The atmospheres of Uranus and Neptune: A review

Plausible models for the atmospheres of Uranus and Neptune are reviewed. Current ideas favor the presence of massive atmospheres above solid cores. Observations of Uranus imply the presence of a visible cloud layer (probably composed of solid methane) beneath about 100 km amagats of hydrogen. A number of other cloud layers far below this upper layer are also possible. However, before any of these conclusions can be considered firm, a number of further crucial observations are required.     

The internal structures and the relative rotation rates of Uranus and Neptune

The difference between the interior structures of Uranus and Neptune is presented, based on models which fit the observed mass, radius, and gravitational moments for the assumed rotation periods of these planets. If Uranus and Neptune are assumed to be as similar in internal structure as they are in mass and radius, the rotation period for Neptune must be shorter than that for Uranus. It is suggested that the true rotation period is given by Neptune's oblateness, while the photometric period corresponds to the motion of Rossby waves in the upper atmosphere.     

New observational constraints on the temperature inversions of Uranus and Neptune

We present 20-μm photometry of Uranus and Neptune which confirms the presence of a temperature inversion in the lower stratospheres in both planets. We find the brightness temperature difference between 17.8 and 19.6 μm to be 0.8 ± 0.5°K for Uranus and 1.8 ± 0.6°K for Neptune. These results indicate that the temperature inversions on both planets are weaker than previously thought. Comparison to model atmospheres by J. Appleby [Ph.D. thesis, SUNY at Stony Brook 1980] indicates that the temperature inversions can be understood as arising from heating by the absorption of sunlight by CH₄ and aerosols. However, the stratospheric CH₄ mixing ratio on Neptune must be higher than that at the temperature minimum.     

The phase diagram of water and the magnetic fields of Uranus and Neptune

The interior of giant planets can give valuable information on formation and evolution processes of planetary systems. However, the interior and evolution of Uranus and Neptune is still largely unknown. In this paper, we compare water-rich three-layer structure models of these planets with predictions of shell structures derived from magnetic field models. Uranus and Neptune have unusual non-dipolar magnetic fields contrary to that of the Earth. Extensive three-dimensional simulations of Stanley and Bloxham (Stanley, S., Bloxham, J. [2004]. Nature 428, 151-153) have indicated that such a magnetic field is generated in a rather thin shell of at most 0.3 planetary radii located below the H/He rich outer envelope and a conducting core that is fluid but stably stratified. Interior models rely on equation of state data for the planetary materials which have usually considerable uncertainties in the high-pressure domain. We present interior models for Uranus and Neptune that are based on ab initio equation of state data for hydrogen, helium, and water as the representative of all heavier elements or ices. Based on a detailed high-pressure phase diagram of water we can specify the region where superionic water should occur in the inner envelope. This superionic region correlates well with the location of the stably-stratified region as found in the dynamo models. Hence we suggest a significant impact of the phase diagram of water on the generation of the magnetic fields in Uranus and Neptune.     

The formation of Uranus and Neptune in solid-rich feeding zones: Connecting chemistry and dynamics

