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U.S. Geological Survey 

Professional Paper 1348

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COVER

The region of the Imbrium and Serenitatis basins during six episodes of 

lunar history. From lower right to upper left, the paintings depict the Moon's 

surface as it may have looked (1) late in the Nectarian Period; (2) in the Early 

Imbrian Epoch, soon after formation of the Imbrium basin; (3) late in the 

Early Imbrian Epoch; (4) in the Late Imbrian Epoch; (5) in the Eratosthenian 

Period; and (6) as it is today Paintings by Donald E. Davis, courtesy of the 

artist.
THE GEOLOGIC HISTORY 

OF THE MOON
ASTRONAUT DAVID SCOTT AT STATION 6, APOLLO 15 LANDING SITE. APOLLO 15 FRAME H-11514.
The Geologic History 

of the Moon

By DON E. WILHELMS

with sections by JOHN F. McCAULEY 

and NEWELL J. TRASK

U.S. GEOLOGICAL SURVEY PROFESSIONAL PAPER 1348

A comprehensive review of lunar science and evolution 

from the viewpoint of historical geology, based on data from 

both photogeologic observations and lunar-sample analysis

UNITED STATES GOVERNMENT PRINTING OFFICE, WASHINGTON : 1987
DEPARTMENT OF THE INTERIOR 

DONALD PAUL HODEL, Secretary

U.S. GEOLOGICAL SURVEY 

Dallas L. Peck, Director

Library of Congress Cataloging-in-Publication Data

Wilhelms, Don Edward, 1930- 

The geologic history of the Moon.

(U.S. Geological Survey Professional Paper 1348)

Bibliography: p. 283-292

Supt. of Docs, no.: I 19.16:1348

1. Lunar geology. I. McCauley, John Francis, 1932- . II. Trask, NewellJefferson, 1930- . III. Title. IV. Series: United

States. Geological Survey. Professional Paper 1348. 

QB592.W5 1987 559.9'1 86-600177

For sale by the Books and Open-File Reports Section, 

U.S. Geological Survey, Federal Center, Box 25425, Denver, CO 80225
CONTENTS

Page

Preface ........................................................... VII

Acknowledgments ................................................. VII

Abstract .......................................................... VIII

Chapter 1. General features ...................................... 1

2. The stratigraphic approach ............................ 15

3. Crater materials ...................................... 25

4. Basin materials Orientale ............................ 55

5. Mare materials ....................................... 83

6. Structure ............................................. 105

7. Relative ages ......................................... 121

8. Pre-Nectarian system ................................. 137

9. Nectarian System ..................................... 161

10. Lower Imbrian Series .................................. 193

11. Upper Imbrian Series .................................. 227

12. Eratosthenian System .................................. 247

13. Copernican System .................................... 263

14. Summary ............................................. 274

References cited .................................................. 283

Index ............................................................ 295

PLATES

[Plates are at end of book]

PLATE 1. Index to photographic coverage

2. Photographic coverage of the Moon, showing best photograph or 

set of photographs of each area

3. Geologic map of ringed basins

4. Maria

5. Structural features

6. Pre-Nectarian system

7. Nectarian System

8. Lower Imbrian Series

9. Upper Imbrian Series

10. Eratosthenian System

11. Copernican System

12. Geologic map of the present Moon

THE GEOLOGIC HISTORY OF THE MOON

By DON E. WILHELMS

PREFACE

The Moon held little interest for most scientists after its basic 

astronomic properties had been determined and before direct exploration appeared likely (Wright and others, 1963; Baldwin, 1978). 

Speculations about its internal structure, composition, and origin 

were only broadly constrained by cosmochemical data from meteorites and solar spectra, and by astronomic data about its size, shape, 

motions, and surficial properties (Urey, 1951, 1952; Kuiper, 1954). 

Most investigators who were active before the space age began in 1957 

believed that significant new advances in lunar knowledge required 

acquisition of additional data.

One analytical technique, however, was insufficiently exploited 

before the 1960's. Few scientists since the geologist Gilbert (1893) had 

studied the lunar surface systematically from the historical point of 

view. Those who did immediately obtained important new insights 

about the Moon's postaccretion evolution (Baldwin, 1949, 1963; 

Kuiper, 1959). Then, the pioneering work of E.M. Shoemaker and R.J. 

