Geology of Mars
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INTRODUCTION
Mars is the fourth planet in our solar system next to our
own planet Earth at a mean distance of 227.94 million km or 1.5
AU. It is the second smallest planet after Mercury and the
outermost of the inner, terrestrial planets. Mars was named
after the Roman god of war because of its red color resembling
the blood. It is often described as the "Red Planet" because the
iron oxide prevalent on its surface gives it a reddish appearance.
Mars is a terrestrial planet with a thin atmosphere, having
surface features similar to both of the impact craters of the
Moon and the volcanoes, valleys, deserts, and polar ice caps of
Earth. It is home to the highest mountain of our solar system:
Olympus Mons and the largest known canyon: Valles Marineris.
The rotational period and seasonal cycles of Mars are similar to those of Earth, as is the tilt that produces the
seasons. Mars has two moons, Phobos and Deimos, which are small and irregularly shaped which may be
captured asteroids.
The mean radius of the planet is almost half of the Earth- some 3390 km. Its volume is also quite less-
about 15% of the Earth’s volume, which comes to be around 1631×107 km3. Mars has a mass of 6417×1019 kg.
It has a weak gravity- about 37% that of the Earth. It also has a very thin atmosphere, which is almost entirely
comprised of Carbon dioxide (about 96%) with small amounts of Nitrogen and Argon. The average temperature
on the planet’s surface is about -63⁰C. Mars can easily be seen from Earth with the naked eye with its reddish
coloring with its apparent magnitude reaches −3.
Until the first successful Mars flyby in 1965 by Mariner 4, many speculated about the presence of liquid
water on the planet's surface. This was based on observed periodic variations in light and dark patches,
particularly in the polar latitudes, which appeared to be seas and continents. Long, dark striations were
interpreted by some as irrigation channels for liquid water. These straight line features were later explained as
optical illusions, though geological evidence gathered by unmanned missions suggests that Mars once had large-
scale water coverage on its surface
at some earlier stage of its life. The
Mars rover Spirit sampled chemical
compounds containing water
molecules in March 2007. The
Phoenix lander directly sampled
water ice in shallow Martian soil on
July 31, 2008.
Mars as seen from Hubble Space Telescope
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Mars is currently host to five functioning spacecraft: three in orbit – the Mars Odyssey, Mars Express, and
Mars Reconnaissance Orbiter – and two on the surface – Mars Exploration Rover Opportunity and the Mars
Science Laboratory Curiosity. Defunct spacecraft on the surface include Mars Exploration rover-A Spirit and
several other inert landers and rovers such as the Phoenix lander, which completed its mission in 2008.
Observations by the Mars Reconnaissance Orbiter have revealed possible flowing water during the warmest
months on Mars. In 2013, NASA's Curiosity rover discovered that Mars' soil contains between 1.5% and 3% water
by mass. This report attempts in summarizing the geological information about Mars.
Comparison between Mars and the Earth
Earth Mars
Average Distance from Sun 150 million km 228 million km
Average Speed in Orbiting Sun 29.8 km per second 23.3 km per second
Average Diameter 12756 km 6791 km
Tilt of Axis 23.5 degrees 25 degrees
Length of Year 365.25 Days 687 Earth Days
Length of Day 23 hours 56 minutes 24 hours 37 minutes
Gravity 2.66 times that of Mars 0.375 that of Earth
Average Surface Temperature 13.8 degree C -63 degree C
Composition of Atmosphere nitrogen, oxygen, argon, others mostly carbon dioxide, some water
vapor
Number of Moons 1 2
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Exploration of Mars
The exploration of Mars has taken place over hundreds of years, beginning in earnest with the invention
and development of the telescope during the 1600s. Increasingly detailed views of the planet from Earth inspired
speculation about its environment and possible life – even intelligent civilizations – that might be found there.
Probes sent from Earth beginning in the late 20th century have yielded a dramatic increase in knowledge about
the Martian system, focused primarily on understanding its geology and habitability potential.
Engineering interplanetary journeys is very complicated, so the exploration of Mars has experienced a
high failure rate, especially in earlier attempts. Roughly two-thirds of all spacecraft destined for Mars failed
before completing their missions, and there are some that failed before their observations could begin. However,
missions have also met with unexpected levels of success, such as the twin Mars Exploration Rovers operating for
years beyond their original mission specifications.
Since 6 August 2012, there have been two scientific rovers on the surface of Mars beaming signals back
to Earth (Opportunity of the Mars Exploration Rover mission, and Curiosity of the Mars Science Laboratory
mission), and three orbiters currently surveying the planet: Mars Odyssey, Mars Express, and Mars
Reconnaissance Orbiter. Two orbiters launched in November 2013, Mars Orbiter Mission of ISRO and MAVEN of
NASA are currently on their way to Mars.
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On 24 January 2014, NASA reported that current studies on the planet Mars by the Curiosity and
Opportunity rovers will now be searching for evidence of ancient life, including a biosphere based on autotrophic,
chemotropic and/or chemo-litho-autotrophic microorganisms, as well as ancient water, including fluvio-
lacustrine environments (plains related to ancient rivers or lakes) that may have been habitable. The search for
evidence of habitability, taphonomy (related to fossils), and organic carbon on the planet Mars is now a primary
NASA objective.
Till date 43 missions have been sent to Mars, of which only 15 have been successful. The following
missions proved to be successful in one or the other sense and provided useful data about the geology and
composition of Mars.
1. Mariner Program:
In 1964 and 1968, NASA sent four probes to Mars- Mariner 3-4, 6-7 and Mariner 8-9 respectively. Mars was
visited by Mariner 4 in 1965 and was photographed by it, becoming the first planet to be photographed. In 1969,
Mariner 9 became the first man-made object to orbit another planet. It photographed the surface of Mars. These
pictures were the first to offer more detailed evidence that liquid water might at one time have flowed on the
planetary surface. They also finally discerned the true nature of many Martian albedo features. For example, Nix
Olympica was one of only a few features that could be seen during the planetary dust-storm, revealing it to be
the highest mountain (volcano, to be exact) on any planet in the entire Solar System, and leading to its
reclassification as Olympus Mons.
2. Viking Program:
The Viking program launched Viking 1 and 2 spacecraft to Mars in 1975; The program consisted of two orbiters
and two landers – these were the first two spacecraft to successfully land and operate on Mars. The Viking
orbiters revealed that large floods of water carved deep valleys, eroded grooves into bedrock, and traveled
thousands of kilometers. Areas of branched streams, in the southern hemisphere, suggest that rain once fell.
3. Mars Pathfinder:
It was a NASA spacecraft that landed a base station with a roving probe on Mars on July 4, 1997. It consisted of a
lander and a small 10.6 kilograms wheeled robotic rover named Sojourner, which was the first rover to operate
on the surface of Mars. Sojourner studied some of the big chunks of rocks spread near its landing site with
respect to their chemistry and mineralogy.
