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Mairan middle dome (MMD), a lunar “red spot” of silicic composition, and the surrounding maria were emplaced in the same two major episodes of volcanism. Both episodes at MMD included eruptions of low-FeO, silica-rich lava, while basaltic lava flooded the surrounding terrain during these episodes. MMD is a composite of, at least, seven small volcanic edifices. Crater counts suggest that the first episode occurred at ∼3.75 ± 0.1 Ga when low FeO, high-silica lavas erupted at MMD, and Mairan T, the small dome 11 km northwest of MMD. At about the same time, basaltic composition lava erupted southeast of MMD. A second major episode of volcanism at MMD occurred at ∼3.35 ± 0.2 Ga when low FeO, and high-silica lavas erupted at the summits of individual small volcanic edifices and a central plateau area between them. During this phase, mare basaltic lavas again flooded the area surrounding MMD and Mairan T. This sequence of events indicates that the emplacement of MMD is more complex than previously thought. In addition, the simultaneous eruption of basaltic composition lavas and low FeO, high-SiO2 lavas in this region supports the underplating model for production of magma to form the “red spots” volcanic complexes on the Moon.

We use topographic data to show that impact craters with pitted floor deposits are among the deepest on Mars. This is consistent with the interpretation of pitted materials as primary crater-fill impactite deposits emplaced during crater formation. Our database consists of 224 pitted material craters ranging in size from ∼1 to 150 km in diameter. Our measurements are based on topographic data from the Mars Orbiter Laser Altimeter (MOLA) and the High-Resolution Stereo Camera (HRSC). We have used these craters to measure the relationship between crater diameter and the initial post-formation depth. Depth was measured as maximum rim-to-floor depth, ( d r ), but we also report the depth measured using other definitions. The database was down-selected by refining or removing elevation measurements from “problematic” craters affected by processes and conditions that influenced their d r /D, such as pre-impact slopes/topography and later overprinting craters. We report a maximum (deepest) and mean scaling relationship of d r = ( 0.347 ± 0.021 ) D 0.537 ± 0.017 and d r = ( 0.323 ± 0.017 ) D 0.538 ± 0.016 , respectively. Our results suggest that significant variations between previously-reported MOLA-based d r vs. D relationships may result from the inclusion of craters that: 1) are influenced by atypical processes ( e.g. , highly oblique impact), 2) are significantly degraded, 3) reside within high-strength regions, and 4) are transitional (partially collapsed). By taking such issues into consideration and only measuring craters with primary floor materials, we present the best estimate to date of a MOLA-based relationship of d r vs. D for the least-degraded complex craters on Mars. This can be applied to crater degradation studies and provides a useful constraint for models of complex crater formation.

Detailed morphologic and morphometric analysis of the radial grooves on Martian layered ejecta indicates that the grooves fall into three groups that correlate with different types of layered ejecta craters. We collected over 16,000 width and length measurements of 2402 radial grooves on the inner ejecta layers of 36 relatively fresh, layered ejecta craters. These include seven single layered ejecta craters (SLE), seven double layered Type-2 (DLE Type-2), ten double ejecta layer Type-1 (DLE-Type-1), and twelve multiple ejecta layer (MLE) test craters. These test craters are in the diameter range 8.3 km – 38.8 km for latitudes ranging from 45oS to 54oN. Radial grooves resembling longitudinal grooves in landslides throughout the solar system occur on all ejecta layers of SLE, DLE Type-2, and MLE craters (i.e., SDM grooves). This leads us to propose that SDM grooves are structures formed in flowing ejecta by the same processes that produce longitudinal grooves in landslides (e.g., shear and divergent flow). In contrast, the radial grooves on the inner and outer ejecta layers of DLE Type-1 are morphologically and morphometrically different from one another and from SDM grooves. The only type of geophysical flows that form grooves similar to grooves on the inner ejecta layers of DLE Type-1 craters are ones produced by high velocity (i.e., supersonic) blast surges associated with explosive volcanic eruptions (e.g., that occurred during the May 18th1980 eruption of Mount St. Helens USA). This leads us to propose that these Martian grooves are produced by high-speed outflow of particles and gas towards the last phase of crater formation. Furthermore, we propose that this high speed flow is generated when volatile-rich target materials slide into the crater cavity to mix with hot impactite. This generates a fuel-coolant explosion that creates a blast surge that subsequently produces the grooves on the inner ejecta layer. In addition, as the velocity of the outflow drops to subsonic outward of the inner ejecta layer, both erosion and deposition processes operate that result in deposition of a thin layer of material that contains channel-like grooves that are unique to this ejecta layer of this type of crater. This thin deposit is found on the outer ejecta layers, and beyond (on the terrain surrounding fresh DLE Type-1 craters). The production of these grooves on DLE Type-1 ejecta layers likely destroys and/or masks SDM grooves that may have initially formed on DLE Type-1 ejecta layers when they were initially emplaced as dense ground hugging flows.

