Marine‐target craters on Mars? An assessment study

Abstract— Observations of impact craters on Earth show that a water column at the target strongly influences lithology and morphology of the resultant crater. The degree of influence varies with the target water depth and impactor diameter. Morphological features detectable in satellite imagery include a concentric shape with an inner crater inset within a shallower outer crater, which is cut by gullies excavated by the resurge of water. In this study, we show that if oceans, large seas, and lakes existed on Mars for periods of time, marine‐target craters must have formed. We make an assessment of the minimum and maximum amounts of such craters based on published data on water depths, extent, and duration of putative oceans within “contacts 1 and 2,” cratering rate during the different oceanic phases, and computer modeling of minimum impactor diameters required to form long‐lasting craters in the seafloor of the oceans. We also discuss the influence of erosion and sedimentation on the preservation and exposure of the craters. For an ocean within the smaller “contact 2” with a duration of 100,000 yr and the low present crater formation rate, only ?1–2 detectable marine‐target craters would have formed. In a maximum estimate with a duration of 0.8 Gyr, as many as 1400 craters may have formed. An ocean within the larger “contact 1‐Meridiani,” with a duration of 100,000 yr, would not have received any seafloor craters despite the higher crater formation rate estimated before 3.5 Gyr. On the other hand, with a maximum duration of 0.8 Gyr, about 160 seafloor craters may have formed. However, terrestrial examples show that most marine‐target craters may be covered by thick sediments. Ground penetrating radar surveys planned for the ESA Mars Express and NASA 2005 missions may reveal buried craters, though it is uncertain if the resolution will allow the detection of diagnostic features of marine‐target craters. The implications regarding the discovery of marine‐target craters on Mars is not without significance, as such discoveries would help address the ongoing debate of whether large water bodies occupied the northern plains of Mars and would help constrain future paleoclimatic reconstructions.     

Pluto—The ninth planet's golden year: A meeting held at New Mexico State University,Las Cruces,New Mexico

The simple-to-complex transition for impact craters on Mars occurs at diameters between about 3 and 8 km. Ballistically emplaced ejecta surround primarily those craters that have a simple interior morphology, whereas ejecta displaying features attributable to fluid flow are mostly restricted to complex craters. Size-dependent characteristics of 73 relatively fresh Martian craters, emphasizing the new depth/diameter (d/D) data of D. W. G. Arthur (1980, to be submitted for publication), test two hypotheses for the mode of formation of central peaks in complex craters. In particular, five features appear sequentially with increasing crater size: first flat floors (3–4 km), then central peaks and shallower depths (4–5 km), next scalloped rims (? km), and lastly terraced walls (～8 km). This relative order indicates that a shallow depth of excavation and an unspecified rebound mechanism, not centripetal collapse and deep sliding, have produced central peaks and in turn have facilitated failure of the rim. The mechanism of formation of a shallow crater remains elusive, but probably operates only at the excavation stage of impact. This interpretation is consistent with two separate and complementary lines of evidence. First, field data have documented only shallow subsurface deformation and a shallow transient cavity in complex terrestrial meteorite craters and in certain surface-burst explosion craters; thus the shallow transient cavities of complex craters never were geometrically similar to the deep cavities of simple craters. Second, the average depths of complex craters and the diameters marking the transition from simple to complex craters on Mars and on three other terrestrial planets vary inversely with gravitational acceleration at the planetary surface, g, a variable more important in the excavation of a crater than in any subsequent modification of its geometry. The new interpretation is summarized diagrammatically for complex craters on all planets.     

Hydrothermal alteration at the Lonar Lake impact structure,India: Implications for impact cratering on Mars

Abstract— The 50,000 year old, 1.8 km diameter Lonar crater is one of only two known terrestrial craters to be emplaced in basaltic target rock (the 65 million year old Deccan Traps). The composition of the Lonar basalts is similar to martian basaltic meteorites, which establishes Lonar as an excellent analogue for similarly sized craters on the surface of Mars. Samples from cores drilled into the Lonar crater floor show that there are basaltic impact breccias that have been altered by post‐impact hydrothermal processes to produce an assemblage of secondary alteration minerals. Microprobe data and X‐ray diffraction analyses show that the alteration mineral assemblage consists primarily of saponite, with minor celadonite, and carbonate. Thermodynamic modeling and terrestrial volcanic analogues were used to demonstrate that these clay minerals formed at temperatures between 130°C and 200°C. By comparing the Lonar alteration assemblage with alteration at other terrestrial craters, we conclude that the Lonar crater represents a lower size limit for impact‐induced hydrothermal activity. Based on these results, we suggest that similarly sized craters on Mars have the potential to form hydrothermal systems, as long as liquid water was present on or near the martian surface. Furthermore, the Fe‐rich alteration minerals produced by post‐impact hydrothermal processes could contribute to the minor iron enrichment associated with the formation of the martian soil.     