4. Mars Global Surveyor (MGS):
It was the first fully successful mission overall, to the red planet in two decades when it was launched on
November 7, 1996, and entered orbit on September 12, 1997. The spacecraft began its primary mapping mission
in March 1999. It observed the planet from a low-altitude, nearly polar orbit. The mission studied the entire
Martian surface, atmosphere, and interior, and returned more data about the red planet than all previous Mars
missions combined. Among key scientific findings, Global Surveyor took pictures of gullies and debris flow
features that suggest there may be current sources of liquid water, similar to an aquifer, at or near the surface of
the planet. Similar channels on Earth are formed by flowing water, but on Mars the temperature is normally too
cold and the atmosphere too thin to sustain liquid water. Nevertheless, many scientists hypothesize that liquid
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groundwater can sometimes surface on Mars, erode gullies and channels, and pool at the bottom before freezing
and evaporating. Magnetometer readings showed that the planet's magnetic field is not globally generated in the
planet's core, but is localized in particular areas of the crust. Data from the spacecraft's laser altimeter gave
scientists their first 3-D views of Mars' north polar ice cap.
5. Mars Odyssey:
In 2001 NASA's Mars Odyssey orbiter arrived at Mars. Its mission is to use spectrometers and imagers to hunt for
evidence of past or present water and volcanic activity on Mars. In 2002, it was announced that the probe's
gamma ray spectrometer and neutron spectrometer had detected large amounts of hydrogen, indicating that
there are vast deposits of water ice in the upper three meters of Mars' soil within 60° latitude of the south pole.
6. Mars Exploration Rovers (MER)- Spirit and Opportunity:
NASA's Mars Exploration Rover Mission is an ongoing robotic space mission involving two rovers, Spirit and
Opportunity, exploring the planet Mars. The mission's scientific objective was to search for and characterize a
wide range of rocks and soils that hold clues to past water activity on Mars. In particular, samples sought include
those that have minerals deposited by water-related processes such as precipitation, evaporation, sedimentary
cementation, or hydrothermal activity; to determine the distribution and composition of minerals, rocks, and
soils surrounding the landing sites; to determine what geologic processes have shaped the local terrain and
influenced the chemistry. Such processes could include water or wind erosion, sedimentation, hydrothermal
mechanisms, volcanism, and cratering. Search for iron-containing minerals, and to identify and quantify relative
amounts of specific mineral types that contain water or were formed in water, such as iron-bearing carbonates.
Characterize the mineralogy and textures of rocks and soils to determine the processes that created them.
7. Mars Reconnaissance Orbiter (MRO):
Mars Reconnaissance Orbiter is a multipurpose spacecraft designed to conduct reconnaissance and exploration
of Mars from orbit which was launched on August 12, 2005, and attained Martian orbit on March 10, 2006. It is
currently imaging the Martian surface at very high resolution.
8. Mars Science Laboratory (MSL)- Curiosity Rover:
The NASA Mars Science Laboratory mission with its rover named Curiosity, was launched on November 26, 2011.
The rover carries instruments designed to look for past or present conditions relevant to the past or present
habitability of Mars. The Curiosity rover landed on Mars on Aeolis Palus in Gale Crater. The geological goals of
this mission are to investigate the chemical, isotopic, and mineralogical composition of the Martian surface and
near-surface geological materials and interpret the processes that have formed and modified rocks and soils.
9. Mars Orbiter Mission (MOM) or Mangalyaan:
It is a Mars orbiter launched in November 2013 by ISRO and it’s India’s first interplanetary mission. Though the
primary objective of the Mars Orbiter Mission is to develop the technologies required for design, planning,
management and operations of an interplanetary mission, the spacecraft will be exploring Mars' surface features,
morphology, mineralogy and Martian atmosphere using indigenous scientific instruments.
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Gravity and Magnetism
The gravity of Mars is nearly 37% of the Earth’s gravity. The gravitational acceleration on the surface of
Mars is 3.71 m/s2 as compared to 9.8 m/s2 of Earth. As we know, gravitational pull of a planet is proportional to
its mass. Mass of Mars is almost 11% that of the Earth, implying the lower gravitational pull than that of the
Earth’s. Gravity of Mars has been measured by many previous and ongoing missions to Mars and this value is
now recognized as standard.
Currently Mars shows almost no magnetism. As there’s no magnetic dipole similar to that of the Earth, it
is suspected that Mars has a cooled, solid core, unlike Earth’s liquid core, which is thought to be responsible for
the generation of the magnetic field. The first indication of the weak magnetic field of Mars was obtained during
the Mariner 4 spacecraft flyby in 1965. But recent studies of magnetism in surface rocks of Martian surface by the
Mars Global Surveyor spacecraft suggest that the Red Planet was magnetized more widely and strongly in its
geologic past. Scientists think Mars had the ability to generate a strong magnetic field in its core during its first
half-billion to 1 billion years. The Martian field flipped polarity (swapping magnetic north and south) just as
Earth's magnetic field has done repeatedly. But perhaps because the Martian core cooled, its magnetic dynamo
shut down within a billion years of the planet's birth.
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Geography of Mars
Many times called as Areography (Ares-Greek for Mars), the geography of Mars includes mapping and
naming the surface of Mars. Martian geography is mainly focused on what is called physical geography on Earth;
that is the distribution of physical features across Mars and their cartographic representations.
The first observations of Mars were from ground-based telescopes. In September 1877, Italian
astronomer Giovanni Schiaparelli published the first detailed map of Mars. These maps notably contained
features he called canali ("channels"), that were later shown to be an optical illusion. Following these
observations, it was a long held belief that Mars contained vast seas and vegetation. It was not until spacecraft
visited the planet during NASA's Mariner missions in the 1960s, that these myths were dispelled. Some maps of
Mars were made using the data from these missions, but it wasn't until the Mars Global Surveyor mission,
launched in 1996 and ending in late 2006, that complete, extremely detailed maps were obtained. These maps
are now available online. Currently we have more detailed maps of Mars than our own planet’s ocean floor.
Albedo Features
The classical albedo features of Mars are the light and dark features that can be seen on the planet Mars
through an Earth-based telescope. Before the age of space probes, several astronomers created maps of Mars on
which they gave names to the features they could see. Today, the improved understanding of Mars enabled by
space probes has rendered many of the classical names obsolete for the purposes of cartography; however, some
of the old names are still used to describe geographical features on the planet. These albedo contrasts rarely
correspond to topographic features and in many cases obscure them. The lighter patches at the poles were
correctly believed to be a frozen substance, either water or carbon dioxide, but the nature of the dark patches
seen against the general reddish tint of Mars was uncertain for centuries. When Giovanni Schiaparelli began
observing Mars in 1877, he believed that the darker features were seas and lakes and named them in Latin
accordingly (mare for sea and lacus for lake). They are now known to be areas where the wind has swept away
High resolution colorized map of Mars based on Viking orbiter images. Surface frost and water ice fog brighten
the impact basin Hellas to the right of lower center; Syrtis Major just above it is darkened by winds that sweep
dust off its basaltic surface. Summer view of North and south polar ice caps are shown at upper and lower right.
