Chapter 4: The Geology of the Moon

Introduction

In many ways the Moon is a geologic Rosetta stone: an airless, waterless body untouched by erosion, containing clues to events that occurred in the early years of the solar system, which have revealed some of the details regarding its origin and providing new insight about the evolution of Earth. Although they also posed new questions, the thousands of satellite photographs brought back from the Moon have permitted us to map its surface with greater accuracy than Earth could be mapped a few decades ago. We now have over 380 kg of rocks from nine places on the Moon, rocks that have been analyzed by hundreds of scientists from many different countries. Data from a variety of experiments have revealed much about the Moon's deep interior. As it turns out, the Moon is truly a whole new world, with rocks and surface features that provide a record of events that occurred during the first billion years of the solar system. This record is not preserved on Earth because all rocks formed during the first 800 million years of Earth's history were recycled back into the interior. The importance of the Moon in studying the principles of geology is that it provides an insight into the basic mechanics of planetary evolution and events that occurred early in the solar system. Much of the knowledge we have of how planets are born and of the events that transpired during the early part of their histories has been gained from studies of the Moon.

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 surface of the Moon can be divided into two major regions: (a) the relatively low, smooth, dark areas called maria (seas) and (b) the densely cratered, rugged highlands, originally called terrae (land).

2. Most of the craters of the Moon resulted from the impact of meteorites, a process fundamental in planetary development.

3. The geologic time scale for the Moon has been established using the principles of superposition and cross-cutting relations. Radiometric dating of rocks returned from the Moon has provided an absolute time scale.

4. The lunar maria are vast plains of basaltic lava, extruded about 4.0 to 2.5 billion years ago. Other volcanic features on the Moon include sinuous rilles and low shield volcanoes.

5. The major tectonic features on the Moon, mare ridges and linear rilles, are products of minor vertical movements.

6. Lunar rocks are of igneous and impact origin. The major types include: (a) anorthosite, (b) basalt, (c) breccia, and (d) glass.

7. The Moon is a differentiated planetary body with a crust about 70 km thick. The lithosphere is approximately 1000 km thick. The deeper interior may consist of a partially molten asthenosphere and a small metallic core.

8. The tectonic and thermal evolution of the Moon was very rapid and terminated more than two billion years ago. The Moon has no surface fluids, so that little surface modification has occurred since the termination of its tectonic activity.

9. The major events in the Moon's history were: (a) accretion of material ejected from Earth after a massive collision with a Mars-sized object, (b) and differentiation with the formation of the lunar crust by crystallization of a magma ocean, (b) intense meteoritic bombardment, (c) extrusion of the mare lavas, and (d) light bombardment.

The Moon as a "Planet"

In July 1969, a human stood for the first time on the surface of another planet, seeing landscape features that were truly alien and returning with a priceless burden of Moon rocks and other information obtainable in no other way. Nonetheless, many of the facts listed in Table 1 were known long before we began to explore space; they represent years of diligent study. For example, it was discovered centuries ago that the Moon revolves about Earth and not the Sun and is thus a natural satellite (the largest in the inner solar system). Long ago the distance from Earth to the Moon was measured and the diameter of the Moon determined. Early astronomers realized that the Moon's rotation period and its period of revolution are the same; thus it keeps one hemisphere facing Earth at all times. Moreover, many of the Moon's surface features have become well known, especially since the days of Galileo, the first to study the Moon through a telescope. Even the density and gravitational field of the Moon had been determined long before our generation. But not until the 1960s---and the inception of space travel with its sophisticated satellites and probes and the eventual Moon landing---did man begin to appreciate the significance of the Moon as a planet. In spite of its small size and forbidding surface, the Moon has revealed secrets that pertain to the ultimate creation of our planet, Earth, and our neighbors beyond.

Moon near side

Figure 4.1

The Moon is geologically different from Earth. The most striking differences are apparent in this photograph, taken through a telescope from Earth. The absence of swirling clouds, oceans, and an atmosphere reveal the water-poor nature of the Moon. Also seen in this photograph are the fundamental differences between the dark lava-covered maria and the lighter highlands, or terrae, which are intensely cratered. This is the familiar hemisphere of the Moon, for the Moon always has the same face turned toward Earth.

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Major Geologic Provinces

When Galileo first observed the Moon through a telescope, he discovered that its dark areas are fairly smooth and its bright areas are rugged and densely pockmarked with craters. He called the dark areas maria, the Latin word for seas, and the bright areas terrae (lands). These terms are still used today, although we know the maria are not seas of water and the terrae are not geologically similar to Earth's continents. The maria and terrae do, however, represent major provinces of the lunar surface, each with different structures, landforms, compositions, and histories.

The maria and terrae can even be distinguished from Earth by the naked eye. As shown in Figure 4.1, the maria on the near side of the Moon appear to be dark and smooth, with only a few large craters. Some maria occur within the walls of large circular basins such as Crisium, Serenitatis, and Imbrium, whereas others such as Oceanus Procellarum occupy much larger, irregular depressions. We know from lunar rock specimens and surface features that the maria are vast layers of thin basaltic lava, which flowed into depressions and flooded large parts of the lunar surface.

Moon farside

Figure 4.2 The far side of the Moon, photographed by orbiting satellites, reveals that most of the Moon's surface consists of the heavily cratered highlands. This part of the eastern limb of the Moon was never seen until the space age. The irregular, dark, smooth area in the lower left is Mare Marginis with Mare Smythii below it; the rest of this vast region lacks large accumulations of lava. Why there are so few maria on the far side of the Moon remains a topic of discussion.

 

The terrae, or highlands, constitute about two-thirds of the near side of the Moon and exhibit a wide range of topographic relief. This is the highest and most rugged topography on the Moon, where local relief in many areas is up to 5000 meters. An important characteristic of the lunar highlands is that they contain abundant craters, many of which range from 50 to 1000 km in diameter. For example, craters larger than 10 km are about 50 times more abundant on the highlands than on the maria.

The far side of the Moon was totally unknown until photographs were first taken by a Soviet spacecraft in 1959. It was a total surprise to learn that, although the details were poorly defined, the far side of the Moon was composed almost entirely of densely cratered highlands. Later, orbiting satellites launched by the United States completely photographed the far side of the Moon with definition sharp enough to map the surface in considerable detail (Figure 4.2). These photographs confirmed that the far side of the Moon is densely cratered; only a few large craters contain mare lavas. Why the maria are largely restricted to the near side of the Moon remains a fundamental problem of lunar geology.

Results from several Apollo experiments demonstrate other fundamental differences between the highlands and the maria. Remote sensing measurements of the composition of the lunar surface indicate that the maria and highlands are composed of distinctly different rock types. Rocks collected by Apollo astronauts show the highlands to be mostly composed of anorthosite and other feldspar-rich rocks, in contrast to the basaltic maria. Tracking Moon-orbiting spacecraft from Earth has confirmed the elevation differences. The highlands may be as much as 5 km above the mean radius of the Moon, whereas the maria lie almost 5 km below the mean radius. (As the Moon has no water, its mean radius can be used as a reference level from which relative elevations or depressions can be measured). Studies of the velocities of seismic waves have shown that the highland crust is also much thicker, in some cases up to 100 km thick, while the crust under some of the near side maria is only around 40 km thick.

The maria and highlands not only represent different types of terrain, they broadly represent two different periods in the history of the Moon. The highlands, which occupy about 80 percent of the entire lunar surface, are composed of rocks that formed very early in the Moon's history. The entire outer portion of the Moon is thought to have been molten at the time the highlands crust began to form 4.6 billion years ago. As light silicate minerals accumulated at the top of this "magma ocean," they formed a crust, which soon became densely cratered by a contemporaneous intense bombardment of meteorites. This initial high cratering rate declined rapidly but the Moon's surface, one of the oldest in the solar system, became covered with craters.

The maria were formed by the extrusion of vast amounts of lava that accumulated in the lowlands of large craters or basins and, in places, overflowed and spread over parts of the lunar highlands. The maria are thus relatively young features of the lunar surface, even though they began to form four billion years ago.

Impact Craters

Impact cratering is a rare event on Earth today, but it has been a fundamental process in planetary development. The Moon is pockmarked with literally billions of craters, which range in size from microscopic pits on the surfaces of rock specimens to huge, circular basins hundreds of kilometers in diameter. The same is true for the surfaces of Mercury, Mars, the asteroids, and most of the satellites of the outer planets as well. Indeed, impact cratering was undoubtedly the dominant geologic process on Earth during the early stages of its evolution. Earth was once heavily cratered like the Moon today. Evaluation of the impact process can provide an important interpretive tool for understanding planets and their development.

The Mechanism of Crater Formation

Although it is difficult to imagine the magnitude of these enormous impact events, which excavated millions of cubic meters of material in seconds, the results of the photographic missions to the planets and studies of experimental craters produced in laboratories have greatly increased our understanding of what happens when a meteorite strikes a planet.

The process of crater formation, like many geologic processes, can be viewed as a transfer of energy of one form to another form. In this case, the kinetic energy of the meteor, which may be moving at velocities of 5 to 50 km second (5 km sec is 18,000 km hr), is imparted to the planetary surface and changed into other types of energy. It has been calculated that a meteorite with a diameter of only 20 m could excavate a crater 500 m across and would be as energetic as the largest volcanic explosions that have occurred on Earth during historic times. However, not all of the energy of an impact goes into excavating a crater. Much of it is converted to heat. If the meteorite is large enough, a small volume of rock may be completely melted and transformed into an igneous impact melt that may come to rest on the floor of the crater or be ejected over its rim during the seconds involved in crater formation. Some materials vaporize completely and other rocks recrystallize as metamorphic rocks. Some of the energy is used in fracturing and pulverizing the target materials and the meteorite itself. Some energy is dispersed as sound or seismic (earthquake) waves.

Crater formation

Figure 4.3 Stages in the formation of a hypothetical meteorite impact crater are depicted in this series of diagrams. During the compression stage, the kinetic energy of the meteorite is almost instantly transferred to the ground as a shock wave, which moves outward, compressing the rock. At the point of impact, the rock is intensely fractured, melted, and partly vaporized by shock metamorphism. After the shock wave passes, the rocks return to normal pressures and expand explosively as a "reflected" rarefaction wave, which throws out large amounts of fragmental debris as ejecta---the excavation stage. Near the edge of the growing crater the solid bedrock is forced upward to form an elevated rim; in some cases an overturned flap may develop. The ejecta are propelled along ballistic pathways away from the crater. Ejecta that fall far from the crater have high velocities and may form chains of secondary craters, which radiate away from the point of impact. These secondary impacts may result in extensive mixing of primary crater ejecta with local surface materials. A large amount of fragmental material also falls back into the crater. During a subsequent modification stage, material may slump or slide off the steep crater walls to form terraces. It is common for some type of uplift to form on the floors of larger craters as well.

The formation of a crater is generally divided into three stages: compression, excavation, and modification (Figure 4.3). During the compression stage, the meteor's kinetic energy is transferred through the ground by a shock wave that expands in a spherical pattern away from the point of impact. The pressures experienced by the target materials are so high that the rock behaves like a fluid for a short time. During this stage, only a small amount of material is removed by vapor jets around the side of the meteorite; the main mass of the material remains to be excavated.

After the passage of the shock wave, the rocks relax and a rarefaction wave is established that travels upward and allows the rocks to expand explosively as the target attempts to return to the same low pressure as its surroundings. The rarefaction wave is in effect a reflection of the compressional wave, because it forces material to flow upward toward the ground surface. The result is similar to an explosion located a few meteorite diameters below the ground. This causes the rock to be lifted upward and bent backward, and material to be ejected from the surface and thrown out along ballistic trajectories. During this excavation stage, the crater grows rapidly, attaining its final diameter even before material thrown from the impact site reaches the ground. The fragmented rock that is blasted away is called ejecta. The path of each particle is affected by the strength of the planet's gravitational field. Initially, the ejecta forms a nearly continuous cone-shaped curtain that appears to enlarge at its base. This ejecta curtain quickly tears apart into filamentous elements, which fall to the ground forming the rays characteristic of young craters. Near the edge of the crater, layers of rock may be overturned in a flap that forms part of the high crater rim. The excavation is accomplished very rapidly and the duration of the event depends mainly on the size of the meteorite and its velocity. For example, Earth's Meteor Crater (Figure 4.4) is thought to have formed in about 10 seconds.

