In September 1977, a crater studies workshop was held for the purpose of developing standardized data analysis and presentation techniques. This report contains the unanimous recommendations of the participants. This first meeting considered primarily crater size-frequency data. Future meetings will treat other aspects of crater studies such as morphologies.

We have conducted a study of the geomorphology, stratigraphy, and composition of Gale crater and its central mound of layered deposits, a region that is being considered as a landing site for the Mars Science Laboratory (MSL) mission. We surveyed the crater for fluvial features and evaluated hypotheses for the origin of the central mound, including aeolian, lacustrine, spring mound, and volcanic processes. The rim of Gale crater is extensively dissected by fluvial channels, and the observed channels on the crater rim appear to flow into the crater with no obvious outlet. Many of the channels are dendritic, some showing third or fourth-order tributaries. Inverted fluvial features occur on the floor and mound, including several meandering channels and channel networks within the proposed MSL landing ellipse. Fractures on the mound are common and are often erosion-resistant, possibly suggesting alteration and/or cementation by fluid. The key geomorphic units of the landing site and mound include: a fan-shaped deposit in the landing site (divided into low and high thermal inertia portions), hummocky plains, a widespread mound-skirting unit, a basal unit that underlies the mound and floor units, a light toned ridge, a phyllosilicate-bearing trough that parallels the ridge, dark-toned layered yardangs, light-toned yardangs, an upper mound unit, a thin mantle unit, and several lobate features on the northern flanks of the mound. Erosional unconformities suggest that the dark-toned layered yardang unit was significantly eroded prior to the deposition of the light-toned yardang-forming unit and the upper mound unit. Fan-shaped deposits of material from the dark-toned layered yardang unit extend onto the mound-skirting unit in some locations, suggesting that the skirting unit was emplaced prior to or during a period of significant erosion of the mound. A fan-shaped unit on the mound near the landing site appears to be an isolated patch of the mound-skirting unit, rather than an alluvial fan. The phyllosilicate-bearing unit has a morphology similar to aeolian bedforms, but appears to be indurated. This unit is visible only in the trough between the light-toned ridge and the mound and may be comprised of altered, aeolian-transported material, trapped by the trough and cemented. The lobate features on the northern flanks of the mound show some evidence of layering in MOC images, and may be due to landslides or glacial/periglacial processes. A streamlined texture on one of the lobate units may be due to fluvial activity or compositional banding in the unit. Some layers in the mound are traceable for >10 km, suggesting that a spring mound origin for the mound, which would produce facies changes over short distances, is unlikely. We were unable to rule out a lacustrine, aeolian or volcanic origin for the lower mound layers. We found apparent large-scale (100s of meters) cross-bedding in the layers of the upper mound unit, suggesting an aeolian origin. We have identified multiple traverses for Mars Science Laboratory that would provide access to the diverse features on the crater floor and the central mound and make full use of the rover's scientific payload to address the question of martian habitability.

Abstract For several decades, most planetary researchers have regarded the impact crater populations on solid-surfaced planets and smaller bodies as predominantly reflecting the direct ('primary') impacts of asteroids and comets. Estimates of the relative and absolute ages of geological units on these objects have been based on this assumption. Here we present an analysis of the comparatively sparse crater population on Jupiter's icy moon Europa and suggest that this assumption is incorrect for small craters. We find that 'secondaries' (craters formed by material ejected from large primary impact craters) comprise about 95 per cent of the small craters (diameters less than 1 km) on Europa. We therefore conclude that large primary impacts into a solid surface (for example, ice or rock) produce far more secondaries than previously believed, implying that the small crater populations on the Moon, Mars and other large bodies must be dominated by secondaries. Moreover, our results indicate that there have been few small comets (less than 100 m diameter) passing through the jovian system in recent times, consistent with dynamical simulations.