We review the methods and data sets used to determine morphometric parameters related to the depth (e.g., rim height and cavity depth) and diameter of Martian craters over the past ~45 yr, and discuss the limitations of shadow length measurements, photoclinometry, Earth-based radar, and laser altimetry. We demonstrate that substantial errors are introduced into crater depth and diameter measurements that are inherent in the use of 128th-degree gridded Mars Orbiter Laser Altimeter (MOLA) topography. We also show that even the use of the raw MOLA Precision Engineering Data Record (PEDR) data can introduce errors in the measurement of craters a few kilometers in diameter. These errors are related to the longitudinal spacing of the MOLA profiles, the along-track spacing of the individual laser shots, and the MOLA spot size. Stereophotogrammetry provides an intrinsically more accurate method for measuring depth and diameter of craters on Mars when applied to high-resolution image pairs. Here, we use 20 stereo Context Camera (CTX) image pairs to create digital elevation models (DEMs) for 25 craters in the diameter range 1.5–25.6 km and cover the latitude range of 25° S to 42° N. These DEMs have a spatial scale of ~24 m per pixel. Six additional craters, 1.5–3.1 km in diameter, were studied using publically available DEMs produced from High-Resolution Imaging Science Experiment (HiRISE) image pairs. Depth/diameter and rim height were determined for each crater, as well as the azimuthal variation of crater rim height in 1-degree increments. These data indicate that morphologically fresh Martian craters at these diameters are significantly deeper for a given size than previously reported using Viking and MOLA data, most likely due to the improvement in spatial resolution provided by the CTX and HiRISE data.

Mons Hansteen is a relatively high-albedo, well-known lunar ``red spot'' located on the southern margin of Oceanus Procellarum (2.3°S, 50.2°W). It is an arrowhead-shaped (∼ 25 km on a side), two-layer mesa with a small cone-shaped massif on its north edge formed by three morphologically and compositionally distinct geologic units. These units were emplaced in three phases over nearly 200 million years. The oldest (∼3.74 Ga), Hilly–Dissected unit, composed of high-silica, and low-FeO content materials formed a low, steep sided mesa. The materials of this unit erupted mainly from vents along northeast- and northwest-trending sets of fractures. The Pitted unit, which comprises the upper-tier mesa, is composed of high-silica and even lower-FeO content materials. This material was erupted at ∼ 3.5 Ga from numerous closely spaced vents (i.e., pits) formed along closely spaced northeast-southwest-trending sets of fractures. At nearly the same time, eruptions of lower silica and higher FeO materials occurred on the north flank of Mons Hansteen at the intersection of two major fractures to produce the North Massif unit. The eruptions that created the North Massif units also produced materials that thinly blanketed small areas of the Hilly-Dissected and Pitted units on the north flank of Mons Hansteen. Also at nearly the same time (i.e., ∼ 3.5 Ga), basalt flows formed the surrounding mare. Each unit of Mons Hansteen appears to be mantled by locally derived ash, which only modestly contaminated the other units. The morphology of Mons Hansteen (especially the Pitted unit) suggests a style of volcanism where only a relatively small amount of material is explosively erupted from numerous, closely spaced vents.

Evidence suggests that the lobate flow feature that extends ~72 km outward from the western rim of Tsiolkovskiy crater is a long runout landslide. This landslide exhibits three (possibly four) morphologically different parts, likely caused by local conditions. All of these, plus the ejecta of Tsiolkovskiy crater, and its mare fill are approximately of the same crater model age, i.e., ~3.55 ± 0.1 Ga. The enormous size of this landslide is unique on the Moon and is a result of a combination of several geometric factors (e.g., its location relative to Fermi crater), and that Tsiolkovskiy crater was an oblique impact that produced an ejecta forbidden zone on its western side (Schultz, 1976). The landslide formed in this ejecta free zone as the rim of Tsiolkovskiy collapsed and its debris flowed across the relatively smooth, flat floor of Fermi crater. In this location, it could be easily identified as a landslide and not ejecta. Its mobility and coefficient of friction are similar to landslides in Valles Marineris on Mars, but less than wet or even dry terrestrial natural flows. This suggests that the Mars landslides may have been emplaced dry. The high density of small craters on the landslide is likely an illusion caused by the effects of age related differences in regolith thickness on crater morphology, and the presence of the abundant young, circular secondary craters produce by debris ejected from distant fresh craters.

The degradation state of an impact crater is an indicator of its age. Previous workers have used crater degradation states to estimate ages of surfaces or geomorphological features; one example is the degree of freshness method developed by Pohn and Offield (1970). Here we attempted to produce an empirical calibration that yields absolute model ages based upon the degree of freshness technique for craters ~8–20 km in diameter. To produce the calibration, we first selected 15 craters with degree of freshness values ranging from 2.5 to 6.3. Next, we used the Kaguya Terrain Camera data to measure crater density on the ejecta of these craters, from which absolute model age could be calculated. The resulting absolute model ages ranged from 0.9 to 4.0 Ga. We used two linear regressions to describe the relationship between the absolute model age and degree of freshness of the craters. We fitted each trend with two linear least-squares regressions, where the first regression represents craters with a degree of freshness from 0.0 to 4.9 and the second regression from 5.0 to 7.0. The 95% confidence belt shows that the calibrations are accurate to ±0.5 Ga to ±1.1 Ga for the fresh crater regression (5.0–7.0) and slightly more accurate, to ±0.3 Ga to ±0.1 Ga, for the degraded crater regression (0.0–4.9). However, the degraded crater regression is likely based upon craters with continuous ejecta that are crater saturated, thus implying that craters with a degree of freshness 3.8 Ga.

