSED REMOTE RAMAN SYSTEM FOR PLANETARY SURFACE EXPLORATION Fernando Rull.
Unidad Asociada UVA-CSIC al Centro de Astrobiología Facultad de Ciencias, Universidad de Valladolid,
in this area but no other carbonate spectral signature.
References: [1] Bibring, J.-P. et al. (2005), Science 307,
1576-1581. [2] Poulet, F. et al. (2005), Nature 438, 623-
627. [3] Roush, T. L. et al. (1986), JGR 102, 1663-1670.
[4] Stockstill, K. R. et al. (2005), JGR 110, DOI:
10.1029/2004JE002353. [5] Morse, J. W. and G. M.
Marion (1999), Am. Jour. of Sc. 299, 738-761. [6] Kahn,
R. (1985), Icarus 62, 175-190. [7] Jouglet, D. et al.
(2007), 7th Mars Conf., Abs #3153. [8] Wagner, C. and
U. Schade (1996), Icarus 123, 256-268. [9] Jouglet, D. et
al. (2007), JGR 112, DOI: 10.1029/2006JE002846.
[10] Gendrin, A. et al. (2005), Science 307, 1587– 1591.
[11] Mangold, N. et al. (2007), JGR 112, in press.
[12] Jouglet, D. et al. (2007), 7th Mars Conf., Abs #3157.
European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
THE SEIS EXPERIMENT : A SEISMIC PACKAGE ON GEP/EXOMARS D. Mimoun
1, D. Mance
2, P.
Lognonne1, D. Giardini
2, W. T. Pike
3, Ulrich Christensen
4, Arie van den Berg
5, P. Schibler
1 and the SEIS team
1IPGP (4 avenue de Neptune, 94107 Saint-Maur cedex, France, [email protected] ),
2ETH (Institute of
Geophysics CH-8093 Zurich), 3Imperial College (Exhibition Road, London SW7 2BT, England ),
4Max-
Planck-Institute for Solar System Research (Max-Planck-Strasse, 237191 Katlenburg-Lindau, Germany), 5Institute of Earth Science (Utrecht University, Budapestlaan 4, 3584 CD Utrecht, NL
Scientific objectives: The SEIS Seismometer
will study the seismic activity of the Planet and
frequency of meteorites impacts. These seismic
events will be characterized by their approximate
distance and azimuth, as well by their magnitude.
The seismometer will also allow also to characterize
shallow and deep interior of the planet, and
especially the water environment as a function of
depth in the deep subsurface, the crustal
thickness of the landing site, the core size and
possibly, if the seismic activity is between the
middle and upper bound of present estimates, the
mantle structure. The sensitivity and noise floor of
the seismometers in the expected Martian
environment are such that the detection of about 20
quakes with Ms magnitude from 4 to 5 and 10-20
impacts per year are expected for a mean model of
seismic activity; our working hypothesis is based on
the thermoelastic cooling of the lithosphere, which
does not consider any tectonic activity possibly
related to volcanoes.
Fig 1. The Seismometer Breadboard (IPGP/CNES/SODERN)
Instrument Configuration: The SEIS
seismometer is based on an hybrid 4 axis
instrument, composed of 2 Very broad Band (VBB)
sensors and 2 Short Period (SP) sensors and has a
mass of about 2200 gr, including all margins. It
includes also highly efficient (24 bits) acquisition
electronics , a deployment system and a wind shield
to allow a deployment outside of he descent module
by the GEP/ExoMars arm. This design reflects a
significant mass reduction compared to design
studied by previous ESA projects (i.e. MarsNet and
InterMarsnet), while offering very little science
return reduction as compared to a more classical 3
VBB +3 SP design.
Fig 2. Instrument architecture
Expected performances
Scientific requirements will be met with a sufficient
signal to noise ratio by the instrument on
- Long term signals : Mars modes
- In bandwidth signals: Marsquakes
- Short Period signals: Asteroid impacts
For each kind of signal, noise ratio are met with a
sufficient margin. The VBB sensor performance (in
red below) is in principle equivalent to a terrestrial
field sensor (STS-2 type), which weights 13 kg. The
black curve presents the target SP noise.
Fig. 3 SEIS noise (Red : VBB self noise Black SP target)
References A1. Lognonné P. & B. Mosser, Planetary
Seismology, 14, 239-302, Survey in Geophysic, 1993.
A2. Lognonné, P., J. Gagnepain-Beyneix, W.B.
Banerdt, S.Cacho, J.F. Karczewski, M. Morand, An Ultra-
Broad Band Seismometer on InterMarsnet, Planetary
Space Sciences, 44, 1237-1249,1996.
A4. P. Lognonné, D. Giardini, B. Banerdt, J.
Gagnepain-Beyneix, A.Mocquet, T. Spohn, J.F.
Karczewski, P. Schibler, S. Cacho, W.T. Pike,C. Cavoit,
A. Desautez, J. Pinassaud, D. Breuer, M. Campillo, P.
Defraigne, V. Dehant, A. Deschamp, J. Hinderer, J.J.
Lévéque, J.P. Montagner, J. Oberst, The NetLander Very
Broad band seismometer, Planet. Space Sc., 48,1289-
1302, 2000.
European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
A Simple Scheme for Batch Processing Atmospheric Corrections of HRSC Colour Images O. J. Stenzel1,
N. Hoekzema1, W. J. Markiewicz
1, H. U. Keller
1 and the HRSC co-investigator team. 1Max-Planck-Institut für
The HRSC camera on board Mars Express has delivered stunning images of Mars for over two and a half years now, greatly enhancing the knowledge about our neighbor planet. However, the study of the Martian surface from orbiter images is hampered by the haziness of the atmosphere; it contains large and variable amounts of aerosols that mainly consist of airborne dust. One should carefully consider the effects of hazes when studying the Martian surface from orbiter images and for many analyses one would like to remove their influence. Our group at the Max Planck Institute for Solar System Research (MPS) is involved in the atmospheric correction of HRSC images since the beginning of the Mars Express mission. We have delivered to the HRSC team a number of tools to estimate the optical thickness of the atmosphere (stereo method, shadow method), and to correct for the contribution of dust (MPAE_ATM_DUST), and dust with high altitude ice (MPAE_ATM_1D). These correction programs work properly for so called ‘IMP aerosols’. The Martian atmosphere however, also contains other types of aerosols, and their properties need to be implemented into the correction routines to optimize the atmospheric correction. To test these in a large number of scenes with different meteorological situations, the optical thicknesses of these scenes need to be retrieved. Current batch processes (MPS_ATM_ST on the above mentioned site) can do this for images where
the Sun is low in the sky. The new scheme presented here is able to estimate the optical thickness independently of solar altitude for all level 3 data. The new scheme has been used with IMP dust (Markiewicz et al., 2002) correction routine MPAE_ATM_DUST to process over 750 colour images composed of the HRSC panchromatic nadir, p1 or p2, green and blue channels. Computation time is about three days on a two processor Intel type machine. Centerpiece of the new batch scheme is the radiative transfer model SHDOM (Evans, 1998). For each scene SHDOM is run iteratively for different values of optical thickness until the albedo of the surface
is within a prescribed range. The resulting is used
for the further correction of the individual channels, scaled appropriately for their absorbance at their
particular wave length. The obtained is not very
accurate but good enough to improve a large number of images and is at this point the only method to estimate the optical thickness for scenes with a high solar altitude. An example of a pair of uncorrected and corrected images is shown in Figure 1.
