Generalised geological map of Mars
Mars as seen by the Hubble Space Telescope
The geology of Mars is the scientific study of the surface, crust, and interior of the planet Mars.
It emphasizes the composition, structure, history, and physical
processes that shape the planet. It is analogous to the field of
terrestrial geology. In planetary science, the term geology is used in its broadest sense to mean the study of the solid parts of planets and moons. The term incorporates aspects of geophysics, geochemistry, mineralogy, geodesy, and cartography. A neologism, areology, from the Greek word Arēs (Mars), sometimes appears as a synonym for Mars's geology in the popular media and works of science fiction (e.g. Kim Stanley Robinson's Mars trilogy).
Geological map of Mars (2014)
Mars - geologic map (USGS; July 14, 2014) (full image)
Composition of Mars
Mars is a differentiated, terrestrial planet. The InSight lander mission is designed to study the deep interior of Mars. The mission landed on 26 November 2018, and will deploy a sensitive seismometer that will enable 3D structure maps of the deep interior.
Global physiography
Most of our current knowledge about the geology of Mars comes from studying landforms and relief features (terrain) seen in images taken by orbiting spacecraft.
Mars has a number of distinct, large-scale surface features that
indicate the types of geological processes that have operated on the
planet over time. This section introduces several of the larger
physiographic regions of Mars. Together, these regions illustrate how
geologic processes involving volcanism, tectonism, water, ice, and impacts have shaped the planet on a global scale.
Hemispheric dichotomy
Mars
Orbital Laser Altimeter (MOLA) colorized shaded-relief maps showing
elevations in the western and eastern hemispheres of Mars. (Left): The
western hemisphere is dominated by the Tharsis
region (red and brown). Tall volcanoes appear white. Valles Marineris
(blue) is the long gash-like feature to the right. (Right): Eastern
hemisphere shows the cratered highlands (yellow to red) with the Hellas basin
(deep blue/purple) at lower left. The Elysium province is at the upper
right edge. Areas north of the dichotomy boundary appear as shades of
blue on both maps.
The northern and southern hemispheres of Mars are strikingly different from each other in topography and physiography. This dichotomy
is a fundamental global geologic feature of the planet. Simply stated,
the northern part of the planet is an enormous topographic depression.
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.
This is a first-order relief feature on par with the elevation
difference between Earth's continents and ocean basins.
The dichotomy is also expressed in two other ways: as a difference in
impact crater density and crustal thickness between the two hemispheres.
The hemisphere south of the dichotomy boundary (often called the
southern highlands or uplands) is very heavily cratered and ancient,
characterized by rugged surfaces that date back to the period of heavy bombardment.
In contrast, the lowlands north of the dichotomy boundary have few
large craters, are very smooth and flat, and have other features
indicating that extensive resurfacing has occurred since the southern
highlands formed. The third 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 (36 mi), whereas crust in the northern lowlands "peaks" at
around 32 km (20 mi) in thickness.
The location of the dichotomy boundary varies in latitude across Mars
and depends on which of the three physical expressions of the dichotomy
is being considered.
The origin and age of the hemispheric dichotomy are still
debated. Hypotheses of origin generally fall into two categories: one,
the dichotomy was produced by a mega-impact event or several large
impacts early in the planet's history (exogenic theories)
or two, the dichotomy was produced by crustal thinning in the northern
hemisphere by mantle convection, overturning, or other chemical and
thermal processes in the planet's interior (endogenic theories). One endogenic model proposes an early episode of plate tectonics producing a thinner crust in the north, similar to what is occurring at spreading plate boundaries on Earth. Whatever its origin, the Martian dichotomy appears to be extremely old. A new theory based on the Southern Polar Giant Impact and validated by the discovery of twelve hemispherical alignments shows that exogenic theories appear to be stronger than endogenic theories and that Mars never had plate tectonics
that could modify the dichotomy. Laser altimeter and radar sounding
data from orbiting spacecraft have identified a large number of
basin-sized structures previously hidden in visual images. Called
quasi-circular depressions (QCDs), these features likely represent
derelict impact craters from the period of heavy bombardment that are
now covered by a veneer of younger deposits. Crater counting studies of
QCDs suggest that the underlying surface in the northern hemisphere is
at least as old as the oldest exposed crust in the southern highlands. The ancient age of the dichotomy places a significant constraint on theories of its origin.