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the surface dust, leaving a darker, rockier surface; their borders change in response to windstorms on the
Martian surface that pick up the dust, widening or narrowing the features. In 1958, the International
Astronomical Union created a list of officially recognized Martian albedo features. Many of the names used for
topographic features on Mars are still based on the classical nomenclature of the feature's location; for instance,
the albedo feature Ascraeus Lacus provides the basis of the name of the volcano Ascraeus Mons. Various albedo
features can be seen in the map of Mars given on the previous page.
Zero Elevation and Zero Meridian
Since Mars has no oceans and hence no sea level, it is convenient to define an arbitrary zero-elevation
level or datum for mapping the surface. The datum for Mars is arbitrarily defined in terms of a constant
atmospheric pressure. From the Mariner 9 mission up until 2001, this was chosen as the point where there exists
the triple point for water which is 6.105 mbar. In 2001, Mars Orbiter Laser Altimeter data led to a new
convention of zero elevation defined as the equipotential surface (gravitational plus rotational) whose average
value at the equator is equal to the mean radius of the planet.
Mars' equator is defined by its rotation, but the location of its Prime Meridian was specified, as was
Earth's, by choice of an arbitrary point which was accepted by later observers. The German astronomers Wilhelm
Beer and Johann Heinrich Mädler selected a small circular feature as a reference point when they produced the
first systematic chart of Mars features in 1830-32. In 1877, their choice was adopted as the prime meridian by the
Italian astronomer Giovanni Schiaparelli when he began work on his notable maps of Mars. After the spacecraft
Mariner 9 provided extensive imagery of Mars in 1972, a small crater (later called Airy-0), located in the Sinus
Meridiani region (Middle Bay or Meridian Bay) along the line of Beer and Mädler, was chosen and accepted
worldwide.
Map of Quadrangles
The following imagemap of the planet Mars is divided into the 30 quadrangles defined by the United
States Geological Survey. The quadrangles are numbered with the prefix MC for Mars Chart. North is at the top; 0
meridian is at the far left. The map images were taken by the Mars Global Surveyor.
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Interior of Mars
The interior of Mars is poorly known. Planetary scientists have yet to conduct a successful seismic
experiment via spacecraft that would provide direct information on internal structure and so must rely on
indirect inferences. Though it is widely accepted that Mars is well differentiated similar to other terrestrial
planets. Thus it has an interior differentiated into an outer crust, middle mantle and inner core. However,
missions to explore Mars have only been successful in understanding its crust. The interior of the Red planet still
remains unknown, but many theories about its origin, structure and composition have been postulated by
planetary and geoscientists.
The moment of inertia of Mars indicates that it has a central core with a radius of about 900–2000 km.
Isotopic data from meteorites determined to have come from Mars demonstrate unequivocally that the planet
differentiated—separated into a metal-rich core and rocky mantle—at the end of the planetary accretion period
4.5 billion years ago. The planet has no detectable magnetic field that would indicate convection in the core
today. Large regions of magnetized rock have been detected in the oldest terrains, however, which suggests that
very early Mars did have a magnetic field but that it disappeared as the planet cooled and the core solidified.
Martian meteorites also suggest that the core may be more sulfur-rich than Earth’s core and the mantle more
iron-rich. The Martian core is probably made of a mixture of iron, sulfur and maybe oxygen.
Like Earth, the mantle of Mars is probably made of silicates; however, it's much smaller, at 1,300 to 1,800
kilometers thick. The mantle is thought to be more iron rich than the Earth’s mantle. There must have been
convection currents active in the mantle at one time. These currents would account for the formation of the
crustal upwarps or bulges, such as the Tharsis region, the Martian volcanoes and the fractures that formed Valles
Marineris.
On Mars, the crust is also thin, but isn't broken into plates like the Earth's crust. Although we do not
know of currently active volcanoes or Mars-quakes, evidence of quakes occurring as recently as a few million
years ago suggest they are possible. The average thickness of Martian crust is thought to be around 100 km,
which is pretty much thicker than that of the Earth.
The future space probes and landers planned to Mars will be carrying high-tech seismometers and other
geophysical instruments which will help improve our understanding of the Martian interior.
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Composition of Mars
Mars is a terrestrial planet, which means that its bulk composition, like Earth's, consists of silicates,
metals and other elements that typically make up a rock. Also like Earth, Mars is a differentiated planet, meaning
that it has a central core made up of metallic iron and nickel surrounded by a less dense, silicate mantle and
crust. The planet's distinctive red color is due to the oxidation of iron on its surface. Much of what we know
about the elemental composition of Mars comes from orbiting spacecraft and landers. Most of these spacecraft
carry spectrometers and other instruments to measure the surface composition of Mars by either remote sensing
from orbit or in situ analyses on the surface. We also have many actual samples of Mars in the form of meteorites
that have made their way to Earth.
Elemental Composition
Based on various data sources, scientists think that the most abundant chemical elements in the Martian
crust, besides silicon and oxygen, are iron, magnesium, aluminum, calcium, and potassium. These elements are
major components of the minerals comprising igneous rocks. The elements titanium, chromium, manganese,
sulfur, phosphorus, sodium, and chlorine are less abundant but are still important components of many accessory
minerals in rocks and of secondary minerals in the dust and soils (or regolith). Hydrogen is present as water (H2O)
ice and in hydrated minerals. Carbon occurs as carbon dioxide (CO2) in the atmosphere and as dry ice at the
poles. An unknown amount of carbon is also stored in carbonates. Molecular nitrogen (N2) makes up 2.7 percent
of the atmosphere. As far as we know, organic compounds are absent except for a trace of methane detected in
the atmosphere. The exact percentage of elemental composition of either Martian surface (crust) or the interior
is unavailable and studies are ongoing on the chemistry by the rovers.
The elemental composition of Mars is different from Earth’s in several significant ways. First, Martian
meteorite analysis suggests that the planet's mantle is about twice as rich in iron as the Earth's mantle. Second,
its core is more rich in sulfur. Third, the Martian mantle is richer in potassium and phosphorus than Earth's, and
fourth, the Martian crust contains a higher percentage of volatile elements such as sulfur and chlorine than the
Earth's crust does. Many of these conclusions are supported by in situ analyses of rocks and soils on the Martian
surface.
Primary Rocks and Minerals
Mars is fundamentally an igneous planet. Rocks on the surface and in the crust consist predominantly of
pyrogenetic minerals. Most of our current knowledge about the mineral composition of Mars comes from
spectroscopic data from orbiting spacecraft, in situ analyses of rocks and soils from six landing sites, and study of
the Martian meteorites. Spectrometers currently in orbit include THEMIS (Mars Odyssey), OMEGA (Mars
Express), and CRISM (Mars Reconnaissance Orbiter). The two Mars exploration rovers each carry an Alpha
Particle X-ray Spectrometer (APXS), a thermal emission spectrometer (Mini-TES), and Mössbauer spectrometer to
identify minerals on the surface.
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On October 17, 2012, the Curiosity rover on the planet Mars
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 feldspars, pyroxenes and olivine, and suggested that the Martian
soil in the sample was similar to the weathered basaltic soils of Hawaiian
volcanoes.