References: Evans, K. F. (1998), Journal of the Atmospheric Sciences 55. Hoekzema, N., et al. (2007), 7
th international conference on Mars, Passadena.
Markiewicz, W. J., et al. (2002), Adv. Space Res. 29 (2).
Figure 1. Uncorrected and corrected near true colour images from HRSC. The images are composed from the nadir, green and blue channels and converted to CIE RGB colours. The frames were taken on orbit h1266 0000. The center of the images is at 64°N and 115°E. Optical thickness of IMP dust in this scene has been estimated to be =1.9.
European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
Simultaneous Measurements of the Martian Plasma Environment from Rosetta and Mars Express. N. J.
T. Edberg1,2
, A. I. Eriksson2, R. Modolo
2, M. Lester
1, S. W. H. Cowley
1, H. Nilsson
3, R. Lundin
3, S. Barabash
3,
A. Boesswetter4, U. Auster
5, KH. Glassmeister
5, I. Richter
5.
1University of Leicester, University Road Leicester
LE1 7RH, UK. 2Swedish Institute of Space Physics, Uppsala, Sweden.
3Swedish Institute of Space Physics,
Kiruna, Sweden. 4Institute for Theoretical Physics, TU Braunschweig, Germany.
5Institute for Geophysics and
Extraterrestrial Physics, TU Braunschweig, Germany. [email protected]
We present results from simultaneous measurements
of the Martian plasma environment by the Mars
Express and Rosetta spacecraft. In February 2007
Rosetta performed a swing-by of Mars as one of its
four gravity assist maneuvers on its way to the
comet 67P Churyomov/Gerasimenko. The trajectory
of Rosetta during the Mars swing-by made it
possible to observe the solar wind parameters far
upstream of the planet before the actual swing-by.
During Rosetta’s approach and entire flyby Mars
Express was in operation in its orbit around Mars
and thus enabled a two-spacecraft investigation of
the plasma environment. For instance, the influence
of specific solar wind parameters on the Martian
plasma environment could be studied and compared
to simulations. The magnetic pileup boundary and
bow shock were detected almost simultaneously at
two different locations around Mars by the two
spacecraft. The results are compared to previous
investigations based on measurements from the
Mars Global Surveyor mission.
European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
Simultaneous Photoelectron and Ion Measurements in the Martian Ionosphere. R. A. Frahm
1, J. D.
Winningham1, J. R. Sharber
1, R. Lundin
2, H. Nilsson
2, S. Barabash
2, A. J. Coates
3, D .R. Linder
3, A. Fedorov
4,
J. –A. Sauvaud4.
1Southwest Research Institute, 6220 Culebra Road, San Antonio, TX 78228, USA.
2Swedish
Institute of Space Physics, Box 812, Kiruna S-981 28, Sweden, 3Mullard Space Science Laboratory, University
College London, Holmbury St. Mary, Dorking RH5 6NT, United Kingdom, 4Centre d'Etude Spatiale des
Rayonnements, 9 Avenue de Colonel Roche, Toulouse 31028, France. [email protected]
The Analyzer of Space Plasmas and Energetic
Atoms (ASPERA-3) experiment on board the Mars
Express spacecraft conducts measurements of
electrons by the Electron Spectrometer (ELS), ions
by the Ion Mass Analyzer (IMA), and neutral
particles by the Neutral Particle Imager (NPI) and
the Neutral Particle Detector (NPD). While orbiting
Mars, the ELS is able to observe peaks in the
photoelectron spectrum due to photoionization of
carbon dioxide and atomic oxygen by Solar Helium
30.4 nm photons. The source of these peaks in the
photoelectron spectrum is the dayside Martian
ionosphere, with the majority of photoelectrons
created at the exobase where the density is greatest.
A fraction of these photoelectrons are transported to
altitudes of the spacecraft. ELS observes
photoelectron peaks in the Martian ionosphere on
nearly every ionospheric transit.
During the times when the Mars Express
spacecraft traveled through the dayside ionosphere
and ELS observed photoelectron peaks, few ions of
any significance were measured. Due to charge
neutrality arguments, when the photoelectrons are
observed, there must be ions present to balance the
electronic charge. Spacecraft charging is often
observed in the dayside ionosphere which is about
-7V, accelerating the ions into IMA and increasing
the probability that ions would have been detected.
The missing observations of significant ions at
the times that photoelectrons are measured lent
support for adjustments to internal voltage setting
within the IMA. These adjustments were carried
out by ESA in the spring of 2007 and were intended
to increase the sensitivity of IMA in the low-energy
ion range. After these adjustments were made, low-
energy ions are observed in the dayside ionosphere
whenever ELS observes photoelectron peaks. The
combined observations of photoelectron peaks and
low-energy ions in the dayside ionosphere are
highlighted by Figure 1. At times when ELS
observes ionospheric photoelectron peaks, IMA now
successfully observes low-energy ions. In this paper
we plan to interrogate dayside ionospheric cases
where photoelectron peaks are observed during
times of increased IMA sensitivity to identify ions.
Figure 1. Observation of Photoelectrons in the Dayside Martian Ionosphere. Photoelectrons are observed as horizontal
lines in the Energy-Time spectrogram at about 22-24 eV and 27 eV of energy (note that the spacecraft is charged to about -7
V in this Figure). At the same time, low-energy ions are observed.
European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
SOIL PREPARATION SYSTEM AND MULTI-FUNCTIONAL DRILL FOR FUTURE SUBSOIL
SAMPLING ACTIVITIES ON PLANET MARS T. C. Ng1, K. L. Yung
2, P. Weiss
2, W. Leung
3, S. Choi
4.
1Dental Surgeon, Room 1605, Medical Floor, Island Center, 1 Great George Street, Causeway Bay, Hong Kong.
