Geology and Surface Processes on Titan

Chapter 5

Geology and Surface Processes on Titan Ralf Jaumann, Randolph L. Kirk, Ralph D. Lorenz, Rosaly M.C. Lopes, Ellen Stofan, Elizabeth P. Turtle, Horst Uwe Keller, Charles A. Wood, Christophe Sotin, Laurence A. Soderblom, and Martin G. Tomasko

Abstract The surface of Titan has been revealed globally, if incompletely, by Cassini observations at infrared and radar wavelengths as well as locally by the instruments on the Huygens probe. Extended dune fields, lakes, mountainous terrain, dendritic erosion patterns and erosional remnants indicate dynamic surface processes. Valleys, small-scale gullies and rounded cobbles such as those observed at the Huygens landing site require erosion by energetic flow of a liquid. There is strong evidence that liquid hydrocarbons are ponded on the surface in high-latitude lakes, predominantly, but not exclusively, at high northern latitudes. A variety of features including extensive flows and caldera-like constructs are interpreted to be cryovolcanic in origin. Chains and isolated blocks of rugged terrain rising from smoother areas are best described as mountains and might be related to tectonic processes. Finally, impact craters are observed but their small numbers indicate that the crater retention age is

R. Jaumann (!) DLR, Institute of Planetary Research, Rutherfordstrasse 2, 12489, Berlin, Germany Freie Universität Berlin, Institute of Geological Sciences, Malteserstr. 74-100, 12249, Berlin, Germany e-mail: [email protected] R.L. Kirk and L.A. Soderblom U.S. Geological Survey, Flagstaff, AZ 86001, USA R.D. Lorenz and E.P. Turtle Space Department, John Hopkins University Applied Physics Laboratory, Laurel, MD 20723, USA R.M.C. Lopes and C. Sotin Jet Propulsion Laboratory, California Institute of Technology, Pasadena CA 91109, USA E. Stofan Proxemy Research, Rectortown VA 20140, USA H. Uwe Keller Max-Planck-Institute for Solar System Research, Katlenburg-Lindau, Germany C.A. Wood Planetary Science Institute, Tucson, AZ 85719, USA M. Tomasko Lunar and Planetary Laboratory, University of Arizona, Tucson, AZ, 85721, USA

R.H. Brown et al. (eds.), Titan from Cassini-Huygens, © Springer Science +Business Media B.V. 2009

very young overall. In general, Titan exhibits a geologically active surface indicating significant endogenic and exogenic processes, with diverse geophysical and atmospheric processes reminiscent of those on Earth.

5.1 Introduction The first detailed discussions about Titan’s surface took place at a workshop at NASA Ames in 1973 (Hunten 1974). Although that workshop was dedicated to Titan’s atmosphere, two contributions in particular addressed Titan’s surface. First, Lewis (1971) considered the formation scenarios for Titan and reasoned that water and ammonia were likely to be the major bulk constituents of Titan. Second, Hunten (1974) recognized that the surface would be the repository for photochemical products of methane photolysis, which could accumulate to a depth of about 1 km over the age of the solar system. This photochemistry-dominated perspective on Titan prevailed through the Voyager encounters in 1980 and 1981 as Voyager images failed to show details of the surface (although a re-analysis of Voyager images 20 years later did show that some surface contrast can be recovered in these data (Richardson et al. 2004). The Voyager radio-occultation experiment revealed the radius of Titan’s solid surface (allowing a more accurate constraint on Titan’s bulk density), and also indicated surface atmospheric conditions of 94 K and 1.5 bar (Tyler et al. 1981). Methane humidity was inferred to be well under 100% (e.g. Flasar 1983), which argued against a surface covered in liquid methane (Fig. 5.1). However, the idea of surface liquids continued to draw interest. Sagan and Dermott (1982) argued that Titan’s orbital eccentricity (0.029) requires tidal dissipation to be low. Noting that, as on Earth, tidal dissipation is high in shallow seas, they argued that seas on Titan should therefore be either deep or non-existent. In addition, they concluded (erroneously, in retrospect) that erosion would obliterate topography on Titan, so that, if deep oceans were present, they would be global, with no islands or continents.

