Vastitas Borealis Formation Research Articles

SHARAD radar sounding of the Vastitas Borealis Formation in Amazonis Planitia

Amazonis Planitia has undergone alternating episodes of sedimentary and volcanic infilling, forming an interleaved sequence with an upper surface that is very smooth at the kilometer scale. Earlier work interprets the near‐surface materials as either young, rough lava flows or ice‐rich sediment layers, overlying a basement comprising the Vastitas Borealis Formation and earlier Hesperian plains. Sounding radar profiles across Amazonis Planitia from the Shallow Radar (SHARAD) instrument on the Mars Reconnaissance Orbiter reveal a subsurface dielectric interface that increases in depth toward the north along most orbital tracks. The maximum depth of detection is 100–170 m, depending upon the real dielectric permittivity of the materials, but the interface may persist at greater depth to the north if the reflected energy is attenuated below the SHARAD noise floor. The dielectric horizon likely marks the boundary between sedimentary material of the Vastitas Borealis Formation and underlying Hesperian volcanic plains. The SHARAD‐detected interface follows the surface topography across at least one of the large wrinkle ridges in north central Amazonis Planitia. This conformality suggests that Vastitas Borealis sediments, at least in this region, were emplaced prior to compressional tectonic deformation. The change in radar echo strength with time delay is consistent with a loss tangent of 0.005–0.012 for the column of material between the surface and the reflector. These values are consistent with dry, moderate‐density sediments or the lower end of the range of values measured for basalts. While a component of distributed ice in a higher‐loss matrix cannot be ruled out, we do not find evidence for a dielectric horizon within the Vastitas Borealis Formation that might suggest an abrupt change from an upper dry layer to an ice‐rich lower deposit.

Two geologic systems providing terrestrial analogues for the exploration of sulfate deposits on Mars: Initial spectral characterization

We present the Messinian evaporite suite (Mediterranean region) and the Solfatara hydrothermal system (Phlegraean Fields volcanic province, Italy), discuss their implications for understanding the origin of sulfates on Mars and show preliminary sets of VNIR laboratory and in situ reflectance spectra of rocks from these geologic systems. The choice was based on a number of evidence relative to Mars: (1) the chemistry of the Martian sulfates, suggesting fluid interactions with possibly alkali-basaltic rocks and/or regolith; (2) close range evidence of sulfates within sedimentary formations on Mars; (3) sulfate spectral signatures associated to large-scale layered patterns interpreted as thick depositional systems on Mars. The Messinian evaporites comprise three units: primary shallow-water sulfates (primary lower gypsum: PLG), shallow- to deep-water mixed sulfates and clastic terrigenous deposits (resedimented lower gypsum: RLG), and shallow-water associations of primary sulfates and clastic fluvio-deltaic deposits (upper evaporites: UE). The onset of the Messinian evaporites records the transition to negative hydrologic budget conditions associated with the Messinian Salinity Crisis, which affected the entire Mediterranean basin and lasted about 640 kyr. The Solfatara is a still evolving hydrothermal system that provides epithermal deposits precipitated from the interaction of fluids and trachybasaltic to phonolitic rocks. Thermal waters include alkali-chloride, alkali-carbonate and alkali-sulfate endmembers. The wide spectrum of sedimentary gypsum facies within the Messinian formation includes some of the depositional environments hitherto identified on Mars and others not found on Mars. The PLG unit includes facies associations correlated over long distances, that could be a possible analog of the stratified rock units exposed from Arabia Terra at least as far as Valles Marineris. The facies cycles within the UE unit can be compared to the sequences of strata observed in craters such as Holden and Eberswalden. The UE unit records paleoenvironmental changes which are ultimately controlled by terrestrial climatic variations. They can be considered as a reliable climatic proxy and may be useful for the reconstruction of climatic events on Mars. The intermediate Messinian RLG unit has not, at present, a well-defined depositional counterpart on Mars, although there are some similarities with the northern lowlands and Vastitas Borealis Formation. The dramatic variation of hydrologic budget conditions at the onset of the Messinian evaporites may provide criteria for the interpretation of similar variations on Mars. The volcanic rocks at the Solfatara bear some similarities with the “alkaline magmatic province” observed at the Gusev crater on Mars, and the assemblages of hydrothermal phases resulting from the Solfatara's parent rocks could be analogues for processes involving Gusev-type rocks. The Messinian sulfates have a prevalent Ca-sulfatic composition and wide textural variability. Preliminary laboratory reflectance spectra of rock samples in the VNIR region reveal the signature of sulfates and mixtures of several Fe-bearing phases. At the Solfatara, in situ reflectance measurements of epithermal minerals close to active fumaroles showed the presence of Fe-bearing sulfates, hematite, Al- and K-sulfates and abundant amorphous fraction. XRD analysis supported this interpretation. The range of depositional facies observed in the Messinian units and the variety of minerals detected in the Solfatara will be useful for the interpretation of close range data of Mars. The spectral characterization at various scales of the Messinian sedimentary facies and the Solfatara hydrothermal minerals will both help in the exploration of Mars from orbit and with close range inspection.