2Department of Industrial and Systems Engineering, The Hong Kong Polytechnic University, Kowloon, Hong
Kong. 3Automation Technology Center, The Hong Kong University of Science & Technology, Clearwater Bay,
Kowloon, Hong Kong. 4COM-X Limited, Suite 1812, 18/F, 113 Argyle Street, Mongkok, Kowloon, Hong
column-density is inferred from the depth of the 2.6-
micron band of H2O (Encrenaz et al., AA 441, L9,
2005 ; Melchiorri et al., Plan. Space Sci. 55, 333,
2007). We selected the same data set as for the
analysis of CO over Hellas (Encrenaz et al., AA
459, 265, 2006). The H2O column density is found
to range from very low or undectable values
(between southern fall and winter) up to about 20
pr-microns during southern spring and summer. The
general behavior of H2O is consistent with the
expected seasonal cycle of water vapor on Mars, as
previously modelled (Forget et al., 1999) and
observed by TES (Smith, 2002, 2004). In particular,
the maximum water vapor content is observed
around southern solstice, and is significantly smaller
than its northern counterpart. However, there is a
noticeable discrepancy around the northern spring
equinox (Ls = 330 – 60 deg.), where the observed
H2O column densities are significantly smaller than
the values predicted by the GCM, as well as the
values measured by the TES instrument, integrated
over longitude.
References:
Encrenaz, T., Melchiorri, R., Fouchet, T. et al. 2005, A
mapping of martian water sublimation during early
northern summer using OMEGA/Mars Express, Astron.
Astrophys. 441, L9-L12.
Encrenaz, T., Fouchet, T., Melchiorri, R. et al. 2006,
Seasonal variations of the martian CO over Hellas as
observed by OMEGA/Mars Express, AA 459, 265-270.
Forget, F., Hourdin, F., Fournier, R. et al. 1999.
Improved general circulation models of the martian
atmosphere from the surface to above 80 km. J. Geophys.
Res. 104, 24155-24176.
Melchiorri, R., Encrenaz, T., Fouchet, T. et al. 2007,
Water vapor mapping on Mars using OMEGA/Mars
Express, Plan. Space Sci. 55, 333-342.
Smith, M. D. 2002. The annual cycle of water vapor on
Mars as observed by the Thermal Emission Spectrometer.
J. Geophys. Res. 107, 1 doi: 10.1029/2001/JE001522,
E11, 5115.
Smith, M. D. 2004. Interannual variability in TES
atmospheric observations of Mars during 1999-2003,
Icarus 167, 148-165.
Figure 2. The water vapor column density over Hellas as
a function of solar longitude, as predicted by the GCM
(Forget et al., 1999).
Figure 1. Red points: the water vapor column density
measured by OMEGA over Hellas as a function of solar
longitude. Blue stars: GCM predictions, extracted from
Fig. 2.
European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
SUBMM WAVE INSTRUMENT (SWI) ON A POTENTIAL EXOMARS ORBITER P. Hartogh
1, P. de
Maagt2.
1MPI für Sonnensystemforschung, Max-Planck-Str. 2, 37191 Katlenburg-Lindau, Germany.
2European
Space Agency, PO Box 299, 2200 AG Noordwijk, The Netherlands. [email protected]
A submm wave sounder concept called MIME
(Microwave Investigation on Mars Express) was
proposed for the Mars Express mission. Based on
MIME, an improved state-of-the-art instrumental
concept has been developed within the framework
of an ESTEC CDF study in order to fit the platform
resources of a potential ExoMars orbiter. The
presentation will address the scientific objectives of
this 10 kg / 50 W class submm wave instrument and
briefly summarize the instrumental specifications.
.
European Space Agency European Mars Science and Exploration Conference: Mars Express & ExoMars ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007 SULFATES ON MARS – FROM THE EYE OF RAMAN-LIBS SYSTEM ON EXOMARS MISSION, Alian Wang. Dept. Earth and Planetary Sciences, McDonnell Center for Space Sciences, Washington University, St. Louis, MO, 63136, USA [email protected]
Recent mission results from Mars – both orbital and landed, have reinforced the importance of sulfates at the surface of Mars as indicators of past geologic environments and as potential hosts for water. Their potential as a near-surface reservoir for water [Vaniman et al., 2004a], especially at mid-latitude and equatorial regions (6-11 wt% from Water-Equivalent Hydrogen, Feldman et al., 2004, 2005), makes this group of minerals extremely important for understanding Mars’ hydrological history. In particular, it is important to understand the exact mineralogy (type of cations & crystallinity), degree of hydration, concentrations, form of deposits, and how to accurately determine these minerals and deposits on the surface of Mars.
Sulfate minerals are especially important record-keepers for the past and current conditions on martian surface and within subsurface, diurnal and seasonal cycles, long-term evolution, and ultimately one of the major records of Mars’ hydrologic history. The hydration state of Mg-sulfates can change rapidly following the changes in temperature (T) and relative humidity (RH) of the environment [Chipera et al.,2005, 2006, 2007, Chou and Seal, 2003, 2005, 2007,Freeman et al., 2007a, 2007b, Vaniman et al., 2004a, 2004b, 2005, 2006, Wang et al., 2006c, 2006d, 2006e, 2007b]. The oxidation state of iron ions in Fe-sulfates will be influenced by the redox condition in the environments where they formed and survived [Morris et al., 2000, Fernande-Remolar et al., 2005]. Cation substitution can occur among different sulfates [Chou et al., 2002]. In a real world, these phase transitions and chemical reactions are dependent upon the structures of starting phases, the kinetics of formation (which can be sluggish), the environment conditions of T & RH variations and the coexisting mineral phases. Even for a pure sulfate, the actual water content is not only determined by its molecular structure, but also controlled by the crystallinity, grain size, and porosity in packing (Wang et al., 2007b).
Because of the ambiguity in some spectral analyses of orbital remote sensing (atmospheric influences, spectral band overlaps) and the instrumentation limits in surface explorations (lack the capability for
definitive identification of sulfates with cations other than Fe and for determination of their hydration states), some discrepancies are found in the publications that report the analysis results of these two sets of data.
Simulation experiments are being conducted in laboratories trying to solve some of these discrepancies. However, the best solution would be on surface exploration with more sophisticated instrumentation. Raman-LIBS system that will be carried by Pasteur rover on ExoMars mission would be one of them. By providing definitive mineral phase identification at molecular level, with the compositional information from the same target, it will open a new window towards Mars surface mineralogy and chemistry, thus to advance our understanding on the surface alteration processes and thus the evolution history of Mars.
Figure 1a &1b compare the spectrum of a mixed hydrous Mg,Ca-sulfate in Vis-NIR spectral range (in OMEGA and CRISM spectral range) with the Raman spectrum of the same sample, in which the individual components in that mixture can be unambiguously distinguished based on their narrow Raman V1 peaks. With chemical composition obtained from the same spot on the sample by LIBS function in Raman-LIBS system, the characterization of this sulfate mixture (types of cations, hydration states, and relative proportions) can be determined.
We are continuing the simulation experiments for sulfates on Mars, and in the same tome developing the synergetic usage of Raman and LIBS spectral data from RLS system.
Acknowledgement: NASA support for Mission of Opportunity for RLS investigation on ExoMars mission. References: Wang et al. (2006), Geochem. Cosmochem. Acta, V70, p6118-6135.
