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Planetary geological processes

Conference Paper · November 2014 DOI: 10.1063/1.4902843

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Rosaly M.C. Lopesa and Anezina Solomonidoua,b

aJet Propulsion Laboratory, California Institute of Technology, 4800 Oak Grove Drive, Pasadena, CA 91109 bLESIA - Observatoire de Paris, CNRS, UPMC Univ. Paris 06, Univ. Paris-Diderot – Meudon, 92195 Meudon Cedex, France

Abstract. In this introduction to planetary geology, we review the major geologic processes affecting the solid bodies of the solar system, namely volcanism, tectonism, impact cratering, and erosion. We illustrate the interplay of these processes in different worlds, briefly reviewing how they affect the surfaces of the Earth’s Moon, Mercury, Venus and Mars, then focusing on two very different worlds: Jupiter’s moon Io, the most volcanically active object in the solar system, and Saturn’s moon Titan, where the interaction between a dense atmosphere and the surface make for remarkably earth-like landscapes despite the great differences in surface temperature and composition. Keywords: impact cratering, volcanism, tectonism, erosion, terrestrial planets, Io, Titan PACS: 96

INTRODUCTION

The solid bodies of the solar system have different surface appearances because of the existence and relative importance of the major planetary geologic processes. The surfaces we see today have been shaped by the interplay of endogenic (volcanism, tectonism) and exogenic (impact cratering, erosion and surficial) processes. Understanding the distribution and interplay of endogenic and exogenic processes on a planet is important for constraining models of the interior, surface-atmosphere interactions and climate evolution. This chapter will review the major types of solid bodies in the solar system in terms of comparing the role of these major geologic processes. The surfaces of the Moon and Mercury are characterized by numerous impact craters, while the surface of Venus and Io are dominated by volcanism, with impact craters being totally absent on Io. Saturn’s moon Titan presents, like the Earth, significant effects from erosion by liquids and wind. To understanding the evolution of these different worlds, we need to understand how planetary geologic processes operate. The chapter aims to give an overall introduction to the very complex field of planetary geology.

Our Geologically Diverse Solar System

Studies of planetary geology can greatly help us to understand how the planets and moons of the solar system were formed, and how their current surfaces came to be. Analyses of rocks, either in-situ or brought to Earth via meteorites or space missions can help determine the composition of the planet or moon as a whole and from this to infer the composition of the material from which those bodies were originally formed. Studies of the geology of other planetary bodies largely rely on studies of the geology of our own planet. However, there are major differences. The Earth’s surface is largely dominated by plate tectonics, in which large plates of the crust move and create mountain chains (where plates collide), subduction zones (where one plate dives under another) and spreading ridges (where plates are moving apart). No other body in the solar system is known to have plate tectonics. One consequence of plate tectonics is the destruction of old crust. If we compare the Earth’s surface with, for example, the Moon’s, we immediately see that much of the Moon’s surface is covered by impact craters, while few impact craters can still be found on the Earth. Other than plate tectonics, the other destructive process on Earth is erosion. The Earth is the only planet with the right combination of atmospheric surface pressure and temperature to allow liquid water to exist and to cover such a large fraction (~70%) of the surface. Impact craters are erased relatively rapidly on Earth by the action of plate tectonics and erosion due to weather, as well as volcanism. The Earth’s atmosphere does offer our planet some protection from small meteorite impacts and, for this reason alone, would have a different crater size and frequency distribution from that of an airless body. The planets of the solar system can be divided into two major compositional groups. Mercury, Venus, Earth, and Mars are known as the terrestrial (Earth-like) planets and are characterized by silicate compositions and iron cores. They formed much closer to the Sun than the outer planets, in the part of the solar nebula which was too warm for ices to condense. The terrestrial planets all have solid surfaces where the major geologic processes have operated. Mercury is heavily cratered, its proximity to the Sun allows for impactors to have high encounter velocities. The lack of any significant atmosphere on Mercury also means that, not only it does not have any protection from oncoming objects, but also it has had no erosion by weather. Similarly, the lack of plate tectonics means that old crust was largely preserved. If we consider impact craters, knowing that the oldest surfaces of the solar system will have the largest impact scars and the largest numbers of craters (1), then next to Mercury is Mars, with also a considerable number of craters, largely because of its proximity to the asteroid belt. However, the surface of Mars shows considerable signs of active geology such as volcanism, tectonism, and erosion. The surface of Venus has fewer impact craters than Mercury or Mars, and is dominated by volcanism. The dense Venusian atmosphere (with a surface pressure of 94 bar) is capable of protecting the surface from some impactors, which break up in the atmosphere. Also, volcanic activity has resurfaced large areas of the planet, erasing the evidence of craters. The small number of impact craters recognized on the Earth’s surface or in the oceans (184 according to (2)) attest to how geologically young the surface of our planet is.

FIGURE 1. Montage of images of planets in the solar system (not to scale). At the top is an image of Mercury, showing its old cratered surface imaged by Mariner 10. Next down is Venus, its volcanic surface revealed by the radar instrument on the Magellan spacecraft. The Earth and the Moon, imaged by the Galileo spacecraft, further illustrate the diversity of geology, the ancient surface of the Moon contrasting with the young surface of the Earth. Mars (imaged by Mars Global Surveyor) shows both ancient cratered surfaces and younger terrains created by volcanism and erosional processes. The Jovian planets (Jupiter imaged by the Cassini spacecraft and Saturn, Uranus, and Neptune imaged by Voyager) are gas giants with no solid surfaces, but have a plethora of moons showing yet more of our solar system’s diverse geology.

The Jovian (also called Jupiter-like or gas giants) planets are Jupiter, Saturn, Uranus, and Neptune. Because of their primarily gaseous compositions, they do not have solid surfaces where the geologic processes described above operate, thus they are thought to have silicate-iron cores. Of particular interest to geologic studies are their moons. In fact, the moons of our solar system display remarkable geologic diversity. The Moon is the best studied extra- terrestrial body so far, and the only one for which we have collected samples and brought them back to Earth for study. Our Moon has a silicate composition similar to the Earth’s mantle, and a small iron core. Mars has two moons, Phobos and Deimos, both are small and irregularly shaped. Their surface compositions appear to be similar to carbonaceous chondrite meteorites in composition. These moons may be captured asteroids and may have formed elsewhere in the solar system. Venus and Mercury have no moons. In contrast, the gas giants offer a large number of moons that, even within one system, show remarkable diversity in the surface geology. Jupiter’s four major satellites are a good example of such diversity, ranging from Io’s volcanically dominated surface to Europa’s icy young surface, Ganymede’s tectonic activity and evidence of resurfacing at some time in the past, to Callisto’s heavily cratered, ancient surface. Saturn’s largest moon, Titan, is the only moon in the solar system with a substantial atmosphere, allowing for erosion processes to be rampant. Tiny Enceladus ejects volcanic plumes from its southern polar region, our only uncontested example so far of cryovolcanism, a type of volcanism in which water, rather than molten rock, is the magma (3). Uranus and Neptune have many icy satellites, with Miranda being particularly unusual because of its complex surface geology while Triton shows evidence of recent cryovolcanism, as well as plumes probably caused by solar heating. Many other small worlds in the solar system display interesting geology, and we look forward to finding out what the surface of Pluto will be like, as the New Horizons spacecraft approaches.

FIGURE 2. The four largest moons of Jupiter, known as the Galilean satellites, show how different geologic processes have shaped their present surfaces. These moons were first seen by Galileo Galilei in 1610. Left to right and in increasing distance from Jupiter are Io, Europa, Ganymede, and Callisto. The distances of these moons from Jupiter help explain some of the differences in their composition and geology. Io is subject to the strongest tidal stresses, which generate enough internal hear to make Io the most volcanically active body in the solar system. Europa has an icy crust with relatively few impact craters, indicating a relatively young age. Ganymede’s surface is characterized by tectonic resurfacing and Callisto, furthest from Jupiter, shows no evidence of internal activity and a heavily cratered, ancient surface. The images were obtained by the Galileo spacecraft. (NASA/DLR PIA01400).