Tharsis and Elysium volcanic provinces
Straddling the dichotomy boundary in Mars's western hemisphere is a massive volcano-tectonic province known as the Tharsis
region or the Tharsis bulge. This immense, elevated structure is
thousands of kilometers in diameter and covers up to 25% of the planet's
surface.[25]
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 enormous volcanoes, Ascraeus Mons, Pavonis Mons, and Arsia Mons (collectively known as the Tharsis Montes), sit aligned NE-SW along the crest of the bulge. The vast Alba Mons (formerly Alba Patera) occupies the northern part of the region. The huge shield volcano Olympus Mons lies off the main bulge, at the western edge of the province. The extreme massiveness of Tharsis has placed tremendous stresses on the planet's lithosphere. As a result, immense extensional fractures (grabens and rift valleys) radiate outward from Tharsis, extending halfway around the planet.
A smaller volcanic center lies several thousand kilometers west of Tharsis in Elysium. The Elysium volcanic complex is about 2,000 kilometers in diameter and consists of three main volcanoes, Elysium Mons, Hecates Tholus, and Albor Tholus.
The Elysium group of volcanoes is thought to be somewhat different from
the Tharsis Montes, in that development of the former involved both
lavas and pyroclastics.
Large impact basins
Several enormous, circular impact basins are present on Mars. The largest one that is readily visible is the Hellas basin
located in the southern hemisphere. It is the second largest confirmed
impact structure on the planet, centered at about 64°E longitude and
40°S latitude. The central part of the basin (Hellas Planitia) is
1,800 km in diameter and surrounded by a broad, heavily eroded annular rim structure characterized by closely spaced rugged irregular mountains (massifs), which probably represent uplifted, jostled blocks of old pre-basin crust. Ancient, low-relief volcanic constructs (highland
paterae) are located on the northeastern and southwestern portions of
the rim. The basin floor contains thick, structurally complex
sedimentary deposits that have a long geologic history of deposition,
erosion, and internal deformation. The lowest elevations on the planet
are located within the Hellas basin, with some areas of the basin floor
lying over 8 km below datum.
The two other large impact structures on the planet are the Argyre and Isidis
basins. Like Hellas, Argyre (800 km in diameter) is located in the
southern highlands and is surrounded by a broad ring of mountains. The
mountains in the southern portion of the rim, Charitum Montes, may have been eroded by valley glaciers and ice sheets at some point in Mars's history.
The Isidis basin (roughly 1,000 km in diameter) lies on the dichotomy
boundary at about 87°E longitude. The northeastern portion of the basin
rim has been eroded and is now buried by northern plains deposits,
giving the basin a semicircular outline. The northwestern rim of the
basin is characterized by arcuate grabens (Nili Fossae) that are circumferential to the basin. One additional large basin, Utopia,
is completely buried by northern plains deposits. Its outline is
clearly discernable only from altimetry data. All of the large basins on
Mars are extremely old, dating back to the late heavy bombardment. They
are thought to be comparable in age to the Imbrium and Orientale basins on the Moon.
Equatorial canyon system
Viking Orbiter 1 view image of Valles Marineris.
Near the equator in the western hemisphere lies an immense system of
deep, interconnected canyons and troughs collectively known as the Valles Marineris.
The canyon system extends eastward from Tharsis for a length of over
4,000 km, nearly a quarter of the planet's circumference. If placed on
Earth, Valles Marineris would span the width of North America. In places, the canyons are up to 300 km wide and 10 km deep. Often compared to Earth's Grand Canyon,
the Valles Marineris has a very different origin than its tinier,
so-called counterpart on Earth. The Grand Canyon is largely a product of
water erosion. The Martian equatorial canyons were of tectonic origin,
i.e. they were formed mostly by faulting. They could be similar to the East African Rift valleys. The canyons represent the surface expression of powerful extensional strain in the Martian crust, probably due to loading from the Tharsis bulge.
