See the martian surface with major features (MOLA Digital Elevation Model)
See the martian surface interactive animation (JAVA required, 220k download)
The Main Division, which is called the crustal dichotomy, on the martian surface is between the ancient highlands and the northern lowlands. Most of the lowlands lie below about -1.5 km elevation and have a very flat surface. There is a gentle planetary wide slope (less than 0.05o) from the ancient highlands down to the north pole. This slope has defined the main drainage pattern of Mars, with ancient floodwaters channelled towards the northern lowlands. The highlands cover approximately 60% of the planet's surface, mainly in the southern hemisphere, and have the highest cratering density on the planet. Extreme topographic variation is a feature of the ancient highlands which contain both high volcanoes and deep impact craters. The Hellas basin (2300 km diameter and 8 km deep) is the largest easily identifiable impact crater and the deepest area on Mars. Isidis Planitia (1600 km diameter) is also the relict of an ancient crater that has undergone subsequent infilling with sediments and volcanic products.
The second major defining feature of Mars's shape is the Tharsis bulge, which is the region of highest ground on the planet, centred around 0o latitude, 105oW. Tharsis covers approximately 10% of the martian surface and is superimposed on the crustal dichotomy between the northern lowlands and ancient highlands. It has average slopes of about 0.2-0.4o downslope to the north and less to the south, where Tharsis merges with the ancient highlands. The Tharsis bulge is covered with radial fractures (these extend to around one hemisphere of the planet) and 3 large shield volcanoes - most notably Olympus Mons which is 21 km high, the largest volcano in the Solar System and nearly 2.5 times the height of Mt. Everest. On the eastern side of Tharsis is the Valles Marineris Canyon system, which extends for 4500 km and is up to 7 km deep (by comparison the Grand Canyon of Arizona, USA is 800 km long and up to 1.6 km deep).
Martian stratigraphy has been deduced from geologic units defined on photographs with ages based on the density of craters with diameter < 1 km. From oldest to youngest, the 3 main divisions are the Noachian (which lasts up to the end of the Late Heavy Bombardment), the Hesperian (during which the Chryse outwash channels formed) and the Amazonian.
Northern Polar Layered Deposits and Ice Cap. MOC image, Malin Space Science Systems.
Layered deposits are present at both poles beyond 80o latitude. They have a notably flat, crater-free surface and cover older, cratered surfaces. They are composed of ice mixed with dust deposits, are seasonally covered by the polar caps and are geologically young, having formed in the late Amazonian. In the north they are up to 4-6 km thick. The figure below shows the northern polar region. The swirl pattern on the actual ice cap (1100 km diameter) is caused by preferential removal of frost on south-facing slopes. The dark band around the ice cap is dunes and layered deposits formed by deposition of volatiles and dust. These deposits have a climatic control e.g during periods of high obliquity 25-35o more dust and sand is released and smaller remnant ice caps occur. This also leads to greater atmosphere pressures. Massive dune fields occur just south of the Northern Layered Polar Deposits. Dune fields are also commonly present within impact craters across Mars and scattered dunes are present across much of the martian surface.
Recently there has been renewed interest in the occurrence of glaciation on the surface of Mars. It has been suggested that gullies within the Olympus Mons shield volcano contain glaciers or have contained glaciers in the past. If glaciers do exist today then they must be protected by metres-thick impermeable layers of dust or other deposits to stop the ice subliming in the low atmospheric pressure. The image to right shows a THEMIS (Mars Odyssey/ASU) image of possible relict glacial terrain on Arsia Mons. The parallel ridges may represent successive moraines left by the retreat of a glacier.
Possible glacial moraine on Arsia Mons. 25 km width, 7oS, 130oW THEMIS/ASU
Tiu Vallis Outflow Channel headwaters and chaotic terrain
The two main types of channels on Mars are runoff and outflow. Runoff channels are relatively small (e.g. less than 20 km long) and often found around crater rims in the highlands. Most of these are thought to have formed when groundwater occasionally reached the surface.
Outflow channels (e.g. Ares Vallis where Pathfinder landed in 1997) are associated with catastrophic floods on a much larger scale. The floods may have started through the melting of ice due to igneous activity. Chaotic zones of disrupted ground are found at the headwaters of these channels. The figure above shows the chaotic zones at the headwaters of the Tiu Vallis channel, south of Chryse Planitia (scale - one degree on Mars is approx. 60 km). Floodplain materials were deposited along outflow channels and where the channels grade into the northern plains.
Layered Deposits 0.15N, 345.6W. These layered deposits were photographed in the Schiaparelli Crater of the ancient highlands. These may have been deposited by crater lakes with subsequent erosion by the wind, or alternatively, they may be exposures of the underlying crust. MOC image EO3-00728, width 3km(MSSS)
Panoramic images taken at the Viking 1 and Pathfinder landing sites show angular, unsorted rocks that were deposited rapidly from floodwaters as they spread across the Chryse Basin. The floods were of enormous scale, with discharges sometimes twice as great as the largest known floods in the geological record of the Earth. Another feature on the martian surface associated with the channelling and ponding of floodwaters is Valles Marineris. The Valles Marineris trough formed on the east side of the Tharsis rise as a result of doming and fracturing of the ground surface. A series of parallel canyons formed along graben faults parallel to the main canyon. Floodwaters later exploited these troughs and were channeled towards the northern plains and the Chryse Basin.
Layered sediments were laid down within the interior of Valles Marineris. Some impact craters also show layered deposits and these have been taken to indicate the past presence of lakes (see image right of layered deposits in Schiaparelli Crater). The presence of channels leading into such craters is also consistent with the ponding of water.
