Chapter 2: Fundamentals of Planetary Science


Earth may seem rather insignificant when viewed in the context of the solar system, especially if one considers the vast expanses of empty space surrounding the planet. Nonetheless, our solar system and Earth take on great importance to our understanding of the universe because they were produced by the same processes that formed other stars and planets. The planet we live on is an accessible product of the evolution of a star; by studying it and our neighbors, the Sun and planets, we can learn much about the mechanics of planet formation and evolution.

At the outset, it is important to note that we assume that the physical and chemical laws that govern nature are constant. For example, we use observations about how chemical reactions occur today, such as the combination of oxygen and hydrogen at specific temperatures and pressures to produce water, and infer that similar conditions produced the same results in the past. This is the basic assumption of all sciences. Moreover, much of what we "know" about the planets, as in all science, is a mixture of observation and theory---a mixture that is always subject to change. Scientific knowledge is pieced together slowly by observation, experiment, and inference. The account of the origin and differentiation of planets we present is such a theory or model; it explains our current understanding of facts and observations. It will certainly be revised as we continue to explore the solar system and beyond, but the basic elements of the theory are firmly established.

Major Concepts

1. The elements, other than hydrogen and helium, were produced by nuclear processes that generally occur within stars. Massive explosions of these stars recharge interstellar space with newly formed elements.

2. Our solar system was probably formed by the gravitational collapse of a nebular cloud composed of gas and dust. The outer planets consist of volatile compounds that condensed far from the Sun, where the nebula was cool. The inner planets are poor in these constituents and rich in refractory silicates and metals that crystallized at high temperatures.

3. As these particles accreted in orbit around the forming Sun, the planets probably became hot and internally differentiated to different degrees.

4. A planet's lithosphere is its rigid outer shell; it consists of many rock types that preserve important clues about their diverse origins.

5. Atmospheres have been formed around some planets by gravitational capture of the gaseous solar nebula or by volcanic outgassing from the planets' warm interiors. Hydrospheres (or cryospheres) form if water is released to planet surfaces under appropriate conditions of temperature and pressure.

6. The thermal history of a planet determines its level of volcanic and tectonic activity. In general, planets appear to have evolved from early periods of enhanced volcanic activity and lithospheric mobility to later periods of declining or inactive volcanism. This is tied to the cooling of the planet and its thickening lithosphere. The evolution of the lithosphere is influenced by planet size, composition, and the nature of internal and external heat sources.

Origin of the Elements

Before we consider the origin and evolution of the planets we must first discuss the nature and origin of the basic material of which they are composed---the elements. An element is composed of a single type of atom, which contains a unique number of protons (positively charged particles that have a mass of one atomic mass unit, or amu). The number of protons helps to determine the distinctive chemical and physical properties of the atom (Figure 2.1). The protons reside in a central nucleus surrounded by orbiting electrons (small, negatively charged particles with almost negligible masses). Most atoms also contain a third particle, called a neutron, which has a mass of about 1 amu but has no electrical charge. Neutrons reside with protons in the nucleus but do not affect the chemical identity of an element; instead they produce isotopes of an element. These have different atomic weights and may have different nuclear characteristics (e.g., some may undergo spontaneous radioactive decay to isotopes of other elements).

Hydrogen and helium atoms

Figure 2.1 The atomic structures of hydrogen and helium illustrate the major particles in atoms. Hydrogen has one positively charged proton (center) and one orbiting, negatively charged electron. Helium has two protons (red), two neutral particles called neutrons (green) and two orbiting electrons. Neutrons and protons contain most of each atom's mass and reside within a central nucleus. As shown here, hydrogen has an atomic mass of 1 amu and helium has an atomic mass of 4 amu.

Hydrogen is the most abundant element in the universe. It is also the simplest, usually consisting only of one proton and one electron. Consequently, the atomic weight of hydrogen is one. Other isotopes of hydrogen also contain these two particles but may also have one or more neutrons, creating isotopes of hydrogen with a variety of atomic masses. The relative abundance and simplicity of hydrogen atoms suggest that they could be the building blocks of other elements. That is, if various particles of hydrogen atoms (protons, electrons, and neutrons) could be combined in different proportions, all of the chemical elements could be formed from primordial hydrogen (Table 2.2). We know that a variety of nuclear transmutations can do just that, but high temperatures and pressures are required---conditions that are found in the interiors of some stars. Once new elements have been generated, they are often ejected into interstellar space by gigantic stellar explosions (novas, supernovas, or planetary nebulas). The dispersed elements may become reconcentrated and recycled to form new stars or planets. Thus, the Sun, the planets, and even our bodies are all composed of elemental "star dust" that saw its origin billions of years ago in another star system. Like the legendary Phoenix, our solar system rose from ashes left on the funeral pyres of an older generation of stars and planets.

The stellar "alchemy" by which new elements are produced may take several forms but commonly involves the fusion of light atomic nuclei to form heavier nuclei. As a by-product of these reactions, energy is released (the life-sustaining heat and light we receive from the Sun) and various other particles are produced (gamma rays, electrons, neutrons, and even hydrogen or helium nuclei). Hydrogen (H) may "burn" to produce helium (He, with two protons) but only at temperatures that exceed 10 million K (Figure 2.2). Helium burning can produce carbon (C, atomic weight 12), which can combine with other helium nuclei to produce oxygen (O, atomic weight 16). Similarly these "ashes" may react during carbon burning to produce oxygen, neon (Ne), sodium (Na), and magnesium (Mg); during neon burning to produce oxygen and magnesium; during oxygen burning, to produce the element magnesium through sulfur (S); and during silicon (Si) burning to produce elements up to iron (Fe, atomic weight 56).

Hydrogen burning

Figure 2.2 Hydrogen burning, the proton-proton chain, is an important energy-producing nuclear reaction that takes place inside stars like our Sun. In this process two hydrogen nuclei (protons in red) combine to form a deuterium nucleus (D, hydrogen with an atomic weight of 2--red and gray) and a positron (electron with a positive charge--white). Neutrons are gray; protons are red. Reaction with an additional hydrogen nucleus (red) produces a helium nucleus with a mass of three (3He). The fusion of two such nuclei results in the production of stable helium (4He, which has two protons and two neutrons) and the ejection of two hydrogen nuclei (red), which can be consumed in other reactions of this sort. Energy is released by these reactions as a small amount of matter is converted into energy in accordance with Einstein's equation, E (energy) = m (mass converted to energy) c2 (the speed of light squared).Credit:

These latter transmutations take place at progressively higher temperatures and pressures, which are produced only in massive stars. For example, the moderate-sized star that forms the center of our solar system is not large enough to sustain the conditions necessary for carbon burning (Figure 2.3). A star more than four times as massive as the Sun may evolve through a series of stages in which the hydrogen-burning reaction moves progressively outward from the core (Figure 2.3). In the wake of this expanding shell, burning of He, C, Ne, O, and Si may be initiated in discrete shells whose outward expansion is controlled by temperature and the production of fuel from other burning reactions. Once this process reaches a stage where the star has a central core of iron (produced by silicon burning), the stage is set for one of nature's most dramatic events, a supernova explosion.

2.3 Internal structures of stars

Figure 2.3 The internal structures of stars change with their age and size or mass. (A) The interior of a small star (less than about 4 times the Sun's mass) changes as it evolves from a small hydrogen-burning star to a large hydrogen- and helium-burning star. (A) Small stars burning hydrogen and helium become cooler at their surfaces and redder, and consequently are called red giants. These giants may be 10 to 20 times the diameter of their precursor. Note how the hydrogen-burning shell (shaded) has expanded outward, leaving in its wake a helium-rich shell; eventually hydrogen-burning may extend to the surface causing the disruption of the star's surface and produce a planetary nebula. (B) The internal structure of a massive star which has evolved past a helium-burning stage. Concentrically arranged shells where burning takes place (shaded) at progressively higher temperatures are separated by unreactive shells (light) where the material is depleted in the fuel being burned in the outer shell and is too cool to participate in the burning reaction of the next inner shell. The "death" of such a massive star is marked by the production of a nova or supernova. Credit: E. Chaisson and S. McMillan, Astronomy Today, Prentice Hall.

Iron does not "burn" to create even heavier elements. Consequently, the dynamic balance between the outward-directed pressure caused by heat release and the inward-directed gravitational pressure is lost. The star begins to contract; as a result, the pressure on the core may become so high that normally unreactive protons and electrons combine to form neutrons. The inner support of the star is effectively removed. Consequently, the star immediately collapses and unburned nuclear fuels n the outer shells suddenly react as they collapse toward the hot interior of the star. The reactions proceed with such violence that the star explodes and produces a temporary beacon in our skies, which we call a nova or supernova (Figure 2.4) depending on its brightness. These explosions are responsible for dispersing many heavy elements into other parts of the galaxy. Our galaxy has seen at least seven supernovas in the last 1000 years; thus, they are not particularly unusual events.

2.4 Crab nebula

Figure 2.4 The Crab nebula is the remnant of a supernova explosion of a massive star. This chaotic mass of expanding gas and dust is correlated with the description of a supernova seen by Chinese astronomers in a.d. 1054. Such explosions are an important method of injecting newly formed elements into interstellar space, where they may eventually be recycled to form other generations of stars and planets. A spinning neutron star embedded in the center of the nebula gives the interior an eerie bluish glow and causes it to pulse 30 times a second. The colors indicate the different elements that were expelled during the explosion. The orange filaments are made mostly of hydrogen from the original star. Blue in the filaments in the outer part of the nebula represents neutral oxygen atoms, green is ionized sulfur, and red reveals ionized oxygen. Courtesy NASA, ESA, J. Hester and A. Loll (Arizona State University).