Meteor Crater

Figure 4.4

Meteor Crater, Arizona, is a recent impact crater similar to those found on the Moon. It is one of the youngest impact craters found on Earth. This small crater formed about 25,000 years ago when an iron meteorite struck Earth. The bowl-shaped crater is a little over 1 km in diameter and is 200 m deep. Its floor is covered by younger lake sediments. An extensive blanket of ejecta surrounds the crater. The polygonal outline of the crater is probably the result of a fracture system that predates the excavation of the crater.

 

During the final stage, after much ejecta has settled to the surface, the crater becomes modified from its initial, excavated form. The floor of the crater may rapidly rebound upward to compensate for the material removed during excavation. Slices of rock from the steep walls may immediately slip back into the crater, forming a series of terraces that partially fill the depression. Terraces form a nearly concentric natural staircase to the crater floor. The modification stage probably continues over a long period as the crust reestablishes a stable configuration in the planet's gravity field (Figure 4.5).

Crater

Figure 4.5

The internal structure of a hypothetical impact crater is shown in this cross section. Note particularly the raised rim of the crater, the lens of impact breccia (and impact melt) that floors the crater, the ejecta blanket, which adds to the height of the crater rim, and the deformation of preexisting rocks. Rocks in the target are both folded and extensively fractured. Terraces form as slices of rock drop along steep faults into the interior of the crater.

 

Features of Impact Craters

Impact craters are typically circular, in contrast to many volcanic craters, which are frequently asymmetrical or elongate. When an impacting body strikes the surface at an angle greater than 15, the crater will nonetheless be circular, because the shock waves spread out with equal velocity in all directions from the point of impact and the rarefaction waves move back toward that point.

As can be seen in Figure 4.5, the rim of the crater is built of deformed bedrock that has been heaved upward and outward, bent back, and even overturned. The material thrown out of the crater accumulates around the crater rim as an ejecta blanket. This is thickest near the crater and thins outward, becoming discontinuous and patchy at a distance of about 1.5 crater diameters from the rim. Larger blocks thrown out from the main crater may impact with enough speed to form secondary craters, which tend to be irregular in shape and are typically grouped in clusters or chains. Many overlap and form distinct linear patterns. Fine powdered material and melted material, which resolidifies into glass beads, is thrown farther and accumulates as a system of long, bright, splashlike rays.

The floor of an impact crater is commonly covered with a lens-shaped deposit of fragmented rock (breccia) and small amounts of lava produced by shock melting during impact. The bedrock below is highly fractured down to a depth equal to about three times the depth of the crater.

It is important to emphasize that the impacts of meteorites produce landforms (craters) and new rock bodies (ejecta blankets), and are therefore similar to other rock-forming processes (such as volcanism and sedimentation) that operate at the surface of a planet. In each process, energy is transferred, material is transported and deposited to form a new rock body, and a new landform (crater, volcanic cone or flow, or delta) is created. The rock-forming processes associated with impact include: fragmentation, ballistic transportation, and deposition of rock particles, in addition to rock modifying processes of shock metamorphism, partial melting, and vaporization.

Types of Impact Craters

The surface of the Moon (Figure 4.6) has been described as being covered by a "forest of craters," and at first glance all craters may look the same. After a few moments of thoughtful study, however, one is able to identify various "trees in the forest" and recognize various types of craters by certain characteristics of their size and form. Detailed study of lunar craters shows that many morphologic differences are associated with size, so it is useful to classify craters into groups that are similar in size and shape and that probably originated from the impact of similar sizes of meteorites.

Shaded relief map

Figure 4.6

A shaded relief map of the Moon shows the topography and major provinces on the Moon.

 

Craters less than 20 km in diameter (Bowl-Shaped Craters). Small lunar craters, generally less than about 20 km in diameter, are almost perfectly circular and are typically bowl-shaped (Figure 4.7).

Shaded relief map

Figure 4.7

Lunar craters smaller than 20 km in diameter are almost perfectly circular and are typically bowl-shaped. In this example from southern Mare Serenitatis, the raised crater rim is well defined and is surrounded by a prominent ejecta blanket composed of light-colored material. The rays formed by secondary cratering extend many crater diameters beyond the rim of the crater.

 

The deepest measure about 2 km from the crest of their rims to their floors. Such a crater usually has a well-defined, raised rim created by ejected material and by the expansion of the materials in a ring very close to the crater rim. The ejecta blanket extends from the rim to a distance approximately equal to the diameter of the crater. The surface of the ejecta blanket consists of a series of hummocky ridges that superficially resemble sand dunes. The crater in Figure 4.8 has some large boulders near its rim that were thrown from the crater as it formed.

Fresh crater

Figure 4.8

Most craters have blocky rims like this crater (about 14 km across) from the highlands of the lunar far side. Long tongues of lava-like material also extend downslope from the crater rim. These flows may consist of fragmented rock debris, but it is perhaps more likely that they are impact-produced melts ejected at low velocities during the late stages of the formation of the impact crater.

 

Another type of fairly small crater possesses a flat floor and is not a simple bowl-shaped hole. The floor of the crater accounts for about 50 percent of the crater diameter. The flat floor may be caused in part by the collection of materials that roll or slump off the walls of the crater and partially fill it. Some prominent flat-floored craters are the result of excavation through distinct layers in the surface materials of a planet (Figure 4.9). The response of materials to the passage of the shock wave depends on their physical properties; if a sharp discontinuity exists in a layered planet, the wave may be reflected by it, creating the flat floors of these small craters. There is substantial evidence that the Moon and several of the other planets are covered by a layer of loose, fragmental material called regolith to a depth of many meters. The regolith of the Moon was probably produced by billions of years of meteoritic bombardment, during which the surface material was constantly fragmented, churned, and mixed to produce a brecciated soil. The depth at which the regolith gives way to more solid rock certainly divides two materials with different strengths and other physical characteristics. It seems quite plausible that the lower, less brecciated bedrock would be much stronger than the regolith, hence would not be as easily ejected, creating a sharp break in the slope of the crater wall and a flat floor. Apparently, some fresh bowl-shaped craters were not produced by events sufficiently energetic to encounter this discontinuity.

Shaded relief map

Figure 4.9

A variety of small generally bowl-shaped craters are present in this view of the lunar highlands. The small crater near the center of the picture is about 1 km across. Its small, flat floor and blocky rim contrast sharply with the older and more degraded craters that dominate the surface. Nonetheless, this crater lacks rays and is probably older than the craters shown in Figures 4.7 and 4.8.

 

Nonetheless, virtually all craters over 30 km in diameter have flat floors and many smaller ones have a distinctive swirl-like texture on their floors that results from the slumping of material from the crater wall in thin sheets toward the floor (Figure 4.10).

Shaded relief map

Figure 4.10

The accumulation of material on the floor of a crater may result from slumping from its walls commonly forming distinctive surface features. The scarp of the right side of the crater is especially high and creates the crater's irregular outline. The ejecta deposit around the crater rim contains dune-like features formed during its deposition..

 

Craters 20 to 200 km in diameter (Terraced Craters with Central Peaks). Impact craters on the Moon with diameters of 20 to 200 km typically have terraces on their inner walls that become well developed with increasing size. The terraces form as slump blocks, when thin slices of rock become unstable on the crater walls and collapse into the depression. All of these craters have more or less flat floors, but some have been arched up, presumably as a result of adjustment during the modification stage. Central peaks that rise abruptly from the floor form prominent features of many terraced craters (Figure 4.11). Almost all lunar craters larger than 50 km in diameter have central peaks.

King crater

Figure 4.11

The lunar crater King illustrates many of the features of impact craters larger than 20 km in diameter. King, about 75 km in diameter and about 4 km deep, is on the far side of the Moon, west of Mendeleev. Among its interesting features are its atypical horseshoe-shaped central peak. Central peaks are common in craters of this size. King crater also has terraced walls that descend to a relatively broad, flat floor. A smooth puddle of solidified impact melt lies above the peak; another large accumulation of impact melt lies outside the crater on its north rim. These pools were fed by lobes of melt that drained into local depressions.

 

Although there are several theories to explain the formation of central peaks, the most likely explanation relates them to a type of elastic rebound. During excavation, material at depth is forcefully pushed toward the center of the crater, then upward. Material pathways mimic this as a result of relaxation after passage of a hemispherical shockwave (Figure 4.3).

Craters 200 to 300 km in diameter (Peak-Ring Basins). Large craters are transitional in morphology to still larger, multiring basins. The terraced walls are retained, but the central peaks change progressively from single promontories, through clusters of peaks, to a ring of peaks halfway between the center and the rim of the crater (Figure 4.12). These impact features are called peak-ring basins.

Shaded relief map

Figure 4.12

Schroedinger Basin lies in the south polar region of the Moon's far side. It is an excellent example of a peak-ring basin, typified by a ring of peaks on its floor in the place of a central peak. The long, narrow valley that extends away from the basin's rim is a chain of closely spaced secondary craters. Schroedinger is 320 km in diameter. (The horizontal lines on the picture are artifacts of the way in which the image was constructed and are characteristic of data collected by the Lunar Orbiter spacecraft.)

 

Basins larger than 300 km in diameter (Multiring Basins). The largest impact features on the planets, called multiring basins, are fringed by a series of concentric ridges and depressions. The youngest and best-preserved multiring basin on the Moon is Orientale (Figure 4.13).

Orientale

Figure 4.13

Impact craters larger than 300 km in diameter take on the appearance of giant bull's-eyes such as that displayed by Orientale Basin. It is the youngest multiring basin on the Moon. The diameter of the outer ring is 900 km. Two well-defined rings lie within the outer ring. The deposition of ejecta has tremendously modified the original nature of the surrounding terrain; radial textures are prominent. The location of the original rim of the impact excavation is controversial. Some argue that the outer ring is a fault scarp that marks the outer edge of a huge terrace; others suggest that the outer ring is itself the rim of the crater. A relatively small patch of much younger mare lava lies within the inner ring.

 

Although it is largely hidden from telescopic view, satellites orbiting the Moon have photographed it in detail. The basin resembles a gigantic bull's-eye with three concentric ridges or scarps and intervening lowlands. The diameter of the outer ring of mountains (the Cordillera) is 900 km; the entire state of Colorado could fit within the second ring. The spacing of the rings increases outward, a feature that is common to many multiring basins. Beyond the outer ring, the most prominent topographic features are low ridges oriented radially to the basin (Figure 4.14).
SE Orientale

Figure 4.14

The southeastern edge of the Orientale Basin is marked by a scarp that in places rises 4 km above the adjacent plain. The radially textured ejecta blanket lies to the right of the scarp; to the left is a terrain marked by numerous small domes. In areas where the scarp is low the radial texture of the ejecta crosses the scarp, suggesting that the outer ring is a terrace blanketed by ejecta rather than the crater's original rim. The lack of evidence of debris piled up along the scarp suggests that it formed after the outward movement of ejecta.

 

A steep scarp, 4 km high in places, forms the outer ring. The texture of the surface just within the outermost ring is distinctly different, consisting mainly of small domes or hummocks. Another steep cliff forms the second ring, but the inner ring is not a well-defined scarp; rather, it consists of a group of peaks that surround the plains in the middle of the basin (Figure 4.15). This ring of peaks is probably very similar to the peak rings formed in smaller craters. The central part of the crater is very dark and smooth and is covered by volcanic lava flows that filled the depression after its formation.

Orientale inner rings

Figure 4.15

The two inner rings of Orientale Basin are much less distinct than the outer ring. Locally the middle ring is bounded by a low scarp, but in general it consists of an annulus of large mounds or peaks that are morphologically similar to those in smaller peak-ring basins. The innermost ring is composed of larger domes. The fractured plains that lie to the north may consist of impact-melt-rich rocks, while the dark, smooth plains in the upper part of the photo are mare basalts.

 

The multiring structure of Nectaris Basin, located near the eastern margin of the near side of the Moon just south of the equator, is somewhat more subtle (Figure 4.1). A definite series of arcuate ridges and intervening lowlands can be seen. This is obviously a much older basin than Orientale because the concentric rings and ejecta blanket are considerably modified by subsequent impact craters. Like Orientale, only the central part of Nectaris is flooded with basalt.