THE TERRESTRIAL PLANETS: MOON, MERCURY, VENUS AND MARS

The Earth’s Moon

The Earth’s Moon is unusual in several ways. It has a low density compared to the inner planets, and a small core. The Moon is thought to have formed as a result of an oblique impact between the Earth and another large, differentiated body sometimes known as Theia (4), which probably had a mass similar to Mars. This collision is thought to have occurred about 4.5 billion years ago (5). The impact released enormous amounts of energy, creating so much heat that many volatile materials, such as water, escaped. The Moon therefore formed mostly from volatile- depleted mantle materials, which explains its current composition. The lack of an atmosphere, water, and plate tectonics mean that the surface of the Moon is, for the most part, ancient. It is estimated that 99% of the lunar surface is older than 3 billion years, and more than 80% is older than 4 billion years (5). In contrast, more than 80% of the Earth’s surface is younger than 200 million years. Impact craters dominate much of the lunar surface and the Moon has provided us with much knowledge about impact cratering as a geological process. The major process modifying the lunar surface (we can consider it a form of erosion) is impact, ranging from micrometer-sized grains to asteroid-sized objects. These many impacts created the lunar regolith, a layer of debris that covers the surface. The material in the regolith ranges from dust to blocks several meters across. The average thickness of the regolith is 4-5 m over the maria (i.e. extensive dark basaltic plains) and about 10 m over the highlands. The Moon has no large tectonic features, though there are tectonic features related to stresses associated with subsidence of the mare basins after they were flooded by lavas. Wrinkle ridges (also called mare ridges) are low- relief, linear to arcuate, broad ridges that tend to form near the edges of maria. They are the result of compressional bending stresses, related to the subsidence of the mare basalts from cooling (5). Some extensional stresses existed on the Moon prior to about 3.6 billion years ago. Rilles, which are similar to terrestrial graben, are found cutting the older maria and highlands. Note that Hadley Rille and other sinuous channels (called sinous rilles) in the maria (Fig. 4) were not formed by tectonic stresses, but by thermal erosion of flowing lava. Planetary geologists are now tending to refer to tectonic lunar rilles as graben to avoid confusion.

FIGURE 3. The Earth’s moon imaged by the Galileo spacecraft in 1992.showing the distinct dark and light areas that characterize the surface. The dark areas (maria) are large expanses of cooled lava that filled impact basins. Oceanus Procellarum is on the left, Mare Imbrium in the center left, Mare Serenitatis and Mare Tranquilitatis are in the center of the image, and Mare Crisium near the right edge. The light-colored areas, known as terrae (highlands) show many impact structures. The distinct crater at the bottom of the image with bright rays is Tycho. The image, taken by Galileo’s Solid State Imaging system, is a composite of images taken through the violet, the 756 nm, and the 968 nm filters; the colors are enhanced to show detail. (NASA PIA00405).

FIGURE 4. A montage of images of the Moon showing possible volcanic features. Clockwise from upper left: False color Clementine mosaic of the Aristarchus Crater and Plateau, showing reddish pyroclastic deposit on the Plateau, and the sinuous rille (lava channel) known as Schröter’s Valley; Apollo 15 image of well-preserved lava flows in Mare Imbrium; Apollo orbiter image of lunar sinuous rilles (lava channels); Apollo orbiter images of the Ina structure of the lunar nearside, thought to represent a recent gasp of lunar degassing ~10 Ma (9); False color Galileo mosaic of the lunar nearside, in which orange units represent low-titanium mare deposits, and blue units represent high-titanium mare deposits, derived from ancient lava flows; image of a dark mantle (pyroclastic) deposit in the near side crater Alphonsus. (Courtesy of David A. Williams, Arizona State University).

The maria, seen easily with the naked eye as dark patches on the lunar surface, were formed by flowing basaltic lavas. They cover about 17% of the Moon’s surface (6) and are very smooth, lacking large constructional forms that would be identified as volcanic mountains. What we see today are massive flood-like lava flow fields that fill impact craters and basins. The maria were mostly emplaced from ~3.9 to ~3.2 billion years ago. Over geologic time the effects of space weathering have destroyed the primary volcanic features such as lava flow boundaries in all but a few cases. Sinous rilles are lava channels that appear as levee-less meandering channels, ~2-300 km in length, meters to 3 km in width, mean depths ~100m, with a crater-like depression at the head of the rille and a fading down slope into the maria at the end of the channel (7). Many volcanic domes, cones, and small shield volcanoes have been preserved in the maria. These are small features, typically a few km in diameter. There are no shield volcanoes on the Moon larger than ~20 km diameter, indicating a lack of shallow buoyancy zones, and demonstrating that lava extrusion did not involve low-volume, short-duration eruptions from shallow reservoirs (8). This is a unique aspect of lunar volcanism compared to the other terrestrial planets.

Mercury

Like the Moon, Mercury has no erosion due to an atmosphere or surface liquids. Mercury is a very dense body due to its large iron core, which is estimated to be 2,020 km in radius (10). Mercury’s surface is somewhat similar to that of the Moon and impact cratering has been a major process. Mercury has heavily cratered highlands and large areas of younger smooth plains that surround and fill impact basins. Like the Moon, the surface of Mercury has a layer of regolith formed by many impacts. The volcanic areas are flooded plains similar to the lunar maria. In addition, there are tectonic features, though Mercury, like all other planets except the Earth, has no plate tectonics. The exploration of Mercury by spacecraft started in the 1970s with the Mariner 10 mission. It would be decades before the next mission, MESSENGER, went back to the planet to make significant new discoveries. Mercury’s geology, like the Moon’s, is dominated by impact cratering, volcanism, and tectonic features. Mercury has several impact basins, including the largest in the solar system: the Caloris basin, with a diameter of 1550 km. The impact that created the Caloris Basin was so powerful that it had global effects on the geology including causing outpourings of lava. At the antipode of Caloris is a large region of unusual, hilly and furrowed terrain, sometimes known as “Weird Terrain”. This terrain may have been formed by shock waves generated during the impact travelling around the planet and converging at the basin’s antipode, fracturing the surface (11). Alternatively, it may have formed by converging ejecta (12). Mercury is thought to have cooled and substantially shrunk in size early in its history. The surface we see today is crisscrossed by steep crustal ridges that reveal the period of rapid contraction. Results from Mariner 10, which only imaged part of the planet, suggested that Mercury had shrunk by about 3 km on average. However, results from the Messenger spacecraft, which has been orbiting Mercury since 2011 mapping the whole planet, show that Mercury shrank more than previously thought, losing about 7 km from its original girth (13). This contraction of the crust might have prevented magma from rising to the surface, however, Mercury’s surface has smooth plains that are volcanic in origin (14). Moreover, the volcanism overlapped with rapid contraction for about 200 to 300 million years (15). This was possible probably because of large meteorite impacts that could have released built-up pressure in the compressed crust and allow lava to pour out while Mercury was shrinking (15). This idea is supported by images of Mercury that show lava flows inside large craters and basins. Mercury does not have a large variety of volcanic features, its volcanic landforms are mostly smooth plains. However, a set of landforms at high northern latitudes resembles surface flow features seen on the Moon, Mars, and Venus (14). Smooth plains cover about 6% of Mercury’s surface and appear to have been formed by low-viscosity lavas flowing at high effusion rates, similar to the flood basalts that formed large basaltic igneous provinces on Earth (16). This interpretation is consistent with compositional data that suggests substantial portions of Mercury’s crust are composed of magnesium, iron-poor materials (17). Two lava compositions are likely, tholeiitic basalt and komatiite (14). These lava compositions erupt on Earth, though komatiites have not erupted for hundreds of millions of years. Mercury’s geology is particularly interesting because it shows us the evolution of a one-plate planet with extreme ranges of surface temperatures, showing volcanic, impact, and tectonic features.

