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The small crater Airy-0 was selected from Mariner 9 images to be the reference for the Mars prime meridian. Initial analyses in the year 2000 tied Viking Orbiter and Mars Orbiter Camera images of Airy-0 to the evolving Mars Orbiter Laser Altimeter global digital terrain model to update the location of Airy-0. Based upon this tie and radiometric tracking of landers/rovers from Earth, new expressions for the Mars spin axis direction, spin rate, and prime meridian epoch value were produced to define the orientation of the Martian surface in inertial space over time. Since the Mars Global Surveyor mission and Mars Orbiter Laser Altimeter global digital terrain model were completed some time ago, a more exhaustive study has been performed to determine the accuracy of the Airy-0 location and orientation of Mars at the standard epoch. Thermal Emission Imaging System (THEMIS) IR image cubes of the Airy and Gale crater regions were tied to the global terrain grid using precision stereo photogrammetric image processing techniques. The Airy-0 location was determined to be about 0.001° east of its predicted location using the currently defined International Astronomical Union (IAU) prime meridian location. Information on this new location and how it was derived will be provided to the NASA Mars Exploration Program Geodesy and Cartography Working Group for their assessment. This NASA group will make a recommendation to the IAU Working Group on Cartographic Coordinates and Rotational Elements to update the expression for the Mars spin axis direction, spin rate, and prime meridian location.

Tharsis Tholus, a more than 3.9 Ga old composite shield volcano to the east of the major Tharsis Montes, has experienced a complex history of growth and destruction. On the basis of new high resolution images we analysed the morphology as well as the tectonic structures of the Tharsis Tholus volcano in detail. From morphological data, cross-cutting relations of the surface structures, and crater modelling ages we propose a chronostratigraphy for the volcano-tectonic history of Tharsis Tholus. The strongly faulted volcano reveals two large-scale landslide events followed by two subsequent shield re-growth phases between 3.8 and 1.7 Ga and two caldera collapses. Tharsis Tholus was also affected by regional extensional tectonics between 1.7 Ga and 0.4 Ga recorded by sub-parallel sets of NE trending graben structures. The steep and up to 5.4 km high landslide scarps on Tharsis Tholus suggest deep faulting of the edifice. In order to confirm this hypothesis we used analogue sand box models in which we demonstrated that gravitational flank movement on top of weak basal substrata may have produced the deformation structures as observed on Tharsis Tholus. (C) 2011 Elsevier B.V. All rights reserved.

Martian gullies show morphologic characteristics that indicate the presence of a flowing liquid during their formation process. However, the current climate in the northern hemisphere of Mars does not allow the existence of substantial amounts of liquid water and, therefore, an assessment of the regional distribution and characteristics of gullies is required in order to understand the most recent climatic development of Mars. We examined High Resolution Stereo Camera (HRSC) and Mars Orbiter Camera—Narrow Angle (MOC-NA) data covering the northern hemisphere in order to determine the distribution, geological setting and orientation of gullied slopes. We used these results to test the consistency of formation models proposed in earlier works and to define the climate dependence. Gully orientations were analyzed only at impact crater walls, because craters have an azimuthally uniform distribution of slopes. This is ideal when investigating the influence of insolation on the distribution and development of landforms. We examined 230 HRSC and ~ 21,000 MOC-NA images and identified more than 2200 gullies on HRSC and more than 3100 gullies in MOC-NA data. Both surveys show the highest densities of gullies between 35°N and 55°N and a shift of gully orientation at 40°N from pole-facing slopes to slopes facing the equator. Orientation trends as well as the latitudinal distribution of identified gullies strongly support the view of gully formation caused by insolation/climatic conditions. This result suggests that formation mechanisms based on the assumption of atmospheric water-ice deposition proposed in earlier work are more likely than processes related to groundwater flow. Further support for a theory related to deposition is gained through observation of the high occurrence of gullies at isolated knobs and hills, which cannot be associated with formation models based upon groundwater flow.

We have investigated 12 prominent regions on Mars in detail with respect to their geologic evolution through time on the basis of detailed geologic mapping exercises and determining age relationships through crater counting techniques using imagery. New data in combination with previously obtained data have been analyzed by way of a refined method of cratering age extraction that also gives fine details of periods of resurfacing. We have found that there has been volcanic and fluvial/glacial geologic activity on the Martian surface at all times from > 4 Ga ago until today. This activity shows episodic pulses in intensity of both volcanic and fluvial/glacial processes at ~ 3.8–3.3 Ga, 2.0–1.8 Ga, 1.6 to 1.2 Ga, ~ 800 to 300 m.y., ~ 200 m.y., and ~ 100 m.y., and a possible weaker phase around ~ 2.5–2.2 Ga ago. In between these episodes, there was relative quiescence of volcanic and/or fluvial/glacial activity. The episodes we find on the Martian surface in the crater frequency analyses of HRSC, MOC and THEMIS data coincide with some age groups of the Martian meteorites (~ 1.3 Ga, ~ 600–300 m.y., ~ 170 m.y.). It appears that the surface activity expressions and their episodicity are related to the interior evolution of the planet when convection in the asymptotic stationary state changes from the so-called stagnant-lid regime to an episodic behavior. Similarities in episodic behavior are found for the other terrestrial planets: Venus, the Earth's moon, and the Earth itself suggesting a common general relationship in the evolutionary tracks.

Four prominent fluvial systems, Dao, Niger, and Harmakhis Valles, and a system of smaller channels between Niger and Harmakhis Valles characterize a wide topographic trough that connects Hesperia Planum and Hellas Planitia (Hesperia–Hellas trough, HHT). The source depressions of the canyons of Dao and Harmakhis Valles are at the lowermost level (about − 5.4 km and lower), the system of Niger Vallis is at medium elevation (about − 3 to − 4 km), and the small channels occur at the highest topographic levels (about − 2.5 to − 3.2 km). All systems are among the youngest features in the region. They formed due to the release of subsurface water and apparently have their sources at different topographic levels below or within the layered suite of volcanic deposits in HHT. Two scenarios can explain the subsurface nature of their sources: (1) flow of liquid ground water from extensive watershed toward the source areas and accumulation of water there and (2) deposition of ice-bearing materials in and around the sources. The first scenario requires the existence of extensive confining layers at different topographic levels corresponding to the levels of the floors of the canyons/channels. Because the presence of such confining layers on Mars is problematic, we consider the first scenario to be less likely. We believe that the deposition of ice-bearing materials at different levels during the growth of the lava plateau in HHT can better explain the key features of the fluvial systems: the subsurface nature of the sources and the relatively young age. In this scenario, the higher topographic position of the ice-bearing deposits corresponds to the younger age of deposition. The possible accumulations of ice were buried by successive volcanic deposits, including lava flows in HHT and later were heated from below by magmatic intrusions that caused melting of the icy deposits, release of water, and formation of the fluvial structures.