Elements not mentioned in the fusion model described above are formed by a variety of processes in stars. Elements heavier than iron are generally thought to be formed as neutrons are captured by other nuclei. These reactions may take place during the He-burning stages of a star if the proper seed nuclei are present. Rarely, reactions involving the capture of protons may result in the production of some of the same heavy elements during supernova explosions.

Yet another process is required to produce the isotopes of the light elements lithium (Li), beryllium (Be), and boron (B). These elements are unstable in the deep interiors of stars but may be produced as energetic protons collide with and fragment atoms of carbon, nitrogen (N), or oxygen. These proton-atom collisions may occur during nova explosions or during the formative stages of stars, when many protons are produced. Chance cosmic ray collisions with the proper atoms can also produce these elements in interstellar space.

Origin of the Solar System

It is believed that the Sun and all of the planetary bodies within the solar system originated from the collapse of a solar nebula. These events occurred billions of years ago; therefore, many details of the process are not completely understood. Nonetheless, observations of other developing star systems and studies of meteorites and the systematic compositional differences between the planets provide fundamental constraints on theories of the formation and evolution of the solar system.

When Did the Solar System Form?

Thought to be leftover raw materials from which the planets formed, meteorites that have fallen to Earth provide a way to look back through time and estimate when the planets formed. Thus, if we can determine when the meteorites formed, we should have a reliable estimate of the date at which the solar system itself formed. Radiometric ages provide a powerful method to determine an absolute age for meteorites and many other types of rocks. A radiometric age can be obtained for a rock, if it contains radioactive isotopes, which decay to form daughter products at a rate that is well understood. The principal elements used to date planetary materials are isotopes of uranium (which decays to lead), potassium (which decays to argon), and rubidium (which decays to strontium). The concentrations of both the parent and daughter isotopes are measured with sensitive instruments, and by using an appropriate decay constant, we can calculate the time at which the rock crystallized. The process of using radioactive decay products to determine the age of a rock is much like that of estimating the amount of time elapsed by using an hourglass filled with sand---just substitute the sand in the upper chamber for the radioactive parent isotope, the sand grains in the lower chamber for the daughter isotopes, and the rate at which sand slips from one chamber to the next for the decay rate. By counting the grains of sand in the lower chamber we could obtain an accurate estimate of the time elapsed since the hourglass was turned over. In planets, melting or recrystallization of a rock effectively resets its radiometric clocks to zero. Since most rocks suitable for dating have several different isotopes decaying at once, multiple clocks are ticking away in each rock. Using these dating techniques on meteorites and samples of lunar rocks provides ample evidence that the solar system originated about 4.6 billion years ago.

Nebular Hypothesis for the Formation of the Solar System

Although, we perceive "outer space" to be completely without elemental material, it is not a perfect vacuum. Throughout the galaxy, gases are thinly dispersed. For each 10 cm3 there may be only one atom. (Near Earth's surface, the atmosphere contains about 1020 atoms in the same volume.) The most important of these interstellar gases consist of the most abundant elements------hydrogen, helium, carbon, nitrogen, and oxygen. There are also a few metallic elements and dust grains composed of metals and silicates. We have seen how the violent deaths of some stars provide a recycling mechanism to charge interstellar space with these materials.

Occasionally, large concentrations of gas and dust, which may have approximately 1000 atoms/cm3 accumulate. Such dense, dusty clouds are called nebulas (Figure 2.5) and have been detected in several places in the galaxy.

2.5 Dusty nebula

Figure 2.5 Dusty, gaseous nebulas, such as this one, are essentially incubators for young stars. The soaring tower of the Eagle nebula is 9.5 light-years high (about twice the distance from our Sun to the next nearest star) and made mostly of cold gas and dust. A torrent of ultraviolet light from a series of massive, hot, young stars [off the top of the image] has sculpted the details of the pillar. Ghostly streamers of gas boil off the tower's envelope. At the top of the tower, thick clouds of hydrogen gas and dust have survived the blast of ultraviolet radiation from the hot, young stars. Inside the gaseous tower, stars may be forming when dense gas collapses under it's own gravitational attraction or where it is rammed together by shock waves moving through the gas. The fingers at the top and center of the tower are stellar birthgrounds that are roughly the size of our solar system. The fledgling stars grow as they feed off the surrounding gas cloud until they are separated by "erosion" from their gas supply. The blue color at the top is from glowing oxygen. The red color in the lower region is from glowing hydrogen. Credit: NASA, ESA, and The Hubble Heritage Team (STScI/AURA)

Several nebulas contain young stars; in part, this is why they are thought to be the birthgrounds of stars. Simple gravitational attraction and contraction can occur when the density of the gas is as low as 20 atoms cm3. The gravitational collapse of a nebular cloud may be fairly rapid (Figure 2.6). Small density differences and gas turbulence can produce several subregions from a large nebula. Each concentration may eventually collapse independently and become a star. Open star clusters and some multiple star systems are thought to result from the fragmentation of contracting nebulas.

2.6 Nebular evolution

Figure 2.6 The evolution of a dusty nebula with a surrounding system of orbiting planets is shown in this schematic diagram. (A) A slowly rotating portion of a large nebula becomes a distinct globule as a mostly gaseous cloud collapses by gravitational attraction. (B) Rotation of the cloud prevents collapse of the equatorial disk while a dense central mass forms. (C) A protostar "ignites" and warms the inner part of the nebula, possibly vaporizing preexisting dust. As the nebula cools, condensation produces solid grains that settle to the central plane of the nebula. (D) The dusty nebula clears either by dust aggregation into larger particles (planets or planetesimals) or by ejection during a T-Tauri stage of the star's evolution. A star energized by fusion and a system of cold bodies remains. Gravitational accretion of these small bodies eventually leads to the development of a small number of major planets.

The exact nature of the process that induces the collapse of diffuse dust and gas to form a nebula remains a mystery. A possible clue comes from studies of the most primitive class of meteorites that have fallen to Earth. These meteorites retain chemical traces of the explosion of a massive star, which injected material into the developing (but still gaseous) solar system no more than a few million years before the meteorites solidified. Some scientists have suggested that the nebula from which the planets formed may have been concentrated by such a supernova "trigger," which swept dispersed atoms closer together. It may be that the supernova that preceded the consolidation of the meteorites marked the death of a massive, short-lived companion of the embryonic Sun that later formed from the same nebula.

The contraction of the cloud of gas and dust guarantees more collisions among the atoms within it, producing heat. Some of this heat can be dissipated by infrared radiation into space; the rest is retained and elevates the temperature of the nebula. Because the interior of the nebula gradually warms, increased gas pressure causes the collapsing of the cloud to slow down. When temperatures exceed 10,000 K a protostar, probably located near the center of the nebula, may begin to radiate the light produced by the release of gravitational and thermal energy (nuclear reactions begin at a later stage). As a result, the inner portions of the nebula become much warmer than its outer reaches.

During this early contraction, the gas cloud begins to rotate, and as it collapses it rotates even faster, like a figure skater who draws in her arms during a spin. Such rapid rotation prevents a flattened disk of material in the equatorial plane from moving inward toward the protostar. The planets eventually formed from the materials in this disk. In fact, disks of dust have been discovered around young nearby stars. Heat loss by convection (radiation is hindered by dust) allows further contraction of the protostar. Rapid and irregular outbursts of light and strong magnetic fields are associated with this relatively slow gravitational contraction. It is thought that during this so-called T-Tauri phase, large amounts of matter are ejected from the nebula in a type of "wind" that sweeps much of the uncondensed gas and even some light dust from the inner part of the evolving nebular disk. This occurs as the star settles onto what is known as the main sequence, a stage in the life of a star where it is relatively stable and long-lived.

As the star continues to contract, critical temperatures (around 8,000,000 K) and pressures are reached at which thermonuclear fusion of hydrogen can be initiated. When the nuclear fires are ignited, the temperature rises farther and essentially halts further contraction. Stars the size of the Sun may maintain this equilibrium state for billions of years, as they gradually consume their budget of hydrogen to form helium. The evolution from stellar nebula to hydrogen-burning star may only take 100,000 years---a short part of the solar system's 4.5-billion-year history.

Differentiation and Condensation of the Solar Nebula

Now let us go back a short time in this grand scenario and try to construct a plausible scheme for the development of the planets in our solar system. To do this, most contemporary theories call on the gravitational coagulation of solid particles that condensed from a dusty nebula such as the one just described.

One important process that occurred in the evolving solar nebula was its separation or differentiation into several physically and chemically unlike products. The gross differences in the sizes and compositions of the inner and outer planets demand that differentiation occurred within what must have been a relatively homogeneous nebula. An important part of this differentiation was the condensation or crystallization of solid particles from the gaseous nebula.

To determine the type of solids that might condense, we must first know the chemical composition and temperature of the nebula. Most planetary scientists assume that the composition of the nebular disk was about the same as that of the present-day Sun. There is some evidence in the primitive class of meteorites called carbonaceous chondrites that temperatures approaching 1800 K may have been approached locally (dust grains may have been vaporized). Using these assumptions regarding temperature and composition and another for pressure within the disk, we can calculate the sequence and composition of the solids condensing from the cooling gas. Figure 2.7 shows the generalized results of such calculations.

2.7 Condensation of the nebula

Figure 2.7 The sequence of condensation from a gas of solar composition as it cools from about 1700 K is shown in this diagram. Shaded bars indicate the interval over which condensation probably could occur; dashed lines indicate the persistence of these materials to lower temperatures in the absence of condensation. The upper axis indicates the possible condensation temperatures of components that produced the planets Mercury, Venus, Earth, Mars, the asteroids, and Jupiter.

