The ultraviolet reflectance of Enceladus: Implications for surface composition

The reflectance of Saturn’s moon Enceladus has been measured at far ultraviolet (FUV) wavelengths (115-190 nm) by Cassini’s Ultraviolet Imaging Spectrograph (UVIS). At visible and near infrared (VNIR) wavelengths Enceladus’ reflectance spectrum is very bright, consistent with a surface composed primarily of H₂O ice. At FUV wavelengths, however, Enceladus is surprisingly dark - darker than would be expected for pure water ice. Previous analyses have focused on the VNIR spectrum, comparing it to pure water ice (Cruikshank, D.P., Owen, T.C., Dalle Ore, C., Geballe, T.R., Roush, T.L., de Bergh, C., Sandford, S.A., Poulet, F., Benedix, G.K., Emery, J.P. [2005] Icarus, 175, 268-283) or pure water ice plus a small amount of NH₃ (Emery, J.P., Burr, D.M., Cruikshank, D.P., Brown, R.H., Dalton, J.B. [2005] Astron. Astrophys., 435, 353-362) or NH₃ hydrate (Verbiscer, A.J., Peterson, D.E., Skrutskie, M.F., Cushing, M., Helfenstein, P., Nelson, M.J., Smith, J.D., Wilson, J.C. [2006] Icarus, 182, 211-223). We compare Enceladus’ FUV spectrum to existing laboratory measurements of the reflectance spectra of candidate species, and to spectral models. We find that the low FUV reflectance of Enceladus can be explained by the presence of a small amount of NH₃ and a small amount of a tholin in addition to H₂O ice on the surface. The presence of these three species (H₂O, NH₃, and a tholin) appears to satisfy not only the low FUV reflectance and spectral shape, but also the middle-ultraviolet to visible wavelength brightness and spectral shape. We expect that ammonia in the Enceladus plume is transported across the surface to provide a global coating.     

Space weathering of small Koronis family asteroids in the SDSS Moving Object Catalog

Following the work of Rivkin et al. (Rivkin, A.S., Thomas, C.A., Trilling, D.E., Enga, M., Grier, J.A. [2011]. Icarus 211, 1294–1297) and Thomas et al. (Thomas, C.A., Rivkin, A.S, Trilling, D.E., Enga, M., Grier, J.A. [2011a]. Icarus 212, 158–166), we investigate space weathering trends in the Koronis family using the larger sample size of the Sloan Digital Sky Survey Moving Object Catalog. We confirm the trend in spectral slope seen in our earlier work and extend our results by investigating the trend in band depth (i ? z color index) to show that Koronis family asteroids smaller than 4 km show the transition from fresh Q-type to weathered S-type surfaces.     

Ground-based near-infrared observations of water vapour in the Venus troposphere

We present a study of water vapour in the Venus troposphere obtained by modelling specific water vapour absorption bands within the 1.18 μm window. We compare the results with the normal technique of obtaining the abundance by matching the peak of the 1.18 μm window. Ground-based infrared imaging spectroscopy of the night side of Venus was obtained with the Anglo-Australian Telescope and IRIS2 instrument with a spectral resolving power of R ～ 2400. The spectra have been fitted with modelled spectra simulated using the radiative transfer model VSTAR. We find a best fit abundance of 31 ppmv (?6 +9 ppmv), which is in agreement with recent results by Bézard et al. (Bézard, B., Fedorova, A., Bertaux, J.-L., Rodin, A., Korablev, O. [2011]. Icarus, 216, 173–183) using VEX/SPICAV (R ～ 1700) and contrary to prior results by Bézard et al. (Bézard, B., de Bergh, C., Crisp, D., Maillard, J.P. [1990]. Nature, 345, 508–511) of 44 ppmv (±9 ppmv) using VEX/VIRTIS-M (R ～ 200) data analyses. Comparison studies are made between water vapour abundances determined from the peak of the 1.18 μm window and abundances determined from different water vapour absorption features within the near infrared window. We find that water vapour abundances determined over the peak of the 1. 18 μm window results in plots with less scatter than those of the individual water vapour features and that analyses conducted over some individual water vapour features are more sensitive to variation in water vapour than those over the peak of the 1. 18 μm window. No evidence for horizontal spatial variations across the night side of the disk are found within the limits of our data with the exception of a possible small decrease in water vapour from the equator to the north pole. We present spectral ratios that show water vapour absorption from within the lowest 4 km of the Venus atmosphere only, and discuss the possible existence of a decreasing water vapour concentration towards the surface.     