The dark areas of Mars are characterized by the mafic rock-forming
minerals olivine, pyroxene, and plagioclase feldspar. Olivine occurs all over
the planet, but some of the largest concentrations are in Nili Fossae, an area
containing Noachian-aged rocks (equivalent of Earth’s Haedean-Archaen
Eons-about 3 to 4 billion years old). Another large olivine-rich outcrop is in
Ganges Chasma, an eastern side chasm of Valles Marineris. Olivine is unstable
at surface pressure-temperature conditions, hence it
weathers rapidly into chloritic and clay minerals in the
presence of liquid water. Therefore, areas with large
outcroppings of olivine-bearing rock indicate that liquid
water has not been abundant since the rocks formed.
Pyroxene minerals are also widespread across the
surface. Both low-calcium i.e. ortho and high-calcium i.e.
clino-pyroxenes are present, with the high-calcium varieties
associated with younger volcanic shields and the low-
calcium forms (enstatite-ferrosilite) more common in the old
highland terrain. Because enstatite melts at a higher
temperature than its high-calcium cousin diopside, some
researchers have argued that its presence in the
highlands indicates that older magmas on Mars had
higher temperatures than younger ones.
Between 1997 and 2006, the Thermal Emission
Spectrometer (TES) on the Mars Global Surveyor (MGS)
spacecraft mapped the global mineral composition of the
planet. TES identified two global-scale volcanic units on
Mars. Type 1 characterizes the Noachian-aged highlands
and consists of unaltered plagioclase and clino-pyroxene-
rich basalts. Type 2 is common in the younger plains
north of the dichotomy boundary and is more silica rich
than the other type. The lavas of Type 2 have been
interpreted as andesites or basaltic andesites, indicating
the lavas in the northern plains originated from more
chemically evolved, volatile-rich magmas. However, other
researchers have suggested that Type 2 represents
First X-ray diffraction view of Martian
rock. Analysis revealed feldspar,
pyroxenes & olivine.
Mars Odyssey THEMIS false-color image of olivine basalts in
the Valles Marineris. Layers rich in olivine appear purple
Sojourner rover analyzing the rock Yogi, photographed by
camera on the Pathfinder lander
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weathered basalts with thin coatings of silica glass or other secondary minerals that formed through interaction
with water- or ice-bearing materials.
True intermediate and felsic rocks are present on Mars, but exposures are uncommon. Both TES and the
Thermal Emission Imaging System (THEMIS) on the Mars Odyssey spacecraft have identified high silica rocks in
Syrtis Major and near the southwestern rim of the Antoniadi crater. The rocks have spectra resembling quartz-
rich dacites and granitic rocks, suggesting that at least some parts of the Martian crust may have a diversity of
igneous rocks similar to Earth's. Some geophysical evidence suggests that the bulk of the Martian crust may
actually consist of basaltic andesites or andesites. The andesitic crust is hidden by overlying basaltic lavas that
dominate the surface composition, but is volumetrically minor.
Rocks studied by Spirit Rover in Gusev crater can be classified in different ways. The amounts and types
of minerals make the rocks primitive basalts—also called picritic basalts. The rocks are similar to ancient
terrestrial rocks called basaltic komatiites. Rocks of the plains also resemble the basaltic shergottites, meteorites
which came from Mars.
In the journal Science from September 2013, researchers described a different type of rock called Jake M,
as it was the first rock analyzed by the Alpha Particle X-ray Spectrometer instrument on the Curiosity rover, and it
was different from other known Martian igneous rocks as it is alkaline (>15% nepheline) and relatively
fractionated. This rock is similar to oligoclase-bearing basalts which are typically found at ocean islands and
continental rifts. Its discovery may mean that alkaline magmas may be more common on Mars than on Earth and
that Curiosity could encounter even more fractionated alkaline rocks (for example, phonolites and trachytes).
Using Curiosity rover’s Sample Analysis at Mars (SAM) mass spectrometer, scientists measured isotopes
of helium, neon, and argon that cosmic rays produce as they go through rock. The fewer of these isotopes they
find, the more recently the rock has been exposed near the surface. The 4-billion-year-old lakebed rock drilled by
Curiosity was uncovered between 30 million and 110 million years ago by winds which sandblasted away 2
meters of overlying rock.
Secondary Minerals
Minerals produced through hydrothermal alteration and weathering of primary basaltic minerals are also
present on Mars. Secondary minerals include hematite, phyllosilicates like clay minerals, goethite, jarosite (a
hydrous sulphate of K and Fe3+), iron sulfate minerals, opaline silica and gypsum. Many of these secondary
minerals require liquid water to form.
Opaline silica and iron sulphate minerals form in acidic (low pH) water. Sulphates have been found in a
variety of locations, including near Juventae Chasma, Ius Chasma, Melas Chasma, Candor Chasma and Ganges
Chasma. These sites all contain fluvial landforms indicating that abundant water was once present. Spirit rover
has discovered sulfates and goethite in the Columbia Hills. On March 18, 2013, NASA reported evidence from
instruments on the Curiosity rover of hydrated minerals, likely hydrated calcium sulfate i.e. Gypsum, in several
rock samples as well as in veins and nodules in other rocks. Analysis using the rover's instruments provided
evidence of subsurface water, amounting to as much as 4% water content, down to a depth of 60 cm.
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Some of the mineral classes detected may have formed in
environments suitable (i.e. enough water and the proper pH) for life. The
mineral smectite (a type of clay mineral) forms in near-neutral waters.
Phyllosilicates and carbonates are good for preserving organic matter, so
they may contain evidence of past life. Sulfate deposits preserve fossils
and fossils of microorganisms form in iron oxides like hematite. The
presence of opaline silica points toward a hydrothermal environment
that could support life. Silica is also excellent for preserving evidence of
microbes.
The most conspicuous of all secondary minerals found on Mars are the hematite spherules (informally
known as blueberries). These are abundant spherical hematite inclusions discovered by the Mars rover
Opportunity at Meridiani Planum. They are found in situ embedded in a sulfate salt evaporite matrix, and also
loose on the surface. The shapes by themselves don't reveal the particles' origin with certainty. Not only are
there spherules on the surface but they are also found deeper in the Martian soil.
Sedimentary Rocks
Layered sedimentary deposits are widespread on
Mars. These deposits probably consist of both- lithified
sediments and semi- or unconsolidated sediments. Thick
sedimentary deposits occur in the interior of several
canyons in Valles Marineris, within large craters in Arabia
and Meridiani Planum and probably comprise much of
the deposits in the northern lowlands. The Mars
Exploration Rover Opportunity landed in an area
containing cross-bedded (mainly eolian) sandstones (see
picture on next page). Fluvial-deltaic deposits are
present in Eberswalde Crater and elsewhere, and photo-
geologic evidence suggests that many craters and low
lying inter-crater areas in the southern highlands contain
Noachian-aged lake sediments.
While the possibility of carbonates on Mars has
been of great interest to exobiologists and geochemists
alike, there was little evidence for significant quantities of
carbonate deposits on the surface. In the summer of 2008,
the Phoenix Mars lander found between 3–5 % by weight
calcite (CaCO3) and an alkaline soil. In 2010, analyses by the
Mars Exploration Rover Spirit identified outcrops rich in
magnesium-iron carbonate (16–34 wt%) in the Columbia
Hills of Gusev crater. The magnesium-iron carbonate
(magnesite-siderite) most likely precipitated from
carbonate-bearing solutions under hydrothermal
Cross-bedded sandstones inside Victoria Crater
Conglomerate as seen by Curiosity rover
Hematite spherules
Cross bedded sandstone in Victoria crater
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conditions at near-neutral pH in association with volcanic activity during the Noachian Period.