The first solids to condense from the nebula were probably small quantities of highly refractory elements such as tungsten, osmium, and zirconium. (Materials that form solids at very high temperatures are called refractory while those that condense or solidify at very low temperatures are called volatile.) Indeed, these elements may never have been vaporized in the nebula. The first compounds to form in significant amounts, however, were crystals of calcium and aluminum oxides, which probably condensed at temperatures of around 1700 K. Metallic iron--nickel compounds precipitated directly from the gas at about 1470 K. With continued cooling to about 1450 K, the oxide minerals reacted with the gas to form silicate minerals (those with silicon and oxygen) of calcium, aluminum, and magnesium. Magnesium-rich olivine and pyroxene are examples of these condensates. Alkali [sodium (Na), potassium (K), and rubidium (Rb)] silicates condensed at around 1000 K, forming feldspars at the expense of some minerals formed earlier. At about 700 K, previously condensed metallic iron reacted with sulfur in the gas to form troilite (FeS), an important mineral in some meteorites. At lower temperatures, some iron combined with oxygen and participated in reactions with magnesium to form silicates (like iron-rich olivine) or oxides (like magnetite). Eventually, below about 400 K, sulfates, carbonates, and hydrated silicates formed by reaction of early formed minerals with the gas.

Where the nebula was cooler than 300 K, sticky carbonaceous compounds precipitated, and at about 185 K water ice formed, probably in a blizzard of snowflakes. At even lower temperatures, volatile substances such as ammonia (NH3), methane (CH4), and nitrogen (N2) which we normally regard to be gases, crystallized as icy solids. As their constituent elements were much more abundant than the refractory elements, the condensation process quite literally snowballed at this point. It is unlikely that more volatile materials such as hydrogen (H2) or helium (He) ever condensed, even in the cold outer reaches of the nebula. Thus, the low-temperature product of condensation from a solar gas consisted of a mixture of carbonaceous materials, hydrated silicates, sulfates, ices (of water, methane, and ammonia) and uncondensed gases---mostly hydrogen and helium. Realizing that the silicate, or rocky, component of the condensed materials would only be a small proportion of the nebular mass, a condensate formed in the coldest regions of the nebula would be little more than a dirty snowball.

2.8 Thermal zonation of nebula

Figure 2.8 A cross section of a hypothetical nebula shows a star forming in its center. Condensation of solids from a solar nebula with a temperature gradient may have given rise to compositional differences in the condensates. At one instant, the condensates in the inner part of the developing solar system would consist of high temperature materials such as silicates, while at the same instant, but farther from the Sun, the nebula may have been cool enough to allow ices to be fully condensed as well. Since water was relatively abundant in the nebular gases more solid matter formed in the cooler outer part of the nebula.

We have already noted that the nebular disk around the Sun must have possessed a strong thermal gradient. It was initially very hot near the proto-Sun and must have been much cooler in its outer parts. At any one instant the composition of the solids in the nebula would be dependent on their distance from the Sun (Figure 2.8). Assume for example that the nebula's temperature near the orbit of Mercury, the innermost planet, was about 1400 K. The condensates that would exist at that temperature have been predicted to consist of metallic grains of iron and nickel as well as silicate minerals rich in calcium, aluminum, and magnesium (Figure 2.7). The density of this assemblage would be quite high. Making some reasonable assumptions, we can calculate the temperature at the orbit of Jupiter to be about 140 K at the very same moment. At this temperature a greater proportion of the nebula would be condensed; the solids would consist of hydrated and oxidized silicates and carbon-rich materials and also a large amount of low-density ice. Thus, if the solids were somehow isolated from further reaction with the gas and instead formed planets, the composition and size of these planets would be dramatically dependent on their distance from the Sun.

This theoretical model, even if it is simple, explains many of the gross features of the planets. It predicts that Mercury should be rich in metallic iron and consequently dense; that Venus, Earth and Mars should be less dense and contain more silicon, sodium, and potassium, and that the outer planets should be large, volatile-rich planets, which, because of their masses, have thick atmospheres that are gravitationally trapped remnants of the nebula. The most obvious compositional differences between the inner and outer planets can be explained by this sort of model of nebular condensation.

The differentiation of the nebula was driven by the temperature-dependent condensation of various chemical species. A requirement of this theory is that the solids and gases became physically separated. We do not know by what process this occurred, but we think that the nebula was swept clean of its uncondensed gases by an appropriately timed T-Tauri phase in the development of the early Sun. This housecleaning event may have left refractory and metal-rich dust close to the Sun, hydrated lower-density dust farther from the Sun, and low-density ices farthest from the Sun. Alternatively, the solids may have become incorporated in planets and hence incapable of communicating or reacting with the gas.

The reasons for the size differences between the inner and the outer planets are suggested by the bulk composition of the nebula itself. Even if the entire refractory or rocky component condensed, it would represent less than 0.5 percent of the total mass available in a nebula of solar composition. Since all of the inner planets are predominantly rocky, over 99.5 percent of their potential mass is missing---it must have been too volatile to condense in the warm inner region of the developing solar system. The outer planets (Jupiter, Saturn, Uranus, Neptune, and Pluto) formed from material that condensed at low temperatures (about 200 to 50 K), where ices of water, ammonia and methane could form in addition to the rocky component. These ices account for about 1.5 percent of the mass of the nebula. Substantial portions of the outer planets, especially of Uranus and Neptune, are postulated to consist of elements from these ices.

Accretion of the Planets

Such differentiation processes may account for the chemical composition of the planets, but it is unlikely that the planets simply crystallized grain by grain and layer by layer from a dusty nebula. What probably did happen is more complex. Once the condensing grains had settled by the force of gravity into the central plane of the nebula (a process that may have taken only a few thousand years), it appears that planetesimals with diameters of a few kilometers formed. Some scientists suggest that this was accomplished by gravitational grain-by-grain accretion to produce streams or clusters of small bodies that moved in nearly coincident circular orbits. Low-velocity collisions within or between the planetesimal swarms eventually led to the accumulation (not destruction) of even larger planetesimals. Some of these planets in embryo may have been as large as the Moon or even Mars, but collisions also produced small fragments that were later accreted. Some collisions may have resulted in total disruption of a body, followed by its reaccretion. For a given distance from the Sun, one body eventually became gravitationally dominant and swept up most of the material near its orbit. This process of accumulation of the planets from smaller bodies is called collisional accretion and was probably complete within a million years. Some terrestrial planets, especially the Moon and Mercury, retain dramatic evidence of the last phase of accretion. Their intensely cratered surfaces were produced by the last infalls of material that lingered in their paths even after they had assumed solid spherical shapes. In short, the elemental material of our solar system evolved from gas to dust to clots in co-orbiting streams that eventually accreted to form planets.

Jupiter (and the other outer planets) grew larger and perhaps faster than the inner planets because of the abundance of icy condensates in the cooler outer nebula. Small nebular disks also formed around the larger planets. As these miniature nebulas were probably very similar to the larger solar nebula, condensation and collisional accretion probably produced the large systems of natural satellites that encircle Jupiter and Saturn. Simultaneously, large quantities of the uncondensable nebular gases surrounding the growing icy planets became hydrodynamically unstable and collapsed onto the planets' cores to form thick, colorful atmospheres. In contrast, the present atmospheres of Venus, Earth, and Mars are largely secondary and were most likely expelled from the interiors of their respective planets rather than inherited from the nebula.

Internal Differentiation of the Planets

We have seen how differentiation of the solar nebula led to important differences among the planets. Another type of differentiation occurred within the interiors of the planets and produced a variety of layered internal structures, depending on variables such as size, density, and composition of the planetary body (Figure 2.9). Internal differentiation of the planets is a very important process in their evolution, although it goes on by a different mechanism and occurs on a much smaller scale than the differentiation processes that occurred in the nebula. The continuing internal differentiation of planets like Venus and Earth drives their dynamic geologic systems and produces, among other things, volcanoes and earthquakes.

2.9 Differentiated planetary interiors

Figure 2.9 The interiors of five planets are compared in this diagram, which illustrates the relative size of various internal components. The densities of the bodies are also given in g cm3. Although some scientists believe these layered structures are the result of layer by layer accretion from the nebula, there is good evidence to suggest that the planets were originally relatively homogeneous and that the layered internal structures are the result of planetary differentiation. Note how the proportions of silicate, ice, and metal change from bodies in the inner solar system to those in the outer solar system.

Differentiation within planets occurs because elements have distinctive physical properties (mostly density) and chemical affinities, which allow them to separate from one another. For example, a large group of elements has an affinity to and behave like metallic iron and are consequently called metals. This group includes relatively dense elements such as iron, nickel, and cobalt. Other geochemical groups have affinities for oxygen and silicon and form rocky materials called silicates. These elements include aluminum, calcium, sodium, and potassium); they form solids that are less dense than metallic iron. Materials that form low density solids only at low temperature are called ices. Of course the most common is water (H2O), but methane (CH4), and nitrogen (N2) are other important ices. Others elements, like hydrogen, helium, neon, carbon, and oxygen, combine to form stable gases and accumulate in the atmosphere of a planet. These are the atmophile elements. Several elements fall in more than one group, depending on external conditions. Distinctive features of the metals include their high densities. Apparently, metals have accumulated in the cores of several planets by gravitational sinking. Churning motions in a molten core of metal may produce a magnetic field that envelopes the planet. If the planets were originally homogeneous, core formation requires considerable redistribution of elements. This mobility of planetary materials may seem at odds with our experience with the "solid" Earth, but under high temperature and pressure rocks become weak and behave much like fluids (although they may not be molten). This plasticity allows the transport of materials during planetary differentiation, but is largely dependent on a critical temperature within the planet. Therefore, planetary differentiation is intimately tied to the thermal history of a planet (a description of temperature variations with time). The fundamental questions may be: From where does a planetary body derive its heat? How is this energy transported from one part of a body to another?