Chaotic terrain and outflow channels
The
terrain at the eastern end of the Valles Marineris grades into dense
jumbles of low rounded hills that seem to have formed by the collapse of
upland surfaces to form broad, rubble-filled hollows. Called chaotic terrain, these areas mark the heads of huge outflow channels that emerge full size from the chaotic terrain and empty (debouch) northward into Chryse Planitia. The presence of streamlined islands and other geomorphic features indicate that the channels were most likely formed by catastrophic releases of water from aquifers
or the melting of subsurface ice. However, these features could also be
formed by abundant volcanic lava flows coming from Tharsis. The channels, which include Ares, Shalbatana,
Simud, and Tiu Valles, are enormous by terrestrial standards, and the
flows that formed them correspondingly immense. For example, the peak
discharge required to carve the 28-km-wide Ares Vallis is estimated to
have been 14 million cubic metres (500 million cu ft) per second, over
ten thousand times the average discharge of the Mississippi River.
Mars Orbital Laser Altimeter (MOLA) derived image of Planum Boreum. Vertical exaggeration is extreme. Note that residual ice cap is only the thin veneer (shown in white) on top of the plateau.
Ice caps
The polar ice caps are well-known telescopic features of Mars, first identified by Christiaan Huygens in 1672.
Since the 1960s, we have known that the seasonal caps (those seen in
the telescope to grow and wane seasonally) are composed of carbon
dioxide (CO2) ice that condenses out of the atmosphere as temperatures fall to 148 K, the frost point of CO2, during the polar wintertime. In the north, the CO2 ice completely dissipates (sublimes) in summer, leaving behind a residual cap of water (H2O) ice. At the south pole, a small residual cap of CO2 ice remains in summer.
Both residual ice caps overlie thick layered deposits of
interbedded ice and dust. In the north, the layered deposits form a
3 km-high, 1,000 km-diameter plateau called Planum Boreum. A similar kilometers-thick plateau, Planum Australe,
lies in the south. Both plana (the Latin plural of planum) are
sometimes treated as synonymous with the polar ice caps, but the
permanent ice (seen as the high albedo, white surfaces in images) forms
only a relatively thin mantle on top of the layered deposits. The
layered deposits probably represent alternating cycles of dust and ice
deposition caused by climate changes related to variations in the
planet's orbital parameters over time. The polar layered deposits are some of the youngest geologic units on Mars.
Geological history
Albedo features
Mollweide
projection of albedo features on Mars from Hubble Space Telescope.
Bright ochre areas in left, center, and right are Tharsis, Arabia, and
Elysium, respectively. The dark region at top center left is Acidalium
Planitia. Syrtis Major is the dark area projecting upward in the center
right. Note orographic clouds over Olympus and Elysium Montes (left and right, respectively).
No topography is visible on Mars from Earth. The bright areas and dark markings seen through a telescope are albedo features. The bright, red-ochre
areas are locations where fine dust covers the surface. Bright areas
(excluding the polar caps and clouds) include Hellas, Tharsis, and Arabia Terra.
The dark gray markings represent areas that the wind has swept clean of
dust, leaving behind the lower layer of dark, rocky material. Dark
markings are most distinct in a broad belt from 0° to 40° S latitude.
However, the most prominent dark marking, Syrtis Major Planum, is in the northern hemisphere. The classical albedo feature, Mare Acidalium (Acidalia Planitia),
is another prominent dark area in the northern hemisphere. A third type
of area, intermediate in color and albedo, is also present and thought
to represent regions containing a mixture of the material from the
bright and dark areas.
Impact craters
Impact craters were first identified on Mars by the Mariner 4 spacecraft in 1965.
Early observations showed that Martian craters were generally shallower
and smoother than lunar craters, indicating that Mars has a more active
history of erosion and deposition than the Moon.
In other aspects, Martian craters resemble lunar craters. Both are products of hypervelocity impacts
and show a progression of morphology types with increasing size.
Martian craters below about 7 km in diameter are called simple craters;
they are bowl-shaped with sharp raised rims and have depth/diameter
ratios of about 1/5.
Martian craters change from simple to more complex types at diameters
of roughly 5 to 8 km. Complex craters have central peaks (or peak
complexes), relatively flat floors, and terracing or slumping along the
inner walls. Complex craters are shallower than simple craters in
proportion to their widths, with depth/diameter ratios ranging from 1/5
at the simple-to-complex transition diameter (~7 km) to about 1/30 for a
100-km diameter crater. Another transition occurs at crater diameters
of around 130 km as central peaks turn into concentric rings of hills to
form multi-ring basins.
Mars has the greatest diversity of impact crater types of any planet in the Solar System.