Carbonates (calcium or iron carbonates) were discovered in a crater on the rim of Huygens Crater,
located in the Iapygia quadrangle. The impact on the rim exposed material that had been dug up from the impact
that created Huygens crater. These minerals represent evidence that Mars once was had a thicker carbon dioxide
atmosphere with abundant moisture. These kind of carbonates only get deposited in marine environments. They
were found with the Compact Reconnaissance Imaging Spectrometer for Mars (CRISM) instrument on the Mars
Reconnaissance Orbiter. Earlier, the instrument had detected clay minerals. The carbonates were found near the
clay minerals. Both of these minerals form in wet environments. It is supposed that billions of years ago Mars was
much warmer and wetter. At that time, carbonates would have formed from water and the carbon dioxide-rich
atmosphere. Later the deposits of carbonate would have been buried. The double impact has now exposed the
minerals. Earth has vast carbonate deposits in the form of limestone.
Dust and Soil
Much of the Martian surface is
deeply covered by dust as fine as talcum
powder i.e. clay sized. The global
predominance of dust obscures the
underlying bedrock, making
spectroscopic identification of primary
minerals impossible from orbit over many
areas of the planet. The red/orange
appearance of the dust is caused by ferric
oxide and the ferric hydroxide mineral-
goethite. The Mars Exploration Rovers
identified magnetite as the mineral
responsible for making the dust
magnetic. It probably also contains some
titanium.
The global dust cover and the presence of other wind-blown sediments has made soil compositions
remarkably uniform across the Martian surface. Analysis of soil samples from the Viking landers in 1976,
Pathfinder, and the Mars Exploration rovers show nearly identical mineral compositions from widely separated
locations around the planet. The soils also consist of finely broken up basaltic rock fragments.
Meteorites found on Mars
The rovers operating on the Martian surface came across many meteorites. Opportunity rover
encountered the most meteorites. The rover found meteorites just sitting on plains near its landing site. The first
one analyzed with Opportunity's instruments was called Heat-shield Rock, as it was found near where
Opportunity's heat shield landed. Examination with the Miniature Thermal Emission Spectrometer (Mini-TES) and
Mossbauer spectrometer lead to its classification as an IAB meteorite. It was determined that it was composed of
93% iron and 7% nickel. Some other meteorites examined were stony, stony-iron and iron meteorites. It is
Composition of Martian soil as analyzed by the rovers
15
estimated that Mars will have a lot more meteorites preserved than the Earth because of the very less active
geological processes.
Martian Meteorites
Martian meteorites refer to the meteorites found on the Earth which were ejected from Mars as asteroid
or comets made impact on its surface and eventually landed on Earth. Of over 61,000 meteorites that have been
found on Earth, 132 were identified as Martian as of March 2014. Martian meteorites should not be confused
with meteorites found on Mars. These meteorites are thought to be from Mars because they have elemental and
isotopic compositions that are similar to rocks and atmosphere gases analyzed by spacecraft on Mars. Martian
meteorites are divided into three rare groups of achondritic (stony) meteorites: shergottites, nakhlites and
chassignites. Consequently, Martian meteorites as a whole are sometimes referred to as the SNC group. They
have isotope ratios that are said to be consistent with each other and inconsistent with the Earth.
(a) Sherghottites: Roughly three-quarters of all Martian meteorites
can be classified as shergottites. They are igneous rocks of mafic to
ultramafic lithology. They fall into three main groups- the basaltic,
olivine-phyric and lherzolitic shergottites, based on their crystal
size and mineral content. The shergottites appear to have
crystallised as recently as 180 million years ago, which is a
surprisingly young age considering how ancient the majority of the
surface of Mars appears to be, and the small size of Mars itself.
Because of this, some have advocated the idea that the
shergottites are much older than this. This "Shergottite Age
Paradox" remains unsolved and is still an area of active research
and debate.
(b) Nakhlites: They are igneous rocks that are rich in augite and were formed from basaltic magma about 1.3
billion years ago. They contain augite and olivine crystals. Their crystallization ages, compared to a crater
count chronology of different regions on Mars, suggest the nakhlites formed on the large volcanic
construct of either Tharsis, Elysium, or Syrtis Major Planum. It has been shown that the nakhlites were
suffused with liquid water around 620 million years ago and that they were ejected from Mars around
10.75 million years ago by an asteroid impact. They fell to Earth within the last 10,000 years.
(c) Chassignites: These are the rarest of all Martian meteorites, only two known specimens have been yet
found. Their composition is almost entirely olivine i.e. of the monomineralic rock dunite with small traces
of feldspars and some oxides.
A Sherghottite meteorite
16
Global Physiography
Most of our current knowledge about the geology of Mars comes from studying landforms and relief
features seen in images taken by orbiting spacecrafts. 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
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 similar to the elevation difference between Earth’s continents and ocean basins. 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
other distinction between the two hemispheres is in their 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), while
crust in the northern lowlands peaks at around 32 km (20 mi) in thickness. The location of the dichotomy
boundary varies in latitude across.
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 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
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
hemispheric dichotomy is clearly visible here.
17
north, similar to what is occurring at spreading plate boundaries on Earth. Whatever its origin, the Martian
dichotomy appears to be extremely old. 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 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. The topographic map of
Mars given on the previous page shows the hemispheric dichotomy clearly.
Crustal Bulges and Volcanic Provinces
Straddling the dichotomy boundary in Mars’
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. 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 buldge. The vast
Alba Mons (formerly Alba Patera) occupies the northern
part of the region. The huge shield volcano Olympus
Mons lies off the main buldge, 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
The Tharsis bulge, showing the Tharsis Montes (right) along their
NE-SW axis and the giant Olympus Mons (upper left corner)
The Hellas basin
18
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, which probably represent uplifted, jostled blocks of
old pre-basin crust. Ancient, low-relief volcanic constructs 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
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. 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.
Equatorial Canyon Systems
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, while 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.
Chaos Terrain
Chaos terrain on Mars is distinctive; nothing on
Earth compares to it. Chaos terrain generally consists of
irregular groups of large blocks, some tens of km across
and a hundred or more meters high. The tilted and flat
topped blocks form depressions hundreds of meters
deep. A chaotic region can be recognized by mesas,
buttes and hills, chopped through with valleys which in
places look almost patterned. Chaos terrain is
presumably formed by sudden melting of subterranean
ice in the form of huge discharge of water. Some parts of
this chaotic area have not collapsed completely—they
are still formed into large mesas, so they may still contain
water ice.
The Hydraotes Chaos terrain (420 km across)
in the Oxia Palus quadrangle
19
Polar Ice Caps
The polar ice caps are well-known telescopic features of
Mars, first identified by Christian 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 to be
synonymous with the "polar ice caps", but the permanent ice (seen as the high albedo, white surfaces in images)
forms only a relatively thin cover 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.