Modification of impact craters in the northern plains of Mars: Implications for Amazonian climate history

Abstract— We measured the depth, wall steepness, and ejecta roughness and surveyed the wall and floor morphology of all craters 10–25 km in diameter within the typical Vastitas Borealis Formation in the northern lowlands of Mars north of 52°N. Two of the 130 craters have unusually rough ejecta; they are deep, have steep walls, and are apparently the youngest in the population. Icy mantles filling the local subkilometer‐scale topographic lows is the main contribution to ejecta smoothing, which occurs at a time scale on the order of tens of Myr. Wall degradation and crater shallowing generally occur at longer time scales, comparable to the duration of the Amazonian period. Many craters are shallow due to filling of the crater with specific ice‐rich material of uncertain origin. We use our collected data to infer the nature of the past climate back through the Amazonian, a period prior to ˜10–20 Myr ago, when orbital parameter solutions are chaotic and one must rely on geological data to infer climate conditions. We conclude that moderately high obliquity and wide obliquity variations were probable during the last 40–160 Myr. We tentatively conclude that high obliquity peaks (>40–45°) may have occurred episodically through the last 210–430 Myr. A sharp step in the frequency distribution of wall steepness at 20° may indicate a geologically long period prior to that time where obliquity never exceeded 40–45°.

Origin and characteristics of the Mars north polar basal unit and implications for polar geologic history

Building upon previous studies, we have used Mars Orbiter Camera and Mars Orbiter Laser Altimeter data to characterize in detail the newly discovered north polar basal unit. Lying stratigraphically between the polar layered deposits, from which it is likely separated by an unconformity, and the Vastitas Borealis Formation, this unit has introduced new complexity into north polar stratigraphy and has important implications for polar history. Exposures of the basal unit in Olympia Planitia and Chasma Boreale reveal relatively dark layers which exhibit differential erosion. Eroded primarily by wind, the basal unit may be the major if not sole source for the north polar dunes and ergs and has contributed material to the lower polar cap layers. We investigate four possible origins for the basal unit (outflow channel/oceanic deposits, basal ice, paleopolar deposits, and eolian deposits). The patchy layering within the unit, its likely sandy grain size, and presence only in the north polar basin suggest that it is primarily an eolian deposit, supporting Byrne and Murray's 2002 earlier conclusion. This implies that at some time during the Early to Late Amazonian, migrating sand was mixed with water ice, forming a relatively dark, sandy deposit. During this time, either no classic polar layered deposits were forming or smaller caps were growing and shrinking, possibly adding material to the basal unit.

