ESA Mars Research Abstracts Part 2

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.

.




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

Denmark, Juliane Maries Vej 30, 2100 Copenhagen, Denmark, 2Centre d’etudes des Environnements Terrestre

et Planétaires, (CETP), France, 3The Blackett Laboratory, Imperial College London, UK,

4CNRS/University of

Nantes, France, 5Earth and Planetary Physics, Niels Bohr Institute, University of Copenhagen, Denmark,

6GeoForschungsZentrum Potsdam, Germany,

7University of Western Brittany, France,

8University of

Edinburgh, UK. [email protected]

Mars Surface Magnetic Observatory (MSMO),

an experiment planned as part of the Geophysical

and Environment Package (GEP) on ExoMars, is

likely to provide the first magnetic field

measurements ever performed at the surface of

Mars. It will provide unique information in a wide

spectrum of scientific applications in accordance

with the ExoMars scientific objectives.

The overall scientific purposes of the

magnetometer are:

• to study the effect on the Martian

environment of the solar wind interaction with the

atmosphere, including atmospheric escape,

• to understand the effect on the Martian

environment of explosive events on the Sun (i.e.

CME’s, flares), including the variability in ionizing

radiation due to solar energetic particles,

• to determine the electrical conductivity of

the planetary interior as a function of depth, in order

to map deep subsurface water reservoirs and

understand the planetary evolution,

• to improve the resolution of the crustal

magnetic field and estimates of its sources.

Figure 1: Average radial magnetic perturbations close

to crustal anomalies at the Martian dayside, as measured

by MGS in 400 km’s altitude. The color scale is in nT.

The contours show the location of the crustal anomalies.

The solar wind interacts with the Martian

atmosphere creating a so-called induced

magnetosphere of draped magnetic field. In this

process currents are created in the day-side

ionosphere that acts to shield the ionosphere and

surface from the magnetospheric field. With the

MSMO we will investigate the efficiency of the

shielding and the morphology of ionospheric

currents. While these processes have been observed

from orbiting spacecraft, the MSMO will provide

the first continuous measurements from a low

altitude vantage point. A description of these

processes is important in order to understand plasma

escape processes at Mars, in particular their

variability with the solar wind and solar activity.

If a landing site close to one of the crustal

magnetic anomalies is selected we will also be able

to study the currents created in the direct interaction

between the solar wind and the crustal field (Figure

1).

The MSMO would have a strong synergy with

simultaneous magnetic and plasma measurements

from orbit.

The magnetometer proposed for the MSMO

experiment derives from instruments flown in

dedicated geomagnetic missions (Ørsted, CHAMP,

SAC-C). The Danish space magnetometers are all

dedicated instruments to accurately map the Earth’s

magnetic field and have more than 26 years total

combined in-orbit operation time. The current

magnetometer is a miniaturised version of the

earlier instruments and is baselined for the ESA

PROBA-2 and Swarm missions. One of the key

parameters of the magnetometer is the zero-level

accuracy of the measured vector field. The fluxgate

transducers use a stable and low-noise (12 pTRMS

in 0.01-10 Hz), stress-annealed amorphous magnetic

material for the ring-cores. The heat treatment of the

ring-cores has been developed in Denmark, starting

with the Ørsted magnetometer. The resolution of the

magnetometer is based on a 22 bit A/D converter.

The zero-level stability of the magnetometer has

been verified in-flight to a level of 0.3 nTRMS over

a period of more than 6 years and in temperature

variations of +/-10°C. By calibration a DC-accuracy

of 0.3 nT (1sigma) can be obtained over a larger

temperature span of about 100°C. In order to

determine the orientation of the measured magnetic

field vector the magnetometer is combined with an

attitude sensor consisting of two gravity sensors.




MARSIS DATA INVERSION APPROACH: Preliminary results

G. Picardi1*, D. Biccari

1, M. Cartacci

1, A.Cicchetti

1, S. Giuppi

1, A. Marini

1 , A..Masdea

1, R.Noschese

1, F.

Piccari1, R. Seu

1, J.J.Plaut

2, W.T.K.Johnson

2, R.L.Jordan

2, A.Safaeinili

2, C.Federico

3, A.Frigeri

3, P.T.Melacci

4,

R. Orosei5, O.Bombaci

6, D.Calabrese

6, E.Zampolini

6, P.Edenhofer

7,D.Plettemeier

8, L. Marinangeli

9,

E.Pettinelli10

, T. Hagfors11

, E. Flamini12

, G.Vannaroni13

, E. Nielsen14

, I.Williams15

, D. A. Gurnett16

, D. L.

Kirchner16

, R. L. Huff16

1Infocom Dept.- University of Rome “La Sapienza”, Via Eudossiana, 18 – 00184 Rome-Italy,

2Jet Propulsion

Laboratory – 4800 Oak Grove Drive - Pasadena, CA-91109 - USA - 3

Dept. of Earth Science - University of

Perugia, 06123 Perugia Italy, 4Computer Science Dept. - University of Perugia- Via Vanvitelli 1, 06123 Perugia

Italy, 5

INAF-IASF. - Via del Fosso di Cavaliere,100 - 00133 Rome – Italy, 6Alcatel Alenia Space - Via

Saccomuro,24 - 00131 Rome –Italy, 7Institut für HochfrequenztechnikArbeitsgruppe Antennen und

Wellenausbreitung Fakultät für Elektrotechnik und Informationstechnik Ruhr-Universität Bochum 44780

Bochum, Germany, 8Fakultaet für Elektrotechnik und Informationstechnik Lehrstuhl und Laboratorium für

Theoretische 9International Research School of Planetary Sciences, Dipartimento di Scienze, Universita'

d'Annunzio, Viale Pindaro 42 - 65127 Pescara – Italy, 10

Physics Dept.–University of Rome“Roma Tre”, Via

della Vasca Navale, 84– 00146 Rome-Italy, 11

Max Plank Institut fur Aeronomie, Germanie, 12

ASI, Viale Liegi,

26 – 00198 Rome, Italy, 13

INAF-IFSI. - Via del Fosso di Cavaliere,100 - 00133 Rome – Italy, 14

Max Planck

Institute for Solar System Research 37191 Katlenburg-Lindau, Federal Republic of Germany 15

Astronomy unit

Queen Mary-University of London-U.K., 16

Dept. of Physics and Astronomy – The University of Iowa – Iowa

City, IA – 52242 – USA

[email protected]

Abstract

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




“ODD” MARTIAN DICHOTOMY AND ITS HARMONIC INTERPRETATION. G.G. Kochemasov.

IGEM RAS, 35 Staromonetny, Moscow 119017, Russia, [email protected]

The long-known looking strange north-south

tectonic and chemical dichotomy of Mars is rather

expected if considered by the wave planetology [1,

2 & others]. This science is based on a

fundamental assertion: “orbits make structures”.

Keplerian orbits stimulate in heavenly bodies

inertia-gravity waves of which the fundamental

wave long 2 R having a stationary character

inevitably produces tectonic dichotomy. This is

followed by a chemical (density) one for leveling

angular momenta of opposed risen and fallen

segments. A sharpness of the martian dichotomy is

due partly to a highly elliptical orbit of the planet.

The structurizing fundamental wave has

overtones of which the first one long R (in

resonance 1 : 1 with the wave granulation also R

size due to the martian orbital frequency [2, 3 &

others]) produces very spectacular effect: Mars

obtains an “surprising” elliptical shape in the

equatorial section. Huge elevations prolong it and

very deep depressions squeeze it in the

perpendicular sense. There is no sense to

continue call Hellas and Argyre basins impact

structures as these deepest hollows on the elevated

southern hemisphere have symmetrical but risen

counterparts on the subsided northern hemisphere

[2]. Hellas is in pair with the far advanced onto the

northern lowlands Alba patera & Tempe terra;

Argyre is in pair with Elysium planum & Phlegra

montes. A such regular (harmonic) construction has

nothing to do with occasional random impacts and

witnesses a major role of regular wave processes in

structurizing planetary surfaces and depths.

The Mars Express’ MARSIS experiment [4]

seems to indicate at the third tectonic pair with the

same regular disposition. Radar measurements of

ice thickness on both pole regions (2 km at the north

and 3.7 km at the south) could indicate that at the

north ice covers slightly elevated background and at

the south ice fills a kind of depression – bowl. This

disposition reminds (with inversion) the terrestrial

case: Antarctica continent on the mainly oceanic

southern hemisphere and a depression of the

Northern Polar ocean on the mainly continental

northern hemisphere. This comparison only

strengthens the case of the planetary wave

structurizations.

A very sharp martian vertical block

differentiation requires significant density difference

of composing them rocks (sharper than at Earth:

tholeiites – andesites) for leveling blocks’ angular

momenta. In 1995 before the Pathfinder landing we

knew about dense Fe-basalts covering deeply

subsided northern lowlands. Thus, very light rocks

as syenites and granites were proposed for sharply

elevated southern highlands [3]. Gradually their

signs were discovered even under obstacles of

plateau-basalts, sills, eolian basalt sands and so on.

Up to now the best ground truth that requires orbital

crafts (and Mars Express) is presented by MERs.

Spirit found very light weathered rocks (like

powder) under dark eolian sands. On a very small

surface there is very sharp compositional difference

of this products of alteration: Fe-S rich, Ca-rich, Si-

rich. This rejects large volumes of open water here

in the past; otherwise a composition difference

would be minimized. There is also rather high

compositional differentiation of floats and outcrops.

Such a sharp chemical differentiation on a very

small expanse and a fine layering is typical for

alkaline massifs. For alkaline rocks, not for other

lithologies, are very typical also poicilitic textures

like observed at Figure (outcrop “Slide”, brushed

circle, 3 cm diameter). Light-colored rocks, smectite

and sulfates indications are widespread on the

southern highlands according to OMEGA [5 &

others]. It seems that a massive involvement of

hydrous sulfites in highland rock compositions plays

on diminishing their density what is required by the

martian very sharp vertical tectonic differentiation.

References: [1] Kochemasov G.G. (1992) 16

th Russian-

American microsymposium on planetology, Abstracts,

Moscow, Vernadsky Inst., 36-37. [2] Kochemasov G.G.

(2004) In Workshop on “Hemispheres apart: the origin

and modification of the martian crustal dichotomy”, LPI

Contribution # 1203, Lunar and Planetary Institute,

Houston, p. 37. [3] Kochemasov G.G. (1995) Golombek

M.P., Edgett K.S., Rice J.W.Jr. (eds) Mars Pathfinder

landing site workshop II: Characteristics of the Ares

Vallis region and Field trips to the Channeled Scabland,

Washington. LPI Tech. Rpt. 95-01, Pt. 1, LPI, Houston,

63 pp. (p. 18-19). [4] Chicarro A.F. (2007) Seventh

International Conference on Mars, Pasadena, Calif., abs. #

3009. [5] Rossi A.P., Neukum G., Poudrelli M. et al.

(2007)LPSC XXXVIII, abs. # 1549.




THE MARTIAN LITHOSPHERE IN THE THARSIS REGION: A COMPARISON BETWEEN MEX

GRAVITY DATA AND THE MOLA TOPOGRAPHY MODEL M. Fels1, M. Pätzold

1, B. Häusler

2.

1Rheinisches Institut für Umweltforschung an der Universität zu Köln, Cologne, Germany.

2Universität der

Bundeswehr München,Neubiberg, Germany. [email protected]

The first European Mars Mission, Mars Express

(MEX), is operating in orbit around Mars since

Januar 2004. The Mars Express Radio-Science

Experiment (MaRS) is performing gravity

measurements above selected target areas during

the pericenter passes at an altitude from 250 km to

350 km. MEX has a much higher sensitivity to

gravity attractions at small scales than the NASA

mission Mars Global Surveyor (MGS) due to this

low pericenter altitude.

A total of 70 Doppler observations above selected

target areas could be recorded at the ESA ground

station in New Norcia and at the antennas of the

Deep Space Network (DSN). Profiles of the

gravitational acceleration could be computed after

low-pass filtering. These residual accelerations will

be compared with the MOLA topography model

from MGS by computing the cross-correlations

between these both datasets for all gravity

operations belonging to the same target area to

make a statement about the compensation status of

the particular local and regional Martian

lithosphere respectively.




The Martian Organic Material Irradiation and Evolution experiment: study of the behavior of organic

molecules at the Mars surface P. Coll1, F. Stalport

1, C. Szopa

2.

1Laboratoire Interuniversitaire des Systèmes

Atmosphériques (LISA), University Paris XII, Créteil France; 2Service d’Aéronomie (SA), University Paris VI, Paris

France. Contact : [email protected]

Among the biomarkers to seek on Mars,

organic molecules are primordial because some are

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.




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

Seventh International Conference on Mars 3354.pdf

mailto:[email protected]


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.


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.


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.





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,

boundary layer phenomena, and climatological

cycles requires simultaneous in-situ meteorological

measurements from networks of stations at the

Martian surface. The scientific payload of the

MetNet Mission encompasses separate instrument

packages for the atmospheric entry and descent

phase and for the surface operation phase. For the

descent phase an imager, accelerometers and

devices for free flow pressure and temperature

observations are envisaged. At the Martian surface

the MML will take panoramic pictures, and perform

measurements of pressure, temperature, humidity,

wind direction and speed, as well as atmospheric

optical depth. The MetNet prototype has been

developed and the critical subsystems have been

qualified for Martian environmental and functional

conditions. Presently a suborbital test launch is

under preparation to test the descent systems of the

MetNet. The first mission step in the MetNet

Mission is to have a MetNet Precursor Mission with

a few MMLs deployed to Mars. The MetNet-type of

mission is what the Martian atmospheric science

currently needs. Detailed characterization of the

Martian atmospheric circulation patterns and

climatological cycles requires simultaneous in situ

atmospheric observations by a network of stations at

the Martian surface. The MetNet mission will

provide the logical next mission tool in the field of

Martian atmospheric science.




MEX/ASPERA-3 NPI DATA STATISTICAL ANALYSIS, A. Milillo

1, A. Mura

1, S. Orsini

1,

1INAF/IFSI,

via del Fosso del Cavaliere 100, 00185, Rome, Italy. [email protected]

Section 1: The Analyser of Space Plasma and

Energetic Atoms (ASPERA-3) Neutral Particle

Imager (NPI) on board Mars Express (MEX) is

devoted to energetic neutral atom (ENA) detection

within the Martian environment. These ENAs

originate from the interaction between the energetic

ions flowing inside the Martian environment and the

exospheric neutral gas, thus providing crucial

information about the dynamics of this interaction.

NPI records the instantaneous angular distribution

of the energy-integrated ENA signal. In order to

identify recurrent ENA signals in the Martian

environment, we have performed a statistical

analysis of the NPI data. Count rates have been

averaged in different ways in order to be able to

discriminate signals coming from the planet, from a

selected direction, or from specific planetographic

regions at the planetary surface. Possible recurrent

ENA signals are from the terminator and the above

atmosphere toward night side, mainly when the

spacecraft is close to the edge of the shadow, while

there is no signal relation to magnetic anomalies.

This study shows that the statistical analysis of the

NPI data does not produce significant results since

the recurrent signals are below the sensor intrinsic

error. In fact, the sensor has some intrinsic

limitations due to non-adequate UV suppression,

difficulties in sector inter-calibrations, and

variations in the sector response versus time.




MICROFLUIDIC CAPILLARY ELECTROPHORESIS LAB-ON-A-CHIP SYSTEM

MICROFABRICATION AND INTEGRATION FOR THE UREY INSTRUMENT P. A. Willis1, B.D.

Hunt2, J. A. Smith

1, V. E. White

1, M. C. Lee

1, H. F. Greer

1, T. Chiesl

3, R. A. Mathies

3, and F. J. Grunthaner

1.

1Jet Propulsion Laboratory, California Institute of Technology, 4800 Oak Grove Dr., Pasadena, CA, USA.

2Atomate Corporation, Simi Valley, CA, USA.

3University of California, Berkeley, CA, USA.

Contact: [email protected]

We have developed an advanced microfluidic

capillary electrophoresis (CE) system integrated

with on-chip Teflon-membrane valves and pumps

[1], as part of the Urey Instrument, scheduled for

inclusion in the Pasteur Payload of the ExoMars

Rover. This work builds on the CE system

developed by Skelley et al. [2], but extends the

capability through the use of bio- and spaceflight-

compatible Teflon-membrane valves rather than a

PDMS-based approach. The CE system is currently

being used for sensitive compositional and chiral

analysis of amino acids with the goal of identifying

past or present life signatures in extraterrestrial

environments. The wafer design utilizes

independent CE channels patterned in glass, along

with a Teflon membrane, a pneumatic manifold

layer, and a fluidic bus layer, as shown in the wafer

cross-sectional view of Figure 1. The valves

provide isolation of the sample and buffer ports, as

well as peristaltic-type pumping in a three-valve

configuration. Electrophoretic separation occurs in

the all-glass channels near the bottom of the

structure. The pumps and fluidic bus channels

deliver and remove buffer, sample, and waste from

the four CE ports, under control of a Labview-

driven pneumatic switching network. The device

configuration is similar to Mars Organic Analyzer

developed at U.C. Berkeley [2], but includes

significant design and process improvements to

enable efficient Teflon valve operation and effective

bonding of the membrane. The completed wafer is

mounted on a fluorescent microscope stage in a

custom fixture, which interfaces the pneumatic and

high voltage lines and has the capability for

controlled atmosphere testing (Figure 2). Typical

electrophoretic separation data for a fluorescamine-

tagged amino acid run is plotted in Figure 3.

References:

[1] Willis, P.A., Hunt, B.D., White, V.E., Lee, M.C.,

Ikeda, M., Bae, S., Pelletier, M.J., and Grunthaner, F.J.

(2007), Lab Chip, DOI:10.1039/b707892g.

[2] Skelley, A.M., Scherer, J.R., Aubrey, A.D., Grover,

W.H., Ivester, R.H.C., Ehrenfreund, P., Grunthaner, F.J.,

Bada, J.L., and Mathies, R.A. PNAS 102 (4), 1041-1046.

Figure 3. (above) Fluorescence intensity vs. time for

electrophoresis of a solution of 50 M L-valine, L-serine,

and L-glutamic acid at two different detection points in the

separation channel.

Figure 2. (left) Photograph of JPL integrated CE system

showing the CE wafer with pneumatic and high voltage inputs above fluorescent excitation/collection optics.

Figure 1. (above) Cross-sectional diagram of CE wafer

stack including integrated Teflon valves and pumps.




MICROMEGA: DESIGN AND STATUS OF A NEAR-INFRARED SPECTRAL MICROSCOPE FOR

IN SITU ANALYSIS OF MARS SAMPLES. V. Leroi1, J.P. Bibring

1, M. Berthe

1.

1Institut d’Astrophysique

Spatiale (IAS), Université Paris-Sud bât. 121, 91405 Orsay, France. Contact: [email protected]

From OMEGA to MicrOmega: The

visible/near-infrared imaging spectrometer

MEx/OMEGA has provided a great change in our

vision of Mars [1,2]. In this context, we are

developing a spectral microscopic imager,

MicrOmega [3] for the ExoMars rover mission. This

instrument will acquire in situ reflectance spectra of

Martian samples, at a scale of the grain size (spatial

sampling of 20 m per pixel), in a non destructive

way. It will work in the spectral range 0.9 to 2.6 m.

MicrOmega will illuminate 5 mm-sized sample

sequently in 1000 contiguous wavelengh channels,

and will take an image on a matrix detector for each

channel. In this way, we get an ‘Image Cube’ in

which the full spectrum of the viewed area is

acquired in each pixel. This will enable us to

retrieve the composition of the different phases

since each mineral exhibits a unique signature in the

near-infrared through specific absorption bands.

This composition is essential to get new clues about

the formation and the evolution of Mars.

The instrument: MicrOmega inherits the

structure and the development of the MEx/OMEGA

[4] and the Rosetta/CIVA [5] instruments. The main

development focuses on the change of the usual

grating technology by an Acousto Optics Tunable

Filter (AOTF) [6]. This optic device is composed of

a cristal in which the light is diffracted by an

acoustic field (under specific conditions): the

wavenumber of the light diffracted is directly linked

to the acoustic frequency. This system does not

weigh more than current technologies (60g) and

presents important improvements: suppression of

mechanism, no second order, increased reliability,

good resolution. This new technology has already

been used in MEx/SPICAM [7] and VEx/SOIR [8]

but the conditions are largely different: lower

temperatures with day/night cycles, atmospheric

environment. Additional developments were

required to qualify this technology in those new

conditions.

Tests already performed: In the design and

conception of MicrOmega, we have a full Labview

piloted test bench. An infrared source, which is an

industrial tungsten filament lamp, is now selected

for our illumination system. An absolute sprectro-

photometry of this lamp has been measured.

We have an operational AOTF (figure 1) in the

spectral range 0.8 m up to 4 m. This AOTF has

been qualified at IAS to run at low temperature

(tested at 140K). This AOTF has also been

calibrated on our Labview test bench and we have

tested all its characteristics (no second order,

diffracted light wavenumber relative to the acoustic

frequency). We have optimized the output beam in

order to remove the non diffracted light that is

considered as noise.

Work in progress: We are now working on an

near-infrared focal plane array detector. This

detector is a 356 per 256 pixels and will be included

in our optical scheme to perform imaging of Martian

analog samples. A demonstrator model should be

ready within the end of the year. We will work on

different ways to illuminate the sample thanks to a

close collaboration with the ExoMars rover design

team.