The large, fresh crater Tycho in the nearside lunar highlands has an extensive system of bright rays covering approximately 560,000 km 2, containing dense clusters of secondary craters. Examination of crater densities in several clusters shows that Tycho produced almost 10 6 secondary craters larger than 63 m diameter. This is a lower limit, because small crater densities are reduced, most likely by mass wasting. We estimate a crater erasure rate of 2 6 cm/Myr, varying with crater size, and consistent with previous results. This process has removed many small craters, and it is probable that the original number of secondary craters formed by Tycho was higher. Also, we can only identify distant secondaries of Tycho where they occur in bright rays. Craters on Mars and Europa also formed large numbers of secondaries, but under possibly ideal conditions for spallation as a mechanism to produce high-velocity ejecta fragments. The results from Tycho show that large numbers of such fragments can be produced even from impact into a heavily fragmented target on which spallation is expected to be less important.

Results by Neukum et al. (2001) and Ivanov (2001) are combined with crater counts to estimate ages of Martian surfaces. These results are combined with studies of Martian meteorites (Nyquist et al. , 2001) to establish a rough chronology of Martian history. High crater densities in some areas, together with the existence of a 4.5 Gyr rock from Mars (ALH84001), which was weathered at about 4.0 Gyr, affirm that some of the oldest surfaces involve primordial crustal materials, degraded by various processes including megaregolith formation and cementing of debris. Small craters have been lost by these processes, as shown by comparison with Phobos and with the production function, and by crater morphology distributions. Crater loss rates and survival lifetimes are estimated as a measure of average depositional/erosional rate of activity. We use our results to date the Martian epochs defined by Tanaka (1986). The high crater densities of the Noachian confine the entire Noachian Period to before about 3.5 Gyr. The Hesperian/Amazonian boundary is estimated to be about 2.9 to 3.3 Gyr ago, but with less probability could range from 2.0 to 3.4 Gyr. Mid-age dates are less well constrained due to uncertainties in the Martian cratering rate. Comparison of our ages with resurfacing data of Tanaka et al. (1987) gives a strong indication that volcanic, fluvial, and periglacial resurfacing rates were all much higher in approximately the first third of Martian history. We estimate that the Late Amazonian Epoch began a few hundred Myr ago (formal solutions 300 to 600 Myr ago). Our work supports Mariner 9 era suggestions of very young lavas on Mars, and is consistent with meteorite evidence for Martian igneous rocks 1.3 and 0.2 0.3 Gyr old. The youngest detected Martian lava flows give formal crater retention ages of the order 10 Myr or less. We note also that certain Martian meteorites indicate fluvial activity younger than the rocks themselves, 700 Myr in one case, and this is supported by evidence of youthful water seeps. The evidence of youthful volcanic and aqueous activity, from both crater-count and meteorite evidence, places important constraints on Martian geological evolution and suggests a more active, complex Mars than has been visualized by some researchers.

Analysis of the Chandrayaan-1 Terrain Mapping Camera image of a 2002km×2702km area in the Mare Imbrium region revealed a cluster of thousands of fresh and buried impact craters in the size range of 20–130002m. A majority of the large fresh craters with diameter ranging from 160 to 127002m exhibit near-circular mounds (30–33502m diameter and 10–4002m height) in the crater floor, and their size depends on the host crater size. The origin of this cluster of secondary craters may be traced to Copernicus crater, based on global lunar image and the analysis of Chandrayaan-1 Hyper Spectral Imager data. Our findings provide further evidence for secondary crater formation by low-velocity impact of a cloud of clustered fragments. The presence of central mounds can also distinguish the secondary craters from the primary craters and refine the chronology of lunar surface based on counting of small craters.Highlights? A distal ray of Copernicus crater contains a cluster of thousands of secondary craters. ? Central mound bearing secondary craters are abundant and are deeper than regolith. ? Central mound craters were produced by low velocity clustered impacts. ? The crater clusters are spectrally similar to Copernicus crater ejecta. ? The central mound craters should be excluded from crater counting dating.

Small craters (less than one kilometer diameter) can be primary craters produced by impact of interplanetary debris, or they can be secondary craters produced by fallback of high-velocity ejecta blocks from much larger but infrequent primary impacts. The prevalent assumption over recent decades has been that primaries are most abundant, so most small craters are independent random events and can be used for dating. However, recent results from Europa and Mars support the early theory that distant secondaries globally dominate the number of small lunar craters; this would invalidate part of production functions that have been widely used for age dating. Crater excavation results in higher mean ejection velocities for smaller fragments, resulting in a steeper size-frequency distribution for secondary craters than is produced by the same size-frequency distribution of interplanetary debris. This review also discusses how small craters can sometimes be used to derive meaningful age constraints.

