Simultaneous mapping of H2O and H2O2 on Mars from infrared high-resolution imaging spectroscopy

New maps of martian water vapor and hydrogen peroxide have been obtained in November-December 2005, using the Texas Echelon Cross Echelle Spectrograph (TEXES) at the NASA Infra Red Telescope facility (IRTF) at Mauna Kea Observatory. The solar longitude Ls was 332° (end of southern summer). Data have been obtained at 1235-1243 cm⁻¹, with a spectral resolution of 0.016 cm⁻¹ (R=8×¹⁰⁴). The mean water vapor mixing ratio in the region [0°-55° S; 345°-45° W], at the evening limb, is 150±50 ppm (corresponding to a column density of 8.3±2.8 pr-μm). The mean water vapor abundance derived from our measurements is in global overall agreement with the TES and Mars Express results, as well as the GCM models, however its spatial distribution looks different from the GCM predictions, with evidence for an enhancement at low latitudes toward the evening side. The inferred mean H₂O₂ abundance is 15±10 ppb, which is significantly lower than the June 2003 result [Encrenaz, T., Bézard, B., Greathouse, T.K., Richter, M.J., Lacy, J.H., Atreya, S.K., Wong, A.S., Lebonnois, S., Lefèvre, F., Forget, F., 2004. Icarus 170, 424-429] and lower than expected from the photochemical models, taking in account the change in season. Its spatial distribution shows some similarities with the map predicted by the GCM but the discrepancy in the H₂O₂ abundance remains to be understood and modeled.     

Infrared imaging spectroscopy of Mars: H2O mapping and determination of CO2 isotopic ratios

High-resolution infrared imaging spectroscopy of Mars has been achieved at the NASA Infrared Telescope Facility (IRTF) on June 19-21, 2003, using the Texas Echelon Cross Echelle Spectrograph (TEXES). The areocentric longitude was 206°. Following the detection and mapping of hydrogen peroxide H₂O₂ [Encrenaz et al., 2004. Icarus 170, 424-429], we have derived, using the same data set, a map of the water vapor abundance. The results appear in good overall agreement with the TES results and with the predictions of the Global Circulation Model (GCM) developed at the Laboratory of Dynamical Meteorology (LMD), with a maximum abundance of water vapor of 3±1.5×¹⁰⁻⁴(17±9 pr-μm). We have searched for CH₄ over the martian disk, but were unable to detect it. Our upper limits are consistent with earlier reports on the methane abundance on Mars. Finally, we have obtained new measurements of CO₂ isotopic ratios in Mars. As compared to the terrestrial values, these values are: (¹⁸O/¹⁷O)[M/E] = 1.03 ± 0.09; (¹³C/¹²C)[M/E] = 1.00 ± 0.11. In conclusion, in contrast with the analysis of Krasnopolsky et al. [1996. Icarus 124, 553-568], we conclude that the derived martian isotopic ratios do not show evidence for a departure from their terrestrial values.     

PFS/MEX observations of the condensing CO2 south polar cap of Mars

The condensing CO₂ south polar cap of Mars and the mechanisms of the CO₂ ice accumulation have been studied through the analysis of spectra acquired by the Planetary Fourier Spectrometer (PFS) during the first two years of ESA's Mars Express (MEX) mission. This dataset spans more than half a martian year, from Ls∼330° to Ls∼194°, and includes the southern fall season which is found to be extremely important for the study of the residual south polar cap asymmetry. The cap expands symmetrically and with constant speed during the fall season. The maximum extension occurs sometime in the 80°-90° Ls range, when the cap edges are as low as −40° latitude. Inside Hellas and Argyre basins, frost can be stable at lower latitudes due to the higher pressure values, causing the seasonal cap to be asymmetric. Within the seasonal range considered in this paper, the cap edge recession rate is approximately half the rate at which the cap edge expanded. The longitudinal asymmetries reduce during the cap retreat, and disappear around Ls∼145°. Two different mechanisms are responsible for CO₂ ice accumulation during the fall season, especially in the 50°-70° Ls range. Here, CO₂ condensation in the atmosphere, and thus precipitation, is allowed exclusively in the western hemisphere, and particularly in the longitudinal corridor of the perennial cap. In the eastern hemisphere, the cap consists mainly of CO₂ frost deposits, as a consequence of direct vapor deposition. The differences in the nature of the surface ice deposits are the main cause for the residual south polar cap asymmetry. Results from selected PFS orbits have also been compared with the results provided by the martian general circulation model (GCM) of the Laboratoire de Météorologie dynamique (LMD) in Paris, with the aim of putting the observations in the context of the global circulation. This first attempt of cross-validation between PFS measurements and the LMD GCM on the one hand confirms the interpretation of the observations, and on the other hand shows that the climate modeling during the southern polar night on Mars is extremely sensitive to the dynamical forcing.     

