Mapping of Titan: Results from the first Titan radar passes

The first two swaths collected by Cassini's Titan Radar Mapper were obtained in October of 2004 (Ta) and February of 2005 (T3). The Ta swath provides evidence for cryovolcanic processes, the possible occurrence of fluvial channels and lakes, and some tectonic activity. The T3 swath has extensive areas of dunes and two large impact craters. We interpret the brightness variations in much of the swaths to result from roughness variations caused by fracturing and erosion of Titan's icy surface, with additional contributions from a combination of volume scattering and compositional variations. Despite the small amount of Titan mapped to date, the significant differences between the terrains of the two swaths suggest that Titan is geologically complex. The overall scarcity of impact craters provides evidence that the surface imaged to date is relatively young, with resurfacing by cryovolcanism, fluvial erosion, aeolian erosion, and likely atmospheric deposition of materials. Future radar swaths will help to further define the nature of and extent to which internal and external processes have shaped Titan's surface.     

Impact craters on Titan

Five certain impact craters and 44 additional nearly certain and probable ones have been identified on the 22% of Titan’s surface imaged by Cassini’s high-resolution radar through December 2007. The certain craters have morphologies similar to impact craters on rocky planets, as well as two with radar bright, jagged rims. The less certain craters often appear to be eroded versions of the certain ones. Titan’s craters are modified by a variety of processes including fluvial erosion, mass wasting, burial by dunes and submergence in seas, but there is no compelling evidence of isostatic adjustments as on other icy moons, nor draping by thick atmospheric deposits. The paucity of craters implies that Titan’s surface is quite young, but the modeled age depends on which published crater production rate is assumed. Using the model of Artemieva and Lunine (2005) suggests that craters with diameters smaller than about 35 km are younger than 200 million years old, and larger craters are older. Craters are not distributed uniformly; Xanadu has a crater density 2-9 times greater than the rest of Titan, and the density on equatorial dune areas is much lower than average. There is a small excess of craters on the leading hemisphere, and craters are deficient in the north polar region compared to the rest of the world. The youthful age of Titan overall, and the various erosional states of its likely impact craters, demonstrate that dynamic processes have destroyed most of the early history of the moon, and that multiple processes continue to strongly modify its surface. The existence of 24 possible impact craters with diameters less than 20 km appears consistent with the Ivanov, Basilevsky and Neukum (1997) model of the effectiveness of Titan’s atmosphere in destroying most but not all small projectiles.     

Cassini SAR, radiometry, scatterometry and altimetry observations of Titan’s dune fields

Large expanses of linear dunes cover Titan’s equatorial regions. As the Cassini mission continues, more dune fields are becoming unveiled and examined by the microwave radar in all its modes of operation (SAR, radiometry, scatterometry, altimetry) and with an increasing variety of observational geometries. In this paper, we report on Cassini’s radar instrument observations of the dune fields mapped through May 2009 and present our key findings in terms of Titan’s geology and climate. We estimate that dune fields cover ∼12.5% of Titan’s surface, which corresponds to an area of ∼10 million km², roughly the area of the United States. If dune sand-sized particles are mainly composed of solid organics as suggested by VIMS observations (Cassini Visual and Infrared Mapping Spectrometer) and atmospheric modeling and supported by radiometry data, dune fields are the largest known organic reservoir on Titan. Dune regions are, with the exception of the polar lakes and seas, the least reflective and most emissive features on this moon. Interestingly, we also find a latitudinal dependence in the dune field microwave properties: up to a latitude of ∼11°, dune fields tend to become less emissive and brighter as one moves northward. Above ∼11° this trend is reversed. The microwave signatures of the dune regions are thought to be primarily controlled by the interdune proportion (relative to that of the dune), roughness and degree of sand cover. In agreement with radiometry and scatterometry observations, SAR images suggest that the fraction of interdunes increases northward up to a latitude of ∼14°. In general, scattering from the subsurface (volume scattering and surface scattering from buried interfaces) makes interdunal regions brighter than the dunes. The observed latitudinal trend may therefore also be partially caused by a gradual thinning of the interdunal sand cover or surrounding sand sheets to the north, thus allowing wave penetration in the underlying substrate. Altimetry measurements over dunes have highlighted a region located in the Fensal dune field (∼5° latitude) where the icy bedrock of Titan is likely exposed within smooth interdune areas. The hemispherical assymetry of dune field properties may point to a general reduction in the availability of sediments and/or an increase in the ground humidity toward the north, which could be related to Titan’s asymmetric seasonal polar insolation. Alternatively, it may indicate that either the wind pattern or the topography is less favorable for dune formation in Titan’s northern tropics.     