This is partly because the presence of both rocky and volatile-rich
layers in the subsurface produces a range of morphologies even among
craters within the same size classes. Mars also has an atmosphere that
plays a role in ejecta emplacement and subsequent erosion. Moreover,
Mars has a rate of volcanic and tectonic activity low enough that
ancient, eroded craters are still preserved, yet high enough to have
resurfaced large areas of the planet, producing a diverse range of
crater populations of widely differing ages. Over 42,000 impact craters
greater than 5 km in diameter have been catalogued on Mars,
and the number of smaller craters is probably innumerable. The density
of craters on Mars is highest in the southern hemisphere, south of the
dichotomy boundary. This is where most of the large craters and basins
are located.
Crater morphology provides information about the physical
structure and composition of the surface and subsurface at the time of
impact. For example, the size of central peaks in Martian craters is
larger than comparable craters on Mercury or the Moon.
In addition, the central peaks of many large craters on Mars have pit
craters at their summits. Central pit craters are rare on the Moon but
are very common on Mars and the icy satellites of the outer Solar
System. Large central peaks and the abundance of pit craters probably
indicate the presence of near-surface ice at the time of impact. Polewards of 30 degrees of latitude, the form of older impact craters is rounded out ("softened") by acceleration of soil creep by ground ice.
The most notable difference between Martian craters and other
craters in the Solar System is the presence of lobate (fludized) ejecta
blankets. Many craters at equatorial and mid-latitudes on Mars have this
form of ejecta morphology, which is thought to arise when the impacting
object melts ice in the subsurface. Liquid water in the ejected
material forms a muddy slurry that flows along the surface, producing
the characteristic lobe shapes. The crater Yuty is a good example of a rampart crater, which is so called because of the rampart-like edge to its ejecta blanket.
Martian craters are commonly classified by their ejecta. Craters with
one ejecta layer are called single-layer ejecta (SLE) craters. Craters
with two superposed ejecta blankets are called double-layer ejecta (DLE)
craters, and craters with more than two ejecta layers are called
multiple-layered ejecta (MLE) craters. These morphological differences
are thought to reflect compositional differences (i.e. interlayered ice,
rock, or water) in the subsurface at the time of impact.
Pedestal crater in Amazonis quadrangle as seen by HiRISE.
Martian craters show a large diversity of preservational states, from
extremely fresh to old and eroded. Degraded and infilled impact craters
record variations in volcanic, fluvial, and eolian activity over geologic time. Pedestal craters are craters
with their ejecta sitting above the surrounding terrain to form raised
platforms. They occur because the crater's ejecta forms a resistant
layer so that the area nearest the crater erodes more slowly than the
rest of the region. Some pedestals are hundreds of meters above the
surrounding area, meaning that hundreds of meters of material were
eroded away. Pedestal craters were first observed during the Mariner 9 mission in 1972.
Volcanism
First X-ray diffraction view of Martian soil - CheMin analysis reveals feldspar, pyroxenes, olivine and more (Curiosity rover at "Rocknest", October 17, 2012).
Volcanic structures and landforms cover large portions of the Martian
surface. The most conspicuous volcanoes on Mars are located in Tharsis and Elysium.
Geologists think one of the reasons volcanoes on Mars were able to grow
so large is that Mars has fewer tectonic boundaries in comparison to
Earth. Lava from a stationary hot spot was able to accumulate at one location on the surface for many hundreds of millions of years.
Scientists have never recorded an active volcano eruption on the surface of Mars. Searches for thermal signatures and surface changes within the last decade have not yielded evidence for active volcanism.
On October 17, 2012, the Curiosity rover on the planet Mars at "Rocknest" performed the first X-ray diffraction analysis of Martian soil. The results from the rover's CheMin analyzer revealed the presence of several minerals, including feldspar, pyroxenes and olivine, and suggested that the Martian soil in the sample was similar to the "weathered basaltic soils" of Hawaiian volcanoes. In July 2015, the same rover identified tridymite
in a rock sample from Gale Crater, leading scientists to conclude that
silicic volcanism might have played a much more prevalent role in the
planet's volcanic history than previously thought.
Sedimentology
Collection of spheres, each about 3 mm in diameter as seen by Opportunity rover
Flowing water appears to have been common on the surface of Mars at
various points in its history, and especially on ancient Mars. Many of these flows carved the surface, forming valley networks and producing sediment. This sediment has been redeposited in a wide variety of wet environments, including in alluvial fans, meandering channels, deltas, lakes, and perhaps even oceans.