A sensitive search for SO2 in the martian atmosphere: Implications for seepage and origin of methane

Mars was observed near the peak of the strongest SO₂ band at 1364-1373 cm⁻¹ with resolving power of 77,000 using the Texas Echelon Cross Echelle Spectrograph on the NASA Infrared Telescope Facility. The observation covered the Tharsis volcano region which may be preferable to search for SO₂. The spectrum shows absorption lines of three CO₂ isotopomers and three H₂O isotopomers. The water vapor abundance derived from the HDO lines assuming D/H = 5.5 times the terrestrial value is 12±1.0 pr. μm, in agreement with the simultaneous MGS/TES observations of 14 pr. μm at the latitudes (50° S to 10° N) of our observation. Summing of spectral intervals at the expected positions of sixteen SO₂ lines puts a 2σ upper limit on SO₂ of 1 ppb. SO₂ may be emitted into the martian atmosphere by seepage and is removed by three-body reactions with OH and O. The SO₂ lifetime, 2 years, is longer than the global mixing time 0.5 year, so SO₂ should be rather uniformly distributed across Mars. Seepage of SO₂ is less than 15,000 tons per year on Mars which is smaller than the volcanic production of SO₂ on the Earth by a factor of 700. Because CH₄/SO₂ is typically 10⁻⁴-10⁻³ in volcanic gases on the Earth, our results show seepage is unlikely to be the source of the recently discovered methane on Mars and therefore strengthen its biogenic origin.     

Photochemistry of the martian atmosphere: Seasonal, latitudinal, and diurnal variations

There is a significant progress in the observational data relevant to Mars photochemistry in the current decade. These data are not covered by and sometimes disagree with the published models. Therefore we consider three types of models for Mars photochemistry. A steady-state model for global-mean conditions is currently the only way to calculate the abundances of long living species (H₂, O₂, and CO). However, our model does not fit the observed CO abundance using gas-phase chemistry and reasonable values of heterogeneous loss of odd hydrogen on the water ice aerosol. The second type of the calculated models is steady-state models for local conditions. The MGS/TES data on temperature profiles, H₂O, and dust are input parameters for these models. The calculations have been made for nine seasonal points spread over the martian year and for twelve latitudes with a step of 10° for each season. The only adopted heterogeneous reaction is a weak loss of H₂O₂ on water ice with probability of 5×¹⁰⁻⁴. The results are in good agreement with the recent observations of the O₂ dayglow at 1.27 μm and the O₃ and H₂O₂ abundances. Global maps of the seasonal and latitudinal behavior of these species have been made. The third type of models is a time-dependent model for local conditions. These models show that odd hydrogen quickly converts to H₂O₂ at the nighttime and the chemistry is switched off while the association of O, the heterogeneous loss of H₂O₂, and eddy diffusion continue. This requires significant changes in the global-mean and local steady-state models discussed above, and these changes have been properly done. The calculated diurnal variations of Mars photochemistry are discussed. The martian photochemistry at low and middle latitudes is significantly different in the aphelion period at LS=10°-130° from that in the remaining part of the year.     

