Gasa impact crater,Mars: Very young gullies formed from impact into latitude-dependent mantle and debris-covered glacier deposits?

The mode of formation of gullies on Mars, very young erosional–depositional landforms consisting of an alcove, channel, and fan, is one of the most enigmatic problems in martian geomorphology. Major questions center on their ages, geographic and stratigraphic associations, relation to recent ice ages, and, if formed by flowing water, the sources of the water to cause the observed erosion/deposition. Gasa (35.72°S, 129.45°E), a very fresh 7-km diameter impact crater and its environment, offer a unique opportunity to explore these questions. We show that Gasa crater formed during the most recent glacial epoch (2.1–0.4 Ma), producing secondary crater clusters on top of the latitude-dependent mantle (LDM), interpreted to be a layered ice-dust-rich deposit emplaced during this glacial epoch. High-resolution images of a pre-Gasa impact crater ～100 km northeast of Gasa reveal that portions of the secondary-crater-covered LDM have been removed from pole-facing slopes in crater interiors near Gasa; gullies are preferentially located in these areas and channels feeding alcoves and fans can be seen to emerge from the eroding LDM layers to produce multiple generations of channel incision and fan lobes. We interpret these data to mean that these gullies formed extremely recently in the post-Gasa-impact time-period by melting of the ice-rich LDM. Stratigraphic and topographic relationships are interpreted to mean that under favorable illumination geometry (steep pole-facing slopes) and insolation conditions, melting of the debris-covered ice-rich mantle took place in multiple stages, most likely related to variations in spin-axis/orbital conditions. Closer to Gasa, in the interior of the ～18 km diameter LDM-covered host crater in which Gasa formed, the pole-facing slopes display two generations of gullies. Early, somewhat degraded gullies, have been modified by proximity to Gasa ejecta emplacement, and later, fresh appearing gullies are clearly superposed, cross-cut the earlier phase, and show multiple channels and fans, interpreted to be derived from continued melting of the LDM on steep pole-facing slopes. Thus, we conclude that melting of the ice-rich LDM is a major source of gully activity both pre-Gasa crater and post-Gasa crater formation. The lack of obscuration of Gasa secondary clusters formed on top of the LDM is interpreted to mean that the Gasa impact occurred following emplacement of the last significant LDM layers at these low latitudes, and thus near the end of the ice ages. This interpretation is corroborated by the lack of LDM within Gasa. However, Gasa crater contains a robustly developed set of gullies on its steep, pole-facing slopes, unlike other very young post-LDM craters in the region. How can the gullies inside Gasa form in the absence of an ice-rich LDM that is interpreted to be the source of water for the other adjacent and partly contemporaneous gullies? Analysis of the interior (floor and walls) of the host crater suggest that prior to the Gasa impact, the pole-facing walls and floor were occupied by remnant debris-covered glaciers formed earlier in the Amazonian, which are relatively common in crater interiors in this latitude band. We suggest that the Gasa impact cratering event penetrated into the southern portion of this debris-covered glacier, emplaced ejecta on top of the debris layer covering the ice, and caused extensive melting of the buried ice and flow of water and debris slurries on the host crater floor. Inside Gasa, the impact crater rim crest and wall intersected the debris-covered glacier deposits around the northern, pole-facing part of the Gasa interior. We interpret this exposure of ice-rich debris-covered glacial material in the crater wall to be the source of meltwater that formed the very well-developed gullies along the northern, pole-facing slopes of Gasa crater.     

Experimental simulation of martian gully forms

Gullies are widespread on slopes on the surface of Mars and have been investigated by numerous authors, yet their formation processes remain elusive. In an attempt to understand the possibility of a water-based origin for these forms, we undertook a series of flume experiments at Earth surface temperatures and pressures. Our objectives were to produce forms that resemble those most commonly observed on Mars, documenting their morphometric characteristics and identifying any statistically significant relationships between form and controlling factors of slope and flow rate. Experiments were conducted in a 1×1.5 m² flume filled with medium grain size sand. The experiments were run over a slope angle range of 10–30°, corresponding to the range for gullies on Mars. Water from a constant-head tank fed through 5 mm silicone hose to a rotameter and then released just below the surface at the top of the slope. Gullies were produced at slope angle values of 10°, 20°, and 30° and flow rate values of 445, 705, 965, and 1260 mL min^?1 at each angle. Eighteen parameters were identified and subsequently measured on each gully produced in the flume. Gully forms were successfully reproduced and displayed development of the fundamental morphological components observed on Mars: alcove, channel, and apron. Slope–gully form relationships for each component revealed the following results: higher slope angles formed shorter gullies with thicker apron deposits. Moreover, longer gullies were seen at higher flow rates. We concluded that forms visually similar to those observed on Mars can be created by water in the laboratory flume under terrestrial conditions. Morphometric parameters can be measured and permit identification of controlling factors. Experimental simulation of gullies appears possible with proper scaling of experimental parameters. Although not directly scalable to Mars, flume gully parameters may be used to develop numerical models in the future.     