Mineralogy of Terra Meridiani and western Arabia Terra from OMEGA/MEx and implications for their formation

Analyses of Mars Express OMEGA hyperspectral data (0.4-2.7 μm) for Terra Meridiani and western Arabia Terra show that the northern mantled cratered terrains are covered by dust that is spectrally dominated by nanophase ferric oxides. Dark aeolian dunes inside craters and dark streaks extending from the dunes into the intercrater areas in mantled cratered terrains in western Arabia Terra have similar pyroxene-rich signatures demonstrating that the dunes supply dark basaltic material to create dark streaks. The dissected cratered terrains to the south of the mantled terrains are dominated spectrally by both low-calcium and high-calcium pyroxenes with abundances of 20-30% each retrieved from nonlinear radiative transfer modeling. Spectra over the hematite-bearing plains in Meridiani Planum are characterized by very weak but unique spectral features attributed to a mixture of a dark and featureless component (possibly gray hematite) and minor olivine in some locations. Hydrated minerals (likely hydrous ferric sulfates and/or hydrous hydroxides) associated with poorly ferric crystalline phases are found in the etched terrains to the north and east of the hematite-bearing plains where erosion has exposed ∼1 km of section of layered outcrops with high thermal inertias. These materials are also found in numerous craters in the northern Terra Meridiani and may represent outliers of the etched terrain materials. A few localized spots within the etched terrain also exhibit the spectral signature of Fe-rich phyllosilicates. The ensemble of observations show that the evidence for aqueous processes detected by the Opportunity Rover in Meridiani Planum is widespread and confirms the extended presence of surface or near-surface water over this large region of Mars. The scenarios of formation of Terra Meridiani (“dirty” acidic evaporite, impact surge or weathering of volcanic ash) cannot satisfactorily explain the mineralogy derived from the OMEGA observations. The formation of the etched terrains is consistent with leaching of iron sulfides and formation of sulfates and hydrated iron oxides, either in-place or via transport and evaporation of aqueous fluids and under aqueous conditions less acidic than inferred from rocks examined by Opportunity.     

Titan: An exogenic world?

All landforms on Titan that are unambiguously identifiable can be explained by exogenic processes (aeolian, fluvial, impact cratering, and mass wasting). Previous suggestions of endogenically produced cryovolcanic constructs and flows have, without exception, lacked conclusive diagnostic evidence. The modification of sparse recognizable impact craters (themselves exogenic) can be explained by aeolian and fluvial erosion. Tectonic activity could be driven by global thermal evolution or external forcing, rather than by active interior processes. A lack of cryovolcanism would be consistent with geophysical inferences of a relatively quiescent interior: incomplete differentiation, only minor tidal heating, and possibly a lack of internal convection today. Titan might be most akin to Callisto with weather: an endogenically relatively inactive world with a cool interior. We do not aim to disprove the existence of any and all endogenic activity at Titan, nor to provide definitive alternative hypotheses for all landforms, but instead to inject a necessary level of caution into the discussion. The hypothesis of Titan as a predominantly exogenic world can be tested through additional Cassini observations and analyses of putative cryovolcanic features, geophysical and thermal modeling of Titan’s interior evolution, modeling of icy satellite landscape evolution that is shaped by exogenic processes alone, and consideration of possible means for supplying Titan’s atmospheric constituents that do not rely on cryovolcanism.     

Preliminary mariner 9 report on the geology of Mars

Mariner 9 pictures indicate that the surface of Mars has been shaped by impact, volcanic, tectonic, erosional and depositional activity. The moonlike cratered terrain, identified as the dominant surface unit from the Mariner 6 and 7 flyby data, has proven to be less typical of Mars than previously believed, although extensive in the mid- and high-latitude regions of the southern hemisphere. Martian craters are highly modified but their size-frequency distribution and morphology suggest that most were formed by impact. Circular basins encompassed by rugged terrain and filled with smooth plains material are recognized. These structures, like the craters, are more modified than corresponding features on the Moon and they exercise a less dominant influence on the regional geology. Smooth plains with few visible craters fill the large basins and the floors of larger craters; they also occupy large parts of the northern hemisphere where the plains lap against higher landforms. The middle northern latitudes of Mars from 90 to 150† longitude contain at least four large shield volcanoes each of which is about twice as massive as the largest on Earth. Steep-sided domes with summit craters and large, fresh-appearing volcanic craters with smooth rims are also present in this region. Multiple flow structures, ridges with lobate flanks, chain craters, and sinuous rilles occur in all regions, suggesting widespread volcanism. Evidence for tectonic activity postdating formation of the cratered terrain and some of the plains units is abundant in the equatorial area from 0 to 120° longitude.Some regions exhibit a complex semiradial array of graben that suggest doming and stretching of the surface. Others contain intensity faulted terrain with broader, deeper graben separated by a complex mosaic of flat-topped blocks. An east-west-trending canyon system about 100–200 km wide and about 2500 km long extends through the Coprates-Eos region. The canyons have gullied walls indicative of extensive headward erosion since their initial formation. Regionally depressed areas called chaotic terrain consist of intricately broken and jumbled blocks and appear to result from breaking up and slumping of older geologic units. Compressional features have not been identified in any of the pictures analyzed to data. Plumose light and dark surface markings can be explained by eolian transport. Mariner 9 has thus revealed that Mars is a complex planet with its own distinctive geologic history and that it is less primitive than the Moon.     