Interannual variability in TES atmospheric observations of Mars during 1999-2003

We use infrared spectra returned by the Mars Global Surveyor Thermal Emission Spectrometer (TES) to retrieve atmospheric and surface temperature, dust and water ice aerosol optical depth, and water vapor column abundance. The data presented here span more than two martian years (Mars Year 24, L_s=104°, 1 March 1999 to Mars Year 26, L_s=180°, 4 May 2003). We present an overview of the seasonal (L_s), latitudinal, and longitudinal dependence of atmospheric quantities during this period, as well as an initial assessment of the interannual variability in the current martian climate. We find that the perihelion season (L_s=180°-360°) is relatively warm, dusty, free of water ice clouds, and shows a relatively high degree of interannual variability in dust optical depth and atmospheric temperature. On the other hand, the aphelion season (L_s=0°-180°) is relatively cool, cloudy, free of dust, and shows a low degree of interannual variability. Water vapor abundance shows a moderate amount of interannual variability at all seasons, but the most in the perihelion season. Much of the small amount of interannual variability that is observed in the aphelion season appears to be caused by perihelion-season planet-encircling dust storms. These dust storms increase albedo through deposition of bright dust on the surface causing cooler daytime surface and atmospheric temperatures well after dust optical depth returns to prestorm values.     

Mars atmospheric water vapor abundance: 1991-1999, emphasis 1998-1999

Our ground-based measurements of martian atmospheric water vapor, made throughout Ls=34° to 249°, 24 September 1998 to 23 November 1999, during Mars year 24 (MY 24), show changes in Mars' humidity on hourly, daily, and seasonal timescales. We made concomitant measurement of nearby CO₂ bands, and when possible, results were corrected for aerosol extinction using aerosol optical depths derived from our own CO₂ analysis. Where there is spatial and temporal overlap, similar results are obtained for water vapor abundances and aerosol opacities as those observed from the Thermal Emission Spectrometer on Mars Global Surveyor. In addition some further discussion of our published earlier water vapor measurements (1991-1995) is included. Six results from this data set are: (1) the measured aerosol opacity in Mars atmosphere was variable but not greater than τ=1, with almost no clear atmosphere being observed, (2) measurements made with the slit crossing many hours of local time on Mars' Earth-facing disk show a diurnal pattern with highest abundances at mid-day and low abundance in very early morning and late afternoon for some but not all measurements, (3) water vapor abundance is patchy on hourly and daily time scales but follows the usual seasonal trends seen by instrumentation on the Mars Atmospheric Water Detector on the Viking Orbiters and by the Thermal Emission Spectrometer on Mars Global Surveyor, (4) there is a slight longitudinal correlation with the ground-ice observed by the Gamma Ray Spectrometer on Mars Odyssey, (5) there is evidence of the Low Southern Latitude Summer Minimum in our water vapor measurements but our data set for southern summer is limited, and (6) MY 24 appears to be wetter than MY 22 and MY 23.     

Mars atmospheric water vapor abundance: 1996-1997

Measurements of martian atmospheric water vapor made throughout L_s = 18.0°-146.4° (October 3, 1996-July 12, 1997) show changes in Mars humidity on hourly, daily, and seasonal time scales. Because our observing program during the 1996-1997 Mars apparition did not include concomitant measurement of nearby CO₂ bands, high northern latitude data were corrected for dust and aerosol extinction assuming an optical depth of 0.8, consistent with ground-based and HST imaging of northern dust storms. All other measurements with airmass greater than 3.5 were corrected using a total optical depth of 0.5. Three dominant results from this data set are as follows: (1) pre- and post-opposition measurements made with the slit crossing many hours of local time on Mars’ Earth-facing disk show a distinct diurnal pattern with highest abundances around and slightly after noon with low abundances in the late afternoon, (2) measurements of water vapor over the Mars Pathfinder landing site (Carl Sagan Memorial Station) on July 12, 1997, found 21 ppt μm in the spatial sector centered near 19° latitude, 36° longitude while abundances around the site varied from as low as 6 to as high as 28 ppt μm, and (3) water vapor abundance is patchy on hourly and daily time scales but follows the usual seasonal trends.     