Breccia dikes and crater‐related faults in impact craters on Mars: Erosion and exposure on the floor of a crater 75 km in diameter at the dichotomy boundary

Abstract— Environmental conditions on Mars are conducive to the modification and erosion of impact craters, potentially revealing the nature of their substructure. On Earth, postimpact erosion of complex craters in a wide range of target rocks has revealed the nature and distribution of craterrelated fault structures and a complex array of breccia and pseudotachylyte dikes, which range up to tens of meters in width and tens of kilometers in length. We review the characteristics of fault structures, breccia dikes, and pseudotachylyte dikes on Earth, showing that they occur in complex network‐like patterns and are often offset along late‐stage crater‐related faults. Individual faults and dikes can undulate in width and can branch and bifurcate along strike. Detailed geological analyses of terrestrial craters show that faults and breccia dikes form during each of the major stages of the impact‐cratering process (compression, excavation, and modification). We report here on the discovery of prominent, lattice‐like ridge networks occurring on the floor of a highly modified impact crater 75 km in diameter near the dichotomy boundary of the northern lowland and southern upland. Interior fill and crater‐floor units have been exhumed by fluvial and eolian processes to reveal a unit below the crater floor containing a distinctive set of linear ridges of broadly similar width and forming a lattice‐like pattern. Ridge exposures range from ?1–4 km in length and ?65–120 m in width, are broadly parallel, straight to slightly curving, and are cross‐cut by near‐orthogonal ridges, forming a box or lattice‐like pattern. Ridges are exposed on the exhumed crater floor, extending from the base of the wall toward the center. On the basis of the strong similarities of these features to terrestrial crater‐related fault structures and breccia dikes, we interpret these ridges to be faults and breccia dikes formed below the floor of the crater during the excavation and modification stages of the impact event, and subsequently exhumed by erosion. The recognition of such features on Mars will help in documenting the nature of impact‐cratering processes and aid in assessment of crustal structure. Faults and breccia dikes can also be used as data for the assessment of post‐cratering depths and degrees of landform exhumation.     

Layered ejecta craters and the early water/ice aquifer on Mars

Abstract— A model for emplacement of deposits of impact craters is presented that explains the size range of Martian layered ejecta craters between 5 km and 60 km in diameter in the low and middle latitudes. The impact model provides estimates of the water content of crater deposits relative to volatile content in the aquifer of Mars. These estimates together with the amount of water required to initiate fluid flow in terrestrial debris flows provide an estimate of 21% by volume (7.6 × 10⁷km³) of water/ice that was stored between 0.27 and 2.5 km depth in the crust of Mars during Hesperian and Amazonian time. This would have been sufficient to supply the water for an ocean in the northern lowlands of Mars. The existence of fluidized craters smaller than 5 km diameter in some places on Mars suggests that volatiles were present locally at depths less than 0.27 km. Deposits of Martian craters may be ideal sites for searches for fossils of early organisms that may have existed in the water table if life originated on Mars.     

Differential melt scaling for oblique impacts on terrestrial planets

Analytical estimates of melt volumes produced by a given projectile and contained in a given impact crater are derived as a function of impact velocity, impact angle, planetary gravity, target and projectile densities, and specific internal energy of melting. Applications to impact events and impact craters on the Earth, Moon, and Mars are demonstrated and discussed. The most probable oblique impact (45°) produces ～1.6 times less melt volume than a vertical impact, and ～1.6 and 3.7 times more melt volume than impacts with 30° and 15° trajectories, respectively. The melt volume for a particular crater diameter increases with planetary gravity, so a crater on Earth should have more melt than similar-size craters on Mars and the Moon. The melt volume for a particular projectile diameter does not depend on gravity, but has a strong dependence on impact velocity, so the melt generated by a given projectile on the Moon is significantly larger than on Mars. Higher surface temperatures and geothermal gradients increase melt production, as do lower energies of melting. Collectively, the results imply thinner central melt sheets and a smaller proportion of melt particles in impact breccias on the Moon and Mars than on Earth. These effects are illustrated in a comparison of the Chicxulub crater on Earth, linked to the Cretaceous–Tertiary mass extinction, Gusev crater on Mars, where the Mars Exploration Rover Spirit landed, and Tsiolkovsky crater on the Moon. The results are comparable to those obtained from field and spacecraft observations, other analytical expressions, and hydrocode simulations.     