We describe the procedure to fit a cumulative production function polynomial to a partial crater size–frequency distribution. The technique is of particular use in deriving ages for surfaces which have undergone partial resurfacing events: namely, erosional or depositional events which have affected a limited diameter range of the crater population. We demonstrate its use in obtaining times for both the surface formation and the resurfacing event. We give a practical outline of the method for making age measurements from crater counts and how to identify resurfacing effects in the results. We discuss the conversion of production function polynomials between common presentations, and the statistical uncertainty of the determined ages with respect to the non-linear chronology function, and a minor refinement of data binning.

An improved technique is presented to determine more accurately deposit thicknesses of any surface units using crater size-frequency distributions (CSFDs). This new approach enables thickness estimates of deposits that completely cover their underlying unit, i.e., where no flooded craters are observed. Here, the crater populations of the unit of interest and its underlying unit (or a representative surface area) are measured and their cratering model ages determined. The CSFD of the younger unit is then treated as if it had the same age as the older unit and the expected cumulative numbers of craters for each bin size are calculated. The crater deficiency for each bin size is then determined by subtracting the expected cumulative crater number from the observed cumulative crater number per bin. The probability is calculated such that the crater deficiency corresponds to the age of the older, underlying unit. The crater bin size where we are 95% confident that these craters are there but just covered by the younger unit is chosen to derive the deposit thickness value using known crater diameter-rim height relationships. We demonstrate the use of the method at two sites of different scales and settings on Mars: Pickering crater and Lunae Planum, and compare our estimates with those obtained by other methods. For the Lunae Planum area a significantly higher minimum thickness of > 840 m is predicted.

[1] We investigate the formation of the outflow channels Dao and Niger Valles near the eastern rim of the Hellas impact basin, Mars. Methods used include image and topography analysis, surface age determination, and finite element modeling. Observations show that deep depressions, source regions for Dao and Niger Valles, are located in an area of shallow subsidence to the south and east of the volcano Hadriaca Patera. Cratering model ages allow for fluvial processes triggered by volcanic loading. Based on the observations, we develop a numerical model of volcanic loading on top of a poroelastic plate leading to flexure and fracturing of the lithosphere. Modeling results show that fracturing may occur up to a depth of about 6 km within an area of shallow subsidence, i.e., the moat surrounding the volcano. Depending on initial aquifer pressurization, groundwater could have reached the surface. Model discharges and channel morphometry suggest that the Dao Vallis channel never reached bankfull flow and that the wetted channel perimeter may have formed during multiple outflow events. The following scenario is proposed: (1) emplacement of a volcanic load on top of a confined, overpressurized aquifer in the early Hesperian, (2) fracturing around the load, possibly reactivated during various stages of volcanic activity, (3) channeling of groundwater to the surface along fractures and outflow channel formation during several events in the Hesperian, and (4) collapse, mass wasting and modification of depressions in the Amazonian.

[1] Basaltic mare materials extruded between ∼ 4 Ga and ∼ 1.2 Ga ago and formed large expanses on the lunar nearside. In contrast to the maria, volcanic shields, domes, and cones also occur but cover only a small area of the total lunar surface. One specific type of dome is characterized (1) by steep slopes, (2) by a high albedo similar to lunar highlands, and (3) by a strong absorption in the ultraviolet. Because of this latter feature they appear spectrally red and therefore were termed red spots. Their morphology suggests that these domes were created by much more viscous, silica-rich lava. Also, red spots which are believed to be possibly volcanic in origin are associated with light plains materials. In order to constrain the formation of these presumed volcanic features in context with the formation of the surrounding mare and highland materials, we carried out (1) geologic and morphologic mapping of high-resolution Lunar Orbiter IV and Apollo 16 frames and (2) measurements of superimposed crater frequencies on geologic units in order to assess stratigraphy and geological history in the following two volcanic provinces in which red spots occur: (1) the Hansteen region, dominated by the dome Hansteen α (also known as the Arrowhead), and (2) the Helmet region, dominated by light plains, possibly of volcanic origin, such as the feature called the Helmet and Darney χ, and the southern Montes Riphaeus. In the Hansteen region, the dome Hansteen α postdates craters Billy (3.88 Ga) and Hansteen (3.87 Ga) but predates most of the mare materials (3.51 Ga). Cratering model ages of the dome range from 3.74 to 3.65 Ga, placing its major activity as extending into the late Imbrian. In the Helmet region, presumed highland volcanism associated with red spots started in the Nectarian with Darney χ (3.94 Ga). Cratering model ages measured on the Helmet light plains range from 3.8 to 2.08 Ga. The origin of these light plains and local resurfacing processes in this unit, however (highland volcanism or impact), remains an open question. Locally, younger processes of Eratosthenian or even Copernican age, such as emplacement of mare materials or mass wasting deposits, could be dated in the higher-resolution Apollo 16 frame.