FIGURE 5. This false-color view of Mercury was produced by using images obtained by the MESSENGER spacecraft. These colors enhance the chemical, mineralogical, and physical differences between the rocks that make up Mercury's surface. Young crater rays, extending radially from fresh impact craters, appear light blue or white. Medium- and dark-blue areas are thought to be rich in a dark, opaque mineral. Tan areas are plains formed by eruption of highly fluid lavas. The giant Caloris Basin is the large circular tan feature located just to the upper right of center of the image. (NASA PIA16853).

FIGURE 6. This region is located diametrically opposite the Caloris impact basin on Mercury, and has been affected by converging seismic waves or ejecta that caused the peculiar "hilly and lineated" texture. Members of the Mariner 10 team informally called this the "weird terrain" when it was first discovered in 1974. Image is about 248 km across.

Venus

Venus is similar to Earth in terms of size and density, but very different in terms of atmospheric conditions and geologic evolution. The greenhouse effect, in which carbon dioxide causes the atmosphere to heat up, was first discovered on Venus from the Mariner 2 spacecraft flyby in 1962, which also measured the surface temperature at 468oC. The surface pressure on the Venusian surface is 95 bars, making it the densest atmosphere of any solid bodies in the solar system. Venus may still be volcanically active today, but it lacks plate tectonics. Like Earth, Venus is large enough for heat (primarily from the decay of radioactive elements, but also from heat loss from accretion) to convect. On Earth, convection is linked to volcanism and tectonism via plate tectonics. Water on Earth serves to reduce the strength of rocks and lubricate plate subduction. The lack of water on Venus may have been a key factor on why plate tectonics didn’t operate. Although the atmosphere of Venus has some water, it would only amount to a layer less than 10 cm thick. The dry Venusian crust was likely to stiff to break up into plates. The exploration of Venus has involved landers as well as orbiters and flyby spacecraft. Soviet Venera probes landed on Venus between 1970 and 1981, sending back data on surface temperature, pressure, and composition, as well as images and geochemical measurements which were consistent with basalts. Rocks at each site are angular, implying minimal erosion. All the sites are consistent with a volcanic origin. The last two landers, Venera 13 and 14, returned color panoramas and drilled into the surface for samples. The exploration of Venus has also seen balloon probes and orbiters. NASA’s Magellan mission obtained Synthetic Aperture Radar (SAR) images of 98% of the Venusian surface between 1990 and 1994. The 120 m resolution SAR images and 1-10 km resolution altimetry data provided a rich data set for geologic studies. The surface of Venus is dominated by plains, which cover at least 80% of the planet (18). There are also major highlands, including plateaus and topographic rises, and ridges and rifts. Venus has few impact craters (~940 identified), indicating that its surface is young. On average, it is comparable to the surfaces of the Earth and Titan. The dense atmosphere causes small impactors (<1km in diameter) to break up before they reach the ground, reducing the number of small craters. The estimated surface age of Venus is ~750 Ma (million years) but, given the uncertainties due to the relatively small number of craters, estimates vary from 300 Ma to 1 Ga (billion years). These ages indicate a relatively young surface and that resurfacing processes are or have happened in the not too distant past. Volcanic products cover much of the Venusian surface. Venus is losing its heat not by plate tectonics, but likely by mantle plumes. Convection in the interior forms hot spots, in a similar way to the hot spot on Earth over which the Hawaiian islands have formed. On Venus, the hot material pushes up on the lithosphere, causing a topographic swell. The heat of the lithosphere causes the crust to melt locally, causing it to thicken and for volcanoes to form. While hot spot volcanism accounts for only ~10% of volcanism on Earth (the rest being controlled by plate boundaries), on Venus this is thought to be the major avenue for volcanoes to form. Volcanic features on Venus have a broad range of morphologies, from vast, sheet-like lava flows to volcanic domes shaped like pancakes. Most, if not all, of the volcanism was effusive. The high atmospheric pressure on Venus makes it less likely that explosive volcanism will occur. Like the Moon, Venus has large plains formed by lava flows. Long, sinuous channels are seen on the lavas, some of them thousands of km long. The Baltis Valley is the longest lava channel in the solar system, about 6,400 km long. As both ends of this channel have been covered by lava flows, it is possible that the channel is even longer. The chemical composition of the lavas must allow them to flow over long distances. They may be ultramafic lavas, or sulfur-rich or more exotic types. We lack compositional information from Venus, so the question remains open. There are many shield volcanoes, most are 20 km in diameter or less. In many areas, they cluster in areas known as shield fields. Some of the shield volcanoes are very large, but there are relatively few of those. For example, Sif Mons (Fig. 7) is about 350 km in diameter and 2 km high. Large shields tend to sit atop broad topographic rises, suggesting that they may be over hotspots. These hotspots appear randomly located on Venus, with no alignments that might suggest tectonic plates. Venus has several types of volcanoes that differ from those on Earth and also those on other planets. Steep-sided domes, nicknamed pancake domes, have flat tops but steep sides. They may be similar to silicic domes on Earth, formed by high viscosity lava that is extruded slowly. However, the Venusian domes are much larger, up to 50 km in diameter, while the terrestrial domes are at least ten times smaller. Pancake domes have smooth rather than rough surfaces, another difference from the terrestrial counterparts. It is still not clear how they formed on Venus. They are probably the result of very pasty lava flows being slowly extruded and spreading out slowly on flat plains. Some pancake domes seem to give rise to other strange volcanic constructs, such as those known as ticks, which appear to be degraded pancake domes. Related to ticks are the arachnoids, volcanic domes surrounded by a cobweb of fractures and crests. Russian scientists first named these features after seeing vast concentric fractures spreading out from volcanic sources in Venera radar images. Their sizes range from 50 to 230 km. Lines radiating beyond the arachnoids may be cracks or ridges resultant of upwelling magma that stretched the surface around them. Another strangely type of volcanic feature on Venus is named anemone. These volcanoes have lava flows arranged in overlapping petals extending outward in flower-like patterns. It is thought that the lava flows usually occur in association with fissure eruptions involving a series of elongated vents at the summit. Coronas are another volcano type on Venus that has no terrestrial analogue, and no similar features are seen on other planets. Coronas are large (> 100 km) circular or oblong systems of fractures and ridges that span hundreds of kilometers. Coronae are probably formed by rising blobs or a stream of hot material that bow up and deform the surface, producing a ring of concentric ridges. As the magma pushes up against the crust, the surface rises forming a dome. The hot lava spreads out under gravity and flattens. As the entire area cools, it sinks and cracks, forming a circular ring around a depression. Some coronae appear to be very old and are sunken almost beyond recognition. The surface of Venus is geologically young, perhaps even active. The European Venus Express spacecraft observations have revealed that some volcanoes on Venus appeared to have erupted recently, between a few hundred years to 2.5 million years ago (19). This study found an excess of thermal emissions, indicating a lack of surface weathering compared to the rest of the planet, which can be explained by relatively recent lava flows. Changes in the level of SO2 have been also detected in the atmosphere, which may indicate active volcanism.

FIGURE 7. Sif Mons on Venus is seen in this computer-generated view using altimetry and Synthetic Aperture Radar (SAR) data from the Magellan spacecraft. Sif Mons has a diameter of about 350 km and is about 2 km high. (Image courtesy of USGS/JPL/NASA, PIA00108).