Looking a gift horse in the mouth: Evaluation of wide-field asteroid photometric surveys

It has recently become possible to do a photometric survey of many asteroids at once, rather than observing single asteroids one (or occasionally a couple) at a time. We evaluate two such surveys. Dermawan et al. (Dermawan et al. [2011]. Publ. Astron. Soc. Jpn. 63, S555–S576) observed one night on the Subaru 8.2 m telescope, and Masiero et al. (Masiero, J., Jedicke, R., Durech, J., Gwen, S., Denneau, L., Larsen, J. [2009]. Icarus 204, 145–171) observed six nights over 2 weeks with the 3.6 m CFHT. Dermawan claimed 83 rotation periods from 127 detected asteroids; Masiero et al. claimed 218 rotation periods from 828 detections. Both teams claim a number of super-fast rotators (P < 2.2 h) among main belt asteroids larger than 250 m diameter, some up to several km in diameter. This would imply that the spin rate distribution of main belt asteroids differs from like-sized NEAs, that there are larger super-fast rotators (monolithic asteroids) in the main belt than among NEAs. Here we evaluate these survey results, applying the same criteria for reliability of results that we apply to all results listed in our Lightcurve Database (Warner, B.D., Harris, A.W., Pravec, P. [2009a]. Icarus 202, 134–146). In doing so, we assigned reliability estimates judged sufficient for inclusion in statistical studies for only 27 out of 83 (33%) periods claimed by Dermawan, and only 87 out of 218 (40%) periods reported by Masiero et al.; none of the super-fast rotators larger than about 250 m diameter claimed by either survey received a reliability rating judged sufficient for analysis. We find no reliable basis for the claim of different rotation properties between main belt and near-Earth asteroids. Our analysis presents a cautionary message for future surveys.     

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.     

General circulation and transport in Saturn’s upper troposphere and stratosphere

We present a model for the general circulation and dynamical transport in Saturn’s upper troposphere and stratosphere and derive the effective advective circulation and eddy transport coefficients required for use in two-dimensional (latitude–altitude) photochemistry–transport models. A three-dimensional Outer-Planet General Circulation Model (OPGCM) is used to generate the transport data. We find that the OPGCM adequately captures the global-scale, pole-to pole temperature contrast, but overestimates mid- and high-latitude temperatures in the summer hemisphere by ～5 K. In addition, the model reproduces the local temperature minimum seen at the equator in Cassini Composite Infrared Spectrometer (CIRS) 0.1-mbar data but not the local maximum in 1-mbar temperatures, suggesting that it is capturing the phase of Saturn’s Semiannual Oscillation associated with a temperature minimum at the equator but not the opposite phase. The meridional circulation at low latitudes is found to be dominated by a seasonally reversing Hadley circulation, characterized by upwelling near the equator, cross-equatorial flow from summer to winter hemisphere, and strong subsidence centered near 25° latitude in the winter hemisphere. The cross-equatorial flow induces an asymmetry in which the equatorial jet is found to be stronger in the winter than in the summer stratosphere. The location of the subsidence near 25°N for L_s ～ 310° coincides with local maxima in acetylene, diacetylene, and methylacetylene mixing ratios measured by Cassini/CIRS (Guerlet, S., Fouchet, T., Bézard, B., Moses, J.I., Fletcher, L.N., Simon-Miller, A.A., Flasar, F.M. [2010]. Icarus 209, 682–695). This result supports the suggestion by Guerlet et al. (2010) that the hydrocarbon abundances are enhanced at this latitude by pronounced downward transport of hydrocarbon-rich air from above. The lateral eddy diffusion coefficient is found to typically be ～10⁵–10⁶ m² s^?1 at mid-latitudes, implying meridional eddy transport time scales of order 100–1000 years.     

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.     