Low-aspect ratio layered ejecta (LARLE) craters are one of the most enigmatic types of martian layered ejecta craters. We propose that the extensive outer layer of these craters is produced through the same base surge mechanism as that which produced the base surge deposits generated by near-surface, buried nuclear and high-explosive detonations. However, the LARLE layers have higher aspect ratios compared with base surge deposits from explosion craters, a result of differences in thicknesses of these layers. This characteristics is probably caused by the addition of large amounts of small particles of dust and ice derived from climate-related mantles of snow, ice and dust in the areas where LARLE craters form. These deposits are likely to be quickly stabilized (order of a few days to a few years) from eolian erosion by formation of duricrust produced by diffusion of water vapor out of the deposits.

Low-Aspect-Ratio Layered Ejecta (LARLE) craters are a unique landform found on Mars. LARLE craters are characterized by a crater and normal layered ejecta pattern surrounded by an extensive but thin outer deposit which terminates in a sinuous, almost flame-like morphology. We have conducted a survey to identify all LARLE craters ⩾1-km-diameter within the ±75° latitude zone and to determine their morphologic and morphometric characteristics. The survey reveals 140 LARLE craters, with the majority (91%) located poleward of 40°S and 35°N and all occurring within thick mantles of fine-grained deposits which are likely ice-rich. LARLE craters range in diameter from the cut-off limit of 1 km up to 12.2 km, with 83% being smaller than 5 km. The radius of the outer LARLE deposit displays a linear trend with the crater radius and is greatest at higher polar latitudes. The LARLE deposit ranges in length between 2.56 and 14.81 crater radii in average extent, with maximum length extending up to 21.4 crater radii. The LARLE layer is very sinuous, with lobateness values ranging between 1.45 and 4.35. LARLE craters display a number of characteristics in common with pedestal craters and we propose that pedestal craters are eroded versions of LARLE craters. The distribution and characteristics of the LARLE craters lead us to propose that impact excavation into ice-rich fine-grained deposits produces a dusty base surge cloud (like those produced by explosion craters) that deposits dust and ice particles to create the LARLE layers. Salts emplaced by upward migration of water through the LARLE deposit produce a surficial duricrust layer which protects the deposit from immediate removal by eolian processes.

We propose a numerical model for the formation of the closely-spaced pits found in the thin, ejecta-related deposits superposed on the floors, interior terrace blocks, and near-rim ejecta blankets of well-preserved martian impact craters. Our model predicts the explosive degassing of water from this pitted material, which is assumed to originally be water-bearing, impact melt-rich breccia at the time of deposition. This process is analogous to what occurred in the fall-out suevite deposits at the Ries impact structure in Germany. At Ries, impact heating of water-bearing target material resulted in the rapid degassing of its water and other volatiles. The martian environment plays an important role in enhancing the effects of this degassing by increasing the flow-speed of the escaping gas. The high flow-rate of gas through particulate materials, such as suevite, tends to quickly form segregation channels or vent pipes, similar to those found in the Ries deposits. These pipes act as conduits for the efficient high-speed escape of the gas and small clasts that it entrains. Escaping gas and entrained clasts abraded and eroded the conduit walls, flaring them to form pits above a network of pipes.

Recently acquired high-resolution images of martian impact craters provide further evidence for the interaction between subsurface volatiles and the impact cratering process. A densely pitted crater-related unit has been identified in images of 204 craters from the Mars Reconnaissance Orbiter. This sample of craters are nearly equally distributed between the two hemispheres, spanning from 53°S to 62°N latitude. They range in diameter from ∼1 to 150 km, and are found at elevations between −5.5 to +5.2 km relative to the martian datum. The pits are polygonal to quasi-circular depressions that often occur in dense clusters and range in size from ∼10 m to as large as 3 km. Pit sizes scale with both the host crater’s diameter and the host deposit size. These pits have subtle raised rims, and unlike primary and secondary impact craters, they lack well-defined ejecta deposits and overlapping stratigraphic relationships. They also lack any sign of any preferential alignment expected of volcanic or tectonic collapse features. Morphologic and stratigraphic evidence in support of an impact origin includes the observation that pitted materials primarily occur as ponded and flow-like deposits on crater floors, behind terraces, and infilling the lowest local topographic depressions atop the ejecta blanket—similar to the distribution of impact melt-bearing bodies on the Moon. Based on the observations and comparisons to terrestrial and lunar analogs, we conclude that the pit-bearing materials are impactite deposits. The presence of these deposits in older craters, where preserved, suggests that they have formed on Mars throughout most of its geologic history; thus, understanding their origin may help to constrain the hydrological and climate history of Mars.