The core accretion theory of planet formation has at least two fundamental problems explaining the origins of Uranus and Neptune: (1) dynamical times in the trans-saturnian solar nebula are so long that core growth can take >15 Myr and (2) the onset of runaway gas accretion that begins when cores reach ∼10M_⊕ necessitates a sudden gas accretion cutoff just as Uranus and Neptune’s cores reach critical mass. Both problems may be resolved by allowing the ice giants to migrate outward after their formation in solid-rich feeding zones with planetesimal surface densities well above the minimum-mass solar nebula. We present new simulations of the formation of Uranus and Neptune in the solid-rich disk of Dodson-Robinson et al. (Dodson-Robinson, S.E., Willacy, K., Bodenheimer, P., Turner, N.J., Beichman, C.A. [2009]. Icarus 200, 672-693) using the initial semimajor axis distribution of the Nice model (Gomes, R., Levison, H.F., Tsiganis, K., Morbidelli, A. [2005]. Nature 435, 466-469; Morbidelli, A., Levison, H.F., Tsiganis, K., Gomes, R. [2005]. Nature 435, 462-465; Tsiganis, K., Gomes, R., Morbidelli, A., Levison, H.F. [2005]. Nature 435, 459-461), with one ice giant forming at 12 AU and the other at 15 AU. The innermost ice giant reaches its present mass after 3.8-4.0 Myr and the outermost after 5.3-6 Myr, a considerable time decrease from previous one-dimensional simulations (e.g. Pollack, J.B., Hubickyj, O., Bodenheimer, P., Lissauer, J.J., Podolak, M., Greenzweig, Y. [1996]. Icarus 124, 62-85). The core masses stay subcritical, eliminating the need for a sudden gas accretion cutoff.Our calculated carbon mass fractions of 22% are in excellent agreement with the ice giant interior models of Podolak et al. (Podolak, M., Weizman, A., Marley, M. [1995]. Planet. Space Sci. 43, 1517-1522) and Marley et al. (Marley, M.S., Gómez, P., Podolak, M. [1995]. J. Geophys. Res. 100, 23349-23354). Based on the requirement that the ice giant-forming planetesimals contain >10% mass fractions of methane ice, we can reject any Solar System formation model that initially places Uranus and Neptune inside of Saturn’s orbit. We also demonstrate that a large population of planetesimals must be present in both ice giant feeding zones throughout the lifetime of the gaseous nebula. This research marks a substantial step forward in connecting both the dynamical and chemical aspects of planet formation. Although we cannot say that the solid-rich solar nebula model of Dodson-Robinson et al. (Dodson-Robinson, S.E., Willacy, K., Bodenheimer, P., Turner, N.J., Beichman, C.A. [2009]. Icarus 200, 672-693) gives exactly the appropriate initial conditions for planet formation, rigorous chemical and dynamical tests have at least revealed it to be a viable model of the early Solar System.     

The solar radiation incident at the top of the atmospheres of Uranus and Neptune

The latitudinal and seasonal variation of the direct solar radiation incident at the top of the atmosphere of Uranus and Neptune has been recalculated by use of updated values for the period of axial rotation and the oblateness. Values for the solar radiation are given in Watt per square meter instead of the unit used in earlier papers (calories per square centimeter per planetary day). The solar radiation averaged over a season and a year as a function of planetocentric latitude has also been reviewed. In addition, attention is made to the ratio of the solar radiation incident on an oblate planet to that incident on a spherical planet.     

Measurements of the H2 4-0 quadrupole bands of Uranus and Neptune

We find that the equivalent widths of the lines of the 4-0 H₂ quadrupole band on Uranus and Neptune are substantially smaller than the values found by some previous observers. An analysis of our results based on a range of atmospheric models yields H₂ abundances of 240 ± 60 km-amagats for Uranus and ?200 km amagats for Neptune.     

Can magnetic fields be generated in the icy mantles of Uranus and Neptune?

It has been shown by us previously that a hydromagnetic dynamo can operate in the core of Uranus but probably not on Neptune. A similar analysis is made for the “icy” liquid mantles of both planets. It is concluded that pressure ionization and the associated increased conductivity of water is probably not enough to satisfy the necessary conditions for a dynamo on Uranus and that it is marginal for Neptune. On the other hand the expected presence of metallic water in a thick layer around the core of Neptune makes the operation of a dynamo on this planet plausible. A similar layer on Uranus might be too thin to play the same role. It appears that if a magnetic field is indeed present on Uranus it is probably generated in the core of the planet, while on Neptune it is more likely operating in the icy mantle.     