Periglacial mass-wasting landforms on Mars suggestive of transient liquid water in the recent past: Insights from solifluction lobes on Svalbard

On Earth, periglacial solifluction is a slow mass-wasting process related to freeze–thaw activity. We compare the morphology of small-scale lobate features on Mars to solifluction lobes in Svalbard to constrain their processes of formation. The analysis is based on high-resolution satellite imagery of Mars (HiRISE, ～25 cm/pxl), aerial images of Svalbard with a similar spatial resolution (HRSC–AX, ～20 cm/pxl) acquired through an air campaign in summer 2008, and ground truth obtained during two summer expeditions in 2009 and 2011 on Svalbard. We present a detailed study of two crater environments on Mars displaying two types of lobate forms, characterized as sorted (clast-banked) and non-sorted lobes. On both Svalbard and Mars such lobes typically occur as clusters of overlapping risers (lobe fronts), pointing to differential velocities in the soil. The martian small-scale lobes have well-defined arcuate risers and lobe treads (surface). Lobe widths range between 14 and 127 m and tread lengths between 13 and 105 m. Riser height is estimated to be approximately 1–5 m. The lobes on Mars share the plan view morphology of solifluction lobes on Svalbard and their morphometry is within the range of values of terrestrial solifluction lobes. The lobes are distinct from permafrost-creep landforms such as rock glaciers. We show the results of a survey of 53 HiRISE images covering latitudes between 59°N and 81°N. Similar to Svalbard, the studied lobate features on Mars occur in close spatial proximity to gullies and thermal contraction polygons. The widespread distribution of the lobate forms in the northern hemisphere and their close association to ground-ice and gullies are best explained by mass-wasting processes related to frost creep, gelifluction and/or plug-like flow. This suggests a protracted process (thousand to several thousands of years) of freeze–thaw activity at the northern high latitudes on Mars. Age constraints on lobe deposits and superposition relationships with gullies and polygons imply a process involving liquid water within the last few million years.     

Rheologies and ages of lava flows on Elysium Mons,Mars

We present results of our study of the rheologies and ages of lava flows in the Elysium Mons region of Mars. Previous studies have shown that the geometric dimensions of lava flows reflect rheological properties such as yield strength, effusion rate and viscosity. In this study the rheological properties of lava flows in the Elysium Mons region were determined and compared to the rheologies of the Ascraeus Mons lava flows. We also derived new crater size-frequency distribution measurements (CSFDs) for the Elysium lava flows to identify possible changes in the rheological properties with time. In addition, possible changes in the rheological properties with the distance from the caldera of Elysium Mons were analyzed.In total, 35 lava flows on and around Elysium Mons were mapped, and divided into three groups, lava flows on the flanks of Elysium Mons, in the plains between the three volcanoes Elysium Mons, Hecates and Albor Tholus and lava flows south of Albor Tholus. The rheological properties of 32 of these flows could be determined. Based on our morphometric measurements of each individual lava flow, estimates for the yield strengths, effusion rates, viscosities, and eruption duration of the studied lava flows were made. The yield strengths of the investigated lava flows range from ～3.8 × 10² Pa to ～1.5 × 10⁴ Pa, with an average of ～3.0 × 10³ Pa. These yield strengths are in good agreement with estimates for terrestrial basaltic lava flows. The effusion rates are on average ～747 m³ s^?1, ranging from ～99 to 4450 m³ s^?1. The viscosities are on average ～4.1 × 10⁶ Pa s, with a range of 1.2 × 10⁵ Pa s to 3.1 × 10⁷ Pa s. The eruption durations of the flows were calculated to be between 6 and 183 days, with an average of ～51 days. The determined rheological properties are generally very similar to those of other volcanic regions on Mars, such as on Ascraeus Mons in the Tharsis region. Calculated yield strengths and viscosities point to a basaltic/andesitic composition of the lava flows, similar to basaltic or andesitic a’a lava flows on Earth.Absolute model ages of all 35 lava flows on Elysium Mons were derived from crater size-frequency distribution measurements (CSFD). The derived model ages show a wide variation from about 632 Ma to 3460 Ma. Crater size-frequency distribution measurements of the Elysium Mons caldera show an age of ～1640 Ma, which is consistent with the resurfacing age of Werner (2009). Significant changes of the rheologies with time could not be observed. Similarly, we did not observe systematic changes in ages with increasing distances of lava flows from the Elysium Mons caldera.     