The 27.2 km diameter Tooting crater is the best preserved young impact crater of its size on Mars. It offers an unprecedented opportunity to study impact-related phenomena as well the geology of the crust in the Amazonis Planitia region of Mars. For example, the nearly pristine condition enables the partial reconstruction of the sequence of events for crater formation, as well as facilitates a comparison to deposits seen at the Ries crater in Germany. High-resolution images taken by the High Resolution Imaging Science Experiment (HiRISE) and Context Camera (CTX) on the Mars Reconnaissance Orbiter spacecraft have revealed a wealth of information on the distribution of features within the crater and beyond the rim: a large central peak, pitted material on the floor and terrace blocks, lobate flows interpreted to be sediment flows, impact melt sheets, four discrete layers of ejecta, and an asymmetric secondary crater field. Topographic data derived from the Mars Orbiter Laser Altimeter (MOLA) and stereo HiRISE and CTX images show that the central peak is ∼1100 m high, the lowest point of the crater floor is 1274 m below the highest part of the rim, and the crater rim has ∼600 m of variability around its perimeter. Layering within the cavity walls indicates ∼260 m of structural uplift of the target material, which constitutes ∼35% of the total relief of the rim. Abundant evidence is found for water flowing down the cavity walls, and on the surface of the ejecta layers, both of which took place sometime after the impact event. Thickness measurements of the ejecta layers reveal that the continuous blanket is remarkably thin (∼3–5 m) in some places, and that the distal ramparts may be ∼60 m high. Crater counts made on the ejecta layers indicate a model age of

A Late Noachian-aged alluvial fan complex within Harris Crater in far western Terra Tyrrhena, Mars, is comprised of two well-defined source regions and associated discrete depositional lobes. Three fan units were recognized based on common morphological characteristics, thermal properties and spectral signatures. Although the entire fan complex has been subjected to extensive erosional degradation, the preserved morphologies record episodic fan formation and indicate the type of flow processes that occurred; the bulk of the fan surface has morphology consistent with fluvial emplacement while one fan unit exhibits a rugged surface texture with boulders consistent with a debris flow. This transition from fluvial to late-stage debris flow(s) suggests a decline in available water and/or change in sediment supply. The thermal inertia values obtained for all three fan surface units (mean values ranged from 318 to 344 J m � 2 K

– Some fresh impact craters on Ganymede have the overall ejecta morphology similar to Martian double-layer ejecta (DLE), with the exception of the crater Nergal that is most like Martian single layer ejecta (SLE) craters (as is the terrestrial crater Lonar). Similar craters also have been identified on Europa, but no outer ejecta layer has been found on these craters. The morphometry of these craters suggests that the types of layered ejecta craters identified by Barlow et al. (2000) are fundamental. In addition, the mere existence of these craters on Ganymede and Europa suggests that an atmosphere is not required for ejecta fluidization, nor can ejecta fluidization be explained by the flow of dry ejecta. Moreover, the absence of fluidized ejecta on other icy bodies suggests that abundant volatiles in the target also may not be the sole cause of ejecta fluidization. The restriction of these craters to the grooved terrain of Ganymede and the concentration of Martian DLE craters on the northern lowlands suggests that these terrains may share key characteristics that control the development of the ejecta of these craters. In addition, average ejecta mobility (EM) ratios indicate that the ejecta of these bodies are self-similar with crater size, but are systematically smaller on Ganymede and Europa. This may be due to the effects of the abundant ice in the crusts of these satellites that results in increased ejection angle causing ejecta to impact closer to the crater and with lower horizontal velocity.

The Mars Crater Morphology Consortium recommends the use of a standardized nomenclature system when discussing Martian impact crater ejecta morphologies. The system utilizes nongenetic descriptors to identify the various ejecta morphologies seen on Mars. This system is designed to facilitate communication and collaboration between researchers. Crater morphology databases will be archived through the U.S. Geological Survey in Flagstaff, where a comprehensive catalog of Martian crater morphologic information will be maintained.