Amplitude of heat flow variations on Mars from possible shoreline topography

Analyses of the effective elastic thickness of the Martian lithosphere have been previously used to calculate surface heat flow on Mars at different places and times. In this work, I use elevation differences in a putative Late Hesperian shoreline, named Deuteronilus shoreline, and the relation between thermal state and buoyancy of the lithosphere, in order to estimate the amplitude of the variations of surface heat flow on Mars, probably related to the time in which this feature was formed. The results suggest that, if the Deuteronilus shoreline is a true paleo‐equipotential surface, the relative amplitude of surface heat flow variations on the shoreline regions in the Late Hesperian were less than present‐day ones in terrestrial continental areas. The results are also roughly valid for the outer contact of the Late Hesperian Vastitas Borealis Formation, if this contact is related to the limits of an ancient ocean. These results could imply that large areas of the Martian lithosphere have been tectonothermally stable since at least that time.

Syrtis Major and Isidis Basin contact: Morphological and topographic characteristics of Syrtis Major lava flows and material of the Vastitas Borealis Formation

The floor of Isidis Basin is covered by materials of the Vastitas Borealis Formation (VBF) that appear to be emplaced essentially as a single unit. Along its western boundary, Isidis Basin is in contact with volcanic flows from Syrtis Major Planum. The contact between the Isidis unit and volcanic flows from Syrtis Major is sharp to gradational and in places is characterized by a high (∼500 m) scarp or by a network of faults that separate pieces of lava plains off the main plateau of Syrtis. Clusters of knobs and mesas, sometimes arranged in flow‐like features, are also typical features of the transition zone. Several important characteristics of the transition from Syrtis Major to Isidis Basin are documented. (1) The small‐scale surface texture seen in MOC images appears to be the same for both the Syrtis lava plateau and the knobs and mesas that characterize the transition. (2) There is strong evidence for the breakup of the coherent surface of Syrtis Major where it is in contact with materials in Isidis Basin. (3) The plateau breakup (the knobby terrain) occurs basinward after the major break of slope of Syrtis Major where it enters the Isidis Basin. (4) There is no evidence for plateau breakup anywhere up on the slopes of Syrtis Major Planum. (5) The lavas of Syrtis remain morphologically intact where they are in contact with other units, such as the Noachian cratered terrain or where lava flows are stacked within Syrtis Major itself. These characteristic features of the transition zone from Syrtis to Isidis are readily explained if the zone of plateau breakup consists of relatively young lava flows that have been superimposed onto the surface of a volatile‐rich substratum, such as the interior unit of Isidis Basin (the Vastitas Borealis Formation). Thus simple superposition of volcanic materials on top of volatile‐bearing sediments can explain the key geological and topographic aspects of the transition zone from Syrtis Major to Isidis Basin. On the basis of our findings, we outline the following scenario for the evolution of this region. In the Early Hesperian, volcanic plains are emplaced in Syrtis Major (the lower part of the Syrtis Major Formation), and wrinkle ridges deform their surfaces soon thereafter. Concurrently, volcanic plains are emplaced on the floor of the Isidis Basin, and wrinkle ridges deform their surfaces soon thereafter. The apparent simultaneity of these units may mean that Syrtis Major was the source of many of the flows in the Isidis Basin. In the early part of the Upper Hesperian, subsequent to the formation of most of the wrinkle ridges, the Vastitas Borealis Formation was emplaced in the Isidis Basin and elsewhere in the northern lowlands. Following the emplacement of the Vastitas Borealis Formation, the upper part of the Syrtis Major Formation was emplaced, erupting from the eastern margins of Syrtis Major Planum and flowing down into the westernmost part of the Isidis Basin on top of the recently emplaced Vastitas Borealis Formation. Modification of the superposed lavas by degradation and evolution of the VBF formed the scarps and unusual morphology of the marginal areas. We found no compelling evidence for massive or sudden erosion from Syrtis Major to produce the plains currently on the surface of the floor of the Isidis Basin (the Vastitas Borealis Formation).