References:. [1] Bibring, J.-P. et al. (2006), Science 312,

400-404. [2] Poulet, F. et al. (2005), Nature 438, 623-627.

[3] Berthe, M. and J.P. Bibring (2006), AGU Fall Meeting

2006, Abs. #P53A-05. [4] Bibing, J.P. et al., ESA pub., sp-

1240. [5] Bibring, J.P. et al. (2007), Space Science

Reviews Vol. 128, 397-412. [6] Goutzoulis, A.P. and D.R.

Pape (1994), Design and fabrication of acousto-optic

devices. [7] Korablev, O. et al. (2002), Adv. Space Res.

Vol. 29, 143-150. [8] Berthaux, J.L. et al. (2006), 36th

COSPAR Scientific Assembly.

Figure 1. The Acousto Optic Tunable Filter (AOTF) with its illumination system and its output optics.




THE MICROMEGA/EXOMARS INVESTIGATION J-P. Bibring

1, F. Westall

2, N. Thomas

3 and the

MicrOmega team, 1IAS, Orsay Campus, France,

2CBM, Orléans, France,

3UoB, Bern, Switzerland

[email protected]

MicrOmega is a set of ultra-miniaturized

microscopes, designed to characterize the samples

acquired and distributed by the ExoMars Sample

Preparation and Distribution System (SPDS). It is

constituted of two distinct units, an optical color

microscope (MicrOmega-V) and a near-infrared

hyperspectral microscope (MicrOmega-I), operated

by a single electronic unit (MicrOmega-E).

MicrOmega-V has the capability of identifying the

texture, the structure and the morphology of each

provided sample, down to a spatial sampling of 3

m: the samples will be sequentially illuminated by

LEDs in a wide spectral range (UV to NIR),

possibly with polarizing filters, and imaged on a

CCD matrix.

MicrOmega-I will characterize the molecular and

mineralogical composition at a scale of ~20 m, by

the identification of diagnostic absorption features

in the near-infrared (0.8 – 2.6 m): for each pixel,

the spectrum will be acquired in up to 1000 spectral

contiguous channels. Specifically, MicrOmega-I has

the capability of identifying and mapping hydrated

phases, such as phyllosilicates, sulfates and

carbonates, if ever present, which constitute unique

tracers of potential habitability.

The MicrOmega fully non destructive sample

characterization will serve as the first step in an

integrated analytical protocol, in which the samples

will be further analyzed by the various instruments

of the ExoMars analytical laboratory.

We will present the development status of our

instrument, and the scientific goals of our

investigation, in the framework of the global

ExoMars mission objectives.




MICROWAVE SOUNDING OF THE MARTIAN ATMOSPHERE FROM AN EXOMARS ORBITER

F. Forget1, M. Capderou

1, G. Beaudin

2, A. Deschamps

2, M. Gheudin

2, A. Maestrini

2, J.M. Krieg

2, E. Lellouch

3

T. Encrenaz3, T. Fouchet

3, P. Ricaud

4, J. Urban

4, P. Hartogh

5, E. Chassefiere

6, F. Lefevre

6, F. Montmessin

6, A.

Emrich7, D. Murtagh

8, M. Janssen

9, R.T. Clancy

10,

1Laboratoire de Météorologie Dynamique, IPSL, Université Pierre & Marie Curie, BP99, 4 place Jussieu, 75252

Paris, Cedex 05, France, 2

LERMA, Observatoire de Paris, France, 3 LESIA, Observatoire de Paris, France,

4

Laboratoire d’Aerologie, Toulouse, France, 5 MPIS, Lindau, Germany,

6 SA, Paris, France,

7 Omnisys, Sweden,

8 Chalmers Univ. Sweden, 9 JPL, Pasadena, USA,

10 SSI, USA, [email protected]

An Exomars Orbiter with a suitable orbit could be a

excellent opportunity to revolutionize our

understanding of the the Martian climate and

meteorology, and of the composition of the Martian

atmosphere. For this purpose, we propose to analyze

the thermal emission of the atmosphere at

microwave frequencies using heterodyne

spectroscopy, for the first time from orbit around

another planet. In practice, the Mars Atmosphere

Microwave Brightness Orbiter MAMBO would

perform measurements at the atmospheric limb and

at nadir using a receiver dedicated to the monitoring

of selected lines of key molecule around 320-350

GHz: H2O, CO, 13

CO, HDO, O3, H2O2.

In such conditions, the instrument performance

would allow the 3D mapping, with an excellent

spatial coverage, of the following characteristics :

• Winds : The high spectral resolution allows to

make use of the line profiles and their Doppler

shift. Limb viewing thus allows the fist direct

measurements of the winds on Mars from orbit

from 20 to 110 km with a vertical resolution

better than 10 km and an accuracy of about 15

m.s-1

Such a measurement, never done before,

would provide key information on the

atmospheric dynamic..

• Temperature : The temperature profile would be

retrieved with high vertical resolution (5 km)

without regard to dust opacity and season. A

unique characteristic of Microwave sounding is

the ability to profile temperature up to 120 km,

compared to 70 km for previous sounders.

• Water Vapour : near the surface up to 60 km,

with a sensitivity and vertical resolution (5 km)

much better than previous experiments, without

regard to dust opacity and season.

• D/H Ratio : This isotopic ratio will be mapped

accurately from 0 to 40 km by simultaneous

spectroscopy of H2O and HDO. Mapping the

variation D/H ratio is a key investigation to

understand the evolution of water on Mars,

escape processes, and Mars cloud microphysics..

• Hydrogen Peroxyde (H202) : This species has

only recently been detected on Mars, but it is

thought to be of key importance for the

photochemistry of the martian atmosphere

(control of H2, O2 and CO) and for its role

(thought to be major) in oxydizing the martian

soil, a problem of key interest for exobiology

• Ozone : Ozone profile will be measured

accurately up to 70 km, simultaneously with

water vapor. This will allow us to better

understand the relationship between the two

species

• Carbon Monoxyde : the variations of this

species will be monitored up to 120 km,

providing important clues on the meridional

transport in the Martian atmosphere.

Overall, the combination of these measurements

provides us with a complete view of Mars

Atmospheric dynamics, Water cycle, and

atmospheric photochemistry. I

In practice, MAMBO would be a 20 kg class

instrument with a peak power of 30 to 50 W. A

detailed concept (phase B) was designed for a

previous project, but it could be revised in order to

benefit from the significant progress made in the

very active microwave technological fields.

Figure 1 : Simulation of limb spectra at 10 km

around 320-350 GHz under typical Martian

conditions.




MIMA, A MINIATURIZED FOURIER INFRARED SPECTROMETER FOR MARS GROUND AND

ATMOSPHERIC STUDIES: PART II, SCIENTIFIC GOALS. L. Zasova1, G. Bellucci

2, B. Saggin

3, S.

Fonti4, A. Grigoriev

1, N. Ignatiev

1, B. Moshkin

1, F. Altieri

2, E. Alberti

3, G. Maurzo

4.

1Space Research Institute,

IKI, 117997 Moscow, Russia. 2INAF, Istituto di Fisica Spazio Interplanetario, 00133 Rome, Italy.

3Politecnico di

Milano, Department of Mechanical Engineering, 23900 Lecco, Italy. 4Università degli Studi di Lecce,

Department of Physics, 73100 Lecce, Italy.. [email protected]

MIMA with its spectral range 400 – 5000

cm-1 and spectral resolution of 5 cm-1 is powerful

instrument for both the surface and atmospheric

studies. Parallel with the mineralogical investigation

and assistance the rover exobiological explorations

MIMA enables to study the boundary layer, aerosols

distribution and optical properties, minor

compound, including biologically important ones.

Comparing to mini TES carried by Spirit and

Opportunity MIMA has higher spectral resolution (5

cm-1, against 10 cm-1 for TES) and in addition,

MIMA spectral range includes the solar reflected

part of the IR spectrum up to 2.0 μm.

1

2

3

ic e d u s t

1 – _ ice = 0.5 (at 825 cm -1), no dust, _ = 10°

2 – no ice, _dust = 1 (at 1075 cm -1)

3 - _ ice = 0.5 (at 825 cm -1), no dust, _ = 75°

H2O

CO2

Fig 1. Synthetic spectrum in TIR spectral range of

MIMA

Martian synthetic spectra of MIMA are shown in

Fig.1a,b for its TIR and NIR spectral ranges.

Temperature profiles in boundary layer are

retrieved with high accuracy and high vertical

resolution (up to 20-30 m) in the lowest 1-2 km

layer from inversion of the intensity in the 15 m

CO2 band.

Aerosol study. A particulate component of the Mars

atmosphere is composed by micron-sized particles,

which are products of soil weathering, and water ice

clouds. They strongly affect the current climate of

the planet. Pointing possibility together with wide

spectral range of MIMA will allow not only

monitoring of opacity and its variation but also

obtain aerosol optical properties, using EPF

observation mode.

Water vapor. Observations from the rover will give

a possibility of measurements of abundance and

water vapor, its diurnal and seasonal variation.

Vertical profile may be obtained from observations

of the near solar sky from zenith to horizon.

Two water vapor bands may be used for water vapor

estimation: 6.3μm (1400-1800 cm-1) and 2.6 μm

(3900 cm-1).

H2Oice dust

1 – _ ice = 0.5 (at 825 cm -1), no dust, _ = 10 °,

2 – no ice, _ dust = 1 (at 1075 cm -1)

CO2

1

2

Fig 2. Synthetic spectrum in NIR spectral range of

MIMA

Carbon monoxide, its abundance, vertical profile

and variation will allow obtained measurements

with MIMA. Biologically important components as

methane and formaldehyde will be measured in

the case of existence of local enhancements of this

species.

Russian channel of MIMA experiment can provide

not only the short-wavelength calibration, but also

Attenuated Total Reflection ("ATR") spectra of

airborne Martian dust and frosts precipitating from

atmosphere. A possibility of cooperation of ExoMars and

Russian mission Phobos-Grunt is under discussion

now. It will allow to produce joint observations by

MIMA from the surface and by mini Fourier

spectrometer AOST on board of Phobos-Grunt from

orbit.




MINERALOGY, MORPHOLOGY, SEDIMENTARY FILLING AND HISTORY OF ARAM CHAOS

ON MARS. M. Massé1, S. Le Mouélic

1, O. Bourgeois

1, C. Sotin

1, J.-P. Bibring

2, B. Gondet

2, Y. Langevin

2.

1Laboratoire de Planétologie et de Géodynamique UMR-CNRS 6112, Université de Nantes, Faculté des

Sciences et Techniques, 2 chemin de la Houssinière BP 92208, 44322 Nantes Cedex 3, France, 2Institut

d’Astrophysique Spatiale, Université Paris 11, 91405 Orsay, France. [email protected]

Introduction: Aram Chaos is a crater 280 km in

diameter located northeast of Valles Marineris. This

depression is connected to the Ares Vallis outflow

channel by a 15 km wide, 2.5 km deep channel cut

across the crater wall, which suggests that

significant amounts of water were present in the

past. Previous global-scale studies of TES and

OMEGA data revealed that the crater is filled by a

dome-shaped, 900 m thick, sedimentary formation

with strong signatures of ferric oxides [1, 2]. The

aim of our study is to describe the nature and

structure of this sedimentary formation using higher

resolution data and to deduce a plausible history for

Aram Chaos.

Methodology: We investigated the detailed

mineralogy of Aram Chaos by using the OMEGA

instrument onboard Mars Express. This imaging

spectrometer has completed a near global coverage

of Mars in 352 spectral channels from 0.3 to 5.1 m

at a spatial resolution ranging from 300 m to 4 km.

After removing of atmospheric contribution, we

computed maps of spectral parameters and maps of

the main mineralogical families derived from a

linear unmixing deconvolution algorithm [3]. These

OMEGA processing products have been integrated

into a GIS.

Morphological, textural and sedimentological

information is provided by available high resolution

images : MOLA for topographic information,

HRSC, MOC, and HiRISE for visible images and

THEMIS for visible and infrared images. These

images are integrated into the same GIS in order to

investigate correlations between mineralogical and

morphological characteristics.

Results: Four spectral units (SU) are identified.

SU1 contains a strong ferric oxide signature, with a

deep absorption band at 0.9 m, and a significant

increase of the reflectance between 0.9 and 1.3 m.

Absorption bands at 1.4 and 1.9 m also indicate the

presence of a hydrated mineral. This SU is found on

different areas corresponding on high resolution

images to large sheets of dark dunes covering

outcrops, too small to be resolved with OMEGA, of

a bright layered sedimentary formation.

SU2 presents the same characteristics as SU1 but

with an additional broad band at 2.1 m, which is

typical of sulfates (kieserite or szomolnokite being

good candidates). SU2 is correlated with clean, wide

and fresh outcrops of the bright sedimentary

formation.

SU3 displays shallower absorption band depths

and presents a negative spectral slope characteristic

of dust. On high-resolution images, SU3

corresponds to the eroded and dust-covered surface

of the bright sedimentary formation.

The spectral characteristics of SU4 are typical of

dusty areas. SU4 appears as chaotic terrains which

are stratigraphically below, and which crop out

around the bright sedimentary formation.

Erosion cliffs, cut across the bright sedimentary

formation, are covered by dark debris fans, which

originate from the bright formation itself. Dark

ferric oxide dunes are located on the bright

sedimentary unit only. We therefore conclude that

the dark ferric oxide dunes and debris fans are

erosional products of the bright sedimentary

formation. This hypothesis is consistent with

observations, by the Opportunity rover in Meridiani

Planum of (1) stratified outcrops containing both

sulfates and ferric oxides spherules and (2)

spherules accumulations in topographic lows [4, 5].

Figure 1. HIRISE image showing the three spectral units

SU1, SU2 and SU3.

Conclusion: We propose the following history

for Aram Chaos, which accounts for the observed

mineralogical and geomorphological contraints. 1-

filling of an pre-existing crater by sediments. 2-

formation of chaotic terrains at the expense of these

rocks, possibly triggered by sudden withdrawal of

the water stored in the sediments themselves [2]. 3-

second infilling by a dome-shaped, stratified and

bright sedimentary formation containing both

sulfates and ferric oxides spherules. 4- wind- and

gravity-driven erosion of this unit, leaving local

accumulations (debris fans on topographic slopes,

dark sand sheets and dunes on topographic flats and

depressions) of ferric oxides spherules.

References: [1] Glotch T.D. and Christensen P.R. (2005),

JGR 110, E09006, doi:10.1029/2004JE002389. [2]

Oosthoek J.H.P.et al. (2007), LPSC XXXVIII, Abs. #1577.

[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.

3MPI Halbleiterlabor, Otto-Hahn-Ring 6, D-81739

München, Germany. 4Von Hoerner&Sulger GmbH, Schwetzingen, Germany.

5NASA Johnson Space Center,

Houston, Texas, USA. [email protected]

Section 1: In January 2004, the first in situ

extraterrestrial Mössbauer spectrum was received

from the Martian surface. At the present time

(August 2007) the two Miniaturized Mössbauer

Spectrometers MIMOS II on board of the two Mars

Exploration Rovers “Spirit” and “Opportunity”

continue to collect valuable scientific data [1-3].

Both spectrometers are operational after more than

3 years of work. Originally, the mission was

expected to last for 90 days. To date more than 600

spectra were obtained with a total integration time

for both rovers exceeding 260 days. The MER mission has proven that Mössbauer

spectroscopy is a valuable technique for the in situ

exploration of extraterrestrial bodies and the study

of Fe-bearing samples. The Mössbauer team at the

University of Mainz has accumulated a lot of

experience and learned many lessons during last

three years. All that makes MIMOS II a feasible

choice for the future missions to Mars and other

targets. Currently MIMOS II is on the scientific

payload of two missions: Phobos Grunt (Russian

Space Agency) and ExoMars (European Space

Agency).

Section 2: Phobos Grunt is scheduled to launch in

2009. The main goals of the mission are: a) Phobos

regolith sample return, b) Phobos in situ study, c)

Mars and Phobos remote sensing. MIMOS II will

be installed on the arm of a landing module.

Currently, we are manufacturing an engineering

model for testing purposes.

The ESA “ExoMars” mission involves the

development of a MER-like rover with more

complex scientific payload (Pasteur exobiology

instruments, including a drilling system). Its aim is

to further characterise the biological environment

in preparation for robotic missions and eventually

human exploration. Data from the mission will

provide invaluable input to the field of exobiology -

the study of the origin, the evolution and

distribution of life in the universe. The launch date

is scheduled for 2013. Like on MER, the MIMOS

II instrument will be mounted on a robotic arm.

Section 3: Advanced and improved version of

MIMOS II instrument is under development for

those and other future missions. The new design

includes additional mass reduction (total mass is

planned to be 320 g). The dimensions of the

electronic-board will be minimized by using state

of the art digital electronics. A new ring-detector

system (Si-Drift detectors) will be used, thus

greatly improving energy resolution. We expect an

energy resolution of around 140-160 eV for

temperatures lower than 250 K. This will increase

the signal to noise ratio by a factor of 10 and,

therefore, integration times will be reduced

significantly. In addition to the Mössbauer data,

simultaneous acquisition of an X-ray fluorescence

spectrum will be possible, thus providing data on a

sample’s elemental composition. New firmware

will be developed to optimize the instrument’s

performance.

References: [1] Klingelhöfer, G. et al., , Hyp. Int. 170

(2006). [2] Morris, R.V., Klingelhöfer, G. et al., J.

Geophys. Res. 111 (2006). [3] Morris R.V.,

Klingelhöfer, G. et al., J. Geophys. Res. 111 (2006).



ESTEC, Noordwijk, The Netherlands, 12-16 November, 2007

MODELING ERUPTIONS FROM APOLLINARIS PATERA: A STRATIGRAPHIC FRAMEWORK

FROM HIGH RESOLUTION IMAGES. L. Kerber and J.W. Head, Department of Geological Sciences,

Brown University, Box 1846, 324 Brook St., Providence, RI 02912

Introduction: The medium-sized Hesperian-aged

volcano Apollinaris Patera (-8ºS, 174ºE) is located in

a unique position close to the global dichotomy

boundary. It also occurs in proximity to the Mars

Exploration Rover (MER) Spirit landing site in Gusev

Crater, where tephra deposits may be exposed in the

Columbia Hills1.

In preparation for modeling of eruptions from the

volcano2, in combination with atmospheric general

circulation models (GCM)3

to assess the distribution

and fate of eruptive products, we have used a

combination of THEMIS, HRSC and CTX images to

examine the volcano in higher resolution than was

previously possible and to identify several notable

features.

A host of pedestal craters have been detected

suggesting that the flanks and summit of Apollinaris

were covered with tephra-like deposits4. To the north,

a field of yardangs (part of the Medusae Fossae

Formation, MFF) (Figure 1, Area A) show several

separate units which erode in very different ways.

Some are more indurated, eroding to form yardangs,

while others form long, undulating dunes. Still others

appear to be more indurated, with a mode of erosion

more like scour or pitting. These differences in

erosional regime could indicate underlying differences

in composition, volatile content, and depositional

history.

A large fan occurs on the south flank of the

volcano. Several hypotheses have been proposed to

explain the origin of the fan, including overlapping

pahoehoe flows5, pyroclastic flows

6, and fluvial

processes contributing to an alluvial fan6. On the

basis of recent work at Ceraunius Tholus7, we find

similarities suggesting that the fan may have a large

fluvial component. Less discussed is the smaller

topographic bulge on the northeast side of the

volcano (Figure 1, Area B). Its true shape is

obfuscated by two large, fresh impact craters, but

from topography alone it resembles the southern fan,

and the erosion pattern at its base shows the same

characteristic branching, fluvial-like morphology.

The two fans may be related, their formation guided

by a north-south zone of weakness defining the

eastern side of the newest caldera at the summit.

On the southwest flank of the volcano at the edge

of the fan there is a feature that has been mapped

previously as a graben8 (Figure 1, Area C). High

resolution images reveal that parts of the graben

resemble a volcanic vent (Figure 2), and could have

been a source for some of the fan material or for the

western now-disrupted volcanic plains. On the east

side of the edifice is a ~46 km

previously unmapped buried

crater. The two wrinkle ridges

that run along the flank of the

volcano deflect around the rim of

the crater, and the eastern half of

its ejecta is visible as part of an

armored remnant (Figure 1). The

crater likely predates the volcanic

edifice, and subsurface fractures

created during the impact could

have served as conduits for

later magma rise. Mapping of

the stratigraphy of these

deposits is providing insight

into modeling the eruption and

dispersal of near- and far-field

tephra. References: [1] Dalton and Christensen (2006) LPSC XXXVII,

Abstract 2430. [2] Wilson and Head (2007) JVGR v. 163, iss. 1-4,

p. 83-97 [3] Forget et al. (1999) JGR v. 104, pg 24,155-24,176 [4]

Kerber and Head (2007) 46th Vernadsky-Brown Symposium, [5]

Robinson and Mouginis-Mark (1993) Icarus 104, 301-323 [6] R.C.

Ghail and J.E. Hutchison (2003) LPSC XXXIV, Abstract 1775, [7]

Fassett and Head (2007) JGR v. 112, iss. 8, [8] Scott et al. (1993)

USGS Map I-2351.

Figure 2. Central

arrow; vent. Upper

arrow: possible

ancient flow. Bottom

arrow; collapse feature or outflow.

Figure 1. Apollinaris Patera. A) Multi-layered unit with

unique erosional patterns. B) Possible secondary fan. C) Feature resembling a volcanic vent. D) Buried impact crater.