We are developing a portable laser desorption resonance ionization mass spectrometer (LDRIMS) for determining the radiometric age of rocks using the 87Rb-87Sr isotope system, as well as constraining lithologic evolution and measuring chemical compositions. The bench-top prototype has been used to assess the capability of LDRIMS to measure 87Sr/86Sr. In this paper, we demonstrate that LDRIMS can measure the isotope ratios of a glass standard with 10 ppm net Sr to a precision of ±0.5% (1σ), with a sensitivity of 1:1010 in less than 1 minute. Increasing the measurement time to 15 minutes improves the precision to 0.1% (1σ). The speed of the LDRIMS measurement allows samples to be measured in significantly shorter periods of time than traditional methods, with little or no sample preparation. Models of the age error derived from isochron dating the SNC meteorites that would be obtained using 100–1000 LDRIMS measurements at ±0.1% (1σ) accuracy show that for ALH84001 and Zagami, which have ages ranging from 4.09 Ga to 165 Ma, dates with analytical uncertainties less than ±50 Ma are possible. These results were obtained using low laser powers (∼10 μJ for resonance

Magellan radar images of 15 percent of the planet show 135 craters of probable impact origin. Craters more than 15 km across tend to contain central peaks, multiple central peaks, and peak rings. Many craters smaller than 15 km exhibit multiple floors or appear in clusters; these phenomena are attributed to atmospheric breakup of incoming meteoroids. Additionally, the atmosphere appears to have prevented the formation of primary impact craters smaller than about 3 km and produced a deficiency in the number of craters smaller than about 25 km across. Ejecta is found at greater distances than that predicted by simple ballistic emplacement, and the distal ends of some ejecta deposits are lobate. These characteristics may represent surface flows of material initially entrained in the atmosphere. Many craters are surrounded by zones of low radar albedo whose origin may have been deformation of the surface by the shock or pressure wave associated with the incoming meteoroid. Craters are absent from several large areas such as a 5 million square kilometer region around Sappho Patera, where the most likely explanation for the dearth of craters is volcanic resurfacing. There is apparently a spectrum of surface ages on Venus ranging approximately from 0 to 800 million years, and therefore Venus must be a geologically active planet.

[1] Based on the depths of pristine complex craters on Mars, we have estimated the final, post-modification phase depth (dr) to diameter (D) function of Martian complex craters. We suggest that the dr/D function determined here holds globally for complex craters >∼12 km diameter. The global nature of this function may be due to the stability of slopes in the highly fractured rock that makes up the steep interior walls of post-modification phase craters. The stability of such slopes should be controlled mainly by such factors as the nature of the fractures and not target strength. This function is an essential benchmark for studies that utilize crater geometry as their basis. For example, the function can be used in gradation history of Mars to predict the magnitude of degradation. For instance, an example 20-km-diameter, Middle/Early Noachian-age crater that initially formed in the mid-latitudes of Mars with a depth of 1800 m (instead of 1404 m as predicted by the fresh craters function) has lost ∼1450 m of its depth over its history, ∼350 m in the Noachian, ∼400 m during an intense episode of degradation at the end of the Noachian, and ∼700 m thereafter. No evidence of the period of intense degradation is found in the high-latitudes, although those regions appear to have been affected by a process that reduced crater depth soon after they form.

[1] The Thermal Emission Imaging System (THEMIS) visible (VIS) images provide new insight into the nature and development of the unique ejecta deposits of Martian craters. This study focuses on double-layered ejecta (DLE) craters. To date, over 100 DLE craters have been examined using mainly THEMIS VIS data. Our observations suggest that emplacement of DLE crater ejecta occurred in two stages, with the inner ejecta layer emplaced similar to single-layered ejecta (SLE) crater ejecta. This may have involved both ballistic and flow processes. In contrast, the outer ejecta layer of DLE craters appears to have been emplaced through the high-velocity outflow of materials from tornadic winds generated by the advancing ejecta curtain or base surge. Remarkably, DLE craters lack secondary craters, which suggests that the large ejecta blocks that normally produce such craters may have either been entrained and/or crushed by these winds or fragmented as a result of the presence of water in the target materials. These observations suggest that volatiles (either trapped in the subsurface or in the atmosphere) have played a key role in the emplacement of the ejecta of DLE craters and leaves open the question as to what role volatiles play in the emplacement of ejecta of other types of fluidized ejecta craters (i.e., SLE and MLE craters). Because DLE craters are found in many different regions of Mars, often in close proximity to other types of craters, conditions (e.g., atmospheric density) that produce DLE craters must fluctuate or the Martian crust must be unexpectedly heterogeneous (laterally and vertically). While the degree of heterogeneity has yet to be recognized, recent suggestions about possible Martian climate change raises the possibility of impact into target materials that are periodically wet or that a significantly higher atmospheric pressure may be periodically present.

The northern lowland plains, such as those found in Acidalia and Utopia Planitia, have high percentages of impact craters with fluidized ejecta. In both regions, the analysis of crater geometry from Mars Orbiter Laser Altimeter (MOLA) data has revealed large ejecta volumes, some exceeding the volume of excavation. Moreover, some of the crater cavities and fluidized ejecta blankets of these craters are topographically perched above the surrounding plains. These perched craters are concentrated between 40 and 70°N in the northern plains. The atypical high volumes of the ejecta and the perched craters suggest that the northern lowlands have experienced one or more episodes of resurfacing that involved deposition and erosion. The removal of material, most likely caused by the sublimation of ice in the materials and their subsequent erosion and transport by the wind, is more rapid on the plains than on the ejecta blankets. The thermal inertia difference between the ejecta and the surrounding plains suggests that ejecta, characterized by a lower thermal inertia, protect the underneath terrain from sublimation. This results in a decreased elevation of the plains relative to the ejecta blankets. Sublimation and eolian erosion can be particularly high during periods of high obliquity.