Origin of the polygons and the thickness of Vastitas Borealis Formation in Western Utopia Planitia on Mars

Many middle scale (about 100m cracking space) polygons are located between the area of the giant polygons and the thumbprint terrains in Western Utopia Planitia. Most of the patterns are located on Vastitas Borealis Formation. The polygonal patterns have frequently been associated with collapsed features, such as a covered trough like features(the width of the underground crack is approximately 400–500 m). The thickness of the surface materials (6–12 m), where polygonal patterns occur, was calculated based on the space of cracking relationship. Polygonal networks with strongly oriented features and parallel to the contour line are likely caused by structural (geological) origin as well as bending forces from the center of the basin.

Oceans on Mars: An assessment of the observational evidence and possible fate

If the large Late Hesperian outflow channels were eroded by extensive floods, as appears likely, then large bodies of water must have once occupied the northern plains during that period. Previous estimates of the sizes of bodies of water in the northern lowlands range up to 3 × 108 km3. Several contacts have been previously mapped around the edges of the northern plains and interpreted to be shorelines remaining from these former standing bodies of water. We examine the elevations and geologic relations along these contacts in detail and find little support for their interpretation as shorelines. Some contacts are clearly of volcanic origin, and all have significant variations in elevation. Better support for the former presence of water over large parts of the northern plains is provided by the Vastitas Borealis Formation (VBF). Most of the post‐Noachian fill within the northern basin is ridged plains of Lower Hesperian age, interpreted to be volcanic in origin. Overlying the ridged plains is the VBF, a thin veneer of material of Upper Hesperian age. The VBF may have been deposited from large floods. Support for this interpretation is the similarity in age between the outflow channels and the VBF, the presence of the VBF at the lower ends of the outflow channels, and identification of numerous features in the outcrop areas of the VBF that are suggestive of basal melting of an ice sheet. To cover all the area over which the VBF is exposed would require ∼2.3 × 107 km3 of water. Spread over the entire surface of Mars, this volume is equal to a global layer (global equivalent layer, or GEL) ∼156 m deep. We find no support for the larger estimates of ocean volumes that range up to 3 × 108 km3 and which imply comparable amounts of water per unit area as are currently present on the surface of the Earth. Under present climatic conditions on Mars an ocean would freeze in a geologically short time period (∼104 years), then would sublimate away at rates strongly dependent on the presence or absence of debris on the ice surface. The present VBF is interpreted as a sublimation residue from the ponded outflow channel effluents. The fate of a volume of water thought to have been emplaced by the outflow channels (∼2.3 × 107 km3) is largely accounted for by the presence of other existing reservoirs on the planet. An approximately 20–30 m GEL of water is estimated to be in the present polar caps, and a 50 m GEL may have escaped to space since the Hesperian, leaving ∼80 m GEL unaccounted for. This amount may be partly trapped in other volatile‐rich deposits on the surface, and a significant amount could have reentered the groundwater system by south polar basal melting and been progressively cold‐trapped at the base of a growing cryosphere. On the basis of our assessment of the Hesperian‐aged deposits, we predict that testing of the Clifford and Parker [2001] hypothesis that a Noachian‐aged ocean covered up to one third of the surface of Mars will be made very difficult by the enhanced degradation rates in the Noachian and subsequent geological events in the northern lowlands.

Resurfacing history of the northern plains of Mars based on geologic mapping of Mars Global Surveyor data

