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.     

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.     

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.     

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.     

Kinetics of methane clathrate formation and dissociation under Mars relevant conditions

Spectral observations have detected methane within the martian atmosphere (Formisano, V., Atreya, S., Encrenaz, T., Ignatiev, N., Giuranna, M. [2004]. Science 306, 1758–1761; Mumma, M.J. et al. [2009]. Science 323, 1041–1045), however, the origin of the methane has not been determined. Methane clathrate (also referred to as methane hydrate) has been suggested as a potential subsurface reservoir, storing and releasing biologic and/or abiogenic methane. In this study, rates of methane hydrate formation and dissociation were measured experimentally at 234–264 K and 1.4–4.7 MPa to test the clathrate reservoir hypothesis. Initial formation rates range from 4.3 × 10^?6 to 8.1 × 10^?5 mol m^?2 s^?1. Results show decreasing rates of formation over time in individual experiments, indicating initial rapid clathration, followed by diffusion-limited transport of methane into the ice through the previously formed hydrate. These experiments indicate increased pressure results in increased formation rates, likely the result of higher concentration gradients, enhancing the methane diffusion flux into the solid phase. Experiments conducted at elevated temperatures produced faster initial rates of formation, resulting from increased kinetic energy of methane molecules and/or thickening of the Quasi-Liquid Layer. Based on this temperature dependence, the activation energy for methane hydrate formation from ice was determined to be 35.9 kJ/mol. Hydrate dissociation experiments initiated by depressurization or warming at conditions between 222 K and 265 K and 0.1–2.0 MPa were conducted following each formation experiment, yielding methane hydrate dissociation rates from 3.01 × 10^?6 to 9.92 × 10^?5 mol m^?2 s^?1. While both hydrate dissociation and formation showed decreasing instantaneous rates over the course of each experiment, the transition between the initial rate of dissociation and the interpreted diffusion-limited period of continued dissociation was more abrupt than that observed in formation experiments, supporting an ice shielding effect. The initial concentration of methane in the solid phase had a significant effect on hydrate dissociation rates. Higher methane concentrations in the solid phase produce faster initial rates, likely due to increased concentration gradients, thus increasing the diffusion component of dissociation. Increased temperatures also produced faster dissociation rates, yielding an activation energy for dissociation of 32.7 kJ/mol. The rates determined within this study suggest that small near-surface methane hydrate reservoirs are a feasible source for recent methane plumes detected on Mars. Rates of methane release from gas hydrates also indicate that gas hydrate dissociation may have played a role in forming ancient chaos terrain and associated outflow channels.     

Dynamical evolution of the Hungaria asteroids

The Hungarias are a stable asteroid group orbiting between Mars and the main asteroid belt, with high inclinations (16–30°), low eccentricities (e < 0.18), and a narrow range of semi-major axes (1.78–2.06 AU). In order to explore the significance of thermally-induced Yarkovsky drift on the population, we conducted three orbital simulations of a 1000-particle grid in Hungaria a–e–i space. The three simulations included asteroid radii of 0.2, 1.0, and 5.0 km, respectively, with run times of 200 Myr. The results show that mean motion resonances—martian ones in particular—play a significant role in the destabilization of asteroids in the region. We conclude that either the initial Hungaria population was enormous, or, more likely, Hungarias must be replenished through collisional or dynamical means. To test the latter possibility, we conducted three more simulations of the same radii, this time in nearby Mars-crossing space. We find that certain Mars crossers can be trapped in martian resonances, and by a combination of chaotic diffusion and the Yarkovsky effect, can be stabilized by them. Therefore, some Hungarias (around 5% of non-family members with absolute magnitudes H < 15.5 and 10% for H < 17) may represent previously transient Mars crossers that have been adopted in this manner.     