Aeolian processes on Venus

In this pre-Magellan review of aeolian processes on Venus we show that the average rate of resurfacing is less than 2 to 4 km/Ga, based on the impact crater size frequency distribution derived from Venera observations, reasonable values of the impact flux, and the assumption of steady state conditions between crater production and obliteration. Viscous relaxation of crater topography, burial by volcanic deposits, tectonic disruption, chemical and mechanical weathering and erosion, and accumulation of windblown sediments probably all contribute to resurfacing. Based on the rate of disappearance of radar-bright haloes around impact craters, the rate of removal of blocky surfaces has been estimated to be about 10^–2 km/Ga. Pioneer-Venus altimetry data show that the average relative permittivity (at 17 cm radar wavelength) of the surface is too high for exposure of soils 10 cm deep, except for ~5% of the planet located primarily in tessarae terrains. The tectonically disrupted tessarae terrains may be sites of soil generation caused by tectonic disruption of bedrock and the presence of relatively steep slopes, or they may be terrains that serve as traps for windblown material. The overall impression is that Venus is a geologically active planet, but one dominated by volcanism and tectonism. On the other hand, theoretical considerations and experimental data on weathering and transport of surface materials suggest rather different conditions. Thermochemical arguments have been advanced that show: (1) CO₂ and SO₂ incorporate into weathering products at high elevation, (2) transport of weathered material by the wind to lower-elevation plains, and (3) re-equilibration of weathered material, releasing both CO₂ and SO₂. In addition, kinetic data suggest a rate of anhydrite formation of 1 km/Ga, a value comparable to the soil erosion rate on Mars, a planet with an active aeolian environment. Experiments and theoretical studies of aeolian processes show that measured surface winds are capable of moving sand and silt on Venus. Assuming that there is a ready sand supply, the flux could be as high as 2.5 × 10^–5 g/cm/s, a value comparable to desert terrains on Earth. In an active aeolian abrasion environment, sand grains could have lifetimes <10³ years. In addition, comminuted debris may be cold-welded to surfaces at the same time as abrasion is occurring. Magellan altimetry and SAR observations should allow assessment of which model for venusian surface modification (active vs. inactive surficial processes) is correct, given the global coverage, high spatial resolution, the calibrated nature of the data, and the potential during extended missions of acquiring multiple SAR views of the surface.Geology and Tectonics of Venus, special issue edited by Alexander T. Basilevsky (USSR Acad. of Sci. Moscow), James W. Head (Brown University, Providence), Gordon H. Pettengill (MIT, Cambridge, Massachusetts) and R. S. Saunders (J.P.L., Pasadena).     

A newly discovered impact crater in Titan's Senkyo: Cassini VIMS observations and comparison with other impact features

Senkyo is an equatorial plain on Titan filled with dunes and surrounded by hummocky plateaus. During the Titan targeted flyby T61 on August 25, 2009, the Cassini Visual and Infrared Mapping Spectrometer (VIMS) onboard the Cassini spacecraft observed a circular feature, centered at 5.4° N and 341°W, that superimposes the dune fields and a bright plateau. This circular feature, which has been named Paxsi by the International Astronomical Union, is 120±10 km in diameter (measured from the outer edge of the crater rim) and exhibits a central bright area that can be interpreted as the central peak or pit of an impact crater. Although there are only a handful of certain impact craters on Titan, there are two other craters that are of similar size to this newly discovered feature and that have been studied by VIMS: Sinlap (Le Mouélic et al., 2008) and Selk (Soderblom et al., 2010). Sinlap is associated with a large downwind, fan-like feature that may have been formed from an impact plume that rapidly expanded and deposited icy particles onto the surface. Although much of the surrounding region is covered with dunes, the plume region is devoid of dunes. The formation process of Selk also appears to have removed (or covered up) dunes from parts of the adjacent dune-filled terrain. The circular feature on Senkyo is quite different: there is no evidence of an ejecta blanket and the crater itself appears to be infilled with dune material. The rim of the crater appears to be eroded by fluvial processes; at one point the rim is breached. The rim is unusually narrow, which may be due to mass wasting on its inside and subsequent infill by dunes. Based on these observations, we interpret this newly discovered feature to be a more eroded crater than both Sinlap and Selk. Paxsi may have formed during a period when Titan was warmer and more ductile than it is currently.     

