Thomas-Keprta^a et al. (>10)*
*Find the extensive, full author and affiliation list on the publishers website.

^aJacobs, NASA Johnson Space Center, Mail Code JE-23, Building 31, Houston TX 77058

Carbonaceous matter on the surfaces of black pyroclastic beads, collected from Shorty crater during the Apollo 17 mission, represents the first identification of complex organic material associated with any lunar sample. We report the chemical, physical and isotopic properties of this organic matter that together support a pre-terrestrial origin. We suggest the most probable source is through the accretion of exogenous meteoritic kerogen from micrometeorite impacts into the lunar regolith. Abiotic organic matter has been continuously delivered to the surfaces of the terrestrial planets and their moons by accretion of asteroidal and cometary material. Determining the nature, distribution and evolution of such matter in the lunar regolith has important implications for understanding the prebiotic chemical inventory of the terrestrial planets.

Reference
Thomas-Keprta et al. (in press) Organic Matter on the Earth’s Moon. Geochimica et Cosmochimica Acta
[doi:10.1016/j.gca.2014.02.047]
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M. M. Woolfson

Physics Department, University of York, Heslington, York, YO10 5DD, UK

We currently do not have a copyright agreement with the publisher and cannot display the abstract here.

Reference
Woolfson MM (2013) A Postulated Planetary Collision, the Terrestrial Planets, the Moon and Smaller Solar-System Bodies. Earth, Moon, and Planets 111:1-14.
[doi:10.1007/s11038-013-9420-8]

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B. B. Ochsendorf¹, N. L. J. Cox², S. Krijt¹, F. Salgado¹, O. Berné⁴, J. P. Bernard⁴, L. Kaper³ and A. G. G. M. Tielens¹

¹Leiden Observatory, Leiden University, PO Box 9513 2300 RA Leiden The Netherlands
²Instituut voor Sterrenkunde, K.U. Leuven, Celestijnenlaan 200D, bus 2401, 3001 Leuven, Belgium
³Sterrenkundig Instituut Anton Pannekoek, University of Amsterdam, Science Park 904, PO Box 94249, 1090 GE Amsterdam, The Netherlands
⁴Université de Toulouse, UPS-OMP, IRAP, 31028 Toulouse, France

Observations obtained with the Spitzer Space Telescope and the WISE satellite have revealed a prominent arc-like structure at 50′′ (≃0.1 pc) from the O9.5V/B0.5V system σOri AB. We measure a total dust mass of 2.3 ± 1.5 × 10^-5  M_⊙. The derived dust-to-gas mass ratio is ≃0.29 ± 0.20. We attribute this dust structure to the interaction of radiation pressure from the star with dust carried along by the IC 434 photo-evaporative flow of ionized gas from the dark cloud L1630. We have developed a quantitative model for the interaction of a dusty ionized flow with nearby (massive) stars where radiation pressure stalls dust, piling it up at an appreciable distance (>0.1 pc), and force it to flow around the star. The model demonstrates that for the conditions in IC 434, the gas will decouple from the dust and will keep its original flow lines. Hence, we argue that this dust structure is the first example of a dust wave created by a massive star moving through the interstellar medium. Our model shows that for higher gas densities, coupling is more efficient and a bow wave will form, containing both dust and gas. Our model describes the physics of dust waves and bow waves and quantitatively reproduces the optical depth profile at 70 μm. Dust waves (and bow waves) stratify dust grains according to their radiation pressure opacity, which reflects the size distribution and composition of the grain material. It is found that in the particular case of σ Ori AB, dust is able to survive inside the ionized region. Comparison of our model results with observations implies that dust-gas coupling through Coulomb interaction is less important than previously thought, challenging our understanding of grain dynamics in hot, ionized regions of space. We describe the difference between dust (and bow) waves and classical bow shocks created by the interaction of a stellar wind with the interstellar medium. The results show that for late O-type stars with weak stellar winds, the stand-off distance of the resulting bow shock is very close to the star, well within the location of the dust wave. In general, we conclude that dust waves and bow waves should be common around stars showing the weak-wind phenomenon, i.e., stars with log(L/L_⊙)<5.2, and that these structures are best observed at mid-IR to FIR wavelengths, depending on the stellar spectral type. In particular, dust waves and bow waves are most efficiently formed around weak-wind stars moving through a high density medium. Moreover, they provide a unique opportunity to study the direct interaction between a (massive) star and its immediate surroundings.