Hackman focused the powerful methods of stratigraphy on lunar 

problems (Hackman and Mason, 1961; Shoemaker, 1962a, b; Shoemaker and Hackman, 1962). Stratigraphy is the study of the spatial 

distribution, chronologic relations, and formative processes of layered 

rocks. Its application to the Moon came relatively late and met resistance, but the fundamental stratigraphic approach (Albritton, 1963) 

was, in fact, readily transferable to the partly familiar, partly exotic 

deposits visible on the lunar surface (Mutch, 1970; Wilhelms, 1970b). *

Stratigraphic methods were applied systematically during the 

1960's in a program of geologic mapping that aimed at reconstructing 

the evolution of the Moon's nearside (McCauley, 1967b; Trask, 1969, 

1972; Mutch, 1970; Wilhelms, 1970b; Wilhelms and McCauley, 1971). 

Order was discovered among the seemingly diverse and random 

landforms of the lunar surface by determining the sequence in which 

they were emplaced. The stratigraphic sequence and the emplacement processes deduced therefrom provided a framework for exploration by the Apollo program and for the task of analyzing the returned 

samples.

During the 1970's, the sophisticated labor of hundreds of analysts 

was brought to bear on the wealth of material returned by the American Apollo and the Soviet Luna spacecraft. Our present perception of 

the Moon has emerged from the interplay between sampling studies 

and stratigraphically based photogeology. These two approaches are 

complementary: Photogeology contributes a historical context by 

viewing the whole Moon from a distant vantage point, whereas the 

samples contain information on rock types and absolute ages unobtainable by remote methods. Neither approach by itself, even the 

most elaborate program of direct surface exploration, could have 

yielded the current advanced state of knowledge within the relatively 

short time of two decades (Greeley and Carr, 1976).

This volume presents a model for the geologic evolution of the 

Moon that has emerged mainly from this integration of photogeologic 

stratigraphy and sample analysis. Other aspects of the vast field of 

lunar science are discussed here only insofar as they pertain to the 

evolution of visible surface features. Chemical data obtained by

Use of the prefix "geo-" for lunar and planetary studies has been criticized, but it is justified by: (1) applicability 

to all other solid bodies of the geologic principles developed for the Earth; (2) elimination of the need for new terms 

for every new world observed at geologically useful scales, whose number now exceeds 20; (3) the Greek etymology, 

which includes the meanings "land" or "ground", and (4) two decades of usage (Shoemaker, 1962a, p. 117; Ronca, 

1965; Mutch, 1970; Wilhelms, 1970b). The prefix "seleno-" is no longer used by professional lunar scientists except in 

some terms referring to coordinates, control points, or the global figure (selenographic, selenodetic). Although 

"astrogeology" was chosen as a convenient and appropriate name (Milton, 1969) for the U.S. Geological Survey's 

branch devoted to lunar and planetary studies, "lunar (planetary) geology" is more commonly used. "Planetology" is 

a broader term that includes such nongeologic sciences as atmospherics and planetary astronomy.

remote sensing supplement the photogeologic interpretations of some 

geologic units (see chap. 5), and geophysical data obtained both from 

lunar orbit and on the surface constrain hypotheses of the origin of 

many internally generated structures and deposits. Studies of the 

same data that treat the Moon as a whole, including speculations 

about the intriguing but unsolved problem of its origin, have been 

adequately covered in other reviews (Hartmann, 1972b, 1983; Taylor, 

1975,1982; French, 1977; Wood, 1979; Basaltic Volcanism Study Project, 1981; Cadogan, 1981; Glass, 1982).

This volume is written primarily for geoscientists and other 

planetologists who have examined some aspect of lunar or planetary 

science and who want a review of lunar science from the viewpoint of 

historical geology. It should also provide a useful summary for the 

advanced student who is conversant with common geologic terms. It 

may, furthermore, interest the geologist who has not studied the 

Moon but who wishes to see how his methodology has been applied to 

another planet.