Polarization studies of the Venus uv contrasts: Cloud height and haze variability

Polarimetry is able to show direct evidence for compositional differences in the Venus clouds. We present observations (collected during $\text{2}\text{1}\text{2}$ Venus years by the Pioneer Venus Orbiter) of the polarization in four colors of the bright and dark ultraviolet features. We find that the polarization is significantly different between the bright and dark areas. The data show that the “null” model of L. W. Esposito (1980, J. Geophys. Res.85, 8151–8157) and the “overlying haze” model of J. B. Pollack et al. (1980, J. Geophys. Res.85, 8223–8231) are insufficient. Exact calculations of the polarization, including multiple scattering and vertical inhomogeneity near the Venus cloud tops, are able to match the observations. Our results give a straightforward interpretation of the polarization differences in terms of known constituents of the Venus atmosphere. The submicron haze and uv absorbers are anticorrelated: for haze properties as given by K. Kawabata et al. (1980, J. Geophys. Res.85, 8129–8140) the excess haze depth at 9350 Å over the bright regions is Δτ_h = 0.03 ± 0.02. The cloud top is slightly lower in the dark features: the extra optical depth at 2700 Å in Rayleigh scattering above the darker areas is Δτ_R = 0.010 ± 0.005. This corresponds to a height difference of 1.2 ± 0.6 km at the cloud tops. The calculated polarization which matches our data also explains the relative polarization of bright and dark features observed by Mariner 10. The observed differential polarization cannot be explained by differential distribution of haze, if the haze aerosols have an effective size of 0.49 μm, as determined by K. Kawabata et al. (1982, submitted) for the aerosols overlying the Venus equator. We propose two models for the uv contrasts consistent with our results. In a physical model, the dark uv regions are locations of vertical convergence and horizontal divergence. In a chemical model, we propose that the photochemistry is limited by local variations in water vapor and molecular oxygen. The portions of the atmosphere where these constituents are depleted at the cloud tops are the dark uv features. Strong support for this chemical explanation is the observation that the number of sulfur atoms above the cloud tops is equal over both the bright and dark areas. The mass budget of sulfur at these altitudes is balanced between excess sulfuric acid haze over the bright regions and excess SO₂ in the dark regions.     

Some dynamical aspects of the accretion of Uranus and Neptune: The exchange of orbital angular momentum with planetesimals

The final stage of the accretion of Uranus and Neptune is numerically investigated. The four Jovian planets are considered with Jupiter and Saturn assumed to have reached their present sizes, whereas Uranus and Neptune are taken with initial masses 0.2 of their present ones. Allowance is made for the orbital variation of the Jovian planets due to the exchange of angular momentum with interacting bodies (“planetesimals”). Two possible effects that may have contributed to the accretion of Uranus and Neptune are incorporated in our model: (1) an enlarged cross section for accretion of incoming planetesimals due to the presence of extended gaseous envelopes and/or circumplanetary swarms of bodies; and (2) intermediate protoplanets in mid-range orbits between the Jovian planets. Significant radial displacements are found for Uranus and Neptune during their accretion and scattering of planetesimals. The orbital angular momentum budgets of Neptune, Uranus, and Saturn turn out to be positive; i.e., they on average gain orbital angular momentum in their interactions with planetesimals and hence they are displaced outwardly. Instead, Jupiter as the main ejector of bodies loses orbital angular momentum so it moves sunward. The gravitational stirring of planetesimals caused by the introduction of intermediate protoplanets has the effect that additional solid matter is injected into the accretion zones of Uranus and Neptune. For moderate enlargements of the radius of the accretion cross section (2–4 times), the accretion time scale of Uranus and Neptune are found to be a few 10⁸ years and the initial amount of solid material required to form them of a few times their present masses. Given the crucial role played by the size of the accretion cross section, questions as to when Uranus and Neptune acquired their gaseous envelopes, when the envelopes collapsed onto the solid cores, and how massive they were are essential in defining the efficiency and time scale of accretion of the two outer Jovian planets.     

The elimination of the 1:2 critical terms of a first order theory of Uranus perturbed by Neptune through Von Zeipel method

In this paper we eliminate in a first order U-N theory the 1 : 2 critical terms up to the third degree with respect to eccentricity — inclination in both parts, main and indirect of the U-N planetary Hamiltonian. We operate the Von Zeipel technique. We adopt, in this theory, the Jacobi-Radau coordinates, and the Poincaré canonical variables. We neglect powers higher than the third in the eccentricity-inclination. This paper is related to the two previous articles (Kamel, 1982; 1983).     