A survey of possible impact structures on 25143 Itokawa

We determined the morphologies and dimensions of possible impact craters on the surface of Asteroid 25143 Itokawa from images taken by the Hayabusa spacecraft. Circular depressions, circular features with flat floors or convex floors, and circular features with smooth surfaces were identified as possible craters. The survey identified 38 candidates with widely varying morphologies including rough, smooth and saddle-shaped floors, a lack of raised rims and fresh material exposures. The average depth/diameter ratio was 0.08±0.03: these craters are very shallow relative to craters observed on other asteroids. These shallow craters are a result of (1) target curvature influencing the cratering process, (2) raised rim not being generated by this process, and (3) fines infilling the craters. As many of the crater candidates have an unusual appearance, we used a classification scheme that reflects the likelihood of an observed candidate's formation by a hypervelocity impact. We considered a variety of alternative interpretations while developing this scheme, including inherited features from a proto-Itokawa, spall scars created by the disruption of the proto-Itokawa, spall scars following the formation of a large crater on Itokawa itself, and apparent depressions due to random arrangements of boulders. The size-frequency distribution of the crater candidates was close to the empirical saturation line at the largest diameter, and then decline with decreasing diameter.     

A new method to determine the direction of impact: Asymmetry of concentric impact craters as observed in the field (Lockne), on Mars,in experiments,and simulations

Most impacts occur at an angle with respect to the horizontal plane. This is primarily reflected in the ejecta distribution, but at very low angle structural asymmetries such as elongation of the crater and nonradial development of the central peak become apparent. Unfortunately, impact craters with pristine ejecta layers are rare on Earth and also in areas with strong past or ongoing surface erosion on other planetary bodies, and the structural analysis of central peaks requires good exposures or even on‐site access to outcrop. However, target properties are known to greatly influence the shape of the crater, especially the relatively common target configuration of a weaker layer covering a more rigid basement. One such effect is the formation of concentric craters, i.e., a nested, deeper, inner crater surrounded by a shallow, outer crater. Here, we show that with decreasing impact angle there is a downrange shift of the outer crater with respect to the nested crater. We use a combination of (1) field observation and published 3‐D numerical simulation of one of the best examples of a terrestrial, concentric impact crater formed in a layered target with preserved ejecta layer: the Lockne crater, Sweden; (2) remote sensing data for three pristine, concentric impact craters on Mars with preserved ejecta layers further constraining the direction of impact; as well as (3) laboratory impact experiments, to develop the offset in crater concentricity into a complementary method to determine the direction of impact for layered‐target craters with poorly preserved ejecta layers.     

Comparison of the crater morphology-size relationship for Mars,Moon, and Mercury

New central peak-crater size data for Mars shows that a higher percentage of relatively unmodified Martian craters have central peaks than do fresh lunar craters below a diameter of 30 km. For example, in the diameter range 10 to 20 km, 60% of studied Martian craters have central peaks compared to 26% for the Moon. Gault et al. (1975, J. Geophys. Res.80, 2444–2460) have demonstrated that central peaks occur in smaller craters on Mercury than on the Moon, and that this effect is due to the different gravity fields in which the craters formed. Similar differences when comparing Mars and the Moon show that gravity has affected the diameter at which central peaks form on Mars. Erosion on Mars, therefore, does not completely mask differences in crater interior structure that are caused by differences in gravity. Effects of Mars' higher surface gravity when compared to the Moon are not detected when comparing terrace and crater shape data. The morphology-crater size statistics also show that a full range of crater shapes occur on Mars, and craters tend to become more morphologically complex with increasing diameter. Comparisons of Martian and Mercurian crater data show differences which may be related to the greater efficacy of erosion on Mars.     

