How the Enceladus dust plume feeds Saturn’s E ring

Pre-Cassini models of Saturn’s E ring [Horányi, M., Burns, J., Hamilton, D., 1992. Icarus 97, 248-259; Juhász, A., Horányi, M., 2002. J. Geophys. Res. 107, 1-10] failed to reproduce its peculiar vertical structure inferred from Earth-bound observations [de Pater, I., Martin, S.C., Showalter, M.R., 2004. Icarus 172, 446-454]. After the discovery of an active ice-volcanism of Saturn’s icy moon Enceladus the relevance of the directed injection of particles for the vertical ring structure of the E ring was swiftly recognised [Juhász, A., Horányi, M., Morfill, G.E., 2007. Geophys. Res. Lett. 34, L09104; Kempf, S., Beckmann, U., Moragas-Klostermeyer, G., Postberg, F., Srama, R., Economou, T., Schmidt, J., Spahn, F., Grün, E., 2008. Icarus 193, 420-437]. However, simple models for the delivery of particles from the plume to the ring predict a too small vertical ring thickness and overestimate the amount of the injected dust.Here we report on numerical simulations of grains leaving the plume and populating the dust torus of Enceladus. We run a large number of dynamical simulations including gravity and Lorentz force to investigate the earliest phase of the ring particle life span. The evolution of the electrostatic charge carried by the initially uncharged grains is treated selfconsistently. Freshly ejected plume particles are moving in almost circular orbits because the Enceladus orbital speed exceeds the particles’ ejection speeds by far. Only a small fraction of grains that leave the Hill sphere of Enceladus survive the next encounter with the moon. Thus, the flux and size distribution of the surviving grains, replenishing the ring particle reservoir, differs significantly from the flux and size distribution of the particles freshly ejected from the plume. Our numerical simulations reproduce the vertical ring profile measured by the Cassini Cosmic Dust Analyzer (CDA) [Kempf, S., Beckmann, U., Moragas-Klostermeyer, G., Postberg, F., Srama, R., EconoDmou, T., Smchmidt, J., Spahn, F., Grün, E., 2008. Icarus 193, 420-437]. From our simulations we calculate the deposition rates of plume particles hitting Enceladus’ surface. We find that at a distance of 100 m from a jet a 10 m sized ice boulder should be covered by plume particles in 10⁵-10⁶ years.     

Geysers of Enceladus: Quantitative analysis of qualitative models

Aspects of two qualitative models of Enceladus’ dust plume—the so-called “Cold Faithful” [Porco, C.C., et al., 2006. Cassini observes the active south pole of Enceladus. Science 311, 1393-1401; Ingersoll, A.P., et al., 2006. Models of the Enceladus plumes. In: Bulletin of the American Astronomical Society, vol. 38, p. 508] and “Frigid Faithful” [Kieffer, S.W., et al., 2006. A clathrate reservoir hypothesis for Enceladus’ south polar plume. Science 314, 1764; Gioia, G., et al., 2007. Unified model of tectonics and heat transport in a Frigid Enceladus. Proc. Natl. Acad. Sci. 104, 13578-13591] models—are analyzed quantitatively. The former model assumes an explosive boiling of subsurface liquid water, when pressure exerted by the ice crust is suddenly released due to an opening crack. In the latter model the existence of a deep shell of clathrates below Enceladus’ south pole is conjectured; clathrates can decompose explosively when exposed to vacuum through a fracture in the outer icy shell. For the Cold Faithful model we estimate the maximal velocity of ice grains, originating from water splashing in explosive boiling. We find that for water near the triple point this velocity is far too small to explain the observed plume properties. For the Frigid Faithful model we consider the problem of momentum transfer from gas to ice particles. It arises since any change in the direction of the gas flow in the cracks of the shell requires re-acceleration of the entrained grains. While this effect may explain the observed speed difference of gas and grains if the gas evaporates from triple point temperature (273.15 K) [Schmidt, J., et al., 2008. Formation of Enceladus dust plume. Nature 451, 685], the low temperatures of the Frigid Faithful model imply a too dilute vapor to support the observed high particle fluxes in Enceladus’ plume.     