] Using THEMIS, MOC and MOLA data, we havefound 51 craters in the diameter range 6–11.8 km withinsouthwestern Utopia Planitia and Isidis Planitia that aresignificantly deeper than typical fresh craters in the northernlowlands of Mars. The restricted geographic distribution ofthese craters, their simple morphology, and data fromimpact and explosion crater studies suggest that unusuallystrong target materials (as much as a factor of 2 greater thanaverage materials in the Martian lowlands) are the cause ofthe excessive crater depth. We propose that the greater targetmaterial strength acts to delay gravity-dominated collapse tolarger crater sizes. Furthermore, we suggest that a regional,olivine-rich mafic to untramafic rock unit identified by TESand THEMIS is a reasonable candidate for these strongmaterials. The unit is exposed on the southern edge of Isidisbasin and in crater ejecta within the basin, and forms layersthat dip toward the Isidis Basin center.

[1] We present depth (d) and diameter (D) data for 2269 Martian impact craters in the diameter range 6–216 km, distributed over ∼11.52 × 106 km2 of terrain that samples several geologic settings and the full range of latitudes and elevations throughout the northern lowlands of Mars. Our data indicate that there are three major crater types that fall into two major populations, those craters that contain Vastitis Borealis formation (VBF) and those that do not. The deepest type 1 craters comprise a Late Hesperian/Early Amazonian-age population that does not contain VBF. Type 2, type 3, and subdued type 1 craters comprise a population of Hesperian-age craters partially buried by the VBF. The unique d/D distribution of type 3 craters and photogeologic evidence suggest an erosional style for the VBF that requires sublimation processes, indicating that the VBF contained a substantial amount of ice. Type 3 craters occur throughout the northern lowland plains at elevations below −2400 m. If the VBF is a sedimentary deposit left by a large body of standing water in the northern lowlands of Mars, then the northern lowland plains contained a body of water with a volume of ∼6 × 107 km3 or the equivalent of a 430 m global ocean during the Late Hesperian/Early Amazonian. This also implies that the VBF contains ∼4 × 106 km3 of material, in agreement with the estimated amount of material eroded from the outflow channels surrounding Chryse basin and washed into the northern lowlands.

[1] We have developed a new interactive computer program that facilitates the easy collection of geomorphic data for Martian impact craters, using the MOLA 128th degree digital elevation model of Mars. We describe the method for ensuring that accurate measurements of crater diameter, depth, rim height, rim volume, cavity volume, ejecta thickness, and ejecta volume are obtained. We compare our measurements of crater diameters and rim heights to results obtained by Garvin et al. [2000], who employed centerline MOLA profiles. Statistical regressions between the two methods give R 2 values of 0.930 for crater depths and 0.984 for crater diameters. The new interactive program facilitates the rapid compilation of large data sets to allow a comparison of crater populations in different settings. Preliminary results are presented for 354 craters on ridged plains materials in Hesperia and Sinai Plana to demonstrate the value of the program for regional comparisons and the analysis of degradational processes on Mars. INDEX TERMS: 5420 Planetology: Solid Surface Planets: Impact phenomena (includes cratering); 5494 Planetology: Solid Surface Planets: Instruments and techniques; KEYWORDS: geomorphology, impact craters, Mars

Impact craters on Triton are scarce owing to the relatively recent resurfacing by icy melts. The most heavily cratered surface has a crater density about the same as the lunar maria. The transition diameter from simple to complex craters occurs at a diameter of about 11 kilometers, and the depth-diameter relationship is similar to that of other icy satellites when gravity is taken into account. The crater size-frequency distribution has a differential -3 slope (cumulative -2 slope) and is the same as that for the fresh crater population on Miranda. The most heavily cratered region is on the leading hemisphere in Triton's orbit. Triton may have a leading-trailing asymmetry in its crater population. Based primarily on the similarity of size distributions on Triton and Miranda and the relatively young surface on Triton, the source of Triton's craters is probably comets. The very peculiar size distribution of sharp craters on the "cantaloupe" terrain and other evidence suggests they are volcanic explosion craters.

The Magellan spacecraft arrived at Venus on August 10, 1990 to begin global high-resolution radar image and altimeter mapping and gravity field determination of the planet. Mapping of this cloud covered planet is done from an elliptical orbit with a 3.15 hour period and a 275 km periapsis. After the initial checkout of the spacecraft, the primary missions lasts for 243 days (one Venus day); the time it takes for Venus to make one revolution beneath the spacecraft. The data collected by Magellan is expected to provide an understanding of the geological and geophysical nature of the planet Venus, and the general processes that control planetary evolution.

