Ever since their recognition, multi-ring basins have fascinated and vexed scientists attempting to reconstruct the geological history of the Moon. As the other terrestrial planets were photographed at high resolution, it became apparent that basins are an important element in the early development of all planetary crusts. This importance spurred research into the basin-forming process and yielded a plethora of models and concepts regarding basin origin and evolution. In this chapter, I outline the general problem areas of basin formation and describe the approach taken by my own work on lunar basins.

Multi-ring basins and their significance

Multi-ring basins are large impact craters. The exact size at which impact features cease to be “craters” and become “basins” is not clear; traditionally, craters on the Moon larger than about 300 km have been called basins (Hartmann and Wood, 1971; Wilhelms, 1987). Basins are defined here as naturally occurring, large, complex impact craters that initially possessed multiple-ring morphology. This definition purposely excludes simulated, multi-ring structures that result from explosion-crater experiments on the Earth and whose mechanics of formation differ from impact events (e.g., Roddy, 1977), although important insights into the mechanics of ring formation may be gained from these studies. The qualification that basins initially possessed multiple rings is in recognition of the fact that many older, degraded basins display only one or two rings, even though their diameters of hundreds of kilometers indicate that they had multiple rings when they originally formed.

The Orientale basin (Figure 3.1) is located on the western limb of the Moon; most of the basin and its deposits extend over the lunar far side. The name “Orientale” (eastern) is derived from the old astronomical practice of displaying telescopic photographs of the Moon with south at the top and east-west convention derived from terrestrial coordinates. The first studies of the Orientale basin utilized Earth-based telescopic photographs; because of the basin's location on the extreme limb, these photographs were geometrically rectified at the Lunar and Planetary Laboratory (University of Arizona) to create a vertical viewing perspective (Hartmann and Kuiper, 1962). These rectified photographs also served as the primary data base for an early geological study of Orientale as the prototype lunar basin (McCauley, 1968).

Because of the spectacular success of the Lunar Orbiter spacecraft, particularly a series of high-quality photographs by Orbiter IV that provide contiguous coverage, the number of detailed geological descriptions of the Orientale basin increased dramatically in the years immediately following the Apollo missions (Hartmann and Wood, 1971; Head, 1974a; Moore et al., 1974; McCauley, 1977; Scott et al., 1977; Spudis et al., 1984b). In this chapter, I review the regional and local geology of Orientale and, in conjunction with data from photogeology and remote sensing, integrate these data into a geological model for the formation and evolution of the basin.

Regional geology of the Orientale impact site

The Orientale basin is sparsely filled by mare basalt and located in rugged highland terrain on the western limb of the Moon (Figure 3.1).

Multi-ring basins are features produced by the collision of solid bodies with the planets, so the basin problem is a subset of the more general problem of impact cratering, a vast field of study. This chapter briefly describes the impact process from theoretical considerations, from the evidence of some well studied terrestrial impact craters, and from the observed morphology of impact craters on the Moon and their systematic changes with increasing crater size.

The cratering process

Impact mechanics

Our understanding of what happens when a solid body hits a planetary surface at high speeds has increased greatly over the past 25 years. The study of the physical processes occurring during impact events is called impact mechanics. Although the details of this complex process are not understood, laboratory experiments, explosion craters, natural impact craters, and computer simulations have given us a general outline of the main stages that characterize the formation of an impact crater.

Solid bodies collide with planetary surfaces at very high speeds; such impact speeds are in the range called hypervelocity. Encounter velocities can vary from lunar escape velocity (about 2.5 km/s) at a minimum, up to many tens of kilometers per second (on the basis of velocities of bodies in heliocentric orbits). On the Moon, the mean impact velocity is about 20 km/s (Shoemaker, 1977). At the moment of contact between an impactor and a planet, the kinetic energy of the impacting body is transferred to the planetary surface target. A shock wave propagates into the target and projectile, resulting in intensive compression of both objects. In hypervelocity impacts, the quantities of energy produced greatly exceed the heat of vaporization for geological materials.

The formation of multi-ring basins dominated the early geological evolution of the Moon. The five basins described in the preceding chapters represent a spectrum of basin ages, sizes and morphologies. By comparing the similarities and differences among these basins, some general inferences may be made regarding the process of formation of multi-ring basins on the Moon. I here synthesize the information described in the previous chapters to develop a model for the formation and geological evolution of multi-ring basins on the Moon. This model is incomplete, but several puzzling aspects of basin geology can be explained satisfactorily through this approach. At various points in the following discussion, please refer to preceding sections in the text.