Radiative transfer analyses of Titan’s tropical atmosphere

Titan’s optical and near-IR spectra result primarily from the scattering of sunlight by haze and its absorption by methane. With a column abundance of 92 km amagat (11 times that of Earth), Titan’s atmosphere is optically thick and only ～10% of the incident solar radiation reaches the surface, compared to 57% on Earth. Such a formidable atmosphere obstructs investigations of the moon’s lower troposphere and surface, which are highly sensitive to the radiative transfer treatment of methane absorption and haze scattering. The absorption and scattering characteristics of Titan’s atmosphere have been constrained by the Huygens Probe Descent Imager/Spectral Radiometer (DISR) experiment for conditions at the probe landing site (Tomasko, M.G., Bézard, B., Doose, L., Engel, S., Karkoschka, E. [2008a]. Planet. Space Sci. 56, 624–247; Tomasko, M.G. et al. [2008b]. Planet. Space Sci. 56, 669–707). Cassini’s Visual and Infrared Mapping Spectrometer (VIMS) data indicate that the rest of the atmosphere (except for the polar regions) can be understood with small perturbations in the high haze structure determined at the landing site (Penteado, P.F., Griffith, C.A., Tomasko, M.G., Engel, S., See, C., Doose, L., Baines, K.H., Brown, R.H., Buratti, B.J., Clark, R., Nicholson, P., Sotin, C. [2010]. Icarus 206, 352–365). However the in situ measurements were analyzed with a doubling and adding radiative transfer calculation that differs considerably from the discrete ordinates codes used to interpret remote data from Cassini and ground-based measurements. In addition, the calibration of the VIMS data with respect to the DISR data has not yet been tested. Here, VIMS data of the probe landing site are analyzed with the DISR radiative transfer method and the faster discrete ordinates radiative transfer calculation; both models are consistent (to within 0.3%) and reproduce the scattering and absorption characteristics derived from in situ measurements. Constraints on the atmospheric opacity at wavelengths outside those measured by DISR, that is from 1.6 to 5.0 μm, are derived using clouds as diffuse reflectors in order to derive Titan’s surface albedo to within a few percent error and cloud altitudes to within 5 km error. VIMS spectra of Titan at 2.6–3.2 μm indicate not only spectral features due to CH₄ and CH₃D (Rannou, P., Cours, T., Le Mouélic, S., Rodriguez, S., Sotin, C., Drossart, P., Brown, R. [2010]. Icarus 208, 850–867), but also a fairly uniform absorption of unknown source, equivalent to the effects of a darkening of the haze to a single scattering albedo of 0.63 ± 0.05. Titan’s 4.8 μm spectrum point to a haze optical depth of 0.2 at that wavelength. Cloud spectra at 2 μm indicate that the far wings of the Voigt profile extend 460 cm^?1 from methane line centers. This paper releases the doubling and adding radiative transfer code developed by the DISR team, so that future studies of Titan’s atmosphere and surface are consistent with the findings by the Huygens Probe. We derive the surface albedo at eight spectral regions of the 8 × 12 km² area surrounding the Huygens landing site. Within the 0.4–1.6 μm spectral region our surface albedos match DISR measurements, indicating that DISR and VIMS measurements are consistently calibrated. These values together with albedos at longer 1.9–5.0 μm wavelengths, not sampled by DISR, resemble a dark version of the spectrum of Ganymede’s icy leading hemisphere. The eight surface albedos of the landing site are consistent with, but not deterministic of, exposed water ice with dark impurities.     

Enceladus: A hypothesis for bringing both heat and chemicals to the surface

The eruptive plumes and large heat flow (～15 GW) observed by Cassini in the South Polar Region of Enceladus may be expressions of hydrothermal activity inside Enceladus. We hypothesize that a subsurface ocean is the heat reservoir for thermal anomalies on the surface and the source of heat and chemicals necessary for the plumes. The ocean is believed to contain dissolved gases, mostly CO₂ and is found to be relatively warm (～0 °C). Regular tidal forces open cracks in the icy crust above the ocean. Ocean water fills these fissures. There, the conditions are met for the upward movement of water and the dissolved gases to exsolve and form bubbles, lowering the bulk density of the water column and making the pressure at its bottom less than that at the top of the ocean. This pressure difference drives ocean water into and up the conduits toward the surface. This transportation mechanism supports the thermal anomalies and delivers heat and chemicals to the chambers from which the plumes erupt. Water enters these chambers and there its bubbles pop and loft an aerosol mist into the ullage. The exiting plume gas entrains some of these small droplets. Thus, nonvolatile chemical species in ocean water can be present in the plume particles. A CO₂ equivalent-gas molar fraction of ～4 × 10^?4 for the ocean is sufficient to support the circulation. A source of heat is needed to keep the ocean warm at ～0 °C (about two degrees above its freezing point). The source of heat is unknown, but our hypothesis is not dependent on any particular mechanism for producing the heat.     