Mars

The surface of Mars is more Earth-like than that of any other body except for Titan. Mars presents remnants of all the major geologic processes. Volcanism has been an important part of the geologic history. Mars has the largest shield volcanoes in the solar system, vast plains of lava flows, many lava channels, domes, and cones, and it has evidence of explosive volcanism and extensive pyroclastic deposits. Because Mars has been the primary focus of planetary exploration since the mid-1990s, and because it has been visited by several spacecraft including orbiters, landers, and rovers, there is a vast quantity of data available on Martian surface geology, including images at resolutions from 30 cm/pixel to several km/pixel, as well as data on surface compositions and the role of volatiles, and atmospheric dynamics and composition, though the composition and nature of the Martian interior has not been well investigated. Rovers have returned field geologist’s views of several areas on the surface. A major feature of the Martian geology is the hemispheric dichotomy: about one-third of the planet’s surface (mostly in the northern hemisphere) lies 3–6 km lower in elevation than the southern two-thirds. The dichotomy is also seen as a difference in impact crater density and crustal thickness between the two hemispheres (20). The terrains south of the dichotomy boundary, the southern highlands, are very heavily cratered and ancient, characterized by rugged surfaces that date back to the period of heavy bombardment. However, the northern lowlands have relatively few large craters, are very smooth and flat, and appear to have undergone extensive resurfacing since the southern highlands formed. Another distinction between the two hemispheres is in crustal thickness. Topographic and geophysical gravity data indicate that the crust in the southern highlands has a maximum thickness of about 58 km, while crust in the northern lowlands "peaks" at around 32 km in thickness (21,22). The origin and age of the dichotomy are still being debated, explanations generally fall into two categories: either the dichotomy is a result of a giant impact, or else the dichotomy is the result of crustal thinning in the northern hemisphere by mantle convection, overturning, or other chemical and thermal processes in the planet’s interior. Mars has the largest volcanoes in the solar system. The main volcanic region is the Tharsis bulge, an elevated structure thousands of kilometers in diameter, which covers up to 25% of the planet’s surface (23). Averaging 7–10 km above datum (Martian "sea" level), Tharsis contains the highest elevations on the planet and the largest known volcanoes in the Solar System. Three of these, Ascraeus Mons, Pavonis Mons, and Arsia Mons (collectively known as the Tharsis Montes), sit aligned NE-SW along the crest of the bulge. Another huge volcano, Alba Mons (formerly Alba Patera) is located further north. The huge shield volcano Olympus Mons, the largest volcano known, is located at the western edge of the Tharsis bulge. Tharsis has placed tremendous stresses on the planet’s lithosphere, which resulted in a series of extensional fractures (grabens and rift valleys) that radiate outward from the bulge, extending halfway around the planet (24). Other large volcanoes on Mars are found in the Elysium volcanic complex several thousand kilometers west of Tharsis. The Elysium complex is about 2,000 kilometers in diameter and consists of three main volcanoes, Elysium Mons, Hecates Tholus, and Albor Tholus. While the Tharsis volcanoes appear to be shields formed by effusive eruptions, the Elysium volcanoes are thought to have erupted both lavas and pyroclastics. The major tectonic feature on Mars is the Tharsis bulge, but Valles Marineris is also noteworthy: it is an enormous system of interconnected canyons in the equatorial region. The canyons are up to 300 km wide and 10 km deep and extend for over 4,000 km. Although popularly compared with the United States Grand Canyon, the origin of Valles Marineris is thought to be tectonic, by faulting, and not erosional, perhaps more similar to a rift valley on Earth (24). The canyons have been subsequently modified by erosion (fluvial and mass wasting). It is possible that Valles Marineris formed as a consequence of the strain put on the crust by the Tharsis bulge, which stressed the lithosphere and caused it to flex under the load of the volcanic pile. Mars has other tectonic features which do not appear to be associated with Tharsis. For example, ridges suggestive of compression can be found in areas such as Hesperia Planum and Syrthis Major, which are far away from Tharsis. Circular fractures around large volcanoes such as Elysium Mons are thought to have formed as a result of bending of the lithosphere under the load of volcanoes such as Elysium Mons. Erosion has happened on Mars on a large scale. Valley networks are thought to have been carved by flowing water when the Martian climate was significantly warmer than it is today, yet how they climate changed is now known (25). There is geologic evidence that huge floods episodically happen during the early history of Mars, carving vast valley networks (26). An ocean may have existed in the northern polar latitudes (27). Ice and snow are also thought to have played an important role in the erosion of the Martian surface. Melting snow during periods of high obliquity is thought to have formed gulleys on steep, poleward-facing slopes. During those periods, water is driven off the polar regions and accumulates as ice at lower latitudes. The ice on poleward-facing slopes may melt in midsummer when the slopes face the sun, causing gulleys to form (24). An additional erosional mechanism on Mars is by the action of wind, which can have substantial effects at local scales. Mass wasting has also played a role on Mars, in some places substantial, such as in the formation of the aureole material around Olympus Mons (28). Impact craters on Mars resemble those on the Moon and Mercury, the patterns of ejecta are quite different on Mars as the ejecta around most fresh-looking Martian craters especially those in the 5-100 km range, have discrete, clearly outlined lobes. These Martian ejecta patterns have been explained in two possible ways. The patterns could be formed by interaction of the ejecta with the atmosphere and there is some experimental work in support of this theory. However, a second explanation is generally preferred: the ejecta contained water and had a mud-like consistency and it continued to flow on the ground after deposition (24). Mars also has large impact basins, the largest of which is Hellas, located in the southern hemisphere, with a diameter of about 1,800 km. Much is known now about the surface geology of Mars because of rovers such as Pathfinder, Spirit, Opportunity and, more recently, Curiosity. These results suggest that Mars was once a habitable environment (e.g. 29, 30) and one of the most likely places in the solar system to find evidence of life.

FIGURE 8. Maps of Mars' global topography. The projections are Mercator to 70° latitude and stereographic at the poles with the south pole at left and north pole at right. Note the elevation difference between the northern and southern hemispheres. The Tharsis volcano-tectonic province is centered near the equator in the longitude range 220° E to 300° E and contains the vast east- west trending Valles Marineris canyon system and several major volcanic shields including Olympus Mons (18° N, 225° E), Alba Patera (42° N, 252° E), Ascraeus Mons (12° N, 248° E), Pavonis Mons (0°, 247° E), and Arsia Mons (9° S, 239° E). (NASA PIA02031).

FIGURE 9. Olympus Mons on Mars is shown as topography draped over a Viking image mosaic. The topography clearly shows the relationship between the volcano's scarp and massive aureole deposit that was produced by flank collapse. The vertical exaggeration is 10:1. (NASA PIA02805).

JUPITER’S MOON IO: EXTREME VOLCANISM

Io (Fig. 10) differs from most solid bodies in the solar system as it has no impact craters, implying a very young surface. It is the most volcanically active body in the solar system and the only place other than the Earth where we see large-scale silicate volcanism. Active volcanism on Io was discovered in 1979 (31) when the camera aboard the Voyager 1 spacecraft showed active plumes erupting and reaching 100-300 kilometers above the surface (Fig. 11). When the infrared instrument showed increased temperatures in some places (dubbed “hot spots”) and one of them coincided with the location of one the plumes, there was no doubt that active volcanism was going on (32). Active volcanism on Io can happen because of its orbit and, in fact, the volcanism was predicted shortly before the Voyager spacecraft arrival (33). Io is about the same size as the Earth’s moon and should have cooled and formed a thick crust long ago. However, Io is in a peculiar orbital dance between Jupiter and the other moons that, along with Io, are known as the Galilean satellites: Europa, Ganymede, and Callisto. Io finds itself in a gravitational tug-of-war between Jupiter, which is massive and the other moons, particularly Europa and Ganymede. All four moons are in orbits such that, for every 2 rotations around Jupiter that Io makes, Europa makes one; for every rotation Ganymede makes, Europa makes two, and so on. Jupiter’s gravitational pull on Io is so large that it creates a tidal bulge on the crust. Europa and Ganymede distort the bulge as they pull it towards them. The friction creates heat, which keeps the interior of Io molten. The colors of Io, mostly yellows, reds, and oranges, as well as patches of white and black, are consistent with the colors of sulfur. Sulfur dioxide was detected from one of the plumes (32), so the question arose whether Io’s volcanism was silicate or sulfur. The two Voyager spacecraft did not carry instruments capable of distinguishing the two compositions. One of the ways of telling sulfur apart from silicate lavas is from their eruption temperature. The highest temperature detected by the Voyager infrared instrument was about 300o C, but most temperatures of hot spots were found to be closer to 100oC, which is about the melting temperature of sulfur. Basalt melts at around 1000oC, but cools quickly once erupted, so there was no way of telling if the Io hot spots were sulfur or cooled basalts. Unfortunately, the infrared instrument aboard the Voyager spacecraft was not designed to be sensitive to the lower infrared wavelengths, which are best at detecting high temperatures. The composition of Io’s lavas remained undetermined for some time, until telescopic observations from Earth, using infrared instruments, detected temperatures around 700oC – too high for sulfur (34). By this time, the Galileo spacecraft was already on its way, with instrumentation better equipped to measure temperature of Io’s hot spots. The Galileo spacecraft’s observations of Io’s surface, from 1996 through 2001, showed many more active volcanoes (hot spots) than had been previously known (35, 36). Data from the Near-Infrared Mapping Spectrometer (NIMS) revealed 71 previously undetected hot spots during this time, others were found by Galileo’s camera and the long-wavelength Photopolarimeter Radiometer (36). We now know that Io has at least 150 active volcanoes. Because Galileo did not make high-resolution observations of the whole moon, and spatial resolution was poor on the Jupiter-facing side, it is estimated that many more active volcanoes exist – probably at least 400 in total.