Reference
Ochsendorf BB, Cox NLJ, Krijt S, Salgado F, Berné O, Bernard JP, Kaper L and Tielens AGGM (2014) Blowing in the wind: The dust wave around σ Orionis AB. Astronomy & Astrophysics 563:A65.
[doi:10.1051/0004-6361/201322873]
Reproduced with permission © ESO

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Catherine Walsh^1,2, Tom. J. Millar², Hideko Nomura^3,4,5, Eric Herbst^6,7, Susanna Widicus Weaver⁸, Yuri Aikawa⁹, Jacob C. Laas⁸ and Anton I. Vasyunin^10,11

¹Leiden Observatory, Leiden University, PO Box 9513, 2300 RA Leiden, The Netherlands
²Astrophysics Research Centre, School of Mathematics and Physics, Queen’s University Belfast, University Road, Belfast BT7 1NN, UK
³Department of Astronomy, Graduate School of Science, Kyoto University, 606-8502 Kyoto, Japan
⁴National Astronomical Observatory of Japan, Osawa, Mitaka, Tokyo 181-8588, Japan
⁵Department of Earth and Planetary Sciences, Tokyo Institute of Technology, 2-12-1 Ookayama, Meguro-ku, 152-8551 Tokyo, Japan
⁶Departments of Physics, Chemistry and Astronomy, The Ohio State University, Columbus OH 43210, USA
⁷Departments of Chemistry, Astronomy, and Physics, University of Virginia, Charlottesville VA 22904, USA
⁸Department of Chemistry, Emory University, Atlanta GA 30322, USA
⁹Department of Earth and Planetary Sciences, Kobe University, 1-1 Rokkodai-cho, Nada, 657-8501 Kobe, Japan
¹⁰Department of Chemistry, University of Virginia, Charlottesville VA 22904, USA
¹¹Visiting Scientist, Ural Federal University, 620075 Ekaterinburg, Russia

Context. Protoplanetary disks are vital objects in star and planet formation, possessing all the material, gas and dust, which may form a planetary system orbiting the new star. Small, simple molecules have traditionally been detected in protoplanetary disks; however, in the ALMA era, we expect the molecular inventory of protoplanetary disks to significantly increase.
Aims. We investigate the synthesis of complex organic molecules (COMs) in protoplanetary disks to put constraints on the achievable chemical complexity and to predict species and transitions which may be observable with ALMA.
Methods. We have coupled a 2D steady-state physical model of a protoplanetary disk around a typical T Tauri star with a large gas-grain chemical network including COMs. We compare the resulting column densities with those derived from observations and perform ray-tracing calculations to predict line spectra. We compare the synthesised line intensities with current observations and determine those COMs which may be observable in nearby objects. We also compare the predicted grain-surface abundances with those derived from cometary comae observations.
Results. We find COMs are efficiently formed in the disk midplane via grain-surface chemical reactions, reaching peak grain-surface fractional abundances ~10^-6–10^-4 that of the H nuclei number density. COMs formed on grain surfaces are returned to the gas phase via non-thermal desorption; however, gas-phase species reach lower fractional abundances than their grain-surface equivalents, ~10^-12–10^-7. Including the irradiation of grain mantle material helps build further complexity in the ice through the replenishment of grain-surface radicals which take part in further grain-surface reactions. There is reasonable agreement with several line transitions of H₂CO observed towards T Tauri star-disk systems. There is poor agreement with HC₃N lines observed towards LkCa 15 and GO Tau and we discuss possible explanations for these discrepancies. The synthesised line intensities for CH₃OH are consistent with upper limits determined towards all sources. Our models suggest CH₃OH should be readily observable in nearby protoplanetary disks with ALMA; however, detection of more complex species may prove challenging, even with ALMA “Full Science” capabilities. Our grain-surface abundances are consistent with those derived from cometary comae observations providing additional evidence for the hypothesis that comets (and other planetesimals) formed via the coagulation of icy grains in the Sun’s natal disk.

Reference
Walsh C, Millar TJ, Nomura H, Herbst E, Weaver SW, Aikawa Y, Laas JC and Vasyunin AI (2014) Complex organic molecules in protoplanetary disks. Astronomy & Astrophysics 563:A33.
[doi:10.1051/0004-6361/201322446]
Reproduced with permission © ESO

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Stephanie C. Werner¹, Anouck Ody², François Poulet³

¹The Centre for Earth Evolution and Dynamics, University of Oslo, Sem Sælandsvei 24, 0371 Oslo, Norway.
²Laboratoire de Géologie de Lyon: Terre, Planètes, Environnement, Université de Lyon 1 (CNRS, ENS-Lyon, Université de Lyon), rue Raphaël Dubois 2, 69622 Villeurbanne, France.
³Institut d’Astrophysique Spatiale, Université Paris Sud 11, Bâtiment 121, 91405 Orsay, France.