Gullies and their relationships to the dust–ice mantle in the northwestern Argyre Basin,Mars

We investigated gullies and their relationships to the atmospherically derived dust–ice mantle and aeolian features in the northwestern part of the Argyre basin. A detailed morphologic map of the Argyre study region allowed us to constrain the stratigraphic relationships and relative ages of gullies. In addition, we investigated the morphologic characteristics and orientations of all gullies in the Argyre study region. Maximum absolute ages for gullies were determined with crater size–frequency distribution measurements of the dust–ice mantle, which is the source material of gullies in the study area. Gullies only evolve from this mantle probably by melting of its ice content. Two different morphologies of pristine and degraded gullies were identified, mostly occurring on pole- and equatorward-facing slopes, respectively. We conclude that the morphologies and orientations were initiated either by a more rapid and extensive erosion of equatorward-facing gullies or by at least two generations of gullies with generally older gullies on equatorward-facing slopes and younger ones on pole-facing slopes. Different intensities of solar insolation on equator- and pole-facing slopes might be responsible for the different development of pristine and degraded gullies. Gullies in the study area generally have ages ?20 Ma. Some uncratered (and thus very young) aeolian dunes are superposed by a few gullies in some locations, indicating another even younger generation of gullies with an upper limit absolute model age of about <500 ka.     

Trapped Ar isotopes in meteorite ALH 84001 indicate Mars did not have a thick ancient atmosphere

Water is not currently stable in liquid form on the martian surface due to the present mean atmospheric pressure of ～7 mbar and mean global temperature of ～220 K. However, geomorphic features and hydrated mineral assemblages suggest that Mars’ climate was once warmer and liquid water flowed on the surface. These observations may indicate a substantially more massive atmosphere in the past, but there have been few observational constraints on paleoatmospheric pressures. Here we show how the ⁴⁰Ar/³⁶Ar ratios of trapped gases within martian meteorite ALH 84001 constrain paleoatmospheric pressure on Mars during the Noachian era [～4.56–3.8 billion years (Ga)]. Our model indicates that atmospheric pressures did not exceed ～1.5 bar during the first 400 million years (Ma) of the Noachian era, and were <400 mbar by 4.16 Ga. Such pressures of CO₂ are only sufficient to stabilize liquid water on Mars’ surface at low latitudes during seasonally warm periods. Other greenhouse gases like SO₂ and water vapor may have played an important role in intermittently stabilizing liquid water at higher latitudes following major volcanic eruptions or impact events.     

Is the primordial crust of Mars magnetized?

A meteorite impact capable of creating a 200 km diameter crater can demagnetize the entire crust beneath, and produce an appreciable magnetic anomaly at satellite altitudes of ～400 km in case the pre-existing crust is magnetized. In this study we examine the magnetic field over all of the craters and impact-related Quasi-Circular Depressions (QCDs) with diameters larger than 200 km that are located on the highlands of Mars, excluding the Tharsis bulge, in order to estimate the mean magnetization of the highland crust. Using the surface topography and the gravity of Mars we first identify those QCDs that are likely produced by impacts. The magnetic map of a given crater or impact-related QCD is derived using the Mars Global Surveyor high-altitude nighttime radial magnetic data. Two extended ancient areas are identified on the highlands, the South Province and the Tempe Terra, which have large number of craters and impact-related QCDs but none of them has an appreciable magnetic signature. The primordial crust of these areas is not magnetized, or is very weakly magnetized at most. We examine some plausible scenarios to explain the weak magnetization of these areas, and conclude that no strong dynamo existed in the first ～100 Myr of Mars’ history when the newly formed primordial crust was cooling below the magnetic blocking temperatures of its minerals.     

Titan’s global crater population: A new assessment

We report a revised crater population for Titan using Cassini RADAR data through January 2010 (flyby T65), and make a size-dependent correction for the incomplete coverage (～33%) using a Monte-Carlo model. Qualitatively, Titan’s landscape is more heavily cratered than Earth, but much less than Mars or Ganymede: the area fraction covered by craters is in fact comparable with that of Venus. Quantitative efforts to interpret crater densities for Titan as surface age have been confounded by widely divergent crater production rates proposed in the literature. We elucidate the specific model assumptions that lead to these differences (assumed projectile density, scaling function for simple crater diameter, and complex crater size exponent) and suggest these are reasonable bounding models, with the Korycansky and Zahnle (2005) model representing a crater retention age of ～1 Ga, and the Artemieva and Lunine (2005) model representing a crater retention age of ～200 Ma. These estimates are consistent with models of Titan’s evolution that predict a thickening of its crust 0.3–1.2 Gyr ago.     