Mapping of Mars O2 1.27 μm dayglow at four seasonal points

The O₂ dayglow at 1.27 μm is formed by high-altitude ozone on Mars and is a sensitive tracer of Mars photochemistry. Mapping of this dayglow using the IRTF/CSHELL long-slit spectrograph requires the extraction of weak emission lines against a strong continuum of the reflected solar light. Some new tools are suggested to improve the data processing. The observed O₂ dayglow intensities at L_S=67°, 112°, 148°, and 173° show a decrease from late spring (aphelion) to fall equinox by a factor of ≈5 at low latitudes (±30°). This decrease agrees with that predicted by a model of Clancy and Nair (1996, J. Geophys. Res. 101 (12) 12785-12790), although the dayglow intensities are weaker than those based on that model. The measured dayglow variations with latitude are rather low at L_S=67°, 112°, and 148° and unexpectedly high at 173°. The dayglow intensity peaks near noon and is smaller at 9:00 and 16:30 LT by a factor of 2. Some data on the ozone profile near aphelion are obtained from a combination of the dayglow and ozone observations. It is hardly possible to detect the O₂ night airglow at 1.27 μm on Mars using the existing ground-based and on-orbit instruments. The O₂ dayglow intensity as a function of latitude and season from aphelion to fall equinox has been obtained. Our goal is to extend this distribution to the full martian year and get a database for Mars photochemistry to complement the MGS/TES observations of water vapor, atmospheric temperature, and dust and ice aerosol.     

Solar infrared occultation observations by SPICAM experiment on Mars-Express: Simultaneous measurements of the vertical distributions of H2O, CO2 and aerosol

The infrared AOTF spectrometer is a part of the SPICAM experiment onboard the Mars-Express ESA mission. The instrument has a capability of solar occultations and operates in the spectral range of 1-1.7 μm with a spectral resolution of ∼3.5 cm⁻¹. We report results from 24 orbits obtained during MY28 at Ls 130°-160°, and the latitude range of 40°-55° N. For these orbits the atmospheric density from 1.43 μm CO₂ band, water vapor mixing ratio based on 1.38 μm absorption, and aerosol opacities were retrieved simultaneously. The vertical resolution of measurements is better than 3.5 km. Aerosol vertical extinction profiles were obtained at 10 wavelengths in the altitude range from 10 to 60 km. The interpretation using Mie scattering theory with adopted refraction indices of dust and H₂O ice allows to retrieve particle size (reff∼0.5-1 μm) and number density (∼1 cm⁻³ at 15-30 km) profiles. The haze top is generally below 40 km, except the longitude range of 320°-50° E, where high-altitude clouds at 50-60 km were detected. Optical properties of these clouds are compatible with ice particles (effective radius reff=0.1-0.3 μm, number density N∼10 cm−3) distributed with variance νeff=0.1-0.2 μm. The vertical optical depth of the clouds is below 0.001 at 1 μm. The atmospheric density profiles are retrieved from CO₂ band in the altitude range of 10-90 km, and H₂O mixing ratio is determined at 15-50 km. Unless a supersaturation of the water vapor occurs in the martian atmosphere, the H₂O mixing ratio indicates ∼5 K warmer atmosphere at 25-45 km than predicted by models.     

Spectral evidence for zeolite in the dust on Mars

Spectral features observed in Mars Global Surveyor Thermal Emission Spectrometer data (∼1670-220 cm⁻¹) of martian surface dust provide clues to its mineralogy. An emissivity peak at ∼1630 cm⁻¹ is consistent with the presence of an H₂O-bearing mineral. This spectral feature can be mapped globally and shows a distribution related to the classical bright regions on Mars that are known to be dust covered. An important spectral feature at ∼830 cm⁻¹ present in a newly derived average spectrum of surface dust likely is a transparency feature arising from the fine particulate nature of the dust. Its shape and location are consistent with plagioclase feldspars and also zeolites, which essentially are the hydrous form of feldspar. The generally favored visible/near-infrared spectral analog for martian dust, JSC Mars-1 altered tephra, does not display the ∼830 cm⁻¹ feature. Zeolites commonly form from the interaction of low temperature aqueous fluids and volcanic glass in a variety of geologic settings. The combination of spectral features that are consistent with zeolites and the likelihood that Mars has (or had) geologic conditions necessary to produce them makes a strong case for recognizing zeolite minerals as likely components of the martian regolith.     