The population of secondary craters - craters formed by the ejecta from an initial impact event - is important to understand when deriving the age of a solid body's surface. Only one crater on Mars, Zunil, has been studied in-depth to examine the distribution, sizes, and number of these features. Here, we present results from a much larger and older Martian crater, Lyot, and we find secondary crater clusters at least 5200 km from the primary impact. Individual craters with diameters >800 m number on the order of 10. Unlike the previous results from Zunil, these craters are not contained in obvious rays, but they are linked back to Lyot due to the clusters' alignment along great circles that converge to a common origin. These widespread and abundant craters from a single impact limit the accuracy of crater age-dating on the Martian surface and beyond.

V-shaped ridge components of the herringbone pattern associated with lunar secondary crater chains have been simulated by simultaneous and nearly simultaneous impact of two projectiles near one another. The impact velocities and angles of the projectiles were similar to those of the fragments that produced secondary craters found at various ranges from large lunar craters. Variables found to affect the included angles of the V-shaped ridges are: relative time of impact of the projectiles, impact angle, relative projectile mass, and azimuth angle of the crater chain relative to the projection of the flight line onto the target surface. The functional relationships between the forms of the ridges and many of these variables are similar to those observed for lunar V-shaped ridges. Comparison of the magnitudes of the ridge angles of both laboratory crater pairs and secondary crater chains of the crater Copernicus implies that material was ejected from Copernicus at angles in excess of 60 , measured from the normal, to form many of Copernicus' satellitic craters. Moreover, other independent calculations presented indicate that many of the fragments that produced secondary craters also ricocheted to produce tertiary craters. Application of the study to identification of isolated secondary craters and to the determination of the origin of large lunar craters is discussed.

Figure 2. A crater in Apollo 17 panoramic camera photograph number 2345 is surrounded by boulders. This crater is likely a distant secondary crater of Burg, a 39.1 km diameter crater on the northeast edge of Mare Serenitatis. The sun shines from the left in this photograph; the right wall of the crater is illuminated, and boulder shadows fall to the right.

A 10-km diameter crater named Zunil in the Cerberus Plains of Mars created 65 10 7 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 (8210 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 10 10 rock fragments 8210 cm diameter, leading to up to 10 9 secondary craters 8210 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 651000 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 650.2 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.

[1] The identification of small (D < a few kilometers) secondary craters and their global distributions are of critical importance to improving our knowledge of surface ages in the solar system. We investigate a technique by which small, distal secondary craters can be discerned from the surrounding primary population of equivalent size based on asymmetries in their ejecta blankets. The asymmetric ejecta blankets are visible in radar circular polarization ratio (CPR) but not as optical albedo features. Measurements with our new technique reveal 94 secondary craters on the Newton and Newton-A crater floors near the lunar south pole. These regions are not in an obvious optical ray, but the orientation of asymmetric secondary ejecta blankets suggests that they represent an extension of the Tycho crater ray that crosses Clavius crater. Including the secondary craters at Newton and Newton-A skews the terrain age inferred by crater counts. It is reduced by few percentages by their removal, from 3.8 to 3.75 Gyr at Newton-A. Because 090008hidden rays090009 like that identified here may also occur beyond the edges of other optically bright lunar crater rays, we assess the effect that similar but hypothetical populations would have on lunar terrains of various ages. The average secondary crater density measured at 1 km diameter is equivalent to the crater density at 1 km on a 3.4 Gyr surface [Neukum et al., 2001]. Younger surfaces (i.e., younger crater ejecta blankets) would be dominated by secondary craters below 1 km if superposed by a hidden ray.