Another basin that is useful in studying multiring structures is the Imbrium Basin, which dominates the near side of the Moon (Figure 4.1). It is almost completely flooded with younger lavas, but the concentric structure can still be identified. The innermost ring, exposed only as chains of islands, is 675 km in diameter. The second ring is exposed mainly as Montes Alpes and the rugged terrain near Archimedes. The third ring is the largest and most conspicuous and corresponds to Montes Carpatus, Montes Apenninus, and Montes Caucasus. This ring marks the beginning of the ejecta blanket characterized by rugged linear ridges best developed at the southeast margin of the basin. Suggestions of still another ring, part of which bounds Sinus Medii on the south, can be seen at a considerable distance southeast of Montes Apenninus.

Megaterrace models

Figure 4.16

Two models for the formation of multiring basins are contrasted in these vertical cross sections of the outer part of the Moon. (A) Megaterraces, the areas between the outer rings of a basin such as Orientale, are interpreted as large fault-bounded terraces, formed during the modification stage of a basin's evolution. The inner rings are uplifts similar to those in peak-ring basins. (B) Nested-crater model explains the multiring structure of large basins as structural discontinuities in the target. Rebound of the deep crater lifts these discontinuities to the surface, where they are expressed as rings. The crater of excavation is much deeper than that inferred for the megaterrace model. Terraces are considered to be minor modifications.

 

There are several ideas about how the rings are formed. Two of the more popular theories are illustrated in Figure 4.16. The megaterrace model for multiring basin formation postulates that the outer ring(s) of a multiring basin are the margins of huge terraces that slumped downward along steep faults. One of the inner rings marks the rim crest and the innermost ring is formed by the same process that forms the peak rings in smaller craters. According to this model, the second or middle ring of Orientale is a remnant of the rim of the crater of excavation. The nested-crater model proposes that the outer ring of the crater is the actual rim crest or limit of excavation for the crater. The inner rings of the basins are produced as shock waves encounter changes in the layered structure of the planet's surface in much the same way as was discussed earlier for the production of small flat-floored craters. Readjustments after excavation allow the deep crater to rebound to the surface, creating the final subdued topography of the basins. Minor slumping of material off the rings may further modify the final configuration.

It is obvious that the nested-crater model calls for a much deeper original crater and for the excavation of larger amounts of material. It is possible, if this is the way multiring basins form, that chunks of the Moon's mantle have been excavated and strewn across the surface. However, no material similar to the mantle has been unequivocally identified from the samples brought back to Earth by the astronauts. In the megaterrace model, the formation of the outer ring is seen as a part of the modification stage of crater formation, whereas in the nested-crater model, it is principally a result of the excavation of the crater.

The ejecta deposits from such large impacts differ considerably from the hummocky rim deposits and rays of smaller craters. Surrounding the Moon's Imbrium Basin, (as well as many others) is a distinctive ridged and furrowed terrain (Figure 4.17). This sculpted pattern is believed to result from the "flowage" of the rapidly moving ejected debris and from pervasive secondary cratering. Even small craters have chains of secondary craters radiating away from them, produced as blocks of ejecta plow into the surface after being thrown from the primary crater. Apparently, the large basin-related events ejected vast amounts of material that eventually fell back to the planets' surfaces and created the furrowed texture.

Imbrium ejecta

Figure 4.17

The formation of a large multiring basin dramatically alters the nature of the landscape for many kilometers outside of the basin itself. The ridges and grooves seen in this photo lie 650 km from Imbrium Basin but were produced by massive secondary cratering related to the emplacement of Imbrium's ejecta. This terrain is called Imbrium sculpture, and similar terrains surround other large basins on the Moon and other planets. The smooth plains within the highly degraded craters may also have been produced by the emplacement of Imbrium ejecta. Herschel (40 km in diameter) is the fresh crater on the right side of the photo. Ptolemaeus is the large crater partly visible in the lower left of the photo. The dark plains to the north are the basaltic lavas of Sinus Medii.

 

The surface upon which these fragments fell became churned up and intimately mixed with ejecta. Figure 4.18 shows the lunar crater Ptolemaeus, located in the middle of the near side of the Moon, just south of the equator. It is partially filled by gently undulating, nearly smooth plains. Numerous pools of similar material fill depressions on the crater rim. Patches of light-colored plains like this were originally thought to be volcanic deposits. Careful inspection of the photos reveals that the topography beneath this sheetlike deposit has been preserved. Filling the craters with lava or ash flows would probably not preserve the old rims. Inspection of the rock samples returned from the Moon has shown the fragmental and brecciated nature of this type of deposit, and it is now believed that these smooth plains are part of the ejecta from a distant basin.

Ptolemaeus

Figure 4.18

The crater Ptolemaeus (150 km in diameter) is filled to about half of its original depth by a younger deposit with a gently undulating, nearly smooth surface. The outlines of several craters on the floor of the crater are preserved beneath this blanketing deposit. Originally, these light plains were thought to be volcanic deposits, but the study of samples returned by the Apollo astronauts has revealed that such plains are blankets of fragmental material ejected from large basins (in this case from Imbrium Basin).

 

 

Crater Degradation

The descriptions of craters in the previous section generally referred to their original shapes produced during the three stages of the impact process. Once a crater is formed, it may then be modified by a number of processes that gradually change the appearance of the crater until it may be totally unrecognizable or obliterated. The changes that occur are collectively called crater degradation.

Gagarin degradation

Figure 4.19

The effect of repeated meteorite impact on the lunar surface is to gradually modify or degrade the original shape of a crater. This fact is helpful in determining the relative ages of craters and of other lunar events. This view of the lunar highlands shows a portion of Gagarin Basin (265 km in diameter), the topographic rim of which is discernible along the right and lower margins of this photo. Innumerable impacts after the formation of Gagarin have removed all signs of its ejecta blanket and destroyed features on its floor. Both large and small craters have been important in this degradational process. Young craters, such as the bowl-shaped, sharp-rimmed crater in the center of the image, are surrounded by a halo of light-colored "fresh" ejecta and secondary craters.

 

 

The effects of crater degradation can be seen in Figure 4.19, a view of the lunar crater Gagarin. The rim of this large crater (265 km in diameter) is barely discernible in this heavily cratered area. Gagarin is shallow, the walls lack terraces and are irregular, and no ejecta patterns can be detected past its rim. Battering by countless impact events has produced these changes. It is obvious from this photo that the appearance of a crater changes with time. Older craters are usually more degraded than younger ones. For example, compare the sharp-rimmed craters that occur on the floor and rim of Gagarin (and are therefore younger) with the irregular, subtle nature of Gagarin itself. Indeed, the rims of the smaller craters superposed on Gagarin show a range from crisp, sharp, unmodified rims to ragged, degraded rims altered by subsequent impact. This concept is important because changes in the morphology of craters make it possible to determine the relative age of specific craters and unravel the sequence of some events in a planet's history.

Crater degradation may proceed in several ways. (1) Later impacts may partly or completely destroy the older crater. (2) The crater may be covered with ejecta from a younger crater. (3) The crater may be partially or completely buried by lava flows or sedimentary deposits. (4) Geologic activity of the atmosphere and hydrosphere of the planet (if it has one) may erode the crater. (5) The crater may undergo tectonic modification by faulting or folding. The sequence of photos in Figure 4.20 shows some of these effects on lunar craters.

Crater degradation Crater degradation Crater degradation Crater degradation

Figure 4.20

Crater degradation and modification can occur by means of slumping, isostatic adjustment, subsequent impact, and burial by ejecta or lava. Examples of the modification of craters are shown in these photographs. Studies of crater degradation and modification are important in reconstructing the sequence of events in lunar history. (A) Slumping of crater walls causes partial filling and the formation of terraces, and isostatic adjustment results in uplift of the crater floor. (B) Crater rims are partly obliterated by subsequent impact. (C) The ejecta from the larger crater at the top partly covers older craters. (D) Some craters are partly covered by lava flows.

 

 

Developing a Lunar Time Scale

The deposits of ejecta from craters, together with lava flows and other volcanic deposits, form a complex sequence of overlapping strata that cover most of the lunar surface. The individual deposits can be recognized by their distinctive topographic characteristics and by their physical properties---such as color, brightness, and thermal and electrical properties determined from measurements made with optical and radio telescopes.

In 1962, geologists from the U.S. Geological Survey developed a geologic time scale for the Moon so that major geologic events can be arranged in their proper chronology. The basic principles used to interpret lunar history are essentially the same as those used to study the history of terrestrial events, the most important of which are the laws of superposition and cross-cutting relations. These principles of determining relative ages are, of course, as valid on the Moon or any other planetary body as they are on Earth. In addition, other methods of determining the relative ages of lunar features were developed based on the abundance of craters (crater frequencies) and crater degradation.

The development of a lunar geologic time scale was a major advancement in the study of planetary geology. For the first time, the sequence of events in the history of another planet was firmly established.

Determining Relative Ages

Although it was recognized long ago that craters and other lunar surface features showed evidence of having been formed at different times, prior to the space program most observers studied features without relating them to their surroundings or their relative ages. Craters were classified according to their dimensions, statistics were calculated on crater frequency, but little effort was made to establish a sequence of events in lunar history.

Copernicus region

Figure 4.21

Relative ages of lunar features in the vicinity of the crater Copernicus are indicated by superposition of ejecta and mare basalts. Ejecta from Copernicus are superposed on all other features and are therefore the youngest materials. Ejecta from Eratosthenes (northeast of Copernicus) rest on the mare basalts and are younger than the mare material but older than Copernicus. Ejecta from the Imbrium Basin are covered partly by mare lavas and rest on the densely cratered highlands. Thus the basin and the lava flows are younger than the complex ejecta deposits of the highlands.

 

 

The first lunar chronology was developed in 1962 by Shoemaker and Hackman, who interpreted the sequence of events in the vicinity of the crater Copernicus (Figure 4.21). They recognized that ejecta from the craters Copernicus and Eratosthenes was superposed upon mare basalt whereas the basalt was superposed upon the craters of the lunar highland. From this kind of reasoning, they established a lunar geologic column. In many ways, what Shoemaker and Hackman did in providing a rationale for interpreting the Moon's history is comparable to what Smith, Lyell, and their contemporaries did in establishing the geologic time scale on Earth during the early 1800s. The planet and nomenclature are different, but the logic remains the same.

To understand the basis for establishing the lunar time scale and the meaning of the major events in lunar history, let us carefully consider the ejecta from the major craters studied by Shoemaker and Hackman and their reasons for recognizing the sequence of events these ejecta represent. As you read the following discussion, study the physiographic map of the Moon (Figure 4.6) and the indicated illustrations, for it is only by recognizing the physical relationships between the features discussed that an appreciation for the relative time involved in their formation can be gained. This discussion serves as an illustration of the method used to establish the sequence of events that shaped various planetary surfaces.

Copernicus. An outstanding feature of the near side of the Moon is the crater Copernicus with its spectacular system of bright rays that extend outward in all directions, in some cases for a distance of more than 300 km. Within the rays, predominantly near the crater, are elongate secondary craters. The rays and ejecta blanket surrounding Copernicus are superposed on essentially every feature in their path (Figure 4.21). During full Moon, when the rays are best observed, it is found that rays extend not only across Mare Procellarum and Mare Imbrium but also up the rim and across the floor of the crater Eratosthenes, 190 km to the northeast. From this superposition, it is clear that Copernicus and its associated system of ejecta and ray material are younger than the mare basalts and are younger than the rayless craters such as Eratosthenes. Other systems of rayed craters similar to Copernicus are centered on slightly smaller craters to the west, such as Kepler and Aristarchus. Their ejecta and ray material overlie everything they come in contact with. In the southern hemisphere, material from similar craters, such as Tycho, likewise overlaps all adjacent features. Although the rayed craters may vary in age, as a group they are younger than all other features on the Moon. The period of time during which rayed craters and their associated rim deposits were formed has been called the Copernican Period. Rocks and landforms of this age are the youngest on the Moon.