FIGURE 10. Jupiter’s moon Io is seen in the highest resolution obtained to date by NASA's Galileo spacecraft. The smallest features that can be discerned are 2.5 kilometers in size. There are rugged mountains several kilometers high, layered materials forming plateaus, and many irregular depressions called paterae, similar to volcanic calderas. Several of the dark, flow-like features correspond to hot lavas. There are no landforms resembling impact craters, as the volcanism covers the surface with new deposits much more rapidly than the flux of comets and asteroids can create large impact craters. The picture is centered on the anti-Jovian hemisphere (NASA PIA00583).

Results from the Galileo mission substantially advanced our understanding of volcanism on Io (37). Galileo showed that Io’s volcanoes cover a wide range of sizes and present varying characteristics such as power output, persistency of activity, and association with plumes. The plumes are composed of sulfur and sulfur dioxide gas, with a small percentage of silicate ash. Most of Io’s volcanoes are caldera-like depressions (38), referred to as paterae (singular patera, meaning ‘saucer- like crater’). Unlike terrestrial volcanoes, Io’s volcanoes rarely build large topographic structures such as shields or stratovolcano-like mountains. There are only a few structures, called tholi (singular tholus, meaning “dome”, though some are more similar to shields), scattered across Io. Some of these features may have been formed by basaltic volcanism building shield-like structures.

FIGURE 11. The Voyager 1 spacecraft obtained this image of the Loki plume in March 1979. The brightness of the plume has been digitally enhanced but the relative color of the plume (greenish white) has been preserved. (NASA PIA1971).

Although the origin of paterae is still somewhat uncertain, they are thought to be similar to terrestrial volcanic calderas, formed by collapse over shallow magma chambers following partial removal of magma. Some paterae show angular shapes that suggest some tectonic control, indicating that they may be structural depressions that were later used by magma to travel to the surface. At least 400 Ionian paterae have been mapped (38). Their average diameter is ~40 km but Loki, the largest patera known in the Solar System, is over 200 km in diameter. In contrast, the largest caldera on Earth, Yellowstone, is ~80 km by 50 km in size. The larger sizes of the Ionian features probably reflect the much larger sizes of magma chambers, which are thought to be relatively shallow. Io’s surface shows some remarkably large lava flows (called fluctii, singular fluctus, meaning ‘flow field’). The lava flow field from Amirani (Fig. 12) is ~300 km long, the largest active flow field known in the Solar System. Io’s large lava flows are possibly analogues of the continental flood basalt lavas on Earth, such as the Columbia River Basalts in the USA (39). Repeated imaging of Amirani during the Galileo fly-bys allowed eruption rates to be estimated to be 50 to 500 m3 s-1 (39). These are considered moderate effusion rates by terrestrial standards. The ability of the Amirani lava flow field to travel large distances at moderate effusion rates suggests not only that lavas had a low viscosity, but also that they were emplaced as insulated (probably pahoehoe like) flows, so that the cooled crust would insulate the hot material underneath. Thermal profiles along the Amirani flow field, and also along another flow field called Prometheus, obtained using the NIMS instrument on Galileo, and high spatial resolution images from the camera suggest that these large flows were similar to terrestrial pahoehoe flows (40).

FIGURE 12. The Amirani lava flow on Jupiter's moon Io appears to be made up of many individual flows; the newest flows are the brightest spots in this infrared image from the NIMS instrument on NASA's Galileo spacecraft. The thermal map is presented on the left, beside a reference picture of the same area from Galileo's camera. The infrared image uses false color to indicate intensity of glowing at a wavelength of 5 microns. White, reds and yellows indicate hotter regions; blues are cold. North is to the top. (NASA PIA03533).

A major question about Ionian volcanism after Voyager was the nature of volcanism - whether sulfur or silicates were predominant. Although temperature measurements from Galileo clearly showed that many hot spots have temperatures far too high for sulfur, the possibility that some sulfur flows occur on the surface cannot be ruled out. Galileo observations showed some locations that may have sulfur flows (41). While most Ionian flows appear dark, a few locations show pale yellow or white flows that may well have been molten sulfur. It is possible that rising silicate magma may melt sulfur-rich rock as it nears the surface, producing ‘secondary’ sulfur flows (as opposed to ‘primary’ flows that originate from molten magmas at depth). This happened in Mauna Loa, Hawaii, around 1950, and produced a small sulfur flow. So far, the presence of sulfur flows on Io remains open, but it is clear that most of the volcanism is silicate because of the high temperatures detected. Inferring the composition of Io’s lavas has been problematic, as no high-resolution spectroscopic measurements could be taken. Most of Io’s surface is covered by sulfur dioxide frost which condensed from the volcanic plumes. Although it was the intent of Galileo scientists to take such measurements using the NIMS instrument, when the spacecraft got sufficiently close to Io to make observations of the relatively small areas not covered by sulfur dioxide, the instrument malfunctioned and lost its capability of obtaining enough wavelengths to make the necessary measurements. With no direct measurements of lava composition, temperatures detected at active hot spots provide the best clues to magma composition. Temperatures can be calculated by measurements made from two Galileo instruments, the Solid-State Imaging System (SSI) and the Near-Infrared Mapping Spectrometer (NIMS). The wavelength range used for the temperature determinations was 1.1-5.1 microns. However, even this is problematic, because of the rapid cooling of lava and the fact that the spatial resolution was of the order of kilometers. The temperature that characterizes the composition of the lava is the melt temperature, and most of the measurements will be dominated by cooled lavas, because the red-hot areas are small relative to the size of the pixels. Therefore, all the temperatures determined have to be considered minimum values. One of Io’s volcanoes, Pillan (Fig. 13), had a major and violent eruption in 1997 (42) that created a major change on Io’s surface. Initial measurements taken with the NIMS and SSI instruments showed temperatures of about 1500oC – too high for basalt. Basalts on Earth rarely exceed 1200oC. Some recent analysis of the observations of Pillan put the temperature lower – around 1300oC (43), still unusually high for basalts. As far as Io’s volcanism is concerned, this is probably the most intriguing aspect. Are all Ionian lavas similarly hot to Pillan’s? What are the compositions?

FIGURE 13. These images of Io’s Pele and Pillan volcanoes show the results of a dramatic event that occurred during a five-month period (April to September 1997). The changes, captured by the solid state imaging (CCD) system on NASA's Galileo spacecraft, show a new dark spot, 400 kilometers in diameter surrounding a volcanic center named Pillan Patera. Pillan's plume deposits appear dark at all wavelengths. This color differs from the very red color associated with Pele and are thought to be due to pyroclastics, while the red deposits are thought to be sulfur. North is to the top of the images. (NASA PIA00744).