MOLECULAR DETECTION OF LIFE ON MARS

Wilfred F.M. Röling1, Hauke Smidt

2, Henk Leeuwis

3, Erik Laan

4, Pascale Ehrenfreund

5.

1Molecular Cell Physiology, Vrije Universiteit, Amsterdam The Netherlands,

2Laboratory for Microbiology,

Wageningen University, The Netherlands, 3LioniX BV, Enschede, The Netherlands,

4Dutch Space, Leiden, The

Netherlands, 5Astrobiology, Leiden University, The Netherlands. [email protected]

Detecting extraterrestrial life is a challenging task.

The payload of the European Exomars mission,

scheduled for 2013, will carry instruments for

characterization of the organic environment and the

detection of life on Mars. The development of

innovative molecular detection strategies is pivotal

for evaluating the possibility of past or present life

on Mars. While current approaches mainly address

simple biomarkers, we propose to base the search

for extraterrestrial life on complex molecules,

especially hereditable information (DNA). We aim

to address hereditable information as this is one of

the general characteristics of life-forms. Here, we

outline a multi-tiered approach, which includes and

extends on experience in molecular microbial

ecology. The developed strategy will be applied to

Martian analogues on Earth and should allow for the

detection of hereditable information that deviates in

its composition somewhat from that found on Earth.

Key issues that will be addressed in the research

include: (I) developing a robust extraction protocol

for biomarkers, with emphasis on nucleic acids,

from Mars-like materials; (II) developing sensitive,

amplification-based methods for the detection and

characterisation of nucleic acids, including nucleic

acids that deviate in their composition from

terrestrial DNA; and (III) characterisation of the

biodiversity of terrestrial Mars analogues, such as

permafrost, deep subsurface, deserts and extremely

acidic environments, by applying the developed

methodology.




A Monte Carlo Model for Planetary Applications of X-Ray Powder Diffraction and Fluorescence

G. M. Hansford1, H. Su

1, R. M. Ambrosi

1, I. Hutchinson

2, L. Marinangeli

3 .

1University of Leicester, Space Research Centre, Department of Physics and Astronomy, University Road,

Leicester LE1 7RH. UK. 2Brunel University, Uxbridge, Middlesex, UB8 3PH. UK.

3International Research School of Planetary Sciences, Dipartimento di Scienze, Universita' d'Annunzio, Viale

Pindaro 42, 65127 Pescara. Italy.

Contact e-mail: [email protected]

A Monte Carlo model for the simulation of x-ray

powder diffraction and fluorescence is presented.

The model is primarily intended as a tool to aid the

development of a compact instrument for in situ

mineralogical and chemical analyses of planetary

surfaces. In the model, x-rays are produced either by

an x-ray tube or a radioactive source and detected

with a CCD, including accurate quantum efficiency

and energy redistribution effects. Given the

appropriate characteristics, further source and

detector types could readily be added. The sample is

assumed to be an ideal powder of any mineral or

mixture of minerals for which the crystal structures

are available. Additional model elements which may

be included are circular/rectangular apertures,

micropore collimators, and söller slits. Any number

of these elements and powder samples may be

included in a model run. Furthermore, any of the

surfaces in the model (x-ray source, sample, and

detector) may be flat or curved, and in the latter case

the curvature can be spherical or cylindrical.

The utility of this model lies particularly in the

flexibility of the geometrical arrangement of the

various elements, and the quantitative accuracy

which allows realistic integration times to be

assessed. Results of the comparison of two

potentially favourable geometries will be presented.

The first geometry is the parafocusing Seeman-

Bolin arrangement (see, for example, Jenkins and

Snyder [1996]), while the second is a non-focusing

geometry involving a collimated x-ray beam. The

merits of each geometry is elucidated in the context

of a planetary instrument with realistic power,

weight and volume budgets. Figure 1 shows an

example of model output compared with

experimental data for a non-focusing geometry.

References:

Jenkins, R. and R. L. Snyder (1996), “Introduction to

Powder Diffractometry”, Chemical Analysis Vol. 138, p.

180, John Wiley and Sons (New York).

Figure 1. Comparison of experimental (left) and modelled (right) diffraction of Cu-K x-rays from a barite (BaSO4) pressed-

powder sample.




MORPHOLOGICAL ANALYSIS OF A SULFATE DOME IN EASTERN TITHONIUM CHASMA ,

MARS. D. Baioni1, F. C. Wezel

1.

1Institute of Earth Science University of Urbino, Campus Scientifico Sogesta

61029 Urbino (PU), Italy. [email protected]

A morphological study has been carried out on an

elevated domal body located inside the eastern part

of Tityhonium Chasma trough. According to

OMEGA mineralogical data the dome appears to be

constituted by magnesium sulfate (kieserite).

Major features of the dome morphology and

morphometry has been investigated using HRSC,

MOC and THEMIS data.

The observed morphological variations are

interpreted as due to the intensity of erosive

processes (flood discharge or landslide events) that

have locally both destroyed or buried the previous

gully morphology. Several morphologies

caracterizing the domal surface (eg, gully

excavation and lobate depositional features) are

regarded to be connected to a slow flowage motion

caused by the partial melting of interstitial ice in a

periglacial environment.

The correlation between morphological features

and the radial fault patterns over the sulfate dome

has been studied too. The martian dome

characterization should provide useful elements for

the identification of Earth analogues.




A MULTI-INSTRUMENT EXOMARS STUDY OF METEOROID EFFECTS ON THE MARTIAN

ENVIRONMENT A. A. Christou1, A. D. Griffiths

2, J. P. McAuliffe

3, D. Koschny

3, M. Pätzold

4, J. Oberst

5,

J.M. Trigo-Rodriguez6, J. Vaubaillon

7, P. Withers

8, J. E. Chappelow

9.

1Armagh Observatory, College Hill,

Armagh BT61 9DG, UK 2

Mullard Space Science Laboratory, Holmbury St Mary, Dorking, Surrey RH5 6NT,

UK 3

Solar System Missions Div., ESA/ESTEC/RSSD, Keplerlaan 1, NL-2201 AZ Noordwijk ZH, NL 4

Rheinisches Inst. für Umweltforschung, Abt. Planetenforschung, Universität zu Köln, 50931 Köln, DE 5German

Aerospace Center, Rutherfordstr. 2, 12489 Berlin, DE 6Inst. of Space Sciences CSIC-IEEC, Campus UAB, Fac.

Sciencies, Torre C-5, 2a planta, 08193 Bellaterra (Barcelona), ES 7Spitzer Science Center, California Inst. of

Technology, 1200 East California Boulvd, Pasadena, CA 91125 USA 8Center for Space Physics, Boston

University, 725 Commonwealth Avenue, Boston, MA 02215, USA 9

Geophysical Institute, University of Alaska

Fairbanks, 903 Koyukuk Drive, PO Box 757320, Fairbanks, AK 99775-7320, USA [email protected]

Mars, like the Earth, encounters meteoroids of

various sizes, composition and origin during its

orbital trek around the Sun. Those meteoroids' mass

and kinetic energy are incorporated into the Martian

environment through: atmospheric ablation and

deposition of meteoroid constituents in the upper

atmosphere; efficient atmospheric braking leading

to a meteorite on the surface; and hard impact,

resulting in luminous flares (and/or plumes), seismic

shaking and crater excavation [7]. These effects

have been modelled theoretically but in situ

measurements needed to test these models have

hitherto been lacking. The Exomars instrument suite

presents an excellent opportunity to carry out such

observations and compare with similar processes

detected at the Earth and Moon. The following

investigations that we advocate promote synergism

between the different instruments, require no

hardware modification or space qualification of

“soft” mission resources such as inflight software

and provide maximum science for the effort.

Meteor activity at Mars would be punctuated by

annually recurring showers and occasional outbursts

with pronounced effects on the Martian atmosphere

and surface [4,5,6,14]. These, mostly cometary,

meteoroids, have been delivering prebiotic material

to Mars for the past 4.5 Gyr. As the present Martian

atmosphere has similarities with that of the early

Earth, the astrobiological relevance of meteor

showers as exogenous sources of organics and water

for both Earth and Mars is obvious.

These events can now be predicted with sufficient

reliability both at Mars [6] and the Earth (eg [8,9])

to justify targeted observational campaigns.

Relevant measurements include: dual-eye

panoramic camera detection of visible meteors in

the Martian sky using existing flight-qualified

change-detection software to minimise data volume

[10]; radio occultation height profiles of ionospheric

electron density during the orbital phase of the

mission [12] and of the total electron content (TEC)

post-landing; and seismic detection of impact event

clusters correlated with Mars' passage through low-

speed meteoroid streams [11].

Decimetre-to-metre size craters are theoretically

expected on the Martian surface due to the influx of

specific meteoroid subpopulations, eg cm-sized M-

type asteroidal fragments [1,13]. Pit-like formations

of this size have been observed by Opportunity

although their origin, whether impact-related or

otherwise, remains a mystery. Observing such pits

would lead to estimates of their area density, and

characterise the mechanisms that destroy them over

time such as dust infilling. A combination of

panoramic and hi-res camera observations is well

suited to this task and will determine the present

hazard from such meteoroids on surface activities.

Meteorites, particularly rare nickel-irons, have

recently been identified on the Martian surface [15].

The area density and size distribution of those and

other, more common, meteorite classes are sensitive

to atmospheric density [2,3] and can be used as

proxies for past climate variations. Identification of

such meteorites using imaging and spectroscopy

during the landed part of the mission will provide a

unique insight on the variation of the Martian

environment with time.

Apart from their role in fulfilling the mission

goals of characterising the biological environment

on Mars in preparation for robotic missions and

human exploration, these investigations hold a

significant potential for communicating to the public

the excitement of exploring Mars and the sense of

"being there". Public release of selected data

products eg images of meteors and fireballs against

the Martian sky are bound to have a positive impact

on the public perception of European planetary

exploration.

References: [1] Chappelow, J. E. and V. L. Sharpton

(2005), Icarus 178, 40-55 [2] Chappelow, J. E. and V. L.

Sharpton (2006), Icarus 184, 424-435 [3] Chappelow, J.

E. and V. L. Sharpton (2006), GRL 33, CiteID L19201 [4]

Christou, A. A. (2004), EM&P 95, 425-431 [5] Christou,

A. A. and K. Beurle (1999), P&SS 47, 1475-1485 [6]

Christou, A. A. et al. (2007), A&A 471, 321-329 [7]

Christou, A. A. et al. (2007), P&SS,

doi:10.1016/j.pss.2007.05.001 [8] Jenniskens, P. and J.

Vaubaillon (2007), WGN 35, 30-34 [9] Jenniskens, P.

(2007), CBET 1049 [10] McAuliffe, J. P., and A. A.

Christou (2006), Proc IMC 2005, 155-160 [11] Oberst, J.

and Y. Nakamura (1991), Icarus 91, 315—325 [12]

Patzold, M. et al. (2005), Science 310, 837-839 [13]

Popova, O. et al. (2003), MP&S 38, 905-925 [14] Selsis,

F. et al. (2004), A&A 416, 783-789 [15] Weitz, C. et al.

(2006), AGU Fall Meeting 2006, Abs #P41B-1268.




MULTI-INSTRUMENT OBSERVATIONS OF AURORA-TYPES EVENT BY MARS EXPRESS

Leblanc F.1, Witasse O.

2, Lilensten J.

3, R. A. Frahm

4, Ali Safaeinili

5, D. A. Brain

6, J. Mouginot

3, J.L. Bertaux

1,

J. D. Winningham4, W. Kofman

3, R. Lundin

7, J. Halekas

6 and M. Holmström

7 1 Service d’Aéronomie du

CNRS/IPSL, Verrières-le-Buisson, France. 2 Research and Scientific Support Department of ESA-ESTEC, The

Netherlands. 3. Laboratoire de Planétologie de Grenoble, France. 4. Southwest Research Institute, San Antonio,

TX 78228-0510, USA. 5. Jet Propulsion Laboratory, Pasadena, CA 91109, USA. 6. Space Sciences Laboratory,

University of California, Berkeley, USA. 7. Swedish Institute of Space Physics, Box 812, S-98 128, Kiruna,

Sweden [email protected]

We present a new set of observations of Martian

aurora obtained by SPICAM UV spectrometer on

board Mars Express (MEX). Several auroral

emissions are identified on the Martian night side

near crustal magnetic fields. During several orbits

consecutive events separated by several tens of

seconds are observed, highlighting the role of

closed and open field line structures in shaping

spatially these events. For most of these events

coordinated observations with MARSIS and

ASPERA-3 on board Mars Express were possible.

ASPERA-3 is composed of an ion mass analyzer

(IMA), of two neutral particle imager (NPI and

NPD) and of one electron spectrometer (ELS). For

these particular events, data from the electron

spectrometer were available so that a simultaneous

measurement of the precipitating electron flux was

possible. MARSIS is a multifrequency synthetic

aperture orbital sounding radar which monitors in

particular the Total Electron Content (TEC) and

which was operating for some of these events. At

the end, SPICAM UVS is a UV spectrograph

covering the spectral range between 110 and 300

nm and which measures the atmospheric glow. It is

this latter instrument which clearly provides the

spectral evidence of the occurrence of an auroral-

event. In order to avoid any ambiguity on the

positions of the simultaneous measurements, we

used orbits of MEX during which SPICAM UVS

field of view was nadir oriented. This new set of

observations shows quite strong coincidences

between the occurrence of energetic precipitating

electrons into the Martian atmosphere, the increase

of the TEC, the presence of crustal magnetic field

anomalies and auroral-type glow. Following the

definition of [1] of open / closed magnetic field

lines, we observe that the aurora detected by

SPICAM UVS occur on open field lines (Figure 1).

This conclusion therefore suggests a significant

relation between aurora events at Mars and the

presence of cusp like magnetic field line structures.

Figure 1: Map of the probability (expressed in

percentage) to be in an open field line region at 400 km

in altitude on the Martian nightside as calculated from

the Electron Reflectometer on board Mars Global

Surveyor (MGS) during its mapping phase orbit.

Electron pitch angle distributions have been recorded by

MGS magnetometer/electron reflectometer every 2-8s

during the spacecraft mapping orbit phase at ~400 km.

Pitch angle distributions recorded in a single instrument

energy channel (115 eV - a channel typically

uncontaminated by photoelectrons) have been classified

according to their shape. This figure shows the

probability, as a function of geographic location on the

Martian night side, of observing a statistically greater

flux of electrons returning from the planet than moving

toward the planet [1]. Also plotted are the trajectories of

Mars Express (white solid line) below 1000 km altitude

for the four orbits with aurora events. The white crosses

indicate the position of the aurora events identified in

SPICAM UVS data. The altitude of the spacecraft at that

time is indicated. The spacecraft was moving from

Northern to Southern hemispheres.

Reference [1] Brain D.A., Lillis R.J., Mitchell D.L.,

Halekas J.S. and R.P. Lin, Electron Pitch Angle

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.





NEW NUMERICAL ESTIMATIONS FOR BOTH IMPACT CONDITIONS AND HYDROTHERMAL

ZONES ON ISIDIS PLANITIA, MARS J. C. Echaurren1.

1Codelco Chile Chuquicamata, North Division,

Pasaje Lince 976, Calama, Chile. [email protected]

Introduction: Mars’s Isidis Planitia [1] is one of

the largest impact Panitias on Mars with a diameter

of about 1,238 km. Isidis is located at N 14.1 deg

and W 271.0 degrees and is the boundary between

ancient highlands and the Northern Plains. Isidis

Planitia is a relatively unique impact basin [2] on

Mars as it has a clear, large gravity anomaly,

reminiscent of lunar mascons, and circumferential

tectonic features consistent with lithospheric

flexure. Curiously, the tectonic signature is limited

to a portion of the basin’s periphery at Nili and

Amenthus Fossae, zones of circumferential

extensional faulting, to the NW and SE,

respectively, of Isidis. Moreover, the basin is

bounded to the west by the Syrtis Major volcanic

province and to the E-NE by Utopia Planitia, each

capable of imposing their influence on the

deformation and state of stress in the lithosphere in

the vicinity of the basin. The aim of this work is to

estimate the impact conditions and predictions in

relation to the generation of hydrothermal systems,

using mathematical models. For the calculations

will be used both diameter and shape of the crater,

and chemical composition of the target rock on

Isidis Planitia.

Numerical results: According the models used for

this basin [3], the diameter of asteroid is calculated

in ~ 186.53 km, with both velocity and impact angle

on the martian surface of ~ 20.3 km/s and 53.26°

respectively. The number of rings on the crater are

calculated in ~ 35.25 with a initial crater profundity

of ~ 4.8 km, the melt volume is ~ 5.75E15 m or ~

5.75E6 km . The number of ejected fragments are

estimated in ~ 2.53E13 or ~ 25,325.4 billion of

fragments, with average sizes of ~ 6.35 m, and a

cloud of dust with diameter of ~ 3.25E14 m. The

total energy in the impact is calculated in ~ 5.14E32

Erg (1.22E10 megatons). Before of the erosion

effects the transient crater is estimated in ~ 828 km,

the hydrothermal zone (hydrothermal systems) is of

~ 97.01 km to 413.98 km from the nucleus of

impact, i.e., a hydrothermal zone of ~ 316.97 km.

The density of this asteroid (or comet) is calculated

in ~ 0.270 g/cm . The seismic shock-wave

magnitude is calculated using linear interpolation in

~ 8.63 in the Richter scale. The temperature peak in

the impact is calculated in ~ 1.18E17 ºC (~ 7.84E9

times the temperature of the solar nucleus), by a

space of time of ~ 25.4 ms. The pressure in the final

crater rim is calculated in ~ 8.88 Gpa, and the

pressure to 1 km of the impact point is ~ 3.41

millions of Gpa. The maximum density for the

fragments is calculated in ~ 0.27 g /cm , and the

combined density for these fragments is calculated

in ~ 0.23 g /cm . A scheme of the hydrothermal

zone is showed in figure 1, according the numerical

results obtained.

Figure 1. Hydrothermal zone according the numerical

estimations obtained.

Future works will be more precise in the determination of numerical results. References: [1] Scott, D., and Tanaka, K, (1986)

Geological Survey Misc. Inv. Map, I-1802-A. [2] J.

Andreas Ritzer and Steven A. Hauck, II, (2007) Lunar

and Planetary Science XXXVIII, 2244.pdf. [3] Echaurren

J., and Ocampo A.C., (2003) EGS-AGU-EUG Joint

Assembly.




NUMERICAL SIMULATION OF THE TEMPERATURE OF MARS INTERIOR TO INFER

THE SUBSURFACE STRUCTURE Fabio Gori1

1Department of Mechanical Engineering, University of

Rome “Tor Vergata”, Rome, Italy [email protected]

Temperature distribution and oscillation inside

Martian regolith is very important for Mars

exploration and for the understanding of the inside

structure. The absence of temperature measurements

in situ can justify the investigation of temperature

variation with depth and time oscillation during the

day by means of numerical simulations.

Theoretical predictions of the temperature

distribution in a layer deep five meters have been

carried out in [1-2]. The boundary condition at the

surface with the Martian atmosphere has been of an

imposed temperature oscillation during the day. A

more real thermal boundary condition is that of an

imposed convection and radiation on the surface.

Also the thermal boundary condition on the bottom

layer of the investigated regolith can be modified

taking into account a geothermal heat flux.

Numerical simulations can be performed for soils

with different porosity and different thermal and

physical properties.

References:

[1] F. Gori and S. Corasaniti, Theoretical prediction of the

thermal conductivity and temperature variation inside

mars soil analogues, Planetary and Space Science, 52 (1-

3) pp. 91-99, 2004.

[2] F. Gori and S. Corasaniti, Thermal Properties and

Temperature Variations in Martian Soil Analogues, In

Maravell N. S., Space Science: New Research, Chapter 6.

New York: Nova Science Publishers (USA), 2006.




NUMERICAL SIMULATION OF THERMAL MEASUREMENTS IN MARTIAN REGOLITH R.

Nadalini1,2

, N. Schmitz1, G.Messina

1, J.Knollenberg

1.

1DLR Institute of Planetary Exploration, Berlin, Germany,

2Active Space Technologies GmbH, Berlin, Germany. [email protected]

Planetary heat flow is an important indicator of

the internal temperature and heat transfer

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.




NUMERICAL SIMULATION OF THERMAL MEASUREMENTS IN MARTIAN REGOLITH R.

Nadalini1,2

, N. Schmitz1, G.Messina

1, J.Knollenberg

1.

1DLR Institute of Planetary Exploration, Berlin, Germany,

2Active Space Technologies GmbH, Berlin, Germany. [email protected]

Planetary heat flow is an important indicator of

the internal temperature and heat transfer

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.




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.,

2005. Astron. Astrophys. 441, L9–L12. [4] Bibring, J.-P.,

2004 ESA-SP 1240, 37–49

Fig 1: Top: Visible RGB reconstruction of the SPR, no

clouds are detected; Middle Top CO2 ice detection

through the 3 m band; Middle Bottom water vapour

detection in the morning (5-9 AM); Bottom, water

vapour detection in the evening (2-5 PM). Saturation

(white regions) in the water vapour detection should be

associated to a possible contamination of data by CO2 ice

and should not be considered.

Fig 2: From left to right three periods are represented:

LS=250° , 260° and 270°. Top: water vapour variability

as a function of the local time (p is the intersection and m

the angular coefficient for the linear best fit). Bottom:

water vapour variability as a function of the incidence

angle (negative values stand for evening).