Methane evolution from UV-irradiated spacecraft materials under simulated martian conditions: Implications for the Mars Science Laboratory (MSL) mission

Fifteen organic and three inorganic compounds were tested for methane (CH₄) evolution under simulated martian conditions of 6.9 mbar; UVC (200-280 nm) flux of 4 W m⁻²; 20 °C; simulated optical depth of 0.1; and a Mars gas composition of CO₂ (95.3%), N₂ (2.7%), Ar (1.7%), O₂ (0.13%), and water vapor (0.03%). All three inorganic compounds (i.e., NaCl, CaCO₃, graphite) failed to evolve methane at the minimum detection level 0.5 ppm, or above. In contrast, all organic compounds evolved methane when exposed to UV irradiation under simulated martian conditions. The polycyclic aromatic hydrocarbon, pyrene, released the most methane per unit of time at 0.175 nmol CH₄ g⁻¹ h⁻¹, and a spectral reflectance target material used for the MER rovers and Phoenix lander released the least methane at 0.00065 nmol CH₄ cm⁻² h⁻¹. Methane was also released from UV-killed bacterial endospores of Bacillus subtilis. Although all organic compounds evolved methane when irradiated with UV photons under martian conditions, the concentrations of residual organics, biogenic signature molecules, and dead microbial cells should be relatively low on the exterior surfaces of the MSL rover, and, thus, not significant sources of methane contamination. In contrast, kapton tape was found to evolve methane at the rate of 0.00165 nmol CH₄ cm⁻² h⁻¹ (16.5 nmol m⁻² h⁻¹) under the UV and martian conditions tested. Although the evolution of methane from kapton tape was found to decline over time, the large amount of kapton tape used on the MSL rover (lower bound estimated at 3 m²) is likely to create a significant source of terrestrial methane contamination during the early part of the mission.     

Ozone abundance on Mars from infrared heterodyne spectra: I. Acquisition, retrieval, and anticorrelation with water vapor

Observations of ozone on Mars were made using the Goddard Space Flight Center's Infrared Heterodyne Spectrometer and Heterodyne Instrument for Planetary Wind and Composition at the NASA Infrared Telescope Facility. Ozone is an important observable tracer of martian photochemistry. Infrared heterodyne spectroscopy with spectral resolution ?¹⁰⁶ is the only technique that directly measures ozone in the martian atmosphere from the surface of the Earth. Ozone column abundances down to the martian surface were acquired in seven data sets taken between 1988 and 2003 at various orbital positions (LS=40°, 74°, 102°, 115°, 202°, 208°, 291°). Ozone abundances are compared with those retrieved using ultraviolet techniques, showing good agreement. Odd hydrogen (HO_X) chemistry predicts anticorrelation of ozone and water vapor abundances. Retrieved ozone abundances consistently show anticorrelation with corresponding water vapor abundances, providing strong confirmation of odd hydrogen activity. Deviation from strict anticorrelation between the observed total column densities of ozone and water vapor suggests that constituent vertical distribution is an additional, significant factor.     

Minimum estimates of the amount and timing of gases released into the martian atmosphere from volcanic eruptions