At present, it is not known whether all lavas on Io are erupted at very high temperatures or not, since it is difficult to detect sufficiently large areas where very fresh (and therefore very hot) materials are exposed. It is also not known what is causing the very high temperatures. Two main hypotheses have been put forward: ultramafic volcanism and superheated basaltic volcanism (42). Ultramafic lavas, such as magnesium-rich komatiites, were erupted on Earth mostly billions of years ago. If the lavas on Io are of similar composition, we could be seeing there a style of volcanism similar to that in the Earth’s distant past. Alternatively, magmas can be superheated by the rapid ascent from a deep, high pressure source. Melting temperatures of dry silicate rock increase with pressure, therefore, the erupted lava can be significantly hotter than its melting temperature at surface pressure. Rapid ascent of basaltic magmas resulting in ~100 Celsius of superheating should be possible. However, no record of such an eruption is known on Earth and the superheating explanation is generally considered problematic. For example, if the ascent of the magma is slow, it will lose heat through the walls of the conduit. At present, the question of whether Io’s lavas are mafic or ultramafic remains open. There may be some other, more exotic but cooler lavas on Io. A possible magma composition on Io, though much rarer, is sulfur dioxide. When Galileo did repeated close flybys of Io, the NIMS instrument, though handicapped by the loss of many wavelengths, was able to map sulfur dioxide distribution at local scales (44). These results showed a nearly pure sulfur dioxide region topographically confined within a caldera (Fig. 14) called Baldur (40, 45). One possibility is that the sulfur dioxide was emplaced as a liquid flow rising from the subsurface. While liquid would normally boil off when exposed to Io’s near-vacuum atmosphere, it is estimated that, given sufficiently large quantities, some of the liquid could freeze to form a layer of sulfur dioxide ice inside the caldera. Only a couple of locations have so far been found on Io where sulfur dioxide may have erupted this way (40, 46).

FIGURE 14. Two paterae on Io showing unusual deposits. On the left is Chaac patera, showing green deposits that may consist of sulfur. On the right is Balder, which shows a nearly-pure SO2 ice deposit topographically confined within the patera walls. (NASA PIA02566).

The majority of Ionian eruptions can be placed in three classes (47) initially designated “Promethean,” “Pillanian,” and “Lokian,” though a single hot spot can exhibit more than one eruption style over time. Explosion- dominated (“Pillanian”) eruptions have an intense, short-lived phase that may correspond to the outbursts detected from Earth. These eruptions originate from either paterae or fissures and produce extensive dark lava flow fields and dark pyroclastic deposits through short-lived, high effusion rate, vigorous activity. These events may or may not include eruption episodes with large (>200 km high) explosive plumes, which can produce large plume deposits such as that around Pillan itself. Less intense but more persistent flow-dominated eruptions (“Promethean”) are named after the Prometheus hot spot, which has a persistent plume about 100 km high, active during both Voyager encounters in 1979, throughout the Galileo mission and spacecraft observations since then. Surprisingly, distant images obtained by Galileo in 1996 showed that the Prometheus plume site had moved about 80 km west since 1979 (48), but its size and appearance had not changed. A new lava flow linked the old and new plume sites. Images and infrared observations obtained in 1999 showed that the main vent of this volcano was near the Voyager plume site and that the plume, not the volcano, had moved west. The plume’s movement was modeled in terms of the interaction between the advancing hot lava and the underlying sulfur dioxide snowfield (49). This type of eruption may be common on Io and, once the flow stops moving, the plume eventually shuts off. The lava flows associated with these eruptions can be quite extensive and are thought to be emplaced through repeated small breakouts of lava, similar to the slowly emplaced pahoehoe flow fields at Kilauea in Hawaii. Intrapatera (“Lokian”) eruptions are confined within the caldera-like paterae. These eruptions are thought to be lava lakes (50), some of which are possibly overturning. Observations from Galileo flybys showed that lava lakes are abundant on Io (50), and they may be a significant mechanism for heat loss from the interior. Io’s most powerful hot spot, Loki, is thought to be a giant lava lake that perhaps undergoes periodic overturning (51), leading to brightenings that have been observed from Earth for decades. Many other hot spots on Io appear to be persistent lava lakes. Although Io’s surface features are dominated by its volcanic activity, tectonism has had a significant effect on the surface as well. Mountains are major structural landforms on Io and tower over the surrounding plains. Ionian mountains are defined as steep-sided landforms rising more than ~1 km over the plains. At least 115 mountains have now been identified and mapped (52). Io’s mountains rise, on average, about 6 km high, with the highest rising 17 km above the surrounding plains. Galileo images revealed that many mountains are partly or completely surrounded by debris aprons, plateaus, and layered plains. Mountains appear to be unstable and are thought to be relatively short-lived features. They are not active volcanoes, but their origin is still uncertain. Various models have been proposed to explain the origin of the mountains (e.g., 53, 54). Their asymmetrical shapes suggest the uplift and rotation of crustal blocks, implying that compressional uplift is probably the dominant mechanism.

FIGURE 15. Loki patera is thought to be a lava lake. On the left is an image obtained by Galileo’s SSI instrument, showing the subset area imaged by the NIMS instrument during a close fly-by of Io. On the right, the NIMS data are shown, the upper map shows the lavas inside Loki using the wavelength of 2.5 microns, in which hot areas appear as white and shades of red, and cooler areas in blues and greens. On the bottom is a thermal map with the temperature scale in Kelvin. (Based on NASA PIA02595 and results from (50)).

Erosion is a minor geologic process on Io, the very tenuous atmosphere does not result in significant interaction with the surface for significant erosion to happen. Between the paterae, mountains, and other major surficial features, Io’s surface appears smooth except for scarps that cut across the plains (55). Some scarps are linear and occur in parallel groups, which suggest a tectonic origin. Other scarps, however, are irregular and appear to be erosional, sometimes forming a series of mesas or large plateaus. The presence of erosional features on Io is somewhat puzzling because of the lack of a significant atmosphere or flow of liquid water. The eroding agent may be the explosive escape of sulfur dioxide from a subterranean SO2 “aquifer” (56). The idea of subterranean SO2 reservoirs gained some support after the analysis of Galileo data of the movement of the Prometheus plume (49). More recently, Io’s geology was summarized in a 1:15M-scale global geologic map (57), based on a set of monochrome and color Galileo-Voyager image mosaics produced at a spatial resolution of 1km/pixel. The surface of Io was mapped into 19 units based on albedo, color and surface morphology, in which additional paterae were identified, showing that they corresponded to 64% of all detected hot spots. This study also showed that 49% of all mountains are lineated and presumably layered, showing evidence of linear structures supportive of a tectonic origin. In contrast, only 6% of visible mountains are mottled (showing hummocks indicative of mass wasting) and 4% are tholi (domes or shields), consistent with a volcanic origin. The latest data on Io were obtained by the New Horizons spacecraft in 2007, which identified many surface changes since the last data from Galileo, and an additional hot spot (58).

TITAN: AN EARTH-LIKE WORLD

Titan is Saturn’s largest satellite and was discovered by the Dutch astronomer Christiaan Huygens on March 25, 1655. It has a mean radius of 2,576 km (2/5 of the Earth’s size) and orbits Saturn at a distance of 1.2 million km in a synchronous orbit. It is the second largest moon in size among all other satellites in our solar system, after Jupiter’s moon Ganymede. It takes almost 16 Earth days to complete a full rotation around its axis and one Titan year equals about 30 terrestrial years. Titan receives only 1% of the solar radiation received on Earth; thus, having a surface temperature of about -179oC. Due to its low atmospheric temperature, the water vapor saturation pressure is extremely low. Thus, any water vapor available appears limited to the stratosphere. Titan is the only natural satellite in our solar system that has a dense atmosphere other than Earth, Venus and Mars, making it a fascinating body to investigate. Despite its small size and weak gravitational field it has a dense atmosphere, unlike Callisto or Ganymede. Its atmosphere contains 98.4% nitrogen and favors photochemical reactions and organic chemistry processes, resembling Earth’s atmospheric environment. The second most abundant constituent is methane, which can be found in solid, liquid and gaseous form due to Titan’s temperature and pressure range. Fractionally, it fluctuates between 4.9% at the surface and 1.5% in the stratosphere and unlike molecular nitrogen it exhibits strong absorption bands in the infrared (IR). Titan formed such a thick atmosphere due to the preservation of a number of constituents that could have been present in the proto-planetary nebula like ammonia, argon, neon, molecular nitrogen and methane (59; 60). Of high importance is that the nitrogen, methane and hydrogen in the atmosphere trigger a positive greenhouse feedback that increases Titan’s surface temperature by trapping 90% of the incoming solar radiation, while at the same time absorbing the incoming IR radiation due to the presence of haze layers in the atmosphere (61), leading to a negative greenhouse feedback (anti-greenhouse effect). The positive and negative greenhouse feedbacks compete resulting in a net warming of Titan’s atmosphere. The interaction of Titan’s dense atmosphere with the surface leads to erosional processes being particularly important. Our knowledge about Titan’s atmospheric and geological conditions/processes come from ground-based and remote-sensing observations and modeling studies, providing critical information about and establishing the need for in situ exploration. The first spacecraft mission to observe Titan was Pioneer 11 on September 2, 1979. Pioneer 11 passed by Titan at a distance of 363,000 km and provided information regarding the moon’s radius. On November 12, 1980 Voyager 1 passed by Titan at a distance of 6,970 km (Fig. 16) and nine months later Voyager 2 flew by Titan at a distance of 663,400 km. Figure 16 shows Titan’s thick and hazy atmosphere that hides the surface.