OMEGA-PFS OBSERVATIONS OF A LOCAL DUST STORM ON MARS. F. Altieri

1, G. Bellucci

1, F.G.

Carrozzo1, D. Grassi

1, L. Zasova

2, J. P. Bibring

3, V. Formisano

1,

1INAF-IFSI, Via Fosso del Cavaliere 00133

Rome, Italy, 2IKI, Moscow, Russia,

3 Institut d'Astrophysique Spatiale, University of Paris-Sud, Orsay Paris,

France. [email protected]

Dust plays an important role in current

Martian climate. The particulate component of the

Mars atmosphere is composed of micron-sized

particles, which are products of soil weathering, and

water ice clouds. In the absence of a dust storm, a

so-called permanent dust haze with opacity

0.05–0.2 in the atmosphere of Mars determines its

thermal structure. Dust loading varies substantially

with the season and geographic location. Opacity

may reach several units during a dust storm.

In this work we report on the observation

of a dust storm in the Atlantis Chaos region on

Mars, observed by the OMEGA [1] instrument on

board of Mars Express. The observation was done

on March 2nd, 2005 at 11.00 LT and Ls = 168° (end

of southern winter), Figure 1 and 2. At the same

time, also the PFS instrument [2] took high

resolution spectra over the region, thus allowing to

retrieve the pressure and thermal profiles with the

altitude. These joint observations constitute a unique

data set which allows to study both the physical and

scattering properties of the suspended dust and the

mechanism of formation of dust storms on Mars.

Moreover, the study of airborne dust can allow to

better constrain its spectral behavior in order to

decouple its effect from the surface mineralogy.

References: [1] Bibring, J-P., et al. (2004), ESA SP-

1240, 37 - 49. [2] Formisano, V., et al. (2004), ESA

SP-1240, 71 – 94.

Figure 1 – a) RGB composite image (R = 0.68 μm, G = 0.53 μm, B = 0.43 μm). Suspended dust can be seen in the

bottom part of the image. – b) Temperature map derived from the 5 μm OMEGA radiance. The dust storm exibhits the

lowest temperature in the scene.

Figure 2 Altitude of the dust storm compared to the surface topography. The temperature along the segment shown in Fig.

1-b) is plotted in red.

a) b)




Physical and chemical characterization of terrestrial carbonates of biotic and abiotic origin in the frame

of the future in situ investigation of the surface of Mars F. Stalport1, P. Coll

1, C. Szopa

2.

1Laboratoire

Interuniversitaire des Systèmes Atmosphériques (LISA), University Paris XII, Créteil France; 2Service d’Aéronomie

(SA), University Paris VI, Paris France. Contact : [email protected]

The collection of independent and various data

with the future space probes to Mars should be

necessary to point out and confirm the presence for

past or present life on Mars, if it ever existed.

Among the targets to consider to reach that goal, we

propose minerals produced from a past biological

activity, named biominerals. Indeed, it seems today,

considering the recent advances done thanks to the

MERs and MEx space probes, that early Mars

owned a denser atmosphere, probably made of CO2,

and a mild climate, allowing liquid water to stand at

the surface for long periods. Similar environmental

conditions led to the origin of life on the Earth more

than 3.5 billion years ago; and led to the production

of large amounts of biominerals such as carbonates.

Such a process could have taken place on Mars,

even if carbonates still not have been detected in

large amounts on the red planet (Bibring et al.,

2005).

Also, we investigated the physical and chemical

properties of terrestrial carbonates, with the aim to

evaluate the possibility to identify biotic carbonates

from their abiotic form with a simple diagnostic

possibly transposable to in situ exploration.

Carbonates are interesting because they are

produced on Earth both from abiotic and biotic

processes. We assumed that crystalline defects and

trace elements in the crystal lattice, as well as the

larger growth speed of biotic calcites, must be

responsible for differences between the physical and

chemical properties of carbonates.

We investigated numerous different terrestrial

carbonate samples, of different structures (calcites,

aragonites), and from various origins (biotic,

diagenetic and abiotic). The minerals were studied

by X-ray diffraction and electron scanning

microscopy to determine their mineralogical and

chemical composition, and differential thermal

analysis coupled to thermogravimetric analysis

(DTA-TG) to determine their thermal behavior.

Our results show that the thermal degradation of

abiotic carbonates occurs at a temperature at least

20°C higher than the degradation temperature of any

biotic carbonate investigated (see figure 1 for calcite

[2]). Consequently, in the case of a Martian in-situ

exploration, or in a sample return mission, the

analysis of Martian minerals by DTA-TG represents

a promising approach to provide one of the clues for

a past biological activity on Mars.

References: [1] Bibring, J.-P. et al. (2005) Science 307,

1576-1581. [2] Stalport F. et al. (2005), GRL 32, L23205.

0

1

2

3

4

830,00 840,00 850,00 860,00 870,00 880,00 890,00 900,00 910,00 920,00 930,00

Température (°C)

Figure 1. Temperatures of degradation of various calcites: in red, abiotic ones; in green, biotic ones; in purple, diagenetic

ones (mix between biotic and abiotic minerals)




POSSIBILITIES OF TERRAFORMING MARS M.K. Shrivastava Department of Geography B.S.P.,

H.S.School Risali Sector, Bhilai Nagar, Durg (C.G.), INDIA [email protected]

“Earth’s active internally produced energy (hot

mantle) constantly sends the submerging seawater

back to the surface (through Mid Oceanic Ridges /

MOR, Volcanoes, hot springs etc.) by vaporization

and thus also keeps effluence of huge amount of

CO2 (Dissolved Inorganic Carbon / DIC) and other

dissolved gases alive, from submerging seawater to

the atmosphere. In this way, Earth’s active

internally produced energy (hot mantle) prevents

the entire surface water from getting submerged

into its subsurface along with the huge amount of

DIC and other dissolved gases and is responsible

for constant existence of surface water, atmosphere

and greenhouse effect on Earth. Diminished

internally produced energy of early Mars would

have resulted into cold mantle. While getting cold

the volume of Martian liquid mantle would have

reduced because of constriction due to

solidification. Then the solid Martian crust might

have had adjusted itself over the cooling mantle

creating many crakes in the crust and gaps at many

places between Martian cold mantle and crustal

base while shifting of crust on the mantle. These

gaps and crakes would have acted as sufficient

reservoir for submerging Martian surface water.

Therefore, diminishment of internally produced

energy of earlier Mars would have resulted in

gradual submersion of the entire Martian surface

water into its subsurface and some interior (which

could not return back to the surface due to cold

Martian mantle) along with a large amount of DIC,

breaking the efflux of CO2 from entire submerging

surface water to the atmosphere, however its influx

remain continued. It would have caused

disappearance of surface water and poorer green

house effect further cooling the Martian

atmosphere. Similarly other dissolved gases might

also have submerged along with Martian surface

water causing thin atmosphere and very low

surface temperature on Mars. Melting of Martian

polar and subsurface ice by increased green house

effect or bombardment of asteroids, etc. would

make liquid water available on Martian surface but

this melted water will again get submerged

gradually, with the dissolved gases into its

subsurface and will not return back due to

diminished internal energy production (cold

mantle) of Mars. Hence terraforming or

revivification of Mars will be possible only when

its diminished internal energy production is got

regenerated or reactivated to make its mantle hot

again. Only then, the submerged water (subsurface

ice), trapped CO2 and other gases will return back

and exist constantly on the Martian surface and in

its atmosphere. Without this all the efforts to

terraform or revive Mars would ultimately result in

failure. But such technology which can regenerate

or reactivate the diminished Martian internal

energy production has not been developed so far

and its possibility in near future also seems to be

negligible. So, to terraform or revive Mars, we

should first think that in future, can we ever

reactivate or regenerate the diminished Martian

internal energy production? As this is an

impossible task with in the present frame of

knowledge. In future Earth will have to encounter

similar conditions like present day Mars, when

Earth’s internally produced energy will also get

diminished”.




POSSIBLE MARTIAN LANDING SITES TO BE CONSIDERED FOR FUTURE EUROPEAN

EXPLORATION MISSIONS P. D. Martin. European Space Agency, European Space Astronomy Centre

(ESAC), Villafranca del Castillo (Madrid), Spain. [email protected]

Introduction: The selection of landing sites for

Mars missions typically follow a roadmap such as

represented in Figure 1, implying a required number

of iterations that must reconciliate landing site

engineering constraints with the scientifically-driven

selection process and lead to the identification of

prime and backup landing sites. Pinpointing with

precision a number of landing sites for future Mars

missions is now possible thanks to the wealth of

scientific data and high-resolution mapping products

resulting from recent and ongoing successful Mars

orbiter missions. The main goal of this work is to

consolidate available mapping products (e.g.,

geological, hyperspectral and compositional) in

order to support the selection process of candidate

landing sites for future European Mars missions.

Figure 1. Mars landing site selection roadmap.

Results: A preliminary investigation was

carried out, assuming a set of landing requirements

[1]. Possible landing regions on Mars resulting from

this preliminary investigation were categorised into

two classes, depending on the level of risk assessed

for the landing, as summarized below:

• Low-risk regions: Amazonis Planitia,

Utopia Planitia, and Elysium Planitia. One

of their potential drawbacks is that most

areas of these regions exhibit a relatively

high dust index [2] which could be

detrimental to the scientific interest of the

in-situ mission.

• Moderate-risk regions:

o Syrtis Major / Nili Fossae, where

phyllosilicates and hydrated

minerals can be found based on

evidence from orbit (Mars

Express/OMEGA [3]).

o Isidis Planitia, in particular

because this region presents a low

vertical roughness [4].

o Chryse/Acidalia Planitia, where

phyllosilicates, hydrated minerals

and sulfates can be found [3].

o The region that spans the terrains

from Sinus Meridiani to Syrtis

Major, between 15ºS and 45ºN.

This region exhibits a high dust

index, and is represented by

rougher, heavily cratered terrains

in many areas.

Within these regions, a more detailed

identification of landing sites has been started by

refining the study (top-down approach) using

higher-resolution geological and compositional

maps coupled with other parameters and constraints.

Preliminary results lead to the following, non-

exhaustive list of candidate landing sites:

• Nilo Syrtis Mensae

• Nili Fossae region

• Mawrth Vallis

• Gale Crater

• North Meridiani

• Candor Chasma

• Eos Chasma

• Juventae Chasma

• South Olympus Mons

These areas may be considered as a first set of

study zones as part of the science-driven and

success-oriented selection process for future Mars

missions such as Exomars. The results shall then be

confronted with bottom-up approaches consisting in

the pre-selection of sites purely based on scientific

goals prior to the assessment of their suitability for

landing.

References: [1] Martin (2007), EPSC abstract #

EPSC2007-A-00305. [2] Ruff and Christensen (2002),

JGR 107. [3] Bibring et al. (2006), Science 312. [4]

Kreslavsky and Head (2002), GRL 29.




Preliminary laboratory XRD/XRF instrument tests and evaluation

H Su1, G. M. Hansford

1, R. M. Ambrosi

1, A. F. Abbey

1, D. Vernon

1, J. Sykes

1, I. B. Hutchinson

2,

L. Marinangeli3, A. Stevoli

4

1University of Leicester, Department of Physics and Astronomy, University Road, Leicester, LE1 7RH, UK.

2Brunel University, Uxbridge, Middlesex, UB8 3PH. UK.

3International Research School of Planetary Sciences, Dipartimento di Scienze, Universita’ d’Annunzio, Viale

Pindaro 42, 65127 Pescara. Italy. 4Thales-Alenia Space, S.S. Padana Superiore 290, Vimodrone, Milano, Italy

Contact e-mail:[email protected]

Section 1:

The results of preliminary laboratory tests carried

out by the University of Leicester are presented.

The tests were carried out in a Mars environment

simulation chamber in support of the MARS-XRD

project. MARS-XRD is a combined X-ray

diffraction and fluorescence instrument. It is also

one of the pre-selected instruments for the ESA

ExoMars mission scheduled for launch in 2013 [1,2].

X-ray diffraction is an analytical technique used to

determine the crystallographic structure of minerals,

while X-ray fluorescence is commonly used to

determine the chemical composition of rock samples.

Although well developed for terrestrial applications,

XRD has not yet been applied in a planetary context

so far.

A Mars environment simulator test chamber,

which consists of X-ray source, source collimator

tube and rotating sample table, was designed and

built by the University of Leicester. The CCD is an

XMM EPIC CCD22 manufactured by e2v

Technologies and is fitted on a motorised rotating

arm. The laboratory experiments were carried out in

a representative Mars environment (in terms of

temperature and CO2 atmospheric pressure).

Theoretical calculations were performed

simultaneously to verify the results.

The tests were aimed at exploring the

spectroscopic performance of the detector and

determine the impact of source intensity, X-ray

energy and geometry on the signal to noise levels

recorded and the measured XRD spatial resolution.

Various samples were used in the tests. Figure 1

shows an example of the XRF spectrum and XRD

pattern of a barite sample.

References:

[1] Marinangeli et al. (2005) Geoph. Res. Abs., Vol. 7.

[2] Marinangeli et al. (2006) Eos Trans. AGU, 87(52),

Fall Meet. Suppl., Abs. P51D-1227.

Figure 1. XRF spectrum of Barite campact powder sample (left); XRD image of Mn-Ka for 2-theta angle 30~50 degree

(middle) and XRD spectrum with calculated resolution for each diffraction bands from selected CCD area x-axis: 300-350

pixels (full width is 0-610 piexels), y-axis:0-602 pixels (right). Details are in the text in the graph.




PRELIMINARY RESULTS OF SURFACE STUDY BY MARSIS J. Mouginot

1, W. Kofman

1, A.

Safaenieli2, J. Plaut

2, G. Picardi

3.

1Laboratoire de Planétologie de Grenoble, CNRS/UJF 38041 Grenoble Cedex,

2Jet Propulsion Laboratory, California Institute of Technology, Pasadena, CA 91109.

3Infocom Departement,

“La Sapienza” University of Rome. [email protected]

Introduction: For two years, the Mars Advanced

Radar for Surface and Ionosphere Sounding

(MARSIS) on board Mars Express is collected data

[1]. This large survey provides an opportunity to

study the surface reflectivity of the Mars planet.

We describe here preliminary results for the

surface study. In the first section, we present our

method to extract from the radar signal, the power

reflected by the surface. In the following one, we

show the different phenomena that control surface

reflectivity. We discuss the method to correct

ionospheric absorption.

Surface detection: After the correction for

ionospheric effects (phase shift correction)[2], we

can start to calibrate MARSIS echoes and study the

first echo reflectivity, which corresponds to the

surface echo.

First we present our method to extract the surface

echo from each frame. In all of our correction, we

use MOLA topography as reference. So we can

easily select the surface echo in the signal as it

correspond to MOLA altitudes (see Figure 1).

Figure 1. Surface echo detection example. Upper figure

is, MARSIS radargram of the orbit 2682 and at the bottom

the position of surface echo selected by our routines is

shown.

Ionospheric absorption: The final aim of this

work is to estimate the dielectric constant of martian

surface at MARSIS frequency. As MARSIS

wavelength is around 100 m, the dielectric constant

corresponds to 100 m depth materials column.

Unfortunately dielectric is not only parameter

that controls surface reflectivity. In fact, the

parameters that controls surface reflectivity are:

material composition (dielectric constant), surface

roughness, local slope, [3] and ionospheric

absorption.

First of all, it’s appears that the surface

reflectivity seen by MARSIS has a strong

dependence on the solar zenith angle. In order to see

this SZA dependence, we must to separate all

surface geometry or roughness effects that affect the

signal. So we have selected a very flat area in the

south polar layered deposits (latitude -81°,-85°;

longitude 180, 210). For this flat area, we have plot

in figure 2 surface reflectivity as function of solar

zenith angles.

Figure 2. Surface reflectivity seen by MARSIS in region

at latitude [-81, -85] and longitude [180, 210] as function

of the solar zenith angles.

As we can see on the figure 2, the reflectivity

can be 15 dB smaller during the day (for small SZA)

than during the night (large SZA > 100).

This absorption is due to a well known effect of

ionospheric absorption [4][5] .

We present here our method to compensate this

effect. We use the region as reference and apply this

compensation on all data.

Finally we show a preliminary surface

reflectivity map. We discuss the reflectivity in terms

of material composition (dielectric constant),

surface roughness, local slope and material density.

References: [1] Picardi G. et al. Science, 2005, 310,

1925-1928. [2] Mouginot J. et al. PSS submitted [3] Ulaby

F.T. et al. Artech House Publishers 1982. [4] Nielsen, E.

et al. PSS , 2007, 55, 864-870. [5] Safaeinili, A. et al. PSS,

2003, 51, 505-515.




PRESERVATION OF ANCIENT ICE IN TROPICAL MOUNTAIN GLACIER DEPOSITS ON MARS:

James W. Head1 and David R. Marchant

2.

1Dept. Geological Sciences, Brown University, Providence, RI 02912

USA. 2Dept. Earth Sciences, Boston University, Boston MA 02215 USA. ([email protected])

Introduction: Analysis of extensive fan-shaped deposits

on the NW flanks of the equatorial Tharsis Montes (Fig. 1)

with new data provide compelling evidence that they repre-

sent the remnants of tropical mountain glaciers (TMG) dat-

ing from the Late Amazonian [1-2]. The distinct geomor-

phology of the deposits, together with updated terrestrial

analogs for glaciation under martian hyper-arid, extremely

cold conditions [3], show that the tropical mountain glaciers

were cold-based. Global climate models show that when

obliquity reaches 45 degrees, water-rich polar air ascends the

flanks of Tharsis, encounters the NW flanks of the Tharsis

volcanoes, undergoes upwelling and adiabatic cooling, pre-

cipitating snow on the northwest flanks [4]. Models of ac-

cumulation and glacial flow show that this scenario can pro-

duce tropical mountain glaciers [5].

On Earth, when glaciers retreat, ablation can result in an

increase of debris on top of the glacier (sublimation till); this

deposit can significantly decrease the sublimation rate and

protect the buried ice from further loss of ice, preserving it

for long periods [6]; in the Antarctic Dry Valleys, ice buried

below sublimation till may be as old as 8 million years [7].

Is there any evidence of similar remnant ice in the tropical

mountain glaciers on Mars?

Description and interpretation: Arsia and Pavonis TMG

deposits (Fig. 1) consist of three basic facies, ridged, knobby

and smooth [1-2]. The proximal smooth facies consists of

lobate, relatively smooth-textured deposits interpreted as the

remnants of individual cold-based glacial lobes (alpine-like

glaciers), emplaced in the waning stages of glaciation. De-

bris-covered cold-based glaciers build up a protective subli-

mation till derived from supraglacial and englacial debris.

As glacial conditions wane, ice is often preserved longest in

the distal portions, where the insulating effect of the till is

greatest, producing thick arcuate lobes. Similar arcuate lobe

configurations are seen at Arsia and Pavonis (Fig. 1).

Could these lobes, morphologically and environmentally

similar to those seen on Earth, still contain remnant glacial

ice from the Late Amazonian glaciation many tens of mil-

lions of years ago? Analysis of high-resolution image and

topography data reveal the presence of several crater-like

depressions in the smooth facies at Pavonis and Arsia. These

feature are shallower than fresh impact craters of similar

diameters and show significant evidence of having under-

gone viscous relaxation. They have several zones: An inner

hummocky, but often oyster-shell like floor with outward-

facing scarps; an intermediate zone beyond the apparent

crater rim of concentric ridges and troughs, a narrow zone of

closely-spaced fractures, and an outer zone of hummocks

oriented along the regional trend of the lobe, sometimes with

superposed secondary craters.

In summary, TMG deposits on Mars record ancient cli-

mates when planetary spin-axis obliquity was in excess of

45°, and polar volatiles were mobilized and transferred equa-

torward. We interpret the set of unusual impact craters su-

perposed on these deposits to indicate that the impact pene-

trated a veneer of sublimation till and excavated buried rem-

nant glacial ice, subsequently undergoing viscous relaxation.

Remaining deposits may be hundreds of meters thick. The

deposits are Late Amazonian in age and the remnant ice may

preserve records of ancient atmospheric gas content and

microbiota, as is seen in terrestrial glacial ice [8].

References: 1) Head and Marchant, Geology, 31, 641, 2003;

2) Shean et al., JGR, 110, 05001, 2005; 3) Marchant and

Head, Icarus, in press, 2007; 4) Forget, et al.. Science, 311,

368, 2006; 5) Fastook et al., LPSC 37, 1794, 2006; 6) Kow-

alewski et al., Ant. Sci., 18, 421, 2006; 7) Marchant et al.,

ISAES, 54, 2007; 8) Bidle et al., PNAS, 104, 13455, 2007.

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

daytime

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/).




PULSED REMOTE RAMAN SYSTEM FOR PLANETARY SURFACE EXPLORATION Fernando Rull.

Unidad Asociada UVA-CSIC al Centro de Astrobiología Facultad de Ciencias, Universidad de Valladolid,

47007-Valladolid (Spain) [email protected]

Raman spectroscopy is a powerful technique for

minerals and organics analysis. Raman in

combination with LIBS (Laser Induced Breakdown

Spectroscopy) is now part of the Pasteur payload in

ExoMars mission1.

This combined spectrometer will perform spectral

analysis in close-contact mode inside at outside of

the rover.