International audience; The Rosetta spacecraft accomplished a flyby of Mars on its way to 67P/Churyumov-Gerasimenko on 25 February 2007. In this paper we describe the measurements obtained by the M channel of the Visual and Infrared Thermal Imaging Spectrometer (VIRTIS-M) and the first scientific results derived from their analysis. The broad spectral coverage of the VIRTIS-M in the IR permitted the study of various phenomena occurring in the Martian atmosphere; observations were further exploited to achieve accurate absolute radiometric calibration. Nighttime data from the VIRTIS-M constrain the air temperature profile in the lower atmosphere (5-30 km), using variations in CO2 opacity at 4.3 mum. A comparison of this data with the global circulation model (GCM) by Forget et al. (1999) shows a trend of slightly higher air temperature in the VIRTIS-M retrievals; this is accompanied by the presence of moderate decreases (˜5 K) in large sections of the equatorial region. This is potentially related to the occurrence of water ice clouds. Daytime data from the VIRTIS-M reveal CO2 non-local thermodynamic equilibrium emission in the high atmosphere. A mapping of emission intensity confirms its strict dependence on solar zenith angle. Additionally, devoted limb observations allowed the retrieval of vertical emission intensity profiles, indicating a peak around 105 km in southern tropical regions. Ozone content can be effectively monitored by the emission of O2 (a1Deltag) at 1.27 mum. Retrieved emission intensity shows that polar regions are particularly rich in ozone. Aerosol scattering was observed in the 1-2.5 mum region above the night region above the night disk, suggesting the occurrence of very high noctilucent clouds.

[1] We report on ages derived from impact crater counts for exposed mare basalt units in the northern part of the lunar nearside hemisphere (Mare Frigoris), the eastern and northeastern part of the nearside hemisphere (Lacus Temporis, Joliot, Hubble, Goddard, Mare Marginis, and Mare Smythii), the central part of the nearside hemisphere (Palus Putredinis, Mare Vaporum, and Sinus Medii), and the southwestern part of the nearside hemisphere (Grimaldi, Cruger, Rocca A, Lacus Aestatis, and Schickard). In Mare Frigoris, we dated 37 basalt units, showing ages from 2.61 to 3.77 Gyr, with most units being formed in the late Imbrian period between 3.4 and 3.8 Gyr ago. In Mare Vaporum we dated six spectrally homogeneous units that show model ages of 3.10 to 3.61 Gyr. Our model ages of basalts in Mare Marginis range from 3.38 to 3.88 Gyr and are mostly older than basalts in Mare Smythii (3.14–3.48 Gyr). The model ages of four units in Sinus Medii indicate that the basalts in this region formed 3.63 to 3.79 Gyr ago. We find an excellent agreement of our crater size-frequency model ages of the Palus Putredinis area, which contains the Apollo 15 landing site, with the radiometric ages of Apollo 15 samples. According to our crater counts, basalts in Palus Putredinis are 3.34 Gyr old and this compares favorably with the radiometric ages of 3.30–3.35 Gyr of the olivine-normative and quartz-normative basalts of the Apollo 15 landing site. Lacus Aestatis is a small irregular-shaped mare patch in the southwestern nearside and shows an Imbrian age of 3.50 Gyr; basalts in Lacus Temporis in the northeastern nearside formed between 3.62 and 3.74 Gyr ago and are, therefore, older than the basalts in Lacus Aestatis. We found that basalts in craters of the southwestern nearside (Schickard, Grimaldi, Cruger, and Rocca A) are also mostly younger than basalts in craters of the northeastern nearside (Hubble, Joliot, and Goddard). While basalt ages vary between 3.16 and 3.75 Gyr in the southwest, basalts in the northeast are 3.60–3.79 Gyr old. These results confirm and extend the general distribution of ages of mare basalt volcanism and further underline the predominance of older mare basalt ages in the eastern and southern nearside and in patches of mare peripheral to the larger maria, in contrast to the younger basalt ages on the western nearside (Oceanus Procellarum).

[1] Geologic evolution of the western portion of Promethei Terra (40–50°S, 90–105°E) consists of two major episodes. The first was the formation of a late Noachian–early Hesperian volcanic province. During this episode a kilometer-thick suite (≈0.3 × 106 km3) of lavas accumulated to the south of Reull and Harmakhis Valles within a regional depression that probably was formed earlier in a massive erosion. A topographic barrier of the cratered terrains separates the volcanic western Promethei plateau from Hesperia Planum. Thus, the lavas could not originate from Hesperia Planum. The topographic position of western Promethei Terra excludes a possible lava feeding from Hadriaca Patera. Dikes from Tyrrhena Patera potentially could be the lava source. The massive erosion before formation of the lava plateau in western Promethei Terra is poorly consistent, however, with the hypothesis that the distal dikes served as the source and favors emplacement of sills or a local source for the lavas. The second episode after the lava plateau emplacement was surface modification by glacial, fluvial, and aeolian processes. The time of the episode is poorly known, and it could continue from the late Hesperian and through the entire Amazonian epoch. Mesas, channels, and esker-like ridges suggest that at the earlier stages of this episode (late Hesperian–early Amazonian) a glacier or ice dust deposits covered the western Promethei Terra area studied. The 50–60 × 103 km3 formation could contain 15–18 to 40–50 × 103 km3 of ice. This is comparable with the ice involved into of lobate apron formation (5–13 × 103 km3) at the later Amazonian stages of western Promethei Terra. Because of the uncertainties in the ice volume estimates, the total amount of water in the earlier glacier and later aprons was approximately the same.

The fretted terrain of the Martian dichotomy boundary is a key region for investigating landforms related to creep of ice and debris as it exhibits landforms comparable to morphologies of periglacial environments. Although features known as lobate debris aprons, lineated valley fills and concentric crater fills have been studied in great detail, basic questions concerned with the composition and the style of emplacement and degradation still remain unanswered. This study focuses on morphologies which are located in a nearA¢Â�Â�circular depression located at the dichotomy escarpment in Deuteronilus Mensae. Analysis of highA¢Â�Â�resolution image data suggests an early formation of these features as the result of backward thermokarstic degradation of highland terrain. Geologically younger processes caused deposition and degradation of an iceA¢Â�Â�rich mantling deposit, which ultimately led to formation of creep morphologies that might have even been active in the geologically recent past. Intermixing of both degradational landform units form complex patterns that cannot be explained by a late stage (glacial) process alone. Morphological comparisons of lineated valley fill units with concentric crater fill landforms in Utopia Planitia strongly suggest comparable emplacement and degradation styles of these features. The sequential development consisting of an initial probably widespread thermokarstic degradation followed by later cyclic deposition of volatile material and continual thermokarstic degradation suggests that the development of parts of the Martian fretted terrain is directly related to climatic variations in the planetA¢Â�Â�s history.