Impact Excavation and the Search for Subsurface Life on Mars

Because of the ubiquity of subsurface microbial life on Earth, examination of the subsurface of Mars could provide an answer to the question of whether microorganisms exist or ever existed on that planet. Impact craters provide a natural mechanism for accessing the deep substrate of Mars and exploring its exobiological potential. Based on equations that relate impact crater diameters to excavation depth we estimate the observed crater diameters that are required to prospect to given depths in the martian subsurface and we relate these depths to observed microbiological phenomena in the terrestrial subsurface. Simple craters can be used to examine material to a depth of ∼270 m. Complex craters can be used to reach greater depths, with craters of diameters ≥300 km required to reach depths of 6 km or greater, which represent the limit of the terrestrial deep subsurface biosphere. Examination of the ejecta blankets of craters between 17.5 and 260 km in diameter would provide insights into whether there is an extant, or whether there is evidence of an extinct, deep subsurface microbiota between 500 and 5000 m prior to committing to large-scale drilling efforts. At depths <500 m some crater excavations are likely to be more important than others from an exobiological point of view. We discuss examples of impacts into putative intracrater paleolacustrine sediments and regions associated with hydrothermal activity. We compare these depths to the characteristics of subsurface life on Earth and the fossil microbiological record in terrestrial impact craters.     

Morphologic classification of Martian craters and some implications

A computer data bank containing information on crater sizes, locations, and morphologies for all craters visible on Mariner 9 wide-angle mapping photography was used to construct a crater morphologic classification. Four general classes were constructed that can be interpreted to represent increasing degrees of crater degradation. Fresh class craters are nearly unmodified and consist of deep bowl-shaped craters and deep, flat-floored craters with terraced walls. The slightly modified class consists of deep flat-floored craters that usually have raised rims, but lack the terracing, central peaks, and hummocky floors indicative of unmodified impact crater morphology. Craters in the modified class are rimless and shallow and those in the ghost class are rimless and extremely shallow. Retention ages for fresh (i.e. unmodified) class craters on equatorial cratered terrain range from millions to billions of years, depending on the impact flux history used. If the trend is toward billions of years, then present degradation rates on Mars are low relative to earlier history and most craters in the degraded classes were probably modified in an early (>3.3 b.y.?) period.     

Martian planetwide crater distributions: Implications for geologic history and surface processes

Population-density maps of craters in three size ranges (0.6 to 1.2 km, 4 to 10 km, and >20 km in diameter) were compiled for most of Mars from Mariner 9 imagery. These data provide: historical records of the eolian processes (0.6 to 1.2 km craters); stratigraphic, relative, and absolute timescales (4 to 10 km craters); and a history of the early postaccretional evolution of the uplands (> 20 km craters).Based on the distribution of large craters (>20 km diameters), Mars is divisible into two general classes of terrain, densely cratered and very lightly cratered—a division remarkably like the uplands-maria dichotomy of the moon. It is probable that this bimodal character in the density distribution of large craters arose from an abrupt transition in the impact flux rate from an early intense period associated with the tailing off of accretion to an extended quiescent epoch, not from a void in geological activity during much of Mars' history. Radio-isotope studies of Apollo lunar samples show that this transition occurred on the moon in a short time.The intermediate-sized craters (4 to 10 km diameter) and the small-sized craters (0.6 to 1.2 km diameter) appear to be genetically related. The smaller ones are apparently secondary impact craters generated by the former. Most of the craters in the larger of these two size classes appear fresh and uneroded, although many are partly buried by dust mantles. Poleward of the 40° parallels the small fresh craters are notably absent owing to these mantles. The density of small craters is highest in an irregular band centered at 20°S. This band coincides closely with (1) the zone of permanent low-albedo markings; (2) the “wind equator” (the latitude of zero net north or south transport at the surface); and (3) a band that includes a majority of the small dendritic channels. Situated in the southermost part of the equatorial unmantled terrain which extends from about 40°N to 40°S, this band is apparently devoid of even a thin mantle. Because this belt is also coincident with the latitutde of maximum solar insolation (periapsis occurs near summer solstice), we suggest that this band arises from the asymmetrical global wind patterns at the surface and that the band probably follows the latitude of maximum heating which migrates north and south from 25°N to 25°S within the unmantled terrain on a 50,000 year timescale.The population of intermediate-sized craters (4–10 km diameter) appears unaffected by the eolian mantles, at least within the ±45° latitudes. Hence the local density of these craters is probably a valid indicator of the relative age of surfaces generated during the period since the uplands were intensely bombarded and eroded. It now appears that the impact fluxes at Mars and the moon have been roughly the same over the last 4 b.y. because the oldest postaccretional, mare-like surfaces on Mars and the moon display about the same crater density. If so, the nearness of Mars to the asteroid belt has not generated a flux 10 to 25 times greater than the lunar flux. Whereas the lunar maria show a variation of about a factor of three in crater density from the oldest to the youngest major units, analogous surfaces on Mars show a variation between 30 and 50. This implies that periods of active eolian erosion, tectonic evolution, volcanic eruption, and possibly fluvial modification have been scattered throughout Martian history since the formation and degradation of the martian uplands and not confined to small, ancient or recent, epochs. These processes are surely active on the planet today.     