Planum Boreum or the Northern Polar Ice Cap on Mars
20
Common Landforms on Mars
The following landforms are very commonly seen on the Martian surface. Some of them are tectonic in
origin, some sedimentary while some are volcanic in origin. Major landforms observed are volcanoes, many types
of impact craters, valleys and canyons, gulies, fans, lava flows, mesas and buttes, chaos terrain, dunes, etc.
Volcanoes
Mars is only about one-half the size of Earth and yet has
several volcanoes that surpass the scale of the largest terrestrial
volcanoes. The most massive volcanoes are located on huge
uplifts or domes in the Tharsis and Elysium regions of Mars.
Located on the northwest flank of Tharsis bulge are three large
shield volcanoes: Ascraeus Mons, Pavonis Mons and Arsia Mons.
Beyond the dome's northwest edge is Olympus Mons, the largest
of the Tharsis volcanoes. Olympus Mons is classified as a shield
volcano. It is 24 km high, 550 km in diameter and is rimmed by a 6
km high scarp. It is one of the largest volcanoes in the Solar
System. By comparison the largest volcano on Earth is Mauna Loa
which is 9 km high and 120 km across. Elysium Planitia has smaller
volcanoes than the Tharsis region, but a more diverse volcanic
history. The three volcanoes include Hecates Tholus, Elysium
Mons and Albor Tholus.
The large shield volcanoes on Mars resemble Hawaiian shield volcanoes. They both have effusive
eruptions which are relatively quiet and basaltic in nature. Both have summit pits or calderas and long lava flows
or channels. The biggest difference between Martian and Terrestrial volcanoes is size. The volcanoes in the
Tharsis region are 10 to 100 times larger than those on Earth. They were built from large magma chambers deep
within the Martian crust. The Martian flows are also much longer. This is probably due to larger eruption rates
and to lower gravity. One of the reasons volcanoes of such magnitude were able to form on Mars is because the
hot volcanic regions in the mantle remained fixed relative to the surface for hundreds of millions of years. On
Earth, the tectonic flow of the crust across the hot volcanic regions prevent large volcanoes from forming. These
volcanoes have a relatively short life time. As the plate moves new volcanoes form and the old ones become
silent.
Not all Martian volcanoes are classified as shields with effusive eruption styles. North of the Tharsis
region lies Alba Patera. This volcano is comparable to Olympus Mons in its horizontal extent but not in height. Its
base diameter is 1,500 km but is less than 7 km high. Ceraunius Tholus is one of the smaller volcanoes. It is about
the size of the Big Island of Hawaii. It exhibits explosive eruption characteristics and probably consists of ash
deposits. Tyrrhena Patera and Hadriaca Patera both have deeply eroded features which indicate explosive ash
eruptions.
Olympus Mons volcano, showing its peripheral scarp
and the summit shows many over-lapping calderas
21
Lava Flows and Volcanic Plains
Volcanic plains are widespread on Mars. Two types of plains are commonly recognized: those where lava
flow features are common, and those where flow features are generally absent but a volcanic origin is inferred by
other characteristics. Plains with abundant lava flow features occur in and around the large volcanic provinces of
Tharsis and Elysium. Flow features include both sheet flow and tube- and channel-fed flow morphologies. Sheet
flows show complex, overlapping flow lobes and may extend for many hundreds of kilometers from their source
areas. Lava flows can form a lava tube when the exposed upper layers of lava cool and solidify to form a roof
while the lava underneath continues flowing. Often, when all the remaining lava leaves the tube, the roof
collapses to make a channel or line of pit craters.
An unusual type of flow feature
occurs in the Cerberus plains south of
Elysium and in Amazonis. These flows have
a broken platey texture, consisting of dark,
kilometer-scale slabs embedded in a light-
toned matrix. They have been attributed
to rafted slabs of solidified lava floating on
a still-molten subsurface. Others have
claimed the broken slabs represent pack
ice that froze over a sea that pooled in the
area after massive releases of
groundwater from the Cerberus Fossae
area.
The second type of volcanic plains
(ridged plains) are characterized by
abundant wrinkle ridges. Volcanic flow
features are rare or absent. The ridged
plains are believed to be regions of
extensive flood basalts, by analogy with
the lunar maria. Ridged plains make up
about 30% of the Martian surface and are
most prominent in Lunae, Hesperia, and
Malea Plana, as well as throughout much
of the northern lowlands. Ridged plains
are all Hesperian in age and represent a
style of volcanism globally predominant
during that time period. The Hesperian
Period is named after the ridged plains in
Hesperia Planum.
Lava flows seen on the western flank of Olympus Mons
Two overlapping lava flows towards west of Olympus Mons showing
wrinkled ridges; the flow on the right is younger
22
Gullies
Gullies are small, incised networks of narrow channels and their associated downslope sediment
deposits, found on Mars. They are named for their resemblance to terrestrial gullies. First discovered on images
from Mars Global Surveyor, they occur on steep slopes, especially on the walls of craters. Usually, each gully has a
dendritic alcove at its head, a fan-shaped apron at its base, and a single thread of incised channel linking the two,
giving the whole gully an hourglass shape. They are believed to be relatively young because they have few, if any
craters. A subclass of gullies is also found cut into the faces of sand dunes, that are themselves considered to be
quite young.
Most gullies occur 30 degrees poleward in each hemisphere, with greater numbers in the southern
hemisphere. Some studies have found that gullies occur on slopes that face all directions; others have found that
the greater number of gullies are found on poleward facing slopes, especially from 30-44 S. Although thousands
have been found, they appear to be restricted to only certain areas of the planet. In the northern hemisphere,
they have been found in Arcadia Planitia, Tempe Terra, Acidalia Planitia, and Utopia Planitia. In the south, high
concentrations are found on the northern edge of Argyre basin, in northern Noachis Terra, and along the walls of
the Hellas outflow channels.
On the basis of their form, aspects, positions, and location amongst and apparent interaction with
features thought to be rich in water ice, many researchers believe that the processes carving the gullies involve
liquid water. However, this remains a topic of active research. Because the gullies are so young, this would
suggest that liquid water has been present on Mars in its very recent geological past, with consequences for the
potential habitability of the modern surface. In 2014, NASA reported that gullies on the surface of Mars were
mostly formed by the seasonal freezing of carbon dioxide, and not by that of liquid water as considered earlier.
Fans and Cones
The role of liquid water on Mars has been a topic of much debate and interest since Mariner 9 returned
images of the surface showing channels resembling terrestrial fluvial channels. Alluvial fans identified on Mars
are one of the more definitive evidences for liquid water flowing on the Martian surface and preserve
Gullies on the rim of the Newton Crater
23
information about the hydrologic conditions at the time of their
formation. These fans are found in three clusters located in Margaritifer
Terra, Terra Sabaea, and Tyrrhena Terra. The fans are located in craters
dated to the Noachian period and the fans themselves are dated to the
Noachian-Hesperian boundary.
