The effect of the oceans on the terrestrial crater size‐frequency distribution: Insight from numerical modeling

Abstract— On Earth, oceanic impacts are twice as likely to occur as continental impacts, yet the effect of the oceans has not been previously considered when estimating the terrestrial crater size‐frequency distribution. Despite recent progress in understanding the qualitative and quantitative effect of a water layer on the impact process through novel laboratory experiments, detailed numerical modeling, and interpretation of geological and geophysical data, no definitive relationship between impactor properties, water depth, and final crater diameter exists. In this paper, we determine the relationship between final (and transient) crater diameter and the ratio of water depth to impactor diameter using the results of numerical impact models. This relationship applies for normal incidence impacts of stoney asteroids into water‐covered, crystalline oceanic crust at a velocity of 15 km s^?1. We use these relationships to construct the first estimates of terrestrial crater size‐frequency distributions (over the last 100 million years) that take into account the depth‐area distribution of oceans on Earth. We find that the oceans reduce the number of craters smaller than 1 km in diameter by about two‐thirds, the number of craters ?30 km in diameter by about one‐third, and that for craters larger than ?100 km in diameter, the oceans have little effect. Above a diameter of ?12 km, more craters occur on the ocean floor than on land; below this diameter more craters form on land than in the oceans. We also estimate that there have been in the region of 150 impact events in the last 100 million years that formed an impact‐related resurge feature, or disturbance on the seafloor, instead of a crater.     

The record of impact cratering on the great volcanic shields of the Tharsis region of Mars

Mariner 9 images of the four great volcanic shields of the Tharsis region of Mars show many circular craters ranging in diameter from 100mm to 20 km. Previous attempts to date the volcanoes from their apparent impact crater densities yielded a range of results. The principal difficulty is sorting volcanic from impact craters for diameters $\text{?1}\text{km}$. Many of the observed craters are aligned in prominent linear and concentric patterns suggestive of volcanic origin. In this paper an attempt is made to date areas of shield surface, covered with high resolution images using only scattered small (?1 km) craters of probable impact origin. Craters of apparent volcanic origin are systematically excluded from the dating counts.The common measure of age, deduced for all surfaces studied, is a calculated “crater age” F′ defined as the number of craters equal to or larger than 1 km in diameter per 10⁶km². The conclusions reached from comparing surface ages and their geological settings are: (1) Lava flow terrain surfaces with ages, F′, from 180 to 490 are seen on the four great volcanoes. Summit surfaces of similar ages, F′ = 360 to 420, occur on the rims of calderas of Arsia Mons, Pavonis Mons, and Olympus Mons. The summit of Ascraeus Mons is possibly younger; F′ is calculated to be 180 for the single area which could be dated. (2) One considerably younger surface, F′ < 110, is seen on the floor of Arsia Mon's summit caldera. (3) Nearly crater free lava flow terrain surfaces seen on Olympus Mons are estimated to be less than half the age of a summit surface. The summit caldera floor is similarly young. (4) The pattern of surface ages on the volcanoes suggests that their eruption patterns are similar to those of Hawaiian basaltic shields. The youngest surfaces seem concentrated on the mid-to-lower flanks and within the summit calderas. (5) The presently imaged sample of shield surface, though incomplete, clearly shows a broad range of ages on three volcanoes—Olympus, Arsia, and Pavonis Mons.Estimated absolute ages of impact dated surfaces are obtained from two previously published estimates of the history of flux of impacting bodies on Mars. The estimated ranges of age for the observed crater populations are 0.5 to 1.2b.y. and 0.07 to 0.2b.y. Areas which are almost certainly younger, less than 0.5 or 0.07b.y., are also seen. The spans of surface age derived for the great shields are minimum estimates of their active lifetimes, apparently very long compared to those of terrestrial volcanoes.     

Boltysh,another end‐Cretaceous impact

Abstract— The Chixculub impact occurred at the Cretaceous/Tertiary (K/T) boundary, and although several other Late Cretaceous and Paleogene impact craters have, at times, been linked with the K/T boundary, isotope geochronology has demonstrated that all have significantly different ages. The currently accepted age of the 24 km diameter Boltysh crater, a K‐Ar whole‐rock age, places it in the Coniacian at 88 ± 3 Ma. However, comprehensive Ar‐Ar dating of a range of melt samples yields a mean age of 65.17 ± 0.64 Ma, within errors of the K/T boundary. Several of the fresh samples exhibit signs of excess argon but this seems to be concentrated in rapidly crystallized glass‐rich samples. The Ar‐Ar age confirms an earlier fission track measurement and thus two dating techniques have yielded an age within errors of the K/T boundary for this crater. Crucially, although the ages of Boltysh and Chixculub are within errors, they may not have formed synchronously. Craters of 24 km diameter occur much more commonly than impacts of Chixculub dimensions, but their proximity does raise the important question of how many impacts there might have been close to the K/T boundary.     