The O2 nightglow in the martian atmosphere by SPICAM onboard of Mars-Express

We present observations of the O₂(a¹Δ_g) nightglow at 1.27 μm on Mars using the SPICAM IR spectrometer onboard of the Mars Express orbiter. In contrast to the O₂(a¹Δ_g) dayglow that results from the ozone photodissociation, the O₂(a¹Δ_g) nightglow is a product of the recombination of O atoms formed by CO₂ photolysis on the dayside at altitudes higher than 80 km and transported downward above the winter pole by the Hadley circulation. The first detections of the O₂(a¹Δ_g) nightglow in 2010 indicate that it is about two order of magnitude less intense than the dayglow (Bertaux, J.-L., Gondet, B., Bibring, J.-P., Montmessin, F., Lefèvre, F. [2010]. Bull. Am. Astron. Soc. 42, 1040; Clancy et al. [2010]. Bull. Am. Astron. Soc. 42, 1041). SPICAM IR sounds the martian atmosphere in the near-IR range (1–1.7 μm) with the spectral resolution of 3.5 cm^?1 in nadir, limb and solar occultation modes. In 2010 the vertical profiles of the O₂(a¹Δ_g) nightside emission have been obtained near the South Pole at latitudes of 82–83°S for two sequences of observations: L_s = 111–120° and L_s = 152–165°. The altitude of the emission maximum varied from 45 km on L_s = 111–120° to 38–49 km on L_s = 152–165°. Averaged vertically integrated intensity of the emission at these latitudes has shown an increase from 0.22 to 0.35 MR. Those values of total vertical emission rate are consistent with the OMEGA observations on Mars-Express in 2010. The estimated density of oxygen atoms at altitudes from 50 to 65 km varies from 1.5 × 10¹¹ to 2.5 × 10¹¹ cm^?3. Comparison with the LMD general circulation model with photochemistry (Lefèvre, F., Lebonnois, S., Montmessin, F., Forget, F. [2004]. J. Geophys. Res. 109, E07004; Lefèvre et al. [2008]. Nature 454, 971–975) shows that the model reproduces fairly well the O₂(a¹Δ_g) emission layer observed by SPICAM when the large field of view (>20 km on the limb) of the instrument is taken into account.     

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.     

Local inversion of magnetic anomalies: Implication for Mars’ crustal evolution

Martian magnetic anomalies have been revealed by the Mars Global Surveyor (MGS) mission in the south hemisphere of Mars. The present study models anomalies located in the ancient Terra Sirenum area between latitudes 26°S and 40°S and longitudes 185°E and 210°E using forward and inverse approaches. While the high-altitude measurements reveal the presence of two main magnetic anomalies, three are detected by low-altitude data. They are modeled as uncorrelated dipolar sources. Forward models predict large magnetizations between 30 and 60 A/m. A generalized non-linear inversion is used to determine the characteristics of the dipoles, based on different subsets of data. Low-altitude measurements inversion leads to more reliable results than those obtained by the inversion of high-altitude measurements only. Inversion of both low- and high-altitude data together provides with three dipoles that explain more than 57% of the signal, within this 10⁶ km² area. All dipoles have large magnetizations. Serpentinization of the early martian crust can explain such remanent magnetizations. Two resulting dipoles are 56 km deep, which suggests a locally thick martian crust. The last one is shallower (31 km). This indicates different origins and/or magnetization processes. Paleomagnetic poles are calculated and located around the Tharsis bulge. It suggests that Tharsis formed at high latitudes and moved toward its present location by polar reorientation.     

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.     

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.     