Absolute ages for planetary surfaces are often inferred by crater densities and only indirectly constrained by the ages of meteorites. We show that the <5 million-year-old and 55-km-wide Mojave Crater on Mars is the ejection source for the meteorites classified as shergottites. Shergottites and this crater are linked by their coinciding meteorite ejection ages and the crater formation age and by mineralogical constraints. Because Mojave formed on 4.3 billion–year-old terrain, the original crystallization ages of shergottites are old, as inferred by Pb-Pb isotope ratios, and the much-quoted shergottite ages of <600 million years are due to resetting. Thus, the cratering-based age determination method for Mars is now calibrated in situ, and it shifts the absolute age of the oldest terrains on Mars backward by 200 million years.

Reference
Werner SC, Ody A and Poulet F (2014) The Source Crater of Martian Shergottite Meteorites. Science 343:1343-1346.
[doi:10.1126/science.1247282]
Reprinted with permission from AAAS

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Krohn^a et al. (>10)*
*Find the extensive, full author and affiliation list on the publishers website.

^aInstitute of Planetary Research, German Aerospace Center (DLR), Berlin, Germany

The Quadrangles Av-11 and Av-12 on Vesta are located at the northern rim of the giant Rheasilvia south polar impact basin. The primary geologic units in Av-11 and Av-12 include material from the Rheasilvia impact basin formation, smooth material and different types of impact crater structures (such as bimodal craters, dark and bright crater ray material and dark ejecta material). Av-11 and Av-12 exhibit almost the full range of mass wasting features observed on Vesta, such as slump blocks, spur-and-gully morphologies and landslides within craters. Processes of collapse, slope instability and seismically triggered events force material to slump down crater walls or scarps and produce landslides or rotational slump blocks. The spur-and-gully morphology that is known to form on Mars is also observed on Vesta; however, on Vesta this morphology formed under dry conditions.

Reference
Krohn et al. (in press) Mass movement on Vesta at steep scarps and crater rims. Icarus
[doi:10.1016/j.icarus.2014.03.013]
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Blewett^a et al. (>10)*
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^aPlanetary Exploration Group, Space Department, Johns Hopkins University Applied Physics Laboratory, MS 200-W230, 11100 Johns Hopkins Road, Laurel, Maryland 20723 USA

As part of systematic global mapping of Vesta using data returned by the Dawn spacecraft, we have produced a geologic map of the north pole quadrangle, Av-1 Albana. Extensive seasonal shadows were present in the north polar region at the time of the Dawn observations, limiting the ability to map morphological features and employ color or spectral data for determination of composition. The major recognizable units present include ancient cratered highlands and younger crater-related units (undivided ejecta, and mass-wasting material on crater floors). The antipode of Vesta’s large southern impact basins, Rheasilvia and Veneneia, lie within or near the Av-1 quadrangle. Therefore it is of particular interest to search for evidence of features of the kind that are found at basin antipodes on other planetary bodies. Albedo markings known as lunar swirls are correlated with basin antipodes and the presence of crustal magnetic anomalies on the Moon, but lighting conditions preclude recognition of such albedo features in images of the antipode of Vesta’s Rheasilvia basin. “Hilly and lineated terrain,” found at the antipodes of large basins on the Moon and Mercury, is not present at the Rheasilvia or Veneneia antipodes. We have identified small-scale linear depressions that may be related to increased fracturing in the Rheasilvia and Veneneia antipodal areas, consistent with impact-induced stresses (  and ). The general high elevation of much of the north polar region could, in part, be a result of uplift caused by the Rheasilvia basin-forming impact, as predicted by numerical modeling ( Bowling et al., 2013). However, stratigraphic and crater size-frequency distribution analysis indicate that the elevated terrain predates the two southern basins and hence is likely a remnant of the ancient vestan crust. The lack of large-scale morphological features at the basin antipodes can be attributed to weakened antipodal constructive interference of seismic waves caused by an oblique impact or by Vesta’s non-spherical shape, or by attenuation of seismic waves because of the physical properties of Vesta’s interior. A first-order analysis of the Dawn global digital elevation model for Vesta indicates that areas of permanent shadow are unlikely to be present in the vicinity of the north pole.