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Fig. 5.1 Methane cycle and deposits of photochemical material (Hunten 1974)

These arguments dominated much of the thinking about Titan’s surface for the following decade. In fact, the tidal dissipation argument was later shown (Sohl et al. 1995) to be irrelevant as a discriminator between surface types because dissipation in the interior (especially if there were a liquid water-ammonia layer, as thermal models began to predict around this time) would in any case make Titan’s eccentricity difficult to understand, whether there were liquid hydrocarbons on the surface or not. Dermott and Sagan (1995) furthermore refuted their own earlier arguments by pointing out that if Titan’s liquid hydrocarbons were confined in small basins such as impact craters, tidal dissipation in those surface liquids would be small. Post-Voyager photochemical models enabled a quantitative treatment of the deposition fluxes of various organic compounds. Notably, the model by Yung et al. (1984) suggested that 100–200 m of solid acetylene and other solids might have been deposited over the age of the solar system, while about 600 m of ethane liquid would be generated. Lunine et al. (1983) suggested that this amount of ethane could resolve the paradox of the low near-surface humidity by the need for a surface reservoir of methane because its inventory in the atmosphere in gaseous form would be depleted in only 10 Myr. The methane vapor pressure above a mixed ‘ocean’ of methane and ethane, which is a rather less volatile molecule and thus has a low vapor pressure, would be less than that above a pure methane ocean. This scenario was consistent with the radio-occultation data and the photochemical model. Formerly, the amount of argon in the atmosphere had not been meaningfully constrained by Voyager data, and the resultant uncertainty about mean molecular weight made for some uncertainty about the surface temperature, which could

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have ranged from 92.5°K to 101°K (Dubouloz et al. 1989). Based on the key assumptions of thermodynamic equilibrium between ocean and atmosphere, and an ethane depth of ~600 m (Yung et al. 1984), this in turn would allow ocean compositions with ethane:methane:nitrogen:argon concentrations of 91:7:2:0 to 5:83:6:6 and ocean depths of 700 m to 9.5 km. Stevenson and Potter (1986) drew attention to the equatorto-pole temperature gradient on Titan and suggested that, somewhat analogously to the Martian seasonal polar caps of CO2, there might be seasonal condensation of methane at the winter pole. In the build-up to the Cassini mission, Lunine and Stevenson (1985, 1987) made a number of calculations regarding Titan’s origin and evolution and the nature of its surface. The first data constraining Titan’s surface was radar reflectivity measurement by Muhleman et al. (1990). Using the Goldstone 70-m antenna as a transmitter, and the very large array (VLA) as a receiver, they detected a Titan echo during three nights. The published reflectivity values were quite high – certainly ruling out a global ocean – and perhaps indicative of a ‘dirty-ice’ surface. In fact, the published values were erroneous – a reduction software error leading to their being reported too high by a factor of 2. If the correct value had been reported, the interpretation might have been less challenging. Soon after radar detection, near-infrared measurements in methane windows began to be made. The realization that Titan’s atmosphere would be sufficiently transparent in the near infrared emerged in the late 1980s. Griffith et al. (1991) suggested that a ‘dirty ice’ composition might be consistent with observations, an interpretation on which 17 years have afforded remarkably little progress, the same conclusion being drawn by McCord et al. (2008) from Cassini VIMS data. Lemmon et al. (1993) showed that Titan’s leading face was brighter than the trailing face, the first detection of surface heterogeneity. Technologically, this discovery could have been made decades before (e.g. McCord et al. 1971; Hunten 1974). Remarkably, Cruikshank and Morgan (1980) searched for spectral variability but failed to detect it although a re-analysis of their data supplemented by new observations did show a light curve (Noll and Knacke 1993 – published shortly before the work by Lemmon et al. (1993)). Several other workers followed up this result (Griffith 1993; Coustenis et al. 1995). All light curves showed essentially the same shape (although differing in amplitude) at all the wavelengths studied – 0.94, 1.07, 1.28, 1.58 and 2 µm. The more challenging 5 µm light curve would not be measured for almost another decade (Lellouch et al. 2004). It was realized that, in principle, the short-wavelength light curves (e.g. 673, 790, 950 nm) could be measured by appropriately equipped amateur astronomers (Lorenz et al. 2003), and indeed, the surface light curve is sufficiently prominent in the small-telescope cloud-monitoring efforts of Schaller et al. (2005) to be taken into account in their analyses.