Composition and structure of the lunar crust

The crust of the Moon is heterogeneous on a local and a regional scale; the impact targets for lunar multi-ring basins were similarly heterogeneous. The crustal thickness at the basin target sites was widely varied, ranging from 50 km thick for parts of the Imbrium basin to over 120 km thick for the Orientale highlands (Bills and Ferrari, 1976). Moreover, lithospheric conditions during the era of basin-forming impacts changed with time in response to rapidly changing thermal conditions within the Moon 4 Ga ago (Hubbard and Minear, 1975; Solomon and Head, 1980). The older basins formed in a relatively thin, easily penetrated lithosphere that gave rise to extensive post-impact modification.

The Serenitatis basin is on the near side of the Moon, east of Mare Imbrium and north of Mare Tranquillitatis (Figure 1.1). The basin is almost completely flooded by mare basalts (Figure 6.1) and displays a mascon gravity anomaly (Sjogren et al., 1974). The Serenitatis basin was recognized as multi-ring in the studies of Hartmann and Kuiper (1962), Baldwin (1963), Hartmann and Wood (1971) and during systematic geological mapping of the Moon (Wilhelms and McCauley, 1971). Because of the large amount of mare flooding and generally degraded appearance of the basin, Serenitatis was once considered to be one of the oldest basins on the Moon (Hartmann and Wood, 1971; Wilhelms and McCauley, 1971). This view has changed, primarily because of ages obtained for some Apollo 17 samples considered to represent impact melt of the Serenitatis basin (James et al., 1978; Wilhelms, 1987). I will describe the regional geology of the Serenitatis basin and some aspects of Apollo 17 site geology that relate to problems in the interpretation of its formation and subsequent evolution.

Regional geological setting and basin definition

The Serenitatis basin is close to the Imbrium basin and the effects of Imbrium on the morphologic evolution of Serenitatis have been significant. Most interpretations of basin geology rely on the well exposed highlands to the east of Mare Serenitatis (Figure 6.1). Thus, the morphological data available for interpreting the geology of the Serenitatis basin are limited compared with those for some of the other basins described in this book.

The Imbrium basin (Figure 7.1) is probably the most studied multi-ring basin on the Moon. Prominently located on the lunar near side, west of Mare Serenitatis and east of the large maria Oceanus Procellarum (Figure 1.1), the Imbrium basin first attracted the attention of G.K. Gilbert (1893) in his historic analysis of lunar craters. Gilbert recognized the extensive pattern of radial texture associated with Imbrium and postulated that Mare Imbrium had formed by the collision of a large meteorite with the Moon. The impact origin of the Imbrium basin was also recognized by Dietz (1946), Baldwin (1949; 1963), Urey (1952), and Hartmann and Kuiper (1962). The landmark paper of Shoemaker and Hackman (1962) proposed a global stratigraphic system for the Moon based on the deposition of ejecta from the Imbrium basin as a marker horizon. The Imbrium impact was considered such a key event in lunar geological history that two Apollo missions (Apollo 14 and 15) were sent to landing sites specifically chosen to address problems of Imbrium basin geology.

Regional geology and setting

Imbrium is one of the youngest major basins on the Moon, but extensively flooded by mare basalt (Figure 7.1). Even so, as one of the largest lunar basins (main topographic rim 1160 km in diameter), it has an ejecta blanket so extensive that almost all of the near side may be dated relatively with respect to the time of the Imbrium impact (Wilhelms, 1970).

The formation of multi-ring basins was an important process in the early histories of Solar System bodies. Thus, study of basins on the other planets potentially can give us insight into the early geological evolution of the planets. Although occurring on all of the terrestrial planets, the most and best preserved basins occur on planets that display remnants of their early crusts, i.e., Mercury, Mars, and the icy satellites of Jupiter and Saturn. In this chapter, I discuss the geology of basins found on the terrestrial planets in relation to the geological model for basin formation and development on the Moon discussed above.

Earth

Most of the recognized impact structures on the Earth are either simple, bowlshaped craters or complex craters displaying central peaks (Grieve and Robertson, 1979; Masaitis et al., 1980; Grieve, 1987). However, several of Earth's larger craters have multiple rings; seventeen craters display at least two rings (Table 9.1; Pike, 1985 and references therein). The paucity of terrestrial multi-ring basins doubtless reflects the relatively youthful average surface age of the Earth, as compared with the more primitive terrestrial planets, such as Mercury and Mars.

Impact craters of the Earth show the morphological transitions with increasing size, as do craters on the planets, but changes in form occur at different diameters (Pike, 1985). Complex craters on the Earth range in size from about 4 km to about 25 km in diameter.

Introduction

The region of the spectrum in the vicinity of 10 μm wavelength is called the thermal infrared. It is important because many materials have strong vibrational absorption bands there (Chapter 3). In most remote-sensing measurements these bands can be detected only through their effects on the radiation that is thermally emitted by the planetary surface being studied. Many substances have overtone or combination bands at shorter wavelengths, and although the latter bands are observed in reflected light, their depths and shapes may be affected by the thermal radiation that is emitted by the material. Hence, even though the primary subject of this book is reflectance, it is important that the effects of thermal emission be discussed. It will be seen that most of the preceding discussions of reflectance also apply to emissivity at the same wavelength because of the complementary relation between the two quantities.