Rotation periods of binary asteroids with large separations – Confronting the Escaping Ejecta Binaries model with observations

Durda et al. (Durda, D.D., Bottke, W.F., Enke, B.L., Merline, W.J., Asphaug, E., Richardson, D.C., Leinhardt, Z.M. [2004]. Icarus 170, 243–257), using numerical models, suggested that binary asteroids with large separation, called Escaping Ejecta Binaries (EEBs), can be created by fragments ejected from a disruptive impact event. It is thought that six binary asteroids recently discovered might be EEBs because of the high separation between their components (～100 > a/R_p > ～20).However, the rotation periods of four out of the six objects measured by our group and others and presented here show that these suspected EEBs have fast rotation rates of 2.5–4 h. Because of the small size of the components of these binary asteroids, linked with this fast spinning, we conclude that the rotational-fission mechanism, which is a result of the thermal YORP effect, is the most likely formation scenario. Moreover, scaling the YORP effect for these objects shows that its timescale is shorter than the estimated ages of the three relevant Hirayama families hosting these binary asteroids. Therefore, only the largest (D ～ 19 km) suspected asteroid, (317) Roxane, could be, in fact, the only known EEB.In addition, our results confirm the triple nature of (3749) Balam by measuring mutual events on its lightcurve that match the orbital period of a nearby satellite in addition to its distant companion. Measurements of (1509) Esclangona at different apparitions show a unique shape of the lightcurve that might be explained by color variations.     

The reflectivity spectrum and opposition effect of Titan's surface observed by Huygens' DISR spectrometers

We determined Titan's reflectivity spectrum near the Huygens' landing site from observations taken with the Descent Imager/Spectral Radiometer below 500 m altitude, in particular the downward-looking photometer and spectrometers. We distinguish signal coming from illumination by sunlight and the lamp onboard Huygens based on their different spectral signatures. For the sunlight data before landing, we find that spatial variations of Titan's reflectivity were only ～0.8%, aside from the phase angle dependence, indicating that the probed area within ～100 m of the landing site was very homogeneous. Only the very last spectrum taken before landing gave a 3% brighter reflectivity, which probably was caused by one bright cobble inside its footprint. The contrast of the cobble was higher at 900 nm wavelength than at 600 nm.For the data from lamp illumination, we confirm that the phase function of Titan's surface displays a strong opposition effect as found by Schröder and Keller (2009. Planetary and Space Science 57, 1963–1974). We extend the phase function to even smaller phase angles (0.02°), which are among the smallest phase angles observed in the solar system. We also confirm the reflectivity spectrum of the dark terrain near the Huygens' landing site between 900 and 1600 nm wavelength by Schröder and Keller (2008. Planetary and Space Science 56, 753–769), but extend the spectrum down to 435 nm wavelength. The reflectivity at zero phase angle peaks at 0.45±0.06 around 750 nm wavelength and drops down to roughly 0.2 at both spectral ends. Our reflectivity of 0.45 is much higher than all previously reported values because our observations probe lower phase angles than others. The spectrum is very smooth except for a known absorption feature longward of 1350 nm. We did not detect any significant variation of the spectral shape along the slit for exposures after landing, probing a 25×4 cm² area. However, the recorded spectral shape was slightly different for exposures before and after landing. This difference is similar to the spectral differences seen on scales of kilometers (Keller et al., 2008. Planetary and Space Science 56, 728–752), indicating that most observations may probe spatially variable contributions from two basic materials, such as a dark soil partially covered by bright cobbles.We used the methane absorption features to constrain the methane mixing ratio near the surface to 5.0±0.3%, in agreement with the 4.92±0.24% value measured in situ by Niemann et al. (2005. Nature 438, 779–784), but smaller than their revised value of 5.65±0.18% (Niemann et al., 2010. Journal of Geophysical Research 115, E12006). Our results were made possible by an in depth review of the calibration of the spectroscopic and photometric data.     