A valley network in the Western Libya Montes region extending from 80°E to 83.4°E and 1.2°S to 3.9°N exhibits significant detail about water release processes and duration of erosional activities. The valley system originates at the southeastern offshoot of Syrtis Major and drains down to Isidis Planitia over a distance of about 300 km. Midstream, the valley network splits into a shorter eastern and a longer western part. For most of its length, the valley exhibits an interior channel which allows constraining discharges ranging from 15,000 m3/s to 430,000 m3/s and yielding sediment volumes up to 250 tons/s. Based on stratigraphic relations, the valley system evolved during a period of about 2.8 billion years with major wet episodes in the Noachian (< 4.1 Ga), the Hesperian (3.6 to 3.0 Ga) and the Amazonian (2.8 to 1.4 Ga). While precipitation dominated the fluvial activity during the Noachian era, as indicated by dendritic drainage pattern, the close correlation of lava deposits and valley source regions suggests that volcanic processes, such as ground ice melting and/or hydrothermal water release, might have played a major role during the Hesperian and Amazonian fluvial activities. In addition, multiple volcanic events in the Hesperian and Amazonian ages show that Syrtis Major was active, at least locally, until 1.4 Ga ago. Discharge estimates demonstrate a significant increase from precipitation-induced to volcanic-triggered water release. Fluvial erosion rates, discharges, sediment transport rates, and the lack of any widespread chemical alteration products, such as sulphates and phyllosilicates that are indicative for large long-lasting standing bodies of water, suggest relatively short valley formation times of only a few thousands of years. Compared with the total age of the valley system, short episodic water release events separated by long dry and inactive periods of some hundreds of millions of years, at least in the Hesperian and Amazonian, seem to be more realistic rather than a continuous flow. The youngest channel segment dates back to less than 1.4 Ga. This indicates the presence of surface runoff at that time, which required atmospheric conditions that support water to be stable on the surface even late in Martian history. In summary, the Western Libya Montes Valley System demonstrates episodic and multi-genetic erosion events over most of the Martian history, with a clear change in the erosion style from precipitation to volcanic-triggered water release during the early Hesperian.

International audience; Volcanism on Mars occurred until recently, but the mineralogy of recent lava plains is poorly known because few regions display fresh outcrops devoid of dust. Using visible and near infrared data of the Mars Express probe, two new volcanic plains in Noctis Labyrinthus have been identified, and the existence of a volcanic plain on the floor of Echus Chasma has been confirmed. Crater retention ages estimated for these three plains range between 50 and 100 My, corresponding to the Late Amazonian. These plains represent an excellent opportunity to constrain the mineralogy of recent volcanic rocks. Results show that basaltic compositions with plagioclase and high calcium pyroxene are predominant. The low olivine proportion suggests that the apparent fluidity of these flat plains is not related to magmas being ultramafic. In addition, a platy-ridged texture is observed in two of the studied regions. Our study shows, for the first time, that this texture is associated with volcanic rocks, and that these rocks are of typical basaltic mineralogy. Finally, these volcanic plains are located more than 1000 km east of previously known late Amazonian volcanic centers of the Tharsis region, an observation to be taken into account when considering models of recent volcanism on Mars. (C) 2009 Elsevier BM. All rights reserved.

We present a large-scale spring hypothesis for the formation of various enigmatic light-toned deposits (LTDs) on Mars. Layered to massive LTDs occur extensively in Valles Marineris, chaotic terrains, and several large craters, in particular, those located in Arabia Terra. Most of these deposits are not easily explained with either a single process or multiple ones, either in combination or occurring sequentially. Spring deposits can have a very wide range of internal facies and exhibit complex architectural variations. We propose the concept of large-scale spring deposits for explaining LTDs on Mars. Stable volcano-tectonic settings, such as the ones typical on Mars, are compatible with a large-scale, long-term, multistage formation of spring deposits. The large-scale spring deposit model can explain the formation of LTDs with a common process, although active in different times and locations, compatible with coeval local or regional processes and deposits, such as volcaniclastic ones. LTDs, if formed as spring deposits derived from subsurface fluids, could potentially offer favorable conditions both to life and to the fossilization of past life forms.

[1] Fluvial landforms are observed on the plateau near Echus Chasma and other locations on the Valles Marineris plateau using high resolution (10 to 50 m/pixel) Mars Express images and topography. Branching valleys have a 20- to 100-m deep V-shape profile typical of fluvial processes. Their incision occurred in a thin (

[1] Circumferential depressions enclosing mesas and plateaus in the northern Kasei Valles and in the Tartarus Colles regions of Mars are interpreted as indicators of the former extent of lobate debris aprons, thought to be mixtures of ice and clastic particles. These former lobate debris aprons existed about 1 Ga ago and were embayed by lavas or other flow deposits. After the lobate debris aprons had been removed by sublimation and deflation, topographic depressions with a depth of 50 m and a width of several kilometers were left behind between the mesa or plateau scarp and the solidified flow materials. These depressions or moats are located equatorward of ±30° at significantly lower latitudes than generally observed for occurrences of modern, intact lobate debris aprons. This observation provides evidence that the paleoclimate at that time was different than today, probably due to a higher averaged obliquity of the planet's rotational axis.

[1] We present analysis of a bright haze observed inside Valles Marineris, which formed in mid northern spring. The data were collected by the High Resolution Stereo Camera (HRSC) and the imaging spectrometer, Observatoire pour la Mineralogie, l'Eau, les Glaces et l'Activite (OMEGA), aboard Mars Express. This study provides a case example of the power of simultaneous multiple emission angle and hyperspectral imaging for study of aerosols and clouds in the Martian atmosphere. The haze appeared thinner after three days and disappeared in nine days. It was limited to a 2-km layer at the bottom of the canyon. The color was redder than the underlying surface. The analysis of the OMEGA spectra indicates that this haze was composed of dust particles. The dust layer appeared brighter with the HRSC stereo channels than the nadir channel due to longer scattering paths. We have estimated the optical depth of the haze by fitting both HRSC and OMEGA data with radiative transfer calculations. The retrieval of the optical depth is very sensitive to the aerosol scattering model used and the reflectance of the surface. Applying an aerosol scattering model derived from sky surveys at a constant elevation by the Imager for Mars Pathfinder, the optical depth of the haze is estimated from HRSC data to be within 1.7 to 2.3 at the wavelength (λ) of 0.675 μm. The wavelength dependence is obtained from OMEGA spectrum. It increases to 2.2–2.6 at λ = 1.35 μm and moderately decreases to 1.2–1.8 at λ = 2.4 μm.