Characterization of hydrated silicate-bearing outcrops in Tyrrhena Terra,Mars: Implications to the alteration history of Mars

The Tyrrhena Terra region of Mars is studied with the imaging spectrometers OMEGA (Observatoire pour la Minéralogie, l’Eau, les Glaces et l’Activité) onboard Mars Express and CRISM (Compact Reconnaissance Infrared Spectrometer for Mars) onboard Mars Reconnaissance Orbiter, through the observation of tens of craters that impacted into this part of the martian highlands. The 175 detections of hydrated silicates are reported, mainly associated with ejecta blankets, crater walls and rims, and central up-lifts. Sizes of craters where hydrated silicates are detected are highly variable, diameters range from less than 1 km to 42 km. We report the presence of zeolites and phyllosilicates like prehnite, Mg-chlorite, Mg-rich smectites and mixed-layer chlorites–smectites and chlorite–vermiculite from comparison of hyperspectral infrared observations with laboratory spectra. These minerals are associated with fresh craters post-dating any aqueous activity. They likely represent ancient hydrated terrains excavated by the crater-forming impacts, and hence reveal the composition of the altered Noachian crust, although crater-related hydrothermal activity may have played a minor role for the largest craters (>20 km in diameter). Most detected minerals formed over relatively high temperatures (100–300 °C), likely due to aqueous alteration of the Noachian crust by regional low grade metamorphism from the Noachian thermal gradient and/or by extended hydrothermal systems associated with Noachian volcanism and ancient large impact craters. This is in contrast with some other phyllosilicate-bearing regions like Mawrth Vallis where smectites, kaolinites and hydrated silica were mainly identified, pointing to a predominance of surface/shallow sub-surface alteration; and where excavation by impacts played only a minor role. Smooth plains containing hydrated silicates are observed at the boundary between the Noachian altered crust, dissected by fluvial valleys, and the Hesperian unaltered volcanic plains. These plains may correspond to alluvial deposition of eroded material. The highlands of Tyrrhena Terra are therefore particularly well suited for investigating the diversity of hydrated minerals in ancient martian terrains.     

Observations of water vapour and carbon monoxide in the Martian atmosphere with the SWC of PFS/MEX

In the history of Mars exploration its atmosphere and planetary climatology aroused particular interest. In the study of the minor gases abundance in the Martian atmosphere, water vapour became especially important, both because it is the most variable trace gas, and because it is involved in several processes characterizing the planetary atmosphere. The water vapour photolysis regulates the Martian atmosphere photochemistry, and so it is strictly related to carbon monoxide. The CO study is very important for the so-called “atmosphere stability problem”, solved by the theoretical modelling involving photochemical reactions in which the H₂O and the CO gases are main characters.The Planetary Fourier Spectrometer (PFS) on board the ESA Mars Express (MEX) mission can probe the Mars atmosphere in the infrared spectral range between 200 and 2000 cm^?1 (5–50 μm) with the Long Wavelength Channel (LWC) and between 1700 and 8000 cm^?1 (1.2–5.8 μm) with the Short Wavelength Channel (SWC). Although there are several H₂O and CO absorption bands in the spectral range covered by PFS, we used the 3845 cm^?1 (2.6 μm) and the 4235 cm^?1 (2.36 μm) bands for the analysis of water vapour and carbon monoxide, respectively, because these ranges are less affected by instrumental problems than the other ones. The gaseous concentrations are retrieved by using an algorithm developed for this purpose.The PFS/SW dataset used in this work covers more than two and a half Martian years from Ls=62° of MY 27 (orbit 634) to Ls=203° of MY 29 (orbit 6537). We measured a mean column density of water vapour of about 9.6 pr. μm and a mean mixing ratio of carbon monoxide of about 990 ppm, but with strong seasonal variations at high latitudes. The seasonal water vapour map reproduces very well the known seasonal water cycle. In the northern summer, water vapour and CO show a good anticorrelation most of the time. This behaviour is due to the carbon dioxide and water sublimation from the north polar ice cap, which dilutes non-condensable species including carbon monoxide. An analogous process takes place during the winter polar cap, but in this case the condensation of carbon dioxide and water vapour causes an increase of the concentration of non-condensable species. Finally, the results show the seasonal variation of the carbon monoxide mixing ratio with the surface pressure.     