Thermal structure and CO distribution for the Venus mesosphere/lower thermosphere: 2001–2009 inferior conjunction sub-millimeter CO absorption line observations

Sub-millimeter ¹²CO (346 GHz) and ¹³CO (330 GHz) line absorptions, formed in the mesosphere and lower thermosphere of Venus (70–120 km), have been mapped across the nightside Venus disk during 2001–2009 inferior conjunctions, employing the James Clerk Maxwell Telescope (JCMT). Radiative transfer analysis of these thermal line absorptions supports temperature and CO mixing profile retrievals, as well as Doppler wind fields (described in the companion paper, Clancy et al., 2012). Temporal sampling over the hourly, daily, weekly and interannual timescales was obtained over 2001–2009. On timescales inferred as several weeks, we observe changes between very distinctive CO and temperature nightside distributions. Retrieved nightside CO, temperature distributions for January 2006 and August 2007 observations display strong local time, latitudinal gradients consistent with early morning (2–3 am), low-to-mid latitude (0–40NS) peaks of 100–200% in CO and 20–30 K in temperature. The temperature increases are most pronounced above 100 km altitudes, whereas CO variations extend from 105 km (top altitude of retrieval) down to below 80 km in the mesosphere. In contrast, the 2004 and 2009 periods of observation display modest temperature (5–10 K) and CO (30–60%) increases, that are centered on antisolar (midnight) local times and equatorial latitudes. Doppler wind derived global (zonal and should be SSAS) circulations from the same data do not exhibit variations correlated with these CO, temperature short-term variations. However, large-scale residual wind fields not fit by the zonal, SSAS circulations are observed in concert with the strong temperature, CO gradients observed in 2006 and 2007 (Clancy et al., 2010). These short term variations in nightside CO, temperature distributions may also be related to observed nightside variations in O₂ airglow (Hueso, H., Sánchez-Lavega, A., Piccioni, G., Drossart, P., Gérard, J.C., Khatuntsev, I., Zasova, L., Migliorini, A. [2008]. J. Geophys. Res. 113, E00B02. doi:10.1029/2008JE003081) and upper mesospheric SO and SO₂ layers (Sandor, B.J., Clancy, R.T., Moriarty-Schieven, G.H., Mills, F.P. [2010]. Icarus 208, 49–60).The retrieved temperature profiles also exhibit 20 K long-term (2001–2009) variations in nightside (whole disk) average mesospheric (80–95 km) temperatures, similar to 1982–1991 variations identified in previous millimeter CO line observations (Clancy et al., 1991). Global average diurnal variations in lower thermospheric temperatures and mesospheric CO abundances decreased by a factor-of-two between 2000–2002 versus 2007–2009 periods of combined dayside and nightside observations. The infrequency and still limited temporal extent of the observations make it difficult to assign specific timescales to such longer term variations, which may be associated with longer term variations observed for cloud top SO₂ (Esposito, L.W., Bertaux, J.-L., Krasnopolsky, V., Moroz, V.I., Zasova, L.V. [1997]. Chemistry of lower atmosphere and clouds. In: Bougher, S.W., Hunten, D.M., Phillips, R.J. (Eds.), VENUS II, 1362pp) and mesospheric water vapor (Sandor, B.J., Clancy, R.T. [2005]. Icarus 177, 129–143) abundances.     