Characterization of hydrated silicate-bearing outcrops in Tyrrhena Terra,Mars: Implications to the alteration history of Mars

The Tyrrhena Terra region of Mars is studied with the imaging spectrometers OMEGA (Observatoire pour la Minéralogie, l’Eau, les Glaces et l’Activité) onboard Mars Express and CRISM (Compact Reconnaissance Infrared Spectrometer for Mars) onboard Mars Reconnaissance Orbiter, through the observation of tens of craters that impacted into this part of the martian highlands. The 175 detections of hydrated silicates are reported, mainly associated with ejecta blankets, crater walls and rims, and central up-lifts. Sizes of craters where hydrated silicates are detected are highly variable, diameters range from less than 1 km to 42 km. We report the presence of zeolites and phyllosilicates like prehnite, Mg-chlorite, Mg-rich smectites and mixed-layer chlorites–smectites and chlorite–vermiculite from comparison of hyperspectral infrared observations with laboratory spectra. These minerals are associated with fresh craters post-dating any aqueous activity. They likely represent ancient hydrated terrains excavated by the crater-forming impacts, and hence reveal the composition of the altered Noachian crust, although crater-related hydrothermal activity may have played a minor role for the largest craters (>20 km in diameter). Most detected minerals formed over relatively high temperatures (100–300 °C), likely due to aqueous alteration of the Noachian crust by regional low grade metamorphism from the Noachian thermal gradient and/or by extended hydrothermal systems associated with Noachian volcanism and ancient large impact craters. This is in contrast with some other phyllosilicate-bearing regions like Mawrth Vallis where smectites, kaolinites and hydrated silica were mainly identified, pointing to a predominance of surface/shallow sub-surface alteration; and where excavation by impacts played only a minor role. Smooth plains containing hydrated silicates are observed at the boundary between the Noachian altered crust, dissected by fluvial valleys, and the Hesperian unaltered volcanic plains. These plains may correspond to alluvial deposition of eroded material. The highlands of Tyrrhena Terra are therefore particularly well suited for investigating the diversity of hydrated minerals in ancient martian terrains.     

Shoreline retreat at Titan’s Ontario Lacus and Arrakis Planitia from Cassini Imaging Science Subsystem observations

Recent observations by Cassini’s Imaging Science Subsystem reveal that part of the shoreline of Titan’s Ontario Lacus has retreated by several kilometers and may indicate that the dark area that appeared at Arrakis Planitia (80°S, 120°W) in late 2004 has subsequently faded. These changes provide constraints on aspects of Titan’s methane cycle, as well as on the properties of Titan’s surface materials.     

