Crater Population Research Articles

The Lunar rayed-crater population — Characteristics of the spatial distribution and ray retention

The global statistics for young impact craters on the Moon is used to unravel potential spatial asymmetries, that may have been introduced by the particular orbital configuration of a synchronously rotating satellite. Only craters that exhibit bright ejecta rays extending for several crater radii were considered in this study. This crater population is younger than about 750Ma. The shape of the crater size-frequency distribution does not show strong dependence on the target properties (mare vs. highlands). However, slightly lower frequencies indicate a shorter retention of the visibility of rays in mare units when their visibility is purely due to immaturity and not due to composition. Rays of small craters fade away much faster. Large, old, rayed craters sustain their visibility longer than the average crater population because of the compositional contrast between rays and mare material, and thus obscure the cratering record when investigated for spatial variations. Using the existence of rays purely based on optical maturity instead of visibility as marker horizon for the Copernican–Eratosthenian boundary, suggests a shift from 1.1Ga to 750Ma.The spatial distribution of lunar rayed craters, namely the latitudinal and longitudinal frequency variations, does not agree with previous analytical and numerical studies. Although there is an apparent hemispherical asymmetry centred close to the apex, the density distribution is patchy and no predicted spatial pattern could be confirmed. Spatial distribution corrections accounting for the lower frequencies in the mare areas did not result in a better agreement with the analytical estimates. Density variations are less than 15% over vast parts of the lunar surface, and the uncertainties for absolute surface ages are similar. However, variations of up to 50% are found even for the more numerous small craters. These extreme values are located at high latitudes. A combination of crater-forming projectile flux distribution and micrometeorite bombardment, which acts on maturation of the ray systems, could prove as an explanation for the contradicting observed rayed crater distribution.An analysis of the older craters is more challenging (on the Moon) because earlier geological processes complicate the setting, and the orbital configuration of the Moon–Earth–Sun–projectile system altered with time.

Confident thickness estimates for planetary surface deposits from concealed crater populations

An improved technique is presented to determine more accurately deposit thicknesses of any surface units using crater size-frequency distributions (CSFDs). This new approach enables thickness estimates of deposits that completely cover their underlying unit, i.e., where no flooded craters are observed. Here, the crater populations of the unit of interest and its underlying unit (or a representative surface area) are measured and their cratering model ages determined. The CSFD of the younger unit is then treated as if it had the same age as the older unit and the expected cumulative numbers of craters for each bin size are calculated. The crater deficiency for each bin size is then determined by subtracting the expected cumulative crater number from the observed cumulative crater number per bin. The probability is calculated such that the crater deficiency corresponds to the age of the older, underlying unit. The crater bin size where we are 95% confident that these craters are there but just covered by the younger unit is chosen to derive the deposit thickness value using known crater diameter-rim height relationships. We demonstrate the use of the method at two sites of different scales and settings on Mars: Pickering crater and Lunae Planum, and compare our estimates with those obtained by other methods. For the Lunae Planum area a significantly higher minimum thickness of >840m is predicted.

A Middle-Late Triassic 40Ar/39Ar age for the Paasselkä impact structure (SE Finland)

Abstract– 40Ar/39Ar dating of recrystallized feldspar glass particles separated from clast-rich impact melt rocks from the approximately 10 km Paasselkä impact structure (SE Finland) yielded a Middle to Late Triassic (Ladinian-Karnian) pseudo-plateau age of 228.7 ± 3.0 (3.4) Ma (2σ). This new age makes Paasselkä the first known Triassic impact structure dated by isotopic methods on the Baltic Shield. The new Paasselkä impact age is, within uncertainty, coeval with isotopic ages recently obtained for the Lake Saint Martin impact structure in Canada, indicating a new Middle to Late Triassic impact crater population on Earth. The comparatively small crater size, however, suggests no relationship between the Paasselkä impact and a postulated extinction event at the Middle/Late Triassic boundary.

Planetary surface dating from crater size–frequency distribution measurements: Partial resurfacing events and statistical age uncertainty

We describe the procedure to fit a cumulative production function polynomial to a partial crater size–frequency distribution. The technique is of particular use in deriving ages for surfaces which have undergone partial resurfacing events: namely, erosional or depositional events which have affected a limited diameter range of the crater population. We demonstrate its use in obtaining times for both the surface formation and the resurfacing event. We give a practical outline of the method for making age measurements from crater counts and how to identify resurfacing effects in the results. We discuss the conversion of production function polynomials between common presentations, and the statistical uncertainty of the determined ages with respect to the non-linear chronology function, and a minor refinement of data binning.