Voyager 2 images of Neptune reveal a windy planet characterized by bright clouds of methane ice suspended in an exceptionally clear atmosphere above a lower deck of hydrogen sulfide or ammonia ices. Neptune's atmosphere is dominated by a large anticyclonic storm system that has been named the Great Dark Spot (GDS). About the same size as Earth in extent, the GDS bears both many similarities and some differences to the Great Red Spot of Jupiter. Neptune's zonal wind profile is remarkably similar to that of Uranus. Neptune has three major rings at radii of 42,000, 53,000, and 63,000 kilometers. The outer ring contains three higher density arc-like segments that were apparently responsible for most of the ground-based occultation events observed during the current decade. Like the rings of Uranus, the Neptune rings are composed of very dark material; unlike that of Uranus, the Neptune system is very dusty. Six new regular satellites were found, with dark surfaces and radii ranging from 200 to 25 kilometers. All lie inside the orbit of Triton and the inner four are located within the ring system. Triton is seen to be a differentiated body, with a radius of 1350 kilometers and a density of 2.1 grams per cubic centimeter; it exhibits clear evidence of early episodes of surface melting. A now rigid crust of what is probably water ice is overlain with a brilliant coating of nitrogen frost, slightly darkened and reddened with organic polymer material. Streaks of organic polymer suggest seasonal winds strong enough to move particles of micrometer size or larger, once they become airborne. At least two active plumes were seen, carrying dark material 8 kilometers above the surface before being transported downstream by high level winds. The plumes may be driven by solar heating and the subsequent violent vaporization of subsurface nitrogen.

Data on the distribution of impact craters on Venus are analyzed in relation to the geologic history of Venus and the effects of high ambient pressures and temperatures on the cratering process itself. The cratering record of Venus is discussed in terms of size, number, spatial distribution, and state of preservation of craters. The morphology of Venusian craters is examined as a function of the diameter, and features unique to these craters are discussed. Attention is also given to several craters that illustrate important features of Venusian crater; surface 'splotches' that are closely related in origin to impact craters; and implications of the distribution of severely modified craters for Venus' recent geologic history.

The reported evaluation of Martian channel characteristics is based on Viking photographs taken from July 1976 to February 1977. The wide variation in crater densities shown by the considered Martian channels strongly implies widely differing ages for both fluviatile and lava channels. Attention is given to age determination methodology, a description of channels and implications for channel formation, surface water under present Martian conditions, surface water under more favorable Martian conditions in the past, channel parameter estimates, and volcanic channels.

The surfaces of the Saturn satellites Tethys, Iapetus and Encedalus display surfaces which indicate active geological processes and therefore suggest a degree of internal evolution. By contrast, the Saturn satellite Hyperion and the coorbitals 1980S1 and 1980S3 show no trace of geological activity and may be fragments of once-larger bodies. Activity on Iapetus appears to have been confined to the dark terrain, and offers no clue as to its timing and extent. The widest terrain type and crater number variations are those of Encedalus, which indicate the most prolonged period of geological activity of any of the satellites studied.

Since the General Assembly at New Delhi in November 1985, the Working Group held two meetings within six weeks of each other; most members of the Working Group and several members of Task Groups were able to attend at least one of these meetings. The thirteenth meeting of the Working Group was held at Toulouse, France on June 30 to July 2,1986; the fourteenth meeting was held at Moscow, USSR, on August 10,1986; the fifteenth meeting took place from August 13 to 15 in Soviet Armenia.

Voyager 2, during its encounter with the Jupiter system, provided images that both complement and supplement in important ways the Voyager 1 images. While many changes have been observed in Jupiter's visual appearance, few, yet significant, changes have been detected in the principal atmospheric currents. Jupiter's ring system is strongly forward scattering at visual wavelengths and consists of a narrow annulus of highest particle density, within which is a broader region in which the density is lower. On Io, changes are observed in eruptive activity, plume structure, and surface albedo patterns. Europa's surface retains little or no record of intense meteorite bombardment, but does reveal a complex and, as yet, little-understood system of overlapping bright and dark linear features. Ganymede is found to have at least one unit of heavily cratered terrain on a surface that otherwise suggests widespread tectonism. Except for two large ringed basins, Callisto's entire surface is heavily cratered.

The cameras aboard Voyager 1 have provided a closeup view of the Jupiter system, revealing heretofore unknown characteristics and phenomena associated with the planet's atmosphere and the surfaces of its five major satellites. On Jupiter itself, atmospheric motions-the interaction of cloud systems-display complex vorticity. On its dark side, lightning and auroras are observed. A ring was discovered surrounding Jupiter. The satellite surfaces display dramatic differences including extensive active volcanismn on Io, complex tectonism on Ganymnede and possibly Europa, and flattened remnants of enormous impact features on Callisto.