The outflow channels have a range of ages (mainly Hesperian-Amazonian) younger than those of the runoff channels (most of which are Noachian to early Hesperian).
However, a few channels have been identified within mid latitude areas of the southern hemisphere which have a very low density of superimposed impact craters, suggesting that they formed within geologically recent times.
Geologically recent gullies on Mars. 40oS, 252oW, MOC04-02479, (MSSS).
Results from the MOC camera (see image to right) have shown that water may have flowed on Mars within the last 106 years, creating gullies (Malin and Edgett, Science, 2000). These 'fluid seepages' are hundreds of metres long and contain a 'head alcove', main channels and depositional fans. They are located in impact crater walls, other pits or valley walls. It has been proposed that the presence of salts could lower the melting point of near-surface ice sufficiently so that orbitally induced variations (ie obliquity cycles, see (See Physical Parameters page)) in surface temperatures might allow the presence of occasional brines. Concentrated brines are able to exist in a liquid state at temperatures of down to -60oC , and as mean surface temperatures on Mars are typically in the range -33oC to -68oC between 60oN and 60oS, it seems that, if as is generally now presumed, these features are formed by flowing water, the fluids were brines. The alternative to liquid water is liquid CO2 (e.g. Musselwhite et al. 2001, Geophys. Res. Lett.). This has the apparent advantage of explaining the existence of fluid activity at the very low temperatures on the current martian surface. However, in an investigation of this theory Stewart and Nimmo (LPSC, 2001) noted that a large reservoir of CO2; ice would be necessary in order to provide the large volumes of fluid necessary to create the gullies. Large reservoirs of pure CO2 ice are prohibited under current martian conditions in the top few metres of crust because they would sublime < 100 m from the martian surface. The occasional action of liquid (brine?) water stands as the consensus opinion for the formation of the gullies.
This MOC image, 3 km width, shows a youthful lava flow with a low crater density that has partially covered an older (lighter), cratered plain. Many such lava flows are found in the Elysium region (MOC 010111b , MSSS).
Much of the northern lowlands are underlain by ridged lava flows similar to the lunar mare. This is shown by the presence of wrinkle ridges. For instance, the Beagle 2 landing site in Isidis Planitia shows subdued wrinkle ridges which are overlain by later sedimentary deposits (See Mars Missions page). The surface rocks in the northern plains include lava flows, pyroclastic (volcanic ash) deposits, dune fields and flood deposits from outflow channels. These deposits have obscured large basins within the northern plains; Chryse, Isidis and Utopia Planitia are ancient (early Noachian) impact basins that have been partially filled. Other features include polygonal fracture patterns up to 20 km across - which may be related to tensional forces resulting from ice freezing and subliming. The crater density on some of the lava flows in the Elysium Planitia region of the northern plains is low enough to indicate that the lavas flowed less than 100 million years ago (Ma) and in some cases could be less than 10 Ma.
These are mainly in the Elysium region and are composed of three main shield volcanoes of which Elysium Mons is the oldest and largest (9 km elevation).
Photographs taken of the sides of channels and Valles Marineris (see Figure to right) show that the ancient highlands are largely composed of metre to tens of metre thick layers. These are generally regarded as lava flows although until they are actually sampled it cannot be ruled out that they include sedimentary layers (e.g. sandstones, limestones). We now know from the results of the Opportunity MER lander (See Mars Missions page) that some sedimentary and evaporite layers do exist on Mars from Noachian rock units. Thus some of the layers we see in Valles Marineris and in other parts of ancient Mars could be sedimentary rocks. This in turn is direct evidence for the warmer and wetter past that many geologists have long thought to have existed at some parts of martian history
This image (a montage of Viking images) shows a typical part of the ancient highlands. It is highly cratered with small runoff channels incising the surface. Image centred at 20oS, 270oW, 450 km dimensions (Carr, Water on Mars, Viking frame 625A27)
The layering within the upper martian crust on a scale of metres to tens of metres is visible on the sides of Coprates Chasma (part of the Valles Marineris system). (MOC image 8003, MSSS, width 9.8 km).>
This is the main Tharsis volcanic region and includes Olympus Mons (600 km in diameter, 21 km height, see wide angle MOC image to right). A linear chain of shield volcanoes extends NE across Tharsis Montes. The Hawaiian shield volcanoes are the most similar terrestrial volcanoes. Much of the volcanism in Tharsis dates from the Hesperian although the region has been volcanically active from the Noachian to Amazonian.
The Hellas crater is the dominant structural feature of the southern highlands and is 2300 km diameter. It has been degraded by runoff channels and also contains lavas and dune fields.
Parts of the ancient highlands contain ridged plains material ie deformed lava flows, similar to those in the northern plains. This formation also includes the Highland Paterae. Many of these are present near the Hellas Basin and are low relief, degraded volcanoes or impact craters. The largest one is however, Alba Patera which is an eroded shield volcano on the north west margin of Tharsis.
Chaotic areas e.g. Noctis Labyrinthus at the head of Valles Marineris are composed of km-sized jumbled blocks formed by the break-up of plateau rocks related to the source of floodwaters.
This is a wide angle, oblique image of the Olympus Mons shield (MSSS, MOC26301/2 red/blue colour filters).
Valles Marineris - image constructed from MOLA topographic data (1/128o data set. V. Marineris is 4500 km long and up to 7 km deep. Image from Adrian Lark (www.mars3d.co.uk.) A high resolution MOC image of the valley sides (Coprates Chasma) can be seen above.