Volcanism has been a major process during most of the geologic history of Mars. Based on data collected from terrestrial basaltic eruptions, we assume that the volatile content of martian lavas was typically ∼0.5 wt.% water, ∼0.7 wt.% carbon dioxide, ∼0.14 wt.% sulfur dioxide, and contained several other important volatile constituents. From the geologic record of volcanism on Mars we find that during the late Noachian and through the Amazonian volcanic degassing contributed ∼0.8 bar to the martian atmosphere. Because most of the outgassing consisted of greenhouse gases (i.e., CO₂ and SO₂) warmer surface temperatures resulting from volcanic eruptions may have been possible. Our estimates suggest that ∼1.1 × 10²¹ g (∼8 ± 1 m m⁻²) of juvenile water were released by volcanism; slightly more than half the amount contained in the north polar cap and atmosphere. Estimates for released CO₂ (1.6 × 10²¹ g) suggests that a large reservoir of carbon dioxide is adsorbed in the martian regolith or alternatively ∼300 cm cm⁻² of carbonates may have formed, although these materials would not occur readily in the presence of excess SO₂. Up to ∼120 cm cm⁻² (2.2 × 10²⁰ g) of acid rain (H₂SO₄) may have precipitated onto the martian surface as the result of SO₂ degassing. The hydrogen flux resulting from volcanic outgassing may help explain the martian atmospheric D/H ratio. The amount of outgassed nitrogen (∼1.3 mbar) may also be capable of explaining the martian atmospheric ¹⁵N/¹⁴N ratio. Minor gas constituents (HF, HCl, and H₂S) could have formed hydroxyl salts on the surface resulting in the physical weathering of geologic materials. The amount of hydrogen fluoride emitted (1.82 × 10¹⁸ g) could be capable of dissolving a global layer of quartz sand ∼5 mm thick, possibly explaining why this mineral has not been positively identified in spectral observations. The estimates of volcanic outgassing presented here will be useful in understanding how the martian atmosphere evolved over time.     

Ozone abundance on Mars from infrared heterodyne spectra: II. Validating photochemical models

Ozone is an important observable tracer of martian photochemistry, including odd hydrogen (HO_x) species important to the chemistry and stability of the martian atmosphere. Infrared heterodyne spectroscopy with spectral resolution ?¹⁰⁶ provides the only ground-based direct access to ozone absorption features in the martian atmosphere. Ozone abundances were measured with the Goddard Infrared Heterodyne Spectrometer and the Heterodyne Instrument for Planetary Wind and Composition at the NASA Infrared Telescope Facility on Mauna Kea, Hawai'i. Retrieved total ozone column abundances from various latitudes and orbital positions (LS=40°, 74°, 102°, 115°, 202°, 208°, 291°) are compared to those predicted by the first three-dimensional gas phase photochemical model of the martian atmosphere [Lefèvre, F., Lebonnois, S., Montmessin, F., Forget, F., 2004. J. Geophys. Res. 109, doi:10.1029/2004JE002268. E07004]. Observed and modeled ozone abundances show good agreement at all latitudes at perihelion orbital positions (LS=202°, 208°, 291°). Observed low-latitude ozone abundances are significantly higher than those predicted by the model at aphelion orbital positions (LS=40°, 74°, 115°). Heterogeneous loss of odd hydrogen onto water ice cloud particles would explain the discrepancy, as clouds are observed at low latitudes around aphelion on Mars.     

Adsorptive fractionation of HDO on JSC MARS-1 during sublimation with implications for the regolith of Mars

A chamber was constructed to simulate the boundary between the ice table, regolith and atmosphere of Mars and to examine fractionation between H₂O and HDO during sublimation under realistic martian conditions of temperature and pressure. Thirteen experimental runs were conducted with regolith overlying the ice. The thickness and characteristic grain size of the regolith layer as well as the temperature of the underlying ice was varied. From these runs, values for the effective diffusivity, taking into account the effects of adsorption, of the regolith were derived. These effective diffusivities ranged from 1.8 × 10⁻⁴ m² s⁻¹ to 2.2 × 10⁻³ m² s⁻¹ for bare ice and from 2.4 × 10⁻¹¹ m² s⁻¹ to 2.0 × 10⁻⁹ m² s⁻¹ with an adsorptive layer present. From these, latent heats of adsorption of 8.6 ± 2.6 kJ mol⁻¹ and 9.3 ± 2.8 kJ mol⁻¹ were derived at ice-surface temperatures above 223 ± 8 K and 96 ± 28 kJ mol⁻¹ and 104 ± 31 kJ mol⁻¹ respectively for H₂O and HDO were derived at colder temperatures. For temperatures below 223 K, the effective diffusivity of HDO was found to be lower than the diffusivity of H₂O by 40% on average, suggesting that the regolith was adsorptively fractionating the sublimating gas with a fractionation factor of 1.96 ± 0.74. Applying these values to Mars predicts that adsorbed water on the regolith is enriched in HDO compared to the atmosphere, particularly where the regolith is colder. Based on current observations, the D/H ratio of the regolith may be as high as 21 ± 8 times VSMOW at 12°S and L_S = 357° if the regolith is hydrated primarily by the atmosphere, neglecting any hydration from subsurface ice.     