Similarity is found in crater densities on the most heavily cratered surfaces throughout the solar system. This is hypothesized to result from a steady-state “saturation equilibrium” being approached or achieved by cratering processes. This hypothesis conflicts with some recent interpretations. However, it accounts for (1) a similarity in maximum relative crater density, below certain theoretically predicted values, on all heavily cratered surfaces; (2) a leveling off at this same relative density among 100-m scale (secondary?) craters in populations on lunar maria and other sparsely cratered lunar surfaces; (3) the approximate uniformity of maximum relative densities on Saturn satellites (in spite of dramatic variations predicted from nonsaturation models assuming heliocentric impactors). The lunar frontside upland crater population, sometimes described as a well-preserved production function useful for interpreting other planetary surfaces, is found not to be a production function. It was modified by intercrater plains formed (at least partly) by early upland basaltic lava flooding, recently confirmed spectrophotometrically. Consistent with this, counts in “pure uplands” (those lacking intercrater plains) match the proposed saturation equilibrium density. Variations among large ( D > 64 km) crater populations are found, but these may involve several hypothesized mechanisms that rapidly obliterate large craters, especially on icy surfaces. Recent models, in which different populations of interplanetary bodies hit different planets, need further appraisal.

We have examined a number of potential landing sites to study effects associated with impact crater populations. We used Mars Global Surveyor high resolution MOC images, and emphasized "ground truth" by calibrating with the MOC images of Viking 1 and Pathfinder sites. An interesting result is that most of Mars (all surfaces with model ages older than 100 My) have small crater populations in saturation equilibrium below diameters D approx. = 60 meters (and down to the smallest resolvable, countable sizes, approx. = 15 m). This may have consequences for preservation of surface bedrock exposures accessible to rovers. In the lunar maria, a similar saturation equilibrium is reached for crater diameters below about 300 meters, and this has produced a regolith depth of about 10-20 meters in those areas. Assuming linear scaling, we infer that saturation at D approx. = 60 m would produce gardening and Martian regolith, or fragmental layers, about 2 to 4 meters deep over all but extremely young surfaces (such as the very fresh thin surface flows in southern Elysium Planitia, which have model ages around 10 My or less). This result may explain the global production of ubiquitous dust and fragmental material on Mars. Removal of fines may leave the boulders that have been seen at all three of the first landing sites. Accumulation of the fines elsewhere produces dunes. Due to these effects, it may be difficult to set down rovers in areas where bedrock is well preserved at depths of centimeters, unless we find cliff sides or areas of deflation where wind has exposed clean surfaces (among residual boulders?) We have also surveyed the PHOBOS 2 Termoskan data to look for regions of thermal anomalies that might produce interesting landing sites. For landing site selection, two of the more interesting types of features are thermally distinct ejecta blankets and thermally distinct channels and valleys. Martian "thermal features" such as these that correlate closely with nonaeolian geologic features are extremely rare, presumably due to reworking of the surface as discussed above, and due to aeolian processes. Thermally distinct ejecta blankets are excellent potential future locations for landers, as well as remote sensing, because they represent relatively dust free exposures of material excavated from depth. However, few, if any meet the current constraints on elevation for Mars '01. Thermally distinct channels, which tend to have fretted morphologies, and are higher in inertia than their surroundings, offer a unique history and probable surface presence of material from various stratigraphic layers and, locations, views of the surrounding walls, and possible areas of past standing water, flowing water, or increased amounts of diffusing water. Any presence of water (e.g., diffusing may have enhanced duricrust formation in the channels, thus increasing the thermal inertias (flowing water may alternatively have enhanced rock deposition, which also could explain the inertia enhancements instead of crust formation). Some of the thermally distinct channels do meet the elevation criteria for '01. We are looking particularly at the relatively flat areas at the northern end of Hydraotes Chaos (eastern end of Valles Marineris), near the beginnings of Tiu and Simud Valles, which appear to meet most all of the current '01 landing criteria. For thermally distinct channels, valleys, and ejecta blankets, we have searched and continue to search for MOC images that may help clarify their characteristics and assist with potential landing site characterization.

The relationship between primary crater diameter and mean maximum secondary crater diameter for lunar craters is essentially linear for primary diameters from 0.5 to 260 km. The mean range from the primary to the largest secondary craters is also nearly linearly related to primary diameter. Extrapolation of these results to lunar basin scale supports the contention that basin secondaries of large diameter are to be found across much of the Moon.