Origin of the structure of the Kuiper belt during a dynamical instability in the orbits of Uranus and Neptune

We explore the origin and orbital evolution of the Kuiper belt in the framework of a recent model of the dynamical evolution of the giant planets, sometimes known as the Nice model. This model is characterized by a short, but violent, instability phase, during which the planets were on large eccentricity orbits. It successfully explains, for the first time, the current orbital architecture of the giant planets [Tsiganis, K., Gomes, R., Morbidelli, A., Levison, H.F., 2005. Nature 435, 459-461], the existence of the Trojans populations of Jupiter and Neptune [Morbidelli, A., Levison, H.F., Tsiganis, K., Gomes, R., 2005. Nature 435, 462-465], and the origin of the late heavy bombardment of the terrestrial planets [Gomes, R., Levison, H.F., Tsiganis, K., Morbidelli, A., 2005. Nature 435, 466-469]. One characteristic of this model is that the proto-planetary disk must have been truncated at roughly 30 to 35 AU so that Neptune would stop migrating at its currently observed location. As a result, the Kuiper belt would have initially been empty. In this paper we present a new dynamical mechanism which can deliver objects from the region interior to ∼35 AU to the Kuiper belt without excessive inclination excitation. In particular, we show that during the phase when Neptune's eccentricity is large, the region interior to its 1:2 mean motion resonance becomes unstable and disk particles can diffuse into this area. In addition, we perform numerical simulations where the planets are forced to evolve using fictitious analytic forces, in a way consistent with the direct N-body simulations of the Nice model. Assuming that the last encounter with Uranus delivered Neptune onto a low-inclination orbit with a semi-major axis of ∼27 AU and an eccentricity of ∼0.3, and that subsequently Neptune's eccentricity damped in ∼1 My, our simulations reproduce the main observed properties of the Kuiper belt at an unprecedented level. In particular, our results explain, at least qualitatively: (1) the co-existence of resonant and non-resonant populations, (2) the eccentricity-inclination distribution of the Plutinos, (3) the peculiar semi-major axis—eccentricity distribution in the classical belt, (4) the outer edge at the 1:2 mean motion resonance with Neptune, (5) the bi-modal inclination distribution of the classical population, (6) the correlations between inclination and physical properties in the classical Kuiper belt, and (7) the existence of the so-called extended scattered disk. Nevertheless, we observe in the simulations a deficit of nearly-circular objects in the classical Kuiper belt.     

Accurate and approximate calculations of Raman scattering in the atmosphere of Neptune

Raman scattering by H₂ in Neptune's atmosphere has significant effects on its reflectivity for λ<0.5 μm, producing baseline decreases of ∼20% in a clear atmosphere and ∼10% in a hazy atmosphere. However, few accurate Raman calculations are carried out because of their complexity and computational costs. Here we present the first radiation transfer algorithm that includes both polarization and Raman scattering and facilitates computation of spatially resolved spectra. New calculations show that Cochran and Trafton's (1978, Astrophys. J. 219, 756-762) suggestion that light reflected in the deep CH₄ bands is mainly Raman scattered is not valid for current estimates of the CH₄ vertical distribution, which implies only a 4% Raman contribution. Comparisons with IUE, HST, and groundbased observations confirm that high altitude haze absorption is reducing Neptune's geometric albedo by ∼6% in the 0.22-0.26 μm range and by ∼13% in the 0.35-0.45 μm range. A sample haze model with 0.2 optical depths of 0.2-μm radius particles between 0.1 and 0.8 bars fits reasonably well, but is not a unique solution. We used accurate calculations to evaluate several approximations of Raman scattering. The Karkoschka (1994, Icarus 111, 174-192) method of applying Raman corrections to calculated spectra and removing Raman effects from observed spectra is shown to have limited applicability and to undercorrect the depths of weak CH₄ absorption bands. The relatively large Q-branch contribution observed by Karkoschka is shown to be consistent with current estimates of Raman cross-sections. The Wallace (1972, Astrophys. J. 176, 249-257) approximation, produces geometric albedo ∼5% low as originally proposed, but can be made much more accurate by including a scattering contribution from the vibrational transition. The original Pollack et al. (1986, Icarus 65, 442-466) approximation is inaccurate and unstable, but can be greatly improved by several simple modifications. A new approximation based on spectral tuning of the effective molecular single scattering albedo provides low errors for zenith angles below 70° in a clear atmosphere, although intermediate clouds present problems at longer wavelengths.     