Floor-fractured crater models of the Sudbury Structure,Canada: Implications for initial crater size and crater modification

Abstract The pattern of radial and concentric offset dikes at Sudbury strongly resembles fracture patterns in certain volcanically modified craters on the Moon. Since the Sudbury dikes apparently formed shortly after the impact event, this resemblance suggests that early endogenic modification at Sudbury was comparable to deformation in lunar floor-fractured craters. Although regional deformation has obscured many details of the Sudbury Structure, such a comparison of Sudbury with lunar floor-fractured craters provides two alternative models for the original size and surface structures of the Sudbury basin. First, the Sudbury date pattern can be correlated with fractures in the central peak crater Haldane (36 km in diameter). This comparison indicates an initial Sudbury diameter of between 100 and 140 km but requires loss of a central peak complex for which there is little evidence. Alternatively, comparison of the Sudbury dikes with fractures in the two-ring basin Schrödinger indicates an initial Sudbury diameter of at least ～ 180 km, which is in agreement with other recent estimates for the size of the Sudbury Structure. In addition to constraining the size and structure of the original Sudbury crater, these comparisons also suggest that crater modification may reflect different deformation mechanisms at different sizes. Most lunar floor-fractured craters are attributed to deformation over a shallow, crater-centered intrusion; however, there is no evidence for such an intrusion at Sudbury. Instead, melts from the evolving impact melt sheet probably entered fractures formed by isostatically-induced flexure of the crater floor. Since most of the lunar floor-fractured craters are too small (<100-km diameter) to induce significant isostatic adjustment, crater modification by isostatic uplift apparently is limited to only the largest of craters, whereas deformation over igneous intrusions dominates the modification of smaller craters.     

Central pit craters on Ganymede

Central pit craters are common on Mars, Ganymede and Callisto, and thus are generally believed to require target volatiles in their formation. The purpose of this study is to identify the environmental conditions under which central pit craters form on Ganymede. We have conducted a study of 471 central pit craters with diameters between 5 and 150 km on Ganymede and compared the results to 1604 central pit craters on Mars (diameter range 5-160 km). Both floor and summit pits occur on Mars whereas floor pits dominate on Ganymede. Central peak craters are found in similar locations and diameter ranges as central pit craters on Mars and overlap in location and at diameters <60 km on Ganymede. Central pit craters show no regional variations on either Ganymede or Mars and are not concentrated on specific geologic units. Central pit craters show a range of preservation states, indicating that conditions favoring central pit formation have existed since crater-retaining surfaces have existed on Ganymede and Mars. Central pit craters on Ganymede are generally about three times larger than those on Mars, probably due to gravity scaling although target characteristics and resolution also may play a role. Central pits tend to be larger relative to their parent crater on Ganymede than on Mars, probably because of Ganymede’s purer ice crust. A transition to different characteristics occurs in Ganymede’s icy crust at depths of 4-7 km based on the larger pit-to-crater-diameter relationship for craters in the 70-130-km-diameter range and lack of central peaks in craters larger than 60-km-diameter. We use our results to constrain the proposed formation models for central pits on these two bodies. Our results are most consistent with the melt-drainage model for central pit formation.     

Shallow and deep fresh impact craters in Hesperia Planum,Mars

The depths of 109 impact craters 2–16 km in diameter, located on the ridged plains materials of Hesperia Planum, Mars, have been measured from their shadow lengths using digital Viking Orbiter images (orbit numbers 417S–419S) and the PICS computer software. On the basis of their pristine morphology (very fresh lobate ejecta blankets, well preserved rim crests, and lack of superposed impact craters), 57 of these craters have been selected for detailed analysis of their spatial distribution and geometry. We find that south of 30°S, craters <6.0 km in diameter are markedly shallower than similar-sized craters equatorward of this latitude. No comparable relationship is observed for morphologically fresh craters >6.0 km diameter. We also find that two populations exist for older craters <6.0 km diameter. When craters that lack ejecta blankets are grouped on the basis of depth/diameter ratio, the deeper craters also typically lie equatorward of 30° S. We interpret the spatial variation in crater depth/diameter ratios as most likely due to a poleward increase in volatiles within the top 400 m of the surface at the times these craters were formed.     