The processes of deposition and transportation are associated with
gravity. Due to gravity, related differences in water fluxes and flow
speeds, inferred from grain size distributions, Martian landscapes were
created by different environmental conditions. Nevertheless, there are other ways of estimating the amount of water on ancient Mars. Groundwater has been implicated in the cementation of aeolian sediments and the formation and transport of a wide variety of sedimentary minerals including clays, sulphates and hematite.
When the surface has been dry, wind has been a major geomorphic agent. Wind driven sand bodies like megaripples and dunes are extremely common on the modern Martian surface, and Opportunity has documented abundant aeolian sandstones on its traverse. Ventifacts, like Jake Matijevic (rock), are another aeolian landform on the Martian Surface.
A wide variety of other sedimentological facies are also present locally on Mars, including glacial deposits, hot springs, dry mass movement deposits (especially landslides), and cryogenic and periglacial material, amongst many others. Evidence for ancient rivers, a lake, and dune fields have all been observed in the preserved strata by rovers at Meridiani Planum and Gale crater.
Common surface features
Groundwater on Mars
One
group of researchers proposed that some of the layers on Mars were
caused by groundwater rising to the surface in many places, especially
inside of craters. According to the theory, groundwater with dissolved
minerals came to the surface, in and later around craters, and helped to
form layers by adding minerals (especially sulfate) and cementing
sediments. This hypothesis is supported by a groundwater model and by
sulfates discovered in a wide area. At first, by examining surface materials with Opportunity Rover, scientists discovered that groundwater had repeatedly risen and deposited sulfates. Later studies with instruments on board the Mars Reconnaissance Orbiter showed that the same kinds of materials exist in a large area that included Arabia.
Interesting geomorphological features
Avalanches
On February 19, 2008, images obtained by the HiRISE camera on the Mars Reconnaissance Orbiter
showed a spectacular avalanche, in which debris thought to be
fine-grained ice, dust, and large blocks fell from a 700-metre
(2,300 ft) high cliff. Evidence of the avalanche included dust clouds
rising from the cliff afterwards. Such geological events are theorized to be the cause of geologic patterns known as slope streaks.
Image of the February 19, 2008 Mars avalanche captured by the Mars Reconnaissance Orbiter.
Closer shot of the avalanche.
Dust clouds rise above the 700-metre (2,300 ft) deep cliff.
A photo with scale demonstrates the size of the avalanche.
Possible caves
NASA scientists studying pictures from the Odyssey spacecraft have spotted what might be seven caves on the flanks of the Arsia Mons volcano on Mars.
The pit entrances measure from 100 to 252 metres (328 to 827 ft) wide
and they are thought to be at least 73 to 96 metres (240 to 315 ft)
deep. See image below: the pits have been informally named (A) Dena, (B)
Chloe, (C) Wendy, (D) Annie, (E) Abby (left) and Nikki, and (F) Jeanne.
Dena's floor was observed and found to be 130 m deep. Further investigation suggested that these were not necessarily lava tube "skylights". Review of the images has resulted in yet more discoveries of deep pits.
It has been suggested that human explorers on Mars could use lava
tubes as shelters. The caves may be the only natural structures
offering protection from the micrometeoroids, UV radiation, solar flares, and high energy particles that bombard the planet's surface. These features may enhance preservation of biosignatures over long periods of time and make caves an attractive astrobiology target in the search for evidence of life beyond Earth.
THEMIS image of cave entrances on Mars.
Inverted relief
Some
areas of Mars show inverted relief, where features that were once
depressions, like streams, are now above the surface. It is believed
that materials like large rocks were deposited in low-lying areas.
Later, wind erosion removed much of the surface layers, but left behind
the more resistant deposits. Other ways of making inverted relief might
be lava flowing down a stream bed or materials being cemented by
minerals dissolved in water. On Earth, materials cemented by silica are
highly resistant to all kinds of erosional forces. Examples of inverted
channels on Earth are found in the Cedar Mountain Formation near Green
River, Utah. Inverted relief in the shape of streams are further evidence of water flowing on the Martian surface in past times.
Inverted relief in the form of stream channels suggest that the climate
was different—much wetter—when the inverted channels were formed.
In an article published in January 2010, a large group of
scientists endorsed the idea of searching for life in Miyamoto Crater
because of inverted stream channels and minerals that indicated the past
presence of water.