European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
SUMMER OBSERVATIONS OF THE MARTIAN NORTH POLAR RESIDUAL CAP BY THE HIGH
RESOLUTION STEREO CAMERA (HRSC) IN 2004/2005 AND 2006/2007 D. Reiss1, H. Hoffmann
2, F.
Scholten2, H. Hiesinger
1, K.-D. Matz
2, G. Neukum
3
1Institut für Planetologie, Westfälische Wilhelms-
Universität Münster, Wilhelm-Klemm-Str. 10, 48149 Münster, Germany, 2Institut für Planetenforschung,
Deutsches Zentrum für Luft- und Raumfahrt (DLR), Rutherfordstr. 2, 12489 Berlin, Germany, 3Institute of
Geosciences, Planetology and Remote Sensing, Freie Universität Berlin, Malteserstrasse 74-100, 12249 Berlin,
of this type of spectroscopy to planetary exploration.
The detection of pigments sited in microbial matter
in a range of samples from extreme environments
(e.g. Villar et al 2005) has supported development
of the technique for space exploration generally, and
Mars exploration in particular (Perez & Martinez-
Frias 2006). A major advantage of conventional
Raman spectroscopy is that the technique can be
applied to characterising bond types in organic and
inorganic materials.
The characterisation of the organic component of
a sample by Raman spectroscopy is best achieved
when the technique is applied in a microscopy
format, and the organic analyte analysed separately
to the mineral matrix. Analyses can easily be
repeated, adjusting the spot size and depth of focus
until a good quality spectra is obtained. Indeed, this
is the approach usually taken when applying the
technique to Carbonaceous Chondrites for example
(Quirico et al., 2003). Data with a high spectral
resolution can be built up and specific spectral
features mapped. In this way a skilled user can
visually sort through an image and target
components of interest. The automated collection of
data in a spatial context is very powerful and can
identify structures that may be of biological origin
(Pasteris et al., 2002).
Surface Enhanced Raman Spectroscopy (SERS)
can readily provide an increase in the Raman signal
greater than 105 (Etchegoin et al., 2003), and has
been shown to overcome the problems created by
the fluorescence of natural materials (Wilson et al.,
2007 – see fig 1). SERS is achieved by adsorbing
the target analyte onto the surface of a roughened
metal surface, which supports localized plasmons
that can have an extremely large EM field
associated with them. Achieving this effect requires
an extra stage of sample processing, but this can be
performed in a microfludic format. We are
combining the additional sample processing
necessary for SERS with sample preparation also
performed in a microfluidic format (including
extraction and sample concentration stages), but
have yet to fully integrate the separate stages
(including the SERS assay) onto a single chip. The
final result will be a very rapid assay that can be
applied to powdered samples, capable of detecting
ppb concentrations of organic analytes.
Figure 1. Spectra obtained for 100 μmolar concentration
solution of scytonemin in DMSO. a) Raman Spectra with
no scytonemin peaks observed. b) Surface Enhanced
Raman Spectra of a 50 nm concentration of scytonemin
acquired with the aid of silver colloid showing enhanced
peaks characteristic of scytonemin. Excitation laser
wavelength was 532 nm and power 10 mW
The payload for the Pastuer EXO-MARS rover
includes a LIBS-Raman instrument that can perform
Raman Spectroscopy as both a first responder probe
and in a microscopy format. But does not have a
SERS capability that would allow for
characterization and detection of very low quantities
of analyte. It would appear logical for the next
generation of Raman Spectroscopy instruments
deployed on the surface of Mars to possess Tele-
Raman, Micro-Raman and LOC-SERS analysis
capabilities and thus maximise the scientific return
from mass dedicated to monochromatic light
sources and Raman spectrometers.
References: Villar et al., (2005) Analyst 130, 730; Perez
& Martinez-Frias (2006) Spectroscopy Europe 18, 18;
Quirico et al., (2003) Meteor Plan. Sci., 38, 795; Pasteris
et al., (2002) Nature, 420, 476; Etchegoin et al., (2003)
Chem. Phys. Letters, 375, 84; Wilson et al., (2007) Anal.
Chem. 79, 7036.
SURFACE PROPERTIES OF MARS’ POLAR LAYERED DEPOSITS AND POLAR LANDING SITES. A.R. Vasavada1 and K. E. Herkenhoff2, 1Department of Earth and Space Sciences, University of California, Los An-geles CA 90095-1567, USA ([email protected]), 2USGS Astrogeology Team, 2255 N. Gemini Drive,Flagstaff AZ 86001, USA.
Introduction: The landed component of the MarsSurveyor 1998 missions, the Mars Polar Lander(MPL), will reach the planet’s south polar regionalong with the Mars Microprobes on Dec. 3, 1999.The spacecraft will land on the south polar layereddeposits, which partially cover the region poleward of70S latitude, and will conduct the first in situ obser-vations of the polar subsurface, surface, and atmos-phere. Like on Earth, the polar regions of Mars arestrongly influenced by seasonal and climatic cycles,and are ideal sites for landed experiments.
The location of MPL’s landing site is limited byatmospheric entry constraints to a latitude of 75+/-2degrees. This latitude range overlaps a contiguous,dissected plateau of layered deposits known as UltimiLobe between 170W and 230W longitude [1]. West of205W, Ultimi Lobe forms a broad plateau with eleva-tions up to ~2 km above the surrounding cratered ter-rain. Elevations gradually decrease east of 205W. Be-cause the area is unexplored at the lander’s scale,properties and processes at that scale can be inferredonly from remote sensing or theoretical results. Inanticipation of the landed mission, here we review thederived surface properties of the southern layered de-posits, and present new determinations of surfacethermal inertia.
Surface Thermal and Optical Properties: D. A.Paige and colleagues have used Viking InfraredThermal Mapper (IRTM) 20-micron measurements toderive thermal inertias poleward of 60S latitude [2].Thermal inertia measures the thermal response of asurface layer to variations in incident energy, and isgiven here in SI units. The results are representativeof the surface down to the diurnal skin depth (a fewcentimeters). We have derived new thermal inertiamaps in a similar fashion to [2], but also includedimportant corrections for Mars’ radiatively active at-mosphere [2,3].
Results indicate that all surfaces poleward of 70Slatitude--excluding the residual ice--are characterizedby very low thermal inertias of ~75-125. These valuesimply that the near-surface is fine-grained, and free ofice and rocks. An apparent particle size of ~10 mi-crons can be inferred from laboratory thermal con-ductivity measurements of well-sorted glass beads atrelevant atmospheric pressures [4].
An analysis of surface color and albedo indicatesthat bright red dust appears to be the major non-volatile component of the layered deposits, possiblyalong with a minor dark component [5]. There is littledetectable color difference between the layered depos-its near the pole and the surrounding cratered terrain,perhaps indicating that a continuous mantle overliesboth units. The composition of the near-surface layeris uncertain. If it is a layer of typical atmospheric dust,an additional cementing agent is probably necessary tosupport observed scarp slopes of up to 20 degrees, andto prevent removal of the material by wind [6].