Characterization of fluvial activity in Parana Valles using different age-dating techniques

Martian valley networks provide the best evidence that the climate on Mars was different in the past. Although these features are located primarily in heavily cratered terrain of Noachian age (>3.7 Ga), the ages of the features and the time when they were active is not well understood. From superposed craters several recent global studies determined that most valley networks formed during the Late Noachian to Early Hesperian; however, there were some disparities between the techniques. In this study, our principal objective was to test the reliability of the different age-dating techniques to better understand their accuracy and limitations. We applied these techniques to Parana Valles using a variety of high-resolution images taken from different instruments that allow us to identify smaller craters ( D > 125 m) while providing sufficient coverage to support a statistically reliable sampling of crater populations, which is necessary to reduce the uncertainties in age determination. Our results indicate that Parana Valles formed during the Early Hesperian Period but that the crater density ( D > 353 m) is heterogeneous inside the Parana Valles basin. The crater population decreases from the headwaters downstream recording a resurfacing event that is most likely related to the erosion of downstream sub-basins. The terrain near the source area is Late Noachian to Early Hesperian in age while terrains closer to the outlet are Early to Late Hesperian in age. Crater densities ( D > 125 m) inside the valley are also heterogeneous and record several resurfacing events on the valley floor. Where the width of the valley network narrows to <2 km we found evidence of an Amazonian age eolian deposit that is a relatively thin layer of only few meters that was probably deposited as a result of topographic influences. Our results validate the reliability of several proposed age-dating techniques, but we also determined the accuracy and applicability of these techniques. Our results also demonstrate that crater populations can be used to not only determine the relative ages of valley networks, but also to map the distribution of sedimentary materials and the extent of resurfacing events that occurred after valley network formation.

Collisional process on Comet 9/P Tempel 1: Mass loss of its dust and ice by impacts of asteroidal objects and its collisional history

We have studied the collisional process on the nucleus of Comet 9/P Tempel 1 (T1) and estimated the mass loss of surface materials due to the impacts of asteroidal objects. The mass loss rate is around ∼2×105 kg yr-1, which is smaller than that of water sublimation of ∼6×109 kg yr-1. We then estimated the number density of craters formed by asteroid impacts to infer the collisional history of T1. We found that the time necessary to accumulate the crater population on T1 is as long as ∼ a few 104 years, suggesting that T1 crossed the asteroid belt region more than ∼ a few 104 years ago, though this estimated time may be reduced to several thousands of years by taking into account the influence of the sublimation process on the crater size. Alternatively, the total period of the recent orbit, in which the perihelion distance is small enough to cause significant sublimation by erasing a large number of small craters, may be less than ∼2×103 years.

Geomorphic analysis of small rayed craters on Mars: Examining primary versus secondary impacts

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 ∼500 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 “air blast” zones, distal ray systems (&gt;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 ± 1.57R) versus primaries (18.07 ± 7.01R), though primary ejecta were less fractal (Fractal Dimension Index (FDI) &lt; 1.30) and more circular on average (Circularity Ratio (CR) = 0.55 ± 0.25 versus 0.27 ± 0.13 for secondaries). Crater rims were remarkably circular (primaries CR = 0.97 ± 0.02, secondaries at 0.94 ± 0.05), though secondaries have the lowest values (CR &lt; 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 &lt; D &lt; 30 m range.

A refined chronology of catastrophic outflow events in Ares Vallis, Mars

We investigate the geomorphology and chronology of catastrophic flooding in a major martian outflow channel, Ares Vallis. We use recently acquired stereo and colour images and derived topographic data with grid-spacing of 50 m from the High Resolution Stereo Camera (HRSC) onboard Mars Express to constrain the detailed flood geomorphology of proximal Ares Vallis. We identify morphologic evidence for 6 distinct erosion surfaces that were likely formed by multiple episodes of catastrophic flooding. These surfaces occur at distinct topographic levels, exhibit multiple sets of flood grooves, and contain distinct post-flood crater populations. Crater count statistics on the flood-eroded surfaces from MRO Context Camera (CTX) and Thermal Emission Imaging System (THEMIS) visible images indicate that flooding was initiated during the Late Noachian/Early Hesperian at ∼ 3.6 Ga. The morphology data confirm that the oldest flooding events occurred on the topographically highest erosion surfaces while subsequent floods abandoned the upper surfaces and continued to modify craters in the topographically lowest, main channel of Ares Vallis into the Late Hesperian/Early Amazonian (∼ 2.5 Ga). Our results suggest that release of water to Ares Vallis likely occurred at repeated intervals over a long period of geological time, thus supporting recent numerical modelling studies that suggest a recharging water source.