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.     

Impact penetration of Europa's ice crust as a mechanism for formation of chaos terrain

Abstract— Ice thickness estimates and impactor dynamics indicate that some impacts must breach Europa's ice crust; and outcomes of impact experiments using ice‐over‐water targets range from simple craters to chaos‐like destroyed zones, depending on impact energy and ice competence. First‐order impacts‐into thick ice or at low impact energy‐produce craters. Second‐order impacts punch through the ice, making holes that resemble raft‐free chaos areas. Third‐order impacts‐into thinnest ice or at highest energy‐produce large irregular raft‐filled zones similar to platy chaos. Other evidence for an impact origin for chaos areas comes from the size‐frequency distribution of chaos+craters on Europa, which matches the impact production functions of Ganymede and Callisto; and from small craters around the large chaos area Thera Macula, which decrease in average size and density per unit area as a function of distance from Thera's center. There are no tiny chaos areas and no craters >50 km diameter. This suggests that small impactors never penetrate, whereas large ones (ÜberPenetrators: >2.5 km diameter at average impact velocity) always do. Existence of both craters and chaos areas in the size range 2–40 km diameter points to spatial/temporal variation in crust thickness. But in this size range, craters are progressively outnumbered by chaos areas at larger diameters, suggesting that probability of penetration increases with increasing scale of impact. If chaos areas do represent impact sites, then Europa's surface is older than previously thought. The recalculated resurfacing age is 480 (‐302/+960) Ma: greater than prior estimates, but still very young by solar system standards.     

The planforms of low‐angle impact craters in the northern hemisphere of Mars

Abstract— We have surveyed Martian impact craters greater than 5 km in diameter using Viking and thermal emission imaging system (THEMIS) imagery to evaluate how the planform of the rim and ejecta changes with decreasing impact angle. We infer the impact angles at which the changes occur by assuming a sin²θ dependence for the cumulative fraction of craters forming below angle θ. At impact angles less than ?40° from horizontal, the ejecta become offset downrange relative to the crater rim. As the impact angle decreases to less than ?20°, the ejecta begin to concentrate in the cross‐range direction and a “forbidden zone” that is void of ejecta develops in the uprange direction. At angles less than ?10°, a “butterfly” ejecta pattern is generated by the presence of downrange and uprange forbidden zones, and the rim planform becomes elliptical with the major axis oriented along the projectile's direction of travel. The uprange forbidden zone appears as a “V” curving outward from the rim, but the downrange forbidden zone is a straight‐edged wedge. Although fresh Martian craters greater than 5 km in diameter have ramparts indicative of surface ejecta flow, the ejecta planforms and the angles at which they occur are very similar to those for lunar craters and laboratory impacts conducted in a dry vacuum. The planforms are different from those for Venusian craters and experimental impacts in a dense atmosphere. We interpret our results to indicate that Martian ejecta are first emplaced predominantly ballistically and then experience modest surface flow.     

Lunar craters evolution and meteoroidal flux in pre-mare and post-mare times

In order to study the geomorphic evolution and lifetimes of lunar craters, data were collected from (i) 32mare andterra provinces of the nearside of the Moon using the L.P.L. catalog; (ii) amare area in Sinus Medii, using direct observations of Lunar Orbiter photos, and (iii) aterra area on the farside using direct observations of Zond-8 photos. The theory presented in a previous publication is expanded and applied to the data.The following conclusions are obtained. (1) Steady-state conditions occur on the studiedmare surfaces for craters of diameter up to approximately 220 m, and on the studiedterra surfaces for craters of diameter up to at least 50 km. (2) The average lifetime of a crater, in addition of being a function of the meteoroidal flux, is a steep function of the diameter of the crater. (3) The correlation is good between a geomorphic classification of craters based on visual comparison with standard craters and a classification of craters based on their depth-diameter ratio, resulting in a coefficient of rank correlation of 0.64. (4) When craters are classified as young, mature, and old, the length of time spent as young is less than a few percent of the total lifetime of the crater; the time spent as mature is 10 to 30%; and as much as 80% is spent as an old crater. Within the error of the calculations, these values are independent of crater diameter and apply to both pre-mare and post-mare craters, indicating that they are also independent of the intensity of the meteoroidal flux. (5) The average lifetime of a 50 km crater in pre-mare times is estimated to be less than 0.3×10⁹ years. (6) The average lifetime of a 50 km crater in post-mare times is estimated to be between 3×10¹¹ and 10¹⁴ years. (7) The average meteoroidal flux in pre-mare times is estimated to be three to six orders of magnitude more intense than in post-mare times.     