Dark Dune-Forming Material: Dark, dune-forming material is distributed over both polar re-gions. In the north, dark material is closely associatedwith erosional scarps in the layered deposits [7]. Thedark, north polar sand sea has very low derived ther-mal inertias near ~75 [8]. In the south, the dark mate-rial appears topographically trapped within depres-sions on the deposits and within impact craters on thesurrounding terrain. Although not well-resolved inthermal inertia maps, the dark material in the southprobably has a similarly low inertia.
The dark material’s low inertia can be reconciledwith its apparently sand-sized grains if it is composedof either basaltic ash fragments or aggregates of a mi-nor, dark dust component of the layered deposits thatforms as a sublimation residue [8, and referencestherein]. Such material may be confined to the ob-served low-albedo patches, or perhaps may be morewidely distributed if under a thermally unimportantlayer of bright dust.
Surface Roughness: In Viking images of thesouthern layered deposits with spatial resolutions>100 m/pixel, the smooth surface of the broad plateaunear 75S and 200W-230W is interrupted only by lowrelief, E-W striking ridges and the rims of partiallyburied impact craters. Ridge slopes are ≤10 degrees asindicated by images taken at low sun angles. Regionswhere the deposits are very thin or absent have km-scale roughness typical of the underlying cratered ter-rain.
At resolutions <100 m/pixel, the surface of thesouthern layered deposits displays considerable tex-ture. Grooves, flutes, and pits have been noted in theanalysis of Mariner 9 images, suggesting mechanicalerosion most likely from wind [9].
SOUTH POLAR SURFACE PROPERTIES: A. R. Vasavada and K. E. Herkenhoff
Summary: Much of the south polar region hassimilar color, albedo, and thermal inertia. The conti-nuity in color and albedo can be explained by thewidespread presence of a few microns of bright dust[5]. However, the thermal inertia results are repre-sentative of a layer at least a few centimeters thick.Accordingly, the south polar region may be mantledby at least a few centimeters of typical Mars dust.However, we speculate that the erosion (sublimation)of the southern layered deposits produces low-inertiamaterial similar to the dark, low-inertia materialthought to form from the sublimation of the northerndeposits. Perhaps such material covers much of thesouth polar region under a thin coating of bright dust.Even if the dark material in the south is confined onlyto observed low-albedo patches, its thermal propertiesare probably similar to those of dark material in thenorth and to those of non-polar dust mantles.
The possibility that a dust mantle or sublimationlag covers the southern layered deposits raises thequestion of whether landed spacecraft will be able toaccess the “pristine”, presumably volatile-rich layereddeposits. The thickness of the surface layer is highlyuncertain. If a sublimation lag, its thickness may beself-limited to the length scale of either vapor orthermal diffusion. Meter-thick, local concentrations ofeolian bright or dark material could also inhibit the
lander’s access to the layered deposits. Unfortunately,these issues cannot be addressed with currently avail-able data.
The MPL’s landing site will most likely be ice-free and relatively rock-free compared to areas such asthe Viking and Pathfinder landing sites. Regionalslopes appear not to pose a major hazard. Rather it issmaller features such as the grooves and texture visi-ble at the ~10-m scale that may be hazardous.
Acknowledgements: The 1-D surface-atmosphere model used to derived thermal inertiaswas developed by David Paige. Pierre Williams, Na-than Bridges, and Deborah Bass helped with imageanalysis. Our ideas have been refined through discus-sions with Bruce Murray and Ron Greeley.
References: [1] Tanaka K. L. and Scott D. H.(1987) U. S. Geol. Surv. Misc. Invest. Map, I-1802-C.[2] Paige D. A. et al. (1994) JGR, 99, 25,993-26,013.[3] Haberle R. M. and Jakosky B. M. (1991) Icarus,90, 187-204. [4] Presley M. A. and Christensen P. R.(1997) JGR, 102, 6551-6566. [5] Herkenhoff, K. E.and Murray B. C. (1990) JGR, 95, 1343-1358. [6]Herkenhoff K. E. and Murray B. C. (1990) JGR, 95,14,511-14,529. [7] Thomas P. and Weitz C. (1989)Icarus, 81, 185-215. [8] Herkenhoff K. E. and Va-savada A. R. (1999) in press. [9] Cutts J. A. (1973)JGR, 78, 4211-4221.
European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
Talus- and Landslide-Derived Mass-Wasting at Olympus Mons, Mars S. van Gasselt
1, E. Hauber , A. Dumke
1,
G. Neukum1.
1Institute of Geological Scienes, Planetary Sciences and Remote Sensing.
Figure 1: Quiraing basalt, Isle of Skye, Scotland.
1 m
European Space Agency European Mars Science and Exploration Conference: Mars Express & ExoMars ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007 TECTONIC AND PALEO-ENVIRONMENTAL CONSTRAINTS OF CLARITAS FOSSAE ON MARS. J. Raitala (1), P. Esestime (1, 2), J. Korteniemi (1), V.-P. Kostama (1), M. Aittola (1), T. Törmänen (1), T. Öhman (1) and G. Neukum (3), 1Astronomy Division, Department of Physical Sci., Univ. of Oulu, Finland, ([email protected]), 2Dipartimento di Scienze della Terra, Università “G. d´Annunzio” Chieti-Pescara, Italy, 3Institut of Geosciences, Dept. of Earth Sciences, Freie Universität, Berlin, Germany.
Introduction: Evolution of Claritas Fossae (CF) on the SSE slope of Tharsis was characterized by tectonics, volcanism and hydrology. Fluvial, erosion and sedimentary features of the CF area were formed within the active rift structure. This added details to their development together with the changes in global climate. The concept of tectonics that co-acted with climate-related events provides a framework to study the area. Morphology details let to identify the interplay between geologic processes and the paleolake [1,2] basin morphology and valley deformation due to climate and tectonics. The area is covered by the maps MC-17 and MC-25. The MEX-HRSC [3], THEMIS [4], MOC [5] and the very first HiRISE [6] images were used together with the MOLA topography [7].
Climate-related factors: Along with changes in Martian climate, water was mobilized from the poles during the high inclination of the rotation axis. Seasonally increased solar radiation evaporated polar caps and accumulated snow and ice on the mid-latitude hills [8-11]. The reverse climate phase due to decrease in inclination melted these ice reservoirs and moved water back to poles.
The hill slope alcoves or amphitheatres (Fig. 1) indicate ice accumulation areas. Glacial U-valleys lead down from them. Release of water from the volatile-rich hilltops eroded the lower slopes and resulted in channels originating from the deposits. An amount of water penetrated the ground and resulted in permafrost, and groundwater that led further to conduit formation along faults and to sapping events. This was repeated along the climate change cycle and resulted in frequent hydrology events that were correlated with tectonics.