An investigation of Martian northern high‐latitude and polar impact crater interiors: Atypical interior topographic features and cavity wall slopes

Abstract– We examine Martian northern high‐latitude and polar impact craters (NPICs) to better understand the north polar materials and polar processes. We examine topographic characteristics for 346 NPICs and compare them to global fit data (e.g., Garvin et al. 2003 ; Boyce and Garbeil 2007 ) as well as to a small set (N = 92) of southern high‐latitude and polar impact craters (SPICs). We find that the NPIC population above 57° N is significantly shallower than the global crater population. This suggests that the NPICs (1) were initially shallow due to target properties of polar geologic units; (2) were once deeper, but have been infilled due to polar processes; or (3) a combination of both. Indeed, many of the NPICs exhibit considerable noncentral peak interior topographic features (IFTs), which may be indicative of infilling processes. The NPIC IFTs also appear to display trends in their preferential orientation within the crater cavity; some SPICs display similar interior features, but do not show a clear preference in their orientation within the crater cavity. In addition, the NPIC population displays cavity wall slope trends that seem to indicate steepening of slopes with increasing crater diameter in comparison to the global slope trend ( Garvin et al. 2003 ). These trends suggest that the NPICs are unique in their geometry when compared to the global data set as well as with the SPICs further indicating that the north polar region may exhibit target properties and polar processes not seen in the south polar region or elsewhere on Mars.     

Evidence for Hesperian impact-induced hydrothermalism on Mars

Several hydrated silicate deposits on Mars are observed within craters and are interpreted as excavated Noachian material. Toro crater (71.8°E, 17.0°N), located on the northern edge of the Syrtis Major Volcanic Plains, shows spectral and morphologic evidence of impact-induced hydrothermal activity. Spectroscopic observations were used to identify extensive hydrated silicate deposits, including prehnite, chlorites, smectites, and opaline material, a suite of phases that frequently results from hydrothermal alteration in terrestrial craters and also expected on Mars from geochemical modeling of hydrothermal environments. When combined with altimetry and high-resolution imaging data, these deposits appear associated predominantly with the central uplift and with portions of the northern part of the crater floor. Detailed geologic mapping of these deposits reveals geomorphic features that are consistent with hydrothermal activity that followed the impact event, including vent-like and conical mound structures, and a complex network of tectonic structures caused by fluid interactions such as fractures and joints. The crater age has been calculated from the cumulative crater size-frequency distributions and is found to be Early Hesperian. The evidence presented here provides support for impact-induced hydrothermal activity in Toro crater, that extends phyllosilicate formation processes beyond the Noachian era.     

Bosumtwi impact structure,Ghana: Evidence for fluidized emplacement of the ejecta

The about 10.5 km diameter Bosumtwi impact crater is one of the youngest large impact structures on Earth. The crater rim is readily noticed on topographic maps or in satellite imagery. It defines a circular basin filled by water (Lake Bosumtwi) and lacustrine sediments. The morphology of this impact structure is also characterized by a circular plateau extending beyond the rim and up to 9–10 km from the center of the crater (about 2 crater radii). This feature comprises a shallow ring depression, also described as an annular moat, and a subdued circular ridge at its outer edge. The origin of this outermost feature could so far not be elucidated based on remote sensing data only. Our approach combines detailed topographic analysis, including roughness mapping, with airborne radiometric surveys (mapping near‐surface K, Th, U concentrations) and field observations. This provides evidence that the moat and outer ring are features inherited from the impact event and represent the partially eroded ejecta layer of the Bosumtwi impact structure. The characteristics of the outer ridge indicate that ejecta emplacement was not purely ballistic but requires ejecta fluidization and surface flow. The setting of Bosumtwi ejecta can therefore be considered as a terrestrial analog for rampart craters, which are common on Mars and Venus, and also found on icy bodies of the outer solar system (e.g., Ganymede, Europa, Dione, Tethys, and Charon). Future studies at Bosumtwi may therefore help to elucidate the mechanism of formation of rampart craters.     