The optical properties of the finest fraction of lunar soil: Implications for space weathering

Abstract— The fine fraction of lunar soils (<45 μm) dominates the optical properties of the bulk soil. Definite trends can be seen in optical properties of size separates with decreasing particle size: diminished spectral contrast and a steeper continuum slope. These trends are related to space weathering processes and their affects on different size fractions. The finest fraction (defined here as the <10 μm fraction) appears to be enriched in weathering products relative to the larger size fractions, as would be expected for surface correlated processes. This <10 μm fraction tends to exhibit very little spectral contrast, often with no distinguishable ferrous iron absorption bands. Additionally, the finest fractions of highland soils are observed to have very different spectral properties than the equivalent fraction of mare soils when compared with larger size fractions. The spectra of the finest fraction of feldspathic soils flatten at longer wavelengths, whereas those of the finest fraction of basaltic soils continue to increase in a steep, almost linear fashion. This compositional distinction is due to differences in the total amount of nanophase iron that accumulates in space weathering products. Such ground‐truth information derived from the <10 μm fraction of lunar soils provides valuable insight into optical properties to be expected in other space weathering environments such as the asteroids and Mercury.     

The 2.5-5.0 micrometers spectra of Io: evidence for H2S and H2O frozen in SO2

Infrared spectra of Io in the region 2.5-5.0 micrometers, including new observational data, are analyzed using detailed laboratory studies of plausible surface ices. Besides the absorption bands attributable to sulfur dioxide frosts, four infrared spectral features of Io are shown to be unidentified. These unidentified features show spatial and temporal band strength variations. One pair is centered around 3.9 micrometers (3.85 and 3.91 micrometers) and the second pair is centered around 3.0 micrometers (2.97 and 3.15 micrometers). These absorptions fall close to the fundamental stretching modes in H2S and H2O, respectively. The infrared absorption spectra of an extensive set of laboratory ices ranging from pure materials, to binary mixtures of H2S and H2O (either mixed at different concentrations or layered), to H2O:H2S:SO2 mixtures are discussed. The effects of ultraviolet irradiation (120 and 160 nm) and temperature variation (from 9 to 130 K) on the infrared spectra of the ices are examined. This comparative study of Io reflectance spectra with the laboratory mixed ice transmission data shows the following: (1) Io's surface most likely contains H2S and H2O mixed with SO2. The 3.85- and 3.91-micrometers bands in the Io spectra can be accounted for by the absorption of the S-H stretching vibration (nu 1) in H2S clusters and isolated molecules in an SO2-dominated ice. The weak 2.97- and 3.15-micrometers bands which vary spatially and temporally in the Io spectra coincide with the nu 3 and nu 1 O-H stretching vibrations of clusters of H2O molecules complexed, through hydrogen bonding and charge transfer interactions, with SO2. (2) The observations are well matched qualitatively by the transmission spectra of SO2 ices containing about 3% H2S and 0.1% H2O which have been formed by the condensation of a mixture of the gases onto a 100 K surface. (3) No new features are produced in the region 2.5 to 5.0 micrometers in the spectrum of these ices under prolonged ultraviolet irradiation or temperature variation up to 120 K. (4) Comparison of the Io spectra to transmission spectra of both mixed molecular ices and layered ices indicates that only the former can explain the shifts and splitting of the absorption bands seen in the Io spectrum and additionally can account for the fact that solid H2S is observed in the surface material of Io at temperature and pressure conditions above the sublimation point of pure H2S.     