Geologic mapping of the northern plains of Mars, based on Mars Orbiter Laser Altimeter topography and Viking and Mars Orbiter Camera images, reveals new insights into geologic processes and events in this region during the Hesperian and Amazonian Periods. We propose four successive stages of lowland resurfacing likely related to the activity of near‐surface volatiles commencing at the highland‐lowland boundary (HLB) and progressing to lower topographic levels as follows (highest elevations indicated): Stage 1, upper boundary plains, Early Hesperian, <−2.0 to −2.9 km; Stage 2, lower boundary plains and outflow channel dissection, Late Hesperian, <−2.7 to −4.0 km; Stage 3, Vastitas Borealis Formation (VBF) surface, Late Hesperian to Early Amazonian, <−3.1 to −4.1 km; and Stage 4, local chaos zones, Early Amazonian, <−3.8 to −5.0 km. At Acidalia Mensa, Stage 2 and 3 levels may be lower (<−4.4 and −4.8 km, respectively). Contractional ridges form the dominant structure in the plains and developed from near the end of the Early Hesperian to the Early Amazonian. Geomorphic evidence for a northern‐plains‐filling ocean during Stage 2 is absent because one did not form or its evidence was destroyed by Stage 3 resurfacing. Remnants of possible Amazonian dust mantles occur on top of the VBF. The north polar layered deposits appear to be made up of an up to kilometer‐thick lower sequence of sandy layers Early to Middle Amazonian in age overlain by Late Amazonian ice‐rich dust layers; both units appear to have outliers, suggesting that they once were more extensive.

Fate of outflow channel effluents in the northern lowlands of Mars: The Vastitas Borealis Formation as a sublimation residue from frozen ponded bodies of water

We analyze the fate of the Hesperian‐aged outflow channel effluents emplaced into the northern lowlands of Mars. We have modeled the evolution of these effluents under the assumption that they were emplaced under a range of atmospheric conditions comparable to those of today and thought to have prevailed in the Hesperian. Under these conditions we find that the evolution of the sediment‐loaded water after it left the channels includes three phases. Phase 1: Emplacement and intensive cooling: Violent emplacement of water followed by a short period of intensive evaporation from the surface and near‐surface layer, and intensive convection. During this phase the water maintained and redistributed its large suspended sediment load. Water vapor strongly influenced the climate, at least for a geologically short time. When the temperature of the water reached the temperature of the maximum density or the freezing point, boiling and intensive convection ceased and the water deposited the sediments. Phase 2: Freezing solid: Geologically rapid freezing of the water body accompanied by weak convective water movement occurred over a period of the order of ∼104years. Phase 3: Sublimation and loss: This period involved sublimation of the ice and lasted longer than the freezing phase. The rate and latitudinal dependence of the sublimation, as well as the location of water vapor condensation, crucially depend on the planetary obliquity, climate, and sedimentary veneering of the ice. Phase 3 would have been very short geologically (∼105–106years) if an insulating sedimentary layer did not build up rapidly. If such an insulating layer did form rapidly, sublimation could cease and residual ice hundreds of meters thick could remain below the surface today. The northern lowlands of Mars are largely covered by the Late Hesperian‐aged Vastitas Borealis Formation (VBF), which has sharp boundaries clearly seen in the map of kilometer‐scale roughness, and a distinctive kilometer‐scale roughness signature. We examine detrended topography data and find evidence that the VBF is underlain at very shallow depths by an Early Hesperian ridged volcanic plains substrate, rather than frozen water deposits remaining from the outflow channel events. Analysis of the VBF roughness characteristics suggests at least 100 m thickness of sediments on top of the underlying volcanic ridged plains. The total inferred volume of the VBF material approximately corresponds to the volume of material removed from the outflow channels. These results support a model in which the Vastitas Borealis Formation predominantly represents the sedimentary residue remaining after the sediment‐laden water effluents of the outflow channels ponded in the northern lowlands and rapidly froze solid and sublimed. The distinctive facies of the VBF are interpreted to be related to the freezing, sublimation, and residue of the outflow channel effluents. Ridged facies are dominantly marginal and are interpreted to be related to ice sheet lateral retreat. Knobby facies occur throughout and are interpreted to be due to a variety of causes, including dewatering, de‐icing and brine‐related processes, and kame‐like ice residues. Grooved facies are interpreted to represent postsublimation uplift and fracturing due to load removal. The mottled facies may represent bright ejecta whose albedo is related to buried sedimentary or evaporitic residues. Presently, the most likely sites to find frozen remnants of the Hesperian ocean are below the floors of stealth craters underlying the VBF. On the basis of these findings and interpretations, we make predictions for the fate of the outflow channel water vapor and the nature of the Martian hydrologic cycle during the Hesperian. If Noachian‐aged oceans existed, conditions at that time would be similar to those described for a Hesperian ocean as soon as a global cryosphere was formed. If any residual deposits remain in the northern lowlands from a proposed Noachian ocean, they would lie below both the VBF and the underlying sequence of Hesperian ridged plains.