FIGURE 16. – Titan as seen by Pioneer 11 (1979), Voyager 1 (1980), and Voyager 2 (1981). (a) First close up image by Pioneer 11 at a distance of 360,000 km on September 2, 1979. The only detection at that time was the possible existence of clouds (Image credit: NASA). (b) Image in visible light from the cameras of Voyager 1 from a distance of 4,000 km in 1980 showing again only Titan’s clouds (Image credit: NASA). (c) Voyager image of Titan, from a distance of 2.3 million km (August 23, 1981) (Image credit: NASA). At this point, the extended haze, composed of submicron-size particles, surrounds the satellite's limb (Image configuration: A. Solomonidou).

In addition to the optical images shown in Fig. 16, Voyager provided information about the orbital and physical characteristics of Titan (59). As mentioned earlier, Titan’s surface radius was estimated about 2,575 ± 2 km, with a surface temperature of –179 ± 2 oC and a pressure of about 1.44 bar; these values were confirmed by the Cassini- Huygens measurements (e.g., 62). After Pioneer 11 and Voyager missions, the Canada/France Hawaii Telescope (CFHT), the Very Large Telescope (VLT), the Hubble Space Telescope (HST), the Infrared Space Observatory (ISO) and the Keck Telescope, providing complementary observational information about Titan’s physical properties (e.g., 63). The Cassini-Huygens mission, a collaborative mission between the European Space Agency (ESA) and the National Astronautics and Space Administration (NASA), involving 17 countries, arrived at Titan 25 years after the Voyager missions providing the first in-situ image of a planetary body of the outer solar system (Figure 17). Cassini- Huygens’ mission provided crucial information about Titan’s surface such as, geology and forming processes. In addition, the Huygens probe returned the first in situ surface image of a planetary body of our outer solar system (Figure 17).

FIGURE 17. Titan’s ground ‘truth’. The first outer solar system in situ image. Small "rocks," possibly made of water ice, at the Huygens landing site. The right panel shows, approximately, the true color of the scene (Image credit: 64) while the left one is in gray scale and shows the approximate sizes of the pebbles. Evidence of erosion at the base of these objects suggests fluvial activity. Solar illumination is obvious on the top left of the images while the brightening in the lower right side is lamp illumination from DISR (Image credit: ESA/NASA/JPL/University of Arizona).

The cameras on board the Cassini orbiter, with the capability to partially or thoroughly penetrate Titan’s thick atmosphere, captured the first surface images (e.g., 65; 66; 67). These observations confirmed previous ground based observations concerning the presence of an extensive large region located at 10ºS and 100ºW officially named Xanadu, which is a quasi-circular and continent-sized albedo feature (e.g., 65; 66). This region is located close to Titan’s equator where the mid-latitude regions were found as uniformly dark with elevated topography (Fig. 18). On the other hand, the poles were seen as bright filled with large dark patches characterized as hydrocarbon lakes, which were deeply investigated by the Cassini/RADAR (e.g., 68). Cassini found that Titan has an interesting topography with mountainous regions ranging from 200 m to 2 km altitude (e.g., 69), while the poles are topographically lower than the equator (e.g., 66). The Distant Imager Spectral Radiometer (DISR) data also led to such distinct characteristics concerning the average hypsometry (altitude) of the dark (lows) and bright regions (highs).

FIGURE 18. Titan’s globe from images acquired by the Cassini Image and Science Subsystem (ISS) camera during the first close flyby over Titan (October 26, 2004). The brighter region on the right side and equatorial region is named Xanadu Regio. On the southern pole, bright clouds are seen (Image credit: NASA/JPL/Space Science Institute/PIA06141).

Cassini’s Visual and IR Mapping Spectrometer allowed us to classify different terrains depending on their spectral behavior (70). The RADAR instrument gives information on the terrains’ shape and morphology, and only in a few cases on their hypsometry. Nevertheless, the Cassini-Huygens mission was the first to report that Titan’s terrain morphology resembles that of Earth’s. Cassini radar recorded several different types of geological features such as mountains, ridges, faults, rectangular drainage patterns and cryovolcanic structures, which are indicative of volcanic and tectonic activity (69; 70; 71; 72). Stable liquid lakes were spotted by both ISS and RADAR, located mainly at the polar regions (e.g., 68), while recently similar lakes were identified at the equatorial latitudes (73) (Fig. 19). Titan is the only extra-terrestrial body to have liquids currently on the surface.

FIGURE 19. Major types of geological features and terrains on Titan from Cassini data characterized as fluvial, aeolian, tectonic, and cryovolcanic. The background map is a mosaic of Cassini visual and infrared data (Map credit: NASA/JPL/Space Science Institute; Image configuration: A. Solomonidou).

Titan’s surface composition has been a mystery even during the Cassini era. Among the constituents we mentioned earlier, water ice, CO2 ice and hydrocarbons were suggested as possible candidates that reside on the surface (e.g., 74; 75). The origin of such materials has been attributed to various sources (e.g., interior) and triggered or deposited by various processes such as fluvial erosion and cryovolcanism. Titan is a dynamic environment that undergoes both short- and long-term variations, due to fluvial, aeolian, atmospheric and possibly endogenic (tectonic and/or volcanic) processes whose signatures can be identified on Titan’s surface (Fig. 20). It is worth noting that the Cassini instruments have to date found no obvious evidence for densely cratered terrains on Titan’s surface, and relatively few features have been interpreted as impact craters. This may suggest that the surface of Titan is young (less than a billion years old) or, highly eroded/resurfaced (76; 77).

FIGURE 20. Surface features on Titan. (a) Impact crater (NASA/JPL-Caltech); (b) Liquid lakes (radar-dark) (Image credit: NASA/JPL-Caltech/ASI). (c) Large fluvial deposit with radar-bright channels (Image credit: Lorenz et al. 2008). (d) Radar dark sand dunes formed by aeolian processes and radar-bright topographic obstacles (Image credit: NASA/JPL). (e) Lobate radar- bright feature and possible cryovolcanic deposits (Image credit: Soderblom et al. 2009). (f) Parallel radar-bright mountain chain (Image credit: NASA/JPL-Caltech) (Image configuration: A. Solomonidou).

Titan has an active methane cycle similar to Earth’s water cycle, characterized by precipitation, formation of channels and evaporation of lakes (e.g., 83) (Fig. 20). The presence of methane is significant in Titan’s thick nitrogenous atmosphere, but its lifetime is about 10-30 Myr, as it is destroyed by photochemical processes and converted to heavier hydrocarbons (84). Because the methane concentration is almost constant in Titan’s atmosphere, there have to be other methane sources to make up for its loss. Previous theories propose that cryovolcanism and lake evaporation are possible methane sources (85; 86). Cryovolcanic processes, along with their influences in Titan’s environment, have been suggested; although, they are considered controversial. The strongest evidence we currently have that support this theory are based on studies of the morphology and spectral emissions of the Sotra Patera region and Hotei Regio (78; 79; 81) and their variability over time (70; 80). Fluvial activity has also been suggested as a geological process that can modify Titan’s surface, based on Cassini images of drainage networks, the first in-situ images of hills (64), and the large mountain chains and fault block mountains (82) (Fig. 20). Other surface terrains observed by the Cassini orbiter include areas covered with dunes similar to the terrestrial ones (e.g., 87) (Fig. 20). The RADAR instrument on board Cassini has also discovered lakes sprinkled over the high northern latitudes of Titan (e.g., 68) – a geological feature that is of crucial importance on the hypothesis of the missing liquid methane or ethane surface reservoir.