[1] The Nili Fossae region located on the northwestern quadrant of the Isidis Basin, Mars, displays superb exposures of bedrock outcrops that reveal mineralogy and composition of the crust. Previous work has shown that this region exhibits the largest exposures of olivine-dominated rock units on Mars. Visible-infrared imaging spectrometer data acquired by OMEGA were calibrated to surface reflectance and analyzed to determine surface mineralogy. The dominant minerals identified are iron-bearing mafic silicates (low- and high-calcium pyroxene, olivine) and water-bearing phyllosilicate (iron-rich smectite clay). The strength and position of mineral absorption features were used to produce mineral indicator maps for the dominant species, and the maps were integrated with high-resolution imaging from the THEMIS, MOC, and HRSC instruments, and MOLA topography. We show that olivine and phyllosilicate occur in spatially distinct outcrops; the olivine-bearing rock unit is a meters- to tens of meters-thick cap unit resting on phyllosilicate-bearing bedrock, and the phyllosilicate units predate the Isidis basin-forming event. On the basis of superposition, crosscutting, and geomorphic relationships, we interpret the emplacement of the olivine-bearing units as having been contemporaneous with the Isidis impact event. By analogy with the Orientale basin on the Moon, we propose that the olivine-bearing unit represents the surface exposure of the impact melt from the Isidis impact event. These results demonstrate that large regions of crust had been altered in the presence of water prior to the date of the Isidis basin-forming event in the Late Noachian (≈3.96 Ga).

[1] Nili Fossae is a region with fresh exposures of old material. OMEGA spectral data have revealed a strong diversity of minerals in the Nili Fossae region, including mafic minerals and phyllosilicates. Phyllosilicates are found to be primarily Fe-rich smectites (a type of clay minerals). These phyllosilicates are observed only on the Noachian crust, in four types of locations: Three are without any signatures other than phyllosilicates, on (1) smooth layered deposits, (2) massive exhumed material, and (3) crater ejecta, and one type of location is associated with olivine (4) over an altered mafic unit. This diversity suggests a widespread liquid water activity and a variety of alteration processes. The formation of phyllosilicates underwent considerable reduction, if not cessation, by the time of formation of Syrtis Major lava flows in the Hesperian period. Phyllosilicates are spatially and chronologically disconnected from the fluvial and depositional landforms observed in the Nili Fossae region. Observations suggest that these landforms formed after the extensive formation of phyllosilicates during short periods of time and/or temperatures close to freezing.

[1] OMEGA/Mars Express has discovered large outcrops rich in phyllosilicates in the region of Mawrth Vallis, Mars (around 20°W, 25°N). The region is located in Noachian highly cratered terrains, close to the limit of the Martian dichotomy, where the outflow channel Mawrth Vallis cuts the highlands. We have examined this region using OMEGA spectra of the surface from 0.9 μm to 2.6 μm, with spatial sampling from 500 m to 3 km, offering a full coverage of the region. OMEGA spectra show two broad bands centered at 1 μm and 2.2 μm, revealing the presence of clinopyroxene on dark surfaces. Phyllosilicates have been identified by absorption bands at 1.4 μm, 1.9 μm, and 2.2 or 2.3 μm. Comparison with laboratory spectra reveals similarities with Al-OH smectites such as montmorillonites, or Fe- or Mg-OH smectites such as nontronite. A precise location of the phyllosilicate-rich areas on visible HRSC images indicates that they are placed exclusively on bright outcrops, mostly on the plateaus, dated to the Noachian period. On HRSC and MOC images the phyllosilicate-rich outcrops reveal strong erosional features such as numerous residual buttes composed of layers a few meters thick. The phyllosilicate-rich unit corresponds to a geological unit more than 100 m thick, over a horizontal extension approximately of 300 km × 400 km. This unit implies a large volume of altered rocks, either in situ or after transport and deposition, in Noachian terrains, revealing a different climatic and geologic environment from the present one.

[1] The High Resolution Stereo Camera (HRSC) on the Mars Express (MEx) spacecraft in orbit about Mars is delivering images of the Mars surface and atmosphere from its high-inclination, elliptical orbit that are intended to cover most of the Mars surface by the end of the mission. These are acquired in four specific spectral passbands that sample the scene at different parts of the extended visual spectrum from near 0.4 μm to about 1μm and in five additional channels, four used for stereo and one nadir, that use the same fifth spectral passband. These channels view from different angles and can be used for photometric and spectral analysis. The data allow a variety of studies of the surface and atmosphere of Mars, but these require detailed understanding of the data characteristics, as there are several effects, including of the observation geometry and the atmosphere, that need to be treated. This article is a report to the science community on the current understanding of the HRSC color data after an initial analysis and also a demonstration of their capabilities for study of Mars, which, from this study, appear to be considerable. We first discuss the characteristics and calibration of the HRSC spectral data. The atmospheric influence on the photometric appearance of the surface materials is explored, and methods of compensating are considered. Reflectances are calculated for surface units by assuming the dark material is basalt and scaling the entire HRSC scene using a laboratory basalt spectrum for the selected dark material area; the results and the correction factors for different orbits are compared. We present several examples of what could become a large volume of future science results utilizing the HRSC spectrophotometric data in conjunction with photo-interpretation. For example, we find relatively few but very important distinct spectral components at the HRSC resolution, including what we interpret as iron oxide-rich material, unoxidized basalt, and polar ice, consistent with past works. The major spectral components are distributed in ways that suggest compositional mixing and the need to revise the current geologic understanding of major regions.