Utopia Basin, Mars: Characterization of topography and morphology and assessment of the origin and evolution of basin internal structure

Recently obtained Mars Orbiter Laser Altimeter (MOLA) topography has permitted a new assessment of the morphology, structure, and history of Utopia Planitia. The new topographic data convincingly demonstrate that the Utopia region is an impact basin, as originally proposed byMcGill[1989], whose major topographic expression is a circular, 1–3 km deep depression ∼3200 km in diameter. Utopia Basin is the largest easily recognizable impact structure in the northern hemisphere of Mars and is the only portion of the northern lowlands that exhibits a distinct large‐scale impact signature; its presence there and its ancient age verify that a significant part of the northern lowlands dates back to the Noachian period. Using slope maps derived from gridded topography, we have mapped a series of linear slope anomalies within the basin and have analyzed their origin. These features can be classified into two groups: one oriented dominantly radial to the basin and the other oriented dominantly circumferential. The circumferential features are tens of kilometers wide and hundreds of kilometers long and appear to remain relatively constant with respect to elevation. We explore end‐member hypotheses (1) that the features formed by the direct action of water/ice or (2) that they are tectonic wrinkle ridges. We interpret the majority of these features to be due to tectonic deformation (primarily wrinkle‐ridge formation), modified by later sedimentation associated with the emplacement of the outflow channels and the formation of the Vastitas Borealis Formation. On the basis of the characteristics of these features and their association with polygonal terrain, thumbprint terrain, smoothness of units, and evidence of sedimentary deposits, we find that our results support earlier hypotheses proposing that a significant volume of water occupied Utopia Basin during the late Hesperian period of outflow channel formation. Utopia Basin differs significantly in depth and degradation state from the deeper and fresher‐appearing Hellas Basin, although the two are both Noachian in age. A significant part of this difference seems reasonably interpreted to be the result of the extensive sedimentary and volcanic infilling history of the northern lowlands.

Kilometer‐scale roughness of Mars: Results from MOLA data analysis

The Mars Orbiter Laser Altimeter (MOLA) data are used to characterize the kilometer‐scale surface roughness of Mars. The median absolute value of the differential slope at a given baseline is proposed as a data‐derived measure of the surface roughness at this scale. Study of the scale dependence of roughness for the smoothest terrains gives an independent estimate of 20 cm for MOLA ranging accuracy. The baseline lengths from 0.6 to ∼20 km are used, and kilometer‐scale roughness is mapped for the entire surface. The maps show that different geological units have distinctive roughness characteristics. Scale dependence of roughness for a number of geological units is presented and discussed. The southern polar cap is rougher than the northern at kilometer and subkilometer scale, which suggests differences in the sublimation/condensation balance. The Vastitas Borealis Formation has a distinctive 3‐km‐scale background surface topography, which suggests a nonvolcanic origin for its upper layer. Young volcanic plains in Amazonis Planitia and the eastern part of Elysium Planitia are very similar to each other in their roughness characteristics and differ from other volcanic plains on Mars, which suggests a distinctive eruption style. There are systematic latitudinal variations of roughness in both the southern highlands and the northern lowlands: terrains at high latitude are smoother at short baselines; the characteristic vertical scale related to this difference is several meters. Processes that could be responsible for formation of this trend include creep of ice‐rich near‐surface material at high latitudes, treatment of the surface with repetitive deposition and sublimation of seasonal frost, climate‐controlled deposition and/or cementation of dust at high latitudes, and repetitive sublimation and accumulation of subsurface ice at low latitudes with climate variations. All mechanisms of origin could operate more effectively under different climate conditions. Relevant morphological observations favor mechanisms involving deposition of smooth blankets at high latitudes.