Thermal inertia and bolometric Bond albedo values for Mimas, Enceladus, Tethys, Dione, Rhea and Iapetus as derived from Cassini/CIRS measurements

Spectra taken by Cassini’s Composite Infrared Spectrometer (CIRS) between 10 and 600 cm⁻¹ (17-1000 μm) of surface thermal emission of Mimas, Enceladus, Tethys, Dione, Rhea and Iapetus have been used to derive the thermal inertia and bolometric Bond albedo values. Only an upper limit for the bolometric Bond albedo of Iapetus’ dark leading side could be determined due to the insensitivity of the thermal model to albedo when albedos are very low. The thermal inertia in this region however is better constrained. The CIRS coverage of Enceladus is extensive enough that the latitudinal variation in these values from 60°S to 70°N has been determined in 10° wide bins. The bolometric Bond albedos determined here are consistent with literature values which show the surface of the saturnian icy moons to be covered in ice contaminated to varying degrees. The thermal inertia of the moons is shown to be in the range 9-, approximately 2-6 times lower than that of the Galilean satellites, implying a less well consolidated and more porous surface. The thermal inertias of Iapetus and Phoebe are somewhat higher, suggesting that the very low thermal inertias of satellites from Rhea inwards may be related to their probable coating of E-ring material. Latitudinal variations on the surface of Enceladus show that the bolometric Bond albedo and thermal inertia increase towards the active plume source at the south pole.     

The E ring in the vicinity of Enceladus: I. Spatial distribution and properties of the ring particles

Saturn's diffuse E ring is the largest ring of the Solar System and extends from about (Saturn radius RS=60,330 km) to at least encompassing the icy moons Mimas, Enceladus, Tethys, Dione, and Rhea. After Cassini's insertion into her saturnian orbit in July 2004, the spacecraft performed a number of equatorial as well as steep traversals through the E ring inside the orbit of the icy moon Dione. Here, we report about dust impact data we obtained during 2 shallow and 6 steep crossings of the orbit of the dominant ring source—the ice moon Enceladus. Based on impact data of grains exceeding 0.9 μm we conclude that Enceladus feeds a torus populated by grains of at least this size along its orbit. The vertical ring structure at agrees well with a Gaussian with a full-width-half-maximum (FWHM) of ∼4200 km. We show that the FWHM at is due to three-body interactions of dust grains ejected by Enceladus' recently discovered ice volcanoes with the moon during their first orbit. We find that particles with initial speeds between 225 and 235 m s⁻¹ relative to the moon's surface dominate the vertical distribution of dust. Particles with initial velocities exceeding the moon's escape speed of 207 m s⁻¹ but slower than 225 m s⁻¹ re-collide with Enceladus and do not contribute to the ring particle population. We find the peak number density to range between 16×¹⁰⁻² m⁻³ and 21×¹⁰⁻² m⁻³ for grains larger 0.9 μm, and 2.1×¹⁰⁻² m⁻³ and 7.6×¹⁰⁻² m⁻³ for grains larger than 1.6 μm. Our data imply that the densest point is displaced outwards by at least with respect of the Enceladus orbit. This finding provides direct evidence for plume particles dragged outwards by the ambient plasma. The differential size distribution for grains >0.9 μm is described best by a power law with slopes between 4 and 5. We also obtained dust data during ring plane crossings in the vicinity of the orbits of Mimas and Tethys. The vertical distribution of grains >0.8 μm at Mimas orbit is also well described by Gaussian with a FWHM of ∼5400 km and displaced southwards by ∼1200 km with respect to the geometrical equator. The vertical distribution of ring particles in the vicinity of Tethys, however, does not match a Gaussian. We use the FWHM values obtained from the vertical crossings to establish a 2-dimensional model for the ring particle distribution which matches our observations during vertical and equatorial traversals through the E ring.     