Planetary Heat Sources

The motion of atoms and molecules in an object gives it a measurable form of energy we call heat or thermal energy. As atoms vibrate or rotate more rapidly in a substance, we perceive it as an increase in temperature. As thermal energy increases, the atoms are forced farther away from one another, and the bonds that hold atoms and molecules together may be broken. The substance may melt and become a liquid or boil and become a gas. Because heat dramatically affects the mechanical and chemical properties of materials, its generation and movement are very important to the differentiation and geologic activity of all the planets. Heat is transferred in three principal ways: (1) conduction, (2) convection, and (3) radiation. Conduction is the process by which the vibrational energy of an atom is transferred to adjacent atoms and is the method by which heat is usually transmitted through rigid, opaque solids. Convection occurs in fluids and nonrigid solids as warm material expands and moves upward, displacing cooler and denser material downward. Convection within planets transfers thermal energy much more efficiently than conduction. Convection transfers matter and may cause planetary differentiation. Radiation involves the emission of electromagnetic waves from the surface of a hot body to its surroundings. The radiation of energy from the Sun to the planets is an obvious means of energy transfer in the solar system.

2.10 Planetary heat sources

Figure 2.10 The variety of planetary heat sources. (A) Accretionary heat comes from the conversion of kinetic energy to heat. This heat is trapped in the planet if accretion is rapid. (B) Core formation converts gravitational potential energy into heat as molten iron drops to the center of a planet. (C) Radiogenic heat is caused by the decay of radioactive atoms dispersed in the interior of a planet. (D) Solar energy is produced by nuclear fusion reactions in a star and is transmitted to planets in electromagnetic waves (light). (E) Tidal heating results when a satellite is repeatedly flexed by the gravitational attraction of its primary.

The major sources of heat involved in planetary dynamics are shown in Figure 2.10. Accretionary heating results from the collisional accretion of two or more bodies. It is likely to have influenced the planets during their formation because some of the kinetic energy (the energy of motion) of falling planetesimals was converted to thermal energy during impacts. If the planets grew rapidly enough, some of this energy may have melted a large part of their exteriors allowing chemical and gravitational differentiation to occur.

Heat produced by loss of gravitational potential energy may have been, in a sense, self-perpetuating by core formation. If the iron-rich cores thought to exist inside some planets formed by accumulation and gravitational segregation from a once-homogeneous planet, then a tremendous amount of heat must have been released. The loss of gravitational potential energy of each iron-rich "droplet" would be accompanied by an increase in planetary temperature. The released heat decreased the strength of surrounding materials and allowed core segregation to go on at a more rapid rate. This feedback cycle may have led to rapid, runaway core formation and the production of a heat pulse during this stage of a planet's formation. Core formation on Earth may have required a few hundred million years.

Radiogenic heat is produced by the spontaneous breakdown of the nucleus of an atom. Although we normally think of the planets as cold bodies, they all contain a certain small proportion of radioactive elements, which decay spontaneously to lighter elements and release energy as they do so. (This is in contrast to nuclear fusion, which occurs in the Sun to produce heavier elements from light ones.) The most important radioactive elements are long-lived---uranium, thorium, and potassium. Short-lived radioactive isotopes of other elements (aluminum and iodine) are not found in significant amounts today but may have been important heat sources during the early stages of planetary differentiation. The nuclear decay of these elements heats the interiors of planets that contain them; rocks may even melt, become less dense, and start on a path to the surface and possible eruption through volcanoes. Each planet once possessed such an internal heat engine, and some may still possess this heat source.

Solar energy, produced from a variety of nuclear fusion reactions, when radiated to Earth, maintains the planet's biologic activity and drives the system of circulating water and gas at its surface. Even so, this energy is not sufficient to raise the temperature of planets enough to allow any internal differentiation. A much smaller amount of solar energy reaches the most distant planets.

A less familiar type of solar energy may have heated the planets or their precursors during a hypothetical T-Tauri phase of the Sun. The strong solar winds associated with this phase may have so distorted and enlarged the solar magnetic field that they induced strong electrical currents in the planets. The outer parts of the innermost planets or planetesimals may have been heated, perhaps even melted, as a result. The outer planets probably would not have interacted as strongly with the magnetic field and would be little affected. This theory of electromagnetic heating is controversial but has been used to explain the heat developed in relatively small bodies that were parents to the meteorites that fall to Earth. If this type of heating occurred, it was restricted to the earliest evolution of the inner planets and would have most strongly affected their surface layers.

Tidal heat, another even less obvious source of heat, is important for the evolution of some planetary bodies. The gravitational attraction of a planet on its satellites may, in special cases, lead to higher interior temperatures and even large-scale melting. On Earth, tides in the oceans are raised by the attraction of the Moon. However, it is not only the oceans that bulge but also the more rigid outer layers of rock. Dissipation of the energy produced during these movements may slightly heat a planet, just as an elastic band heats up during repeated stretching. Imagine the size of the tides that might arise on the satellites of Jupiter, the largest planet in the solar system. The mutual gravitational tugs of its four largest moons cause predictable variations in the distance between the satellites and their parent planet that result in tidal variations large enough to generate sufficient heat to melt the interior of at least the innermost satellite, Io. This tidal heat, produced by a sort of mechanical heat pump, appears to have played an important role in the differentiation and evolution of the small moons of Jupiter and Saturn.

In summary, there are several different forms of energy and types of processes involved in the thermal evolution of the planets. The most important heat source for the original differentiation of the planets may have been accretionary heating. Short-lived radioactive isotopes of aluminum and iron may have substantially heated small planetesimals, but it seems unlikely that these elements persisted much past the planetesimal stage. Likewise, electromagnetic heating associated with a T-Tauri phase probably occurred before the planets accreted. Long-lived radioactivity, although important in sustaining the differentiation and geologic activity of the planets, probably contributed as little as 300 K to the temperature of the primitive undifferentiated planets. In contrast, thousands of degrees of heat may have been generated by collisional accretion. It is estimated that the average temperature of Earth would have been about 30,000 K if all of this heat had been retained (1300 K for the Moon, 4000 K for Mercury, 6000 K for Mars, and 25,000 K for Venus). Of course, much of this heat was quickly radiated away into space, but since iron- and magnesium-rich silicates melt at temperatures of about 1400 to 4000 K (increasing with pressure) and ices melt at temperatures below 300 K, retention of even a small fraction of the accretionary heat could lead to melting or even vaporization of planetary materials. Once initiated by accretionary heating, the process of differentiation may have been invigorated by core formation.

Thermal History

Planetary scientists attempt to reconstruct the thermal history of a planet from facts (and assumptions) about its surface features, its chemical composition (the abundance of radioactive elements and metals, ices, and silicates), its original temperature distribution, and its accretionary history. These factors help to determine the change of temperature with depth and give us an idea about the internal temperature at different times during the planet's history. Indeed, the geologic history of a planet is a reflection of its thermal history. The size of a planet is an extremely important factor, which influences how rapidly a planet can lose the heat released from its interior. All the heat generated within a planetary body must eventually be transported to the surface by conduction or convection and then radiated away. Thus, small planets with large surface area to mass ratios cool faster as heat is readily radiated away into space. Large planets lose their heat much more slowly. For example, a large, shallow pan of water placed in a refrigerator cools more rapidly than the same amount of water in a glass. This is a result of the larger surface area to mass ratio of the water body in the pan. The planets Mercury, Mars, and Earth, given the same initial temperature and composition, would not cool at the same rate. Mercury, being the smallest, would cool the fastest (Figure 2.11).

2.11Surface area

Figure 2.11 The cooling rate of a planet depends on the surface area to mass ratio. A body with a large value for this ratio would cool relatively rapidly because of the large area through which heat may be lost to space relative to its mass.

As a planet's temperature varies, either by heating during accretion or subsequent cooling, substantial changes occur in the nature of the geologic processes that shape its interior and surface. A short scenario of a possible thermal history for a terrestrial type of planet shows some of these principles (Figure 2.12). Shortly after (or during) the formation of a planet, its outer shell may become entirely molten as the result of the heat generated by many impacts of planetesimals during accretion. As this magma ocean cools, a crust (or more precisely a lithosphere---the solid, rigid outer shell of a planet) composed of light silicate minerals rich in silicon, aluminum, and oxygen, may form at the surface. It is easy to imagine that the lithosphere would be very mobile at this stage and could slide across the surface as a silicate scum. Numerous small plates, slabs, or "rockbergs" would form as the melt cooled. Collisions and rifts, driven by vigorous convection, must have been fairly many during this period of crustal accumulation and high surface temperature. With continued cooling, a dense mantle composed of silicate minerals rich in iron and magnesium probably formed beneath the chemically distinct crust, adding to the thickness of the mechanically distinct lithosphere. Extensive melting of at least the outer portions of a planet may have allowed the separation and gravitational accumulation of metallic iron and other dense metals into distinct droplets. Once blobs about 100 km in diameter formed, they could penetrate the lower, previously undifferentiated, part of the planet and form its core. Core formation probably accompanied the accretion of the planets. For example, many planetesimals that collided to form Earth may have previously developed cores. Differentiation caused a tremendous redistribution of mass and should release more heat. Further differentiation of the planet could occur as silica-rich magmas were formed and extruded onto the surface to form volcanoes, or cooled beneath the surface to form plutons. During early differentiation some of the volatile, atmophile elements would be released from the rocky materials as fluids or gases and accumulated at the surface to form atmospheres or hydrospheres. Once core formation is completed, internal geologic activity is maintained by heat released from long-lived radioactive decay. However, in most cases the heat generated by this process is exceeded by heat lost to space, causing the lithosphere to become progressively thicker. Solidification of silicate minerals may begin simultaneously at the core-mantle boundary, producing an ever diminishing volume of partially molten or solid-but-weak material called the asthenosphere. This convecting plastic zone allows the more rigid lithosphere to slide and shift about and is a ready source of magma for volcanoes.