Nevertheless for surface analysis in planetary

missions with landers and rovers the possibility of

remote characterisation of materials show clear

advantages over the contact mode. These advantages

seem to be particularly important in the case of the

future sample return missions in which reliable

identification of the potential samples will surely

become a crucial task.

Remote Raman spectroscopy has demonstrated its

potential in several applications2,3

up to 217 mtrs!4.

The remote Raman system will be also useful for

detecting hydrocarbon plumes and gas hydrates on

planetary surfaces.

We report in this work the concept, principles of

design and results obtained in the field with a

remote Raman prototype working in the range 5 to

25 meters.

The system consists in a compact spectrograph

fibre-optic (FO) coupled with a telescope. The laser

excitation is performed by a frequency-doubled

Nd:YAG pulsed laser (20 Hz, 4ns, 532 nm, 35

mJ/pulse) in coaxial geometry with the telescope.

The detection is made by a gated intensified charged

couple device (ICCD) detector. Gated mode (in the

range 20-80 ns) shows particular advantages over

the continuous (CW) mode of operation in reducing

the background signal and eliminating long-lived

fluorescence signals from the Raman spectra. Gated

mode is also very useful for daylight operation.

The remote Raman system is full computer

controlled for laser pointing, sample focusing, laser

shooting and spectra acquisition.

Results obtained in field operation in Rio Tinto

(Spain) and recently in AMASE 2007 expedition at

the Artic (Svalbard Islands) are presented and

discussed

Figure 1. The remote Raman prototype deployed at

Rio Tinto (Spain)

References: 1- Science Management Plan for ExoMars, EUROPEAN SPACE AGENCY, HUMAN SPACEFLIGHT, MICROGRAVITY AND EXPLORATION PROGRAMME BOARD, Paris, 6

March 2007. 2- Sharma S.K., Angel M.S., Ghosh M.,

Hubble H.W., and Lucey P.G. (2002). Applied

Spectroscopy, 56, 699-705. 3- Sharma S.K., Lucey P.G.,

Ghosh M., Hubble H.W., and Horton K.A. (2003).

Spectrochim. Acta A59, 2391-2407. 4- Chen T, Madey

TJM, FRANK M. Price FM, Sharma SK and Lienert B.

Applied Spectroscopy (2007) 61, 624-629.




QUANTIFYING PHOTOCHEMICAL CATALYTIC CYCLES NEAR THE MARTIAN

ATMOSPHERIC SURFACE J. L. Grenfell1, R. Lehmann

2, H. Rauer.

1,3

1 Institut für Planetenforschung,

Extrasolare Planeten und Atmosphären, Deutsches Zentrum für Luft- und Raumfahrt (DLR), Rutherford Str. 2,

12489 Berlin, Germany. 2Alfred-Wegener-Institut für Polar- und Meeresforschung, Telegrafenberg A43, 14473

Potsdam, Germany. 3Zentrum für Astronomie und Astrophysik,

Technische Universität Berlin (TUB), Hardenbergstr. 36, 10623 Berlin, Germany. [email protected].

We have applied a unique tool to the Martian

atmosphere which automatically identifies and

quantifies the photochemical catalytic cycles. This

tool is called the pathway analysis program (PAP)

and was originally developed for use in the Earth’s

atmosphere. We have applied PAP to determine the

cycles affecting carbon monoxide (CO) as well as

species associated with habitability i.e. ozone (O3),

nitrous oxide (N2O) and water (H2O) on Mars.

Results could reproduce some of the established

Martian cycles e.g. for CO which already appear in

the literature but we have also identified additional

cycles which, to our knowledge are new. Numerous

cycles appear to be rather complex mixes

consisting of many steps, catalysed by hydrogen

oxides, nitrogen oxides and oxygen species. Such

complexity seems to justify applying a tool such as

PAP which automates the analysis procedure and

removes subjective identification of cycles.




RADIATIVE TRANSFER SIMULATIONS FOR SPEX: AN IN-ORBIT SPECTROMETER C. Aas

1,2, D. M. Stam

1,2, E. Laan

3

1Aerospace Engineering, TU Delft, Kluyverweg 1, 2629 HS Delft, the Netherlands

2Space Research Organization Netherlands (SRON), Sorbonnelaan 2, 3584 CA Utrecht, the Netherlands

3Dutchspace, PO Box 32070, 2303 DB Leiden, the Netherlands.

[email protected]

Abstract

We present numerical simulations of the flux and

the state of polarization of light reflected by Mars,

illustrating the observing strategy of the

Spectropolarimeter for Planetary EXploration

(SPEX). For our calculations we use realistic, non-

spherical, Martian-analog dust particles, a range of

dust optical thicknesses. Our results show the

strength of polarimetry for the characterization of

optical properties of Martian dust particles, both on

the surface and in the atmosphere. Section 1: Martian dust particles

Despite the detailed observations of Mars obtained

during the last years, little is known about the

microphysical properties (size, shape, composition)

of the dust particles on the Martian surface and in

the atmosphere. Because microphysical properties

determine how particles scatter, absorb, and emit

radiation, knowledge of these properties and their

spatial and temporal distribution is crucial for

understanding the dust particles’ role in Mars’

radiation balance and hence in its weather and

climate, including the development of dust devils

and dust storms. Dust also influences the chemical

balance in the Martian atmosphere, e.g. reactions

involving methane (Farrell et al., 2006), and

transports condensates.

Section 2: Spectropolarimetry

Unpolarized light, such as incident sunlight, that is

scattered by particles in a planetary atmosphere or

that is reflected by the underlying surface will

generally get polarized. The degree of polarization

of the light that emerges from the top or bottom of

the planetary atmosphere, depends, like the flux, on

the illumination and viewing geometries, the optical

properties of the atmospheric particles, their spatial

distribution, the optical properties of the surface,

and the wavelength (see Hovenier et al. [2004] for

theory). The degree of polarization has been shown

to be more sensitive to particle microphysics (size,

shape, and composition) than the flux, and its

angular dependence is less sensitive to multiple

scattering than that of the flux (e.g. Hansen and

Travis, 1974). The Spectropolarimeter for Planetary

EXploration (SPEX) is being developed for

polarimetry of Martian dust and cloud particles (see

the contribution by Laan et al.) and is foreseen to be

placed on an orbiter.

Section 3: SPEX simulations

We use advanced radiative transfer algorithms (de

Haan et al, 1987), taking into account realistic

irregular Martian analogue palagonite dust particles,

to simulate the flux and state of polarization of light

reflected by the Martian atmosphere and surface, as

it will be observed by SPEX, across the wavelength

region of 0.35 to 0.8 microns. Our simulations, for a

range of dust optical thicknesses and illumination

and viewing geometries, clearly show the added

value of polarimetry for the retrieval of dust

particles.

References:

Banin, Han, Kan, Cicelsky (1997), J.Geophys. Res. 102,

13341-13356, 1997.

De Haan, Bosma, Hovenier (1987), Astron. Astrophys.

183, 371-391.

Farrell, Delory, Atreya (2006), Geophys. Res. Lett. 33,

CiteID L21203.

Foster, et al., (1998), JGR 103, 25839-25850.

Hansen and Travis (1974), Space Sci. Rev., 16, 527-610,

doi:10.1007/BF00168069.

Hovenier, van der Mee, Domke (2004), Kluwer,

Dordrecht (Springer, Berlin), 2004.

Figure 1. Images of non-spherical particles. The white bars indicate the scale. Left: Martian analogue palagonite particles

(sample 91-16 described by Banin et al. [1997], provided by T. Roush). The length of the bar (at the bottom) is 10 μm.

Right: CO2 ice crystals (bi-pyramids) made in a laboratory [Foster et al., 1998]. The length of the bar (at the left) is 3.0 μm.




RAMAN SPECTRAL CHARACTERIZATION OF A TERRESTRIAL SCENARIO WITH

IMPLICATIONS FOR MARS EXPLORATION: RIO TINTO, SPAIN. P. Sobron1, A. Sanz

1, J. Medina

1, F.

Sobron1, F. Rull

1 and C. J. Nielsen

2.

1Unidad Asociada Universidad de Valladolid-Centro de Astrobiología

(CAB-CSIC). Cristalografía y Mineralogía. Facultad de Ciencias. Paseo Prado de la Magdalena s/n. 47011

Valladolid, Spain. 2 Department of Chemistry, University of Oslo, Blindern, N-0372 Oslo, Norway.

[email protected]

Sulfate minerals may be the best means of

identifying potentially habitable sites on Mars and

thus studying them in terrestrial analog sites is of

critical importance to Astrobiology. Rio Tinto

(Huelva, Spain) is an example of an acidic, iron-rich

soil drainage site formed through mining. Sulfates

are also abundant on the surface of Mars. How they

formed is unknown, but they are thought to be

associated with aqueous processes. Terrestrial

studies of sulfate-rich sites are essential in order to

characterize the processes by which sulfates occur

on Earth and further to determine constraints on the

Martian mineralogy and surface processes. The Rio

Tinto environment is widely recognized as such an

analog site for potential sulfate-forming processes

on Mars.

Vibrational spectroscopy techniques such as

Raman and Infrared -which require little or no

sample preparation prior to spectra collection- are of

great importance for potential field analysis of

sulfate-rich sites. Both techniques have been

successfully applied to the analysis of synthetic

solutions that mimic the concentration of molecular

species in acid sulfate waters [1-3]. Jarosites, sulfate

minerals and efflorescent salts, both synthetic and

natural, have also been analyzed by using

spectroscopic techniques [4, 5].

In this work we report the X-ray diffraction and

Raman spectroscopy of aqueous solutions and

associated precipitates of Rio Tinto. Particularly,

Raman spectroscopy is a noninvasive and

nondestructive technique and both aqueous and

solid samples can be readily analyzed without any

preparation. Besides, the Raman spectrometer is an

instrument that can be used for identification of

biogenic and a-biogenic materials, different types of

ices, organic, and inorganic materials on planetary

surfaces. This is probably one of the reasons why

the compact Raman/LIBS instrument is regarded as

the highest priority instrument for mineral analysis

within the ExoMars mission roadmap.

Figure 1 show the Raman spectra of an aqueous

sample collected in Rio Tinto area. The species in

solution are readily identified through band-fitting

of the Raman spectra. Sulfate and bisulfate ions

concentration can be accurately computed. The

Raman spectrum of an efflorescent salt also

collected in the site is plotted in Figure 2. A

database match-search process is used in order to

identify the nature of the sample.

Figure 1. Raman spectrum of an aqueous sample of Rio

Tinto. Sulfate bands can be identified at 450, 625, 982 and

1105 cm-1. The band at 1650 cm-1 is characteristic of

water molecules bending vibration.

Figure 2. Raman spectrum of an efflorescent salt of Rio

Tinto, unambiguously identified as coquimbite

[Fe2(SO4)3·(H2O)9]

This is a first step in the development of

instrumental and analytical tools for the analysis of

Rio Tinto area with the objective of understanding

links between sulfate minerals and their

environment.

References: [1] Majzlan, J. and Myneni, S.C.B. (2005),

Environ. Sci. Technol. 39, 188-194. [2] Sobron, P. et al.

(2007), J. Raman Spectrosc., 38, 1127–1132. [3] Sobron,

P. et al. (2007), Spectrochim. Acta A,

DOI:10.1016/j.saa.2007.06.044. [4] Chio, C.H. et al.

(2005) Spectrochim. Acta A, 61, 2428-2433. [5] Frost.

R.L. et al. (2005) Spectrochim. Acta A, 62, 176-180

7000

7500

8000

8500

9000

9500

10000

10500

11000

11500

12000

300 500 700 900 1100 1300 1500 1700

Raman shift/cm-1

Inte

ns

ity

/a.u

.

0

2000

4000

6000

8000

10000

12000

14000

16000

100 200 300 400 500 600 700 800 900 1000 1100 1200 1300

Raman shift/cm-1

Inte

ns

ity

/a.u

.




RESULTS FROM THE 7

TH INTERNATIONAL CONFERENCE ON MARS C. J. Budney

1, D. Beaty

1, M.

A. Meyer2, R. W. Zurek

1 and the attendees of the 7

th International Conference on Mars.

1Mars Exploration

Program, Jet Propulsion Laboratory, California Institute of Technology, 4800 Oak Grove Drive, Pasadena, CA

91109, USA. 2 Mars Exploration Program, Planetary Science Division, Space Science Directorate, National

Aeronautics and Space Administration, Washington, DC 20546-0001, USA. [email protected]

The 7th International Conference on Mars, held

July 9-13, 2007 in Pasadena, California, was

organized into nine theme-oriented oral sessions

(listed below). Each of the oral sessions was

configured with a set of papers of relevance to the

theme of the session, along with some introductory

comments by the session chairs to frame the session,

and a concluding discussion session that was

moderated by the session chairs. The following

master discussion prompts were used for each of the

concluding discussion sessions:

• What do we know?

• What do we need to learn next?

• Are we doing the right things to find the

answers?

For each of the nine discussion sessions, the

comments from the audience were documented and

the session chairs pulled together summaries. We

present the results of those exchanges.

Technical Sessions and Key Findings

The distribution and context of water-related

minerals on Mars

The key hypothesis framing the session was that

phyllosilicates were dominantly formed during

Mars’ earliest period (Noachian), followed by

sulfates in the late Noachian to Hesperian. There has

been little evidence for formation of water-related

minerals since the Hesperian. This is despite the

abundance of morphologic features thought to be

water related (e.g. young valley networks and

outflow channels, gullies, volcano-ice interactions,

and ice-related features).

Geology of the martian surface: Lithologic

variation, composition, and structure

We now know that the crust of Mars is

dominated by basaltic volcanism, and that aqueous

alteration (chemical weathering) occurred early in

Mars history but weathering in more recent times

has been mostly physical.

Water through Mars’s geologic history

The history of water on Mars appears to span,

literally, the whole of geologic history. It pervades

the magmas, has formed evaporates and clastic

sediments as well as nearly pure salt and silica

concentrations, yet has only slightly reacted with the

widespread contemporaneous fine-grained soils.

Volatiles and interior evolution

Several models now exist for early accretion and

differentiation, production and longevity of a core

dynamo, production of basaltic crust, and

partitioning of initial and secondary volatile

contents into the atmosphere (with their effects on

climate) and loss to space.

The Martian climate and atmosphere:

variations in time and space

Recent evidence shows that the Mars lower and

upper atmospheres are coupled thermally,

dynamically, and chemically. GCM modeling

frameworks are evolving to properly capture the

“whole atmosphere” coupling processes that are

required to explain these observed variations in the

Martian upper atmosphere.

Modern Mars: Weather, atmospheric

chemistry, geologic processes, and water cycle

It is apparent that the martian dust cycle is highly

variable. Recent observations of the water vapor

content of the martian atmosphere suggest it may be

drier than previously assumed. Trace gases remain a

significant outstanding question.

The north and south polar layered-deposits,

circumpolar regions, and changes with time

As the planet’s principal cold traps, the Martian

polar regions have accumulated extensive mantles

of ice and dust that cover individual areas >106 km

2

and total as much as 3–4 km thick. From the small

number of superposed craters found on their surface,

these layered deposits are thought to be

comparatively young. Radar sounding investigations

have provided a first look at the basal topography

and internal layered structure of both caps.

Mars astrobiology and upcoming Missions

There is now strong evidence that habitable

environments might have existed on Mars at least

intermittently in the distant past.

The Phoenix mission and the MSL site selection

process were discussed. Needed for future missions

are information for site selection, access to more

sites, and capable in situ science instruments to

identify the best samples.

Martian stratigraphy and sedimentology:

Reading the sedimentary record

The recent explosion of high-resolution data both

from orbit and from the ground allows probing the

third dimension of martian crust, thereby permitting

the stratigraphic architecture of sedimentary

deposits to be better constrained and interpreted.

Recent results include evidence of eolian

crossbedding at Victoria crater and suggestions for a

volcaniclastic origin of sediments at Home Plate.

References: Beaty, D.W., Budney, C.J., and McCleese,

D.J. (2007). Session Summaries, 7th International

Conference on Mars, July 9-13, 2007. Unpublished white

paper, 22 p, posted August 2007 by the Mars Exploration

Program Analysis Group (MEPAG) at

http://mepag.jpl.nasa.gov/reports/index.html.




THE RHEOLOGY OF YOUNG LAVA FLOWS ON ARSIA, PAVONIS AND ASCRAEUS MONS,

MARS H. Hiesinger1, D. Reiss

1, S. Dude

1, C. Ohm

1, G. Neukum

2.

1Institut für Planetologie, Westfälische

Wilhelms-Universität, Wilhelm-Klemm-Str. 10, 48149 Münster, Germany. 2Freie Universität Berlin,

Malteserstr. 74-100, 12249 Berlin. [email protected]

Introduction: The Tharsis Montes, Arsia Mons,

Pavonis Mons, and Ascraeus Mons, are large

volcanic constructs that are located on the Tharsis

bulge and are the locations of some of the youngest

volcanic deposits on Mars [1,2]. From previous

studies it is known that in principle, the dimensions

of flows reflect rheological properties such as yield

strength, effusion rates and viscosity. We expand on

our previous study of the rheological properties of

lava flows on Ascraeus Mons [3] in order to

investigate possible similarities and differences

between lava flows on the the Tharsis Montes.

Data: We used new high-resolution images

obtained by the High Resolution Stereo Camera

(HRSC) on board ESA’s Mars Express spacecraft in

combination with Mars Orbiter Laser Altimeter

(MOLA) data to constrain these rheological

properties. We made use of several HRSC orbits

with spatial resolutions of about 10-20 m/pixel in

order to measure the length and width of the studied

lava flows. Individual MOLA profiles were used to

determine the height of the lava flows and gridded

MOLA topography was used to measure the slope

on which these flows occur. Compared to earlier

studies, HRSC and MOLA data allowed us to map a

larger number of late-stage lava flows and to

measure their dimensions, as well as their

morphological characteristics in greater detail.

Method: We modeled the investigated lava

flows as a Bingham plastic controlled by two

parameters, the yield strength and the plastic

viscosity [e.g., 4]. The yield strength of lava flows

(Pa) is related to the flow dimensions by the

following equations [e.g., 5]

= g sin h (1)

= g h2/w (2)

= g sin2

2wl (3)

= g sin2

(w-wc) (4)

where is the density (kg m-3

), g is the

gravitational acceleration (m s-2

), is the slope

angle (degree), h is the flow height (m), w is the

flow width (m), wl is the total levee width (m), and

wc is defined as the width of a leveed channel (m).

The effusion rates Q (m3/s) can then be

calculated as

Q = Gz x w/h (5)

where Gz is the dimensionless Graetz number,

is the thermal diffusivity (m2

s-1

), x is the flow

length (m), and w and h are defined as above [e.g.,

4; 6].

The viscosities (Pa-s) were calculated using

the relationship given for example by [7,8].

h = (Q / g)1/4

(7)

Jeffrey's equation also relates the viscosity of a

flow to its effusion rate and its dimensions [e.g., 9-

11].

= ( g h3 w sin )/nQ (8)

In this equation n is a constant equal to 3 for

broad flows and 4 for narrow flows.

Results: Our estimates of the yield strengths for

flows on Ascraeus Mons range from ~2.0 x 102 Pa

to ~1.3 x 105 Pa, with an average of ~2.1 x 10

4 Pa.

These values are in good agreement with estimates

for terrestrial basaltic lava flows. The effusion rates

are on average ~185 m3s

-1, ranging from ~23 m

3 s

-1

to ~404 m3

s-1

. While these results are higher than

earlier findings that indicate effusion rates of 18-60

m3

s-1

, with an average of 35 m3

s-1

, they are similar

to terrestrial effusion rates of Kilauea and Mauna

Loa and other Martian volcanoes. Viscosities were

estimated to be on average ~4.1 x 106 Pa-s, ranging

from ~1.8 x 104 Pa-s to ~4.2 x 10

7 Pa-s. On the basis

of our effusion rates and the flow dimensions, we

calculated that the time necessary to emplace the

flows is on average ~26 days.

Conclusions: On the basis of our investigation

we find our results for the yield strength, effusion

rate, eruption duration, and viscosity to be in good

agreement with previously published results for

Martian and terrestrial flows. The strength of our

study is that we investigated a much larger number

of flows than in previous studies [e.g., 5,6,8,12-14]

and therefore provides a more complete foundation

of our understanding of Martian lava rheologies. In

a next step, we plan to put our results into a

temporal context in order to study whether the

rheology of the Tharsis Montes lava flows

systematically varied with time.

References: [1] Scott and Tanaka, I-1802-A, 1986;

[2] Neukum et al., Nature, 432, 971-979, 2004; [3]

Hiesinger, et al., J. Geophys. Res., 112,

10.1029/2006JE002717, 2007; [4] Wilson and Head,

Nature, 302, 663-669, 1983; [5] Moore et al., Proc.

Lunar Planet. Sci. Conf., 9, 3351-3378, 1978; [6]

Zimbelman, Proc. Lunar Planet. Sci. Conf., 16, 157-162,

1985; [7] Fink and Griffiths, J. Fluid Mech., 221, 485-

500, 1990; [8] Warner and Gregg, J. Geophys. Res., 108,

doi:10.1029/2002JE001969, 2003; [9] Nichols, J. of

Geology, 47, 290-302, 1939; [10] Gregg and Fink, J.

Geophys. Res., 101, 16891-16900, 1996; [11] Gregg and

Zimbelman, Env. Effects on Volcanic Eruptions: From

Deep Oceans to Deep Space, (Zimbelman, Gregg, eds.)