Rugged crater ejecta as a guide to megaregolith thickness in the southern nearside of the Moon

The southern highlands of the Moon comprise superposed ejecta layers, individually as thick as a few kilometers, from the major basins. Smaller (1–16-km-diameter) impact craters that penetrate this layered megaregolith and excavate material from depth have radar properties that provide insight into the variability of megaregolith thickness above a postulated basement of large crustal blocks. We observe a significant difference in the population of radar-bright craters, 1–16 km and larger in diameter, between regions of the southeastern near-side highlands north and south of ~lat 48°S. There are about one-third more radar-bright craters north of this line than to the south, broadly coincident with the mapped boundary between southern deposits mapped as pre-Nectarian age and those of Nectarian–Imbrian age to the north. The radar-bright crater population is consistent with a megaregolith thickness of ~1.5 km in the north and ~2.5 km in the south, a difference we attribute to South Pole–Aitken basin ejecta.

Effects of ejecta accumulation on the crater population of asteroid 433 Eros

The crater population of asteroid 433 Eros exhibits a deficit in small crater diameters that has been suggested to result from impact‐induced seismic shaking initiating downslope movements of regolith material, covering these small craters. As in lunar maria, saturation equilibrium was expected to characterize the crater population of Eros, but was surprisingly not shown by the data set. The surface of Eros displays evidence of burial by regolith especially for boulders, suggesting that ejecta coverage erases the craters in addition to seismic shaking erasure. In this work we investigate the production and erasure of craters by impact ejecta and compare derived crater size distributions with those measured for Eros. We simulate a bombardment of Eros by an impactor population derived from the Main Asteroid Belt and estimate the crater and ejecta characteristics with a scaling law, allowing ejecta to progressively create a regolith blanket. Assuming the contribution of the ejecta blanketing process only, we find a good agreement between the simulated and the observed population of 250 m to 4 km diameter craters for exposure times of 600 Ma and 400 Ma. This suggests a major impact or breakup that occurred about 500 Ma ago, inducing a surface reset. A mismatch for craters with a diameter smaller than ∼100 m remains, indicating that seismic shaking (or another erasure process) is still necessary to explain their low number. Our simulations emphasize the importance of an accurate modeling of both processes to fully understand and interpret the small body size‐frequency crater curves.

The Yarkovsky effect is not responsible for small crater depletion on Eros and Itokawa

Abstract The near-Earth Asteroids Eros and Itokawa show a pronounced lack of small (≲100 m) craters, the vast majority of which were formed during their time in the main belt, and this has been cited as possible evidence that small (≲10 m) impactors are efficiently removed from the main belt by the Yarkovsky effect. Using well-tested models for the evolution of the main-belt size distribution and the evolution of crater populations on asteroid surfaces, I show that a pronounced lack of small impactors would require size-dependent removal far stronger than can result from the Yarkovsky effect (or any other known process). Furthermore, such strong removal would lead to wavelike perturbations in the main-belt and near-Earth asteroid size distributions that are inconsistent with their observed size distributions, as well as the cratering records on asteroid surfaces. A more likely explanation is that processes on asteroid surfaces, such as seismic shaking, are responsible for erasing small craters after they form.