Search for methane and upper limits to ethane and SO2 on Mars

The IRTF/CSHELL observations in February 2006 at L_S = 10° and 63–93°W show ～10 ppb of methane at 45°S to 7°N and ～3 ppb outside this region that covers the deepest canyon Valles Marineris. Observations in December 2009 at L_S = 20° and 0–30°W included spectra of the Moon at a similar airmass as a telluric calibrator. A technique for extraction of the martian methane line from a combination of the Mars and Moon spectra has been developed. The observations reveal no methane with an upper limit of 8 ppb. The results of both sessions agree with the observations by Mumma et al. (Mumma, M.J. et al. [2009]. Science 323, 1041–1045) at the same season in February 2006 and are smaller than those in the PFS and TES maps. Production and removal of the biological methane on Mars do not significantly change the redox state of the atmosphere and the balance of hydrogen. A search for ethane at 2977 cm^?1 results in an upper limit of 0.2 ppb. However, this limit does not help to establish the origin of methane on Mars. Reanalysis of our search for SO₂ using TEXES confirms the recently established upper limit of 0.3 ppb. Along with the lack of hot spots and gas vents with endogenic heat sources in the THEMIS observations, the very low upper limit to SO₂ on Mars does not favor geological methane that is less abundant than SO₂ in the outgassing from the terrestrial planets.     

Mars mesospheric zonal wind around northern spring equinox from infrared heterodyne observations of CO2

We present direct observations of Mars zonal wind velocities around northern spring equinox (L_S = 336°, L_S = 355°, L_S = 42°) during martian year 27 and 29. Data was acquired by means of infrared heterodyne spectroscopy of CO₂ features at 959.3917 cm^?1 (10.4232 μm) and 957.8005 cm^?1 (10.4405 μm) using the Cologne Tuneable Heterodyne Infrared Spectrometer (THIS) at the McMath–Pierce telescope of the National Solar Observatory on Kitt Peak in Arizona and the NASA Infrared Telescope Facility on Mauna Kea, Hawaii between 2005 and 2008. Winds were measured on the dayside of Mars with an unprecedented spatial resolution allowing sampling of up to nine independent latitudes over the martian disk. Retrieved wind velocities depend strongly on latitude and season with values ranging from 180 m/s prograde to ?94 m/s retrograde. A comparison of the observational results to predicted values from the Mars Climate Database yield a reasonable agreement between modeling and observation.     

Spatial and temporal variations in Titan's surface temperatures from Cassini CIRS observations

We report a wide-ranging study of Titan's surface temperatures by analysis of the Moon's outgoing radiance through a spectral window in the thermal infrared at 19 μm (530 cm^?1) characterized by lower atmospheric opacity. We begin by modeling Cassini Composite Infrared Spectrometer (CIRS) far infrared spectra collected in the period 2004–2010, using a radiative transfer forward model combined with a non-linear optimal estimation inversion method. At low-latitudes, we agree with the HASI near-surface temperature of about 94 K at 10°S (Fulchignoni et al., 2005). We find a systematic decrease from the equator toward the poles, hemispherically asymmetric, of ～1 K at 60° south and ～3 K at 60° north, in general agreement with a previous analysis of CIRS data (Jennings et al., 2009), and with Voyager results from the previous northern winter. Subdividing the available database, corresponding to about one Titan season, into 3 consecutive periods, small seasonal changes of up to 2 K at 60°N became noticeable in the results. In addition, clear evidence of diurnal variations of the surface temperatures near the equator are observed for the first time: we find a trend of slowly increasing temperature from the morning to the early afternoon and a faster decrease during the night. The diurnal change is ～1.5 K, in agreement with model predictions for a surface with a thermal inertia between 300 and 600 J m^?2 s^?1/2 K^?1. These results provide important constraints on coupled surface–atmosphere models of Titan's meteorology and atmospheric dynamic.     