Obliquity tides do not significantly heat Enceladus

Recently, Tyler [Tyler, R.H., 2009. Geophys. Res. Lett. 36, L15205; Tyler, R., 2011. Icarus, 211, 770-779] proposed that the tide due to an obliquity of greater than 0.1° might drive resonant flow in a liquid ocean at Enceladus, and that dissipation of the ocean’s kinetic energy may be an alternate source for the observed global heat flux. While there is currently no measurement of Enceladus’ obliquity, dissipation is expected to drive the spin pole to a Cassini state. Under this assumption, we find that Enceladus should occupy Cassini state 1 and that the obliquity of Enceladus should be less than 0.0015° for values of the degree-2 gravity coefficient C_2,2 between 1.0 × 10⁻³ and 2.5 × 10⁻³. Unless there is a significant free obliquity or the gravity coefficient C_2,2 has been significantly overestimated, it is unlikely that obliquity-driven flow in a subsurface ocean is the source of the extreme heat on Enceladus.     

New laboratory measurements of CH4 in Titan's conditions and a reanalysis of the DISR near-surface spectra at the Huygens landing site

Laboratory spectra of methane-nitrogen mixtures have been recorded in the near-infrared range (1.0-1.65 μm) in conditions similar to Titan's near surface, to facilitate the interpretation of the DISR/DLIS (DISR—Descent Imager/Spectral Radiometer) spectra taken during the last phase of the descent of the Huygens Probe, when the surface was illuminated by a surface-science lamp. We used a 0.03 cm⁻¹ spectral resolution, adequate to resolve the lines at high pressure (p_N2∼1.5 bar). By comparing the laboratory spectra with synthetic calculations in the well-studied ν₂+2ν₃ band (7515-7620 cm⁻¹), we determine a methane absorption column density of 178±20 cm atm and a temperature of 118±10 K in our experiment. From this, we derive the methane absorption coefficients over 1.0-1.65 μm with a 0.03 cm⁻¹ sampling, allowing for the extrapolation of the results to any other methane column density under the relevant pressure and temperature conditions. We then revisit the calibration and analysis of the Titan “lamp-on” DLIS spectra. We infer a 5.1±0.8% methane-mixing ratio in the first 25 m of Titan's atmosphere. The CH₄ mixing ratio measured 90 s after landing from a distance of 45 cm is found to be 0.92±0.25 times this value, thus showing no post-landing outgassing of methane in excess of ∼20%. Finally, we determine the surface reflectivity as seen between 25 m and 45 cm and find that the 1500 nm absorption band is deeper in the post-landing spectrum as compared to pre-landing.     

Cassini RPWS observations of dust in Saturn's E Ring

The Cassini radio and plasma wave science (RPWS) instrument is sensitive to few-micron dust grains impacting on the spacecraft at relative speeds of order 10 km/s. Through the first year or so of operations in orbit at Saturn, the RPWS has made a number of both inclined and equatorial crossings of the E ring, particularly near the orbit of Enceladus. Assuming water ice grains, the typical size particle detected by the RPWS has a radius of a few microns. Peak impact rates of about 50 s⁻¹ are found near the orbit of Enceladus corresponding to densities of order 5×10⁻⁴ m⁻³. The variation of dust fluxes as a function of height above or below the equator is well described by a Gaussian distribution with a scale height of about 2800 km although there is usually some non-Gaussian variation near the peak fluxes suggesting some structure in the core of the ring. Offsets of the peak number densities are typically of the order of a few hundred km from the geometric equator. A near-equatorial radial profile through the orbit of Enceladus shows a sharply peaked distribution at the orbit of the moon. A size distribution averaged over several passes through the orbit of Enceladus is determined which varies as m^−2.80. The peak in dust number density at the orbit of Enceladus is consistent with previous optical measurements and strongly supports the suggestion that Enceladus is a primary source for E ring particles.     