The remarkable surface homogeneity of the Dawn mission target (1) Ceres

Dwarf-planet (1) Ceres is one of the two targets, along with (4) Vesta, that will be studied by the NASA Dawn spacecraft via imaging, visible and near-infrared spectroscopy, and gamma-ray and neutron spectroscopy. While Ceres’ visible and near-infrared disk-integrated spectra have been well characterized, little has been done about quantifying spectral variations over the surface. Any spectral variation would give us insights on the geographical variation of the composition and/or the surface age. The only work so far was that of Rivkin and Volquardsen ([2010], Icarus 206, 327) who reported rotationally-resolved spectroscopic (disk-integrated) observations in the 2.2–4.0 μm range; their observations showed evidence for a relatively uniform surface.Here, we report disk-resolved observations of Ceres with SINFONI (ESO VLT) in the 1.17–1.32 μm and 1.45–2.35 μm wavelength ranges. The observations were made under excellent seeing conditions (0.6″), allowing us to reach a spatial resolution of ～75 km on Ceres’ surface. We do not find any spectral variation above a 3% level, suggesting a homogeneous surface at our spatial resolution. Slight variations (about 2%) of the spectral slope are detected, geographically correlated with the albedo markings reported from the analysis of the HST and Keck disk-resolved images of Ceres (Li et al. [2006], Icarus 182, 143; Carry et al. [2008], Astron. Astrophys. 478, 235). Given the lack of constraints on the surface composition of Ceres, however, we cannot assert the causes of these variations.     

Absolute magnitudes of asteroids and a revision of asteroid albedo estimates from WISE thermal observations

We obtained estimates of the Johnson V absolute magnitudes (H) and slope parameters (G) for 583 main-belt and near-Earth asteroids observed at Ond?ejov and Table Mountain Observatory from 1978 to 2011. Uncertainties of the absolute magnitudes in our sample are <0.21 mag, with a median value of 0.10 mag. We compared the H data with absolute magnitude values given in the MPCORB, Pisa AstDyS and JPL Horizons orbit catalogs. We found that while the catalog absolute magnitudes for large asteroids are relatively good on average, showing only little biases smaller than 0.1 mag, there is a systematic offset of the catalog values for smaller asteroids that becomes prominent in a range of H greater than ～10 and is particularly big above H ～ 12. The mean (H_catalog ? H) value is negative, i.e., the catalog H values are systematically too bright. This systematic negative offset of the catalog values reaches a maximum around H = 14 where the mean (H_catalog ? H) is ?0.4 to ?0.5. We found also smaller correlations of the offset of the catalog H values with taxonomic types and with lightcurve amplitude, up to ～0.1 mag or less. We discuss a few possible observational causes for the observed correlations, but the reason for the large bias of the catalog absolute magnitudes peaking around H = 14 is unknown; we suspect that the problem lies in the magnitude estimates reported by asteroid surveys. With our photometric H and G data, we revised the preliminary WISE albedo estimates made by Masiero et al. (Masired, J.R. et al. [2011]. Astrophys. J. 741, 68–89) and Mainzer et al. (Mainzer, A. et al. [2011b]. Astrophys. J. 743, 156–172) for asteroids in our sample. We found that the mean geometric albedo of Tholen/Bus/DeMeo C/G/B/F/P/D types with sizes of 25–300 km is p_V = 0.057 with the standard deviation (dispersion) of the sample of 0.013 and the mean albedo of S/A/L types with sizes 0.6–200 km is 0.197 with the standard deviation of the sample of 0.051. The standard errors of the mean albedos are 0.002 and 0.006, respectively; systematic observational or modeling errors can predominate over the quoted formal errors. There is apparent only a small, marginally significant difference of 0.031 ± 0.011 between the mean albedos of sub-samples of large and small (divided at diameter 25 km) S/A/L asteroids, with the smaller ones having a higher albedo. The difference will have to be confirmed and explained; we speculate that it may be either a real size dependence of surface properties of S type asteroids or a small size-dependent bias in the data (e.g., a bias towards higher albedos in the optically-selected sample of asteroids). A trend of the mean of the preliminary WISE albedo estimates increasing with asteroid size decreasing from D ～ 30 down to ～5 km (for S types) showed in Mainzer et al. (Mainzer, A. et al. [2011a]. Astrophys. J. 741, 90–114) appears to be mainly due to the systematic bias in the MPCORB absolute magnitudes that progressively increases with H in the corresponding range H = 10–14.     