Tectonics vs. hydrology: The CF tectonics has included several deformation phases. The E-W grabens belong to the oldest phase. They are still visible on the elongated NWW-SEE antiforms associated with the N-S Claritas Rupes (CR) fault on its western side. The wide set of conjugate N-S, NNE-SSW and NNW-SSE grabens were formed in several deformation events. The CR fault and the CF grabens form a rift zone on the main CF bulge. The multi-temporal tectonic events were accompanied by changes in climate and hydrology over a period of time as seen from the fact that channels were frequently re-arranged by tectonics. Some of the channels pre- and other post-date the faults of the very same set. Some basins provided
temporal volatile reservoirs. The southern CF paleolake [1,2] resembles that in the Morpheos basin [12]. An outflow carved a channel out of the lake to Icaria Planum while tectonic activity still continued as seen from the channels that do not follow the present topography. The few young faults on the basin floors can be used to identify some of the last hydrologic and tectonic re-surfacing types.
Fig. 1. The local hills and slopes display glacial
amphitheatres eroded by ice and water. Further consideration: The interwoven activity
phases of CF includes the rift development that had its driver in the Martian interior. Tectonics was complicated by volcanism [13] and hydrology. The faults provided aquifers for a substantial part of the water that originated from the high mid-latitude hills, and even water from the Tharsis volcanoes [9,11] may have utilized the CF rift. The broken uppermost surface allowed water to erode flow channels and channel networks. Groundwater has affected faults by erosion and fault lubrication. It followed faults carving conduits and cavities, and welling in places to the surface to form sapping structures. Repeated aquifer activation may also have provided humid shelters to support the increase and evolution of life forms - if they ever existed on Mars.
Acknowledgements: The HRSC Team, Academy of Finland and the Erasmus program supported the study.
References: [1] Raitala et al. (2004) Vernadsky-Brown Microsymposium 40, Abstr. #51. [2] Mangold and Ansan (2005) Icarus 180, 75-87. [3] Jaumann et al. (2007) PSS 55, 928-952. [4] Christensen et al. (2004) Space Sci. Rev. 110, 85-130. [5] Malin and Edgett (2001) JGR 106, 23429-23570. [6] McEwen et al. (2007) JGR in press. [7] Zuber et al. (1992) JGR 97, 7781-7797. [8] Laskar and Robutel (1993) Nature 361, 608-612. [9] Head et al. (2003) Nature 426, 797-802. [10] Raitala et al. (2005) LPSC XXXVI, Abstr. #1307. [11] Head et al. (2005) Nature 434, 346-351. [12] Kostama et al. (2007) JGR in press. [13] Dohm et al. (2001) USGS Map I-2650.
European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
THERMAL EVOLUTION OF MARS WITH PHASE TRANSITIONS O. Forni
1, D. Breuer
2.
1CESR-
CNRS, 9, av. du Colonel Roche, BP 44346, 31028 Toulouse Cedex 4, France. 2DLR, Rutherfordstraße 2
Breuer, D., and T. Spohn (2003), JGR., 108 (E7), 5072,
doi:10.1029/2002JE001999. [3] Roberts, J. H., and S. Zhong (2006), JGR, 111, E06013,
doi:10.1029/2005JE002668. [4] Spohn T. et al. (2001),
Space Science Reviews, 96, 231. [5] Nimmo F. and D. J.
Stevenson (2000), JGR, 195(E5), 11969. 33, [6] R. J.
Lillis et al. (2006), GRL, L03202, doi:10.1029/2005GL024905.
Figure 1. Evolution of the spherically averaged temperature in the lower mantle. The temperature moves from a convective
profile to a conductive profile around 900 Ma. The dashed line represents the position of the Spinel to Perovskite phase
transition.
THEMIS OBSERVES POSSIBLE CAVE SKYLIGHTS ON MARS. G. E. Cushing1,2, T. N. Titus1, J. J. Wynne1,2, P. R. Christensen3, 1U.S.G.S. 2255 N. Gemini Dr. Flagstaff, AZ 86001, [email protected], 2Northern Arizona University, Flagstaff, AZ 86011, 3Arizona State University, Tempe, AZ 85287.
Introduction: Here we report the discovery of seven candidate skylight entrances into subterranean caverns (Figure 1). All seven are located on the flanks of Arsia Mons (southernmost of the massive Tharsis-ridge shield volcanoes), a region with widespread col-lapse pits and grabens which may indicate an abun-dance of subsurface void spaces [1,2].
Motivation: Subterranean void spaces may be the only natural structures on Mars capable of pro-tecting life from a range of significant environmental hazards. With an atmospheric density less than 1% of the Earth’s and practically no magnetic field, the Mar-tian surface is essentially unprotected from micro-meteoroid bombardment, solar flares, UV radiation and high-energy particles from space [3,4,5,6]. Addi-tionally, intense dust storms occur planet wide, and some regions exhibit temperature ranges that can dou-ble over each diurnal cycle [7]. Besides general geo-logical interest, there is a strong motivation to find and explore Martian caves to determine what advantages these structures may provide future explorers. Fur-thermore, Martian caves are of great interest for their biological possibilities because they may have pro-vided habitat for past (or even current) life [5,6,8].
Preserved evidence of past or present life on Mars might only be found in caves [5,6,8], and such a discovery would be of unparalleled biological signifi-cance [3]. Cave deep zones on Earth generally main-tain constant climate conditions [9,10] which are ideal for the preservation of organic material. Accordingly, Martian caves are among the most desirable targets for astrobiological exploration [11,12,13,14].
Observations: The Mars Odyssey Thermal Emission Imaging System (THEMIS) collected the majority of data for this study [15]. From a nadir per-spective, THEMIS observes both visible and thermal-infrared wavelengths during the afternoon (~ 1500-1700 hrs), and IR wavelengths only for early-morning observations (~ 0300-0500 hrs.) [15].
The inspection of dark, circular pit-like fea-tures at visible wavelengths (VIS band 3, ~.654 μm) gave our first indication of potential skylight openings (nadir-pointing observations prevent us from determin-ing whether these are caverns or deep vertical shafts). To aid in visualization, we have informally named these ‘seven sisters’ on Arsia Mons as: Dena, Chloë, Wendy, Annie, Abbey, Nikki and Jeanne (Figure 1). Most of the candidates are adjacent to collapse pits or are directly in-line with collapse-pit chains, and appear to have formed by similar processes. They are visibly
distinct from collapse pits, however, by a lack of visi-ble (sunlit) walls or floors. These proposed skylights also lack the visible characteristics (such as raised rims or ejecta patterns) that would associate them with im-pact craters. Thermal behaviors furthermore confirm they are not misidentified surface features such as dark sand or rock.