Kilometer‐scale slopes on Mars and their correlation with geologic units: Initial results from Mars Orbiter Laser Altimeter (MOLA) data

Martian surface slopes were calculated at baselines from 0.4 to 25 km using profiles obtained by the Mars Orbiter Laser Altimeter (MOLA) instrument during the aerobraking phase of the Mars Global Surveyor mission. Median slope is proposed as a characteristic measurement of the typical surface roughness at each corresponding scale. Median slope is favored over RMS slope because it is not influenced by the small number of higher slopes at the upper end of the slope‐frequency distribution tail. Median slope complements interquartile scale roughness characterization in that it is more sensitive to smaller baseline slopes. A map of the median slope of the northern hemisphere is presented. Median slopes and their scale dependences are used to characterize typical kilometer‐scale roughness for a set of geologic units mapped in the northern hemisphere. This analysis demonstrates that many individual units and groups of units are characterized by distinctive surface slopes and that these characteristics are sufficiently different that they hold promise for use in the definition and characterization of units. Characterization of the slope properties of geologic units provides information useful in the interpretation of their origin and evolution. For example, the generally smooth topography of the diverse Vastitas Borealis Formation subunits is dominated by about 3 km, 0.3° steep features almost indistinguishable in Viking images. The roughness characteristics of this unit differ from those of other geologic units on Mars and suggest some distinctive process(es) of formation and/or modification of kilometer‐scale topography common for all subunits. The similarity of roughness characteristics of the several highland plateau units suggests that kilometer‐scale topography was largely inherited from the period of heavy bombardment. The northern polar cap and layered terrains are largely very smooth at small scale. The long, steep‐sloped tails of the slope‐frequency distributions are compared for the dominant terrain types in the northern hemisphere of Mars and are compared with Earth continents. The Vastitas Borealis Formation found in the northern lowlands differs significantly from both cratered uplands and volcanic plains, and these latter two units can be distinguished from each other on the basis of median slope. Terrestrial continents are smoother than cratered highlands but rougher than the Vastitas Borealis Formation and most volcanic plains. An inventory of all observed slopes much steeper than the angle of repose is presented. Steep slopes occur in the upper parts of tectonic scarps, most likely representing bedrock exposures. The presence of extremely steep slopes in the presumably ductile polar cap provides evidence for geologically recent and/or ongoing formation of these slopes.

Ice in the northern lowlands and southern highlands of Mars and its enrichment beneath the Elysium lavas

A comprehensive study of the Elysium region has revealed evidence of widespread ground‐ice, whose concentration and depth has varied considerably. Ejecta mobility and crater morphology analyses of a database of 7289 craters (which range in diameter from 1.875 to over 160 km) have been used to investigate the distribution of ice between longitudes 245°–155° and latitudes 20°S–45° N, an area incorporating all of the Elysium Quadrangle and large portions of Amenthes, Amazonis, Mare Tyrrhenum, Aeolis, Memnonia, Cebrenia and Diacria Quadrangles. Assuming that the observed variations in the crater characteristics are the result of the incorporation of varying concentrations of subsurface volatiles, the ice distribution is shown to be strongly dependent upon latitude, though its concentration varies with depth and is strongly influenced by geological setting. Important differences have been detected in the distribution of southern and northern lowland ground‐ice. While there is substantial evidence of near‐surface ice in the northern lowlands, where near‐surface limits are strongly influenced by latitude, the greatest concentrations of subsurface ice occur at much greater depths in the southern highlands. In addition, the ice appears to be enriched at depth in the Elysium Lavas, and nearer to the surface in deposits located to the northwest of Elysium Mons and in the Vastitas Borealis Formation. This distribution suggests that several ice‐emplacement mechanisms have operated on Mars, including the enrichment of the deeper ice by juvenile water from beneath the Elysium volcanic province, and the redistribution to the lowlands of a highly concentrated ice reservoir from a great depth within the highlands.