FIGURE 21. Schematic view of the dynamic internal-surface-atmospheric processes and the methane cycle of Titan from evidence of Cassini-Huygens data (Image credit: 88). Methane (CH4) is released into the atmosphere from Titan's interior stores through volcanic action, and evaporates from the lakes of methane and ethane (C2H6) identified by the Cassini spacecraft on the satellite's surface. Chemical reactions in the atmosphere convert it to ethane; complex organic aerosols consisting of carbon, hydrogen and nitrogen; and hydrogen gas (H2), which escapes into space. Ethane and methane partly condense, forming clouds and hazes that precipitate, replenishing the lakes and bearing many organic species in solution.

Titan’s continuous changing surface can be seen through the recent and past lake-level changes, changes of fluvial deposits in the dunes, and possibly on the cryovolcanic sites (80; 89; 90; 91). The presence of all the aforementioned surface features and the processes that form them shows that Titan is an active world and changes on the surface environment are expected to occur. Endogenic processes (tectonism and cryovolcanic) may be present on Titan, though an endogenic origin for candidate features is still open for debate. A number of tectonic patterns have been discussed for Titan, although much different than the terrestrial ones, as indicated by the bright elevated crustal features on Cassini/RADAR images (Fig. 20f) (69; 92; 93). These patterns include compressional (i.e., shortening and thickening of the crust/lithosphere) extensional tectonism (i.e., stretching of the crust/lithosphere). Extensional tectonism is common and observed in almost all icy moons in the Solar system but this is not the case for contractional/compressional. Thus, if the formation of mountains on Titan is triggered due to crustal compressions then Titan is the only icy satellite where contractional tectonism occurs and is the predominant style (92). Cryovolcanism has also been suggested as another methane source on Titan, although it has not yet been proven and other processes may have created the candidate features. Pre-Cassini models predicted the activity of cryovolcanism on Titan and possible cryovolcanic features include lobate flows (Fig. 22 for the strongest cryovolcanic candidate on Titan namely Sotra Patera) and possible calderas or pits (e.g. 71; 79). However, alternative hypotheses based on spectral similarities of some cryovolcanic candidate features with dry lakebeds at the poles, rank them as evaporitic or fluvial/lacustrine deposits with no interior connection (94; 95; 96).

FIGURE 22. The Sotra Patera area, currently the strongest cryovolcanic candidate on Titan. The area consists of on high peak named Doom Mons (1,450 m), the additional peak a crater Erebor Mons, the crater-depression Sotra Patera (1,700 m) and the finger-like flows named Mohini Fluctus (Image credit: NASA/JPL-Caltech/ASI/USGS/University of Arizona, PIA13695).

The triggering mechanism that leads to dynamic processes, such as tectonism and cryovolcanism is possibly Saturn’s tides. The maximum shear stresses are concentrated in the polar regions, while the maximum tensile stresses predominate in the near-equatorial and mid-latitude regions of the sub- and anti-saturnian hemispheres (e.g., 97). Most of the geophysical models that predict the activity of tectonism and cryovolcanism on Titan require the presence of a subsurface liquid water layer/ocean (e.g., 98). The Cassini data recordings suggested that an undersurface liquid water ocean exists on Titan. The Huygens Atmospheric Structure Instrument - Permittivity Wave and Altimetry (HASI-PWA) experiment indicated the presence of a Schumann resonance1 between the ionosphere and an internal ocean that lies 50 km beneath the surface, through observations made by the extremely low-frequency electric signal recorded (99). Furthermore, gravity measurements performed by the Cassini Radio Science experiment showed that the gravitational field significantly varies along Titan’s terrain, indicating that Titan is subjected to tides of about 10 m in height. Such strong tides confirm the deformability of the interior, while in addition to Titan’s obliquity, confirm the presence of an internal liquid ocean roughly 100 km below the surface (100; 101). In addition, thermal evolution models testable with Cassini–Huygens measurements suggest that Titan may have an icy shell 50 to 150 km thick, lying atop a liquid water ocean a couple of hundred kilometers deep, with some amount (up to ~10%) of ammonia dissolved in it, acting as an antifreeze, in addition to a layer of high pressure ice that could exist underneath (e.g., 98) (Fig. 23).

1 The Schumann Resonances are quasi-standing electromagnetic waves that exist in the Earth's 'electromagnetic' cavity (the space between the surface of the Earth and the Ionosphere).

FIGURE 23. Artistic view of a possible internal layering in Titan as hypothesized by Cassini data (Image credit: A. Tavani/101).

Current Titan interior models primarily consider the presence of a subsurface liquid water layer in the satellite while some of suggest the presence of methane clathrates in the ice shell in Titan’s interior. The models take into account heat transfer, convection, tidal dissipation, clathrate dissociation and cooling of the subsurface ocean, suggest cryovolcanic and tectonic phenomena as a possible ‘pathway’ for methane to resupply the atmosphere through outgassing, and possibly for water and ammonia to be ejected covering patches of the surface around cryovolcanic sites on the other (e.g. 97; 98). According to these models, cryovolcanism and deposition of water and ammonia ice on the surface of Titan is plausible. Additional theories for the methane replenishment in the atmosphere exist; although they are less favorable, because they explain only a portion (about 1%) of the methane in the atmosphere (87). Two of these processes include impacts (103) and methane sources from surface lakes. Given the information we currently have, Titan is a good candidate for habitability and astrobiological studies as its atmospheric composition and pressure has similar properties with that of the Earth’s today (and even more so in the past) than any other Solar system body (e.g., 60). Recent Cassini-Huygens findings revolutionized our understanding of Titan’s environment and its potential for harboring the “ingredients” necessary for life such as, a rich organic budget and sufficient energy sources to initiate and sustain organic chemical evolution (104). The environment on Titan’s surface is not too distinct from the natural environment we experience here on Earth, especially the geological processes and the interactions of the Earth’s interior with its surface and its atmosphere (Figs. 24 and 25). This makes Titan an excellent terrestrial analogue, and studying its dynamic environment allows to delve into physical processes that occur on Earth and other planetary bodies.

FIGURE 24. ‘Twin’ lakes on Titan (left) and Earth (right). Ontario Lacus on Titan is considered a depression that drains and refills from below, exposing liquid areas ringed by materials such as saturated sand or mudflats. The Etosha salt pan, Kalahari basin, North Namibia on Earth, that shares some common morphological characteristics with Ontario Lacus, is a lake bed that fills with a shallow layer of water from groundwater levels that rise during the rainy season (Image credit: NASA/JPL-Caltech and NASA/USGS; information credit: 105).

FIGURE 25. Huygens’ landing site on Titan (left) resembles terrestrial terrains (right). The image on the right portrays a part of the Firiplaka coastal side in Milos Island, Cyclades, Greece (Image credit: left- ESA/NASA/University of Arizona; right- A. Solomonidou).

The Cassini-Huygens mission, an extremely successful mission, has enormously advanced our knowledge of the Saturnian system and the satellites in orbit. Future planetary investigations will use the Cassini-Huygens mission scientific findings and engineering developments as reference points. However, the key contribution to planetary science of Cassini may be the questions raised rather than those answered. Without a doubt, a new mission focusing on the seasonal, astrobiological and geological aspects of several of the satellites of Saturn is needed in order to improve our understanding of these unique planetary bodies.

ACKNOWLEDGMENTS

This work was carried out at the Jet Propulsion Laboratory, California Institute of Technology, under contract with NASA.

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