Methane emissions from Earth’s degassing: Implications for Mars

The presence of methane on Mars is of great interest, since one possibility for its origin is that it derives from living microbes. However, CH₄ in the martian atmosphere also could be attributable to geologic emissions released through pathways similar to those occurring on Earth. Using recent data on methane degassing of the Earth, we have estimated the relative terrestrial contributions of fossil geologic methane vs. modern methane from living methanogens, and have examined the significance that various geologic sources might have for Mars.Geologic degassing includes microbial methane (produced by ancient methanogens), thermogenic methane (from maturation of sedimentary organic matter), and subordinately geothermal and volcanic methane (mainly produced abiogenically). Our analysis suggests that ～80% of the “natural” emission to the terrestrial atmosphere originates from modern microbial activity and ～20% originates from geologic degassing, for a total CH₄ emission of ～28.0×10⁷ tonnes year^?1.Estimates of methane emission on Mars range from 12.6×10¹ to 57.0×10⁴ tonnes year^?1 and are 3–6 orders of magnitude lower than that estimated for Earth. Nevertheless, the recently detected martian, Northern-Summer-2003 CH₄ plume could be compared with methane expulsion from large mud volcanoes or from the integrated emission of a few hundred gas seeps, such as many of those located in Europe, USA, Mid-East or Asia. Methane could also be released by diffuse microseepage from martian soil, even if macro-seeps or mud volcanoes were lacking or inactive. We calculated that a weak microseepage spread over a few tens of km², as frequently occurs on Earth, may be sufficient to generate the lower estimate of methane emission in the martian atmosphere.At least 65% of Earth’s degassing is provided by kerogen thermogenesis. A similar process may exist on Mars, where kerogen might include abiogenic organics (delivered by meteorites and comets) and remnants of possible, past martian life. The remainder of terrestrial degassed methane is attributed to fossil microbial gas (～25%) and geothermal-volcanic emissions (～10%). Global abiogenic emissions from serpentinization are negligible on Earth, but, on Mars, individual seeps from serpentinization could be significant. Gas discharge from clathrate-permafrost destabilization should also be considered.Finally, we have shown examples of potential degassing pathways on Mars, including mud volcano-like structures, fault and fracture systems, and major volcanic edifices. All these types of structures could provide avenues for extensive gas expulsion, as on Earth. Future investigations of martian methane should be focused on such potential pathways.     

A one-year seasonal analysis of martian gravity waves using MCS data

Temperature measurements from the Mars Climate Sounder on NASA’s Mars Reconnaissance Orbiter are examined for gravity wave signals using the Stockwell transform, a technique previously applied to terrestrial temperature profiles. An analysis is presented for internal-gravity waves throughout the martian atmosphere for the period July 2007–May 2009, representing a full martian year of data, divided by season. Momentum fluxes observed in the altitude range 200–20 Pa are measured as ～10^?5–10^?3 Pa, significantly higher than at equivalent altitudes on Earth, and are primarily peaked around the tropics. Observed orographic effects are minimal; the primary observed effects are seasonal. Waves are shown to vary significantly in wavelength with latitude, but are generally approximately zonally symmetric. Horizontal wavelengths are consistently much shorter in the northern hemisphere than the southern.     

Erosive flood events on the surface of Mars: Application to Mangala and Athabasca Valles

We address key factors involved in determining water flow conditions in outflow channels on Mars, including the temperature of the sub-surface water being released and the environmental conditions of low temperature, low atmospheric pressure, and low acceleration due to gravity. We suggest how some of the assumptions made in previous work may be improved. Our model considers the thermodynamic effects of simultaneous evaporation and freezing of water, and fluid dynamical processes including changes in flow rheology caused by assimilation of cold rock and ice eroded at the channel bed, and ice crystal growth due to water freezing. We model how far initially turbulent water could flow in a channel before it erodes and entrains enough material to become laminar, and subsequently ceases to erode the bed. An ice raft will begin to form on the flood while transition occurs between turbulent and laminar flow. Estimates are given for water transit times, ～17–19 h, initial water depths, 50–62 m, and average flow speeds, 5–12 m s^?1, in the Mangala and Athabasca Valles. We show that these two outflow channels, and by implication others like them, could plausibly have been formed in single water release events. Resulting mean erosion rates are approximately 0.7 mm s^?1, a factor of three greater than previous estimates based on combinations of estimates of flood duration and required water volumes. This is explained by the consideration of the effects of eroded ice and the physics of thermal erosion in the present study.     