The volume's organization reflects the stratigraphic practice of 

first identifying and interpreting geologic units, then ranking them 

in chronologic sequence over the whole planet. Chapter 1 presents 

what may be considered the raw data by briefly describing the major 

surface and subsurface lunar features and defining some common 

terms. Chapter 2 explains the general philosophy of stratigraphy by 

showing how it has been applied to some critical past and contemporary lunar problems. Chapters 3 through 6 are devoted to the four 

major classes of lunar features craters, basins, maria, and tectonic 

structures with emphasis on craters and basins. These four chapters contain descriptions and interpretations that apply to the rest of 

the volume; except for some new hypotheses concerning basin origin 

and lithospheric thickness, these chapters are based mostly on existing literature and should be of use mainly to the reader unfamiliar 

with lunar geology. The rest of this volume is organized around the 

dimension of time. Chapter 7 reviews general guidelines for ranking 

lunar deposits in order of relative age. Chapters 8 through 13 trace the 

histories of the craters, basins, and igneous materials belonging to 

each of the six main divisions of the lunar stratigraphic column. 

These six chapters constitute the core of the original contribution of 

this volume and contain information for all of its intended audience. 

The volume is summarized in chapter 14 and in 12 maps of the two 

lunar hemispheres (pis. 1-12) at the end of this volume.

ACKNOWLEDGMENTS

This volume is dedicated to Eugene M. Shoemaker, who, more 

than any other individual, was responsible for the incorporation of 

geology into the American space program. First, in his study of 

Meteor Crater in Arizona, he developed a physical theory of the 

impact process on which subsequent cratering investigations have 

been based, and established the structural and petrographic indicators of impact later used on the Earth and the Moon (Shoemaker, 

1960,1963; Chao and others, 1960). With his colleague Edward C.-T. 

Chao, he applied these criteria in establishing an impact origin for 

the previously enigmatic Ries crater (Rieskessel), Germany (Shoemaker and Chao, 1961), which has become a model for larger lunar 

craters. Then, in four major papers written over two years (1960-61), 

he (1) put lunar geologic mapping on the firm stratigraphic footing 

that still supports it (Shoemaker and Hackman, 1962); (2) championed the impact interpretation of large fresh lunar craters by the 

most convincing combination of arguments yet advanced (Shoemaker, 1962b); (3) accurately appraised the objectives and future 

course of lunar exploration, including the forerunner role of the small, 

dry, and airless Moon in comparative planetology (Shoemaker,

VII
VIII THE GEOLOGIC HISTORY OF THE MOON

1962a); and (4) attacked the problem of extraterrestrial absolute ages 

by means of the impact-cratering rate on the Earth (Shoemaker and 

others, 1962a). As chairman of the Joint Working Group responsible 

for recommending a detailed program of scientific exploration (Hall, 

1977, p. 163) and by daily contacts at the U.S. National Aeronautics 

and Space Administration (NASA) headquarters in 1962 and 1963, he 

was able to implement many of his sound and visionary concepts of 

exploration strategy. He has followed up his earlier interests by 

definitive studies of the lunar regolith, primary- and secondary-crater 

populations, and the impact flux from interplanetary space (Shoemaker, 1965,1966,1971,1981; Shoemaker and others, 1967a, b, 1968, 

1969a, b, 1970, 1979; Shoemaker and Morris, 1970). He was the 

geologist on the experimenter team for the Ranger program, the 

principal experimenter for the Surveyor television experiment, and 

the leader of the field geology teams for Apollos 11,12, and 13. He has 

continued to guide and inspire those of us fortunate enough to have 

been associated with him during the development of lunar geologic 

investigations.

The manuscript was prepared with the help of numerous other 

scientists who are or were active in lunar research as members of the 

U.S. Geological Survey's Branch of Astrogeology. John F. McCauley 

contributed much of chapter 4, was codeveloper with me of many of 

the concepts of stratigraphy and basin formation embodied in this 

volume, and furnished valuable reviews of chapters 1 through 8. 

Odette B. James and Paul D. Spudis helped fill gaps left by my 

ignorance of lunar petrology and kindly supplied photographs of 

lunar rocks and thin sections. James submitted the first version of the 

section on mare-basalt samples and suggested major beneficial 

changes in all the sections on petrology. Spudis contributed drafts of 

the sections on Apollo 15 in chapter 10, read two versions of the entire 

manuscript, and contributed many valuable suggestions to every 

part. Richard J. Pike and Joseph M. Boyce contributed drafts of the 

sections on complex craters (chap. 3) and the DL method (chap. 7), 

respectively. Baerbel K. Lucchitta and David H. Scott contributed 

drafts of much of chapter 6. Perceptive reviews were furnished by 

Maurice J. Grolier and Carroll Ann Hodges of chapters 3 and 4, 

respectively. Jay L. Inge prepared the shaded-relief base for the twohemisphere maps (pis. 1-12). Donald E. Davis prepared most of the 

artistic illustrations. Gary D. Clow helped with the mathematics. The 

manuscript furthermore benefited from the first-hand experience of 

Survey geologists Gordon A. Swann, George E. Ulrich, and Edward W. 