Diameters generally range between 100-252 m (estimated from THEMIS VIS at 18 m/pixel for most images). Only minimum depths can be calculated (because the floors are not illuminated by the sun in THEMIS observations) and range between 73-96 m (diameter ÷ tan(incidence angle)). However, a fortu-nate MOC observation of Dena at ~2 p.m. (R0800159) actually does show an illuminated floor, allowing us to tightly constrain the depth using a 1-D photoclinome-try routine. This routine returns a depth of ~130 m for the illuminated floor, while the minimum depth esti-mated from the THEMIS observation is only ~80 m.
Because THEMIS IR observes at 100-m reso-lution, cavern skylights with diameters much smaller than that are probably not thermally distinguishable from regular temperature variations on the surface.
Discussion: Analyses of the candidates sug-gest they are not of impact origin, not patches of dark surface material, and are likely skylight openings into subsurface cavernous spaces. Visible observations show dark holes with sufficient depth that no illumi-nated floors (incidence angles ≥ 61.5°) can be seen from a nadir perspective (Thermal-infrared data sug-gest temperatures inside these features remain nearly constant throughout each diurnal cycle. Figure 2 shows afternoon temperatures for Annie that are warmer than the shadows of adjacent collapse pits, and cooler than sunlit portions. Meanwhile, nighttime temperatures for this candidate are warmer than all nearby surfaces. Such is the behavior we would expect of a cavern floor that receives little or no daily solar insolation [9,10].
Wendy, Dena, Annie and Jeanne are the strongest candidates because they have the most com-plete data sets; i.e., they have both VIS and diurnal IR coverage, and they are large enough to be clearly iden-tified at 100-m resolution. Chloë, Abbey and Nikki are also strong candidates because they have the same visible and thermal characteristics as the other candi-dates. Their minimum depths could not be constrained, however, because of late-afternoon observations when the sun is too low to shine deeply into the pits.
Conclusion: Additional observations are necessary—particularly those at different times of day
Lunar and Planetary Science XXXVIII (2007) 1371.pdf
and from an off-nadir perspective. These candidates cannot be physically explored with our current state of technology because the targets are too small and spe-cific, and the atmosphere at these elevations is too thin for a lander to slow down or maneuver sufficiently to reach them. The astrobiological significance may also be reduced at these elevations because microbial life, if it ever existed on Mars, may not have occurred at these elevations. However, possible evidence of liquid water at the Martian surface was recently identified by Ma-lin, et al. (2006) [16]. If liquid water does exist at or near the surface, then caves at lower elevations could hold natural reservoirs, greatly improving the possi-bilities for past or present microbial life.
The discovery of potential skylight openings into Martian caves is an exciting step towards future exploration and discovery. New spacecraft orbiting Mars, with greater observational capabilities, can ob-serve these candidates at higher resolutions, at differ-ent times of day, from different perspectives and in
different wavelengths. Future observations will pro-vide more substantial information about the character-istics and history of these features. A planet-wide search for similar targets is currently underway—particularly for those existing at lower elevations. This discovery presents us with new insights and new chal-lenges for the future of Mars exploration. References: [1] Ferrill, et al. (2003) LPSC XXXIV; [2] Wyrick, et al. (2004) JGR, 109(E6); [3] Mazur et al. (1978) Space Sci. Rev. v.22, 3-34; [4] Kuhn and Atreya (1979) J. Mol. Evol. v.14, 57-64; [5] Boston, et al. (2004) STAIF v.699 1007-1018; [6] Schulze-Makuch et al. (2005) JGR, 110(E12); [7] Cushing and Titus (2005) GRL, v. 32; [8] Fre-derick (2000) Concepts and App. for Mars Exp. 114; [9] Tuttle and Stevenson (1978) Nat. Cave Mgmt. Symp. Proc.; [10] Howarth (1980) Evolution v.34; [11] Grin et al. (1998) LPSC XXIX; [12] Boston (2000) Geotimes 45(8) 14-17; [13] Boston et al. (2001) LPSC XXXII; [14] Parnell et al. (2002) Astrobio. v.2(1), 43-57; [15] Christensen (2004) Space Sci. Rev. v.110(1); [16] Malin (2006) Science v.314 1574-1577.
Figure 1: Seven proposed cave skylights. Clockwise from upper-left: Dena, Chloë, Wendy, Annie, Abbey, Nikki and Jeanne. Arrows signify direction of solar illumination (I) and direction of North (N). Respective image IDs are: 18053001, 13448001, 17716001, 18340001, 14334002 and 18315002. To facilitate our photoclinometry routine, each candidate has been map-projected with the sun coming from the 9 o’clock direction.
Figure 2: THEMIS VIS and IR images show diurnal thermal behavior of a candidate cave skylight. [A] is the visible image, [B] is an afternoon IR image observed concurrently with the VIS (~1500 hrs), and panel [C] is an early-morning observation at 0400 hrs. This example represents the typical thermal behavior for all of our candidates.
Lunar and Planetary Science XXXVIII (2007) 1371.pdf
European Space Agency European Mars Science and Exploration Conference: Mars Express & ExoMars ESTEC, Noordwijk, The Netherlands, 12 ‐ 16 November, 2007
TITHONIUM CHASMA SALT BEARING OUTCROPS, STRATIGRAPHIC MARKER FOR
MARTIAN WATER SPAN C. Popa1, F. Esposito1, L. Colangeli1, G.G.Ori2; 1OAC (Osservatorio Astronomico di Capodimonte, salita Moiariello 16, Napoli 80131, Italy), 2IRSPS (International Research School of Planetary Sciences, Pescara, Italy)
Magnesium sulfate hydrated salts are present in the Internal Light toned outcrops of Tithonium Chasma (TC), the northern trough in western Valles Marineris (VM). Major part of formational paths for the formation of magnesium sulfates requires water presence in quantities large enough to pond in topographic depressions on Mars surface. Evaporation from brine derived from pristine rock alteration is a primary candidate for the formation of these outcrops. Morphological evidences prove a very likely situation of post depositional disturbance of the initial horizontal deposition for TC case. The TC outcrops have also a unique morphology amongst the VM magnesium hydrates salt bearing deposits, having an elongated attitude parallel to main tectonic lineation (the same of the trough) and an almost symmetric position in respect with Chasma walls, with a positive topography standing up to 2000 m above the chasm floor. A geologic analysis approach for this area is performed using available remote sensed data from the Mars Express ESA mission, in order to characterize the morphology and mineral distribution in the area. HRSC and OMEGA C channel data are used to establish the relationship between the topography and the mineral composition, (within the capabilities of the spectral range used). Seven OMEGA orbits (431, 887, 997, 1008, 1345, 1889, and 1911) were used for the spectral mapping of the area using the characteristic absorptions for the hydrated magnesium sulfates.
The study is focused on the establishment of the process(es) that could have emplaced the salt bearing outcrops, taking into account each possible candidate mechanism of formation from those synthesized in [1]. Lacustrine and dry depositions are good candidates for the outcrop emplacement, but can hardly explain the amount and the spatial confined emplacement of the outcrops emplacement in TC.