The Circum-Hellas Volcanic Province,Mars: Overview

Building on previous studies of volcanoes around the Hellas basin with new studies of imaging (High-Resolution Stereo Camera (HRSC), Thermal Emission Imaging System (THEMIS), Mars Orbiter Camera (MOC), High-Resolution Imaging Science Experiment (HiRISE), Context Imager (CTX)), multispectral (HRSC, Observatoire pour la Minéralogie, l’Eau, les Glaces et l’Activité (OMEGA)), topographic (Mars Orbiter Laser Altimeter (MOLA)) and gravity data, we define a new Martian volcanic province as the Circum-Hellas Volcanic Province (CHVP). With an area of >2.1 million km², it contains the six oldest central vent volcanoes on Mars, which formed after the Hellas impact basin, between 4.0 and 3.6 Ga. These volcanoes mark a transition from the flood volcanism that formed Malea Planum ～3.8 Ga, to localized edifice-building eruptions. The CHVP volcanoes have two general morphologies: (1) shield-like edifices (Tyrrhena, Hadriaca, and Amphitrites Paterae), and (2) caldera-like depressions surrounded by ridged plains (Peneus, Malea, and Pityusa Paterae). Positive gravity anomalies are found at Tyrrhena, Hadriaca, and Amphitrites, perhaps indicative of dense magma bodies below the surface. The lack of positive-relief edifices and weak gravity anomalies at Peneus, Malea, and Pityusa suggest a fundamental difference in their formation, styles of eruption, and/or compositions. The northernmost volcanoes, the ～3.7–3.9 Ga Tyrrhena and Hadriaca Paterae, have low slopes, well-channeled flanks, and smooth caldera floors (at tens of meters/pixel scale), indicative of volcanoes formed from poorly consolidated pyroclastic deposits that have been modified by fluvial and aeolian erosion and deposition. The ～3.6 Ga Amphitrites Patera also has a well-channeled flank, but it and the ～3.8 Ga Peneus Patera are dominated by scalloped and pitted terrain, pedestal and ejecta flow craters, and a general ‘softened’ appearance. This morphology is indicative not only of surface materials subjected to periglacial processes involving water ice, but also of a surface composed of easily eroded materials such as ash and dust. The southernmost volcanoes, the ～3.8 Ga Malea and Pityusa Paterae, have no channeled flanks, no scalloped and pitted terrain, and lack the ‘softened’ appearance of their surfaces, but they do contain pedestal and ejecta flow craters and large, smooth, bright plateaus in their central depressions. This morphology is indicative of a surface with not only a high water ice content, but also a more consolidated material that is less susceptible to degradation (relative to the other four volcanoes). We suggest that Malea and Pityusa (and possibly Peneus) Paterae are Martian equivalents to Earth's giant calderas (e.g., Yellowstone, Long Valley) that erupted large volumes of volcanic materials, and that Malea and Pityusa are probably composed of either lava flows or ignimbrites. HRSC and OMEGA spectral data indicate that dark gray to slightly red materials (often represented as blue or black pixels in HRSC color images), found in the patera floors and topographic lows throughout the CHVP, have a basaltic composition. A key issue is whether this dark material represents concentrations of underlying basaltic material eroded by various processes and exposed by aeolian winnowing, or if the material was transported from elsewhere on Mars by regional winds. Understanding the provenance of these dark materials may be the key to understanding the volcanic diversity of the Circum-Hellas Volcanic Province.     

Stochastic variability of luminous blue variables

Using the archives of the American Association of Variable Stars Observers and our own data, we analyse the long-term variability of several well-studied luminous blue variables (LBVs) aiming on a general picture of stochastic variability of these objects. The power density spectra of all the selected objects may be generally described by a single power law contaminated by observational noise at higher frequencies. The slopes of the power-law component are close to p = 2 (where PDS ∝ f^?p, and f is frequency) for strongly variable flaring objects like AG Car and significantly smaller (p ～ 1.3) for P Cyg where brightness variation amplitude is ? 1^m and dominated by slow low-amplitude variability. The slope holds for about two orders of magnitude in the frequency domain, though peaks and curvatures are present at f ? 10^?2–10^?3 d^?1. We show that pseudo-photosphere approach to variability may explain the power-law shape of the variability spectrum at higher frequencies. However, the observed spectra are actually rather “red” than “brown”: flux variations are correlated up to tens of years that is much longer than the characteristic refreshment time scales of the pseudo-photosphere. We propose that several stochastic noise components produce the power spectra of LBVs.     