[1] We report on estimates of the rheological properties of late-stage lava flows on the eastern flank of Ascraeus Mons, Mars. From previous studies it is known that the dimensions of flows reflect rheological properties such as yield strength, effusion rates, and viscosity. Our estimates are based on new high-resolution images obtained by the High Resolution Stereo Camera (HRSC) on board the European Space Agency's Mars Express spacecraft in combination with Mars Orbiter Laser Altimeter (MOLA) data. Compared to earlier studies, the high spatial resolution of the HRSC and MOLA data allowed us to map 25 late-stage lava flows and to measure their dimensions, as well as their morphological characteristics, in greater detail. Our estimates of the yield strengths for these flows range from ∼2.0 × 102 Pa to ∼1.3 × 105 Pa, with an average of ∼2.1 × 104 Pa. These values are in good agreement with estimates for terrestrial basaltic lava flows and are comparable to previous estimates derived for a small number of lava flows on Ascraeus Mons. Our investigation indicates that the effusion rates for the studied Ascraeus Mons flows are on average ∼185 m3 s−1, ranging from ∼23 m3 s−1 to ∼404 m3 s−1. 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. However, our effusion rates are similar to terrestrial effusion rates of Kilauea and Mauna Loa and other Martian volcanoes. On the basis of our estimates of the effusion rates and the measured dimensions of the flows, we calculated that the time necessary to emplace the flows is on average ∼26 days. Viscosities were estimated on the basis of yield strengths and effusion rates, yielding average values of ∼4.1 × 106 Pa-s and ranging from ∼1.8 × 104 Pa-s to ∼4.2 × 107 Pa-s. On the basis of newly available data sets (e.g., HRSC, MOLA) we are now able not only to identify possible differences in eruptive behavior between Ascraeus Mons and Elysium Mons but also to study such differences over time.

[1] Lava flow mapping was conducted on a north-south transect of Olympus Mons using the European Space Agency's Mars Express High Resolution Stereo Camera (HRSC) image H0037. The HRSC image was coregistered to Mars Orbiter Laser Altimeter, Thermal Emission Imaging System, and Mars Orbiter Camera data, enabling lava flow structures to be differentiated and mapped consistently across the shield. Because different structures develop as a result of different effusive conditions, their abundance and distribution provide insight into the eruptive history of a shield volcano. Results show that lava channels are the dominant flow structure, whereas tabular sheets are most common beyond the basal scarp. A hummocky unit dominates the summit area and likely represents a combination of (1) volcanic lava flows, (2) pyroclastic deposits, (3) a dust mantle, and (4) frozen volatiles, all of which have been suggested to exist on Olympus Mons in the past. Lava fans are typically associated with lava tubes, indicating that they represent tube outbreaks as was previously suggested as one possible formation mechanism. No vents were identified, suggesting that major rift zones have not developed on the north or south flank. Younger channel-fed flows typically embay older tube-related flows, which they outnumber by a ratio of 5:1. Therefore Olympus Mons likely experienced a change in eruptive style from longer-lived, stable, tube-forming eruptions to shorter-lived, less stable, channel-forming eruptions in the late Amazonian. A similar trend exists for the Hawaiian volcanoes in which a decrease in the magma production rate drives a change to dominantly channel forming eruptions associated with increased shield age.

[1] We use spaceborne image and topographic data of Mars for a geomorphological study of the origin, development, and postemplacement modification of a 30-km-long spatulate landform in the Hellas Montes area, east of Hellas Planitia. The area is characterized by an abundance of lobate debris aprons, all of which have been interpreted as a result of creep and viscous deformation of rock and ice mixtures that make the eastern Hellas Planitia assemblages a type location for possible Martian rock glaciers. The spatulate landform has been discussed controversially in the past. Explanations range from a rock glacier to a wet debris avalanche or debris flow origin. This work scrutinizes arguments provided in earlier works and presents evidence that a landslide origin of the spatulate landform, connected with a sector collapse of a volcanic construct, is conceivable. We also find extensive observational evidence that the landslide has recently undergone or is still undergoing significant postemplacement modifications, which are characteristic of rock glaciers in periglacial environments. Creep deformation as well as disintegration of the original surface by thermokarstic degradation is considered to have obliterated impact craters and, in consequence, create a seemingly younger age. Similar processes are considered to be common on Mars under current climatic conditions. Our results imply that the strong and possible long-term modification of landforms on Mars by cold-climate processes often hinders the identification of their original nature and complicates the assessment of true ages.

We present observations of seasonal variations in polygonal crack patterns located in a polar trough on the south polar cap of Mars; previously, there was no direct observation showing that these patterns change. Polygonal patterns on Mars are attributed to thermal contraction cracking, which is commonly observed in periglacial environments on Earth. In this paper we discuss observations based upon the high-resolution image data of the Mars Orbiter Camera and focus on the reconstruction of the seasonal development. The image-based observations are further supported by temperature data. We show that the south polar trough pattern is located in an active geologic unit, which undergoes seasonal variations and annual crack formation. Furthermore, there are strong indications showing these contraction-crack processes take place in a thin layer that might be composed of water-ice and is located beneath the seasonal carbon dioxide ice cover.

[1] The High Resolution Stereo Color camera (HRSC) data over Hesperia Planum and its surroundings reveal important details of geologic episodes and water-related processes in this region. (1) The Noachian fluvial events of Hesperia Planum depression included accumulation of water and formation of a water/ice reservoir there. Later, the reservoir was depleted in several phases reflecting diminishing amounts of water. Climate changes and/or volcanism were important in these volatile releases. (2) The massive, 0.45–1.5 × 106 km3, erosion from the Hesperia depression before the main lava eruption possibly resulted in thick, 0.5–1.5 km, deposits in Hellas Planitia. (3) Measurements of the flooded craters within Hesperia Planum provide the estimates of the thickness of lavas there, about 250–500 m. The final volume of lavas within Hesperia Planum (0.4–0.7 × 106 km3) is comparable with the range of some terrestrial igneous provinces such as Columbia River Basalts. (4) Extended magmatism possibly triggered formation of the outflow channel in a few locations after the lava emplacement. During this episode, about 0.04 × 106 km3 of material (about 4.5–8.9% of the volume eroded in the episode of massive erosion) were removed. The thickness of the composite lava layer exposed on the walls of the outflow channels, a few hundreds of meters, corresponds well to the thickness estimates made by the measurements of the flooded craters. (5) Dispersed viscous flows (debris aprons, flow-like deposits) reflect the final fluvial events. Viscous flows from the subsurface sources in the Southwestern trough associate with Dao, Niger, and Harmakhis Valles. These flows represent the final volatile discharge from the Hesperia reservoir that mostly was depleted by the earlier events of massive erosion and formation of the outflow channels. Viscous surface flows are mostly associated with Reull Vallis and probably reflect redistribution of volatiles related to the late episodes of evolution of this outflow channel.