Geological context of water‐altered minerals in Valles Marineris, Mars

Greater than 15,000 km2 of the layered deposits within Valles Marineris are associated with water‐altered minerals, yet their origin and history of alteration remain a mystery. There are numerous competing hypotheses for the formation of the interior layered deposits including aeolian, lacustrine, and volcanic. Recent orbiter spectroscopic data have indicated that water has played a role in their geological history. Thermal Emission Spectrometer (TES) measurements have revealed significant crystalline hematite‐bearing deposits within Valles Marineris, typically related to interior layered deposits. These hematite deposits, found with a wide range of albedo values, are associated with relatively steep bedrock exposures but can also be seen downslope on flat surfaces where they may be a lag deposit. More recently, Observatoire la Minéralogie, l'Eau, les Glaces, et l'Activité (OMEGA) data have shown hydrated sulfates covering more than 13,000 km2 area of Valles Marineris. Sulfates are found in numerous topographic settings and geological units, but are typically located along the flanks of interior layered deposits and nearby low‐lying floor units. Here we study the detailed morphologies of hematite and sulfate‐bearing units such as mantled wall units, mass‐wasting blocky deposits, massive floor deposits, and tectonically altered floor units. All of these terrains have diverse erosional styles and varied crater populations. In both hematite‐ and sulfate‐bearing units, occasionally found in conjunction with one another, formation processes require contributions from water. The results indicate a wide range of diversity within an individual mineral class, between mineral classes, and also among morphological types. The diversity of geological settings and properties suggest that any single, unified formation mechanism is improbable.

Theoretical analysis of secondary cratering on Mars and an image-based study on the Cerberus Plains

The issue of crater retention age estimates on planetary surfaces is discussed with an attempt to quantify the effect of overlapping primary and secondary impact crater populations in restricted crater diameter ranges. The approach to this problem is illustrated with a simple model production function where the secondary crater input is artificially enhanced. Extrapolation of such a secondary crater model distribution to a global record results in extraordinarily high crater frequencies that do not exist on Mars, and implies the need of detailed studies of the size-frequency distribution for remote secondary craters, to date poorly known. A key case, the martian crater Zunil and its secondary crater field, illustrate that reasonable predictions for the secondary crater size-frequency distribution at small (<100 m) crater diameters affected the standard model crater retention age for the Cerberus plains less than the statistical uncertainty. These observations show that age determination based on appropriate crater counting statistics is valid in a wide primary crater diameter range.

Polygonal impact craters in Argyre region, Mars: Implications for geology and cratering mechanics

Abstract— Impact craters are not always circular; sometimes their rims are composed of several straight segments. Such polygonal impact craters (PICs) are controlled by pre‐existing target structures, mainly faults or other similar planes of weakness. In the Argyre region, Mars, PICs comprise ˜ 17% of the total impact crater population (&gt;7 km in diameter), and PICs are relatively more common in older geologic units. Their formation is mainly controlled by radial fractures induced by the Argyre and Ladon impact basins, and to a lesser extent by the basin‐concentric fractures. Also basin‐induced conjugate shear fractures may play a role. Unlike the PICs, ridges and graben in the Argyre region are mostly controlled by Tharsis‐induced tectonism, with the ridges being concentric and graben radial to Tharsis. Therefore, the PICs primarily reflect an old impact basin‐centered tectonic pattern, whereas Tharsis‐centered tectonism responsible for the graben and the ridges has only minor influence on the PIC rim orientations.According to current models of PIC formation, complex PICs should form through a different mechanism than simple PICs, leading to different orientations of straight rim segments. However, when simple and complex PICs from same areas are studied, no statistically significant difference can be observed. Hence, in addition to enhanced excavation parallel to the strike of fractures (simple craters) and slumping along the fracture planes (complex craters), we propose a third mechanism involving thrusting along the fracture planes. This model is applicable to both simple and small complex craters in targets with some dominating orientations of structural weakness.

Martian pedestal craters: Marginal sublimation pits implicate a climate‐related formation mechanism

Pedestal craters on Mars are defined by an outward‐facing scarp forming a plateau perched tens of meters above the surrounding terrain. Their origin has been attributed to impact armoring of the surface and subsequent removal of inter‐crater terrain by either eolian deflation or sublimation of an ice‐rich substrate. We identified 2696 pedestal craters between ∼60°N and 60°S latitude; 98% are poleward of 33°N and 40°S. The majority of pedestal crater margins are smoothly sloped, but ∼3%, concentrated in Utopia Planitia and Malea Planum, display distinctive marginal pits. These pedestal crater scarps are anomalously tall (usually &gt;80–100 m) and the pits resemble sublimation depressions seen on Earth and elsewhere on Mars, providing evidence for sublimation of volatiles in the scarp, where the armored surface has tapered. The pitted scarps provide insight into the origin of the general pedestal crater population, favoring formation via deposition of a volatile‐rich substrate, impact armoring, and sublimation of intervening volatiles. Crater densities and overlapping pedestal craters suggest multiple periods of emplacement and loss of these climate‐related, latitude‐dependent deposits throughout the Amazonian.