Dust devil sediment flux on Earth and Mars: Laboratory simulations

Laboratory simulations using the Arizona State University Vortex Generator (ASUVG) were run to simulate sediment flux in dust devils in terrestrial ambient and Mars-analog conditions. The objective of this study was to measure vortex sediment flux in the laboratory to yield estimations of natural dust devils on Earth and Mars, where all parameters may not be measured. These tests used particles ranging from 2 to 2000 μm in diameter and 1300 to 4800 kg m⁻³ in density, and the results were compared with data from natural dust devils on Earth and Mars. Typically, the cores of dust devils (regardless of planetary environment) have a pressure decrease of ∼0.1-1.5% of ambient atmospheric pressure, which enhances the lifting of particles from the surface. Core pressure decreases in our experiments ranged from ∼0.01% to 5.00% of ambient pressure (10 mbar Mars cases and 1000 mbar for Earth cases) corresponding to a few tenths of a millibar for Mars cases and a few millibars for Earth cases. Sediment flux experiments were run at vortex tangential wind velocities of 1-45 m s⁻¹, which typically correspond to ∼30-70% above vortex threshold values for the test particle sizes and densities. Sediment flux was determined by time-averaged measurements of mass loss for a given vortex size. Sediment fluxes of ∼10⁻⁶-10⁰ kg m⁻² s⁻¹ were obtained, similar to estimates and measurements for fluxes in dust devils on Earth and Mars. Sediment flux is closely related to the vortex intensity, which depends on the strength of the pressure decrease in the core (ΔP). This study found vortex size is less important for lifting materials because many different diameters can have the same ΔP. This finding is critical in scaling the laboratory results to natural dust devils that can be several orders of magnitude larger than the laboratory counterparts.     

In situ measurements of particle load and transport in dust devils

In situ (mobile) sampling of 33 natural dust devil vortices reveals very high total suspended particle (TSP) mean values of 296 mg m⁻³ and fine dust loadings (PM10) mean values ranging from 15.1 to 43.8 mg m⁻³ (milligrams per cubic meter). Concurrent three-dimensional wind profiles show mean tangential rotation of 12.3 m s⁻¹ and vertical uplift of 2.7 m s⁻¹ driving mean vertical TSP flux of 1689 mg m⁻³ s⁻¹ and fine particle flux of ∼1.0 to ∼50 mg m⁻³ s⁻¹. Peak PM10 dust loading and flux within the dust column are three times greater than mean values, suggesting previous estimates of dust devil flux might be too high. We find that deflation rates caused by dust devil erosion are ∼2.5-50 μm per year in dust devil active zones on Earth. Similar values are expected for Mars, and may be more significant there where competing erosional mechanisms are less likely.     

Trapping and dynamical evolution of interplanetary dust particles in Earth’s quasi-satellite resonance

We used numerical simulations to model the orbital evolution of interplanetary dust particles (IDPs) evolving inward past Earth’s orbit under the influence of radiation pressure, Poynting–Robertson light drag (PR drag), solar wind drag, and gravitational perturbations from the planets. A series of β values (where β is the ratio of the force from radiation pressure to that of central gravity) were used ranging from 0.0025 up to 0.02. Assuming a composition consistent with astronomical silicate and a particle density of 2.5 g cm⁻³ these β values correspond to dust particle diameters ranging from 200 μm down to 25 μm. As the dust particle orbits decay past 1 AU between 4% (for β = 0.02, or 25 μm) and 40% (for β = 0.0025, or 200 μm) of the population became trapped in 1:1 co-orbital resonance with Earth. In addition to traditional horseshoe type co-orbitals, we found about a quarter of the co-orbital IDPs became trapped as so-called quasi-satellites. Quasi-satellite IDPs always remain relatively near to Earth (within 0.1–0.3 AU, or 10–30 Hill radii, R_H) and undergo two close-encounters with Earth each year. While resonant perturbations from Earth halt the decay in semi-major axis of quasi-satellite IDPs their orbital eccentricities continue to decrease under the influence of PR drag and solar wind drag, forcing the IDPs onto more Earth-like orbits. This has dramatic consequences for the relative velocity and distance of closest approach between Earth and the quasi-satellite IDPs. After 10⁴–10⁵ years in the quasi-satellite resonance dust particles are typically less than 10R_H from Earth and consistently coming within about 3R_H. In the late stages of evolution, as the dust particles are escaping the 1:1 resonance, quasi-satellite IDPs can have deep close-encounters with Earth significantly below R_H. Removing the effects of Earth’s gravitational acceleration reveals that encounter velocities (i.e., velocities “at infinity”) between quasi-satellite IDPs and Earth during these close-encounters are just a few hundred meters per second or slower, well below the average values of 2–4 km s⁻¹ for non-resonant Earth-crossing IDPs with similar initial orbits. These low encounter velocities lead to a factor of 10–100 increase in Earth’s gravitationally enhanced impact cross-section (σ_grav) for quasi-satellite IDPs compared to similar non-resonant IDPs. The enhancement in σ_grav between quasi-satellite IDPs and cometary Earth-crossing IDPs is even more pronounced, favoring accretion of quasi-satellite dust particles by a factor of 100–3000 over the cometary IDPs. This suggests that quasi-satellite dust particles may dominate the flux of large (25–200 μm) IDPs entering Earth’s atmosphere. Furthermore, because quasi-satellite trapping is known to be directly correlated with the host planet’s orbital eccentricity the accretion of quasi-satellite dust likely ebbs and flows on 10⁵ year time scales synchronized with Earth’s orbital evolution.     