Kluwer Academic/Plenum Publishers, New York, 75-

112, 2000; [12] Keszthelyi, J. Geophys. Res., 100, 20411-

20420, 1995; [13] Sakimoto et al., J. Geophys. Res., 102,

6597-6613, 1997; [14] Cattermole, Proc. Lunar Planet.

Sci. Conf., 17, 553-560, 1987.




RIFTING ON MARS: STRUCTURAL GEOLOGY AND GEOPHYSICS P. Kronberg

1, E. Hauber

2,

M. Grott2, J.-M. Ilger

1.

1Institut für Geologie, TU Clausthal, Leibnizstr. 10, 38678 Clausthal-Zellerfeld,

Germany. 2Institut für Planetenforschung, DLR, Rutherfordstr. 2, 12489 Berlin, Germany. peter.kronberg@tu-

clausthal.de

Introduction: Rift-like extensional features on

Mars have been identified on the basis of Mariner 9

and Viking Orbiter images [e.g., 1-3], but little

detailed work has been done before accurate

topographic data were obtained by the MOLA laser

altimeter. Over the last years, we have investigated

several rift-like structures (Fig. 1) with respect to

their topography and structural geology and found

that they can directly be compared to terrestrial

analoga like the Kenya Rift [4-7].

Figure 1. A comparison of several rift-like structural

features on Mars. Clockwise from top left: Tempe Fossae

Rift [4], Thaumasia Double Rift [6], unnamed rifts in

western Tempe Terra (“X”-Rift in Fig. 2), Acheron

Fossae Rift [7].

Summary and Discussion: Crater counts

revealed that at least two of the rifts formed in Late

Noachian- to Early Hesperian time. Using

measurements of rift flank uplift, we determined the

thickness of the elastic lithosphere at that time to be

in the order of ~10-15 km. The corresponding heat

flux would range between ~55 and 80 mW m-2

. In

contrast to most of the long and narrow linear

graben sets (e.g., Mareotis, Memnonia, and Icaria

Fossae; Fig. 2), which dominate the tectonics of

Tharsis and, indeed, almost the entire western

hemisphere of Mars, the rifts are not consistently

oriented radial to Tharsis center(s). It seems possible

that plume tectonics might not be the only plausible

scenario that could account for rift formation on

Mars. The combination of regional stresses, local

magmatism, and high elevations that generate

stresses associated with horizontal gradients of the

gravitational potential energy [9] might account for

passive rifting [10]. We will present a synthesis of

our work on Martian rifts, and will present the

results for the first time in the context of all

observed rifts.

References: [1] Masson, P. (1980) Moon & Planet. 22,

211-219. [2] Tanaka, K. et al. (1991) JGR 96, 15,617-

15,633. [3] Banerdt, B. et al. (1992) Mars, by H. Kieffer

et al. (Eds.), pp. 249-297, Uni. Ariz. Press. [4] Hauber, E.

and Kronberg, P. (2001) JGR 106, 20,587-20602.

[5] Hauber, E. and Kronberg, P. (2005) JGR 110,

DOI:10.1029/2005JE002407. [6] Grott, M. et al. (2005)

GRL 32, DOI:10.1029/2005GL023894. [7] Kronberg, P.

et al. (2007) JGR 112, DOI: 10.1029/2006JE002780.

[8] Dimitrova, L. et al. (2006) GRL 33, DOI :

10.1029/2005GL025307. [9] Grott, M. et al. (2007) JGR

112, DOI: 10.1029/2006JE002800.

Figure 2. Rifts (red) and other major extensional features

in Tharsis. Note that the orientation of some rifts (e.g., the

Thaumasia Double Rift [6] or Acheron Fossae [7]) is not

consistently radial to Tharsis. The stereographic

projection is centered near Pavonis Mons to highlight

Tharsis-centered axial symmetries (large volcanic edifices

marked by grey circles, Valles Marineris and the large

Thaumasia graben also marked as grey areas).




SCALLOPED TERRAIN SOUTH OF THE HELLAS BASIN: RESULTS FROM HRSC, MOC AND

MOLA. M. Zanetti1, H. Hiesinger

1, D. Reiss

1, E. Hauber

2, G. Neukum

3.

1Institute für Planetologie, Westfälische

Wilhelms-Universität Münster, Wilhelm-Klemm-Str. 10, 48149 Münster, Germany. 2Institute of Planetary

Research, German Aerospace Center (DLR), Rutherfordstr. 2, 12489 Berlin, Germany. 3Freie Universität

Berlin, Maltesertr. 74-100, 12249 Berlin. [email protected]

Introduction: We performed an investigation of

proposed sublimation landforms in the Amphitrites

and Peneus Paterae region of Malea Planum, on the

southern rim of the Hellas Basin. A latitude

dependent, several meters thick, surface mantle

presumably composed of dust and water ice is found

throughout this region, and is thought to be related

to obliquity-driven ice activity as recently as 2.1-0.4

Myr [1]. Scallop-shaped asymmetrical depressions

(Fig. 1), a type of dissected mantle terrain described

by [2], have formed in this mantle deposit, and have

been mapped in detail in order to characterize these

unique features. Milliken [2] and Plescia [3] have

proposed that they were formed by interstitial ice

sublimating from the mantle material. Scallops are

also found in northern latitudes, in Utopia Planitia

[4,5]. Morgenstern et al. [4] propose an insolation-

driven model of scallop formation in Utopia. We

studied these features in an area that extends from

50°E to 70°E longitude and 50°S to 70°S latitude in

order to determine the recent geologic history of the

region and to study if the regional climate of the

Hellas basin has had an impact on the formation of

these scallops.

Data and Methods: We used high-resolution

images obtained by the High Resolution Stereo

Camera (HRSC) on board ESA’s Mars Express

spacecraft in combination with Mars Orbiter

Camera-Narrow Angle (MOC-NA) and THEMIS-

VIS data to map scalloped depressions in the study

area. Mars Orbiter Laser Altimeter (MOLA) data

were used to obtain depths of individual scallops.

Results: The study area is extensively scalloped

through the entire longitude range of 50°E-70°E, but

only between 52°S and 59°S latitude. Here the

mantle appears smooth with very few small craters.

The area with scallops contours the slope of the

southern rim of the Hellas Basin, as they grade from

small isolated scallops in the northern latitudes and

lower elevations to large coalesced bands in the

higher regions of the Hellas rim and on the caldera

rims of Amphitrites and Peneus paterae. The

scallops abruptly terminate near 59°S after a small

(~1°) transition zone from a thick smooth top-most

mantling layer, to a degraded lower level of

mantling material. Scallops that are large enough to

be measured with MOLA data show a strong slope

asymmetry, with steeper south-facing slopes and

gentler north-facing slopes. This shape can also be

inferred for smaller scallops from imaging data.

Depths vary with the area of the depression but

typically do not exceed 30 m. We performed a

systematic southern hemisphere survey of HRSC

images centered at ~55°S latitude to locate scallop

features as a follow-up investigation to that done by

Milliken [2] using MOC images. This was done to

identify any longitudinal dependence in the

distribution of scallops or effects of different

geologic units or topography. Our survey revealed

that scallops are found almost exclusively in the

area along the southern slope of the Hellas Basin

(30°E - 110°E), with a few isolated scallops

occurring in the southern latitudes near Argyre

Basin (300°E - 335°E), hence pointing toward a

strong effect of basin morphology on scallop

formation.

Conclusions: The presence of the studied

scallops in the southern hemisphere appears to be

closely linked to the Hellas Basin. We interpret the

results of our observations as evidence of the

influence of the Hellas Basin on the development of

scalloped depressions. We propose that weather

conditions, solar radiation insolation, subsurface

water-ice availability and seasonal variation in

temperature have allowed scallops to degrade the

mantle in this region.

Figure 1. Scallop features in mantle material. Scallops

appear to coalesce as they grow larger. HRSC image

h2448 (centered at -57.60 lat, 65.60 long).

References: [1] Head et al. (2003), Nature, 426, 797-802

[2] Milliken, R.E. and Mustard J.F. (2003), Sixth

International Conference on Mars, Abs. #3240.

[3] Plescia J.B (2003) LPSC XXXIV, Abs. #1478.

[4] Morgenstern et al. (2007) JGR 112, E6. [5] Lefort et

al. (2007), LPSC XXXVIII Abs. #1796.




SEARCH FOR CARBONATES ON MARS WITH THE OMEGA / MARS EXPRESS DATA. D. Jouglet,

F. Poulet, J. P. Bibring, Y. Langevin, B. Gondet. Institut d’Astrophysique Spatiale, Université Paris-Sud, Orsay,

France. Contact : [email protected]

Introduction: Key minerals for the study of the

past of Mars, such as phyllosilicates, sulfates or

hydroxides, have been discovered by the Mars

Express / OMEGA experiment [e.g. 1,2]. This

abstract focuses on the search for carbonates from

the OMEGA dataset. Previous orbital missions,

like IRS [3] or TES [4], failed in finding large

amounts of carbonates on the surface of Mars.

The presence or the lack of carbonates on the

Martian surface is very important to 1) better

understand the climatic and geological past of the

planet, since carbonates easily form in aqueous

media [e.g. 5]; 2) get new elements about the

evolution of a primitive thicker CO2 atmosphere,

since dissolved carbon dioxide precipitates in

carbonate minerals [e.g. 6].

Detection tool: This study is based on the

detection of the strong 3.4 m and 3.9 m

absorption bands present in carbonate reflectance

spectra [7]. Since the 3.9 m area is subject to

thermal contrast reduction [8] and to atmospheric

absorptions influence, the detection focuses on the

3.4 m band depth. This band is inside a broad 3

m hydration band [9] and therefore requires a

continuum removal, as illustrated by figure 1 [7].

Application on the OMEGA dataset: This

method is tested first on old terrains where clays

and sulfates were detected: Mawrth Vallis [2],

Terra Meridiani [10] or Nili Fossae [11]. These

areas reveal no spectral features of carbonates.

Then the detection tool is applied on every

OMEGA spectra, through an automatic and quick

program. Spectra are recorded if their 3.4 m band

depth is greater than 1%. Since water ice may

influence the 3.4 m spectral area, spectra

exhibiting water ice features (a characteristic 1.5

m absorption band) are removed from the study

[12]. Spectra with low signal to noise ratio are also

excluded through the value of the flux received by

OMEGA. In order to avoid isolated spurious pixels,

detections are recorded only if at least one other

detection is done in its neighborhood. This method

is described in [7].

140 million albedo spectra are tested by this tool

(orbit 0 to 1989), covering about 80% of the

Martian surface between 80°N and 80°S. A few

candidate spectra have been recorded and require a

visual diagnosis based on the analysis of spectral

signatures, spatial distribution and comparison with

other observations on the same area. This diagnosis

concludes that no obvious carbonate spectrum is

present in the OMEGA dataset, suggesting

carbonates are not widely spread on the surface of

Mars. A few areas are discussed because they

exhibit a small 3.4 m feature and a geographical

clustering but other important carbonate features

are missing. An example is given in figure 2.

Future work: In the future the threshold for the

3.4 m band depth detection will be adapted to the

signal to noise ratio. The 3.9 m band depth will be

also recorded as additional information for

carbonate detection.

Figure 1. Application of the 3.4 m carbonate detection

tool on a laboratory spectrum obtained during the ground

calibration of OMEGA. The blue dashed lines borders

the areas used for polynomial fitting, the yellow dashed

lines the area used for band depth calculation once the

continuum is removed. In this example the Band Depth is

2.6%. Black: radiance spectrum of a 80% palagonite / 20

% calcite mixture divided by a reference radiance

spectrum (MgO). Red: polynomial interpolation to

simulate the spectrum continuum at 3.4 m.

Figure 2. Black curve: OMEGA reflectance spectrum at

132.04E, 66.75N (orbit 1059), red curve: spectrum

continuum, green curve: laboratory carbonate spectrum

of fig.1. This pixel was recorded by the carbonate

detection tool. Visual diagnosis reveals spatial clustering

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.




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.




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

Sonnensystemforschung, Max-Planck-Straße 2, 37191 Katlenburg-Lindau, Germany. [email protected]

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.




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.




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.




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

Kong. [email protected]

The search of signs of past life on the red planet

continues to be the focus for several major future

missions where analyses of the Martian surface and

subsurface soil composition could provide further

insights towards the search and the better

understanding of Mars’ morphology. Efficient and

reliable tools are necessary to support these

exploration activities and sophisticated experiments

to be carried out effectively. The payload

constraints demand microscopic yet multi-functional

tool sets to adapt to a wide range of tasks.

A team from Hong Kong has been appointed by

the United Kingdom to design and manufacture the

sampling tools on board of the Beagle 2 Lander.

Based on past experience, the team has designed

small and lightweight soil preparation system

(weighing only 230g and grinding to size of 1mm)

for the Russian Phobos-Grunt Mission planned to be

launched in 2009. The device is developed to

function under an environment of practically no

gravity. These recent developments can be adapted

for future applications on Mars.

An overview on different sampling methods and

sample preparation techniques will be presented

here based on experiences acquired during past

missions to our solar system’s planets. Using the

ExoMars vehicle as a baseline, an advanced

sampling concept of integrated downhole

hammering sampler will be presented. The

proposed sampler is designed for fine sand

sampling, as well as for rock coring and gripping.

Implemented onto a long drill, this system is able to

function several feet below the Martian surface.

The correlation between the scientific objectives of

future Mars missions to the design of the proposed

novel sampling strategy will be illustrated. The

paper will conclude with an outline of the

prototyping efforts and the other future development

of the tool.

References: .

Figure 1. Drill End of Proposed Integrated Sampler




SOLAR FORCING AND MARTIAN PLANETARY ION OUTFLOW – NEW MARS

EXPRESS FINDINGS R. Lundin1, S. Barabash1, M. Holmström1, Hans Nilsson1, M. Yamauchi. B1.

J.-A. Sauvaud2, A. Fedorov2.

1Swedish Institute of Space Physics (IRF), Kiruna, Sweden. 2Centre

d’Etude Spatiale des Rayonnements, BP-4346, F-31028 Toulouse, France. [email protected]

The solar wind and the solar XUV/EUV radiation

constitute a highly variable forcing of the upper

atmosphere of Mars. Solar forcing leads to heating-,

ionization, energization, and loss of planetary

atmospheric atoms and molecules. The extent to

which solar forcing governs planetary ion escape is

still debated. The solar forcing variability leads in

the Earth's case to outflow variability by up to three

orders of magnitude for O+.

New energy settings, implemented in May 2007,

enables us to analyze cold ionospheric ions by the

ASPERA-3 Ion Mass Analyzer (IMA) in greater

detail. After some four months of data taking a

revised picture of the planetary ion escape emerges.

Low energy ionospheric ions expanding into the tail

with velocities in the 5-20 km/s range dominates the

outflow (Fig. 1). The expansion/outflow is comet-

like, the low-energy ions forming a mantle of

variable thickness connecting to the dayside/flank

high-altitude ionosphere. The outer bound of this

comet-like mantle lies well inside the induced

magnetosphere boundary (IMB).

We present results from a statistical study based

on data from 42 pre-selected orbits before and 30

orbits after the change of energy settings. A

preliminary analysis indicates that the escape rate

from Mars is substantially higher than those

previously reported from MEX. The variability of

the low-energy planetary ion outflow is compared

with solar forcing conditions. Assuming that the

XUV/EUV flux and the solar wind dynamic

pressure are the main drivers for solar forcing we

find that solar forcing variability leads to outflow

rates varying by up to three orders of magnitude.

The outflow varies substantially from orbit to orbit,

even during stable solar wind conditions. This

implies a highly variable solar forcing by primarily

solar XUV/EUV, alternatively a non-linear response

to solar forcing variability.

Figure 1. ASPERA-3 ion data from two orbit illustrating the expansion/outflow of low-energy ionospheric ions (O+, O2

+,

CO2+) into the Martian tail. Heavy purple line along the orbit marks the extension of ionospheric ions into the tail.

O+, O2+, CO2

+

O+, O2+, CO2

+

H+, He++

H+, He++

IMB

IMB

IMB IMB

IMB IMB

MEX orbitMEX orbit




SOUNDING OF THE MARTIAN EXOSPHERE WITH SPICAM ON MARS EXPRESS

J-Y. Chaufray1, F. Leblanc

2., J-L. Bertaux

1, E. Quémerais

1

1Service d’Aéronomie du CNRS/IPSL, Reduit de Verrieres BP3 Route des Gatines 91371 Verrieres-le-Buisson,

FRANCE. 2Temporarily at Osservatorio astronomico di Trieste, Via Tepolo 11 34131 Trieste, ITALY

[email protected]

Section 1: Atomic hydrogen and oxygen are

important tracers of the global behaviour and

evolutionary processes of the water. A series of

exospheric observations performed in 2005 with the

ultraviolet spectrometer SPICAM on board Mars

Express are studied. Two types of observation are

analyzed : observations of the Lyman- line in the

upper atmosphere and observations of the O I 130.4

nm triplet in the lower exosphere. We will present

the data processing and the methodology, based on a

model of thermospheric and exospheric profiles

coupled to a radiative transfer model used to

analysis these optically thick emissions. The oxygen

and hydrogen densities deduced from these

observations will be presented and compared to the

Mariner’s results. The analysis of the Lyman-

emission above the exobase suggests the presence of

a hot component of the exospheric Martian

hydrogen population. However this conclusion will

need to be confirmed by new measurements,

development of sophisticated radiative transfer

approaches and by simultaneous measurements with

other instruments. The analysis of the oxygen

component remains limited by the range of altitude

where SPICAM UVS can observe the oxygen

emissions. However, as for the case of the Lyman

alpha, a first analysis suggests that a single

component fit to the observation cannot reproduce

the observed measurements without introducing

exospheric temperature much larger than previously

measured by previous space missions or through the

analysis of other emission lines.

.




THE SOURCES OF METHANE ON MARS V. Formisano

1, A. Geminale

1.

1IFSI-INAF, Via Fosso del

Cavaliere 100 , Rome, Italy. [email protected]

ABSTRACT

The presence of methane in the atmosphere of

Mars is well established ( Formisano et al 2004).

The previous study has shown the average

properties of the methane mixing ratio ( Geminale et

al 2007) . On average methane has a mixing ratio of

15 +- 5 ppbv . There are , however, several orbits in

which the mixing ratio is substantially higher. On

the basis of the GCM study of F.Forget, these orbits

may help identify the locations where the possible

sources of methane are located. Out of 20 000

measurements, organized in more than 300 orbits,

we have searched for the orbit averages having the

methane line depth larger than 2 % of the

continuum, while on average is of the order of 0.5 %

of the continuum. The selected orbits are 25 : they

group naturally in 4 groups, which are then studied

separately.

The first group of 8 orbits are strongly confined

in time and space: they are located above 60o

Northern latitude during the advanced spring. In

other words these orbits point to the sublimating

northern polar cap as one source of methane.

The second group of 9 orbits are confined in

space ( not in time) over Arabia Terra broadly

spiking. They point, therefore to the possible water

ice source identified by Odissey Neutron monitor

experiment and the Gamma Ray Spectrometer

experiment.

The third and fourth group are made of 4 orbits

each and they point , but with less accuracy than the

first two groups, to the Vallis Marineris region, and

to the Elisium planum region ( where Marsis has

found extended regions with underground water ice,

see the report to EGU …..).

These findings link the source of methane and

the water ice deposits on Mars, where methane

could be included as clathrate hydrates.

We propose for discussion two possible source

mechanisms, one biological and another abiotic.

It has been demonstrated on Earth that Archea

can live in ices and produce methane.

The abiotic possibility comes from possible

analogy of Mars environment with comets. Comets

do contain methane. Methane on comets is thought

to be formed by bombardment with energetic

particles of water ice containing CO2 molecules.

The annual sublimation and condensation of the

polar cap poses a number of constraints to the two

mechanisms, which , if well modeled, can

eventually bring us to selection of one of them.

References:

Formisano, V., Atreya, S., Encrenaz, T., Ignatiev,

N., Giuranna, M., 2004. Detection of methane in the

atmosphere of Mars. Science 306, 1758-1761.

Geminale, A., Formisano V., Giuranna M :

Methane in Martian atmosphere : average spatial,

diurnal , and seasonal behaviour.. This conference .

(2007).

Forget F, Haberle B., Montmessin M. : Spatial

variation of methane and other trace gases detected

on Mars:interpretation with a GCM. Lunar and

Planetary Science conf. XXXVI ( 2005).

Figure 1 :- a) Left. The average spectrum measured by PFS showing an average quantity of methane .

b) Right. A special orbit average spectrum showing rather high mixing ratio of methane.




THERMAL STRUCTURE OF MARTIAN SOIL AND THE MEASURABILIY OF THE PLANETARY

HEAT FLOW M. Grott1, J. Helbert

1 and R. Nadalini

2.

1Institute of Planetary Research, German Aerospace

Center (DLR), Berlin, Germany, 2Active Space Technologies GmbH, Berlin, Germany. [email protected]

Introduction: The planetary heat flow is one of

the key quantities characterizing a planets' thermal

state and significantly influences tectonic,

magmatic and geological processes on the surface.

Furthermore, it is one of the few constraints we

have for thermal evolution models and is closely

connected to the concentration and distribution of

radioactive isotopes in the planetary interior. Apart

from the Earth, in-situ heat flow measurements

have only been performed on the Moon and

indirect methods had to be relied on to estimate the

Martian planetary heat flow.