Reference
Blewett et al. (in press) Vesta’s North Pole Quadrangle Av-1 (Albana): Geologic Map and the Nature of the South Polar Basin Antipodes. Icarus
[doi:10.1016/j.icarus.2014.03.007]
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David A. Williams^a, R. Aileen Yingst^b and W. Brent Garry^c

^aSchool of Earth & Space Exploration, Arizona State University, Tempe, Arizona 85287-1404.
^bPlanetary Science Institute, Tucson, Arizona
^cNASA Goddard Spaceflight Center, Greenbelt, Maryland

The purpose of this paper is to introduce the Geologic Mapping of Vesta Special Issue/Section of Icarus, which includes several papers containing geologic maps of the surface of Vesta made to support data analysis conducted by the Dawn Science Team during the Vesta Encounter (July 2011-September 2012). In this paper we briefly discuss pre-Dawn knowledge of Vesta, provide the goals of our geologic mapping campaign, discuss the methodologies and materials used for geologic mapping, review the global geologic context of Vesta, discuss the challenges of mapping the geology of Vesta as a small airless body, and describe the content of the papers in this Special Issue/Section. We conclude with a discussion of lessons learned from our quadrangle-based mapping effort and provide recommendations for conducting mapping campaigns as part of planetary spacecraft nominal missions.

Reference
Williams DA, Yingst RA and Garry WB (in press) Introduction: The Geologic Mapping of Vesta. Icarus
[doi:10.1016/j.icarus.2014.03.001]
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Cyrena Anne Goodrich^a,b, George E. Harlow^c, James A. Van Orman^d, Stephen R. Sutton^e, Michael J. Jercinovic^b, Takashi Mikouchi^f

^aPlanetary Science Institute, 1700 E. Ft. Lowell, Suite 106, Tucson, AZ 85719 USA
^bDepartment of Geosciences, University of Massachusetts, 611 North Pleasant Street, Amherst, MA 01003 USA
^cAmerican Museum of Natural History, Department of Earth and Planetary Sciences, Central Park West at 79[th] Street, New York, NY 10024 USA
^dDept. of Earth, Environmental and Planetary Sciences, Case Western Reserve University, Cleveland, OH 44120 USA
^eDept. of Geophysical Sciences and Center for Advanced Radiation Sources, University of Chicago, Chicago, IL 60637 USA
^fDepartment of Earth and Planetary Sciences, University of Tokyo, Tokyo 113-0033 Japan

Ureilites are ultramafic achondrites thought to be residues of partial melting on a carbon-rich asteroid. They show a trend of FeO-variation (olivine Fo from ~74 to 95) that suggests variation in oxidation state. Whether this variation was established during high-temperature igneous processing on the ureilite parent body (UPB), or preserved from nebular precursors, is a subject of debate. The behavior of chromium in ureilites offers a way to assess redox conditions during their formation and address this issue, independent of Fo. We conducted a petrographic and mineral compositional study of occurrences of chromite (Cr-rich spinel) in ureilites, aimed at determining the origin of the chromite in each occurrence and using primary occurrences to constrain models of ureilite petrogenesis. Chromite was studied in LEW 88774 (Fo 74.2), NWA 766 (Fo 76.7), NWA 3109 (Fo 76.3), HaH 064 (Fo 77.5), LAP 03587 (Fo 74.9), CMS 04048 (Fo 76.4), LAP 02382 (Fo 78.6) and EET 96328 (Fo 85.2).
Chromite occurs in LEW 88774 (~5 vol.%), NWA 766 (<1 vol.%), NWA 3109 (<1 vol.%) and HaH 064 (<1 vol.%) as subhedral to anhedral grains comparable in size (~30 μm to 1 mm) and/or textural setting to the major silicates (olivine and pyroxenes[s]) in each rock, indicating that it is a primary phase. The most FeO-rich chromites in these sample (rare grain cores or chadocrysts in silicates) are the most primitive compositions preserved (fe# = 0.55-0.6; Cr# varying from 0.65 to 0.72 among samples). They record olivine-chromite equilibration temperatures of ~1040-1050°C, reflecting subsolidus Fe/Mg reequilibration during slow cooling from ~1200-1300°C. All other chromite in these samples is reduced. Three types of zones are observed. 1) Inclusion-free interior zones showing reduction of FeO (fe# ~0.4→0.28); 2) Outer zones showing further reduction of FeO (fe# ~0.28→0.15) and containing abundant laths of eskolaite-corundum (Cr₂O₃-Al₂O₃); 3) Outermost zones showing extreme reduction of both FeO (fe# <0.15) and Cr₂O₃ (Cr# as low as 0.2). The grains are surrounded by rims of Si-Al-rich glass, graphite, Fe,Cr-carbides ([Fe,Cr]₃C and [Fe,Cr]₇C₃), Cr-rich sulfides (daubréelite and brezinaite) and Cr-rich symplectic bands on adjacent silicates. Chromite is inferred to have been reduced by graphite, forming eskolaite-corundum and carbides as byproducts, during impact excavation. This event involved initial elevation of T (to 1300-1400°C), followed by rapid decompression and drop in T (to <700°C) at 1-20°C/hr. The kinetics of reduction of chromite is consistent with this scenario. The reduction was facilitated by silicate melt surrounding the chromites, which was partly generated by shock-melting of pyroxenes. Symplectic bands, consisting of fine-scale intergrowths of Ca-pyroxene, chromite and glass, formed by reaction between the Cr-enriched melt and adjacent silicates.
Early chromite also occurs in a melt inclusion in olivine in HaH 064 and in a metallic spherule in olivine in LAP 02382. LAP 03587 and CMS 04048 contain ⩽μm-sized chromite+pyroxene symplectic exsolutions in olivine, indicating high Cr valence in the primary olivine. EET 96328 contains a round grain of chromite that could be a late-crystallizing phase. Tiny chromite grains in melt inclusions in EET 96328 formed in late, closed-system reactions.
For 7 of the 8 ureilites we conclude that the relatively oxidizing conditions evidenced by the presence of primary or early chromite pertain to the period of high-T igneous processing. The observation that such conditions are recorded almost exclusively in low-Fo samples supports the interpretation that the ureilite FeO-variation was established during igneous processing on the UPB.