On airless bodies such as the Moon and Mercury, secondary craters on the continuous secondaries facies of fresh craters mostly occur in chains and clusters. They have very irregular shapes. Secondaries on the continuous secondaries facies of some Martian and Mercurian craters are more isolated from each other in distribution and are more circular in shape, probably due to the effect of target properties on the impact excavation process. This paper studies secondaries on the continuous secondaries facies of all fresh lunar complex craters using recently-obtained high resolution images. After a global search, we find that 3 impact craters and basins on the Moon have circular and isolated secondaries on the continuous secondaries facies similar to those on Mercury: the Orientale basin, the Antoniadi crater, and the Compton crater. The morphological differences between such special secondaries and typical lunar secondaries are quantitatively compared and analyzed. Our preliminary analyses suggest that the special secondaries were probably caused by high temperature gradients within the local targets when these craters and basins formed. The high-temperature of the targets could have affected the impact excavation process by causing higher ejection angles, giving rise to more scattered circular secondaries.

[1] Twenty confirmed impacts over a 7-year time period on Mars were qualitatively and statistically compared to 287 secondary craters believed to originate from Zunil, an 090804500 ka, 10-km diameter, primary crater. Our goal was to establish criteria to distinguish secondaries from primaries in the general crater population on the basis of their horizontal planforms. Recent primary impacts have extensive 090008air blast090009 zones, distal ray systems (>100 crater radii, R), and ephemeral ejecta. Recent primaries formed clusters of craters from atmospheric fragmentation of the meteoroid body. Secondary craters have ejecta blankets with shorter rays that are consistent with emplacement by low-impact velocities (near 1 km/s). The mean extent of the continuous ejecta blankets was less distal for secondaries (5.38 00± 1.57R) versus primaries (18.07 00± 7.01R), though primary ejecta were less fractal (Fractal Dimension Index (FDI) < 1.30) and more circular on average (Circularity Ratio (CR) = 0.55 00± 0.25 versus 0.27 00± 0.13 for secondaries). Crater rims were remarkably circular (primaries CR = 0.97 00± 0.02, secondaries at 0.94 00± 0.05), though secondaries have the lowest values (CR < 0.9). Secondary crater rims were elongated toward or orthogonal to their primary of origin. Uprange source directions for most secondaries, determined by ejecta planform and crater rim ellipticity, point toward Zunil, although contamination from other primaries is considered in some areas. Ejecta blanket discrepancies between recent primaries and Zunil secondaries are attributable to differences in impact velocity and retention age. After removal of the ejecta blanket, crater rims are generally not diagnostic for determining crater origin. Fragmentation of primaries may play some role in steepening the size-frequency distribution of crater diameters in the 5 m < D < 30 m range.

The crater size-frequency distribution curve for the lunar surface has a bending point near 4 km in diameter. It is suggested that the steepening in the smaller diameter range is caused by the effects of secondary craters and/or size-frequency distribution of the impactors. In order to extract the fraction of the primary craters created by Near Earth Objects, among the primary-secondary mixture of km-size craters, we have examined the crater shape statistics using Clementine images of the lunar surface. We found that the size-frequency distribution of the craters with ellipticity smaller than 1.2 can be considered primary craters generated by impactors from interplanetary space.

The Mars Exploration Rovers investigated numerous craters in Gusev crater and Meridiani Planum during the first ~400 sols of their missions. Craters vary in size and preservation state but are mostly due to secondary impacts at Gusev and primary impacts at Meridiani. Craters at both locations are modified primarily by eolian erosion and infilling and lack evidence for modification by aqueous processes. Effects of gradation on crater form are dependent on size, local lithology, slopes, and availability of mobile sediments. At Gusev, impacts into basaltic rubble create shallow craters and ejecta composed of resistant rocks. Ejecta initially experience eolian stripping, which becomes weathering-limited as lags develop on ejecta surfaces and sediments are trapped within craters. Subsequent eolian gradation depends on the slow production of fines by weathering and impacts and is accompanied by minor mass wasting. At Meridiani the sulfate-rich bedrock is more susceptible to eolian erosion, and exposed crater rims, walls, and ejecta are eroded, while lower interiors and low-relief surfaces are increasingly infilled and buried by mostly basaltic sediments. Eolian processes outpace early mass wasting, often produce meters of erosion, and mantle some surfaces. Some small craters were likely completely eroded/buried. Craters >100 m in diameter on the Hesperian-aged floor of Gusev are generally more pristine than on the Amazonian-aged Meridiani plains. This conclusion contradicts interpretations from orbital views, which do not readily distinguish crater gradation state at Meridiani and reveal apparently subdued crater forms at Gusev that may suggest more gradation than has occurred.