Comparison of PFS and TES observations of temperature and water vapor in the martian atmosphere

The interval from L_s = 330° in Mars Year (MY) 26 until L_s = 84° in MY 27 has been used to compare and validate measurements from the Mars Global Surveyor Thermal Emission Spectrometer (TES) and the Mars Express Planetary Fourier Spectrometer (PFS). We studied differences between atmospheric temperatures observed by the two instruments. The best agreement between atmospheric temperatures was found at 50 Pa between 40°S and 40°N latitude, where differences were within ±5 K. For other atmospheric levels, differences as large as ∼25 K were observed between the two instruments at some locations. The largest temperature differences occurred mainly over the Hellas Planitia, Argyre Planitia, Tharsis and Valles Marineris regions.On this basis we report on the variability of the martian atmosphere during the 5.5 martian years of Mars climatology obtained by combining the two data sets from TES and PFS. Atmospheric temperatures at 50 Pa responded to the global-scale dust storms of MY 25 and in MY 28 raising temperatures from ∼220 K to ∼250 K during the daytime. An atmospheric temperature of ∼140 K at 50 Pa was observed poleward of 70°N during northern winter and poleward of 60°S during southern winter each year in both the PFS and TES results. Water vapor observed by the two spectrometers showed consistent seasonal and latitudinal variations.     

“Methane on Mars: A product of H2O photolysis in the presence of CO”: Response to V.A. Krasnopolsky

The detection of CH₄ in the martian atmosphere, at a mixing ratio of about 10 ppb, prompted Krasnopolsky et al. [Krasnopolsky, V.A., Maillard, J.P., Owen, T.C., 2004. Icarus 172, 537-547] and Krasnopolsky [Krasnopolsky, V.A., 2006. Icarus 180, 359-367] to propose that the CH₄ is of biogenic origin. Bar-Nun and Dimitrov [Bar-Nun, A., Dimitrov, V., 2006. Icarus 181, 320-322] proposed that CH₄ can be formed in the martian atmosphere by photolysis of H₂O in the presence of CO. We based our arguments on a clear demonstration that CH₄ is formed in our experiments, and on thermodynamic equilibrium calculations, which show that CH₄ formation is favored even in the presence of oxygen at a mixing ratio 1.3×¹⁰⁻³, as observed on Mars. In the present comment, Krasnopolsky [Krasnopolsky, V.A., 2007. Icarus, in press (this issue)] presents his arguments against the suggestion of Bar-Nun and Dimitrov [Bar-Nun, A., Dimitrov, V., 2006. Icarus 181, 320-322], based on the effect of O₂ on CH₄ formation, the absence of kinetic pathways for CH₄ formation and on the inadequacy of thermodynamic equilibrium calculations to describe the martian atmosphere. In this rebuttal we demonstrate that experiments with molecular oxygen at a ratio of O₂/CO₂=(8.9-17)×¹⁰⁻³, exceeding the martian ratio, still form CH₄. Thermodynamic equilibrium calculations replicate the experimental CH₄ mixing ratio to within a factor of 1.9 and demonstrate that CH₄ production is favored in the martian atmosphere, which is obviously not in thermodynamic equilibrium. Consequently, we do not find the presence of methane to be a sign of biological activity on Mars.     