Petrogenesis of lunar impact melt rock meteorite Oued Awlitis 001

Oued Awlitis 001 is a highly feldspathic, moderately equilibrated, clast‐rich, poikilitic impact melt rock lunar meteorite that was recovered in 2014. Its poikilitic texture formed due to moderately slow cooling, which judging from textures of rocks in melt sheets of terrestrial impact structures, is observed in impact melt volumes at least 100 m thick. Such coherent impact melt volumes occur in lunar craters larger than ~50 km in diameter. The composition of Oued Awlitis 001 points toward a crustal origin distant from incompatible‐element‐rich regions. Comparison of the bulk composition of Oued Awlitis 001 with Lunar Prospector 5° γ‐ray spectrometer data indicates a limited region of matches on the lunar farside. After its initial formation in an impact crater larger than ~50 km in diameter, Oued Awlitis 001 was excavated from a depth greater than ~50 m. The cosmogenic nuclide inventory of Oued Awlitis 001 records ejection from the Moon 0.3 Ma ago from a depth of at least 4 m and little mass loss due to ablation during its passage through Earth's atmosphere. The terrestrial residence time must have been very short, probably less than a few hundred years; its exact determination was precluded by a high concentration of solar cosmic ray‐produced ¹⁴C. If the impact that excavated Oued Awlitis 001 also launched it, this event likely produced an impact crater >10 km in diameter. Using petrologic constraints and Lunar Reconnaissance Orbiter Camera and Diviner data, we test Giordano Bruno and Pierazzo as possible launch craters for Oued Awlitis 001.     

Crater size-shape profiles for the moon and mercury: Terrain effects and interplanetary comparisons

New crater size-shape data were compiled for 221 fresh lunar craters and 152 youthful mercurian craters. Terraces and central peaks develop initially in fresh craters on the Moon in the 0–10 km diameter interval. Above a diameter of 65 km all craters are terraced and have central peaks. Swirl floor texture is most common in craters in the size range 20–30 km, but it occurs less frequently as terraces become a dominant feature of crater interiors. For the Moon there is a correlation between crater shape and geomorphic terrain type. For example, craters on the maria are more complex in terms of central peak and terrace detail at any given crater diameter than are craters in the highlands. These crater data suggest that there are significant differences in substrate and/or target properties between maria and highlands. Size-shape profiles for Mercury show that central peak and terrace onset is in the 10–20 km diameter interval; all craters are terraced at 65 km, and all have central peaks at 45 km. The crater data for Mercury show no clear cut terrain correlation. Comparison of lunar and mercurian data indicates that both central peaks and terraces are more abundant in craters in the diameter range 5–75 km on Mercury. Differences in crater shape between Mercury and the Moon may be due to differences in planetary gravitational acceleration (gMercury=2.3gMoon). Also differences between Mercury and the Moon in target and substrate and in modal impact velocity may contribute to affect crater shape.     

Doublet craters on Venus

Of the impact craters on Earth larger than 20 km in diameter, 10-15% (3 out of 28) are doublets, having been formed by the simultaneous impact of two well-separated projectiles. The most likely scenario for their formation is the impact of well-separated binary asteroids. If a population of binary asteroids is capable of striking the Earth, it should also be able to hit the other terrestrial planets as well. Venus is a promising planet to search for doublet craters because its surface is young, erosion is nearly nonexistent, and its crater population is significantly larger than the Earth's. After a detailed investigation of single craters separated by less than 150 km and “multiple” craters having diameters greater than 10 km, we found that the proportion of doublet craters on Venus is at most 2.2%, significantly smaller than Earth's, although several nearly incontrovertible doublets were recognized. We believe this apparent deficit relative to the Earth's doublet population is a consequence of atmospheric screening of small projectiles on Venus rather than a real difference in the population of impacting bodies. We also examined “splotches,” circular radar reflectance features in the Magellan data. Projectiles that are too small to form craters probably formed these features. After a careful study of these patterns, we believe that the proportion of doublet splotches on Venus (14%) is comparable to the proportion of doublet craters found on Earth (10-15%). Thus, given the uncertainties of interpretation and the statistics of small numbers, it appears that the doublet crater population on Venus is consistent with that of the Earth.     