Mid-infrared reflectance spectra and optical constants of six iron oxide/oxyhydroxide phases

We have determined the real and imaginary indices of refraction (n and k) for six iron oxide/oxyhydroxide phases—magnetite, maghemite, goethite, lepidocrocite, akaganéite, and ferrihydrite. A single crystal of magnetite was used to derive bulk n and k values from 100-2000 cm⁻¹ (5-100 μm). Synthetic nanocrystalline samples of maghemite, goethite, lepidocrocite, akaganéite, and ferrihydrite were pressed into compact pellets used to determine bulk n and k values from 100-1200 cm⁻¹ (8.33-100 μm). All values of n and k (the optical constants) were determined from specular reflectance spectra acquired at 2 cm⁻¹ spectral sampling using classical Lorentz-Lorenz dispersion theory. In this paper, we present the optical constants of all six minerals and the oscillator parameters with which they were modeled. Use of these optical constants could aid in radiative transfer models of terrestrial dust as well as Mars, the Moon, and airless bodies in the Solar System.     

Dust production from the hypervelocity impact disruption of the Murchison hydrous CM2 meteorite: Implications for the disruption of hydrous asteroids and the production of interplanetary dust

More than half of the C-type asteroids, the dominant type of asteroid in the outer half of the main-belt, show evidence of hydration in their reflectance spectra. In order to understand the collisional evolution of asteroids and the production of interplanetary dust and to model the infrared signature of small particles in the Solar System it is important to characterize the dust production from primary impact disruption events, and compare the disruption of hydrous and anhydrous targets. We performed a hypervelocity impact disruption experiment on an ∼30 g target of the Murchison CM2 hydrated carbonaceous chondrite meteorite, and compared the results with our previous disruption experiments on anhydrous meteorites including Allende, a CV3 carbonaceous chondrite, and nine ordinary chondrites. Murchison is significantly more friable than the ordinary chondrites or Allende. Nonetheless, on a plot of mass of the largest fragment versus specific impact energy, the Murchison disruption plots within the field of the anhydrous meteorites points, suggesting that Murchison is at least as resistant to impact disruption as the anhydrous meteorites, which require about twice the energy for disruption as terrestrial anhydrous basalt targets. We determined the mass-frequency distribution of the debris from the Murchison disruption over a nine order-of-magnitude mass range, from ∼10⁻⁹ g to the mass of the largest fragment produced in the disruption. The cumulative mass-frequency distribution from the Murchison disruption is fit by three power-law segments. For masses >10⁻² g the slope is only slightly steeper than that of the corresponding segment from the disruption of most anhydrous meteorites. Over the range from ∼10⁻⁶ to 10⁻² g the slope is significantly steeper than that for the anhydrous meteorites. For masses <10⁻⁶ g the slopes of both the Murchison and the anhydrous meteorites are almost flat. Thus the Murchison disruption significantly over-produced small fragments (10⁻⁶-10⁻³ g) compared to anhydrous meteorite targets. If the Murchison results are representative of hydrous asteroids, the hydrous asteroids may dominate over anhydrous asteroids in the production of interplanetary dust >100 μm in size, the size of micrometeorites recovered from the polar ices, while both types of asteroids might produce comparable amounts of ∼10 μm interplanetary dust. This would explain the puzzle that polar micrometeorites (>100 μm in size) are similar to hydrous meteorites, while the majority of the ∼10 μm interplanetary dust particles are anhydrous.     