2.12 Thermal evolution

Figure 2.12 The thermal evolution of a terrestrial planet shows the changing temperature inside the planet. The time scale is relative. (Compare this diagram to those in chapters 3, 4, 5, and 6, which include absolute time estimates.) The occurrences, timing, and relative importance of these processes are unique to each planet and are determined by the planet's composition, mass, heat budget, and other characteristics.

Most planets loose their internal heat not only by conduction but also by convection of their interiors. Cylindrical plumes of material rise or fall depending on their temperatures and densities. Hot rising plumes, which bring heat to the surface where it can be radiated away, may cause the lithosphere of a planet to bulge, and cold sinking plumes may depress the surface of a planet. On some planets with thin lithospheres, a different style of shallow convection may develop wherein large lithospheric slabs or plates are continually created and consumed as they shift about driven by gravity. Linear belts of volcanoes, folded mountains, and earthquakes result from collisions or rifts of these plates. This is called plate tectonics and is characteristic of Earth's present stage of development. Tectonics denotes the large-scale deformation of the lithosphere. Earth's plate-tectonic system, which produces continental drift, is a reflection of the continued differentiation of Earth. Eventually, as a planet continues to cool, the lithosphere may become so thick that the lateral motion of separate sections of lithosphere is no longer possible. Separate plates coalesce to form a single, planet-encircling, solid, rigid lithosphere. Vertical movements and lithospheric flexures related to plumes may be the only expressions of the convective flow of mantle materials at this stage. Even later, as the lithosphere continues to thicken, hot magmas, cannot penetrate through the lithosphere before they cool. The geologic evolution of a planet is nearly complete by this stage, as its internal heat source (generally natural radioactivity) gradually dies and internal activity ceases to be expressed at its surface. This inexorable cooling may lead to slow planetary contraction. Chemical differentiation slows and eventually ceases, as there is insufficient heat to cause, or even allow, matter to migrate or segregate any longer.

William Kaula, paraphrasing a famous passage from Shakespeare, has summarized this grand evolutionary scheme.


Our system is a stage,
And both the Sun and planets merely players.
They had their birth and'll have their fiery end.
A planet in its time plays many parts,
Its acts being seven ages.
The first of these is condensation: dust grains drifting to
The nebula plane in chondrite clods. And then
The planetesimals: breaking sometimes, but
Most growing, though the Sun's hot breath blows gas
Away. And then formation: sweeping up
The bodies in its way, in fierce infalls
To bring them full convective vigor, too hot
For crust to form, though iron may sink and seas
Outgas, by radioactive energy driven.
And then comes plate tectonics: cooling leads
To lithosphere, with many marginal breaks.
Convective thrusts a crust create in belts
Complex. But heating slows; the sixth age shifts
Into final volcanism: no
More lithospheric spreading, only vents
For magma, Nix Olympica* or mare**
To surface, ending fractionation. Last scene
That ends this history is quiescence: time
Sans melt, sans plate, sans almost everything.

*A large volcano on Mars, usually called Olympus Mons.
**Vast volcanic plains on the Moon.

Planetary bodies whose evolution and differentiation are not limited by their budget of radioactive elements may continue to be dynamic long after their radioactive heat is dissipated. Tidally flexed Io, for example, is only slightly larger than Earth's Moon, yet Io continues to be volcanically active while the Moon ceased to be active almost 3 billion years ago. Nonetheless, the extent, duration, and style of planetary differentiation in most planets are complex functions of their sizes and compositions. For example, small bodies lose heat so rapidly that they quickly move through all these stages in a billion years or so (e.g., the Moon). Very small bodies such as the moons of the outer planets, may not have become warm enough to differentiate at all, and may preserve their primitive compositions to the present day. Larger planets, which cool slowly, evolve at a more leisurely pace, allowing much time for extensive chemical separations to occur (e.g., Earth).

The central theme of this brief summary of the origin and differentiation of planets is that each has a life cycle related to its thermal history. The internal differentiation of planets is a result of this interaction of heat with planetary matter.

Lithospheres, Hydrospheres, and Atmospheres:
Products of Planetary Differentiation

The part of a planet that is visible and that may be directly sampled consists entirely of its outermost layers (its atmosphere, hydrosphere, or lithosphere---all products of planetary differentiation). Only the surface, the very skin of the lithosphere, is visible from space, and even this is sometimes obscured from our view by thick or opaque atmospheres (e.g., Venus and Titan) whereas others have no atmosphere at all (e.g., Moon and Ganymede). Before we delve into a study of each planet, let us turn to examine these accessible parts of the planets generally .

Planetary Lithospheres

A planetary lithosphere is commonly conceived of as a mechanically strong, rigid outer layer that overlays a plastic or ductile asthenosphere. Planetary asthenospheres are commonly partially molten. From this point of view, then, the lithosphere is not a chemically distinct unit; instead, it is usually composed of a planet's crust and the upper part of its mantle (Figure 2.13). Obviously, the temperature within the interior determines the thickness of the rigid lithosphere and hence, the depth to the top of the asthenosphere. Thus, the origin and evolution of a lithosphere is strongly tied to the thermal history of a planet. If a planet formed hot, an early lithosphere would be thin and mobile and fragments would slide about the slippery asthenosphere below. But as a planet cools, its lithosphere thickens and slowly becomes immobile. The link between the thermal evolution of a planet and its surface geologic features is, therefore, the lithosphere.

2.13 Outer layers of planets

Figure 2.13 The terminology used to describe the outer layers of a planet depends on one's point of view as shown in these cross sections of a planet. Chemically a planet is differentiated into a crust and mantle of distinct compositions. Mechanically the crust and upper mantle of a planet may behave as a single "rigid" unit called the lithosphere. Underlying the lithosphere is the asthenosphere, a semiplastic zone that yields to flow much more easily than does the lithosphere. In some cases the asthenosphere may extend to the base of the crust.

The lithosphere of a differentiated planet consists of its solid outer shell and is composed of a variety of solid materials that we call rocks. The nature of these rocks provides important information about the evolution of a planet because their composition and physical character change in response to external physical conditions (pressure, temperature, or composition of the surroundings). A fundamental laws of nature holds that all materials attempt to achieve a balance with the chemical and physical forces exerted upon them and arrive at a state of equilibrium. This is also the state of lowest total energy. This effort results in progressive changes when planetary materials are exposed to environments different from that in which they formed. Although this equilibrium state is the preferred state of all systems, there are many intermediate or metastable states. All geologic materials experience various metastable states, but all changes tend toward achieving a physical or chemical state that is more stable in the present environment.

Because rocks may retain metastable characteristics in their physical and chemical nature, they record information about their mode of formation. Therefore, the composition of a rock preserves some information about its past. For example, the temperature and pressure of a volcanic lava can be deduced from careful chemical analyses of its now-solid products. Even the internal structures of a sand dune, stream channel, or lava flow are commonly preserved in resultant rocks, to inform geologists of the environment in which the rock body was originally formed. Rocks are records of past events and in their composition, structure, and sequence relations to other rock bodies are records of their formation and evolution.

2.14 Gabbro

Figure 2.14 A hand specimen (shown 4 times it's actual size) of a dark iron- and magnesium-rich rock called basalt, composed mostly of the silicate minerals olivine, plagioclase, and pyroxene. The individual mineral grains of olivine (green) are large but the other minerals have very fine grain sizes; the aggregate nature of the rock is clearly demonstrated in the inset, which shows a magnified view of a thin slice of such a rock. Discrete mineral grains of different compositions interlock in this typical igneous fabric, which developed as crystals formed from molten rock.

To put it less abstractly, rocks are simply aggregates of smaller entities called minerals (Figure 2.14). By definition, these naturally occurring solids contain a distinctive set of elements arranged in a well-defined internal structure. As a result, each mineral has definite physical and chemical properties by which it can be identified. Differences between minerals arise from the kinds of atoms they contain and the way in which they are arranged (Figure 2.15). Many minerals grow when atoms from a surrounding liquid or gas are added to its crystal structure. Mineral grains may also grow from other solids in the absence of either gases or liquids by recrystallization in the solid state.

2.15 Silicate tetrahedron

Figure 2.15 The silicon-oxygen tetrahedron is the basic building block of silicate minerals (top). Four large oxygen atoms are arranged in the form of a tetrahedron with a small silicon atom in the small central space. This basic tetrahedron can be combined with others in a variety of ways to create the complex internal atomic structures of silicate minerals. Chains or sheets are created by sharing oxygen ions between two tetrahedral units.

It is useful to divide the major planetary materials into three major groups---metals, silicates, and ices. (With the addition of gases, these are the major constituents we mentioned in the discussion of the solar nebula.) Silicate minerals and ice probably comprise over 95 percent of all planetary lithospheres. Dense metallic minerals (principally iron compounds) are rare (but economically important resources) in the lithospheres of differentiated planets, but the bulk of a planet's metallic minerals probably lie in its core. However, metallic minerals may be important at the surfaces of some fragmented asteroids whose deeper interiors are now exposed. The most important silicates are compounds of iron, magnesium, aluminum, calcium, sodium, and potassium. These elements are linked to tetrahedra of silicon and oxygen to form a variety of crystal patterns and form minerals with moderate density and relatively high melting temperatures (about 1000 to 1300 K). Silicate minerals make up the bulk of the lithospheres of the inner planets. Icy minerals are important at the surfaces of the satellites of the outer planets and in the cold polar regions of Mars and Earth. (Some investigators balk at the use of lithosphere---rocky sphere---to describe an icy planet's outer layer and prefer cryosphere---ice sphere.) The principal ices consist of water, ammonia, and methane. Water is the most abundant of these ices, but on most icy planetary bodies it is probably admixed with ammonia. Such a mixture may allow liquids to persist to temperatures below the freezing point of water (273 K at 1 bar). Methane ice forms at 109 K (at 1 bar) and is only stable at the surfaces of planets in the far outer reaches of the solar system (e.g., on Pluto and on Neptune's satellite, Triton). These disparate compounds are grouped together as ices because of their low freezing temperatures (300 to 100 K) and uniformly low densities (less than 1.0 gm cm3). Depending on differentiation histories, other types of minerals may also be important on specific planets (e.g., carbonate minerals are an important but small part of Earth's lithosphere, and various sulfur-based minerals are important at the surface of Io).