A theoretical relationship between the slope of alluvial fans and
the water to sediment discharge ratio tested against laboratory and field
data under terrestrial conditions is utilized to determine overland runoff volumes and minimum flow durations
required for the formation of large alluvial fans on Mars. These volumes were calculated for both gravel and sand
fans forming by either expanding sheet-floods or channelized flow. Where possible, bankfull water discharge at
the apex of fans was determined from the width of feeder channels. The large volumes of water required to form
the fans and the large discharge at the apex of the fan suggest that groundwater would not be an adequate
source for water to form these alluvial fans.
Valley Networks and Outflow Channels
Valley networks are branching networks of valleys on
Mars that superficially resemble terrestrial river drainage
basins. They are found mainly incised into the terrain of the
Martian southern highlands, and are typically - though not
always - of Noachian age (approximately four billion years
old). The individual valleys are typically less than 5 kilometers
wide, though they may extend for up to hundreds or even
thousands of kilometers across the Martian surface.
The form, distribution, and implied evolution of the
valley networks are of great importance for what they may
tell us about the history of liquid water on the Martian
surface, and hence Mars' climate history. Some authors have
argued that the properties of the networks demand that a
A fan (right)
and cones
(left) on the
rim of the
Gusev crater
A valley network on Mars showing dendritic pattern
24
hydrological cycle must have been active on ancient Mars,
though this remains contentious. Objections chiefly arise from
repeated results from models of Martian paleoclimate suggesting
high enough temperatures and pressures to sustain liquid water
on the surface have not ever been possible on Mars.
The advent of very high resolution images of the surface
from the HiRISE, THEMIS and CTX satellite cameras as well as the
Mars Orbital Laser Altimeter (MOLA) digital terrain models have
drastically improved our understanding of the networks in the
last decade.
Outflow channels are extremely long, wide swathes of
scoured ground on Mars, commonly containing the streamlined
remnants of pre-existing topography and other linear erosive
features indicating sculpting by fluids moving downslope.
Channels extend many hundreds of kilometers in length and are
typically greater than one kilometer in width; the largest valley
(Kasei Vallis) is around 3,500 km (2,200 mi) long, greater than
400 km (250 mi) wide and exceeds 2.5 km (1.6 mi) in depth cut
into the surrounding plains. These features tend to appear fully
sized at fractures in the Martian surface, either from chaos
terrains or from canyon systems or other tectonically controlled,
deep graben, though there are exceptions. Besides their
exceptional size, the channels are also characterized by low
sinusitis and high width: depth ratios compared both to other
Martian valley features and to terrestrial river channels. Crater
counts indicate that most of the channels were cut since the
early Hesperian age, though the age of the features is variable
between different regions of Mars. Some outflow channels in the
Amazonis and Elysium Planitiae regions have yielded ages of only
tens of million years, extremely young by the standards of
Martian topographic features.
On the basis of their geomorphology, locations and
sources, the channels are today generally thought to have been
carved by outburst floods (huge, rare, episodic floods of liquid
water), although some authors still make the case for formation
by the action of glaciers, lava, or debris flows. Calculations
indicate that the volumes of water required to cut such channels
at least equal and most likely exceed by several orders of
magnitude the present discharges of the largest terrestrial rivers,
and are probably comparable to the largest floods known to
have ever occurred on Earth. Such exceptional flow rates and the
View of the massive Ares Vallis showing the
huge outflow channels which are emerging
from the chaotic terrains
The erosional, streamlined features in Ares Vallis
25
implied associated volumes of water released could not be sourced by
precipitation but rather demand the release of water from some long-term
store, probably a subsurface aquifer sealed by ice and subsequently breached
by meteorite impact or igneous activity.
The outflow channels contrast with the Martian valley networks,
which much more closely resemble the dendritic planform more typical of
terrestrial river drainage basins. Outflow channels tend to be named after the
names for Mars in various ancient world languages, or more rarely for major
terrestrial rivers. These outflow channels often show erosional features such
as streamlined island like landmasses.
Mesas and Buttes
Though not very common, mesa and butte like erosional landforms
are seen on Mars by the orbiting spacecrafts. Such landmasses can be found
in the huge outflow channels and chaotic terrains. They are thought to have
formed in similar fashion as their terrestrial siblings. The huge outbursts of
floods which carved the outflow channels shaped them as they stood out
beacause of their resistivity towards erosion. In chaos terrains, they are
thought to be remnant ground-ice-sheets, yet to be melted.
Dunes
Many locations on Mars have sand dunes. A sand sea, made up of aeolian dune fields referred to as the
Circumpolar Dune Field surrounds most of the north polar cap. The dunes are covered by a seasonal carbon
dioxide frost that forms in early autumn and remains until late spring. Many Martian dunes strongly resemble
terrestrial dunes but images acquired by the High-Resolution Imaging Science Experiment on the Mars
Reconnaissance Orbiter have shown that Martian dunes in the north polar region are subject to modification via
grainflow triggered by seasonal CO2 sublimation, a process not seen on Earth. Many dunes are black because
they are derived from the dark volcanic rock basalt. Extraterrestrial sand seas such as those found on Mars are
referred to as "undae" from the Latin for waves.
Mesas and buttes on Mars
Longitudinal dunes in Noachis region (left) and Barchan dunes in Hellespontus region (right) as seen by HiRISE onboard MRO
26
Impact Craters
An impact crater is an approximately circular depression
in the surface formed by the hypervelocity impact of a smaller
body with the surface. In contrast to volcanic craters, which result
from explosion or internal collapse, impact craters typically have
raised rims and floors that are lower in elevation than the
surrounding terrain. Impact craters range from small, simple,
bowl-shaped depressions to large, complex, multi-ringed impact
basins. Impact craters are the dominant geographic features on
Mars. The cratering on very old surfaces on Mars, witness a period
of intense early bombardment in the inner Solar System around
3.9 billion years ago. Mars is host to a wide variety of impact
craters from few meters to few thousand kilometers in diameter.
The largest known impact crater on Mars is the Hellas planitia
which is about 2000 km in diameter and about 7 km deep.
Some of the oldest known impact craters on Mars are nearly 4 billion years old. The lack of active geology
on Mars leads to the preservation of these craters, unlike those on the Earth. Although some craters in northern
planes on Mars are found to be buried underneath a veneer of fine wind blown sediments while some are
featuring gullies.
Gale crater with its central peak- Aeolis Mons,
ellipse in the upper left part of the crater shows
landing site of the Curiosity rover
Various types of impact craters on Mars, rampart crater (left) showing its ejecta blanket around it & gullies on the rim,
crater with remnant water ice in central depression (centre) and a crater with concentric ridges (right)
27
Geological History
Much of a planet's history can be deciphered by looking at its surface and asking what came first and
what came next. For example, a lava flow that spreads out and fills a large impact crater is clearly younger than
the crater, and a small crater on top of the same lava flow is younger than both the lava and the larger crater.
This principle, called the law of superposition, and other principles of stratigraphy, first formulated by Nicholas
Steno in the 17th century. The same principles are use for Mars.