Influence of methane concentration on the optical indices of Titan’s aerosols analogues

This work deals with the optical constant characterization of Titan aerosol analogues or “tholins” produced with the PAMPRE experimental setup and deposited as thin films onto a silicon substrate. Tholins were produced in different N₂–CH₄ gaseous mixtures to study the effect of the initial methane concentration on their optical constants. The real (n) and imaginary (k) parts of the complex refractive index were determined using the spectroscopic ellipsometry technique in the 370–1000 nm wavelength range. We found that optical constants depend strongly on the methane concentrations of the gas phase in which tholins are produced: imaginary optical index (k) decreases with initial CH₄ concentration from 2.3 × 10^?2 down to 2.7 × 10^?3 at 1000 nm wavelength, while the real optical index (n) increases from 1.48 up to 1.58 at 1000 nm wavelength. The larger absorption in the visible range of tholins produced at lower methane percentage is explained by an increase of the secondary and primary amines signature in the mid-IR absorption. Comparison with results of other tholins and data from Titan observations are presented. We found an agreement between our values obtained with 10% methane concentration, and Imanaka et al. (Imanaka, H., Khare, B.N., Elsila, J.E., Bakes, E.L.O., McKay, C.P., Cruikshank, D.P., Sugita, S., Matsui, T., Zare, R.N. [2004]. Icarus, 168, 344–366) values, in spite of the difference in the analytical method. This confirms a reliability of the optical properties of tholins prepared with various setups but with similar plasma conditions. Our comparison with Titan’s observations also raises a possible inconsistency between the mid-IR aerosol signature by VIMS and CIRS Cassini instruments and the visible Huygens-DISR derived data. The mid-IR VIMS and CIRS signatures are in agreement with an aerosol dominated by an aliphatic carbon content, whereas the important visible absorption derived from the DISR measurement seems to be incompatible with such an important aliphatic content, but more compatible with an amine-rich aerosol.     

Atomic oxygen on the Venus nightside: Global distribution deduced from airglow mapping

The Visible and Infra-Red Thermal Imaging Spectrometer (VIRTIS) instrument on board the Venus Express spacecraft has measured the O₂(a¹Δ) nightglow distribution at 1.27 μm in the Venus mesosphere for more than two years. Nadir observations have been used to create a statistical map of the emission on Venus nightside. It appears that the statistical 1.6 MR maximum of the emission is located around the antisolar point. Limb observations provide information on the altitude and on the shape of the emission layer. We combine nadir observations essentially covering the southern hemisphere, corrected for the thermal emission of the lower atmosphere, with limb profiles of the northern hemisphere to generate a global map of the Venus nightside emission at 1.27 μm. Given all the O₂(a¹Δ) intensity profiles, O₂(a¹Δ) and O density profiles have been calculated and three-dimensional maps of metastable molecular and atomic oxygen densities have been generated. This global O density nightside distribution improves that available from the VTS3 model, which was based on measurements made above 145 km. The O₂(a¹Δ) hemispheric average density is 2.1 × 10⁹ cm^?3, with a maximum value of 6.5 × 10⁹ cm^?3 at 99.2 km. The O density profiles have been derived from the nightglow data using CO₂ profiles from the empirical VTS3 model or from SPICAV stellar occultations. The O hemispheric average density is 1.9 × 10¹¹ cm^?3 in both cases, with a mean altitude of the peak located at 106.1 km and 103.4 km, respectively. These results tend to confirm the modeled values of 2.8 × 10¹¹ cm^?3 at 104 km and 2.0 × 10¹¹ cm^?3 at 110 km obtained by Brecht et al. [Brecht, A., Bougher, S.W., Gérard, J.-C., Parkinson, C.D., Rafkin, S., Foster, B., 2011a. J. Geophys. Res., in press] and Krasnopolsky [Krasnopolsky, V.A., 2010. Icarus 207, 17–27], respectively. Comparing the oxygen density map derived from the O₂(a¹Δ) nightglow observations, it appears that the morphology is very different and that the densities obtained in this study are about three times higher than those predicted by the VTS3 model.     

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.     

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.     