Besides these compositional groups, three broad groups of rocks can be distinguished by their mechanism of formation: (1) Igneous rocks form by cooling of a melt; (2) sedimentary rocks form by erosion, transportation, and deposition of rock particles; and (3) metamorphic rocks form by changes resulting from heat, pressure, or the introduction of new elements without the intervention of a melt.

Igneous Rocks

Igneous rocks form from molten rock material called magma. A magma may include crystals (solid), melt (liquid), and volatiles (gases). Several types of magmas are important on the planets. The lithospheres of the terrestrial planets have been built by silicate magmas. Figure 2.16 shows the broad range of silicate magma compositions. They may range from basalt or gabbro (typically hot---about 1500 K---highly fluid, containing up to about 50 percent silica (SiO2) to rhyolite or granite (typically cooler---about 1000 K---and viscous, containing up to about 77 percent silica). On the satellites of the outer planets, water magmas were important. Although on Earth we do not generally think of water as a magma, the processes of melting at temperatures above normal surface temperature, mobilization, and cooling of watery liquids lead us to regard water as a magma on these cold bodies, whose surface temperatures are only about 50 to 100 K. (Temperatures as high as 275 K are required for the existence of pure water melts on these bodies.) Ammonia and other chemicals may also be dissolved in these melts and subsequent solidification products. Sulfur magmas may be important on Io, the innermost satellite of Jupiter, where extensive internal differentiation has led to the accumulation of an outer sulfur-rich shell. Melting of this material would produce volcanic eruptions of fluid sulfurous lavas with temperatures of about 400 K.

2.16 Igneous classification

Figure 2.16 The classification of igneous rocks is based on composition and texture. Granite is the most abundant intrusive rock, whereas basalt is the most abundant extrusive rock on Earth.

The classification of igneous rocks is based upon texture (mostly the size of the constituent mineral crystals) and composition (Figure 2.16). Most volcanic or extrusive rocks, igneous rocks that cooled at or very near the surface of a planet, are fine-grained because they cooled rapidly. In contrast, many plutonic or intrusive rocks have larger grains because their magmas cooled slowly beneath the ground, allowing extensive growth of crystals. High-silica (silicic) magmas produce rocks of the rhyolite family, with the minerals quartz, potassium feldspar, and sodium plagioclase (Figure 2.16). Low-silica magmas, rich in iron and magnesium (mafic), produce rocks of the basalt family, with the minerals calcium plagioclase, pyroxene, and olivine. Magmas with an intermediate composition produce rocks of the andesite family with compositions between that of rhyolite and basalt. Along with anorthosites (plagioclase feldspar rocks), these igneous rocks are the most important constituents of the crusts of the terrestrial planets. In contrast, the mantles of these planets are thought to be formed from dense olivine-rich rocks such as peridotite.

2.17 Volcanic landforms

Figure 2.17 A large variety of volcanic landforms is found on the planets; their appearance depends upon the composition of the magma, its eruption rate, and its eruption style. Basaltic lavas are fluid; they are commonly extruded along fissures and flood the surrounding area. Magmas rich in silica are more viscous and form large volcanic mountains or erupt to form vast sheets of volcanic ash.

Eruptions of magma may occur because the hot molten magma is less dense than the surrounding solids. The magma thus tends to rise, and reach the surface to erupt through a crack or through a pipe-like conduit. The eruption style of a magma depends on many complex variables, including the composition, viscosity, gas content, and volume of the magma. For example, basaltic magmas are fluid and are commonly extruded in quiet eruptions issuing from fissures in the lithosphere. They produce successions of thin flows that cover broad areas. Basalt eruptions commonly form small cinder cones and large and small shield volcanoes (Figure 2.17). Water magmas might erupt in a fashion similar to geysers along a fissure, but because they are so fluid the flows would be very thin and topographic volcanoes would be difficult to detect. The fluid eruptions of basaltic or icy lavas may resurface a very large part of the planet, burying the older landforms beneath it. In contrast, andesitic and rhyolitic magmas are extremely viscous and, for silicate magmas, contain large amounts of dissolved gas. They therefore explode violently and extrude lava flows in thick, pasty masses or turbulent rapidly moving ash flows and sheetlike ash falls. These eruptions may produce lava domes, composite volcanoes, or large calderas with ash-flow shields (Figure 2.17). Masses of igneous rock formed by the cooling of magma beneath the surface are called plutons or intrusions. Although volcanic rocks and landforms are much more obvious, plutonic rocks comprise the bulk of many planetary lithospheres. They may be especially common on icy planets because of water's unique properties---its liquid state is more dense that its solid state---which could allow large magma chambers to collect in the lithosphere without rising and erupting at the surface. Some of the variety of intrusive rock bodies are shown in Figure 2.18. Their shapes and sizes are strongly dependent on the stress regime (compression or extension) in the lithosphere.

2.18 Intrusions

Figure 2.18 Intrusive rock bodies may assume a variety of shapes. Large bodies are elliptical, smaller bodies are formed as magma is squeezed into cracks and zones of weakness in the surrounding rocks.

Magma originates by partial melting of the interior of planets. Magmas that reach the surface were probably generated in the lithosphere or asthenosphere of a planet. It is extremely unlikely that planetary magmas arise from their metallic cores. The heat necessary for lithospheric melting is commonly thought to be transported by "primary" magmas, derived in an asthenosphere, into a lithosphere. Lithospheric and especially crustal rocks may be brought to their melting temperatures by the heat added in this way. Since the first fraction of melt formed usually has a composition very different from its parent material, the processes of partial melting and magma migration promote the internal chemical differentiation of the planets.

Sedimentary Rocks

Sedimentary rocks originate from fragments of other rocks, from minerals precipitated chemically, and, at least on Earth, from accumulation of organic matter. Although the natures of these materials contrast sharply with one another in many regards, all accumulate at or near the surfaces of planets under the temperatures and pressures characteristic of the planet. Sedimentary rocks are volumetrically minor components of planetary lithospheres but are important surface materials. On Earth, less than 5 percent of the lithosphere consists of sedimentary rocks, and yet they cover 75 percent of the exposed land masses.

Four major processes are involved in the genesis of sedimentary rocks. Disintegration of rocks can occur because both mechanical and chemical processes break down rock material at a planet's surface. Transportation moves these particles and allows them to accumulate elsewhere during deposition processes. Finally compaction and cementation convert loose particles into solid rock. It is easy to visualize these processes---rock debris is produced on a volcano's slopes, carried off by streams that drain the mountain, subsequently collected as particles of sand or clay in a delta at the mouth of a river, and lithified, as cements form between grains. The same general processes also occur to produce deposits of impact breccia associated with impact craters. (Breccias are rocks composed of angular fragments of older rocks.) In this case, disintegration is caused by the transfer of a meteor's kinetic energy to the planet's surface and resultant fragmentation as the strength of mineral crystals is exceeded. The debris is transported along arching ballistic pathways far from the site of origin and deposited in a sheetlike accumulation around the newly created crater. Subsequent burial or other processes help to cement the rock fragments together creating what might be called a sedimentary rock. Another class of sedimentary rocks, however, have much in common with igneous rocks. Called evaporites, these rocks form as minerals precipitate from briny water. These include deposits of salt (NaCl) and gypsum (CaSO4.2H2O). In principle, there is little difference between the crystallization of a brine and that of a magma; both are governed by the same chemical laws. Evaporites may be important types of rocks or intergranular cements on Mars, as they are in terrestrial deserts. Other chemical precipitates are also important. Limestone and other carbonate rocks and minerals precipitate from terrestrial water, sometimes aided by biologic processes. Deposits of terrestrial coal and petroleum originate as organic matter accumulates, as plants and animals die.

2.19 Grand Canyon

Figure 2.19 Sequences of sedimentary rock in the Grand Canyon show the characteristic stratification or layering of this rock type (A). The cross section in (B) emphasizes the layering and shows the difference between the sedimentary rocks and the older igneous and deformed metamorphic rocks near the floor of the canyon.

Reflecting their nature as thin units of surface rock, sedimentary rocks typically occur in layers. In a sequence of sedimentary layers, the youngest is at the top and the oldest is at the bottom, reflecting the law of superposition (Figure 2.19). Sedimentary rocks cover large parts of some planetary surfaces; if impact breccias are included, such would be the case for most planetary bodies. Other important structures visible in sedimentary rocks (such as crossbedding, graded bedding, ripple marks, and mud cracks) reveal much about their formation. All of these items tell about the environment of deposition (Figure 2.20); for example, marine limestones prove that oceans once inundated large parts of Earth's continents millions of years ago. Unfortunately, detecting these features requires examination of samples and outcrops---information not usually available to a planetary geologist.

2.20 Sedimentary environments

Figure 2.20 The major sedimentary environments on the planets are represented in these idealized diagrams. Sedimentary environments can be divided into four types---(A) Impact produces an ejecta blanket of breccia. (B) Volcanic eruptions produce deposits of volcanic ash that are carried by the wind on planets with atmospheres. (C) Wind transports and deposits sediments in sand dunes and layers of dust. and (D) fluvial marine.