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Geology and Surface Processes on Titan

Finally, Cassini VIMS measured several light curves during the approach phase in 2004 when Titan was too distant to be resolved (Buratti et al. 2006). The spatial variability implied by the light curve was more graphically revealed by the first near-infrared maps of Titan made with data from the repaired Hubble space telescope in 1994 (Smith et al. 1996) (Fig. 5.2). The 940 nm map showed leading-edge brightness to be concentrated in a region measuring about 4,000 × 2,500 km, now known as Xanadu. The map, made from 14 HST observations, had a coverage gap at 0° longitude. This gap was a factor in the decision to publish the Titan maps centered on the anti-Saturn prime meridian (i.e. 180°W) rather than the 0° prime meridian as on Earth: this convention has been largely but not universally followed since. Another set of HST maps (using a longer-wavelength camera) was generated with data acquired in 1997 (Meier et al. 2000). Ground-based near-infrared imaging with adaptive optics (AO) began to show features in 1993 (Saint-Pe et al. 1993; Combes et al., 1997; Coustenis et al. 2001; Hirtzig et al. (2007); Negraõ et al. (2007)). The same holds true for speckle imaging (Gibbard et al. 1999, 2004). Since then, AO has progressively improved (e.g. Gendron et al. 2004; Roe et al. 2004). By 2004 the 8- and 10-m-class telescopes such as Keck, Gemini and VLT began to surpass the HST resolution. Furthermore, the superior clarity of the 2 µm window used by these telescopes allowed surface features to be detected better. However, the nature of the bright and dark features remained unknown. Attempts to interpret the surface ‘spectrum’ were hampered by measurement uncertainties and differences in the models used to interpret the spectrum. In particular, given a specific observed albedo, the inferred surface reflectivity was highly dependent on the assumed methane absorption coefficient. These difficulties led to uncertainties of a factor of 2–3 in reflectivity, rendering efforts at compositional interpretation largely futile. The almost universal model, ‘dirty ice’, still holds (e.g. Soderblom et al. 2007a). New radar observations of Titan’s surface were enabled in 2002 by the upgrade of the Arecibo radio telescope and

Fig. 5.2 HST Map from Smith et al. (1996), showing brightness heterogeneities

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the changing declination of Titan, which brought it into the view of that facility. This 300-m dish could achieve radar observations of Titan with a much higher signal-to-noise ratio than the Goldstone-VLA work, and allowed the radar spectrum to be measured with more accuracy. The Doppler broadening of the monochromatic transmitted signal was easily seen, allowing the tracking of specific reflecting regions as they moved with Titan’s rotation, and on several occasions, striking specular reflections were seen at the subtelescope point (Campbell et al. 2003). These bright glints required very flat, although actually somewhat dark, surfaces. These results were interpreted as being possibly due to surface liquids, although in retrospect, since these echoes were from latitudes of about −22°, it is possible that flat areas surfaces could have been responsible. It had been noted that no near-infrared specular glints had been seen (West et al. 2005; Fussner 2006) which would have been indicative of liquid surfaces. In anticipation of Cassini-Huygens’s observations, various landforms and surface processes on Titan were considered theoretically, with varying degrees of success (Fig. 5.3). Rainfall had been suggested as a cleansing mechanism that might render elevated terrain optically brighter than lowlands (Griffith et al. 1991; Smith et al. 1996), especially since elevated terrain might receive more rainfall (Lorenz 1993a). But methane raindrops on Titan would fall slowly and might not reach the ground at all before evaporating in the unsaturated lower atmosphere (Lorenz 1993a). Lorenz and Lunine (1996) initially argued that fluvial erosion would be weak on Titan, especially since the meager sunlight does allow a vigorous hydrological cycle amounting to only ~1 cm of methane rainfall per Earth year. However, subsequent thinking (Lorenz 2000) noted that even though desert regions on Earth receive little rainfall on average, rain and rivers can substantially erode the landscape if that rainfall comes in rare but violent storms. That revised paradigm appears to have been borne out. It was realized early (e.g. Greeley and Iversen 1985; Allison 1992; Grier and Lunine 1993) that Titan’s thick