Line broadening in the neutral and ionized mercury spectra

The neutral, singly, doubly and triply ionized mercury (Hg I–IV, respectively) spectral line shapes and line center positions have been investigated in the laboratory helium plasma at electron densities ranging between 9.3 × 10²² m^?3 and 1.93 × 10²³ m^?3 and electron temperatures around 19,500 K, both interesting for astrophysics. The mercury (natural isotope composition) atoms were sputtered from the cylindrical amalgamated gold plates located in the homogenous part of the pulsed helium discharge operating at a pressure of 665 Pa in a flowing regime. The mercury spectral line profiles were recorded using the McPherson model 209 spectrograph and the Andor ICCD camera as the detection system. This research presents Stark broadening parameters, the width (W) and the shift (d), of one Hg I, 19 Hg II, 6 Hg III and 4 Hg IV lines, not investigated so far. Our experimental W values were compared with the data calculated applying various approaches. The shape and intensity of astrophysically important 398.4 nm Hg II spectral line was discussed taking into account the isotope shift, hyperfine structure and Penning effects. At the mentioned plasma parameters the Stark broadening is found to be a main line broadening mechanism of the lines (λ > 200 nm) in the Hg I–IV spectra.     

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.     

Regional differences in gully occurrence on Mars: A comparison between the Hale and Bond craters

The observation of gullies on Mars raised questions about the presence of liquid water in the recent past. In some regions like Hale and Bond crater, gullies occur in one crater (Hale) but do not in another crater nearby (Bond). These regional differences have been interpreted as an argument for a formation of the gullies related to groundwater. The formation of gullies on Earth depends on rainfall and/or melting of snow as well as on several parameters such as the presence of steep slopes and sufficient amounts of fines and debris. We investigated the Hale/Bond region for differences in crater wall morphology and texture, slopes, and thermal properties to determine whether the gully formation is dependent on factors such as steep slope angles and availability of fine-grained material. Morphologically there exist two kinds of gullies in the Hale crater: Gullies on the south- and east-facing crater slopes have a pristine appearance with deep channels eroded into the talus material and well-preserved aprons. Gully-like features on the north- and west-facing slopes are degraded and superposed by craters, indicating that they are old in comparison to the pristine ones. However, their formation process is unclear and might be due to debris flows, surface runoff or dry mass wasting processes or a combination of these processes. The crater walls of Bond do not show gullies. Their morphology is most likely consistent with a degraded mantle deposit. Slope measurements reveal that the gullies in Hale crater occur on slopes between ～20° and ～30° in contrast to the slopes without gullies in Bond that are between ～10° and ～20° steep. Mean thermal inertia values on slopes with younger gullies are ～175 J m^?2 K^?1 s^?1/2 corresponding to higher amounts of fine-grained material. At slopes with older gully-like features mean thermal inertia values are ～315 J m^?2 K^?1 s^?1/2 corresponding to higher amounts of bedrock or possibly indurated grain sizes. Mean thermal inertia values of the Bond crater walls are ～230 J m^?2 K^?1 s^?1/2 indicating more consolidated terrain possibly due to the cementation of the dissected mantle material. From our investigation we conclude that the occurrence of gullies in the Hale/Bond region most likely depends on the distribution of unconsolidated material and steep slopes. The regional and local gully distribution on Mars likely varies due to differences in topography and surface material properties. Their proposed clustered distribution on Mars is not an argument for a groundwater formation mechanism of the gullies.     

The effect of high temperatures on the mid-to-far-infrared emission and near-infrared reflectance spectra of phyllosilicates and natural zeolites: Implications for martian exploration