A catalog of crater dimensions that were compiled mostly from the new Apollo-based Lunar Topographic Orthophotomaps is presented in its entirety. Values of crater diameter, depth, rim height, flank width, circularity, and floor diameter (where applicable) are tabulated for a sample of 484 craters on the Moon and 22 craters on Earth. Systematic techniques of mensuration are detailed. The lunar craters range in size from 400 m to 300 km across and include primary impact craters of the main sequence, secondary impact craters, craterlets atop domes and cones, and dark-halo craters. The terrestrial craters are between 10 m and 22.5 km in diameter and were formed by meteorite impact.

Continuous ejecta deposits on Ganymede consist of two major units, or facies: a thick inner hummocky pedestal facies, and a relatively thin outer radially scoured facies defined also by the inner limit of the secondary crater field. Both ejecta facies have a well-defined power-law relationship to crater diameter for craters ranging from 15 to ~600 km across. This relationship can be used to estimate the nominal crater diameter for impact features on icy satellites (such as palimpsests and multiring basins) for which the crater rim is no longer recognizable. Ejecta deposits have also been mapped on 4 other icy satellites. Although morphologically similar to eject deposits on the Moon, ejecta deposits for smaller craters are generally significantly broader in extent on the icy satellites, in apparent defiance of predictions of self-similarity. A greater degree of rim collapse and enlargement on the Moon may explain the observed difference.

We find evidence, by both observation and analysis, that primary crater ejecta play an important role in the small crater (less than a few km) populations on the saturnian satellites, and more broadly, on cratered surfaces throughout the Solar System. We measure crater populations in Cassini images of Enceladus, Rhea, and Mimas, focusing on image data with scales less than 500m/pixel. We use recent updates to crater scaling laws and their constants (Housen, K.R., Holsapple, K.A. [2011]. Icarus 211, 856–875) to estimate the amount of mass ejected in three different velocity ranges: (i) greater than escape velocity, (ii) less than escape velocity and faster than the minimum velocity required to make a secondary crater (vmin), and (iii) velocities less than vmin. Although the vast majority of mass on each satellite is ejected at speeds less than vmin, our calculations demonstrate that the differences in mass available in the other two categories should lead to observable differences in the small crater populations; the predictions are borne out by the measurements we have made to date. In particular, Rhea, Tethys, and Dione have sufficient surface gravities to retain ejecta moving fast enough to make secondary crater populations. The smaller satellites, such as Enceladus but especially Mimas, are expected to have little or no traditional secondary populations because their escape velocities are near the threshold velocity necessary to make a secondary crater. Our work clarifies why the Galilean satellites have extensive secondary crater populations relative to the saturnian satellites. The presence, extent, and sizes of sesquinary craters (craters formed by ejecta that escape into temporary orbits around Saturn before re-impacting the surface, see Dobrovolskis, A.R., Lissauer, J.J. [2004]. Icarus 169, 462–473; Alvarellos, J.L., Zahnle, K.J., Dobrovolskis, A.R., Hamill, P. [2005]. Icarus 178, 104–123; Zahnle, K., Alvarellos, J.L., Dobrovolskis, A.R., Hamill, P. [2008]. Icarus 194, 660–674) is not yet well understood. Finally, our work provides further evidence for a “shallow” size–frequency distribution (slope index of 652 for a differential power-law) for comets a few kilometers diameter and smaller.