Figure 13.1 shows the spectrum of sunlight reflected from a surface with a diffusive reflectance of 10%, compared with the spectrum of thermal emission from a black body in radiative equilibrium with the sunlight, at various distances from the sun. Clearly, thermal emission can be ignored at short wavelengths, and reflected sunlight at long, but at intermediate wavelengths in the mid-infrared the radiance received by a detector viewing the surface includes both sources.

Scientific rationale

The subject of this book is remote sensing, that is, seeing “with clearer eyes.” In particular, it is concerned with how light is emitted and scattered by media composed of discrete particles and what can be learned about such a medium from its scattering properties.

If you stop reading now and look around, you will notice that most of the surfaces you see consist of particulate materials. Sometimes the particles are loose, as in soils or clouds. Sometimes they are embedded in a transparent matrix, as in paint, which consists of white particles in a colored binder. Or they may be fused together, as in rocks, or in tiles, which consist of sintered ceramic powder. Even vegetation is a kind of particulate medium in which the “particles” are leaves and stems. These examples show that if we wish to interpret quantitatively the electromagnetic radiation that reaches us, rather than simply form an image from it, it is necessary to consider the scattering and propagation of light within nonuniform media.

One of the first persons to use remote sensing to learn about the surface of a planet was Galileo Galilei.

Introduction

Chapter 8 treated the bidirectional reflectance of an optically thick, plane-parallel particulate medium in which the particles were randomly oriented and could be regarded as embedded in a vacuum. In this chapter we will discuss the effects on the reflectance when each of these restrictions is removed.

Diffuse reflectance from a medium with a specularly reflecting surface

The upper surfaces of many particulate materials may be sufficiently smooth on a scale comparable to the wavelength that light is scattered both quasi-specularly from the surface and diffusely from below the surface. The specular component is known as regular reflection. Because a surface effectively becomes more optically smooth at large angles of incidence (Chapter 6), the regular component may become especially important at large phase angles.

The most familiar example of the combination of diffuse reflection and regular reflection is water containing suspended solids, as in rivers, lakes, and oceans. If a body of water is examined in the geometry for specular reflection from the surface, a bright glare is seen, which is the reflected image of the sun. However, if the same body is examined in an off-specular configuration, it looks dark and may be colored blue, brown, or green, depending on the nature of the suspended solids.

Introduction

One of the objectives of studying a planet by reflectance is to infer certain properties of the surface by inverting the remote measurement. In the laboratory, the objective of a reflectance measurement is usually to determine the spectral absorption coefficient of the material or, at least, some quantity proportional to it, by inversion of the reflectance.

There are at least three reasons why reflectance spectroscopy is a powerful technique for measuring the characteristic absorption spectrum of a particulate material. First, the dynamic range of the measurement is extremely large. Multiple scattering amplifies the contrast within very weak absorption bands in the light transmitted through the particles, while very strong bands can be detected by anomalous dispersion in radiation reflected from the particle surfaces. Hence, the measurement of a single spectrum can give information on the spectral absorption coefficient over a range of several orders of magnitude in α. Second, sample preparation is convenient and simply requires grinding the material to the desired degree of fineness and sieving it to constrain the particle size. Third, reflectance techniques are effective in the range k ∼10−3–10−1, where both transmission- and specular-reflection techniques are very difficult. By contrast, if α(λ) is measured by transmission, the sample must be sliced into a thin section that must then be polished on both sides; also, the range by which α(λ) can vary is limited to about one order of magnitude.

Introduction

The expressions for reflectance developed in previous chapters of this book implicitly assume that the apparent surface of the particulate medium is smooth on scales large compared with the particle size. Although that assumption may be valid for surfaces in the laboratory, it is certainly not the case for planetary regoliths. In this chapter the expressions that were derived in Chapters 8–10 to describe the light scattered from a planet with a smooth surface will be modified so as to be applicable to a planet with large-scale roughness.

In calculations of this type we are immediately faced with the problem of choosing an appropriate geometric model to describe roughness. Some authors have chosen specific shapes, such as hemispherical cups (Van Diggelen, 1959; Hameen-Anttila, 1967), that approximate impact craters on the surface of a planet. However, such models may not be applicable to other geometries. To make the expressions to be derived as general as possible, it will be assumed that the surfaces are randomly rough. There is a large body of literature that treats shadowing on such surfaces — see, for example, Muhleman (1964), Wagner (1967), Saunders (1967), Hagfors (1968), Lumme and Bowell (1981), and Simpson and Tyler (1982), as well as the references cited in those papers — although many of those papers deal only with specular reflection, such as is involved in analyses of sea glitter or backscattered lunar radar signals.