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.     

HOLLOW LUNAR SPHERULES AND MICROCRATERING

Since thin-walled hollow glass spherules exist in the lunar regolith and perhaps as a component of cosmic dust, laboratory simulations of impacts by and upon such spherules were done to determine identifying features of the resulting craters and perforations. The targets were soda-lime glass, stainless steel, and hollow glass beads. Craters were generated in the first two targets by the normal impact of thin-walled hollow glass spheres with masses and velocities between eight and 240 pg and 1.8 and 10 km/s, respectively. With increasing impact velocity, the crater morphology in glass progresses as follows: 1, a dent; 2, a narrow lip around the depression; and 3, spallation around the pit that may carry away all of part of the lip. The craters differ from those formed by solid spherical projectiles in that the central pit is an annular rather than a cup-shaped depression. The craters in steel display a typical outer lip and an additional concentric inner lip which is subdued to an annular mound as the impact velocity increases. In both targets, shattered remnants of the projectiles remain in the craters at low impact velocities. At higher velocities, melting of the projectile material occurs. The annular features distinguish these craters from craters generated by solid spheres or irregular projectiles', and the existence of such a crater morphology on a surface exposed to cosmic dust would indicate the presence of thin-walled hollow spherules. Contrary to common opinion, hollow spheres do not adequately simulate cratering by low density materials because of the mass distribution. Penetrations of thin-walled hollow glass beads by high velocity, solid, micrometer-size spheres are characterized by inward and outward flowing lips that show asymmetries dependent on the angle of impact. The morphology is sufficient to discriminate against other mechanisms that cause perforations in the one to 10 μm size range in hollow lunar spherules. The identifying lip may break away by fragmentation in the impact of larger size projectiles.     

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.     

Impact craters on Titan

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

Catalogue of large alluvial fans in martian impact craters

We conducted a systematic, global survey using Thermal Emission Imaging System Infrared (THEMIS IR) coverage (∼100 m/pixel) to search for large alluvial fans in impact craters on Mars. Our survey has focused on large fans (apron areas greater than ∼40 km², usually located in craters greater than 20 km in diameter) due to the resolution of the THEMIS images and Mars Orbiter Laser Altimeter (MOLA) coverage. We find that the host craters are found to have a distinctive diameter range from 30-150 km. The fans generally cluster in three geographic areas—southern Margaritifer Terra, southwestern Terra Sabaea, and southwestern Tyrrhena Terra, however several outliers do exist. The alluvial fans do not form in a particular orientation along the crater rim nor are they associated with the location of current high rim topography. Fan area magnitude and variability increase with crater diameter while fan concavity magnitude and variability increase with decreasing crater diameter. Smaller fan aprons in general have higher, more variable concavity. The source of the water forming these fans is uncertain given the challenges of accommodating the global distribution pattern and formation patterns within the craters.     

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.     

Crater geometry and ejecta thickness of the Martian impact crater Tooting

Abstract— We use Mars Orbiter Laser Altimeter (MOLA) topographic data and Thermal Emission Imaging System (THEMIS) visible (VIS) images to study the cavity and the ejecta blanket of a very fresh Martian impact crater ?29 km in diameter, with the provisional International Astronomical Union (IAU) name Tooting crater. This crater is very young, as demonstrated by the large depth/diameter ratio (0.065), impact melt preserved on the walls and floor, an extensive secondary crater field, and only 13 superposed impact craters (all 54 to 234 meters in diameter) on the ?8120 km² ejecta blanket. Because the pre‐impact terrain was essentially flat, we can measure the volume of the crater cavity and ejecta deposits. Tooting crater has a rim height that has >500 m variation around the rim crest and a very large central peak (1052 m high and >9 km wide). Crater cavity volume (i.e., volume below the pre‐impact terrain) is ?380 km³ the volume of materials above the pre‐impact terrain is ?425 km³. The ejecta thickness is often very thin (<20 m) throughout much of the ejecta blanket. There is a pronounced asymmetry in the ejecta blanket, suggestive of an oblique impact, which has resulted in up to ?100 m of additional ejecta thickness being deposited down‐range compared to the up‐range value at the same radial distance from the rim crest. Distal ramparts are 60 to 125 m high, comparable to the heights of ramparts measured at other multi‐layered ejecta craters. Tooting crater serves as a fresh end‐member for the large impact craters on Mars formed in volcanic materials, and as such may be useful for comparison to fresh craters in other target materials.