Upcoming in-situ geophysical experiments [1]

will measure the Martian surface heat to learn

about the thermal state of the planet. However, the

near surface thermal gradient and therefore the near

surface heat flow are determined by a number of

processes, most of which are exogenic. At shallow

depth, soil temperatures are driven by insolation,

and diurnal as well as seasonal cycles have a

dominant influence. Furthermore, climatic

variations like ice ages can have a significant

influence on the surface heat flow. Only part of the

temperature gradient near the surface is determined

by the planetary heat flow and a measurement of

this quantity can potentially pose severe problems.

We have investigated how the near surface heat

flow is influenced by exogenic processes like the

diurnal and seasonal temperature cycles and how

meaningful measurements of the planetary heat

flow can be obtained [1]. Furthermore, we have

estimated how long-term climate and interannual

temperature variations disturb the surface heat flow

and assess the resulting errors.

Model: We have investigated the thermal

structure of dry Martian soil assuming that it is

determined by insolation and planetary heat flow.

Soil temperatures are then determined by the one-

dimensional heat conduction equation and we

consider a model with depth dependent thermal

conductivity (cp. Fig.1) and density:

4

3

2

1

)(

)(

)()(

cz

czz

cz

czkzk

z

Tzk

zt

Tcz p

+

+=

+

+=

=

The surface boundary condition is given in

terms of temperature, which is taken from the

NASA/MSFC Global Reference Atmospheric

model [2]. The parameters used in this study are

summarized in Table 1.

Figure 1. Thermal conductivity k as a function of

depth for the two end-member cases considered. The

surface conductivity is limited by the conductivity of

CO2 and taken to be 0.01 W m-1

K-1

. At a depth of 10 m,

the models reach 95 % of the final conductivity k of

0.02 W m-1

K-1

(solid line) and 0.1 W m-1

K-1

(dashed

line), respectively.

Table 1. Parameters used in this study

Conclusions: The influence of seasonal

temperature changes was found to be efficiently

removed if measurements are extended over the

period of at least a full Martian year. Interannual

variability due to, e.g., eolian driven surface albedo

changes typically alters the heat flow by less than

15%, although errors may be larger if the soil's

thermal conductivity and the albedo variations are

both large. Heat flow perturbations caused by long-

term climatic changes are found to stay below 15%.

In order to also determine the soil's thermal

conductivity with an accuracy of 20 % or better, a

direct conductivity measurement is required. We

conclude that a measurement of the Martian

planetary heat flow is possible with an accuracy of

30 % or better if measurements are extended over

the period of at least a full Martian year and

thermal conductivity is directly measured.

Temperature sensors should have a precision of 0.1

K and measurements should be conducted up to a

depth of 3-5 m.

References: [1] Spohn, T., et al. (2001), PSS, 49, 1571-

1577. [2] Grott, M., et al. (2007), JGR, in press. [3]

Justus, C.G., et al. (2002), Adv. Space Res., 29, 193-202.

Variable Physical Meaning Value Units

k0 Surface thermal conduct. 0.01 W m-1 K-1

0 Surface density 1000 kg m-3

Asymptotic density 1750 kg m-3

cp Soil specific heat 600 J kg-1 K-1

F Planetary heat flow 0.02 W m-2




STUDY OF AEROSOL PHASE FUNCTION WITH PFS DATA, G. Rinaldi

1, V. Formisano

1

1IFSI-INAF,

Rome, Italy [email protected]

The aim of our work is to analyze the aerosol phase

function inside the 2.7 μm band by means of a set

of PFS spectra taken in nadir and spot pointing

mode and evaluate the radiative properties of the

aerosol atmospheric particles. Planetary Fourier

Spectrometer (PFS), a payload instrument of the

European Mars Express Mission, is a Fourier

interferometer with a spectral range of 250-8000 cm-

1 and with a spectral resolution

of 1.3 cm-1

(Formisano et al. 2005).

The understanding of aerosol properties is necessary

to constrain the aerosol influence on heating and

cooling of atmosphere.

This work consists in two parts: in the first part the

PFS data set used consist of spectra acquired in

nadir mode, in two different seasons (Ls=330o and

Ls=40o) and along different orbits with different

geometries (i.e. incident and emission angle).

Spectra with the same phase angle (10o

bins) and in

the same season have been groupped and averaged.

In this way we are able to obtain the radiance factor

at 3700 and 3600 cm-1

in a wide phase angle range

(from 10o to 120

o).

These aerosol phase functions have been retrieved

under the following assumptions: properties are

uniform along one orbit; no local dust storm is

present.

The properties of the mean spectra have been

modeled by a radiative code implementing line by

line calculations of gaseous and aerosol opacities in

LTE conditions (N.I. Ignatiev et al., 2005) to study

particle size distribution (reff),albedo and shape of

airborne particles.

In the second part we compare these phase functions

with the one obtained from orbits taken in spot

pointing mode to have the phase functions at the

same wavenumber (Ls= 58o). From our preliminary

studies about phase functions we infer that there are

two types of dust particles, in agreement with

Clancy and Wolff, 2003.

These dust particles indicate agreement of EPF-

derived dust single scattering albedos (0.92-0.94)

with results from Viking lander and

TES observations (Pollack et al. 1995 and Clancy

and Wolff, 2003). The two set of phase functions

measured from PFS show also how the solid

components of the Martian atmosphere are not yet

well understood.

References:

Clancy, R. T., Wolff, M. J., Christensen, P. R.

Mars aerosol studies with the MGS TES emission phase

function observations: Optical depths, particle sizes, and

ice cloud types versus latitude and solar longitude

Journal of Geophysical Research, Volume 108, Issue E9,

pp. 2-1, CiteID 5098, DOI 10.1029/2003JE002058

Formisano, V. et al., The Planetary Fourier Spectrometer

onboard the European Mars Express mission. Planet.

Space Sci. , Vol 53, p. 963, ( 2005).

Ignatiev, N. I., Grassi, D., Zasova, L. V.,

Planetary Fourier spectrometer data analysis: Fast

radiative transfer models, Planetary and Space Science,

Volume 53, Issue 10,p. 1035-1042, 2005

Pollack, J.B. and Ockert-Bell, M.~E. and Shepard, M.K.,

Viking Lander image analysis of Martian atmospheric

dust, Journal of Geophysical Research (ISSN 0148-0227),

Vol. 100, no. E3, p. 5235-5250 , 1995.




A STUDY OF WATER VAPOR OVER HELLAS USING OMEGA/MARS EXPRESS T. Encrenaz

1, R. Melchiorri

1, T. Fouchet

1, P. Drossart

1, B. Gondet

2, Y. Langevin

2, J.-P. Bibring

2, F. Forget

3, L.

Maltagliati4, D. Titov

4.

1LESIA, Paris Observatory, Meudon, France,

2 IAS, Orsay, France,

3LMD, Paris,

France, 4MPS, Katlenburg-Lindau, Germany. [email protected]

We have used the OMEGA imaging spectrometer

aboard Mars Express to study the evolution of the

water vapor abundance over Hellas basin, as a

function of the seasonal cycle. The water vapor

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.




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.




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,

Germany. [email protected]

Introduction: Monitoring the polar caps is

important for understanding the current climate on

Mars. The northern residual cap (NRC) is an

important source of atmospheric water vapour and it

is unclear if there is a net sublimation or deposition

of water vapour under present climatic conditions.

The recession of the seasonal polar cap in spring

uncovers the bright NRC, which is composed of

H2O ice [1]. First observations suggested

interannual changes of the NRC in extent [2, 3]. [4]

showed that these observed changes are most likely

due to seasonal variations. Detailed examination of

the interior cap showed that the albedo varies

spatially but is generally repeatable interannual [5].

Hyperspectral data from OMEGA acquired between

LS ~93° and ~127° in 2004/2005 revealed that the

albedo decrease on the polar cap is due to the

sublimation of fine-grained frost which exposes

older large-grained ice while in outlying regions,

dominated by large-grained ice, the albedo increases

with time [6]. We tracked the seasonal and possible

interannual albedo changes of the NRC using the

Lambert albedo derived from High Resolution

Stereo Camera (HRSC) image data [7] of the

summer seasons 2004/2005 (LS ~120°-~160°) and

2006/2007 (LS ~90°-~150).

Dataset and Method: For our analysis we used

44 (22 in each summer) HRSC images north of

75°N. Although the temporal and spatial coverage is

not complete, there are several areas of yearly and

multi-year repeated coverage with high resolution.

Lambert albedos were derived for the panchromatic

nadir channel (675 ± 90 nm), the blue channel (440

± 45 nm), the green channel (530 ± 45 nm), the red

channel (750 ± 20 nm), and the near-infrared

channel (970 ± 45 nm). After radiometric correction,

HRSC image data are given in units of I/F

corresponding to the ‘‘radiance factor’’ [8], i.e. the

ratio of the surface reflectance as measured and the

reflectance of a perfectly diffuse surface illuminated

at 0° incidence [7]. To derive the Lambert albedo,

incidence angles were determined for each pixel

separately to account for the large variance of

illumination within and between the image scenes.

The emission angles of the red and infrared channel

are -15.9° and +15.9, respectively and the large

viewing angle offsets are likely to have an influence

on the measurements due to different scattering

contributions of the atmosphere and surface.

Simultaneous observations at identical atmospheric

conditions with the OMEGA imaging spectrometer

on Mars Express [9] revealed largest discrepancies

with the HRSC red and infrared channel whereas the

blue and green channel with emission angles of -

3.3° and +3.3°, respectively, generally agree well

with the OMEGA observations [7, 10]. Therefore,

our analysis is mainly based on the green channel

(less influenced by atmospheric conditions than the

blue channel) and, if not available, on the nadir

channel. To constrain the evolution of the residual

cap, we made Lambert albedo mosaics of images

which were acquired within a time span of ~5° LS

and compared their percental albedo changes. For

interannual changes, we compared single images

which were acquired within ~2° LS in each year.

Results: Albedo changes in the first year

(2004/2005) from LS ~120° to ~130° of the cap are

vary by about 15% (-5% to +10%) while the cap

edges (e.g., 68°E/79.5°N) brighten in this period by

20% to 50%. From LS ~130° to ~140° the albedo of

the cap decreases by ~10%. Changes at the cap

edges vary regionally. Some areas brighten by up to

~20% (e.g., 265°E/82°N), others darken by up to

~20% (e.g. Olympia Planitia). Between LS ~140° to

~160° the albedo of the cap as well as of the cap

edge decreases by about 20% but visible inspection

of the image data indicates an increased atmospheric

dust load. Currently, we are analyzing the

observations from the second Martian year

(2006/2007) to track the evolution during summer

and to compare it to the first year of observation.

Conclusions: The relatively stable cap albedo

from LS ~120° to ~130° is in agreement with the

results of Region B (42.5°E, 85.1°N) in [6]. Both

OMEGA measurements at LS 117.4° and 127.6° are

near the final spectrum indicating large-grained ice.

The following albedo decrease of the cap until LS

~160° is similar to the measurements of [11] close

to the geographic north pole in 1999, 2001 and

2003. They observed a strong drop in albedo

starting at LS 134°. The observed differences in

albedo changes of the cap and of the cap edges from

LS ~120° to ~140° shows the complex regional

variability as also reported by [5]. References: [1] Kieffer, H.H. et al. (1976) Science 194, 1341-

1344. [2] James, P.B and L. Martin (1985) Bull. Am. Astron. Soc

17, 735. [3] Kieffer, H.H. (1990) JGR 95, 1481-1493. [4] Bass,

D. et al. (2000) Icarus 144, 382-396. [5] Hale A. S. et al. (2005)

Icarus 174, 502-512. [6] Langevin Y. et al. (2005) Science 307,

1581-1584. [7] Jaumann, R. et al. (2007) PSS 55, 928-952. [8]

Hapke, B. (1993) Cambridge University Press, 455 p. [9] Bibring

J.-P. et al. (2004) ESA SP 1240, 37-49. [10] McCord T.B. et al.

(2007) JGR 12, DOI:10.1029/2006JE002769. [11] Benson J.L.

and P.B. James (2005) Icarus 174, 397-409.




SURFACE ENHANCED RESONANCE RAMAN SPECTROSCOPY AS A COMPLIMENTARY TOOL

FOR PLANETARY EXPLORATION R. Wilson1, S. A. Bowden

2, J.M. Cooper

1 and J. Parnell

2

1 Dept.

Electronics & Bioelectronics, University of Glasgow, UK. 2 Dept Geol. & Petrol Geol., University of Aberdeen,

UK. [email protected]

Studies using conventional laser Raman

instruments have made a good case for application

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.




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.

2German Aerospace Center, Institute

of Planetary Research. [email protected]

Abstract: Slope deposits are prominent types of

large-scale landforms on Mars related to the

gravitational movement of debris masses and rock-

slope material in high-relief terrain, such as scarps

and massifs of the dichotomy boundary [1-2],

impact craters, volcanic edifices and tectonic

structures [3-4]. Such landforms occur at all

latitudes on Mars and cover essentially all geologic

eras of Martian history. Some of them are

considered to be related to the release of water/ice

and are connected to the past and/or present Martian

climatic environment.

We here focus on a system of tongue-shaped

features located at the footslope of Olympus Mons

at 221.5°E, 19°N. These features consist of two

overlapping morphological units which are

superimposed on an older (basal) flow unit. They

are framed by and partly superimposed on two

massive landslide deposits extending towards the

western volcanic plains. The tongue-shaped

landfroms are superimposed on or are part of a large

spatulate flow feature controversially discussed as

glacial [5], cold-based glacial [6,7] or non-glacial

[8] in origin. The same controversy also applies to

the small tongue-shaped flows discussed herein for

which a periglacial [9], or a (alpine) glacial [10]

origin were proposed.

It is shown here that the tongue-shaped features

have a mass-wasting origin related to the

disintegration of basal talus material at the footslope

of Olympus Mons as well as to the destabilization of

avalanche-deposit margins. The tongue-shaped units

are considered to represent the youngest sequence of

scarp-related mass-wasting processes. Talus

disintegration might be related to the release of

water or thawing of near-surface permafrost bodies,

indicated by closed and debris-filled depressions

suggestive of thermokarstic degradation. Although

the overall shape of these landforms roughly

resembles certain rock-glaciers [9], textural and

structural properties (e.g., distribution of ridges and

furrows as well as rocky and fine-grained material)

and the interrelationship of these features to

terminal and marginal areas imply a landslide

origin. Consequently, although the climatic

boundary conditions were favourable to have

facilitated rock-glacier formation, a paraglacial

origin related to slope destabilization as a

consequence of the proposed glacial degradation

and retreat at the western scarp of Olympus Mons

seems more appropriate.

References: [1] Squyres, Icarus, 34, 1978; [2]

Squyres,JGR, 84, 1979; [3] Lucchitta, JGR, 84, 1979 ; [4]

Quantin, Icarus, 172, 2004; [5] Lucchitta, Icarus, 45,

1981; [6] Head, Geology, 31,2003; [7] Milkovich, Icarus,

181, 2006; [8] Carr & Schaber, JGR, 82, 1977; [9] Head,

Nature, 434, 2005; [10] Basilevsky, Sol. Sys. Res., 39,

2005.

Figure 1. Geomorphologic map of the western scarp of Olympus Mons showing flow units (blue) situated at the footslope

and framed by debris avalanche deposits. Flow units have been controversially discussed as either glacial or periglacial in

nature.




TARGETING THE SEARCH FOR LIFE: EVALUATION OF LIFE INDICATING PARAMETERS ON

THE QUIRAING BASALT, ISLE OF SKYE, SCOTLAND S. J. M. Phillips & J. Parnell. Department of

Geology and Petroleum Geology, School of Geosciences, College of Physical Sciences, Meston Building,

King’s College, Aberdeen, AB24 3UE, Scotland, UK. [email protected]

The Quiraing basalt (~60 Ma) [1] located 4 km east

of Staffin on the Isle of Skye, Scotland (Fig. 1) was

selected as a test site to evaluate a range of positive

indicators to look for when deciding which rocks to

target to analyse for life. Life indicating parameters

include evidence of water, sedimentary rocks, a

suitable matrix for organic molecules (chemical

sediments, clays), evidence for carbon, (including

pigments), possible energy gradients (including

signs of alteration such as colour), shielding from

irradiation (actual, or in recent past as in debris

flows) and organised structures (such as

stromatolites) [2]. The Quiraing basalt outcrop

contained several variations that were detected

including fracturing, hydrothermal veining, red

weathering surfaces, shielding and bedding. A 5 m2

area was divided into 25 cells and each cell was

analysed for life indicating parameters. A 100 gram

rock sample was collected from each cell and

returned to the laboratory for analysis using

Pyrolysis-Gas Chromatography and an ATP

(adenosine trisphosphate) assay to positively

determine if life was present. This was then

correlated back to the life indicating parameters to

give an optimum set of parameters. A preliminary

evaluation of sample sites will help increase the

possibility of targeting a more favourable area to

search for life when analysing planetary surfaces

such as on Mars.

References: [1] Anderson, F.W. & Dunham K.C. (1966).

Memoirs of the Geological Survey of Great Britain. Her

Majesty’s stationery office, Edinburgh. [2] Gorbushina,

A.A., Krumbein, W.E. & Volkmann, M. (2002),

Astrobiology, 2, 203-213.

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.




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

12489 Berlin, Germany. [email protected]

Introduction: We investigate the effect of the

high-pressure phase transitions on the thermal

evolution of Mars, focusing mainly on the role

played by the Spinel-Perovskite phase transition. The earlier calculations of Martian mantle

convection with phase transitions (e.g. [1], [2], [3]),

however, neglected important aspects of an evolving

planet, i.e. core cooling and the temperature

dependence of the phase transitions. Instead, these

models assumed that the temperature at the base of

the mantle and the mean depth of the phase

transitions was constant in time, assumptions often

made for simplicity in numerical mantle convection

calculations. It has been speculated by [4] that if the

core is allowed to cool that the plumes decrease in

strength on a time scale of a few 100 Ma depending

on the initial superheat of the core.

Model: The thermal evolution has been

calculated with a 2D axi-symmetrical convection

model which considers core cooling. We solve the

time-dependent compressible equations of thermo-

convection for an infinite Prandtl number fluid. This

is done by solving both the coupled vorticity-stream

function system and the temperature in an axi-

symmetric spherical domain. The buoyancy terms

and latent heat terms associated with the phase

transitions are also taken into account as well as the

wandering of the phase transitions, the decay of the

radiogenic heat sources with time and the core

cooling.

The core size of Mars being unknown, we

consider two core sizes in these modeling. The core

sizes we choose allow the presence of the deepest

phase transition namely the Spinel-Perovskite phase

transition. This endothermic transition, which is

temperature and pressure dependent, moves

downwards due to the cooling mantle.

Results: Depending on the initial thickness of the

Perovskite layer and the efficiency of mantle

cooling, basically two different scenarios can be

observed: For an early perovskite layer of more than

about 150 km, this lower mantle layer initially

convects separately from the upper mantle and

becomes conductive later in the evolution as the

lower mantle layer shrinks (fig 1).

This transition towards a conductive layer

prevents the core from cooling and consequently

stops any active dynamo. This process is very

similar to the one advocated by [5] who explain the

cessation of the dynamo by the transition from a

tectonic plate regime to a stagnant lid regime.

Moreover this transition convective-conductive is

correlated with the growth of a huge hot plume that

induces a degree-1 convection pattern in the upper

mantle.

For an early Perovskite layer of less than about

50 km, the layer is initially conductive and

disappears during the cooling of the planet. As the

lower mantle disappears, the cooling of the core

becomes more efficient and may reactivate a

dynamo. Some younger magnetised features as the

one described by [6] may favour such reactivation.

Generally, one important finding of our modeling

is that strong mantle plumes are most likely not

stable for a long period of time in the Martian

evolution if one considers the temperature

dependence of the Spinel to Perovskite phase

transition.

References: [1] Harder, H., 1998, JGR, 103, 16775. [3]

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.




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,

[email protected] 2

University of Potsdam, Department of Geography, Geomatics Section

Geographic Information Systems (GIS) are

powerful tools for integration of different planetary

datasets, e.g. images, spectral data, and digital

terrain models which are typically given in different

formats like vector and raster. We are currently

involved in a project to import large volumes of data

from the recent Mars missions into a planetary GIS

database.

Before working in GIS with such datasets, it is

necessary to prepare them for import. Using

ArcOBJECTS, a collection of ArcGIS programming

objects, and object oriented programming languages

like Visual Basic .NET, we create ESRI shape files

according to a suitable specification. Regular shape

files are not sufficient, because data points often

have large numbers of attributes associated with

them in the original ASCII dataset. Here, the MOLA

(Mars Orbiting Laser Altimeter) dataset is a typical

example with over 33 attributes per Laser shot.

These have to be imported using a .dbf database file.

Once this is accomplished, it is possible to combine

all these different datasets with raster information,

such as HRSC (High Resolution Stereo Camera), or

MOC (Mars Orbiter Camera) images, or MDIM 2.1

maps for joint analysis.

In addition we implemented window front ends to

access a planetary MySQL Database for creating

specific shape files. So it is possible to search in

huge datasets for attributes or label points from

MOLA, TES, HRSC, the USGS crater catalogue for

instance using ArcOBJECTS and MySQL

Connector Net 1.0.8. This results in smaller datasets

which facilitate the handling in ArcGIS. Also we

began studying the widespread Martian drainage

networks using ESRI`s “ModelBuilder” to automate

the time-consuming step by step workflow. The goal

is to find pour points for runoff water and

watersheds.