Most phyllosilicates on Mars appear to be associated with ancient terrains. As such, they may have experienced shock heating produced by impacts and could have been significantly altered or melted. We characterized the effects of high temperatures on the mid-to-far-infrared (mid-to-far-IR) emission (100–1400 cm^?1; 7.1–100 μm) and near-infrared (NIR) reflectance (1.2–2.5 μm) spectra of phyllosilicates by measuring experimentally calcined (100–900 °C) phyllosilicates and also two zeolites. Correlated differential scanning calorimetry (DSC) measurements were also performed on each sample to provide insight into the thermal activities of the phyllosilicates and natural zeolites. Our results indicate that all phyllosilicates exhibit characteristic degradations in both NIR and mid-to-far-IR spectral properties between 400 and 800 °C, mainly attributable to the dehydroxylation and recrystallization processes as temperature increases. Spectral features of natural zeolites persist to higher temperatures compared to features of phyllosilicates during heating treatments. The thermal behaviors of phyllosilicate infrared (IR) properties are greatly influenced by the compositions of the octahedral cations: (1) changes in both the NIR and mid-to-far-IR spectra of phyllosilicates tend to occur at lower temperatures (300–400 °C) in the Fe³⁺-rich samples as compared to the Al³⁺-rich types (400–600 °C); (2) Mg²⁺-trioctahedral phyllosilicates hectorite, saponite, and sepiolite all display major mid-to-far-IR spectral changes at 700 °C, corresponding to the formation of enstatite; (3) phyllosilicates that have minor replacement of Mg²⁺ for Al³⁺ in octahedral positions (e.g. cheto-type montmorillonite and palygorskite) show an absorption band at ～920 cm^?1 that becomes strong at 900 °C. Inconsistency between spectral behaviors in the mid-to-far-IR and NIR regions is also discussed for phyllosilicates. Results from this study have provided suggestive evidence for the scenario that some phyllosilicates could lose all original spectral features in mid-to-far-IR region while maintaining their characteristic hydration bands in NIR region in the same temperature range.     

Observation of DCl and upper limit to NH3 on Venus

To search for DCl in the Venus atmosphere, a spectrum near the D³⁵Cl (1–0) R4 line at 2141.54 cm^?1 was observed using the CSHELL spectrograph at NASA IRTF. Least square fitting to the spectrum by a synthetic spectrum results in a DCl mixing ratio of 17.8 ± 6.8 ppb. Comparing to the HCl abundance of 400 ± 30 ppb (Krasnopolsky [2010a] Icarus, 208, 314–322), the DCl/HCl ratio is equal to 280 ± 110 times the terrestrial D/H = 1.56 × 10^?4. This ratio is similar to that of HDO/H₂O = 240 ± 25 times the terrestrial HDO/H₂O from the VEX/SOIR occultations at 70–110 km. Photochemistry in the Venus mesosphere converts H from HCl to that in H₂O with a rate of 1.9 × 10⁹ cm^?2 s^?1 (Krasnopolsky [2012] Icarus, 218, 230–246). The conversion involves photolysis of HCl; therefore, the photochemistry tends to enrich D/H in HCl and deplete in H₂O. Formation of the sulfuric acid clouds may affect HDO/H₂O as well. The enriched HCl moves down by mixing to the lower atmosphere where thermodynamic equilibriums for H₂ and HCl near the surface correspond to D/H = 0.71 and 0.74 times that in H₂O, respectively. Time to establish these equilibriums is estimated at ～3 years and comparable to the mixing time in the lower atmosphere. Therefore, the enriched HCl from the mesosphere gives D back to H₂O near the surface. Comparison of chemical and mixing times favors a constant HDO/H₂O up to ～100 km and DCl/HCl equal to D/H in H₂O times 0.74.Ammonia is an abundant form of nitrogen in the reducing environments. Thermodynamic equilibriums with N₂ and NO near the surface of Venus give its mixing ratio of 10^?14 and 6 × 10^?7, respectively. A spectrum of Venus near the NH₃ line at 4481.11 cm^?1 was observed at NASA IRTF and resulted in a two-sigma upper limit of 6 ppb for NH₃ above the Venus clouds. This is an improvement of the previous upper limit by a factor of 5. If ammonia exists at the ppb level or less in the lower atmosphere, it quickly dissociates in the mesosphere and weakly affects its photochemistry.