Cratering on Mars with almost no atmosphere or volatiles: Pangboche crater

Pangboche crater (17.2°N, 226.7°E; 10.4 km dia.) lies close to the summit of Olympus Mons volcano, Mars, at an elevation of ~20.9 km above the datum. Given a scale height of 11.1 km for the atmosphere, this relatively large fresh crater most likely formed at an atmospheric pressure <1 mbar in essentially volatile‐free young lava flows. Detailed analysis of Pangboche crater from High Resolution Imaging Science Experiment (HiRISE) and Context Camera (CTX) images reveals that volatile‐related features (e.g., fluidized ejecta layers and pitted floor material) are absent. In contrast, abundant impact melt occurs on the floor, inner walls, and rim of the crater, and there is an extensive field of secondary craters that extend up to approximately 45 km from the rim crest. All of these attributes argue that it was the absence of volatiles in the target rocks at the time of crater formation, rather than the thin atmosphere, which had a controlling influence on crater morphology. Digital elevation data derived from the CTX images reveal that Pangboche crater has a depth of about 954 m (depth/diameter = approximately 0.092) and that uplifted target rocks comprise about 58% of the relief of the 180 m‐high north rim. As the target material comprised a sequence of layered lava flows, Pangboche crater may well represent the best crater on Mars for direct comparison with craters formed on the Moon (permitting variations in gravitational effects to be investigated) or on Mercury (allowing the role of the atmosphere to be studied).     

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.     

Impact cratering in H2O‐bearing targets on Mars: Thermal field under craters as starting conditions for hydrothermal activity

Abstract– We present a case modeling study of impact crater formation in H₂O‐bearing targets. The main goal of this work was to investigate the postimpact thermal state of the rock layers modified in the formation of hypervelocity impact craters. We present model results for a target consisting of a mixture of H₂O‐ice and rock, assuming an ice/water content variable with depth. Our model results, combined with results from previous work using dry targets, indicate that for craters larger than about 30 km in diameter, the onset of postimpact hydrothermal circulation is characterized by two stages: first, the formation of a mostly dry, hot central uplift followed by water beginning to flow in and circulate through the initially dry and hot uplifted crustal rocks. The postimpact thermal field in the periphery of the crater is dependent on crater size: in midsize craters, 30–50 km in diameter, crater walls are not strongly heated in the impact event, and even though ice present in the rock may initially be heated enough to melt, overall temperatures in the rock remain below melting, undermining the development of a crater‐wide hydrothermal circulation. In large craters (with diameters more than 100 km or so), the region underneath the crater floor and walls is heated well above the melting point of ice, thus facilitating the onset of an extended hydrothermal circulation. These results provide preliminary constraints in characterizing the many water‐related features, both morphologic and spectroscopic, that high‐resolution images of Mars are now detecting within many Martian craters.     

The Whitecourt meteorite impact crater,Alberta, Canada

Abstract– The <1,100 yr old Whitecourt meteorite impact crater, located south of Whitecourt, Alberta, Canada, is a well‐preserved bowl‐shaped structure having a depth and diameter of approximately 6 and 36 m, respectively. There are fewer than a dozen known terrestrial sites of similar size and age. Unlike most of these sites, however, the Whitecourt crater contains nearly all of the features associated with small impact craters including meteorites, ejecta blanket, observable transient crater boundary, raised rim, and associated shock indicators. This study indicates that the crater formed from the impact of an approximately 1 m diameter type IIIAB iron meteoroid traveling east‐northeast at less than approximately 10 km s^?1, striking the surface at an angle between 40° and 55° to horizontal. It appears that the main mass survived atmospheric transit relatively intact, with fragmentation and partial melting during impact. Most meteoritic material has a jagged, shrapnel‐like morphology and is distributed downrange of the crater.     

Ejecta emplacement and modes of formation of martian fluidized ejecta craters

From an analysis of 1173 craters possessing single (Type I) and double (Type 2) concentric ejecta deposits, Type 2 craters are found to occur most frequently in areas that have also been described as possessing periglacial features. The frequency of occurence of central peaks and wall failure (terraces plus scallops) within the craters indicate that, by analogy with previous analyses, Type 1 craters form in more fragmental targets than Type 2 craters. The maximum range of the outer ejecta deposits of Type 2 craters, however, consistently extends ～0.8 crater radii further than ejecta deposits of Type 1 craters, suggesting a greater degree of ejecta fluidization for the twin-lobed Type 2 craters. Numerous characteristics of Ries Crater, West Germany, show similarities to craters on Mars, indicating that Martian fluidized ejecta craters may be closer analogs to this terrestrial crater than are lunar craters.