Characterizing Io’s Pele,Tvashtar and Pillan plumes: Lessons learned from Hubble

Hubble Space Telescope/Wide Field and Planetary Camera 2 (HST/WFPC2) images of Io obtained between 1995 and 2007 between 0.24 and 0.42 μm led to the detection of the Pele plume in reflected sunlight in 1995 and 1999; imaging of the Pele plume via absorption of jovian light in 1996 and 1999; detection of the Prometheus-type Pillan plume in reflected sunlight in 1997; and detection of the 2007 Pele-type Tvashtar plume eruption in reflected sunlight and via absorption of jovian light. Based on a detailed analysis of these observations we characterize and compare the gas and dust properties of each of the detected plumes. In each case, the brightness of the plumes in reflected sunlight is less at 0.26 μm than at 0.33 μm. Mie scattering analysis of the wavelength dependence of each plume’s reflectance signature suggests that range of particle sizes within the plumes is quite narrow. Assuming a normal distribution of particle sizes, the range of mean particle sizes is ～0.035–0.12 μm for the 1997 Pillan eruption, ～0.05–0.08 μm for the 1999 Pele and 2007 Tvasthar plumes, and ～0.05–0.11 μm for the 1995 Pele plume, and in each case the standard deviation in the particle size distribution is <15%. The Mie analysis also suggests that the 2007 Tvashtar eruption released ～10⁹ g of sulfur dust, the 1999 Pele eruption released ～10⁹ g of SO₂ dust, the 1997 Pillan eruption released ～10¹⁰ g of SO₂ dust, and the 1995 Pele plume may have released ～10¹⁰ g of SO₂ dust. Analysis of the plume absorption signatures recorded in the F255W filter bandpass (0.24–0.28 μm) indicates that the opacity of the 2007 Tvashtar plume was 2× that of the 1996 and 1999 Pele plume eruptions. While the sulfur dust density estimated for the Tvashtar from the reflected sunlight data could have produced 61% of the observed plume opacity, <10% of the 1999 Pele F255W plume opacity could have resulted from the SO₂ dust detected in the eruption. Accounting for the remaining F255W opacity level of the Pele and Tvasthar plumes based on SO₂ and S₂ gas absorption, the SO₂ and S₂ gas density inferred for each plume is almost equivalent corresponding to ～2–6 × 10¹⁶ cm^?2 and 3–5 × 10¹⁵ cm^?2, respectively, producing SO₂ and S₂ gas resurfacing rates ～0.04–0.2 cm yr^?1 and 0.007–0.01 cm yr^?1; and SO₂ and S₂ gas masses ～1–4 × 10¹⁰ g and ～2–3 × 10⁹ g; for a total dust to gas ratio in the plumes ～10^?1–10^?2. The 2007 Tvashtar plume was detected by HST at ～380 ± 40 km in both reflected sunlight and absorbed jovian light; in 1999, the detected Pele plume altitude was 500 km in absorbed jovian light, but in reflected sunlight the detected height was ～2× lower. Thus, for the 1999 Pele plume, similar to the 1979 Voyager Pele plume observations, the most efficient dust reflections occurred in the region closest to the plume vent. The 0.33–0.42 μm brightness of the 1997 Pillan plume was 10–20× greater than the Pele or Tvashtar plumes, exceeding by a factor of 3 the average brightness levels observed within 200 km of 1979 Loki eruption vent. But, the 0.26 μm brightness of the 1997 Pillan plume in reflected sunlight was significantly lower than would be predicted by the dust scattering model. Presuming that the 0.26 μm brightness of the 1997 Pillan plume was attenuated by the eruption plume’s gas component, then an SO₂ gas density ～3–6 × 10¹⁸ cm^?2 is inferred from the data (for S₂/SO₂ ratios ?4%), comparable to the 0.3–2 × 10¹⁸ cm^?2 SO₂ density detected at Loki in 1979 (Pearl, J.C. et al. [1979]. Nature 280, 755; Lellouch et al., 1992), and producing an SO₂ gas mass ～3–8 × 10¹¹ g and an SO₂ resurfacing rate ～8–23 cm yr^?1. These results confirm the connection between high (?10¹⁷ cm^?2) SO₂ gas content and plumes that scatter strongly at nearly blue wavelengths, and it validates the occurrence of high density SO₂ gas eruptions on Io. Noting that the SO₂ gas content inferred from a spectrum of the 2003 Pillan plume was significantly lower ～2 × 10¹⁶ cm^?2 (Jessup, K.L., Spencer, J., Yelle, R. [2007]. Icarus 192, 24–40); and that the Pillan caldera was flooded with fresh SO₂ frost/slush just prior to the 1997 Pillan plume eruption (Geissler, P., McEwen, A., Phillips, C., Keszthelyi, L., Spencer, J. [2004a]. Icarus 169, 29–64; Phillips, C.B. [2000]. Voyager and Galileo SSI Views of Volcanic Resurfacing on Io and the Search for Geologic Activity at Europa. Ph.D. Thesis, Univ. of Ariz., Tucson); we propose that the density of SO₂ gas released by this volcano is directly linked to the local SO₂ frost abundance at the time of eruption.     