[1] Tyrrhena Terra, located just north of the roughly 200 km Hellas impact basin, is a typical region of the ancient cratered southern highlands on Mars. Its base material is a mixture of Hellas ejecta, prebasin remnants, and highland terrain, all later saturated with smaller (

[1] Transverse dunes cover the valley floor of Nirgal Vallis, a 670 km long valley network at 318°E and 29°S. The dunes are superposed by small undeformed impact craters, which implies that the dunes are inactive under present atmospheric conditions. The last active phase of dune movement (absolute age) can be determined by crater size frequency distributions. The derived absolute ages from high-resolution Mars Global Surveyor Mars Orbiter Camera imagery range from approximately 300,000 years to 1.4 Myr B.P. (error of factor of 2). The dunes are also superposed by gully deposits (aprons), which implies a younger age of the gullies. These superposition relationships show that the gullies in this midlatitude region are at the most 3 Myr old but might be even younger than 300,000 years.

[1] Comprehensive studies of the stratigraphic relationships in the Athabasca Valles region have been performed on imagery from the Viking Orbiters and the Mars Orbiter Camera (narrow and wide angle) on board Mars Global Surveyor as well as by making use of topographic data derived from the Mars Global Surveyor Laser Altimeter (MOLA) instrument. From superimposed crater frequencies we have derived absolute surface ages by applying the Hartmann and Neukum [2001] cratering chronology model. The main channel has been incised into the Cerberus volcanic plains with an average plains age of 3.4 Ga. The main fluvial or glacial erosion processes ended 2.6 Ga ago. This result shows that the age of the valley system itself is older than commonly believed. One major possible fluvial event occurred in the topographically lower volcanic plains southwest of the valley 1.6 Ga ago. The valley itself was covered by lava episodically and ending 0.9 Ga and in case of a few younger episodes 30 Ma ago. The surface texture south of the valley system suggests a younger, possibly fluvial overprinting of the volcanic texture 30 Ma ago. The latest volcanic activity is dated to about 3 Ma ago. With the latter ages we have been able not only to confirm earlier age estimates by Berman and Hartmann [2002] but also to show that the valley system itself has undergone over a period of two billion years active geological cycles which were dominated by volcanic processes at the end.

Accurate estimates of mare basalt ages are necessary to place constraints on the duration and the flux of lunar volcanism as well as on the petrogenesis of lunar mare basalts and their relationship to the thermal evolution of the Moon. We performed new crater size-frequency distribution measurements in order to investigate the stratigraphy of mare basalts in Oceanus Procellarum and related regions such as Mare Nubium, Mare Cognitum, and Mare Insularum. We used high-resolution Clementine color data to define 86 spectrally homogeneous units within these basins, which were then dated with crater counts on Lunar Orbiter IV images. Our crater size-frequency distribution measurements define mineralogical and spectral surface units and offer significant improvements in accuracy over previous analyses. Our data show that volcanism in the investigated region was active over a long period of time from ~3.93 to 1.2 b.y., a total of ~2.7 b.y. Volumetrically, most of the basalts erupted in the Late Imbrian Period between ~3.3 and 3.7 b.y., and we see evidence that numerous units have been resurfaced. During the Eratosthenian Period, significantly less basalt was erupted. Depending on the absolute model ages that one can assign to the lunar chronostratigraphic systems, five units might be of Copernican age. Younger basalts are generally exposed in the center of the investigated area, that is, closer to the volcanic centers of the Aristarchus Plateau and Marius Hills. Older basalts occur preferentially along the northwestern margin of Oceanus Procellarum and in the southeastern regions of the studied area, i.e., in Mare Cognitum and Mare Nubium. Combining the new data with our previously measured ages for basalts in Mare Imbrium, Serenitatis, Tranquillitatis, Humorum, Australe, and Humboldtianum, we find that the period of active volcanism on the Moon lasted ~2.8 b.y., from ~4 b.y. to ~1.2 b.y. On the basis of the basalts dated so far, which do not yet include the potentially young basalts of Mare Smythii [e.g., Schultz and Spudis, 1983], we conclude that Oceanus Procellarum not only exhibits the widest range of ages of all investigated basins but probably also is the location of some of the youngest basalts on the lunar surface.

The multispectral images of the lunar limb and farside obtained by the solid state imaging (SSI) system on board the Galileo spacecraft provide the first new pulse of compositional data of the Moon by a spacecraft in well over a decade. The wavelength range covered by SSI filters (0.4-1.0 }m) is particularly sensitive to the composition of mare basalts, the abundance of mafic (ferrous) minerals, and the maturity of the regolith. To a first order, the limb and farside material is consistent with previous characterization of nearside lunar spectral types for mare and highland soils and craters. Most basalts are of an intermediate TiO 2 composition and most of the highland crust is feldspathic with local variations in mafic content identified principally at impact craters. Dark manfling material on the farside can be interpreted in terms of known properties of lunar pyroclastic glass. Regions of cryptornate are shown to have spectral properties intermediate between those of highland and mare soils, as would be expected from a mixture of the two. There are several important exceptions and surprises, however. Unlike the basalt types identified on the nearside, limb and farside basalts exhibit an exceptionally weak 1 [xm ferrous absorption band. This may indicate a compositionally distinct lunar basalt group that, for example, is more Mg-rich than most basalts of the nearside. Some of the most notable compositional anomalies are associated with South Pole-Aitken Basin. This large region has a much lower albedo than surrounding highlands. The inner, darkest, portion of the basin exhibits optical properties indistinguishable low-Ti basalts. Deposits to the south exhibit unique properties with a strong and broad ferrous 1 [xrn absorption, most consistent with abundant olivine. The unusual compositions associated with South Pole-Aitken and their spatial e3tent suggests the impact creating this huge lunar basin excavated mafic-rich lower crust or perhaps mantle material. 1. INTRO OUCrI'ION