We review the radiometric ages of the 16 currently known Martian meteorites, classified as 11 shergottites (8 basaltic and 3 lherzolitic), 3 nakhlites (clinopyroxenites), Chassigny (a dunite), and the orthopyroxenite ALH84001. The basaltic shergottites represent surface lava flows, the others magmas that solidified at depth. Shock effects correlate with these compositional types, and, in each case, they can be attributed to a single shock event, most likely the meteorite's ejection from Mars. Peak pressures in the range 15 – 45 GPa appear to be a "launch window": shergottites experienced 6530 – 45 GPa, nakhlites 6520 ± 5 GPa, Chassigny 6535 GPa, and ALH84001 6535 – 40 GPa. Two meteorites, lherzolitic shergottite Y-793605 and orthopyroxenite ALH84001, are monomict breccias, indicating a two-phase shock history in toto : monomict brecciation at depth in a first impact and later shock metamorphism in a second impact, probably the ejection event.Crystallization ages of shergottites show only two pronounced groups designated S 1 (65175 Myr), including 4 of 6 dated basalts and all 3 lherzolites, and S 2 (330 – 475 Myr), including two basaltic shergottites and probably a third according to preliminary data. Ejection ages of shergottites, defined as the sum of their cosmic ray exposure ages and their terrestrial residence ages, range from the oldest (6520 Myr) to the youngest (650.7 Myr) values for Martian meteorites. Five groups are distinguished and designated S Dho (one basalt, 6520 Myr), S L (two lherzolites of overlapping ejection ages, 3.94 ± 0.40 Myr and 4.70 ± 0.50 Myr), S (four basalts and one lherzolite, 652.7 – 3.1 Myr), S DaG (two basalts, 651.25 Myr), and S E (the youngest basalt, 0.73 ± 0.15 Myr). Consequently, crystallization age group S 1 includes ejection age groups S L , S E and 4 of the 5 members of S, whereas S 2 includes the remaining member of S and one of the two members of S DaG . Shock effects are different for basalts and lherzolites in group S/S 1 . Similarities to the dated meteorite DaG476 suggest that the two shergottites that are not dated yet belong to group S 2 . Whether or not S 2 is a single group is unclear at present. If crystallization age group S 1 represents a single ejection event, pre-exposure on the Martian surface is required to account for ejection ages of S L that are greater than ejection ages of S, whereas secondary breakup in space is required to account for ejection ages of S E less than those of S. Because one member of crystallization age group S 2 belongs to ejection group S, the maximum number of shergottite ejection events is 6, whereas the minimum number is 2.Crystallization ages of nakhlites and Chassigny are concordant at 651.3 Gyr. These meteorites also have concordant ejection ages, i.e., they were ejected together in a single event (NC). Shock effects vary within group NC between the nakhlites and Chassigny.The orthopyroxenite ALH84001 is characterized by the oldest crystallization age of 654.5 Gyr. Its secondary carbonates are 653.9 Gyr old, an age corresponding to the time of Ar-outgassing from silicates. Carbonate formation appears to have coincided with impact metamorphism, either directly, or indirectly, perhaps via precipitation from a transient impact crater lake.The crystallization age and the ejection age of ALH84001, the second oldest ejection age at 15.0 ± 0.8 Myr, give evidence for another ejection event (O). Consequently, the total number of ejection events for the 16 Martian meteorites lies in the range 4 – 8.The Martian meteorites indicate that Martian magmatism has been active over most of Martian geologic history, in agreement with the inferred very young ages of flood basalt flows observed in Elysium and Amazonis Planitia with the Mars Orbital Camera (MOC) on the Mars Global Surveyor (MGS). The provenance of the youngest meteorites must be found among the youngest volcanic surfaces on Mars, i.e., in the Tharsis, Amazonis, and Elysium regions.

Abstract—In oblique impacts with an impact angle under 45°, the bilateral shape of the distal ejecta blanket is used as the strongest indicator for an impact vector. This bilateral symmetry is attenuated and is superimposed by radial symmetry towards the crater rim, which remains circular for impact angles down to 10–15°. The possibility that remnants of bilateral symmetry might still be present in the most proximal ejecta, the overturned flap and the crater rim was explored with the intention of deducing an impact vector. A model is presented that postulates bilateral patterns using proximal ejecta trajectories and predicts these patterns in the orientation of bedding planes in the crater rim. This model was successfully correlated to patterns described by radial grooves in the proximal ejecta blanket of the oblique Tooting crater on Mars. A new method was developed to detect structural asymmetries by converting bedding data into values that express the deviation from concentric strike orientation in the crater rim relative to the crater center, termed "concentric deviation." The method was applied to field data from Wolfe Creek crater, Western Australia. Bedding in the overturned flap implies an impactor striking from the east, which refines earlier publications, while bedding from the inner rim shows a correlation with the crater rim morphology.