Small-scale topography of 25143 Itokawa from the Hayabusa laser altimeter

The surface topography of Asteroid 25143 Itokawa is explored using the LIght Detection And Ranging instrument (LIDAR). The data confirm the presence of a rough highland and a smooth lowland. The highland is dominated by boulders, but also possesses topography associated with surface lineaments and broad surface facets. The boulders ensure that the roughness of the highlands over short distances is typically greater relative to most surfaces on 433 Eros. Over larger distances, Itokawa is always smoother than Eros possibly because of its smaller size and weak rubble-pile structure. The lowlands of Itokawa are very smooth, and are typically devoid of boulders. Some transitional regions midway between the highlands and lowlands also exist. In these areas, craters that retain their regolith fill possess flat floors and resemble “ponds” seen on 433 Eros. Analyses of surface elevation, imagery and a quantitative measure of surface roughness are consistent with regolith flowing downhill from the highlands to fill in the low areas of Itokawa, probably covering up any pre-existing rough terrain. Using this interpretation, we find a minimum 2.3 ± 0.4 m thick layer of regolith in the lowlands, which, if spread evenly across the entire asteroid, corresponds to a 42 ± 1 cm thick layer. It is very difficult to generate this amount of regolith with the population of craters seen on Itokawa. However, an Itokawa composed of several large masses may have retained this regolith during its formation. The presence of such large masses could account for the observed lineaments and what appear to be exposures of bedrock on the largest steep slope observed.

The Borealis basin and the origin of the martian crustal dichotomy

The most prominent feature on the surface of Mars is the near-hemispheric dichotomy between the southern highlands and northern lowlands. The root of this dichotomy is a change in crustal thickness along an apparently irregular boundary, which can be traced around the planet, except where it is presumably buried beneath the Tharsis volcanic rise. The isostatic compensation of these distinct provinces and the ancient population of impact craters buried beneath the young lowlands surface suggest that the dichotomy is one of the most ancient features on the planet. However, the origin of this dichotomy has remained uncertain, with little evidence to distinguish between the suggested causes: a giant impact or mantle convection/overturn. Here we use the gravity and topography of Mars to constrain the location of the dichotomy boundary beneath Tharsis, taking advantage of the different modes of compensation for Tharsis and the dichotomy to separate their effects. We find that the dichotomy boundary along its entire path around the planet is accurately fitted by an ellipse measuring approximately 10,600 by 8,500 km, centred at 67 degrees N, 208 degrees E. We suggest that the elliptical nature of the crustal dichotomy is most simply explained by a giant impact, representing the largest such structure thus far identified in the Solar System.

The timing of martian valley network activity: Constraints from buffered crater counting

Valley networks, concentrations of dendritic channels that often suggest widespread pluvial and fluvial activity, have been cited as indicators that the climate of Mars differed significantly in the past from the present hyperarid cold desert conditions. Some researchers suggest that the change in climate was abrupt, while others favor a much more gradual transition. Thus, the precise timing of valley network formation is critical to understanding the climate history on Mars. We examine thirty valley network-incised regions on Mars, including both cratered upland valley networks and those outside the uplands, and apply a buffered crater counting technique to directly constrain when valley network formation occurred. The crater populations that we derive using this approach allow assessment of the timing of the last activity in a valley network independent of the mapping of specific geological units. From these measurements we find that valley networks cluster into two subdivisions in terms characteristics and age: (1) valley network activity in the cratered highlands has an average cessation age at the Noachian–Hesperian boundary and all valleys that we crater counted are Early Hesperian or older. No evidence is found for valley networks in the cratered uplands of Late Hesperian or Amazonian age. The timing of the cessation of cratered upland valley network activity at the Noachian–Hesperian boundary also corresponds to a decline in the intensity of large crater formation and degradation and to the apparent end of phyllosilicate-type weathering. (2) A few valley network-incised regions formed outside of the cratered uplands on volcanic edifices, in association with younger impact craters, and on the rim of Valles Marineris. We applied our buffered crater counting technique to four such valleys, on the volcanoes Ceraunius Tholus, Hecates Tholus, and Alba Patera and on the rim of Echus Chasma, and find that each has distinctive and different Late Hesperian or Early Amazonian ages, indicating that valley networks formed from time to time in the post-Noachian period. Unlike the cratered upland valley networks, these isolated occurrences are very local and have been interpreted to represent local conditions (e.g., snowpack melted during periods of intrusive volcanic activity). In contrast to a gradual cessation in the formation of valley networks proposed by some workers, our new buffered crater counting results indicate a relatively abrupt cessation in the formation of the widespread cratered upland valley networks at approximately the end of the Noachian, followed only by episodic and very localized valley network formation in later Mars history, very likely due to specific conditions (e.g., local magmatic heating). These valley network ages and correlations are thus consistent with a major change in the near-surface aqueous environment on Mars at approximately the Noachian–Hesperian boundary. The Noachian environment supported surface running water and fluvial erosion across Mars in the cratered uplands, enhanced crater degradation, and a weathering environment favoring the formation of phyllosilicates. The Hesperian–Amazonian environment was more similar to the hyperarid cold desert of today, with valley networks forming only extremely rarely and confined to localized special conditions. Sources of water for these latter occurrences are likely to be related to periodic mobilization and equatorward migration of polar volatiles due to variations in spin-axis orbital parameters, and to periodic catastrophic emergence of groundwater.