Eratosthenes. About half of the craters larger than 10 km in diameter that occur on the maria are rayed craters. Most other craters of this size range found on the maria are similar, but their ejecta blankets are dark, and they lack rays. Consider, for example, the crater Eratosthenes just northeast of Copernicus (Figure 4.21). It has terraced walls, a roughly circular floor, a central peak, a hummocky rim, and a distinctive pattern of secondary craters similar to Copernicus. However, unlike Copernicus, it does not have a visible ray system, and the secondary craters are noticeably more subdued than those around Copernicus. Eratosthenes and similar craters, together with their ejecta, are superposed on the maria and are, therefore, younger than the mare lava on which they are formed. However, they are older than the rayed craters, as shown by the fact that the ray material, secondary craters, and ejecta deposits of rayed craters are superposed on the dark rim of Eratosthenes. The period of time when deposits of these dark-rimmed craters formed is referred to as the Eratosthenian Period. Deposits of Copernican and Eratosthenian craters are easily recognized on the maria, but it is difficult, in some cases, to discriminate between Eratosthenian and older crater deposits on the lunar highlands where crater densities are much higher. In these instances, relative ages may be established using the principles of crater degradation discussed earlier.

Imbrium Basin. In the northwest part of the near side of the Moon is an enormous multiring basin, now largely filled with lava flows, called Imbrium Basin. We have seen in previous discussions that Imbrium Basin is the largest multiring basin on the Moon and, like other craters, was formed by impact. Imbrium Basin is surrounded by ejecta deposits similar to those formed by smaller craters---the best exposures being the Montes Apenninus, which extend outward from the southwest rim. The ejecta deposits are called by various names, depending on their topography and location. The most important of these is the Fra Mauro Formation, which can be traced as far as 400 km from the mountains surrounding the basin.

The ejecta from Imbrium Basin are partly covered with lava, as is most of the interior of the basin (Mare Imbrium). The time during which the lava and the ejecta deposits were formed is called the Imbrian Period. It is apparent that some impacts occurred after Imbrium Basin and its ejecta were formed, but before the extrusion of the lava flows. Some of these deposits around larger craters such as Archimedes, Plato, and Sinus Iridum are very extensive. It is calculated that about four times as many craters larger than 10 km in diameter were formed on the Moon during the Imbrian Period as during all the time since the last lava flows in Mare Imbrium were extruded.

The Ancient Terrae. The Imbrium ejecta partly overlie a complex sequence of craters and ejecta found in the lunar highlands. Ejecta from the upper part of this densely cratered terrain formed during what is termed the Nectarian Period, named after a large subdued basin on the eastern side of the Moon. The ejecta deposit from the Nectaris Basin can almost be traced to the far side of the Moon and serves as an important marker for this part of the Moon. The time before Nectaris Basin formed is called Pre-Nectarian and rocks of this age constitute the oldest materials on the Moon's surface. The geologic structure of these deposits is very complex. Large craters are closely spaced and modified by impact. Apparently the surface of the terrae was churned by repeated formation of large craters early in lunar history. Humorum and Serenitatis Basins and their ejecta deposits may all be of Pre-Nectarian age. Of course, the lavas that fill these depressions are much younger, mostly formed during the Imbrian Period. In many places in the highlands, Nectarian and Pre-Nectarian materials are not distinguished from one another and are called Pre-Imbrian deposits.

In light of these observations, geologists have mapped most of the Moon and outlined some of the major events in lunar history. These are summarized in the time chart (Figure 4.22).

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Figure 4.22

The geologic time scale of the Moon is shown in this diagram, along with some of the events which occurred during these periods.

 

 

Crater Frequencies

As we have seen in previous sections, cratering is the major geologic process that has operated on the Moon; the form and number of craters provide a wealth of information concerning the history of the planet and the relative ages of surface features. Cratering can be used to determine the relative age of a terrain unit or reference area because of the simple fact that a greater number of craters will have developed on older terrains than on younger ones. This relationship holds true regardless of whether the rate of cratering is constant, steadily decreasing, or erratic. Comparison of the number of craters in different regions enables estimates to be made of the relative ages of the cratered surfaces. A simple example illustrates this point. Assume that it has been snowing for several days and that the snow is 1 m deep on undisturbed lawns throughout the neighborhood. If the snow is 1 m deep on the sidewalks of a house, you can conclude that nobody has shoveled the snow since the storm began. If the snow is about half a meter deep in front of another house, it would be obvious that the walk had been shoveled, possibly midway through the storm. If another walk has only a few centimeters of snow, it would be obvious that it had been shoveled a short time before. A walk with only a few scattered snowflakes apparently would have just been shoveled.

crater frequency

Figure 4.6

Crater frequency can be used to determine relative ages of lunar surfaces, because an older surface has more craters than a younger one does. In the photographs, (A) has the fewest craters and is therefore the youngest surface. (B) and (C) are progressively older, and (D) is the oldest.

 

To understand the basis for using cratering to determine relative ages, one needs only to mentally substitute a planet's surface for the sidewalk and cratering projectiles for the snowflakes. The idea can be explained more directly with reference to Figures 4.23 and 4.24. On a hypothetical planet, a new surface subjected to bombardment will have a variety of crater sizes. Studies of the Moon and other terrestrial planets show that the number of craters is inversely proportional to their diameters. That is, the number of craters decreases dramatically as their size increases. This distribution of craters is shown in Figure 4.24. With time, the number of craters in each size range increases. There is a limit, however, because eventually further impact does not increase the number of craters, it simply makes new craters out of old ones. This "steady state" occurs first for small craters because they are more numerous.

Crater frequency graph

Figure 4.24

The number of craters accumulated on a planetary surface is a good indication of its age. If these crater densities can be related to rock units that have radiometric ages, reliable absolute ages can be assigned to geologic units exposed at the surface---a concept illustrated in this diagram, where the number of craters in a specific size range is plotted against the crater diameter. Lines, or isochrons (equal-time lines), show the number of craters accumulated on surfaces of several different ages. Note that the lunar highlands plot is above and to the left of the isochron for the lunar maria because of the marked difference in the number of superposed craters on the two surfaces. The curve labeled saturation indicates the point at which the total number of craters cannot increase. This steady state is reached when the number of craters formed by these impacts equals the number of craters obliterated.

 

 

An example of determining relative age by crater populations is shown in Figure 4.24. A curve lying above or to the right of another describes an older surface. Changes in the slope of a line are produced as a steady-state condition occurs in the smaller size classes, but "new" larger craters continue to be recorded. Crater counts on the rims of Orientale, Imbrium and Humorum basins show that although Humorum is approaching steady-state values at small crater diameters, it has a greater total frequency of craters and is thus the oldest of the three basins. This can be easily seen by comparing the number of craters of a given size, for example 10 km, on each rim. Humorum has less than 100, Imbrium around 25, and Orientale only 10.

Radiometric Dates for Lunar Events

Some of the most critical information about the geology of the Moon was obtained from isotopic age determinations of lunar rocks. The passage of time is recorded in rocks by the accumulation of the products of radioactive decay. Most rocks contain several elements (usually potassium, uranium, or rubidium) that decay to other elements, called daughter products. If the rate of decay is known, careful measurements of the amount of daughter elements in a rock reveal its absolute (as opposed to relative) age---the time at which it started to accumulate decay products. Many expected that the lunar surface was old, but the fresh lava of the maria and the bright rayed craters appeared as if they formed as recently as the Ice Age on Earth, which ended only 20,000 years ago. The first lunar rocks to be dated were basaltic lavas from one of the maria; to the amazement of many, these gave an age of 3.65 billion years---older than almost all rocks found on Earth. Additional radiometric ages of other mare lavas from other spacecraft landing sites indicate that the extrusions of lava that formed the maria began about 4.0 billion years ago and continued episodically for at least 800 million years. The lunar highlands, of course, are older, and samples collected by Apollo 17 astronauts show that some rocks crystallized over 4.5 billion years ago. This is nearly the age of the oldest meteorites. The lunar crust developed very shortly after the accretion of the Moon. The isotopic ages of most highlands rocks cluster around 3.9 to 4.0 billion years ago and probably reflect "resetting" of their radiometric clocks by intense bombardment that tapered off about this time.

Another important age was determined from ejecta thrown out of the Imbrium Basin. Fra Mauro breccias collected by the Apollo 14 astronauts were emplaced about 3.9 billion years ago. The age of Copernicus was determined from ray material collected from the Apollo 12 landing site. The age of this event, one of the most recent in lunar history, is 0.8 to 0.9 billion years. Other important ages determined from samples are shown in Figures 4.24 and 4.25.

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Figure 4.25

The variation in the number of craters formed on the Moon's surface during different periods of time shows that the rate of cratering on the Moon has not been constant but has decreased dramatically from pre-Imbrian times to the present. The most dramatic decline occurred 4 billion years ago.

 

 

By integrating these and other radiometric dates into the relative geologic time scale of the Moon determined from superposition and crater counts, it is possible to construct an absolute time scale for the Moon plus a graph showing rate of cratering (Figure 4.25). This curve shows that early in lunar history the rate of cratering was hundreds or even thousands of times greater than today. Moreover, the rate of decline in impact was very rapid until about 3 billion years ago when it reached a low level. Since then the rate of cratering has been relatively constant.

It is believed other planets and satellites experienced broadly similar variations in cratering and that crater frequency may be a crude way to correlate planetary events. However, this notion is complicated by the large uncertainties in estimates of the rates of impact on individual bodies in the solar system. Impact rates may have varied dramatically from planet to planet or satellite system to satellite system.

Volcanic Features

Volcanic activity is especially important in studies of planetary evolution because volcanic products are a type of window into the planet's interior and provide valuable insight into how the planet operates. Many geologists long suspected that the lunar maria were composed of volcanic rocks, but it wasn't until orbiting satellites photographed individual flow fronts and the Apollo astronauts brought back rock samples that we knew for certain that the maria were formed by vast floods of basaltic lava and that the Moon had a spectacular volcanic history. The highlands, too, have some volcanic rocks, but volcanic landforms are not as obvious there.

Lobate lava

Figure 4.26

The lobate fronts of lava flows in southern Mare Imbrium are similar to the flow margins of flood basalts on Earth. These vast sheets were probably erupted from long fissures.

 

 

The general nature of the lunar maria can be seen in Figures 4.1 and 4.26. The floods of lava fill the large multiring basins on the near side of the Moon and commonly overflow the rims, spilling out over the surrounding areas. From distant views it may appear that the maria represent one huge flood with the basins all filled to the same level, like the oceans on Earth, but upon closer examination we find that the mare flows are not at the same level and were extruded over a relatively long period of time. The radiometric ages of lunar samples show that the lavas filled the large basins several hundred million years after the basins were formed. This is of considerable significance in considering the origin of the lava. Obviously the lavas cannot be melts generated by impact. Instead, they were produced by melting deep within the Moon. They were formed, however, during a specific interval of time of lunar history extending from about 4.0 to at least 3.2 billion years ago and represent a thermal "event" that lasted roughly a billion years. Although the youngest samples returned from the Moon yield radiometric ages of 3.2 billion years, crater frequencies on some unsampled lava flows in northern Oceanus Procellarum indicate that some flows may be much younger, around 2.5 billion years old. If this age is correct, the length of the episode of volcanism may be much longer than indicated by the samples we now have.

When we consider the relatively small volume of lava and the time span during which eruptions occurred, plus the fact that a given eruption occurred at a very high rate, the total period of time of lunar volcanic activity must have been interrupted by long dormant periods.

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Figure 4.27

The lunar maria are covered with small craters, as is shown by the long shadows cast by the setting Sun. Small craters are abundant and lava flow fronts appear as ridges.

 

 

Although the maria appear to have flat, featureless surfaces, photos taken at low sun angle (Figures 4.27 and 4.28) show that there are numerous tiny craters which dot their surfaces. In places the mare surface appears to undulate in broad swells and depressions only a few meters high. The basalt in the interior of the mare basins completely covers the old cratered surfaces, and therefore must be in excess of 1500 to 2000 meters thick. It is interesting to note, however, that some rims of old craters project above the mare material near the margins of the basins and show that here the lava is relatively thin. These "ghost" craters show that a considerable part of the margins of the basins are covered with only 200 to 400 m of basalt (Figure 4.20).