Apollo video photogrammetry estimation of plume impingement effects

Future missions to the Moon may require numerous landings at the same site. Since the top few centimeters are loosely packed regolith, plume impingement from the Lander ejects the granular material at high velocities. Much work is needed to understand the physics of plume impingement during landing to protect hardware surrounding the landing sites. While mostly qualitative in nature, the Apollo Lunar Module landing videos can provide a wealth of quantitative information using modern photogrammetry techniques. The authors have used the digitized videos to quantify plume impingement effects of the landing exhaust on the lunar surface. The dust ejection angle from the plume is estimated at 1°-3°. The lofted particle density is estimated at 10⁸-10¹³ particles/m³. Additionally, evidence for ejection of large 10-15 cm sized objects and a dependence of ejection angle on thrust are presented. Further work is ongoing to continue quantitative analysis of the landing videos.     

Energy balance of the Deep Impact experiment

We present results on the energy balance of the Deep Impact experiment based on analysis of 180 infrared spectra of the ejecta obtained by the Deep Impact spacecraft. We derive an output energy of 16.5 (+9.1/−4.1) GJ. With an input energy of 19.7 GJ, the error bars are large enough so that there may or may not be a balance between the kinetic energy of the impact and that of outflowing materials. Although possible, no other source of energy other than the impactor or the Sun is needed to explain the observations. Most of the energy (85%) goes into the hot plume in the first few seconds, which only represents a very small fraction (<0.01%) of the total ejected mass. The hot plume contains 190 (+263/−71) kg of H₂O, 1.6 ± 0.5 kg of CO₂, 8.2 (+11.3/3.1) kg of CO (assuming a CO/H₂O ratio of 4.3%), 27.9 (+25.0/−8.9) kg of organic material and 255 ± 128 kg of dust, while the ejecta contains ∼10⁷ kg of materials. About 12% of the energy goes into the ejecta (mostly water) and 3% to destroy the impactor. Volatiles species other than H₂O (CO₂, CO or organic molecules) contribute to <7% of the energy balance. In terms of physical processes, 68% of the energy is used to accelerate grains (kinetic energy), 16% to heat them, 6% to sublimate or melt them and 10% (upper limit) to break and compress dust and/or water ice aggregates into small micron size particles. For the hot plume, we derive a dust/H₂O ratio of 1.3 (+1.9/−1.0), a CO₂/H₂O ratio of 0.008 (+0.009/−0.006), an organics/H₂O ratio of 0.15 (+0.29/−0.11) and an organics/dust ratio of 0.11 (+0.30/−0.07). This composition refers to the impact site and is different from that of the bulk nucleus, consistent with the idea of layers of different composition in the nucleus sub-surface. Our results emphasize the importance of laboratory impact experiments to understand the physical processes involved at such a large scale.     