Metamorphic Rocks

The third major group of rocks are those produced by metamorphic changes that occur deep within a planetary interior or at its surface as a result of impact. Metamorphism results mainly from changes in temperature and pressure or from changes in the chemical environment of a mineral, and excludes those changes that result in melting.

These changes cause new minerals to grow, and also cause changes in textures and structural elements of the rock. The diagnostic features of the original sedimentary and igneous rocks are greatly modified or completely obliterated by metamorphism. The deeper parts of most planetary lithospheres probably consist of metamorphic rocks after original igneous rocks. As in other rock groups, metamorphic rocks are classified on the basis of texture and composition (Figure 2.21).

2.20 Sedimentary environments

Figure 2.21 Metamorphic rocks form by the recrystallization of other rocks, commonly in the presence of deforming stresses. This sample of gneiss, composed mainly of quartz and feldspar, shows deformed layers---evidence of internal deformation that occurred deep within a planet's crust.

Metamorphism can result from a variety of processes, typically by deep burial or from heat released from igneous intrusions. Most metamorphic rocks of regional extent develop in the deeper parts of the lithosphere. This type of metamorphic terrain is only exposed at the surfaces of planets where lithospheric deformation is sufficient to expose such deep levels to erosional stripping. For example, these types of rocks are probably rare at the surface of Mars, but large areas of Earth's continents are underlain by complexly deformed metamorphic rocks that formed in the roots of mountain belts and became exposed by extensive erosion (Figure 2.22). Metamorphism also accompanies meteorite impact. Shock metamorphism changes the internal structure of many mineral species. The metastable persistence of these minerals at normal surface pressure is one of the strongest evidences for meteorite impact and must produce the most abundant type of metamorphic rock on the surfaces of cratered planets like the Moon and Mercury.

2.22 Canadian shield

Figure 2.22 The characteristics of metamorphic rocks on Earth's continents are shown in this photograph of the Canadian Shield, taken from an altitude of 15 km. These rocks have been compressed and deformed to such an extent that many original features such as horizontal bedding have been obliterated. Metamorphism and folding occurred at great depths. The metamorphic rocks were then intruded by granite plutons (light tones) and cut by fractures. The area was than eroded to expose the complex rock sequence.

The evolution of planetary lithospheres is a function of their composition and the physical conditions (pressure, temperature, and state of stress) that are imposed upon them. Rocks result from the flow of energy in this system which drives planetary differentiation. Figure 2.23 shows how rock types change in response to a planet's physical and chemical environment. This cycle can be envisioned as beginning with the crystallization of a magma and proceeding, through interaction with the atmosphere and hydrosphere, to the production of sediments, which lithify to become rocks. Under conditions of higher pressure and temperature, sedimentary rocks may become metamorphic rocks, or, with the introduction of large amounts of heat, magma may again be generated, to start the cycle again. Obviously, this is a highly schematic version of the rock-generating processes in planetary lithospheres. Sediments may be derived from other sedimentary rocks or from metamorphic rocks, and not all metamorphic rocks were originally sediments. The atmosphere of a planet is itself the result of processes similar to those that give rise to igneous rocks. These observations are expressed by appropriate lines on the diagram. Placed centrally within this scheme is the energy for the modification of planetary lithospheres. Here we have included not only mantle sources for energy but also the energy provided by accreting planetary material. Accretion (which is continuing even today) was probably the ultimate source of all planetary material, but, more important in this context, it helped form large volumes of shock metamorphic rocks and may have initiated magma generation on the forming planets, important processes in the geologic development of planetary lithospheres.

2.23 Rock cycle

Figure 2.23 The possible changes of one rock type to another are shown in this diagram. By interaction with the atmosphere and hydrosphere of a planet (surface processes), igneous rocks can be broken down into sediments. These may then be compacted and cemented to form sedimentary rock. Igneous and sedimentary rocks may be metamorphosed. Metamorphic rocks may be melted to form magma giving rise to igneous rocks.

Planetary Atmospheres and Hydrospheres

A geologist is principally interested in the nature and evolution of the various features of the lithosphere. However, an important class of planetary landforms and materials (mostly sedimentary rocks) is the result of the interaction of the atmosphere or hydrosphere with the lithosphere. For example, the ubiquitous river valleys of Earth, the vast sheets of sand and dust on Mars, the chemical precipitates of Titan, and the decomposition of rocks on Venus---all are the result of fluids decomposing, abrading, transporting, or combining with planetary surface materials. In fact, the rocks and oceans of Earth contain a large proportion of the elements that once resided in its atmosphere. In the larger scheme of things, many planetary atmospheres are only the outer layer (albeit gaseous) of a differentiated planet, and naturally fall under the purview of the planetary geologist.

We have already pointed out that there are two types of planetary atmospheres (Figure 2.24). The first type, called a primary or inherited atmosphere, consists of gases gravitationally trapped from the solar nebula. The outer planets appear to have atmospheres of this sort. It is likely that these planets attained large sizes before the gaseous nebula was dispersed. Once the growing planets reached a critical size, the nebular gases became hydrodynamically unstable in their dispersed condition and collapsed around the planets to form dense and often colorful atmospheres. The composition of these atmospheres should be almost the same as that of the nebula, rich in the light elements. Jupiter and Saturn consist predominantly (more than two-thirds by mass) of hydrogen and helium, which appear to have been added to the planets in this way. Uranus and Neptune have smaller proportions of these light elements (less than one-fourth). Apparently the original rock-and-ice cores of Uranus and Neptune did not grow large enough to trap as much of the surrounding nebular gas.

It seems reasonable that the inner planets once had similar "inherited" atmospheres. However, careful measurements of the composition of these atmospheres demonstrate that they retain little or none of this primitive component. Instead, strong solar winds may have stripped away the primitive atmospheres of the already accreted inner planets. This may have been part of the T-Tauri phase of the early Sun, which also cleared away the nebular gases and terminated condensation. An alternative hypothesis is that the nebula was swept away before the inner planets (but after the outer planets) had accreted from parental materials. We have little evidence to help us establish which theory is correct, but it seems clear that the gases that now surround the inner planets were formed after the accretion of the planet and are thus secondary atmospheres. Many investigators believe that they were derived from the interior of the planets, by a process commonly called outgassing.

The most important constituents of secondary planetary atmospheres are various combinations of the elements carbon, hydrogen, and nitrogen, which were incorporated into the planets during accretion. As pointed out earlier, the condensation of the solar nebula led to the formation of hydrated silicates at temperatures of about 400 K. Carbonaceous compounds formed at slightly lower temperatures. Thus some of the elements (hydrogen, oxygen, and carbon) that eventually formed atmospheres probably became chemically incorporated into solid materials. At these temperatures, small amounts of other unreactive atmophile elements (such as nitrogen, argon, and neon), though not chemically bound in the solids, may have become trapped into the structures of minerals condensing from the nebular gases. Alternatively, ices with some of these elements may have formed at temperatures approaching 100 K and become incorporated into the planets by accretion. Thus, the elements that eventually formed the atmospheres may have resided in nebular condensates in different ways: (1) as chemical constituents of rocky minerals, (2) as occluded elements in other solids, and (3) as ices.

Earlier in this chapter we arrived at the preliminary conclusion that Earth accreted from materials that condensed at temperatures of about 600 K, well above the temperature at which water could be incorporated into solids condensing from the solar nebula (Figure 2.7), and yet over 70 percent of the surface of Earth is now composed of water. Likewise, Venus possesses a dense, carbon dioxide-rich atmosphere, but it should consist of materials that condensed at temperatures above the stability point of carbon-containing solids. Finally Mars, which must have accreted from materials that formed at lower temperatures than the materials that formed Earth, has only a tenuous carbon dioxide-rich atmosphere and small quantities of polar ice (water and carbon dioxide). These inconsistencies are the result of many factors (which will be discussed in subsequent chapters), but certainly they suggest that the planets probably accreted from a variety of materials condensed at different temperatures. Following this interpretation, Earth may have formed from a mix of planetesimals, some that formed beyond the orbit of Mars and, hence, were hydrated, and drier planetesimals that condensed closer to Venus. Thus, each planet may contain components that represent a relatively broad range of nebular equilibration temperatures, but the dominant material probably came from near the planet's present orbit. Some scientists have suggested that the accretion of atmophile substances was a late-stage event. They envision a cloak of volatile-rich planetesimals and comets, formed in the outer solar system, being added to the already formed planets. The influence of Jupiter's gravity field disrupted the original orbits of these planetesimals and led them to eventual collision with the inner planets.

Whatever the mechanism by which these atmophile elements became incorporated into the planets, their segregation and concentration into a gaseous envelope surrounding the planet resulted from its internal differentiation. At the temperatures of planetary differentiation, some volatile elements are released from the rocky materials which bind them. This may be accomplished by melting, or simply by metamorphism at elevated temperature (which may, for example, drive water out of hydrated minerals to produce anhydrous solids and a fluid). Because the resultant fluids or gases have relatively low densities, they rise through fractures and along grain boundaries to the surface and accumulate there. In contrast to primary atmospheres, those produced by outgassing are relatively poor in helium and most hydrogen is combined with oxygen or carbon to form water vapor (H2O) or methane (CH4) and accompanied by carbon-oxygen gases such as carbon monoxide (CO) or carbon dioxide (CO2) and nitrogen (N2). Argon, other noble gases, and sulfurous compounds are also released during the outgassing of nebular condensates. Free oxygen (O2) was probably only a very minor part of the gases vented to the surface of the differentiating planets. Ices, in planets that also incorporated them, also melted during differentiation. The satellites of the outer planets, for example, contain outer icy shells many kilometers thick, composed principally of water and ammonia (NH3), which were probably formed from accumulations of melted ices. Nonetheless, it seems unlikely that the differentiation temperatures for these small bodies were high enough to expel water bound in hydrous silicate minerals; these materials probably constitute their cores.