Layered mantling deposits in northeast Arabia Terra, Mars: Noachian‐Hesperian sedimentation, erosion, and terrain inversion

Thick, layered mantling deposits of different ages occur in several nonpolar regions of Mars and are thought to represent volcanic ash and/or climate‐related ice‐dust deposition. One such deposit is a layered mantling unit that unconformably blanketed highlands terrain in northeast Arabia Terra during the Late Noachian and/or earliest Hesperian. Shortly thereafter, by the mid‐Hesperian, this deposit was substantially eroded; on the basis of its superposed crater population, it appears to have subsequently retained its approximate morphology and distribution in the ∼3.5 Gyr since. When the erosion occurred, in the Early Hesperian, the mantling unit was more resistant in some areas than its surroundings, resulting in inversion of relief: craters were transformed into highstanding buttes, and valleys were transformed into isolated ridges. On the basis of the scale of the observed inversion of relief, the magnitude of erosion in northeast Arabia Terra was substantial, up to hundreds of meters. We have used newly available data to assess the nature of the mantling material and probable mechanisms for its deposition and removal. We find that (1) the mantle unit is layered at scales of meters to tens of meters; (2) the deposit has a median thickness of ∼60 m, and the thickness of the deposit was locally controlled by the topography at the time of its deposition; (3) the mantling unit thins toward the south; and (4) characteristics of the mantling unit, such as its induration, as well as observed pits and fractures, suggest that volatiles may have been incorporated into the mantling unit either during or following its emplacement. Taken together, our observations lead us to favor models for the deposition of the mantling unit involving airfall of dust or ash. Either climate‐driven dust deposition or plinian volcanic eruptions from the nearby Syrtis Major volcanic province are plausible candidate models for the emplacement of the unit in time and space.

Chemical composition of lunar meteorites and the lunar crust

The paper presents the first analyses of major and trace elements in 19 lunar meteorites newly found in Oman. These and literature data were used to assay the composition of highland, mare, and transitional (high- land-mare interface) regions of the lunar surface. The databank used in the research comprises data on 44 mete- orites weighing 11 kg in total, which likely represent 26 individual falls. Our data demonstrate that the lunar highland crust should be richer in Ca and Al but poorer in mafic and incompatible elements than it was thought based on studying lunar samples and the first orbital data. The Ir concentration in the highland crust and the analysis of lunar crater population suggest that most lunar impactites were formed by a single major impact event, which predetermined the geochemical characteristics of these rocks. Lunar mare regions should be dom- inated by low-Ti basalts, which are, however, enriched in LREEs compared to those sampled by lunar missions. The typical material of mare-highland interface zones can contain KREEP and magnesian VLT basalts. The composition of the lunar highland crust deduced from the chemistry of lunar meteorites does not contradict the model of the lunar magma ocean, but the average composition of lunar mare meteorites is inconsistent with this concept and suggests assimilation of KREEP material by basaltic magmas. The newly obtained evaluations of the composition of the highland crust confirm that the Moon can be enriched in refractory elements and depleted in volatile and siderophile elements.