It is shown that Mercury's surface has only 70% as many large craters (of at least 200 km in diameter) as the moon. The density of Mercurian impact craters having diameters over 400 km is 30% of that of the moon, and for craters with diameters between 400 and 700 km, Mercurian density is 21% of that of the moon. The size-frequency distribution curve of Mercury is the same as the lunar cumulative -2 slope. The Mercurian curve, however, lies well below the 10% surface saturation level of the lunar curve. This may indicate that the old, heavily-cratered Mercurian terrain is not presently in a state of cratering equilibrium. The differences in crater and basin densities observed between Mercury and the moon may be functions of crater-production rates or of different crustal histories. The total isostatic compensation of impact craters having diameters of about 800 km suggests that the average viscosity of the Mercurian crust during approximately the past 4 eons was the same as that of the moon.

The Saturn icy satellites Rhea, Dione, Mimas and Tethys were imaged by Voyager 1 with sufficiently high resolution to permit crater counts that will determine the character and magnitude of crater density variations and the relative ages of the different terrain types. The data show that the surfaces of these satellites are not uniformly cratered, and therefore not of uniform age. Density variations are found to occur at all diameters, especially in the distribution of large-diameter craters (greater than 50 km). It is concluded that this satellite group, despite their small size and mass and low density, have evolved beyond their original, accretionary state. The alteration processes appear in light of crater density determinations to have been active beyond the period of heavy bombardment. The energy source for the evolution of these bodies has yet to be determined.

Voyager 2 photography has complemented that of Voyager I in revealing many additional characteristics of Saturn and its satellites and rings. Saturn's atmosphere contains persistent oval cloud features reminiscent of features on Jupiter. Smaller irregular features track out a pattern of zonal winds that is symmetric about Saturn's equator and appears to extend to great depth. Winds are predominantly eastward and reach 500 meters per second at the equator. Titan has several haze layers with significantly varying optical properties and a northern polar "collar" that is dark at short wavelengths. Several satellites have been photographed at substantially improved resolution. Enceladus' surface ranges from old, densely cratered terrain to relatively young, uncratered plains crossed by grooves and faults. Tethys has a crater 400 kilometers in diameter whose floor has domed to match Tethys' surface curvature and a deep trench that extends at least 270 degrees around Tethys' circumference. Hyperion is cratered and irregular in shape. Iapetus' bright, trailing hemisphere includes several dark-floored craters, and Phoebe has a very low albedo and rotates in the direction opposite to that of its orbital revolution with a period of 9 hours. Within Saturn's rings, the "birth" of a spoke has been observed, and surprising azimuthal and time variability is found in the ringlet structure of the outer B ring. These observations lead to speculations about Saturn's internal structure and about the collisional and thermal history of the rings and satellites.

Voyager 2 images of the southern hemisphere of Uranus indicate that submicrometersize haze particles and particles of a methane condensation cloud produce faint patterns in the atmosphere. The alignment of the cloud bands is similar to that of bands on Jupiter and Saturn, but the zonal winds are nearly opposite. At mid-latitudes (-70 degrees to -27 degrees ), where winds were measured, the atmosphere rotates faster than the magnetic field; however, the rotation rate of the atmosphere decreases toward the equator, so that the two probably corotate at about -20 degrees . Voyager images confirm the extremely low albedo of the ring particles. High phase angle images reveal on the order of 10(2) new ringlike features of very low optical depth and relatively high dust abundance interspersed within the main rings, as well as a broad, diffuse, low optical depth ring just inside the main rings system. Nine of the newly discovered small satellites (40 to 165 kilometers in diameter) orbit between the rings and Miranda; the tenth is within the ring system. Two of these small objects may gravitationally confine the e ring. Oberon and Umbriel have heavily cratered surfaces resembling the ancient cratered highlands of Earth's moon, although Umbriel is almost completely covered with uniform dark material, which perhaps indicates some ongoing process. Titania and Ariel show crater populations different from those on Oberon and Umbriel; these were probably generated by collisions with debris confined to their orbits. Titania and Ariel also show many extensional fault systems; Ariel shows strong evidence for the presence of extrusive material. About halfof Miranda's surface is relatively bland, old, cratered terrain. The remainder comprises three large regions of younger terrain, each rectangular to ovoid in plan, that display complex sets of parallel and intersecting scarps and ridges as well as numerous outcrops of bright and dark materials, perhaps suggesting some exotic composition.

As Voyager 1 flew through the Saturn system it returned photographs revealing many new and surprising characteristics of this complicated community of bodies. Saturn's atmosphere has numerous, low-contrast, discrete cloud features and a pattern of circulation significantly different from that of Jupiter. Titan is shrouded in a haze layer that varies in thickness and appearance. Among the icy satellites there is considerable variety in density, albedo, and surface morphology and substantial evidence for endogenic surface modification. Trends in density and crater characteristics are quite unlike those of the Galilean satellites. Small inner satellites, three of which were discovered in Voyager images, interact gravitationally with one another and with the ring particles in ways not observed elsewhere in the solar system. Saturn's broad A, B, and C rings contain hundreds of "ringlets," and in the densest portion of the B ring there are numerous nonaxisymmetric features. The narrow F ring has three components which, in at least one instance, are kinked and crisscrossed. Two rings are observed beyond the F ring, and material is seen between the C ring and the planet.