Slow degradation of ATP in simulated martian environments suggests long residence times for the biosignature molecule on spacecraft surfaces on Mars

Prelaunch planetary protection protocols on spacecraft are designed to reduce the numbers and diversity of viable bioloads on surfaces in order to mitigate the forward contamination of planetary surfaces. In addition, there is a growing appreciation that prelaunch spacecraft cleaning protocols will be required to reduce the levels of biogenic signature molecules on spacecraft to levels that will not compromise life-detection experiments on landers. The biogenic molecule, adenosine triphosphate (ATP) was tested for long-term stability under simulated Mars surface conditions of high UV flux, low temperature, low pressure, Mars atmosphere, and clear-sky dust loading conditions. Data on UV-induced ATP degradation rates were then extrapolated to a diversity of global conditions using a radiative transfer model for UV on Mars. The UV-induced degradation of ATP tested at 4.1 W m⁻² UVC (200-280 nm), −10 °C, 7.1 mb, 95% CO₂ gas composition, and an atmospheric opacity of τ=0.1 yielded a half-life for ATP of 1342 kJ m⁻²; or extrapolated to approximately 22 sols on equatorial Mars with an atmospheric opacity of τ=0.5. Temperature was found to moderately affect ATP degradation rates under martian conditions; tests at −80 or 20 °C yielded ATP half-lives of 2594 or 1183 kJ m⁻², respectively. The ATP degradation rates reported here are over 10 orders of magnitude slower than the UV-induced biocidal rates reported in the literature on the inactivation of strongly UV-resistant bacterial spores from Bacillus pumilus SAFR-032 [Schuerger, A.C., Richards, J.T., Newcombe, D.A., Venkateswaran, K.J., 2006. Icarus 181, 52-62]. Extrapolating results to global Mars conditions, residence times for a 99% reduction of ATP on spacecraft surfaces ranged from 158 sols on Sun-exposed surfaces to approximately 32,000 sols for the undersides of landers similar to Viking. However, spacecraft materials greatly affected the survival times of ATP under martian conditions. Stainless steel was found to enhance the UV degradation of ATP by over 2 orders of magnitude compared to ATP-doped iridited aluminum, graphite, and astroquartz coupons. Extrapolating these results to global conditions, ATP on stainless steel might be expected to persist between 2 and 320 sols for upper and lower surfaces of landers. Liquid chromatography-mass spectrometry data supported the conclusion that UV irradiation acted to remove the γ-phosphate group from ATP, and no evidence was observed for the UV-degradation of d-ribose or adenine moieties. Long residence times for ATP on spacecraft materials under martian conditions suggest that prelaunch cleaning protocols may need to be strengthened to mitigate against possible ATP contamination of life-detection experiments on Mars landers.     

MARCI and MOC observations of the atmosphere and surface cap in the north polar region of Mars

We used MGS-MOC and MRO-MARCI daily mapping images of the North Polar Region of Mars from 16 August 2005 (L_s = 270°) to 21 May 2009 (L_s = 270°), covering portions of three consecutive martian years (MY 27-MY 29), to observe the seasonal behavior of the polar ice cap and atmospheric phenomena. The rate of cap regression was similar in MY 28 and MY 29, but was advanced by 3.5° of L_s (∼7-8 sols) in MY 29. The spatial and temporal behaviors of dust and condensate clouds were similar in the two years and generally in accord with prior years. Dust storms (>100 km²) were observed in all seasons, with peak activity occurring at L_s = 10-20° from 50°N to 70°N and at L_s = 135-140° from 70°N to 90°N. The most active quadrant was 0-90°W in MY 28, shifting to 180-270°W in MY 29. The majority of regional storms in both years developed in longitudes from 10°W to 60°W. During late summer the larger storms obscure the North Polar Region in a cloud of dust that transitions to north polar hood condensate clouds around autumnal equinox.Changes in the distribution of perennial ice deposits, especially in Olympia Planum, were observed between the 2 years, with the MY 29 ice distribution being the most extensive observed to date. Modeling suggests that the small, bright ice patches on the residual cap are not the result of slope or elevation effects. Rather we suggest that they are the result of local meteorological effects on ice deposition. The annual darkening and brightening of peripheral areas of the residual cap around summer solstice can be explained by the sublimation of a brighter frost layer revealing an underlying darker, ice rich layer that itself either sublimes to reveal brighter material below or acts as a cold trap, attracting condensation of water vapor that brightens the surface. An alternative explanation invokes transport and deposition of dust on the surface from the cap interior, and later removal of that dust. The decrease in cap albedo and accompanying increase in near surface atmospheric stability may be related to the annual minimum of polar storm activity near northern summer solstice.