Long lunar lava flows

Figure 4.28

The spectacular lengths of some lava flows in southern Mare Imbrium demonstrate that the lavas were highly fluid and were erupted at high rates. The flows, which terminate in the upper right corner of this photograph, are over 400 km long. The arcuate mare ridges that cross this region must not have existed when the lavas were flowing, otherwise the lavas would have ponded behind the ridges. A portion of a lava channel is visible near the lower right corner of the photo. The "islands" formed by peaks on the floor of Imbrium Basin attest to the relative thinness of the lava accumulations in this region of the Moon. The sprays of irregular craters were formed by ejecta from the crater Copernicus, which lies 460 km to the south.

 

 

With the extensive lava fields of the maria, one might expect to find spectacular stratovolcanoes like Earth's Vesuvius, Shasta, and Fuji, together with numerous cinder cones and other volcanic landforms. If such features were present they would be easily discernible from the excellent photography we have, because we can clearly see details of tiny impact craters across the surface of the maria; however, we see only a few volcanic shields and little evidence of the zones of fissures from which some of the lava was erupted. One of the reasons for this is that lavas in the maria were far more fluid than any lavas found on Earth. Instead of flowing a short distance and developing well-defined margins, they flowed great distances and ponded in depressions, almost like water accumulating in lakes. There are few visible margins to the individual flows or other characteristic flow features seen on cooled lavas on Earth. The high fluidity of lunar lavas was confirmed when samples of the maria were melted in the laboratory. Their viscosity was similar to that of engine oil.

In the southern part of Mare Imbrium, however, lavas were viscous enough to develop well-defined flow margins. Excellent details of many individual flow units are apparent on photographs from the Apollo missions, which show that the flows are similar in many respects to the fluid basaltic flows on Earth. The flows shown in Figure 4.28 are typical. You will note that many flow units were extruded from a single source region and flowed down the regional slope. Each flow is characterized by lobate margins similar to basaltic flows on Earth. The Columbia River Plateau in the northwestern United States and the Deccan Plateau of India consist, in part, of a similar series of lava floods that were extruded quietly from a network of fissures. The individual lava flows on the Moon are much longer than those found on Earth. Some traveled as much as 600 km over slopes with gradients of probably less than 1 in 100. The distance lava flows depends on slope, viscosity and, to a considerable degree, on the rate of eruption. We can conclude, therefore, that the eruption rate for the mare lavas on the Moon was very high---much higher than any known on Earth. The flows were apparently fed from huge fissure systems. As the lavas were very fluid, they may have ponded over their vents and drowned the fissures from which they issued.

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Figure 4.29

Sinuous rilles east of the Aristarchus plateau in Oceanus Procellarum are thought to be enlarged lava flow channels. Several of the channels head in low volcanic cones or shields which surround volcanic craters. In some cases it appears that the hot lava flowing through these channels has eroded the underlying surface, allowing the channel to incise itself into the surrounding terrain. These central vents contrast with the fissures, which are thought to have fed many of the mare lavas. The branching system of channels near the center of the photo are distributaries, which mark the end of a lava channel. Lunar rilles are distinct from terrestrial river channels but no close volcanic analog has been found on Earth. The rugged mountain chains in the upper part of the photo are old highland massifs, which protrude through the lava cover.

 

 

Other features associated with the maria and believed to be of volcanic origin are the meandering channels known as sinuous rilles (Figure 4.29). Most occur along the shallow edges of maria and on the flat floors of larger mare-filled craters. Many sinuous rilles begin at a volcanic crater and, when traced downslope, become progressively smaller until they disappear. Some have V-shaped profiles, but flat floors complicated by inner channels, craters, and irregular hummocks are more typical. Apollo 15 astronauts visited Hadley Rille, a large sinuous rille on the edge of the Imbrium Basin. Layered basalts crop out along its rim and its V-shaped profile is produced as loose material falls from its walls. It is typical of many rilles in size, 135 km long, 1.2 km wide and 370 m deep. At first glance, some sinuous rilles appear to be very similar to terrestrial stream valleys, but they lack many features characteristic of stream-cut channels, such as tributary systems, increase in channel size downslope, and such associated depositional features as deltas, flood plains, and alluvial fans. In all probability, the rilles originated in a variety of ways. Some are similar to collapsed lava tubes. Lava tubes are formed because a thin crust typically develops over the liquid interior of a lava flow. Pipelike zones of movement develop within the interior of the flow and may drain to form a long, hollow, cylindrical tube in the middle of the flow. The roofs of terrestrial lava tubes sometimes collapse, forming steep-sided troughs or chains of circular pits. Meandering lava channels with steep, high walls also form when the outer edges of a lava flow cool and solidify while the interior of the flow moves on downslope. This produces a lava channel or lava "gutter" (Figure 4.30). Sinuous lunar rilles are probably produced in similar fashion. Some rilles may be fault-troughs modified by flowing lavas. Many lunar rilles are, therefore, similar to common volcanic features on Earth.

Hadley Rille

Figure 4.30

A sinuous rille, called Hadley Rille, found at the base of the Apennine Mountains in the eastern part of Mare Imbrium, is interpreted to be a lava channel. The channel was visited by the Apollo 15 astronauts. The rille starts from a small volcanic crater and follows a sinuous course downhill. The margins of the lava flow(s) from this crater are not apparent in this photo.

 

 

A number of smooth low domes occur in the lunar maria and are considered by most geologists to be the lunar equivalent of low-shield volcanoes (Figure 4.31). The small volcanic features are at first difficult to distinguish from impact craters because they sometimes have nearly circular craters similar in size to many impact craters. However, close examination indicates that they can be distinguished from impact craters. (1) They are positive geomorphic features rising above the surrounding surface. Impact craters are depressions. (2) Their slopes are commonly breached, presumably by lava flows. (3) As can be seen in Figure 4.31, some occur in groups to form a broad mound of domes and flows and others occur as single isolated features. Foremost among the volcanic complexes are those in Oceanus Procellarum.

Domes of Marius Hills

Figure 4.31

Low domes in this oblique view of Oceanus Procellarum are thought to be low shield volcanoes. The Marius Hills are one of several such volcanic complexes in Oceanus Procellarum.

 

 

The Rumker Hills (Figure 4.32) are an array of domes and cones that are similar in many respects to volcanic features on Earth. In addition, rilles emanate from some of these low domes, further suggesting their volcanic origin. These domes may be a type of shield volcano and may be the source of the mare fill in some areas, particularly Oceanus Procellarum. This type of volcanism, associated with central vents and low shields, has also been identified on Earth but is not generally well-known to the layman.

Rumker Hills

Figure 4.32

The Rumker Hills are believed to be another volcanic complex in Oceanus Procellarum. The low domes are low shield volcanoes that fed lava flows both younger and older than the surrounding mare lavas. The volcanoes appear to have formed on a high patch of pre-mare ejecta from the Imbrium Basin.

 

 

For example, in the Snake River Plain of southern Idaho, the apparently flat surface of these basaltic plains commonly consists of a series of coalescing lava cones or low shield volcanoes (Figure 4.33). This type of volcanism occurs when magma supplies are small and are extruded from central vents, in contrast to the large volumes extruded from long fractures typical of flood basalt eruptions. Both types of volcanism (flood and plains) are found in northwestern United States and produce edifices similar to those seen on the Moon.

SRP model

Figure 4.33

Accumulations of lavas erupted from low-shield volcanoes and short fissures typify basaltic plains like those of the Snake River Plain in southern Idaho. Each shield is less than 10 to 30 km across---much smaller than the large shields of the Hawaiian islands. Some of the volcanic complexes of Oceanus Procellarum are probably composed of plains-style volcanism.

 

The Moon has fewer types of volcanic landforms than does Earth. One reason may be that the physical environment on the Moon inhibits the development of many volcanic features found on Earth. If a typical cinder-cone eruption occurred on the Moon, the lack of atmospheric drag on the ejected particles and the lower gravitational attraction on the Moon would allow the ejected particles to fly far from the center of eruption. As a result, the ejected material would tend to form a thin, low, widespread blanket of ash rather than a steep cinder cone. Some distinctive patches of dark material surrounding craters with low rims (Figure 4.34) are possibly the lunar equivalents of cinder cones on Earth. Another reason for the lack of a diversity of lunar volcanic features lies in the small variability of volcanic rock compositions found there. No rhyolites or even andesites, which usually form stratocones on Earth, have yet been identified. It may be that water, extremely rare inside the Moon, plays an important role in the generation of silicic magmas which produce most of the high volcanoes on Earth.

Dark halo crater

Figure 4.34

Dark-halo craters in this photograph are considered by some scientists as evidence of landforms produced by the accumulation of basaltic pyroclastic material---the lunar equivalents of cinder cones. These craters are centered on a fracture system that cuts the floor of the crater Alphonsus. The craters cap broad cones, and one has an irregular outline---features atypical of impact craters.

 

 

In summary, the basalts of the lunar maria represent a major thermal period in lunar history that extended over a period of about one billion years, beginning about 4.0 billion years ago. During this period, heat sources from within the Moon partially melted parts of the interior, generating basaltic magma that migrated upward to the surface and was extruded through the fractured crust. There are a number of distinctly volcanic features associated with the mare flood basalts, but there are no spectacular stratocones or huge shields topped with calderas such as we find on Earth and Mars. The extrusions occurred through low shield volcanoes or as quiet fissure eruptions and appear to have terminated several billion years ago. As we will see, this type of thermal event occurred on other planetary bodies as well.

Tectonic Features

Tectonic landforms are created by the deformation of a planet's outer layers and surface. They arise from forces inside the planet. Young tectonic features show that a planet still possesses enough energy to be dynamic. If no young tectonic features are found on a planet it is a sure indication that its internal heat energy has been dissipated. In general, there are two important types of tectonic forces--extension and compression. Extension causes stretching and fracturing of the shallow lithosphere. The lithosphere may fracture to form high-angle faults along which movement of blocks occurs. These extensional faults are usually long and straight (Figure 4.35). Rifts and grabens, fault-bounded valleys, are common expressions of extension.

Linear rilles

Figure 4.35

Linear scarps seen in this picture are interpreted as expressions of faulting of the Moon's surface layers. These fault-bounded valleys, called linear rilles or grabens, are generally produced by tectonic extension, which may accompany the subsidence of a basin floor or global expansion.

 

Compression is caused by shortening. If under compression, the lithosphere of a planet may buckle and fold, or break into thin sheets bounded by low-angle faults. These low-angle faults are descriptively named thrust faults and they usually form sinuous, overlapping ridges quite different from grabens (Figure 4.36).

Wrinkle ridges

Figure 4.36

Wrinkle or mare ridges are visible in this view of one of the lunar maria. These ridges are thought to be the result of thrust faulting produced by compressional forces exerted at the Moon's surface. The compression may be the result of global cooling and contraction superimposed on local adjustments to the load imposed by the mare basalts. Fault-bounded linear rilles contrast sharply with these ridges.

 

There has been almost no significant tectonic activity on the Moon during the last 2.5 billion years, or since the extrusion of the flood basalts. We know this because the millions of craters that cover the lunar surface provide an excellent reference system for even the most subtle structural deformation of the crust. Just like lines on graph paper, the circular craters that cover the lunar surface would be deformed by any significant crustal movement, compression or extension, and would record even the slightest disturbances. The network of young craters, however, is essentially undeformed and the lunar crust thus appears to have been nearly fixed throughout time. There is no evidence of intense folding or thrust faulting, and no indication of major rifts. The major features that may be attributed to crustal deformation are linear rilles and wrinkle ridges.

Linear Rilles

Many linear rilles are sharp, linear depressions, which generally take the form of flat-floored, steep-walled troughs ranging up to several kilometers in width and hundreds of kilometers in length (Figures 4.35). The valley floors lie a few hundred meters below the surrounding region. The valley walls are straight or arcuate and stand at the same elevation as though they had been pulled apart while the floor subsided as a graben. Linear rilles are parallel or arranged in an echelon pattern. Some intersect and others form zigzag patterns.

There is little doubt that linear rilles are grabens that resulted from normal faulting. However, the origins are undoubtedly quite different from those that produce the great rift systems on Earth, which are related to the pulling apart of relatively thin lithospheric plates by movement over a plastic asthenosphere. Some rilles on the Moon may be the result of an early expansion of the entire sphere as it slowly heated. Others were probably caused by relatively local stress systems set up around lava-filled impact basins that subsided due to the added weight of the basalt. Still other linear rilles formed as the result of the massive impacts that created the large lunar basins. From careful studies of the terrain crossed by these grabens, it appears that few formed after about 3.5 billion years ago.