Titan’s cloud seasonal activity from winter to spring with Cassini/VIMS

Since Saturn orbital insertion in July 2004, the Cassini orbiter has been observing Titan throughout most of the northern winter season (October 2002–August 2009) and the beginning of spring, allowing a detailed monitoring of Titan’s cloud coverage at high spatial resolution with close flybys on a monthly basis. This study reports on the analysis of all the near-infrared images of Titan’s clouds acquired by the Visual and Infrared Mapping Spectrometer (VIMS) during 67 targeted flybys of Titan between July 2004 and April 2010.The VIMS observations show numerous sporadic clouds at southern high and mid-latitudes, rare clouds in the equatorial region, and reveal a long-lived cloud cap above the north pole, ubiquitous poleward of 60°N. These observations allow us to follow the evolution of the cloud coverage during almost a 6-year period including the equinox, and greatly help to further constrain global circulation models (GCMs). After 4 years of regular outbursts observed by Cassini between 2004 and 2008, southern polar cloud activity started declining, and completely ceased 1 year before spring equinox. The extensive cloud system over the north pole, stable between 2004 and 2008, progressively fractionated and vanished as Titan entered into northern spring. At southern mid-latitudes, clouds were continuously observed throughout the VIMS observing period, even after equinox, in a latitude band between 30°S and 60°S. During the whole period of observation, only a dozen clouds were observed closer to the equator, though they were slightly more frequent as equinox approached.We also investigated the distribution of clouds with longitude. We found that southern polar clouds, before disappearing in mid-2008, were systematically concentrated in the leading hemisphere of Titan, in particular above and to the east of Ontario Lacus, the largest reservoir of hydrocarbons in the area. Clouds are also non-homogeneously distributed with longitude at southern mid-latitudes. The n = 2-mode wave pattern of the distribution, observed since 2003 by Earth-based telescopes and confirmed by our Cassini observations, may be attributed to Saturn’s tides.Although the latitudinal distribution of clouds is now relatively well reproduced and understood by the GCMs, the non-homogeneous longitudinal distributions and the evolution of the cloud coverage with seasons still need investigation. If the observation of a few single clouds at the tropics and at northern mid-latitudes late in winter and at the start of spring cannot be further interpreted for the moment, the obvious shutdown of the cloud activity at Titan’s poles provides clear signs of the onset of the general circulation turnover that is expected to accompany the beginning of Titan’s northern spring. According to our GCM, the persistence of clouds at certain latitudes rather suggests a ‘sudden’ shift in near future of the meteorology into the more illuminated hemisphere. Finally, the observed seasonal change in cloud activity occurred with a significant time lag that is not predicted by our model. This may be due to an overall methane humidity at Titan’s surface higher than previously expected.     

Global stratigraphy of Venus: analysis of a random

The age relations between 36 impact craters with dark paraboloids and other geologic units and structures at these localities have been studied through photogeologic analysis of Magellan SAR images of the surface of Venus. Geologic settings in all 36 sites, about 1000 × 1000 km each, could be characterized using only 10 different terrain units and six types of structures. These units and structures form a major stratigraphic and geologic sequence (from oldest to youngest): 1) tessera terrain; 2) densely fractured terrains associated with coronae and in the form of remnants among plains; 3) fractured and ridged plains and ridge belts; 4) plains with wrinkle ridges; 5) ridges associated with coronae annulae and ridges of arachnoid annulae which are contemporary with wrinkle ridges of the ridged plains; 6) smooth and lobate plains; 7) fractures of coronae annulae, and fractures not related to coronae annulae, which disrupt ridged and smooth plains; 8) rift-associated fractures; 9) craters with associated dark paraboloids, which represent the youngest 10% of the Venus impact crater population (Campbellet al., 1992), and are on top of all volcanic and tectonic units except the youngest episodes of rift-associated fracturing and volcanism; surficial streaks and patches are approximately contemporary with dark-paraboloid craters.Mapping of such units and structures in 36 randomly distributed large regions (each 10⁶ km²) shows evidence for a distinctive regional and global stratigraphic and geologic sequence. On the basis of this sequence we have developed a model that illustrates several major themes in the history of Venus. Most of the history of Venus (that of its first 80% or so) is not preserved in the surface geomorphological record. The major deformation associated with tessera formation in the period sometime between 0.5–1.0 b.y. ago (Ivanov and Basilevsky, 1993) is the earliest event detected. In the terminal stages of tessera formation, extensive parallel linear graben swarms representing a change in the style of deformation from shortening to extension were formed on the tessera and on some volcanic plains that were emplaced just after (and perhaps also during the latter stages of the major compressional phase of tessera emplacement. Our stratigraphic analyses suggest that following tessera formation, extensive volcanic flooding resurfaced at least 85% of the planet in the form of the presently-ridged and fractured plains. Several lines of evidence favor a high flux in the post-tessera period but we have no independent evidence for the absolute duration of ridged plains emplacement. During this time, the net state of stress in the lithosphere apparently changed from extensional to compressional, first in the form of extensive ridge belt development, followed by the formation of extensive wrinkle ridges on the flow units. Subsequently, there occurred local emplacement of smooth and lobate plains units which are presently essentially undeformed. The major events in the latest 10% of the presently preserved history of Venus (less than 50 m.y. ago) are continued rifting and some associated volcanism, and the redistribution of eolian material largely derived from impact crater deposits.Detailed geologic mapping and stratigraphic synthesis are necessary to test this sequence and to address many of the outstanding problems raised by this analysis. For example, we are uncertain whether this stratigraphic sequence corresponds to geologic events which were generally synchronous in all the sites and all around the planet, or whether the sequence is simply a typical sequence of events which occurred in different places at different times. In addition, it is currently unknown whether the present state represents a normal consequence of the general thermal evolution of Venus (and is thus representative of the level of geological activity predicted for the future), or if Venus, has been characterized by a sequence of periodic global changes in the composition and thermal state of its crust and upper mantle (in which case, Venus could in the future return to levels of deformation and resurfacing typical of the period of tessera formation).     