The heat provided by accretion may have been sufficient to release some of these gases, which, because of their very low densities, then accumulated at the growing planet's surface. Core formation may have been the impetus for rapid outgassing of some planetary bodies (Figure 2.24). In either case it seems that the secondary atmospheres (or for that matter oceans and icy mantles) of many planetary bodies formed very early in the planets' histories. In addition, Earth and some other planets continue to expel (and perhaps recycle) gases from their interiors to their surface as temperatures are sustained by long-lived heat sources.

2.24 Atmospheres

Figure 2.24 The development of the two basic types of planetary atmospheres is shown in these diagrams. A) Inherited atmospheres formed as nebular gases were gravitationally trapped around a growing planetary nucleus of icy materials (left). Eventually the gases became gravitationally unstable and collapsed to form a thick atmosphere around the core (right). (B) Secondary atmospheres developed on some of the terrestrial planets that retained no nebular gases (left). Magmatism and thermal metamorphism released gases to the surface of the planet where they accumulated (and eventually escaped into space if the planet was small). Secondary atmospheres are thin and compositionally different than inherited atmospheres (right). Activity in the biosphere and lithosphere may continue to modify the composition of the atmosphere.

Just as there are sources for planetary volatiles, there are also sinks or traps where these atmophile elements may be temporarily or permanently removed from their respective atmospheres (Figure 2.25). The light elements, particularly hydrogen and helium, are subject to permanent loss to space as they gradually leak out of the top of most planetary atmospheres. One reason for this is that a small proportion of all atmospheric atoms have velocities that exceed a particular planet's mass-dependent escape velocity. For example, hydrogen (as H2) has an average lifetime on Earth of only 1 million years. The process of hydrogen loss is aided by the decomposition of water vapor and other hydrogen-bearing compounds by the action of sunlight in the upper atmospheres of the planets.More dramatic losses of atmospheric gases may have occurred when a hot magma ocean warmed and accelerated the overlying gas to escape velocity. This wind or blowoff may have dissipated a lot of gas around the primordial Earth. In addition, the impact of large meteorites may also eject atmospheric gases from a planet.

2.25 Atmospheric sources and sinks

Figure 2.25 Temporary or even permanent sinks where volatiles may be trapped are abundant on the planets and include atmospheres, lakes, oceans, rivers, ice caps and glaciers, reactions with other materials to form rocks (carbonates, red beds), and trapping in soils.

Other types of volatile traps involve precipitation of gases into solids or liquids, which accumulate on the surfaces of planets. For example, not all of the water vapor released to Earth's atmosphere is now present there. Some water is tied up as ice, some is in rocks, and some has decomposed, allowing hydrogen to escape to space. Carbonate minerals like calcite also form at the surfaces of some planets, removing carbon dioxide and locking it away to form solid rocks. This type of volatile trap may be particularly important for Earth, where thick layers of carbonate rock have formed in its oceans. The atmosphere of Mars, also, may have lost carbon dioxide in this way.

As temperatures declined in the early terrestrial atmosphere, water condensed and showered onto the surface of the planet and accumulated in craters and basins. This condensed portion of a planet's fluid envelope is called its hydrosphere. If retained in a liquid state, it can substantially modify a planet's surface by eroding some features away and redepositing the transported material elsewhere to form new sedimentary rock bodies. The solid icy shells of the outer satellites (and the polar caps of Earth and Mars) can be considered as volatile traps, which are similar in origin to planetary hydrospheres. In fact, on some planetary bodies liquid water may be sandwiched beneath a solid outer rind in a fashion similar to that which produces asthenospheres of partially molten silicates on the larger inner planets.

Water was not the only icy substance to have become incorporated in the planets. Nitrogen (N2) and methane (CH4) should also be important in bodies formed in the cool outer solar system. It has been postulated that Titan, the largest satellite of Saturn, contains a nitrogen-rich atmosphere with methane clouds from which rain or snow falls onto the surface, producing what may be the solar system's only methane seas. In short, there are important sinks for many volatile substances---including water, methane, carbon dioxide, sulfur, and oxygen. We will return to these later; for now, it is important to realize that even though gases are released by differentiation to an atmosphere, they may escape completely---by loss into space or by incorporation into a variety of surface materials---thereby considerably reducing the mass of the atmosphere. Consequently, the density and composition of atmospheres may change radically as the result of climate change or even as a result of biologic processes. We note finally that the oxygen-rich atmosphere of Earth appears to be the result of the accumulation of oxygen produced through billions of years of photosynthesis by living organisms in Earth's distinctive biosphere.

Just as the interiors of planets evolve by differentiation, so do the atmospheres of planets, reflecting the unique compositions, sizes, and physical settings of each planet.


The planets and the Sun are made of elements, which, other than hydrogen and helium, were formed by nuclear processes in earlier generations of stars. Accumulations of gas and dust, called nebulas, occasionally form within interstellar space. Gravitational contraction of part a nebula may have led to the formation of our solar system. As the nebula collapsed by gravitational contraction, the proto-Sun formed at its center; because of rotation the nebula maintained a thin outer disk of gas and dust. New grains of dust condensed from the gas and settled toward the central plane of the nebular disk. As a result of a strong temperature gradient, the compositions of the grains were controlled by their distance from the still-forming Sun. Metals and silicates were important condensates in the inner solar system and ices of water, ammonia, and methane were important in the outer solar system. Eventually, the orbiting dust accreted by collision resulting in a system of planetary bodies whose size and composition were determined by their distance from the center of the nebula. (A T-Tauri phase of the early Sun may have dispersed the gaseous part of the nebula at this time or a bit earlier.) The accretionary heating of the planets caused internal differentiation, which produced their layered internal structures. Crusts and rigid lithospheres composed of rocks with widely varying compositions formed the outer solid shells of most planets. Atmospheres composed of gaseous elements were trapped from the nebula by the gravity of the giant outer planets. Smaller planets developed atmospheres through outgassing, by which volatile elements were released from their interiors during accretion and differentiation. The geologic histories of the planets are reflected by features preserved on their surfaces and in their atmospheres and are largely the result of their thermal evolution.

Review Questions

1. How and where were the elements formed?

2. Is there any reason to believe that the elements that Earth and other planets are made of once resided in stars?

3. Describe the solar nebula from which the planets are thought to have formed.

4. Contrast the behavior of volatile and refractory materials. Give an example of each.

5. When did the solar system originate? How can we date an event that happened so long ago?

6. Why did planets with compositions as different as those of Earth and Jupiter form from a solar nebula that may have had a uniform chemical composition?

7. What prevented Mars from obtaining an ice-rich outer shell encasing a rocky interior as the temperature of the nebula continued to drop?

8. Contrast the origin and evolution of an inherited atmosphere with that of a secondary atmosphere and give real examples of each type.

9. Why doesn't Earth have a thick atmosphere of hydrogen and helium if these are the two most abundant elements in the universe and in the solar nebula from which Earth is presumed to have formed?

10. How is the nature of a planet's interior reflected in the appearance of its surface?

11. Can solid rock flow?

12. Describe the differences across the crust-mantle boundary in a planet like Earth and contrast those differences with those found at the asthenosphere-lithosphere boundary.

13. Explain why the rocks found at the surface of a planet, say the Moon, are so different in elemental composition from the meteoritic material from which the planet formed.

14. Are all heavy elements concentrated in the cores of differentiated planets? (Hint: Consider uranium, the heaviest natural element.)

15. How do we know that the number of meteorites striking Earth in a year has varied tremendously over the last 4.5 billion years? Is this consistent with our ideas about how the planets formed?

16. Explain the important characteristics of igneous, sedimentary, and metamorphic rocks.

17. Why does magma tend to rise upward toward the surface of a planet?

18. Compare the surfaces of two hypothetical planetary bodies with diameters of 5000 km; the first is found in the inner solar system, the second is in the outer solar system.

19. What are the most important controls on the thermal (heating or cooling) history of a planet?

Important Terms


accretionary heating

ash-flow shield






carbonaceous chondrite


cinder cone

collisional accretion


composite volcano










escape velocity







gravitational collapse




inherited atmosphere



lava dome




magma chamber

magma ocean











planetary nebula


plate tectonics


primary atmosphere






radiogenic heat

radiometric age




secondary atmosphere

sedimentary rock

shield volcano

shock metamorphism



solar energy



T-Tauri phase

tidal heat

thermal energy

thermal history




Additional Reading

Allegre, C. H. 1992. From Stone to Star. Cambridge, MA: Harvard University Press.

Berman, L. and J. C. Evans. 1986. Exploring the Cosmos. Boston: Little, Brown and Company.

Beatty, J. K., C.C. Peterson, and A. Chaiken (eds). 1999. The New Solar System, 4th edition. London: Cambridge University Press. (Chapters by J. A. Wood, W. K. Hartmann, J. S. Lewis, and J. B. Pollack cover material discussed in this section.)

Hartmann, W. K. 2005. Moons and Planets, 5th edition. Belmont, CA: Wadsworth Publishing Company.

Hamblin, W.K., and E.H. Christiansen. 2003. Earth's Dynamic Systems, 10th edition. Upper Saddle River, New Jersey: Prentice Hall.

Morrison, D., Owen, T. C. 1987. The Planetary System. Reading, MA: Addison-Wesley Publishing Company.

Taylor, S. R. 1992, Solar System Evolution: A New Perspective. Cambridge, England: Cambridge University Press.

Wood, J. A. 1979. The Solar System. Englewood Cliffs, NJ: Prentice-Hall, Inc.

Zeilik, M., and J. Gaustad. 1983. Astronomy: The Cosmic Perspective. New York: Harper and Row.

Chapter 3: Meteorites and Asteroids

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