Wrinkle Ridges

Some of the most conspicuous structural features on the mare surface are long narrow ridges, sometimes called "wrinkle" or mare ridges (Figure 4.36). Typically, they have sinuous outlines and extend discontinuously for great distances across the maria. In some places, they transect highland surfaces as well. Individual segments may be several kilometers wide, a few hundred meters high, and hundreds of kilometers long. Systems of mare ridges commonly parallel the margins of the major mare basins, although some parallel structural trends in the highlands. In contrast to lunar grabens, these ridges disrupt even the youngest mare surfaces and must have continued to form until after the maria were completely formed less than 3.0 billion years ago.

Several explanations have been proposed for the origin of the ridges, and, as is the case for rilles, wrinkle ridges may originate in several different ways. Originally, some geologists thought these ridges marked the location of the fissures through which the mare lavas were erupted or that they were produced by intrusions of magma beneath a solidified crust, buckling and folding the overlying flows. These explanations are presently discounted for several reasons. Some circular ridges apparently formed as lavas settled or compacted over preexisting crater rims. Other wrinkle ridges cross geologic units of different age and origin, for example the highlands-maria boundaries. Where they do cross into adjacent highlands terrain, they appear as simple scarps and are not associated with lava flows. Others cut craters like simple faults.

The evidence suggests that most mare ridges were produced by compression related to simple vertical adjustments of the lunar crust. It appears that these movements occurred along faults around and in the mare basins as the crust gradually adjusted to the load of the accumulating lava. A compressive stress system may have been established in the central part of the mare as the down-warping occurred, buckling the crust and producing the ridge systems. An instructive example, Mare Serenitatis, experienced earlier rille or graben formation followed by more recent mare ridge development (Figure 4.37). Some lunar scientists think that the change in the style of tectonic deformation from extensional (graben formation) to compressive (ridge formation) stresses may mark a shift in global stress patterns, in turn accentuating the production of one or the other of these features. A change from extension to compression could occur when the Moon started to cool and contract after an earlier slight expansion caused by internal heating. This shift may have occurred when the last linear rilles formed, about 3.5 billion years ago.

Mare Serenetatis

Figure 4.37

Two distinctive types of mare surfaces are shown in this view of southeastern Mare Serenitatis. Older maria occur in the lower right part of the photo and are elevated relative to the younger maria within the basin interior. The older maria are cut by numerous linear rilles that are parallel to the basin margin, and are also deformed by a number of broad mare ridges. The younger maria are not cut by linear rilles but mare ridges are prominent. These observations, also made elsewhere on the Moon, suggest that lunar tectonism changed from broadly extensional to compressional about 3.5 billion years ago---the age of the break between the two types of mare surfaces shown here.

 

In short, it appears that the preserved record of lunar tectonism is the result of small-scale vertical adjustments of the lithosphere to produce mare ridges and linear rilles. Many of these features can be related to the gradual subsidence of lava-filled impact basins superimposed on the effects of global expansion followed by contraction early in the Moon's history. The Moon does not have an active tectonic system like that of Earth and represents a primitive body that has experienced relatively little tectonic activity.

Lunar Rocks

To some people the rocks returned from the Moon were a disappointment. Few exotic minerals were found. Instead, Moon rocks are like the common rocks found on Earth. Yet these lunar rocks hold the key to understanding the thermal and chemical evolution of the Moon. Moreover, these rocks chronicle events in the early days of planetary evolution---lunar rocks are all much older than most rocks found on Earth.

Landing site geology

Figure 4.38

The structure of the Apollo landing sites are shown in these schematic block diagrams. (A) Apollo 11 landed on the flood basalt of Mare Tranquillitatis. Here, the maria consist of numerous thin layers of basalt flows with an aggregate thickness of several kilometers that were extruded about 3.7 billion years ago. These mare basalts rest on older rocks of the cratered highlands. (B) Apollo 12 landed on a ray of Copernicus, approximately 400 km south of the crater. The ray material rests on a sequence of basalt, which forms Oceanus Procellarum. Below the basalts is a layer of ejecta that was formed during the early stage of abundant meteorite impact and development of the lunar highlands. (C) Apollo 14 landed on the Fra Mauro Formation, which is composed of material ejected by the Imbrium impact event. Basalts from Oceanus Procellarum lap up against the ejecta, proving that the mare lavas are younger than the highlands. The rocks of the Fra Mauro Formation were thrown out of Imbrium Basin about 3.9 billion years ago. (D) Apollo 15 landed near Hadley Rille at the base of the Montes Apenninus, which form the rim of Imbrium Basin. Hadley is a sinuous rille of volcanic origin. The mare basalts lap up against the ancient rocks of the lunar highlands, which have been dated as more than 4 billion years old. (E) Apollo 16 landed in the highlands of the Descartes region. The surface material is composed of debris churned up by North Ray and South Ray craters and overlies layers of breccia, which were formed by more ancient meteorite impacts. (F) Apollo 17 landed in the Taurus-Littrow valley, which cuts the rim of Serenitatis Basin. The high massifs are composed of ejecta thrown from the basin 3.9 to 4 billion years ago. The valley is floored by younger mare lava flows which were erupted 3.5 to 3.7 billion years ago. A prominent wrinkle ridge transects both the mare and the highland materials

 

The samples brought back from the Moon were obtained from a variety of geologic settings (Figure 4.38) in an effort to provide maximum information. They have been studied by hundreds of scientists from many countries and are still the subject of thorough and sophisticated analyses of their physical and chemical properties. This lunar material consists of several types of igneous rocks as well as rocks created by meteoritic bombardment.

Anorthosite

Anorthosite is a coarse-grained igneous rock composed almost entirely of the mineral plagioclase (Figure 4.39). This rock type was collected from the lunar highlands and is an important constituent of the soils and breccia described below. Along with other plagioclase-rich rocks, it forms a group of rocks that are the most abundant and oldest (greater than 4.4 billion years old) on the lunar surface. Part of the significance of anorthosite is that it records a thermal event very early in the Moon's history, long before the development of mare basalts and even before the period of intense bombardment that formed the craters of the lunar highland.

Studies of the lunar samples indicate that shortly after accretion of the Moon, its outer layers melted to form a global "ocean" of molten rock, perhaps as deep as 100 to 1000 km. Crystallization within this "magma ocean" produced plagioclase feldspar crystals, which floated to the surface because they were lighter than the melt, much as ice cubes rise in a glass of water. Accumulation of these crystals produced masses of floating anorthosite "rockbergs," which coalesced to form the lunar crust. Most of these rocks have been significantly altered by subsequent metamorphism and melting caused by impact, and it is doubtful that much original, primordial crustal material remains, and then only as small grains in other rocks.

One of the most exciting finds of recent years was the discovery that small pieces of lunar anorthosites occur in meteoritic material found on Earth, specifically in Antarctica. The distinctive characteristics of lunar anorthosite make their identification certain. Apparently, this material was delivered to Earth by an energetic impact event. Careful searches of meteorite collections may yield samples of lunar rocks that come from areas not visited by any spacecraft.

Anorthosite

Figure 4.39

Lunar anorthosite, seen through a microscope, consists of a meshwork of plagioclase crystals. These are characteristically lath-shaped, with some pyroxene occupying the spaces between the plagioclase crystals. Olivine occurs in amounts as great as one percent, with small traces of opaque minerals and glass. This is an older lunar rock type which is abundant in soil and breccia from the highlands. Anorthosites are important because they record a major thermal event early in the Moon's history which formed its crust. The width of the field of view is approximately 4 mm.

 

Lunar Basalt

Most of the igneous rocks collected from the Moon's maria are very similar to terrestrial basalt, the most common rock in Earth's crust. Like anorthosite, basalts were once totally molten, as is indicated by their vesicles (gas bubbles), interlocking crystalline textures, and compositions (Figure 4.40). Basalts differ from anorthosites in their mineral constituents. The principal minerals found in lunar basalt are plagioclase, pyroxene, ilmenite, and olivine, all found in terrestrial basalt. Only minor amounts of a few minerals previously unknown on Earth were found. The most significant difference between lunar and terrestrial basalts is that the lunar basalts contain greater concentrations of refractory elements (titanium, zirconium, and chromium). Lunar basalts are also devoid of water and have much lower amounts of relatively volatile elements such as sodium and potassium than do terrestrial basalts. These chemical characteristics are significant in that they suggest that the material that forms the Moon is poor in volatile elements---perhaps because it condensed at higher temperatures than did the material that formed Earth. This could explain both the concentration of refractory elements and the low proportion of volatile elements. The texture and composition of the basalts, along with melting experiments, suggest that most of the basalts crystallized at 1500 K to 1900 K and indicate that they were more fluid than their earthly counterparts. The high fluidity (low viscosity) of lunar basalts is also reflected in the broad sheets of basalt that fill the maria basins.

Lunar basalt

Figure 4.40

Lunar basalt is similar to terrestrial basalt but contains greater amounts of refractory elements (titanium, zirconium, and chromium). Lunar basalts are generally younger than the anorthosites. The width of the field of view is approximately 4 mm.

 

There are two basic types of lunar basalt. The older, less extensive basalts have been called KREEP basalts because of relatively high concentrations of potassium (K), rare earth elements (REE), and phosphorous (P). Radiometric studies indicate that they were extruded at the surface around 4 billion years ago, prior to the formation of the Imbrium Basin and are found mainly in highly fragmented rocks from the highlands. Indeed, it is difficult to determine conclusively that they represent lavas at all. These basalts may have evolved from the primitive molten layer that surrounded the Moon. As the magma ocean cooled, crystallization and removal of various minerals gradually changed the composition of the remaining melt. A crust of floating feldspar was created over the molten zone and an iron and magnesium-rich mantle accumulated beneath it. Residual liquids trapped between these thickening layers eventually became similar to the KREEP composition and were erupted on the surface through the fractured crust. Another theory for the origin of KREEP basalt holds that the entire magma ocean solidified but remelted around 4.0 billion years ago as the Moon warmed up by radiogenic heat. Partial melting of rocks believed to be part of the lower crust yields liquids very similar to the samples of KREEP basalt returned from the Moon.

The other basalts which are more common on the lunar surface are the mare basalts. Although there is evidence that mare basalts and KREEP basalts may have been erupted during overlapping epochs, mare basalts are generally less than 3.9 billion years old and are chemically distinct from earlier magmas. The simplest explanation for the origin of mare basalts employs remelting of portions of the mantle (formed earlier by crystal accumulation at the base of the magma ocean). Melting may have occurred in zones that deepened with time, but most mare basalts appear to have been produced at a depth of around 400 km.

Breccia

Breccia is a fragmental rock in which the individual particles are angular rather than rounded like particles of sand and gravel. Lunar breccias consist of fragments of rock and glass from a variety of sources (Figure 4.41). They result from meteorite impact and fragmentation. Some are consolidated regolith that have high proportions of glassy fragments. The mechanism by which regolith is consolidated into a coherent mass of breccia presents a problem. Two possibilities are immediately apparent: (1) shock lithification (compression of the grains together as a strong shock wave passes), and (2) welding of the deposits as they accumulate in a hot state after being ejected from impact craters.

Lunar breccia

Figure 4.41

Lunar breccia consists of angular fragments of broken rocks from a variety of sources. Typically, the fragments are angular and show essentially no evidence of modification by abrasion. Some lunar breccias contain large amounts of glass, with particles that are remarkably spherical. The dark gray and black materials are glass particles. Breccia results from fragmentation and subsequent compression due to repeated impact. Breccias are the main component of the lunar regolith.

 

Regolith

The Moon's surface is covered in most places by a thin layer of relatively loose, unconsolidated fragments of rock, crystals, and glass; particles vary in size from large boulders to fine powder. This layer is called the lunar soil or regolith. The presence of regolith was demonstrated on all lunar landings as dust thrown up by rockets of the landers, and the ever-present dust clinging to the astronauts and lunar rover. The average thickness of the regolith depends upon the age of the surface on which it has been formed. The regolith on ejecta from very young craters such as Tycho may be only about 10 cm thick. On the maria, the average thickness is 5 m. The regolith in the older highlands is possibly more than 10 m thick. As a general rule, then, the older the surface, the thicker the regolith.