Crater counting dating of the outcrops would place them at the top of the stratigraphic chart for the geologic units, way recent compared to the proposed water span period [2], unless buried by geologic processes (igneous activity) posteriori their formation and subsequently exhumed by various processes in recent geologic periods.
Based on the appearance observation in visible MOC NA images [1,3] there is a debate for the stratigraphic sequence for IDL with respect to the wall rock along VM.
A very likely process for exhumation is salt diapirism. TC present all required tectonic and mineral conditions for diapirism in thin-skinned condition to occur. Also the surface morphology sustains diapirism process as primary exhuming process of a possibly early to medium Hesperian depositional process.
TC system on Mars can offer the means of assessing the Martian water time span, especially the superior limit of considered wet-dry boundary climatic transition [2]. We consider the most likely hypotheses of formation of LTO, and each specific hypothesis implication to the configuration of Tithonium Chasma, sorting the best fitting one according to the observed geomorphology and mineralogy aspects. We isolate and determine the spatial distribution of the LTO and LLO, as well as the water mineral bearing of each unit in order to assign a correct formational process that will better fit the existing mineral, temporal, and water span constrains. Water related mineral distribution partly overlaps the LTO, and bears various mineral hydration states that match the spectral signatures of magnesium sulfate (kieserite, epsomite). Their morphology rules out the posteriori formation of the outcrops in an eventuality of water filling the chasm after its opening whatever the formation mechanism may be, calling for alternative processes of emplacement. Diapirism hypothesis led to a model of stratigraphic time evolution of the area that fit the current general water span currently accepted.
References: [1] Catling et al. (2006) Icarus; [2] Bibring J.-P. et al. (2006) Science, 312, 400-404; [3]Malin and Edgett (2000).
European Space Agency European Mars Science and Exploration Conference: Mars Express & ExoMars ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007 Topographic Mapping and Rover Localization during the 2003 Mar Exploration Rover Mission Operations and New Developments for Future Landed Missions. R. Li, K. Di, B. Wu, W. Chen, J. Wang, S. He, J. Hwangbo, Y. Chen. The Ohio State University, Dept. of Civil & Env. Engineering & Geodetic Science, 470 Hitchcock Hall, 2070 Neil Avenue, Columbus, OH 43210, U.S.A. E-mail: [email protected].
MER mission operations In support of Mars Exploration Rover (MER)
mission operations, researchers at the Mapping & GIS Lab of the Ohio State University (OSU) have been collaborating with JPL and many other mission teams in performing rover localization and topographic mapping on a daily basis since the landing of the two rovers in January 2004 [1, 2]. From thousands of Pancam and Navcam ground images, we have produced a) rover localization products including accurate traverse maps, horizontal and vertical traverse profiles, plus the Spirit drive metric; b) regular topographic products including DTMs and ortho maps, 3D crater models, and 3D maps of large topographic features; and c) special topographic products such as north-facing slope maps and solar energy maps. These topographic maps and rover localization data have been extensively used in tactical and strategic planning and operations as well as various scientific investigations.
On-board rover localization is performed using wheel odometry, IMU (Inertial Measurement Unit), and a Sun positioning technique using Pancam imagery. A visual odometry technique is applied in order to correct errors caused by wheel slippage and ensure safe drives over difficult terrain as well as to provide high precision approaches to science targets within a relatively short distance [3]. In order to achieve high accuracy over long distances, incremental bundle adjustment (BA) of an image network formed by Pancam and/or Navcam images is carried out on Earth at the OSU Lab. After BA, 2D accuracy generally ranges from sub-pixel to 1.5 pixels while 3D accuracy is at a centimeter to sub-meter level (based on consistency check of the BA results). It has been demonstrated that BA-based rover localization technology has corrected wheel slippage, IMU drift and other navigation errors as large as 10.5% in the Husband Hill area of the Gusev Crater landing site (Spirit) and 21% in Eagle Crater at the Meridiani Planum landing site (Opportunity) [1, 2].
Autonomous rover localization Recently we developed an innovative method to
automate cross-site tie-point selection so that BA-based rover localization can be autonomously performed onboard the rover [4]. This new method consists of algorithms for rock extraction, rock modeling, and rock matching from multiple rover sites. Rocks are extracted from 3D ground points generated by dense matching of stereo images, and then modeled using analytical surface models such
as hemispheroid, semi-ellipsoid, cone and tetrahedron. Rocks extracted and modeled from two adjacent rover sites are matched by a combination of rock-model matching and rock-distribution-pattern matching. We have tested our software using a 337m traverse (20 pairs of sites) taken by Spirit at the Husband Hill summit area and a 206m traverse (13 pairs of sites) obtained at a Silver Lake test site on Earth. Test results show the proposed method is effective for medium-range (up to 26m) traverse segments; success rates for the number of site pairs are 65% and 76% (or 81% and 85% after prescreening) for the Spirit and Silver Lake data, respectively. We are further improving our methods and are performing tests using the entire 5-km traverse acquired at Silver Lake, CA, in January, 2007. At the same time, the onboard incremental BA technology we are developing will be integrated with JPL’s visual odometry technology to achieve long-range autonomous rover localization.
Enhanced topographic mapping With the support of the NASA Applied
Information System Research (AISR) Program, we are developing a method for the integration of orbital and ground images for enhanced topographic mapping. In this ongoing research, a combined bundle adjustment of orbital and ground imagery will be used to achieve the best possible accuracy for topographic mapping. We have developed a rigorous photogrammetric model and bundle-adjustment software for MOC NA and HiRISE stereo image processing and achieved sub-pixel accuracy at the MER sites. We have also developed a hierarchical stereo-matching process for DTM generation from stereo orbital images and for tie point selection.
Next, we will develop landmark (e.g., mountain peaks and crater rims) extraction methods for automatic linking between orbital and ground images. Consequently, the combined orbit-ground bundle adjustment will improve the precision of the image orientation parameters. The integration of orbital and ground images will enhance high-precision topographic mapping and rover localization in support of such planetary exploration tasks as pre-landing target selection, high-precision lander localization, and onboard navigation for the rovers.
References: [1] Li et al. (2006), JGR 111, DOI: 10.1029/2005JE002483. [2] Li et al. (2006), JGR 112, DOI: 10.1029/2006JE002776. [3] Maimone et al. (2007), JFR 24, 169-186. [4] Li et al. (2007), JFR 24, 187-203.
European Space Agency
European Mars Science and Exploration Conference: Mars Express & ExoMars
ESTEC, Noordwijk, The Netherlands, 12 - 16 November, 2007
TWO YEARS MARS GIS DEVELOPMENT
P. Saiger1,2
, F. Preusker1, A. Nass
1, M.Waehlisch
1, H. Asche
2, J. Oberst
1, R. Jaumann
1
1 Institute of Planetary Research, German Aerospace Center, Rutherfordstr. 2, D-12489 Berlin,