Millochau crater, Mars: Infilling and erosion of an ancient highland impact crater

The geology and stratigraphy of Millochau crater (21.4° S, 275° W), located in the highlands of Tyrrhena Terra, Mars, are documented through geomorphic analyses and geologic mapping. Crater size-frequency distributions and superposition relationships are used to constrain relative ages of geologic units and determine the timing and duration of the geologic processes that modified Millochau rim materials and emplaced deposits on Millochau's floor. Crater size-frequency distributions show a Middle Noachian age for rim materials and Middle Noachian to Early Hesperian ages for most of the interior deposits. Valley networks and gullies incised within Millochau's rim materials and interior wall, respectively, indicate fluvial activity was an important erosional process. Millochau contains an interior plateau, offset northeast of Millochau's center, which rises up to 400 m above the surrounding crater floor and slopes downward to the south and west. Layers exposed along the northern and eastern scarp boundaries of the plateau are tens to hundreds of meters thick and laterally continuous in MOC images. These layers suggest most materials within Millochau were emplaced by sedimentary processes (e.g., fluvial or eolian), with the potential for lacustrine deposition in shallow transient bodies of water and contributions of volcanic airfall. Mass wasting may have also contributed significant quantities of material to Millochau's interior, especially to the deposits surrounding the plateau. Superposition relationships combined with impact crater statistics indicate that most deposition and erosion of Millochau's interior deposits is ancient, which implies that fluvial activity in this part of Tyrrhena Terra is much older than in the eastern Hellas region. Eolian processes mobilized sediment to form complicated patterns of long- and short-wavelength dunes, whose emplacement is controlled by local topography. These deposits are some of the youngest within Millochau (Amazonian) and eolian modification may be ongoing.     

Anomalous radar backscatter from Titan’s surface?

Since Cassini arrived at Saturn in 2004, its moon Titan has been thoroughly mapped by the RADAR instrument at 2-cm wavelength, in both active and passive modes. Some regions on Titan, including Xanadu and various bright hummocky bright terrains, contain surfaces that are among the most radar-bright encountered in the Solar System. This high brightness has been generally attributed to volume scattering processes in the inhomogeneous, low-loss medium expected for a cold, icy satellite surface. We can test this assumption now that the emissivity has been obtained from the concurrent radiometric measurements for nearly all the surface, with unprecedented accuracy (Janssen et al., and the Cassini RADAR Team [2009]. Icarus 200, 222-239). Kirchhoff’s law of thermal radiation relates the radar and radiometric properties in a way that has never been fully exploited. In this paper we examine here how this law may be applied in this case to better understand the nature of Titan’s radar-bright regions. We develop a quantitative model that, when compared to the observational data, allows us to conclude that either the reflective characteristics of the putative volume scattering subsurface must be highly constrained, or, more likely, organized structure on or in the surface is present that enhances the backscatter.     

The spatial distribution of volatiles in the Martian hydrolithosphere

In order to quantify the spatial distribution of volatiles on Mars, 2600 fluidized ejecta craters have been systematically measured, classified and mapped over the planet Mars, using 1 : 2 M scale USGS photomosaics. The latitudinal distribution of ejecta craters reveals that flower ejecta deposits (Type 1), together with low mobility ejecta, are frequently observed in the equatorial region and on ridged plains. Rampart craters (Type 2), with high mobility ejecta, occur at mid latitudes and exhibit a spatial relationship with polygonal patterns and pseudocrater areas. The increase of ejecta mobility with latitude attests for a concentration of volatiles at high latitudes. Statistical analysis shows that cratered uplands and ridged plains contain less volatile material near the surface than the underlying materials. In Chryse Planitia and Utopia Planitia the statistical study and the spatial relationships between polygonally fractured patterns, pseudocraters and the great number of high mobility ejecta deposits suggest the presence of a water-rich alluvial deposit close to the surface near the mouth of Chryse and Elysium channels. This result explains, on a more quantitative basis, the idea that fractured patterns were preferentially developed in a volatile-rich sedimentary deposits. The behaviour of volatiles, at 41 S, 257 W near Reull Vallis, exhibits a strong anomaly, with the presence of an abnormally volatile rich layer close to the surface.     