Reference
Goodrich CA, Harlow GE, Van Orman JA, Sutton SR, Jercinovic MJ and Mikouchi T (in press) Petrology of Chromite in Ureilites: Deconvolution of Primary Oxidation States and Secondary Reduction Processes. Geochimica et Cosmochimica Acta
[doi:10.1016/j.gca.2014.02.028]
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J.-P. Williams^a, A.V. Pathare^b and O. Aharonson^c

^aDept. Earth and Space Sciences, University of California, Los Angeles, CA 90095, USA
^bPlanetary Science Institute, Tucson, AZ 85719, USA
^cHelen Kimmel Center for Planetary Science, Weizmann Institute Of Science, Rehovot, 76100 Israel

We model the primary crater production of small (D < 100 m) primary craters on Mars and the Moon using the observed annual flux of terrestrial fireballs. From the size-frequency distribution (SFD) of meteor diameters, with appropriate velocity distributions for Mars and the Moon, we are able to reproduce martian and lunar crater-count chronometry systems (isochrons) in both slope and magnitude. We include an atmospheric model for Mars that accounts for the deceleration, ablation, and fragmentation of meteors. We find that the details of the atmosphere or the fragmentation of the meteors do not strongly influence our results. The downturn in the crater SFD from atmospheric filtering is predicted to occur at D ~ 10-20 cm, well below the downturn observed in the distribution of fresh craters detected by the Mars Global Surveyor (MGS) Mars Orbiter Camera (MOC) or the Mars Reconnaissance Orbiter (MRO) Context Camera (CTX). Crater counts are conducted on the ejecta blanket of Zunil crater and the interior of Pangboche crater on Mars and North Ray and Cone craters on the Moon. Our model isochrons produce a similar slope and age estimate for the formation of Zunil crater as the Hartmann production function (~1 Ma). We derive an age of 35.1 Ma for Pangboche when accounting for the higher elevation (>20 km higher than Zunil), a factor ~2 younger than estimated using the Hartmann production function which assumes 6 mbar surface pressure. We estimate ages of 52.3 Ma and 23.9 Ma for North Ray and Cone crater respectively, consistent with cosmic ray exposure ages from Apollo samples. Our results indicate that the average cratering rate has been constant on these bodies over these time periods. Since our Monte Carlo simulations demonstrate that the existing crater chronology systems can be applied to date young surfaces using small craters on the Moon and Mars, we conclude that the signal from secondary craters in the isochrons must be relatively small at these locations, as our Monte Carlo model only generates primary craters.

Reference
Williams J-P, Pathare AV and O. Aharonson O (in press) The Production of Small Primary Craters on Mars and the Moon. Icarus
[doi:10.1016/j.icarus.2014.03.011]
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