The lunar red spots, Helmet, Hansteen Alpha, and the NW quadrant of the crater Copernicus, were selected for a complex comparative investigation of their characteristics measured by the spacecraft Clementine, LRO, and Chandrayaan-1. For the analysis we used the following parameters: the reflectanceA(750nm), color-ratioA(750nm)/A(415nm), parameter of optical micro-roughness (LRO WAC), parameters deduced from LRO Diviner data, optical maturityOMAT, abundance of FeO and TiO2(Clementine UVVIS and LRO WAC data), oxygen content determined using Lunar Prospector data, and parameters characterizing the 0.95- m and 2.2- m bands of Fe2+ions (crystal field bands), and 2.8- m band of H2O/OH and/or Fe2+ions. The red spots Helmet and Hansteen Alpha are considered to be extrusions of rhyolite composition, which can be attributed to the Nectarian period; we did not find contradictions of this assumption. As for the Copernicus red spot, this, perhaps, is a similar formation that has been destroyed by the impact. We demonstrate that the material of the Copernicus red spot probably has the same composition as the classical red spots Helmet and Hansteen Alpha. The distributions of the parameter of optical micro-roughness and optical maturityOMATshow that the Copernicus red anomaly was not formed during the long evolution of the lunar surface, but results from crater formation. We find several confirmations of the hypothesis that the Copernicus red spot can be a residual of a red material (possibly rhyolite) extrusion that was involved in the impact process. The red material could have been partially melted, crushed, and ejected to the crater's north-western vicinity. The described red asymmetry of the Copernicus ejecta can be related to the eccentricity, relative to the extrusion, of the impact and/or to the inclination of the impactor trajectory. The latter also is confirmed by an analysis of the region, which is based on the geological map shown in this paper.

The Moon Mineralogy Mapper (M3) acquired high spatial and spectral resolution data of the Aristarchus Plateau with 140 m/pixel in 85 spectral bands from 0.43 to 3.0 m. The data were collected as radiance and converted to reflectance using the observational constraints and a solar spectrum scaled to the Moon-Sun distance. Summary spectral parameters for the area of mafic silicate 1 and 2 m bands were calculated from the M3 data and used to map the distribution of key units that were then analyzed in detail with the spectral data. This analysis focuses on five key compositional units in the region. (1) The central peaks are shown to be strongly enriched in feldspar and are likely from the upper plagioclase-rich crust of the Moon. (2) The impact melt is compositionally diverse with clear signatures of feldspathic crust, olivine, and glass. (3) The crater walls and ejecta show a high degree of spatial heterogeneity and evidence for massive breccia blocks. (4) Olivine, strongly concentrated on the rim, wall, and exterior of the southeastern quadrant of the crater, is commonly associated the impact melt. (5) There are at least two types of glass deposits observed: pyroclastic glass and impact glass. Copyright 2011 by the American Geophysical Union.

[1] Kepler is a 31 km diameter Copernican age complex impact crater located on the nearside maria of the Moon. We used Lunar Reconnaissance Orbiter imagery and topographic data in combination with Kaguya terrain camera and other image data sets to construct a new geomorphologic sketch map of the Kepler crater, with a focus on impact melt-rich lithologies. Most of the interior melt rocks are preserved in smooth and hummocky floor materials. Smaller volumes of impact melt were deposited in rim veneer, interior and exterior ponds, and lobe-like overlapping flows on the upper crater wall. Based on shadow lengths, typical flows of melt-rich material on crater walls and the western rim flank are 09080410900095 m thick, and have yield strengths of 090804109000910 kPa. The melt rock distribution is notably asymmetric, with interior and exterior melt-rich deposits concentrated north and west of the crater center. This melt distribution and the similarly asymmetric ray distribution imply a slightly less than 4500° impact trajectory from the southeast. The exposed wall of Kepler displays distinct layering, with individual layers having typical thicknesses of 09080430900095 m. These are interpreted as flows of Procellarum mare basalts in the impact target. From the point of view of exploration, numerous fractures and pits in the melt-rich floor materials not only enable detailed studies of melt-related processes of impact crater formation, but also provide potential shelters for longer duration manned lunar missions.