Another stratigraphic principle used on planets where impact craters are well preserved is that of crater
number density. The number of craters greater than a given size per unit surface area provide a relative age for
that surface. Heavily cratered surfaces are old, and sparsely cratered surfaces are young. Old surfaces have a lot
of big craters, and young surfaces have mostly small craters or none at all. These stratigraphic concepts form the
basis for the Martian geologic timescale.
Absolute and Relative Ages
By using stratigraphic principles, rock units' ages can usually only be determined relative to each other.
For example, knowing that Mesozoic rock strata making up the Cretaceous System lie on top of (and are
therefore younger than) rocks of the Jurassic System reveals nothing about how long ago the Cretaceous or
Jurassic Periods were. Other methods, such as radiometric dating, are needed to determine absolute ages in
geologic time. Generally, this is only known for rocks on the Earth.
Assigning absolute ages to rock units on Mars is much more problematic. Numerous attempts have been
made over the years to determine an absolute Martian chronology (timeline) by comparing estimated impact
cratering rates for Mars to those on the Moon. If the rate of impact crater formation on Mars by crater size per
unit area over geologic time (the production rate or flux) is known with precision, then crater densities also
provide a way to determine absolute ages. Unfortunately, practical difficulties in crater counting and
uncertainties in estimating the flux still create huge uncertainties in the ages derived from these methods.
Martian meteorites have provided datable samples that are consistent with ages calculated thus far, but the
locations on Mars from where the meteorites came (provenance) are unknown, limiting their value as
chronostratigraphic tools. Absolute ages determined by crater density should therefore be taken with some
skepticism.
Crater density timescale
Studies of impact crater densities on the Martian surface have delineated three broad periods in the planet's
geologic history. The periods were named after places on Mars that have large-scale surface features, such as
large craters or widespread lava flows, that date back to these time periods. The absolute ages given here are
only approximate. From oldest to youngest, the time periods are:
Pre-Noachian: It represents the interval from the accretion and differentiation of the planet about 4.5
billion years ago (Gya) to the formation of the Hellas impact basin, between 4.1 and 3.8 Gya. Most of the
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geologic record of this interval has been erased by subsequent erosion and high impact rates. The crustal
dichotomy is thought to have formed during this time, along with the Argyre and Isidis basins.
Noachian Period (named after Noachis Terra): Formation of the oldest extant surfaces of Mars between
4.1 and about 3.7 billion years ago (Gya). Noachian-aged surfaces are scarred by many large impact
craters. The Tharsis bulge is thought to have formed during the Noachian, along with extensive erosion
by liquid water producing river valley networks. Large lakes or oceans may have been present.
Hesperian Period (named after Hesperia Planum): 3.7 to approximately 3.0 Gya. Marked by the
formation of extensive lava plains. The formation of Olympus Mons probably began during this period.
Catastrophic releases of water carved extensive outflow channels around Chryse Planitia and elsewhere.
Ephemeral lakes or seas formed in the northern lowlands.
Amazonian Period (named after Amazonis Planitia): 3.0 Gya to present. Amazonian regions have few
meteorite impact craters but are otherwise quite varied. Lava flows, glacial/periglacial activity, and minor
releases of liquid water continued during this period.
The date of the Hesperian/Amazonian boundary is particularly uncertain and could range anywhere from 3.0
to 1.5 Gya. Basically, the Hesperian is thought of as a transitional period between the end of heavy bombardment
and the cold, dry Mars seen today.
Mineral alteration timescale
In 2006, researchers using data from the OMEGA Visible and Infrared Mineralogical Mapping
Spectrometer on board the Mars Express orbiter proposed an alternative Martian timescale based on the
predominant type of mineral alteration that occurred on Mars due to different styles of chemical weathering in
the planet’s past. They proposed dividing the history of the Mars into three eras: the Phyllocian, Theiikian and
Siderikan.
Phyllocian (named after phyllosilicate or clay minerals that characterize the era) lasted from the
formation of the planet until around the Early Noachian (about 4.0 Gya). OMEGA identified outcropping
of phyllosilicates at numerous locations on Mars, all in rocks that were exclusively Pre-Noachian or
Noachian in age (most notably in rock exposures in Nili Fossae and Mawrth Vallis). Phyllosillicates require
a water-rich, alkaline environment to form. The Phyllocian era correlates with the age of valley network
formation on Mars, suggesting an early climate that was conducive to the presence of abundant surface
water. It is thought that deposits from this era are the best candidates in which to search for evidence of
past life on the planet.
Theiikian (named after sulphurous in Greek, for the sulphate minerals that were formed) lasted until
about 3.5 Gya. It was an era of extensive volcanism, which released large amounts of sulphur dioxide
(SO2) into the atmosphere. The SO2 combined with water to create a sulphuric acid-rich environment
that allowed the formation of hydrated sulphates (notably kieserite and gypsum).
Martian crater density time-scale, in millions of years
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Siderikan (named for iron in Greek, for the iron oxides that formed) lasted from 3.5 Gya until the present.
With the decline of volcanism and available water, the most notable surface weathering process has
been the slow oxidation of the iron-rich rocks by atmospheric peroxides producing the red iron oxides
that give the planet its familiar colour.
Modern Geological Processes
Mars, though thought to be geologically inactive, shows some signs of geological processes. Landslides,
glacial activities, avalanches and aeolian processes are active today on Mars.
Slight changes in local terrain observed in images
captured by many spacecrafts are inferred to be results
of landslides. Canyons, valleys and crater rims are the
most prone areas to landslides, sometimes referred as
debris flows. The HiRISE onboard the NASA’s Mars
Reconnaissance Orbiter (MRO) captured the adjoining
image showing a debris field formed after a landslide
took place on the northern slopes in the central region of
Valles Marineris.
Scientists have observed avalanches also on
Mars. Martian polar caps contain a semi-permanent
residual cap beneath a surface seasonal cap that waxes
and wanes. The residual caps are largely water, but each year as the winter cold deepens the water caps become
surfaced with frozen CO2, which was thought to sublime gently in spring. This photograph proved otherwise. The
close-ups show two separate CO2/dust avalanches cascading down the side of a single 700-meter scarp. Four
different avalanches were observed in
the same MRO HiRISE shot; the upper
fall is 160 m wide. Analysis suggests the
falls originated on the sides of the
scarp rather than the top and were
triggered by spring sublimation of dry
ice. Mars is still a planet of dramatic
changes.
Martian mineral alteration time-scale, in millions of years
A landslide in Valles Marineris
Avalanches seen in an
MRO HiRISE shot
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Geological Map of Mars
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References
http://www.Marspedia.org
http://www.solarviews.com/eng/Marsvolc.htm
http://nineplanets.org/Mars.html
http://spaceref.com/Mars
http://csep10.phys.utk.edu/astr161/lect/Mars/Mars.html
http://astroengine.com
www.google.co.in
MRO HiRISE website:
http://www.uahirise.org
NASA websites:
www.nasa.gov
science.nasa.gov
Mars.nasa.gov
http://solarsystem.nasa.gov/planets
http://www.jpl.nasa.gov
Euripian Space Agency website:
http://www.esa.int/ESA
Planetary Science and Exploration Program Newsletter (January 2014, Volume 4, Issue 1)
Mars Express: The Scientific Investigations (ESA SP-1291, June 2009)