Crater modification and geologic activity in Enceladus' heavily cratered plains: Evidence from the impact crater distribution

Although we can observe current activity on Saturn's satellite Enceladus with Cassini, insight into past activity is best achieved (for now) through studying the impact crater distributions. Furthermore, approximation of terrain ages can only be attained through calculations using crater densities and estimations of impact rates in the saturnian system. Here we focus on what the impact crater distribution in Enceladus' heavily cratered plains can tell us about Enceladus' geologic history. We use Cassini ISS images to count craters in the heavily cratered plains on Enceladus, along with Rhea, Dione, Tethys and Mimas as references, to develop and compare their size-frequency distributions. Comparisons of our counts show that Enceladus' cratered plains distribution is unique in that it appears to have a relative deficiency of craters for diameters ?2 km and ?6 km compared to the other satellites' heavily cratered plains. Our data also indicates that the impact crater density within the cratered plains changes with latitude. Specifically, both the north and south mid-latitude regions have approximately three times higher density than the equatorial region. We hypothesize that the “missing” small and large craters in Enceladus' cratered plains is due to a combination of viscous relaxation of the larger craters, and burial of the relaxed large craters and small craters by south polar plume and possibly E-ring material. We also conclude that the spatial density distribution is not consistent with recent polar wander.     

Geology of the Selk crater region on Titan from Cassini VIMS observations

Observations of Titan obtained by the Cassini Visual and Infrared Mapping Spectrometer (VIMS) have revealed Selk crater, a geologically young, bright-rimmed, impact crater located ∼800 km north-northwest of the Huygens landing site. The crater rim-crest diameter is ∼90 km; its floor diameter is ∼60 km. A central pit/peak, 20-30 km in diameter, is seen; the ratio of the size of this feature to the crater diameter is consistent with similarly sized craters on Ganymede and Callisto, all of which are dome craters. The VIMS data, unfortunately, are not of sufficient resolution to detect such a dome. The inner rim of Selk crater is fluted, probably by eolian erosion, while the outer flank and presumed ejecta blanket appear dissected by drainages (particularly to the east), likely the result of fluvial erosion. Terracing is observed on the northern and western walls of Selk crater within a 10-15 km wide terrace zone identified in VIMS data; the terrace zone is bright in SAR data, consistent with it being a rough surface. The terrace zone is slightly wider than those observed on Ganymede and Callisto and may reflect differences in thermal structure and/or composition of the lithosphere. The polygonal appearance of the crater likely results from two preexisting planes of weakness (oriented at azimuths of 21° and 122° east of north). A unit of generally bright terrain that exhibits similar infrared-color variation and contrast to Selk crater extends east-southeast from the crater several hundred kilometers. We informally refer to this terrain as the Selk “bench.” Both Selk and the bench are surrounded by the infrared-dark Belet dune field. Hypotheses for the genesis of the optically bright terrain of the bench include: wind shadowing in the lee of Selk crater preventing the encroachment of dunes, impact-induced cryovolcanism, flow of a fluidized-ejecta blanket (similar to the bright crater outflows observed on Venus), and erosion of a streamlined upland formed in the lee of Selk crater by fluid flow. Vestigial circular outlines in this feature just east of Selk’s ejecta blanket suggest that this might be a remnant of an ancient, cratered crust. Evidently the southern margin of the feature has sufficient relief to prevent the encroachment of dunes from the Belet dune field. We conclude that this feature either represents a relatively high-viscosity, fluidized-ejecta flow (a class intermediate to ejecta blankets and long venusian-style ejecta flows) or a streamlined upland remnant that formed downstream from the crater by erosive fluid flow from the west-northwest.     

Martian cratering 8: Isochron refinement and the chronology of Mars

This paper reviews and refines the technique of dating martian surfaces by using impact-crater isochrons (defined as size distributions of impact craters on undisturbed martian surfaces of specified ages). In the 1970s, this system identified not only abundant ancient martian volcanic surfaces, but also extensive lava plains with ages of a few 10⁸ y-old; this dating was initially controversial but confirmed in the 1980s and 90s by martian meteorites. The present update utilizes updated estimates of the Mars/Moon cratering ratio (the most important calibration factor), improves treatment of gravity and impact velocity scaling effects, combines aspects of the crater size distribution data from earlier work by both Neukum and Hartmann, and for the first time applies a correction for loss of small meteoroids in the martian atmosphere from Popova et al. (2003, Meteorit. Planet. Sci. 38, 905-925). The updated isochrons are not radically different from the previous “2002 iteration” but fit observed data better and give somewhat older model ages for features dated from small craters (diameter D<100 m). Crater counts from young lava flows in various areas give good fits to the new isochrons over as much as 3 orders of magnitude in D, confirming the general isochron shape and giving crater retention ages in the range of some 10⁶ to some 10⁸ y, interpreted as lava flow ages. More complex, older units are also discussed. Uncertainties are greatest if only small craters (D?100 m) are used. Suggestions by other workers of gross uncertainties, due to local secondary craters and deposition/exhumation, are discussed; they do not refute our conclusions of significant volcanic, fluvial, and other geologic activity in the last few percent of martian geologic time or the importance of cratering as a tool for studying processes such as exhumation. Indeed, crater count data suggest certain very recent episodes of deposition, exhumation, and ice flow, possibly associated with obliquity cycles of ∼¹⁰⁷ y timescale. Evidence from ancient surfaces suggests higher rates of volcanism, fluvial activity, glaciation, and other processes in Noachian/Hesperian time than in Amazonian time.     