On the origin of south polar folds on Enceladus

Recent high-resolution Cassini images of the south polar terrain of Enceladus reveal regions of short-wavelength deformation, inferred to be compressional folds between the Baghdad and Damascus tiger stripes (Spencer, J.R., Barr, A.C., Esposito, L.W., Helfenstein, P., Ingersoll, A.P., Jaumann, R., McKay, C.P., Nimmo, F., Waite, J.H. [2009a]. Enceladus: An active cryovolcanic satellite. In: Saturn after Cassini-Huygens. Springer, New York, pp. 683-722). Here, we use Fourier analysis of the bright/dark variations to show that the folds have a dominant wavelength of 1.1 ± 0.4 km. We use the simple model of lava flow folding from Fink (Fink, J. [1980]. Geology 8, 250-254) to show that the folds could form in an ice shell with an upper high-viscosity boundary layer of thickness <400 m, with a driving stress of 40-80 kPa, and strain rate between 10⁻¹⁴ s⁻¹ and 10⁻¹² s⁻¹. Such deformation rates imply resurfacing of the SPT in 0.05-5 Myr, consistent with its estimated surface age. Measurements of fold topography and more sophisticated numerical modeling can narrow down the conditions of fold formation and provide valuable constraints on the thermal structure of the ice shell on Enceladus.     

Keck and VLT AO observations and models of the uranian rings during the 2007 ring plane crossings

We present observations of the uranian ring system at a wavelength of 2.2 μm, taken between 2003 and 2008 with NIRC2 on the W.M. Keck telescope in Hawaii, and on 15–17 August 2007 with NaCo on the Very Large Telescope (VLT) in Chile. Of particular interest are the data taken around the time of the uranian ring plane crossing with Earth on 16 August 2007, and with the Sun (equinox) on 7 December 2007. We model the data at the different viewing aspects with a Monte Carlo model to determine: (1) the normal optical depth τ₀, the location, and the radial extent of the main rings, and (2) the parameter Aτ₀ (A is the particle geometric albedo), the location, and the radial plus vertical extent of the dusty rings. Our main conclusions are: (i) The brightness of the ? ring is significantly enhanced at small phase and ring inclination angles; we suggest this extreme opposition effect to probably be dominated by a reduction in interparticle shadowing. (ii) A broad sheet of dust particles extends inwards from the λ ring almost to the planet itself. This dust sheet has a vertical extent of ∼140 km, and Aτ₀ = 2.2 × 10⁻⁶. (iii) The dusty rings between ring 4 and the α ring and between the α and β rings are vertically extended with a thickness of ∼300 km. (iv) The ζ ring extends from ∼41,350 km almost all the way inwards to the planet. The main ζ ring, centered at ∼39,500 km from the planet, is characterized by Aτ₀ = 3.7 × 10⁻⁶; this parameter decreases closer to the planet. The ζ ring has a full vertical extent of order 800–900 km, with a pronounced density enhancement in the mid-plane. (v) The η_c ring is optically thin and less than several tens of km in the vertical direction. This ring may be composed of macroscopic material, surrounded by clumps of dust.     

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.     

Waves in Cassini UVIS stellar occultations: 2. The C ring

We performed a complete wavelet analysis of Saturn’s C ring on 62 stellar occultation profiles. These profiles were obtained by Cassini’s Ultraviolet Imaging Spectrograph High Speed Photometer. We used a WWZ wavelet power transform to analyze them. With a co-adding process, we found evidence of 40 wavelike structures, 18 of which are reported here for the first time. Seventeen of these appear to be propagating waves (wavelength changing systematically with distance from Saturn). The longest new wavetrain in the C ring is a 52-km-long wave in a plateau at 86,397 km. We produced a complete map of resonances with external satellites and possible structures rotating with Saturn’s rotation period up to the eighth order, allowing us to associate a previously observed wave with the Atlas 2:1 inner Lindblad resonance (ILR) and newly detected waves with the Mimas 6:2 ILR and the Pandora 4:2 ILR. We derived surface mass densities and mass extinction coefficients, finding σ = 0.22(±0.03) g cm⁻² for the Atlas 2:1 ILR, σ = 1.31(±0.20) g cm⁻² for the Mimas 6:2 ILR, and σ = 1.42(±0.21) g cm⁻² for the Pandora 4:2 ILR. We determined a range of mass extinction coefficients (κ = τ/σ) for the waves associated with resonances with κ = 0.13 (±0.03) to 0.28(±0.06) cm² g⁻¹, where τ is the optical depth. These values are higher than the reported values for the A ring (0.01-0.02 cm² g⁻¹) and the Cassini Division (0.07-0.12 cm² g⁻¹ from Colwell et al. (Colwell, J.E., Cooney, J.H., Esposito, L.W., Srem?evi?, M. [2009]. Icarus 200, 574-580)). We also note that the mass extinction coefficient is probably not constant across the C ring (in contrast to the A ring and the Cassini Division): it is systematically higher in the plateaus than elsewhere, suggesting smaller particles in the plateaus. We present the results of our analysis of these waves in the C ring and estimate the mass of the C ring to be between3.7(±0.9) × 10¹⁶ kg and 7.9(±2.0) × 10¹⁶ kg (equivalent to an icy satellite of radius between 28.0(±2.3) km and 36.2(±3.0) km with a density of 400 kg m⁻³, close to that of Pan or Atlas). Using the ring viscosity derived from the wave damping length, we also estimate the vertical thickness of the C ring between 1.9(±0.4) m and 5.6(±1.4) m, comparable to the vertical thickness of the Cassini Division.     