Furthermore, we implemented an algorithm to

calculate surface roughness from calibrated MOLA

shots joined with slopemaps, generated from HRSC

digital elevation models.

We have also developed an ArcGIS toolbar with

several tasks for better handling huge datasets for

reprojecting and displaying raster information.

Further, there are a lot of calculate functions for area

measurement, attribute write outs or for joining the

different raster and vector datasets to derive new

scientific results. Also, we programmed modules of

easier handling of map layouts in ArcGIS.

We applied our GIS tools for various geologic

mapping and interpretation tasks as well as for 2d

and 3d visualisation and analysis.

We will demonstrate several examples of importing,

making measurements and reprojecting in large data

sets in different formats with ESRI’s object model

for ArcGIS 9.X.

Figure 1. GIS results around Gusev crater




THE TYRRHENA-MALEA VOLCANIC PROVINCE, MARS R. Greeley

1, D. A. Williams

1, R. L.

Fergason1, R. O. Kuzmin

2, J. Raitala

3, G. Neukum

4, D. Baratoux

5, P. Pinet

5, L. Xiao

6, and the HRSC Team.

1Arizona State University, SESE, Box 871404, Tempe AZ 85287-1404, USA.

2Vernadsky Institute, Russian

Academy of Sciences, Moscow, Russia. 3University of Oulu, Oulu, Finland.

4Institute of Geological Sciences,

Freie Universität, Berlin, Germany. 5Observatoire Midi-Pyrenees, UMR 5562, Toulouse, France.

6Research

Center for Space Science and Technology, China University of Geosciences, Wuhan, 430074 China.

[email protected]

Tyrrhena, Hadriaca, Amphitrites, and Peneus are

highland paterae, which are characterized as broad,

low-profile “central-vent” volcanoes of mafic

composition (Pinet and Chevrel, 1990). Tyrrhena

and Hadriaca are on the northeast side of the Hellas

impact basin, while Amphitrites and Peneus Paterae

are part of Malea Planum, which is superposed on

the south southwest rim of the basin. South and

southwest lie Malea and Pityusa Paterae. Although

all of these features were inferred from Mariner 9

data to be volcanic, the generally poor atmospheric

"seeing" conditions at the high southern latitudes

precluded high quality imaging until recently. New

synoptic data from HRSC, THEMIS and local

"sample" images from MOC and HiRISE are

providing new insight into the volcanic features of

Malea Planum, including the volcanic paterae that

are the eruptive sources. Ages estimated from

impact crater frequencies suggest that initial

eruptions began at about the same time for the four

structures Tyrrhena (3.9 Ga), Hadriaca (3.9 Ga),

Amphitrites (3.7 Ga), and Peneus (3.7 Ga). (Crater

counts have not yet been completed for Malea and

Pityuse Paterae). We suggest that collectively, the

six volcanic patera and Malea Planum can define a

major volcanic region, here termed the Tyrrhena-

Malea Volcanic Province (TMVP), which could be

tectonically linked to the Hellas impact structures.

In all, TMVP covers 2.6 x 106 km

2, stretching some

5000 km from Tyrrhena Patera to the southern

extent of flows from Pityusa Patera. Although not

continuous across the floor of Hellas, TMVP is

comparable in extent to the well-known Tharsis

volcanic province.

References: Pinet, P. and S. Chevrel (1990), JGR 95,

DOI: 10.1029/90JB00703.

Figure 1. Fig. 1. Topographic rendition from MOLA data showing the proposed Tyrrhena Malea Volcanic Province (in red)

superposed on the Hellas impact basin. Tyrrhena, Hadriaca, Amphitrites, and Peneus Paterae all have ages >3.6 Ga;

together with Malea and Pityusa Paterae, these were the primary eruptive centers for the province and might have utilized

structures associated with Hellas as vent foci.




VOLCANIC EDIFICES AT THE MARTIAN NORTH POLE – NEW INVESTIGATIONS T. Kneissl

1, G.

Neukum1. 1

Institute of Geosciences, FU Berlin, 12249 Berlin, Germany. [email protected]

Small-scale volcanic features in the Martian north

polar area were identified and investigated by

several authors in the past (e.g. [1 - 4]). These

investigations were mostly made on the basis of

image data of the Viking mission and topography

data from the Mars Global Surveyor (MGS) Mars

Orbiter Laser Altimeter (MOLA). MOLA data

made it possible to distinguish possible volcanic

cones and domes from other features like impact

craters. Domes were generally considered to be

broad and flat shield-like constructs, while cones

were considered to be steeper volcanic landforms

with a central crater. These volcanic edifices occur

in the height-range of several tens of meters up to

more than hundred meters with diameters up to

more than 20 km.

MOLA data helped to conduct first studies

regarding the possible formation mechanisms.

Basaltic effusive and explosive volcanisms are

discussed in connection with the Martian volatile

distribution in the subsurface that may play a

significant role. The already identified possible

volcanic edifices are situated mainly in two

regions, in a dark polar dune field between 240° to

300° east and 75° to 85° north and in the area

between 195° to 215° east and 72° to 80° north.

The High Resolution Camera (HRSC) onboard

ESA’s Mars Express spacecraft has achieved

almost a full coverage of the north polar region

with resolutions between 12 m/px and 200 m/px.

Both image data and derived digital terrain models

are an excellent basis for a re-investigation of the

distribution and characteristics of possible volcanic

edifices in the north polar region. Thus far, detailed

information has been obtained regarding the shape

and morphometry of major edifices. The general

young age as obtained from age determinations on

the basis of high-resolution MOC image data

suggest recent or even ongoing geologic activity

[5]. Additionally, the distribution of these

landforms has been re-investigated in more detail

which helped to identify many more features

clustered in the two regions mentioned above. References: [1] Garvin, J.B. et al., (2000), Icarus, 145,

648-652. [2] Sakimoto, S.E.H. et al., (2000), LPSC XXXI,

Abs.#1971. [3] Wright, H.M. et al., (2000) LPSC XXXI,

Abs. #1894. [4] Sakimoto, S.E.H. et al., (2001), LPSC

XXXII, #1808. [5] Neukum, G. et al., (2006), EPSC,

p.621.




WATER ICE AT LOW LATITUDES ON MARS F. G. Carrozzo

1, G. Bellucci

1, F. Altieri

1, J-P Bibring

2.

1INAF – Istituto di Fisica dello Spazio Interplanetario, 00133, via del Fosso del Cavaliere 100, Rome, Italy.

2Institut d'Astrophysique Spatiale, Orsay Campus, France. [email protected]

Introduction: Although water vapour is one of

the minor constituents in the Martian atmosphere, it

plays a key role in the climate of the planet, together

with the carbon dioxide and the dust. The key factor

in the stability of water ice is the amount of

humidity in the atmosphere: the more it increases,

the higher the probability that water vapour

saturates. On Mars, water ice is found in the polar

cap deposits, in the clouds in the form of ice crystals

and at high latitudes in the form of surface ice. This

work reports on the identification of superficial

water ice at low latitudes on Mars with OMEGA

instrument on board of Mars Express.

Discussion: The identification of water ice is

based on detection of three bands: ~1.5 m, ~2.0

m and ~2.5 m. The water ice is mostly found on

the slopes, along the walls of numerous craters,

scarps and feet of hills [1]. The areas where the ice

is observed are in shadow. This shows a clear

relationship with the local illumination condition,

which favours the deposition of water ice on the

surface when the temperatures are very low. In the

orbits analysed, water ice is found on 74 tracks

between 30°S and 15°S (2 during fall and 72 during

winter). Water ice can deposit on the surface if the

atmosphere is saturated. In order to verify this, we

have implemented a thermal model, including the

effect of the illumination on the slopes. Here we

report as example the case of cube 1221_4 (Ls 136°)

from the OMEGA dataset (fig 1). We fixed two

values for thermal inertia: 150 J m-2

k-1

s-1/2

(mean

values in the terrains outside the scarp and craters)

350 J m-2

k-1

s-1/2

(mean values in the terrains inside

the scarp and craters). Thermal inertia data have

been taken from TES map [2]. The bolometric

albedo is fixed at 0.24 [3]. According to the model,

for thermal inertia of 350 J m-2

k-1

s-1/2

(fig. 2), the

atmosphere becomes saturated the whole day during

the end fall and early in the winter as soon as the

slope is increased up to 20°. For a slope of 25° the

atmosphere is saturated during all the day for Ls

between 50° and 120°. Decreasing the slope (20°)

we got saturation all the day in a shorter seasonal

period (Ls=75°÷100°). In the remaining cases the

saturation never occurs between 11:00 and 15:00 in

local time. The period of daily saturation is longer

for the high thermal inertia terrains compared to

lower ones and in the last case we have 24 hours

saturation only for 25° slope.

Conclusions: The ice observed by OMEGA

can be: 1) ice that deposits during the period in

which saturation occurs and then sublimes for a

short period. 2) Stable ice during the period in

which all day saturation occurs. At the moment, the

thermal model does not allow us to discriminate

between them because it does not account for the

sublimation/deposition rate. Moreover, the presence

of ice changes the thermal properties of the surface

and sub-surface. In particular, its presence increases

the thermal inertia [4], which in turn favours its

stability.

References: (Times New Roman, 9pt.) [1] Carrozzo et al.

(2007), LPSC XXXVIII, Abs. #2096. [2] Putzig et al.

(2005), Icarus 173, 325-341. [3] Christensen et al. (2001),

JGR, 106, 23,823-23,872. [4] Mellon and Putzig (2007),

LPSC XXXVIII, Abs. #2184.

Figure 1. Figure 1a shows the OMEGA cube 1221_4

centred at 139.8°W and 26.8°S, Ls=136°. The pixels with water

frost are colored according to the 1.5 m band depth. Figure 1b

shows some examples of icy OMEGA spectra from the scarp

with band depth 0.02 (black), 0.04 (blue), 0.07 (green) and 0.10

(red).

Figure 2. In the figures are reported, for different slopes

(5°, 15°, 25°), the saturation state (grey area) and no-saturation

state (white area) during the martian day (y-axis) as a function the

solar longitude between 0° and 180° (x-axis). The latitude is that

of figure 1. We assumed a thermal inertia 350 J m-2 k-1 s-1/2,

bolometric albedo: 0.24, slope azimuth: 180° (south oriented).




THE ORIGIN OF PERENNIAL WATER ICE AT THE SOUTH POLE OF MARS. F. Montmessin

1,2, R.

M. Haberle1, F. Forget

3, R. T. Clancy

4, J.-P. Bibring

5, and Y. Langevin

5 1NASA Ames Research Center, Moffett

Field, USA, 2now at Service d’Aéronomie, Verrières le Buisson, France,

3Laboratoire de Météorologie

Dynamique, Paris, France, 4Space Science Institute, Boulder, Colorado, USA,

5Institut d’Astrophysique

Spatiale, Orsay Campus, France. [email protected]

The poles of Mars are known to have recorded

recent (<107 years) climatic changes. While the

South Polar Region appears to have preserved its

million-year-old environment from major

resurfacing events, except for the small portion

containing the CO2 residual cap, the discovery of

residual water ice units in areas adjacent to the cap

provides compelling evidence for recent

glaciological activity. The mapping and

characterization of these H2O-rich terrains by

OMEGA onboard Mars Express, which have

supplemented earlier findings by Mars Odyssey and

Mars Global Surveyor, have raised a number of

questions related to their origin. We propose that

these water ice deposits are the relics of Mars' orbit

precession cycle and that they were laid down when

perihelion was synchronized with northern summer;

i.e. more than 10,000 years ago. We favor

precession over other possible explanations because

(1) as shown by our General Circulation Model

(GCM) and previous studies, current climate is not

conducive to the accumulation of water at the South

Pole due to an unfavorable volatile transport and

insolation configuration, (2) the residual CO2 ice

cap, which is known to cold-trap water molecules

on its surface and which probably controls the

current extent of the water ice units, is geologically

younger, (3) our GCM shows that 21,500 years ago,

when perihelion occurred during northern spring,

water ice at the North Pole was no longer stable and

accumulated instead near the South Pole with rates

as high as 1 mm/year. This could have led to the

formation of a meters-thick circumpolar water ice

mantle. As perihelion slowly shifted back to the

current value, southern summer insolation

intensified and the water ice layer became unstable.

The layer recessed poleward until the residual CO2

ice cover eventually formed on top of it and

protected water ice from further sublimation (see

Fig. 1). In this polar accumulation process, water ice

clouds play a critical role since they regulate the

exchange of water between hemispheres. The so-

called “Clancy Effect”, which sequesters water in

the spring/summer hemisphere coinciding with

aphelion due to cloud sedimentation, is

demonstrated to be comparable in magnitude to the

circulation bias forced by the north-to-south

topographic dichotomy.

Figure 1. Illustration summarizing the sequence of events in the south polar region since the last “reversed perihelion”

regime of the precession cycle. (1) At that time, water was extracted off the north polar cap and was deposited over the south

PLD terrains thanks to a favourable summer insolation gradient between the poles. (2) The passage to present-day

configuration, with perihelion argument now entering a northern spring regime, reversed the orientation of the insolation

gradient and forced water to progressively return back to the North Pole. (3) In a third act, erosion process stopped as

permanent CO2 ice slabs formed and kept water from subliming further.




WATER VAPOR IN THE MARTIAN ATMOSPHERE BY SPICAM IR/MARS-EXPRESS: TWO

YEARS OF OBSERVATIONS A.Y. Trokhimovskiy1, A.A. Fedorova

1, O.I. Korablev

1, J.-L. Bertaux

2,3, E.

Villard2,3

, A.V. Rodin1, L. Joly

2,4.

1Space Research Institute (IKI), 84/32 Profsoyuznaya, 117810 Moscow,

Russia. 2Service d'Aéronomie du CNRS/IPSL, BP.3, 91371, Verrières-le-Buisson, France.

3Institut Pierre Simon

Laplace, Université de Versailles-Saint-Quentin, 78 Saint-Quentin en Yvelines, France. 4Groupe de

Spectrométrie Moléculaire et Atmosphérique, Moulin de la Housse 51687, Reims, France. [email protected]

Introduction: SPICAM experiment along with

PFS and OMEGA spectrometers on Mars Express

has a capability to sound the water vapor in the

atmosphere. The results of H2O measurements have

been intensively published during last two years [1,

2, 3, 4]. Here we present the new analysis of

SPICAM IR water vapour measurements, covering

almost two Martian years. The near-IR channel of

SPICAM experiment on Mars Express spacecraft is

a 800-g acousto-optic tunable filter (AOTF)-based

spectrometer operating in the spectral range of 1-

1.7 μm with resolving power of ~2000 [5, 6]. The

nadir measurements of H2O in the 1.37-μm spectral

band is one of the main objectives of the

experiment.

Data treatment: As compared with previous

analysis of water vapor presented in [4] we used

the spectroscopic database HITRAN2004 [7]

instead of HITRAN 2000 and the most recent

measurements of the water line-width broadening

in CO2 atmosphere.

The new version of the Martian Climate

Database V4.2 [8] was adopted for modeling of

synthetic spectra and a scenario based on TES

MY24 was used.

Figure 1. A portion of SPICAM spectrum, showing

the vicinity of H2O absorption band at 1.38 μm, and the

adjacent CO2 band at 1.43 μm. 10 subsequent spectra of

orbit 30 are averaged. The synthetic model assumes 8

pr.μm of atmospheric water. SPICAM spectrum is

measured on orbit 30, LS = 335.7, latitude 65S, longitude

58W, and local time of 09:20 [6].

The spare model of SPICAM IR instrument was

recalibrated in June 2007 in Reims, to analyze

specifically the sensitivity to the H2O vapor band.

According to laboratory measurements, a leakage

from the AOTF is responsible up to 5% of signal in

sharp absorption features, making the apparent

depth of the H2O band lower. However, corrected

SPICAM results remain lower than results of other

experiments on Mars-Express that could not be

explained longer by instrumental problems of

SPICAM.

Radiative transfer modeling: Sensitivity of

retrieval to aerosol scattering and different vertical

distributions of aerosol and water vapor was

analyzed for H2O absorption band at 1.38 and 2.56

μm. The aerosol scattering will be accounted for in

further analysis of the bulk of SPICAM data.

Results: We present the results from January

2004 (Ls = 330°, MY26) to August 2007 (Ls =

290°, MY28), i.e. almost two Martian years. The

seasonal trend of water vapor obtained by SPICAM

IR is consistent with TES results and disagrees

with MAWD South pole maximum measurements.

The maximum abundance is 50-55 pr. μm at the

north pole (during MY28 data are missing) and 13-

16 pr.μm at the south pole. The northern tropical

maximum amounts to 12-15 pr μm.

Figure 2. Map of water vapor for one and half Martian

year of SPICAM IR observation

Acknowledge: Russian team acknowledges support

from RFBR grants 07-02-00850 and 06-02-72563.

References: [1] Fouchet, T., (2007), Icarus 190, 32-49.

[2] Melchiorri, R. (2007), PSS 55, 333-342. [3]

Encrenaz, Th. (2005), A&A 441, L9-L12. [4] Fedorova,

A. et al. (2006), JGR 111, DOI:10.1029/2006JE002695. [5] Bertaux, J.-L. et al. (2006), JGR 111,

DOI:10.1029/2006JE002690. [6] Korablev, O. et al.

(2006), JGR 111, DOI:10.1029/2006JE002696. [7]

Rothman, L.S. et al. (2005), JQSRT, 96, 139-204. [8]

Forget, F. et al. (2007), LPICo1353.3098F.




WESTERN PROMETHEI TERRA SMOOTH PLAINS REGION ON MARS: IMPLICATIONS OF

VOLCANIC ORIGIN AND ANALYSIS OF SURFACE FEATURES. J. Raitala1, V.-P. Kostama

1, J.

Korteniemi1, M. A. Ivanov

1,2, T. Törmänen

1 and G. Neukum

3;

1Astronomy, Department of Physical Sciences,

PO BOX 3000, FI-90014 University of Oulu, Finland; 2Vernadsky Institute, Moscow, Russia;

3Institute of

Geological Sciences, Freie Universität, Berlin, Germany. <[email protected]>

The study area within the W Promethei Terra

(36-50°S, 90-106°E) is ~700 km across. It is a

distinct area on the smoothened eastern Hellas basin

rim (Fig.1), and consists of two topographic parts: a

regional slope of ~0.07° eastward of ~97°E and a

steeper ~0.88° westward of ~97°E. To the NE, E,

and S the region is confined by Noachian cratered

terrain, and the central area is cut by the large

canyons of Harmakhis, Reull and Teviot Valles. The

western and central areas exhibit smooth Hesperian

plains. The plains are multi-layered, which is seen

on the walls of the canyons that cut them. The

average observed thickness of the Promethei layers

is ~70-80 m and the typical measured slope of the

walls of the canyons is ~25-30o. This gives an

estimate of the thickness of the layers, which is ~35-

45 m. The true thickness of the stack of layers is

estimated from observations to be ~1-1.3 km.

Similar stacks of layers are seen in other regions of

Mars where the interiors of lava plains are exposed

(i.e. Lunae Planum and Syrtis Major).

Besides the observed layering, the plains have a

variety of surface features, some genetically related

to the basement material and others due to

deposition/modification of younger materials

(mesas, channels [1-3]). In many places, but mostly

in the eastern portion of the area, the preferentially

E-W oriented wrinkle ridges deform the surface of

the plains. Another set of long straight narrow

ridges (widths < km, heights 10s m, lengths 10s km)

are also seen on the surface of the plains. They

occur in preferentially NE-SW-oriented groups. The

regional topography does not control the distribution

of the long ridges. Their morphologic

characteristics, areal distribution, and close

association with the lava plains are consistent with

and suggest that the straight ridges may represent

exhumed dikes, which have served as feeders for the

lava plains. These two ridge types and the layered

structure suggest that the regional basement material

is of volcanic origin. The volume of the layered

material in this region is estimated to be ~0.3 x 106

km3

and the time of emplacement of the material

may correspond to the Late Noachian-Early

Hesperian epochs. The topographic characteristics

of the area of the smooth plains in Promethei Terra

collectively suggest that if the plains were emplaced

in the void on the rim of Hellas, it likely was a

steep-sided trough-like depression [4]. The thermal

erosion of an ice-saturated megaregolith may

explain this and therefore must precede the phase of

the massive volcanism [5]. Therefore, the first two

major episodes in the evolution of the eastern

portion of the Hellas rim likely are: (1) Erosion of a

large amount (~0.5 x 106 km

3) of regolith from the

central-western portion of Promethei Terra and (2)

emplacement of volcanic basement material within

the created void and formation of the possibly

independent volcanic province of Western

Promethei Terra [4,6]. The abundant fluvial-related

features and the younger age of terrains to the south

of Reull also suggest that there was also a third

episode related to relatively late resurfacing [1-4].

References: [1] Kostama et al. (2007) LPSC 38;

[2] Korteniemi et al. (2007) 7th Mars; [3]

Korteniemi et al. (2007) Vernadsky-Brown 46; [4]

Raitala et al. (2007) submitted. [5] Tanaka et al.

(2002) GRL; [6] Kostama et al. (2007) Vernadsky-

Brown 46.

Figure 1. The location of the probable volcanic province (white ellipse) in regional context, plotted on MOLA DTM.

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ESA Mars Research Abstracts Part 2