The infrared spectra of nitriles and related compounds frozen in Ar and H2O

We present the 2320-2050 cm-1 (4.31-4.88 micrometers) infrared spectra of 16 solid-state nitriles, isonitriles, and related compounds in order to facilitate the assignment of absorption features in a spectral region now becoming accessible to astronomers for the first time through the Infrared Space Observatory (ISO). This frequency range spans the positions of the strong C triple bond N stretching vibration of these compounds and is inaccessible from the ground due to absorption by CO2 in the terrestrial atmosphere. Band positions, profiles, and intrinsic strengths (A values) were measured for compounds frozen in Ar and H2O matrices at 12 K. The molecular species examined included acetonitrile, benzonitrile (phenylcyanide), 9-anthracenecarbonitrile, dimethylcyanamide, isopropylnitrile (isobutyronitrile), methylacrylonitrile, crotononitrile, acrylonitrile (vinyl cyanide), 3-aminocrotononitrile, pyruvonitrile, dicyandiamide, cyanamide, n-butylisocyanide, methylisocyanoacetate, diisopropylcarbodiimide, and hydrogen cyanide. The C triple bond N stretching bands of the majority of nitriles fall in the 2300-2200 cm-1 (4.35-4.55 micrometers) range and have similar positions in both Ar and H2O matrices, although the bands are generally considerably broader in the H2O matrices. In contrast, the isonitriles and a few exceptional nitriles and related species produce bands at lower frequencies spanning the 2200-2080 cm-1 (4.55-4.81 micrometers) range. These features also have similar positions in both Ar and H2O matrices, and the bands are broader in the H2O matrices. Three of the compounds (pyruvonitrile, dicyandiamide, and cyanamide) show unusually large shifts of their C triple bond N stretching frequencies when changing from Ar to H2O matrices. We attribute these shifts to the formation of H2O:nitrile complexes with these compounds. The implications of these results for the identification of the 2165 cm-1 (4.62 micrometers) "XCN" interstellar feature and the 4550 cm-1 (2.2 micrometers) feature of various objects in the solar system are discussed.     

The solar system cratering record: Voyager 2 results at Uranus and implications for the origin of impacting objects

The cratering record at Uranus shows two different crater populations of different ages. The old crater population occurs on the heavily cratered surfaces of Oberon, Umbriel, and Miranda, while the younger one is found on Titania, Ariel and the resurfaced areas of Miranda. Since only the young population occurs on Titania, this satellite must have experienced a global resurfacing event which obliterated the older population prior to the impact of objects causing the younger one. The old crater population is characterized by an abundance of large craters and a relative paucity of small ones. The young crater population, however, has an abundance of small craters and a paucity of large ones relative to the old population. Furthermore, the abundance of small craters and the paucity of large craters increases with decreasing density. This change in the size distribution is consistent with a population of impactors that evolved with time by mutual collision, and therefore was probably in planetocentric orbits. In fact, both crater populations may be the result of accretional remnants in planetocentric orbits that evolved with time by mutual collisions. If so, then the higher crater density on Miranda compared to Oberon and Umbriel suggests that both Oberon and Umbriel were also resurfaced early in their histories.A comparison of the Solar System cratering record from Mercury to Uranus (19 AU) shows different crater populations at different locations in the Solar System. Computer simulations using a modified Holsapple-Schmidt crater scaling and short-period comet impact velocities to recover the projectile diameters from the cratering record produce different projectile populations in different parts of the Solar System. Furthermore, adjusting the Jovian crater curve to match that in the inner Solar System requires differences in the impact velocities that are unrealistic for objects in heliocentric orbits. These results suggest that the Solar System cratering record cannot be explained by a single family of objects in heliocentric orbits, e.g., comets. One possible explanation is that the cratering record is the result of different families of objects (possibly accretional remnants) indigenous to that region of the Solar System in which the different crater populations are found. Thus, in the inner Solar System, the impactors responsible for heavy bombardment were in heliocentric orbits with semimajor axes less than 3 AU. In the outer Solar System, they may have been in planetocentric orbits around each of the Jovian planets.     