Influence of high abundances of aerosols on the electrical conductivity of the Titan atmosphere

Observations of optical depth and scattering by instrumentation onboard the Huygens probe have been used by Tomasko et al. [Tomasko et al., 2005. Rain, winds and haze during Huygens probe's descent to Titan's surface. Nature 438 (8), 765-778] to deduce that the size and abundance of Titan aerosols could be nearly independent of altitude. Here we show that by assuming a constant mass flux with altitude and using the measured optical depth as a constraint, we obtain more realistic size and abundance distributions. In particular, the calculated abundance decreases from 3.5×10⁷ m⁻³ at 100 km to 8×10⁶ m⁻³ near the surface while the particle radius varies from 0.25 μm at 150 km to 1.1 μm at the surface. These distributions are consistent with the reported measurements for these quantities. Our results are then employed to compute electron and ion densities and conductivities for various solar UV photoelectron emission thresholds. Our model shows that to get agreement with the published (preliminary) conductivity measurements, photoemission cannot be an important source of electrons and ions. To get agreement with the electron and ion conductivity observations, both an additional population of aerosol embryos above 50 km and a trace amount of an electrophillic molecular species below 50 km are needed.     

Dust extinction and X-ray emission from the starburst galaxy NGC 1482

We present the results based on multiwavelength imaging observations of the prominent dust lane starburst galaxy NGC 1482 aimed to investigate the extinction properties of dust existing in the extreme environment. (B-V) colour-index map derived for the starburst galaxy NGC 1482 confirms two prominent dust lanes running along its optical major axis and are found to extend up to ∼11 kpc. In addition to the main lanes, several filamentary structures of dust originating from the central starburst are also evident. Though, the dust is surrounded by exotic environment, the average extinction curve derived for this target galaxy is compatible with the Galactic curve, with R_V = 3.05, and imply that the dust grains responsible for the optical extinction in the target galaxy are not really different than the canonical grains in the Milky Way. Our estimate of total dust content of NGC 1482 assuming screening effect of dust is ∼2.7 × 10⁵ M_⊙, and provide lower limit due to the fact that our method is not sensitive to the intermix component of dust. Comparison of the observed dust in the galaxy with that supplied by the SNe to the ISM, imply that this supply is not sufficient to account for the observed dust and hence point towards the origin of dust in this galaxy through a merger like event.Our multiband imaging analysis reveals a qualitative physical correspondence between the morphologies of the dust and Hα emission lines as well as diffuse X-ray emission in this galaxy. Spatially resolved spectral analysis of the hot gas along outflows exhibit a gradient in the temperature. Similar gradient was also noticed in the measured values of metallicity, indicating that the gas in the halo is not yet enriched. High resolution, 2-8 keV Chandra image reveals a pair of point sources in the nuclear region with their luminosities equal to 2.27 × 10³⁹ erg s⁻¹ and 9.34 × 10³⁹ erg s⁻¹, and are in excess of the Eddington-limit of 1.5 M_⊙ accreting source. Spectral analysis of these sources exhibit an absorbed-power law with the hydrogen column density higher than that derived from the optical measurements.