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Fig. 5.3 Sketch of Titan surface processes from Lunine (1990)

atmosphere and low gravity should make aeolian transport easy, in the sense that no more than weak winds would be necessary to move fine-grained surface material. Thus, it was expected that dunes might exist on Titan. Conversely, Lorenz et al. (1995) argued that the weak sunlight and the thick atmosphere, which can transport a great deal of heat at low wind speeds, would not permit strong enough winds for even this low of saltation threshold to be exceeded. Furthermore, consideration of sand-generation processes suggested that the budget of sand-sized sediments on Titan would probably be small. Glacial processes were considered by Lorenz and Lunine (1996) and ruled out for two reasons. First, the accumulation rate of materials would be rather low, so that the driving stresses behind a glacier would be small. Second, the thermodynamic conditions at Titan’s surface did not permit solid methane or ethane to be in equilibrium with the observed atmosphere. Several workers considered possible tides (Sagan and Dermott 1982; Lorenz 1993b; Sears 1995; Ori et al. 1998; Dermott and Sagan 1995) in bodies of liquid: in particular, Lorenz 1994 considered tides in lakes on Titan and noted that the location-dependent tidal tilt could be as high as 2 × 10−5 radians. Lunine (1990) noted that the sedimentation of aerosol particles Tomasko and Smith (1982) or other sediment in Titan liquids would be very slow. In anticipation of possible radar measurements or the floatation dynamics of the Huygens probe, several workers considered waves on the

surface of Titan seas. Especially Srokosz et al. (1992) noted that gravity waves could be large and of relatively long period compared to Earth. Ghafoor et al. (2000) suggested that waves of about 0.2 m amplitude could be generated by winds of ~1 m/s, although this estimate was based on scaling terrestrial relations by gravity, without taking air density or other effects into account. Lorenz et al. (2005) showed by wind tunnel experiments that capillary waves in hydrocarbons grow larger than waves in water at the same wind speed, and also demonstrated that wave growth rates had a strong and somewhat nonlinear dependence on air density. Lorenz (1996b) considered possible cryovolcanism on Titan and noted that the low solubility of methane in water or a water-ammonia solution would be such as to make volatile-enriched cryo-lavas unlikely. Thus, explosive eruptions, plinian or strombolian, for example, could not be expected, and cinder cones and large strato-volcanos resembling Mt. Fuji would not occur, with pancake-dome surface flows appearing more likely instead. Lorenz (1996b) noted that, expressing 10% of the likely geothermal heat flow as latent heat, the approximate fraction of surface eruptions on Earth would yield a global resurfacing rate of 2 mm/year, equivalent to an eruption rate of about 2 km3 per year. Lorenz (2002) considered the possibility of methane geysers on Titan, noting some interesting similarities with Earth with lower temperature gradients being compensated for by the higher volatility of methane compared to water on Earth.

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Several workers investigated impact cratering. Important early considerations were the suggestion of impact shock synthesis in Titan’s atmosphere by Jones and Lewis (1987) and the astrobiological role of ‘impact oases’ containing liquid water created by impact melt (Thompson and Sagan 1992; see also O’Brien et al. 2005). Impact cratering was realized to be a potential tool for detecting variation over time in Titan’s atmosphere, since small craters (