Constraints on the Formation Regions of Comets from their D:H Ratios

Studies of the D:H ratio in H₂O within the Solar nebula provide a relationship between the degree of enrichment of deuterium and the distance from the young Sun. In the context of cometary formation, such models suggest that comets which formed in different regions of the Solar nebula should have measurably different D:H ratios. We aim to illustrate how the observed comets can give information about the formation regions of the reservoirs in which they originated. After a discussion of the current understanding of the regions in which comets formed, simple models of plausible formation regions for two different cometary reservoirs (the Edgeworth–Kuiper belt and the Oort Cloud) are convolved with a deuterium-enrichment profile for the pre-solar nebula. This allows us to illustrate how different formation regions for these objects can lead to great variations in the deuterium enrichment distributions that we would observe in comets today. We also provide an illustrative example of how variations in the population within a source region can modify the resulting observational profile. The convolution of a deuterium-enrichment profile with examples of proto-cometary populations gives a feel for how observations could be used to draw conclusions on the formation region of comets which are currently fed into the inner Solar system from at least two reservoirs. Such observations have, to date, been carried out on only three comets, but future work with instruments such as ALMA and Herschel should vastly improve the dataset, leading to a clearer consensus on the formation of the Oort cloud and Edgeworth–Kuiper belt.     

Optical properties of Titan and early Earth haze laboratory analogs in the mid-visible

Scattering and absorption of sunlight by aerosols are integral to understanding the radiative balance of any planetary atmosphere covered in a haze, such as Titan and possibly the early Earth. One key optical parameter of an aerosol is its refractive index. We have simulated both Titan and early Earth organic haze aerosols in the laboratory and measured the real and imaginary portion of their refractive index at λ = 532 nm using cavity ringdown aerosol extinction spectroscopy. This novel technique allows analysis on freely-floating particles minutes after formation. For our Titan analog particles, we find a real refractive index of n = 1.35 ± 0.01 and an imaginary refractive index k = 0.023 ± 0.007, and for the early Earth analog particles we find n = 1.81 ± 0.02 and k = 0.055 ± 0.020. The Titan analog refractive index has a smaller real and similar imaginary refractive index compared to most previous laboratory measurements of Titan analog films, including values from Khare et al. (Khare, B.N., Sagan, C., Arakawa, E.T., Suits, F., Callcott, T.A., Williams, M.W. [1984]. Icarus 60, 127-137). These newly measured Titan analog values have implications for spacecraft retrievals of aerosol properties on Titan. The early Earth analog has a significantly higher real and imaginary refractive index than Titan analogs reported in the literature. These differences suggest that, for a given amount of aerosol, the early Earth analog would act as a stronger anti-greenhouse agent than the Titan analog.     

Speckle imaging of Titan at 2 microns: surface albedo, haze optical depth, and tropospheric clouds 1996-1998

We present results from 14 nights of observations of Titan in 1996-1998 using near-infrared (centered at 2.1 microns) speckle imaging at the 10-meter W.M. Keck Telescope. The observations have a spatial resolution of 0.06 arcseconds. We detect bright clouds on three days in October 1998, with a brightness about 0.5% of the brightness of Titan. Using a 16-stream radiative transfer model (DISORT) to model the central equatorial longitude of each image, we construct a suite of surface albedo models parameterized by the optical depth of Titan's hydrocarbon haze layer. From this we conclude that Titan's equatorial surface albedo has plausible values in the range of 0-0.20. Titan's minimum haze optical depth cannot be constrained from this modeling, but an upper limit of 0.3 at this wavelength range is found. More accurate determination of Titan's surface albedo and haze optical depth, especially at higher latitudes, will require a model that fully considers the 3-dimensional nature of Titan's atmosphere.