Stardust-NExT,Deep Impact,and the accelerating spin of 9P/Tempel 1

The evolution of the spin rate of Comet 9P/Tempel 1 through two perihelion passages (in 2000 and 2005) is determined from 1922 Earth-based observations taken over a period of 13 year as part of a World-Wide observing campaign and from 2888 observations taken over a period of 50 days from the Deep Impact spacecraft. We determine the following sidereal spin rates (periods): 209.023 ± 0.025°/dy (41.335 ± 0.005 h) prior to the 2000 perihelion passage, 210.448 ± 0.016°/dy (41.055 ± 0.003 h) for the interval between the 2000 and 2005 perihelion passages, 211.856 ± 0.030°/dy (40.783 ± 0.006 h) from Deep Impact photometry just prior to the 2005 perihelion passage, and 211.625 ± 0.012°/dy (40.827 ± 0.002 h) in the interval 2006–2010 following the 2005 perihelion passage. The period decreased by 16.8 ± 0.3 min during the 2000 passage and by 13.7 ± 0.2 min during the 2005 passage suggesting a secular decrease in the net torque. The change in spin rate is asymmetric with respect to perihelion with the maximum net torque being applied on approach to perihelion. The Deep Impact data alone show that the spin rate was increasing at a rate of 0.024 ± 0.003°/dy/dy at JD2453530.60510 (i.e., 25.134 dy before impact), which provides independent confirmation of the change seen in the Earth-based observations.The rotational phase of the nucleus at times before and after each perihelion and at the Deep Impact encounter is estimated based on the Thomas et al. (Thomas et al. [2007]. Icarus 187, 4–15) pole and longitude system. The possibility of a 180° error in the rotational phase is assessed and found to be significant. Analytical and physical modeling of the behavior of the spin rate through of each perihelion is presented and used as a basis to predict the rotational state of the nucleus at the time of the nominal (i.e., prior to February 2010) Stardust-NExT encounter on 2011 February 14 at 20:42.We find that a net torque in the range of 0.3–2.5 × 10⁷ kg m² s^?2 acts on the nucleus during perihelion passage. The spin rate initially slows down on approach to perihelion and then passes through a minimum. It then accelerates rapidly as it passes through perihelion eventually reaching a maximum post-perihelion. It then decreases to a stable value as the nucleus moves away from the Sun. We find that the pole direction is unlikely to precess by more than ～1° per perihelion passage. The trend of the period with time and the fact that the modeled peak torque occurs before perihelion are in agreement with published accounts of trends in water production rate and suggests that widespread H₂O out-gassing from the surface is largely responsible for the observed spin-up.     

The role of ejecta in the small crater populations on the mid-sized saturnian satellites

We find evidence, by both observation and analysis, that primary crater ejecta play an important role in the small crater (less than a few km) populations on the saturnian satellites, and more broadly, on cratered surfaces throughout the Solar System. We measure crater populations in Cassini images of Enceladus, Rhea, and Mimas, focusing on image data with scales less than 500 m/pixel. We use recent updates to crater scaling laws and their constants (Housen, K.R., Holsapple, K.A. [2011]. Icarus 211, 856–875) to estimate the amount of mass ejected in three different velocity ranges: (i) greater than escape velocity, (ii) less than escape velocity and faster than the minimum velocity required to make a secondary crater (v_min), and (iii) velocities less than v_min. Although the vast majority of mass on each satellite is ejected at speeds less than v_min, our calculations demonstrate that the differences in mass available in the other two categories should lead to observable differences in the small crater populations; the predictions are borne out by the measurements we have made to date. In particular, Rhea, Tethys, and Dione have sufficient surface gravities to retain ejecta moving fast enough to make secondary crater populations. The smaller satellites, such as Enceladus but especially Mimas, are expected to have little or no traditional secondary populations because their escape velocities are near the threshold velocity necessary to make a secondary crater. Our work clarifies why the Galilean satellites have extensive secondary crater populations relative to the saturnian satellites. The presence, extent, and sizes of sesquinary craters (craters formed by ejecta that escape into temporary orbits around Saturn before re-impacting the surface, see Dobrovolskis, A.R., Lissauer, J.J. [2004]. Icarus 169, 462–473; Alvarellos, J.L., Zahnle, K.J., Dobrovolskis, A.R., Hamill, P. [2005]. Icarus 178, 104–123; Zahnle, K., Alvarellos, J.L., Dobrovolskis, A.R., Hamill, P. [2008]. Icarus 194, 660–674) is not yet well understood. Finally, our work provides further evidence for a “shallow” size–frequency distribution (slope index of ～2 for a differential power-law) for comets a few kilometers diameter and smaller.