Regolith on the Moon consists of debris thrown out of craters and, at any given place, most of the debris have been derived from the local underlying substratum. As a result, the composition and texture vary considerably from place to place and reflect the history and the processes of the area where it is formed. For example, samples of regolith from the maria contain considerable amounts of basalt fragments and generally less than 50 percent glass, whereas regolith on the rays of the crater Copernicus contains 70 to 90 percent glass.

We can estimate the rate at which regolith forms on the Moon by measuring its thickness and determining the age of the bedrock beneath. At Tranquillitatis, the lunar basalt is about 4 billion years old and the regolith is about 4 m thick. This gives an average rate of regolith formation of 1 mm per million years. This accumulation rate is somewhat misleading because regolith production must have proceeded at much higher rates during the early bombardment of the Moon and has since slowed to very low rates.

The surface of the Moon is modified by the churning action of impact which is the major factor in fragmenting solid rock and developing a regolith. There is, in addition, a type of microscopic weathering that occurs on the Moon. Since the surface of the Moon is unprotected by an atmosphere or global magnetic field, the regolith is continually bombarded by micrometeorites, solar winds, and galactic cosmic rays. The net effect of this bombardment is to slowly change the regolith with time. One important change is the welding of particles together by glass generated by impact forming glass-bonded aggregates.

Thus, the outer few centimeters of the lunar surface is where interactions between space processes and the Moon take place. The core-tube samples of the lunar regolith returned by the astronauts are of particular interest in that the complex layers of brecciated rock record such interactions in the distant past. The Moon may therefore hold an important record of the Sun's activity billions of years ago.

Glass beads

Figure 4.42

Glass particles are common in the lunar regolith. They are formed by the shock melting of rock fragments during impact.

 

Glass

Glass particles are abundant in nearly all samples of lunar regolith and are one of the features which distinguish lunar from terrestrial soils. The glass occurs as beautifully formed spheres and tear drops, and as spatter on other fragments (Figure 4.42). The lunar glass is distributed with ejecta and is believed to have formed by shock melting of rock debris during the process of impact. Samples from the rays of Copernicus, for example, consist of 70 to 90 percent glass. Under high magnification, the glass "beads" show pits, grooves, and spatter that results from micrometeorites striking the glass surfaces after the particles had cooled and accumulated in the regolith (Figure 4.43). Thus, practically every rock fragment in the regolith appears to have been involved in one or more impact events. Other glass beads, orange and green in color, were probably produced by lava fountaining near the fissures from which the mare basalts were extruded.

Small crater

Figure 4.43

Micrometeorite craters on the surface of crystal fragments in lunar breccia show that the products of impact on the Moon range from the large multiring basins to microscopic pits. Practically every rock fragment on the Moon appears to have been involved in one or more impact events.

 

In summary, a variety of rock types has been found on the Moon, but the diversity is not as great as that found on Earth. Only trivial quantities of granites, peridotites, and so on, have been identified. Apparently, the rock samples returned from the Moon tell exactly what one would expect from studies of the surface features. They record the details of impact and basaltic volcanism, the two major geologic processes that had operated on the Moon. Perhaps their greatest value is the information they provide concerning absolute dates of major lunar events, establishing a radiometric time scale for another planet and informing us of its chemical and thermal evolution.

The Internal Structure of the Moon

Many of the major geologic processes that shape the surface of a planet are driven by forces from within the planet itself. Volcanic activity, tectonic movements, faulting and folding, and the generation of a magnetic field are examples of these internally derived processes. A thorough understanding of the interior structure of the Moon reveals much about its present state and past geologic history.

Our present understanding of the internal structure of the Moon is based on a variety of physical observations such as density, magnetism, and seismicity, as well as on the rock types and landforms present at the surface. Much remains uncertain, but several facts place significant constraint on what the internal structure may or may not be. First, the bulk density of the Moon is 3.34 g/cm3, whereas the mean density of lunar surface rock is about 3.3 g/cm3. Thus, the density of the surface material is only slightly less than that of the Moon as a whole, and there is little

possibility for a significant increase in density with depth. (Earth, in contrast, has a mean density of 5.5 g/cm3 with the density of surface rocks being only 2.7 g/cm3, clearly indicating that the interior of Earth is much denser than the crustal material.) Nevertheless, from studies of lunar rocks, surface features, and lunar seismicity, it is clear that the Moon is layered and that the composition of the interior differs from that of the rocks exposed at the surface. In addition, seismic energy on the Moon is 1 million times less than that on Earth, implying that internal convection is not occurring at shallow depths in the mantle.

Moon's interior

Figure 4.44

The internal structure of the Moon is interpreted from measurements of density, gravity, magnetism, and seismic properties. The thickness of the crust ranges from 60 to 100 km. The rigid mantle extends to a depth of about 1000 km and together with the crust makes up the lunar lithosphere. Most moonquakes originate in the region near the base of the lithosphere. The center of the Moon appears to consist of asthenospheric mantle, possibly surrounding a small iron-sulfide core.

 

The favored model of the Moon's interior is shown in Figure 4.44. The major units are (1) a crust 60 to 100 km thick, (2) a rigid mantle extending down to a depth of 1000 km, and (3) an asthenosphere which may surround a metallic core.

The Crust

Using instruments called seismometers, Apollo experiments recorded the strength of "moonquakes" at various points on the Moon, providing a way of "x-raying" the Moon and seeing its internal structure. Variations in the velocities of seismic waves with depth show that the crust is layered (Figure 4.45). In the maria the surface layer is composed of a thin, 2- to 5-km layer of basaltic lavas. Below this, the seismic velocities are low but gradually increase to depths of almost 25 km. This zone of shattered rock is the result of extensive fragmentation and fracturing of bedrock, caused by repeated impact of meteorites. The rapid increase in the velocity from the surface to 25 km is believed to be the result of fewer fractures in the anorthositic crust. From 25 km to 60 km, the velocities are similar to those determined for unfractured anorthosites. Results from Apollo experiments show that the crust varies in thickness from approximately 50 km on the near side to perhaps 100 km on the far side. The average thickness may be about 70 km.

Lunar crust

Figure 4.45

The structure of the lunar crust in the Oceanus Procellarum region has been interpreted from variations in the velocity of seismic (earthquake) waves. Discontinuities in seismic velocities indicate that the crust is layered. Seismic velocities increase very rapidly with depth to about 1 or 2 km beneath the surface. A very sharp increase occurs at a depth of about 25 km. Between 25 and 60 km below the surface, the velocities are nearly constant. A significant increase in velocity marks the base of the crust. From comparisons of these velocities with the velocities of seismic waves in major rock types, it appears that, from the near surface to a depth of about 2 or 3 km, the crust is composed of impact breccia and basaltic lava flows. Below this, to a depth of about 25 km, is a layer composed largely of brecciated and fractured anorthositic and gabbroic rocks like those exposed in the lunar highlands. Below a depth of about 25 km rock fractures are not abundant and seismic velocities in the anorthositic crust increases markedly. Fractures produced by impact may have been "healed" at this great depth or, alternatively, may never have formed. Below 60 km is the mantle of the Moon, which is believed to be rich in olivine and pyroxene.

 

The Mantle

At a depth of 60 to 100 km, a sharp increase in velocity occurs, marking the contact between the lunar crust and mantle (Figure 4.45). The mantle rocks show a higher seismic velocity than crustal rocks and are believed to be rich in olivine, pyroxene, iron, and magnesium. The thick mantle is rigid and, together with the crust, makes up the Moon's lithosphere. The lunar lithosphere is much thicker (1000 km versus 100 km) and more rigid than Earth's. This thick lithosphere makes it nearly impossible for molten lava to reach the surface and prohibits lateral movements like those that have produced continental drift on Earth. In essence, the Moon is a "one-plate" planet. Most moonquakes occur at the bottom of this layer (800 to 1000 km deep) and are apparently triggered by Earth tides. Just as the Moon's gravitational attraction produces tides in Earth's oceans, Earth exerts a similar but stronger pull on the Moon's surface, slightly deforming the lithosphere.

The Core

As shown in Figure 4.44, an inner zone may exist beneath the moonquake zone. S-waves traversing the deep interior (below about 1000 km) are much weaker, suggesting that the rocks are partly molten, perhaps somewhat like Earth's asthenosphere. This lunar asthenosphere may extend to the center of the Moon or it may form a thin shell surrounding a solid metallic core, which may have a radius of about 400 to 700 km. As yet, evidence does not demand an iron core, but the detection of remanent magnetism in lunar rocks suggests that there may be a core composed of iron or iron sulfide. Remanent magnetism is acquired when hot magmas or rocks cool in the presence of a magnetic field. Geochemical studies of lunar samples also suggest that the Moon has a small metallic core. Even though the present-day Moon has no internally produced magnetic field, it appears that at one time thermal convection within the core produced a magnetic dynamo which set up a field that magnetized ancient igneous rocks as they cooled. Detailed studies show strong remanent magnetism occurs only in rocks between 3.9 and 3.1 billion years old. This suggests that a magnetic field may have been produced at about the same time as the mare basalts were erupted. The lunar magnetic field may have been caused by heat released from the core once it was cool enough to begin to crystallize. Subsequently, as the entire core cooled movement of material slowed, decreasing the strength of the magnetic field.

The Geologic History of the Moon

Each planetary body has a time-varying set of dynamic geologic systems that modify and shape its surface and deep interior. On Earth, these include the hydrologic system, in which water deposits new rocks and erodes away older ones, and the tectonic system, which produces shifting lithospheric plates, crustal deformation, mountain building, volcanic activity, and the growth of continents. An entirely different set of geologic processes shaped the Moon's history and its surface. We have seen that billions of years ago huge meteorites impacted the surface, at other times great lava flows spread across its surface, and intermittently the surface was warped, cracked, or faulted. At present, only occasional small meteorites strike the surface, changing only very small areas, so the Moon is essentially geologically inactive---a striking contrast to dynamic Earth.

Utilizing data obtained from studies of the relative ages of lunar landforms and rocks, the absolute age determinations of lunar rock samples, and the composition of lunar rocks and soils, a sequence of events outlining the major events in the geologic evolution of the Moon can be constructed.

Stage I: Formation of the Moon (4.6 To 4.5 Billion Years Ago)

The Moon, like other planets in the solar system, is believed to have been formed by accretion of many smaller objects in a short period of time about 4.55 billion years ago. One scenario for the formation of the Moon calls for the accretion of material in orbit around a larger, but still accreting, Earth. This theory has difficulties in explaining why Earth is volatile-rich and the Moon volatile-poor. Alternatively, the Moon may have formed in another part of the solar system, only to be captured by Earth's gravity. This theory has severe dynamical problems regarding the Moon's present orbit. Another, presently popular, theory suggests that the Moon accreted near the ancient Earth from material ejected from Earth's already differentiated interior by the impact of a Mars-sized object. It is speculated that the impactor and part of Earth's mantle vaporized at temperatures between 1500 and 2000 K. This material cooled, condense to solids, and reaccreted in Earth orbit. The giant impact hypothesis is attractive because it explains why the Moon has a composition similar to Earth's mantle for most elements but is poor in volatile materials. The volatile elements simply remained vapors after the impact leaving the solids depleted in such materials. This hypothesis also explains why the Moon's metallic core is so small. The material that was ejected came principally from Earth's mantle and the mantle of the impactor. Apparently, core formation had already occurred inside both bodies.

Moon's thermal history

Figure 4.46

The thermal history of the Moon is summarized in this schematic time versus depth diagram. Accretion of the Moon led to widespread shallow melting and the creation of a magma ocean hundreds of kilometers deep from which the anorthositic crust crystallized. Subsequent heating by radioactive decay led to partial melting of the mantle (producing an asthenosphere) at increasingly greater depths. The differentiation of the interior is shown as including the formation of a small lunar core overlain by the asthenosphere. The thickness of the lithosphere steadily increased with time, achieving a present-day thickness of 1000 km. The early history of the Moon appears to have been marked by slight expansion, followed by cooling and later contraction.