Io’s atmosphere: Constraints on sublimation support from density variations on seasonal timescales using NASA IRTF/TEXES observations from 2001 to 2010

We present an analysis of 19 μm spectra of Io’s SO₂ atmosphere from the TEXES mid-infrared high spectral resolution spectrograph on NASA’s Infrared Telescope Facility, incorporating new data taken between January 2005 and June 2010 and a re-analysis of earlier data taken from November 2001 to January 2004. This is the longest set of contiguous observations of Io’s atmosphere using the same instrument and technique thus far. We have fitted all 16 detected blended absorption lines of the ν₂ SO₂ vibrational band to retrieve the subsolar values of SO₂ column abundance and the gas kinetic temperature. By incorporating an existing model of Io’s surface temperatures and atmosphere, we retrieve sub-solar column densities from the disk-integrated data. Spectra from all years are best fit by atmospheric temperatures <150 K. Best-fit gas kinetic temperatures on the anti-Jupiter hemisphere, where SO₂ gas abundance is highest, are low and stable, with a mean of 108 (±18) K. The sub-solar SO₂ column density between longitudes of 90–220° varies from a low of 0.61 (±0.145) × 10^?17 cm^?2, near aphelion in 2004, to a high of 1.51 (±0.215) × 10¹⁷ cm^?2 in 2010 when Jupiter was approaching its early 2011 perihelion. No correlation in the gas temperature was seen with the increasing SO₂ column densities outside the errors.Assuming that any volcanic component of the atmosphere is constant with time, the correlation of increasing SO₂ abundance with decreasing heliocentric distance provides good evidence that the atmosphere is at least partially supported by frost sublimation. The SO₂ frost thermal inertias and albedos that fit the variation in atmospheric density best are between 150–1250 W m^?2 s^?1/2 K^?1 and 0.613–0.425 respectively. Photometric evidence favors albedos near the upper end of this range, corresponding to thermal inertias near the lower end. This relatively low frost thermal inertia produces larger amplitude seasonal variations than are observed, which in turn implies a substantial additional volcanic atmospheric component to moderate the amplitude of the seasonal variations of the total atmosphere on the anti-Jupiter hemisphere. The seasonal thermal inertia we measure is unique both because it refers exclusively to the SO₂ frost surface component, and also because it refers to relatively deep subsurface layers (few meters) due to the timescales of many years, while previous studies have determined thermal inertias at shallower levels (few centimeters), relevant for timescales of ～2 h (eclipse) or ～2 days (diurnal curves).     

Sedimentary deposits in Xanthe Terra: Implications for the ancient climate on Mars

A variety of sedimentary deposits is observed in Xanthe Terra, Mars, including Gilbert-type deltas, fan deltas dominated by resedimentation processes, and alluvial fans. Sediments were provided through deeply incised valleys, which were probably incised by both runoff and groundwater sapping. Mass balances based on High-Resolution Stereo Camera (HRSC) digital terrain models show that up to ～30% of the material that was eroded in the valleys is present as deltas or alluvial fan deposits. Stratigraphic relationships and crater counts indicate an age of ～4.0 to ～3.8 Ga for the fluvial activity. Hydrologic modeling indicates that the deposits were probably formed in geologically very short time scales. Our results point to episodes of a warmer and wetter climate on early Mars, followed by a long period of significantly reduced erosion rates.     

The surface energy balance at the Huygens landing site and the moist surface conditions on Titan

The Huygens Probe provided a wealth of data concerning the atmosphere of Titan. It also provided tantalizing evidence of a small amount of surface liquid. We have developed a detailed surface energy balance for the Probe landing site. We find that the daily averaged non-radiative fluxes at the surface are 0.7 W m^?2, much larger than the global average value predicted by McKay et al. (1991) of 0.037 W m^?2. Considering the moist surface, the methane and ethane detected by the Probe from the surface is consistent with a ternary liquid of ethane, methane, and nitrogen present on the surface with mole fractions of methane, ethane, and nitrogen of 0.44, 0.34, and 0.22, respectively, and a total mass load of ～0.05 kg m^?2. If this liquid is included in the surface energy balance, only a small fraction of the non-radiative energy is due to latent heat release (～10^?3 W m^?2). If the amount of atmospheric ethane is less than 0.6×10^?5, the surface liquid is most likely evaporating over timescales of 5 Titan days, and the moist surface is probably a remnant of a recent precipitation event. If the surface liquid mass loading is increased to 0.5 kg m^?2, then the liquid lifetime increases to ～56 Titan days. Our modeling results indicate a dew cycle is unlikely, given that even when the diurnal variation of liquid is in equilibrium, the diurnal mass variation is only 3% of the total liquid. If we assume a high atmospheric mixing ratio of ethane (>0.6×10^?5), the precipitation of liquid is large (38 cm/Titan year for an ethane mixing ratio of 2×10^?5). Such a flux is many orders of magnitude in excess of the photochemical production rate of ethane.