Residual meteoritic material has been detected on the surface of crater interiors on lunar samples 60315,29 and 65315,68. Iron-nickel micrometeoroid residues are present in the form of mettallic spherules embedded in or attached to the crater glass-linings; stony-iron meteorite residual material is homogeneously mixed with the glass-linings. Crater simulation experiments show the dependence of crater diameter to depth ratio on projectile density. On the other hand, the projectile velocities exceeding 4 km/s have no measurable influence on the D/T ratios of microcraters. As a result, diameter-to-depth measurements on lunar microcraters yield three groups of micrometeorites in the size range between 1 μm and 1 mm: iron-nickel, stony-iron and low-density particles. The measured D/T values correspond directly to the kind of determined projectile residues: craters showing residues of iron-nickel meteorites have ratios of D/T = 1.3–1.4 and craters with stony-iron residues have ratios ofD/T = 1.9–2.1. Craters with diameters ≲30 μm seem to have been formed predominantly by iron-nickel micrometeorites, whereas craters with diameters ≳80 μm predominantly by stony meteorites.

Readily applicable equations are derived for calculating the impact velocities, collision intervals, and cumulative crater frequencies in a satellite system for planetocentric and heliocentric impactors. Observed cumulative crater frequencies in the Saturnian satellite system and a sometimes observed lack of apex-antapex asymmetry of crater frequencies favor crater-producing projectiles orbiting initially in elliptic orbits round the parent planet (planetocentric impactors). On the other hand, a sometimes observed asymmetry of crater size distribution and short collision intervals (

We present a compilation of crater-scaling laws suitable for planetological applications, connecting crater diameter to impact energy.

Discrepancies in published crater frequency data prompted this study of lunar crater distributions. Effects modifying production size distributions of impact craters such as surface lava flows, blanketing by ejecta, superposition, infilling, and abrasion of craters, mass wasting, and the contribution of secondary and volcanic craters are discussed. The resulting criteria have been applied in the determination of the size distributions of unmodified impact crater populations in selected lunar regions of different ages. The measured cumulative crater frequencies are used to obtain a general calibration size distribution curve by a normalization procedure. It is found that the lunar impact crater size distribution is largely constant in the size range 0.3 km ⩽D ⩽ 20 km for regions with formation ages between ≈ 3 × 109 yr and ≳ 4 × 109 yr. A polynomial of 4th degree, valid in the size range 0.8 km ⩽D ⩽ 20 km, and a polynomial of 7th degree, valid in the size range 0.3 km ⩽D ⩽ ⩽ 20 km, have been approximated to the logarithm of the cumulative crater frequencyN as a function of the logarithm of crater diameterD. The resulting relationship can be expressed asN ∼D α(D) where α is a function depending onD. This relationship allows the comparison of crater frequencies in different size ranges. Exponential relationships with constant α, commonly used in the literature, are shown to inadequately approximate the lunar impact crater size distribution. Deviations of measured size distributions from the calibration distribution are strongly suggestive of the existence of processes having modified the primary impact crater population.

In this paper a theory is evaluated to describe the development of the lunar crater population with time under the bombardment by meteoroids and solar wind. Starting from a general mass distribution law a differential equation has been established and solved separately for meteoroid impact and solar wind bombardment. The theory permits the calculation of absolute formation ages of the lunar surface as well as the particle flux, supposing the crater distributions on the moon have been measured. As an important result it includes a D−2 equilibrium crater distribution law (D =crater diameter), actually measured in Mare Tranquillitatis and Oceanus Procellarum. Additionally, the exponential decrease of particle flux with time is confirmed.

Experiments on impact craters and impact ionization of dust particles have been performed using iron microparticles from a 2-Mv Van de Graaff accelerator. The particle velocity υ ranged from 0.2 to 40 km/sec and the mass m ranged from 1 × 10−15 g to 5 × 10−10 g. The analysis and mass determination of iron microspheres on smooth targets and of iron layers of the projectile in the middle of impact craters have been done with an electron microprobe. The Kα radiation emitted by the iron layer in the crater has been measured as a function of mass and velocity of the projectile. The radiation measurement gives, in combination with the crater diameter D, a means for the determination of the projectile parameters m and υ. The targets used were Ag, Al, Cu, Cd, and W. Within an error of approximately 20% the total mass of the iron projectile has been found inside the craters in W, Cu, and Al targets at velocities of ≤13 km/sec. The impact ionization has been studied for impact velocities of up to 40 km/sec and projectile masses of down to 10−15g. The yield of either ions or electrons, normalized to the incident mass m, can be described by an empirical relation of the form Q = const ƒ(θ, υ) · mα · υβ, where θ is the angle of incidence. Analysis of impact ionization has been applied to extremely sensitive detectors of cosmic dust particles. The impact cratering and ionization are discussed in terms of shock effects by applying the Rankine-Hugoniot theory. The energy partition in the form of kinetic energy and internal energy (e.g., elastic compression and irreversible heating) is discussed as a function of the velocity of iron particles impacting the targets Al, Cu, W, and Au. As a result, W and Au are targets that transform a greater fraction of the primary energy into heating and ionization of projectile material.

The NASA lunar surface samples 10019 and 10046 from the Apollo 11 mission were searched for micrometeorite impact craters by means of an optical reflecting microscope and a Stereoscan electron microscope. Craters with diameters between 2 microns and 0.8 mm were observed. Approx. 1 crater larger than 0.15 mm in diameter per 1 mm 2 and 1 crater larger than 2 μm per 500 μm 2 could be found. The time of exposure is calculated from crater number densities and cosmic dust influx values to be 10 000 yr and 100 yr for layers of respectivethicknesses. On lunar glassy spherules impact features were found which yield the same crater number densities. The phenomenology of high velocity impacts was studies by shooting microspheres on targets similar to lunar material using a 2 MV dust accelerator.