 

In any case, the Moon probably heated up tremendously during the process of accretion as the kinetic energy of impacting bodies was converted to heat. Additional heat may have come from the decay of short-lived radioactive elements. This may have resulted in complete melting of the outer several hundred kilometers of the Moon, and the creation of an ocean of molten rock (Figure 4.46). As this global magma ocean cooled, minerals began to crystallize and separate according to their densities. Heavy mafic minerals rich in iron and magnesium, like the mineral olivine, sank, and lighter minerals, mainly plagioclase, migrated upward and were concentrated in a surface layer. This layer of anorthositic rock formed the lunar crust and is presently exposed in the highlands. It is easy to imagine a very active, mobile crust sliding across the convecting magma ocean during the period of crustal accumulation. The oldest anorthosites have ages of about 4.5 billion years. Some of the earliest volcanic rocks, the KREEP basalts, were erupted several hundred million years after the formation of the crust (Figure 4.46). They may have formed as residual pockets of magma trapped beneath the thickening crust and eventually broke through to the surface along fractures. These early developments led to a crust composed of anorthosite, gabbro, and KREEP basalt. From a geologic point of view, the development of a global crust, which is now exposed in the lunar highlands, was the first important event in the evolution of the Moon.

Gradually the crust cooled and a rigid lithosphere, which thickened with time, formed. As the lithosphere thickened, lateral movements were inhibited and eventually halted completely. Tectonic disturbances must have been obvious until this time, but because of the intense meteoritic bombardment that the Moon experienced, no evidence of them has been preserved. The development of a lunar plate tectonic system was hindered by the rapid development of this thick globe-encircling shell of material. The uniformly low density of anorthosite precluded Earth-style subduction of lithosphere back into the mantle; lithospheric subduction appears to be driven by the gravitational sinking of dense lithospheric plates.

Lunar history

Figure 4.47

The major events in lunar history include: intense meteorite bombardment during an early period, formation of multiring basins, extrusion of mare basalts, and, subsequently, light meteorite bombardment. (A) Stage I. Formation of Moon by accretion (about 4.55 billion years ago) created densely cratered terrain over the entire surface of the Moon. The outer layers of the Moon may have been completely molten before this surface was shaped. (B) Stage II. The formation of multiring basins (Imbrium Basin formed 3.9 billion years ago) is attributed to the impact of asteroid-sized bodies. The infall of these meteoritic bodies may represent the final stages of accretion. Remnants of this cratered surface are preserved in the lunar highlands. (C) Stage III. Extrusion of the mare basalts (from about 4 billion to perhaps 2.5 billion years ago) was a manifestation of a major thermal event in lunar history, which occurred when the lithosphere was still relatively thin. Lava flows filled many of the multiring basins on the Moon's near side and in some areas they covered parts of the highlands that lack obvious multiring structures. (D) Stage IV. Relatively light meteorite bombardment (from 3.2 billion years ago to the present) formed some craters with bright rays, but the rate of cratering has been greatly reduced. The lunar landscape has changed little during the last 3 billion years.

 

Stage II: Pre-Nectarian and Nectarian Periods (4.6 To 3.9 Billion Years Ago)

The next major event (perhaps a continuation of accretion) was a period of intense bombardment of large and small bodies. The impact of these bodies formed a densely cratered terrain over the entire surface of the Moon, a terrain that is now preserved in the lunar highlands (Figure 4.47a). Meteorite impact may have churned the ancient magma ocean or later fragmented the crust to facilitate the eruption or exposure of KREEP basalts. Events in Stage II include the impact of large asteroid-size bodies which produced most of the multiring basins. These large objects may have been small proto-moons formed at the same time as the Moon, their infall resulting from the larger gravitational pull of the Moon, or they may have been asteroids nudged from beyond Mars by the gravitational affects of Jupiter. In any event, their collisions significantly modified the surface of the Moon (Figure 4.47b). At least 40 basins with diameters ranging from 300 km to over 1000 km are older than Imbrium Basin. During and after the formation of the multiring basins, impacts of smaller meteorites formed craters, such as Plato, Archimedes, and Sinus Iridum, on the large basins. These craters show that there was a significant time lapse between the formation of the multiring basins and their filling with lava.

During this stage of development some planets outgassed envelopes of fluids at their surfaces. Because of its bulk composition and relatively small mass, the Moon developed neither an atmosphere nor a hydrosphere. No water has been found in any of the lunar rocks, and other volatile elements are found only in very low abundances. By contrast, terrestrial rocks found at the surface commonly contain several percent water by weight. Possibly, the Moon accreted from materials that were depleted in volatiles as a result of vaporization during the hypothetical impact mentioned above. The low lunar gravity, determined by its mass, is also important for the present lack of an atmosphere. Even if tremendous amounts of volatiles had been released from the interior during an early melting episode, they would have been able to escape quickly into space. Larger planets with larger gravitational attractions have retained these gases to form atmospheres and hydrospheres which continually change the surface landforms and create new rock bodies.

Stage III: Imbrian Period (3.9 To 3.2 Billion Years Ago)

The next major events in lunar history were the extrusions of the mare lavas that cover large areas of the Moon (Figure 4.47c). The lavas erupted episodically during an interval of at least 800 million years and probably over a much longer time. There is even some evidence for pre-Imbrian basaltic volcanism. The basalts were generated at depths of about 400 km in the lunar mantle. Heat was most likely provided by radioactive decay. The zone in which basaltic magma was generated migrated inward as the Moon slowly cooled and its lithosphere thickened. Similar plains are found on Mercury and Mars. It is therefore possible that an early thermal event and the eruption of basaltic lavas represent basic elements in the evolution of all terrestrial planets.

A fundamental problem in lunar geology is the distribution of the mare lavas in near-side basins underlain by thin anorthositic crust (less than 60 km thick) while the far side of the Moon has fewer maria and typically thicker crust, from 60 to 100 km thick. The near side crust may have been preferentially thinned by the excavation of several basins formed coincidentally in the same region. As the Moon heated by radioactive decay and mare basalts were produced, the magmas would have been able to escape to the surface much more easily through these areas of already thin crust. By the time the basalts formed, the lithosphere was too thick to be much deformed but could still be punctured by these hot liquids. Although the lithosphere bent slightly beneath the loads exerted by accumulations of these lavas (producing ridges and rilles), it was strong and thick enough to prohibit their complete isostatic compensation. Calculations indicate that a switch from early graben and rille formation to compressive faulting and mare ridge formation occurred about 3.5 billion years ago. This change is attributed to the Moon's cooling and contraction.

Stage IV: Eratosthenian and Copernican Periods (3.2 Billion Years Ago To Present)

The internal differentiation of the Moon was essentially complete by the end of Stage III, and most of the Moon's present features were developed. The most significant events to occur since that time were the impacts of meteorites to form post-maria craters (Figure 4.47d). The influx of meteorites was greatly reduced, and most authorities consider the post-maria craters to have been formed by bodies from the asteroid belt or from comet nuclei.

Minor local volcanic activity, such as the domes in the Marius Hills, has probably occurred since the maria formed. As the lithosphere thickened, the depth of basalt sources also migrated downward; by about 2.5 to 2.0 billion years ago this rigid layer was too thick to allow the extrusion of lavas. It appears that the tectonic evolution of the Moon was then complete. Volcanism ceased and the Moon entered a terminal quiet state. Today the lithosphere is probably 1000 km thick, so the Moon is geologically active only at great depths. Aside from the occasional impact of a small meteorite, landscape evolution is largely complete.

Conclusions

The vast amount of new knowledge obtained about the Moon during the period of lunar exploration permitted geologists to study, for the first time, the details of another planet. In a sense, the Moon is a controlled experiment that shows a planet evolving without a hydrologic system and with a single global tectonic plate. Cratering, which is responsible for the formation and modification of most of the lunar landscape, was the most important surface process. Major thermal events did occur on the Moon, however, during differentiation and when basaltic lavas were extruded in a series of eruptions to form the lunar maria. Only minor tectonic features have been found that appear to result from relatively small down-warping of the lithosphere combined with modest expansion followed by contraction of the Moon.

A geologic time scale has been developed for the Moon using the principles of superposition originated by geologists studying Earth. Radiometric dates of lunar rock samples provide a base of absolute time for events in lunar history. Perhaps the most important aspect of the Moon's geologic evolution is that most of it occurred during the early history of the solar system, before even the oldest rocks on Earth were formed. The Moon thus provides important insight into planetary evolution unobtainable from studies of Earth.

Soon after the planet was formed by accretion, the a crust of anorthositic rock and minor basalt formed and is preserved in the highlands. A mantle and perhaps a small iron core formed at the same time. The crust was then subjected to a period of intense bombardment. The lunar maria formed from vast floods of lavas subsequent to the heavy cratering. Since then (about 3.0 billion years ago) the Moon has been inactive, and the surface has been only slightly modified by relatively few impact craters.

Today, the Moon does not have a significant source of internal energy nor a tectonic system like Earth. It has no continents nor ocean basins and no deformed rocks resulting from mountain building. Moreover, because of its volatile-poor character, it has no atmosphere nor surface fluids, so it lacks a hydrologic system to modify its surface. The Moon apparently lacked sufficient mass to experience an extended differentiation history. Small bodies radiate away their internal heat at a much higher rate than larger bodies because they have larger surface-area mass ratios. This is because all heat must eventually escape through the surface. Small planetary bodies like the Moon cooled much faster than larger ones. Their thermal and tectonic evolution proceeded at an accelerated pace and terminated when the lithosphere became so thick that it could no longer be deformed.

Review Questions

1. Draw a diagram of the internal structure of the Moon and briefly describe the core, mantle, asthenosphere, lithosphere, and crust.

2. Outline the stages in the production of a crater by the impact of a meteorite. What geologic features are produced by impact?

3. Contrast the origin and features of an impact crater with one produced by volcanic activity.

4. Explain two ideas for the origin of a multiring basin.

5. How are lunar craters modified as time passes?

6. Why are there fewer craters on the lunar maria than on the highlands?

7. What are the major differences between the highlands and maria on the Moon?

8. In what way is the Moon asymmetric?

9. Why are there so few impact craters on Earth as compared to the Moon?

10. What do linear rilles and wrinkle ridges tell us about the tectonic history of the Moon?

11. Discuss some of the possibilities why cinder cones are so rare on the Moon.

12. If the Moon's crust formed from a magma ocean about 4.5 billion years ago, why are rocks that have radiometric dates this old so rare on the Moon?

13. The Moon's highland crust consists largely of plagioclase feldspar. Why does this imply that a lunar magma ocean once existed? How was plagioclase separated from pyroxene and olivine crystals forming in the magma at the same time? 1

4. Could Earth and the Moon have once been part of a larger body?

15. Outline the major events in the history of the Moon.

16. Explain how a geologic time scale was developed for events in the Moon's history.

Key Terms

anorthosite

ballistic

basalt

basaltic plains

breccia

compression stage

Copernican Period

crater degradation

ejecta

Eratosthenian Period

excavation stage

Imbrian Period

impact melt

KREEP basalt

lava channel

lava tube

linear rille

mare basalt

maria

megaterrace

modification stage

multiring basin

Nectarian Period

nested crater

peak-ring basin

Pre-Nectarian period

rarefaction

rays

regolith

remanent magnetism

sinuous rille

terrae

wrinkle ridge

Additional Reading

Bowker, D. C., and J. K. Hughes. 1971. Lunar Orbiter Photographic Atlas of the Moon. NASA SP-206. Washington, DC: National Aeronautics and Space Administration.

 Hartmann, W. K. 1977. "Cratering in the Solar System." Scientific American. Vol. 235, No. 1, pp. 84--99.

Mutch, T. A. 1972. Geology of the Moon. Princeton, NJ: Princeton University Press.

Masursky, H, G. W. Colton and F. El-Baz. 1978. Apollo over the Moon: A view from orbit. Washington, DC: National Aeronautics and Space Administration.

Schultz, P. H. 1976. Moon morphology. Austin: University of Texas Press. Taylor, S. R. 1975. Lunar Science: A Post-Apollo View. New York: Pergamon Press, Inc.

Taylor, S. R. 1982. Planetary Science: A Lunar Perspective. Houston: Lunar and Planetary Institute.

Wilhelms, D. E. 1987. The Geologic History of the Moon. U. S. Geological Survey Professional Paper 1348, Washington DC, U.S. Geological Survey.

Wood, J. A. 1975. "The Moon." The Solar System. New York: W. H. Freeman and Co., pp. 69-80.

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