Intra-crater glacial processes in central Utopia Planitia, Mars

We describe and interpret a series of previously unidentified glacial-like lobes (34-43°N; 107-125°E) that were discovered as part of a survey of large (D > 5 km) impact craters in Utopia Planitia, one of the great northern plains of Mars. The lobes have characteristics that are consistent with a glacial origin. Evidence includes curvilinearity of form, lineations and ridges, and surface textures that are thought to form by the sublimation of near-surface volatiles. The lobes display morphologies that range from wedge-shaped to near-circular to elongate. The flow directions are towards the northern walls in the case of craters with large single lobes, and in all directions in the case of the largest (D > 30 km) craters. Concentric crater fill is also interspersed within craters of our study region, with such craters having much higher filling rates than those with flow lobes. We suggest that the impact crater population in south-west Utopia Planitia demonstrates a spectrum of glacial modifications, from low levels of filling in the case of craters with elongate lobes at one extreme, to concentric crater fill in highly-filled craters at the other.     

The rayed crater Zunil and interpretations of small impact craters on Mars

A 10-km diameter crater named Zunil in the Cerberus Plains of Mars created ∼¹⁰⁷ secondary craters 10 to 200 m in diameter. Many of these secondary craters are concentrated in radial streaks that extend up to 1600 km from the primary crater, identical to lunar rays. Most of the larger Zunil secondaries are distinctive in both visible and thermal infrared imaging. MOC images of the secondary craters show sharp rims and bright ejecta and rays, but the craters are shallow and often noncircular, as expected for relatively low-velocity impacts. About 80% of the impact craters superimposed over the youngest surfaces in the Cerberus Plains, such as Athabasca Valles, have the distinctive characteristics of Zunil secondaries. We have not identified any other large (?10 km diameter) impact crater on Mars with such distinctive rays of young secondary craters, so the age of the crater may be less than a few Ma. Zunil formed in the apparently youngest (least cratered) large-scale lava plains on Mars, and may be an excellent example of how spallation of a competent surface layer can produce high-velocity ejecta (Melosh, 1984, Impact ejection, spallation, and the origin of meteorites, Icarus 59, 234-260). It could be the source crater for some of the basaltic shergottites, consistent with their crystallization and ejection ages, composition, and the fact that Zunil produced abundant high-velocity ejecta fragments. A 3D hydrodynamic simulation of the impact event produced 10¹⁰ rock fragments ?10 cm diameter, leading to up to 10⁹ secondary craters ?10 m diameter. Nearly all of the simulated secondary craters larger than 50 m are within 800 km of the impact site but the more abundant smaller (10-50 m) craters extend out to 3500 km. If Zunil is representative of large impact events on Mars, then secondaries should be more abundant than primaries at diameters a factor of ∼1000 smaller than that of the largest primary crater that contributed secondaries. As a result, most small craters on Mars could be secondaries. Depth/diameter ratios of 1300 small craters (10-500 m diameter) in Isidis Planitia and Gusev crater have a mean value of 0.08; the freshest of these craters give a ratio of 0.11, identical to that of fresh secondary craters on the Moon (Pike and Wilhelms, 1978, Secondary-impact craters on the Moon: topographic form and geologic process, Lunar Planet. Sci. IX, 907-909) and significantly less than the value of ∼0.2 or more expected for fresh primary craters of this size range. Several observations suggest that the production functions of Hartmann and Neukum (2001, Cratering chronology and the evolution of Mars, Space Sci. Rev. 96, 165-194) predict too many primary craters smaller than a few hundred meters in diameter. Fewer small, high-velocity impacts may explain why there appears to be little impact regolith over Amazonian terrains. Martian terrains dated by small craters could be older than reported in recent publications.