Wolfe as leaders or members of the field geology teams for Apollos 14, 

15, 16, and 17. Non-Survey experts also generously helped with sections outside my fields of research. Discussions with Gunther W. 

Lugmair (University of California, San Diego) and S. Ross Taylor 

(Australian National University, Canberra) clarified many aspects of 

geochronology and petrology. Carle M. Pieters (Brown University) 

contributed figures and a valuable review of the section on remote 

sensing in chapter 5.

I am most deeply indebted to Albert Estrada and David B. 

Snyder. Estrada provided indispensible help in preparing the photographic illustrations and gave valuable counsel about scope. Snyder 

read two versions of the manuscript in sequence from beginning to 

end and greatly improved its clarity by his careful scrutiny and 

insight.

This research was originally commissioned by the U.S. Geological Survey. Otherwise, the Survey's involvement in the lunar 

program has been entirely funded by NASA under some 30 contracts. 

Studies under Contracts R-66, T-1167B, T-5874A, T-66353G, 

W13,130, and W13,709 led most directly to this volume. The NASA 

personnel who monitored these contracts extended understanding 

support of the geologic approach to deciphering the flood of data 

returned so spectacularly from the Moon. Special thanks go to NASA 

contract monitor Robert P. Bryson for his perception and support of 

our approach during the most active phases of the program. Preparation of this manuscript was partly supported in 1978-80 by the 

Planetary Geology Program Office of NASA, Stephen H. Dwornik and 

Joseph M. Boyce, chiefs, under Contract W13,709. In 1977-80, both 

the Lunar and Planetary Programs Office and the Planetary Geology 

Program Office of NASA supported my research on chronology and 

ringed basins that is incorporated into this volume. The Geochemistry and Geophysics Office of NASA, William L. Quaide, chief, supported the Lunar Geoscience Consortium under Contract W13,130, 

which included my study of the maria also incorporated here.

ABSTRACT

More than two decades of study have established the major 

features of lunar geologic style and history. The most numerous and 

significant landforms belong to a size-morphology series of simple 

craters, complex craters, and ringed basins that were formed by 

impacts. Each crater and basin is the source of primary ejecta and 

secondary craters that, collectively, cover the entire terra. The largest 

impacts thinned, weakened, and redistributed feldspathic terracrustal material averaging about 75 km in thickness. Relatively small 

volumes of basalt, generated by partial remelting of mantle material, 

were erupted through the thin subbasin and subcrater crust to form 

the maria that cover 16 percent of the lunar surface. Tectonism has 

modified the various stratigraphic deposits relatively little; most 

structures are confined to basins and large craters.

This general geologic style, basically simple though complex in 

detail, has persisted longer than 4 aeons (1 aeon = 109 yr). Impacts 

began to leave a visible record about 4.2 aeons ago, after the crust and 

mantle had differentiated and the crust had solidified. At least 30 

basins and 100 times that many craters larger than 30 km in diameter 

were formed before a massive impact created the Nectaris basin about 

3.92 aeons ago. Impacts continued during the ensuing Nectarian 

Period at a lesser rate, whereas volcanism left more traces than 

during pre-Nectarian time. The latest basin-forming impacts created 

the giant and still-conspicuous Imbrium and Orientale basins during 

the Early Imbrian Epoch, between 3.85 and 3.80 aeons ago. The rate 

of crater-forming impacts continued to decline during the Imbrian 

Period. Beginning in the Late Imbrian Epoch, mare-basalt flows 

remained exposed because they were no longer obscured by many 

large impacts. The Eratosthenian Period (3.2-1.1 aeons ago) and the 

Copernican Period (1.1 aeons ago to present) were times of lesser 

volcanism and a still lower, probably constant impact rate. Copernican impacts created craters whose surfaces have remained brighter 

and topographically crisper than those of the more ancient lunar 

features.
I
FIGURE 1.1. Lunar nearside photographed 4 days after full Moon from the Pic du Midi 

Observatory, France, in August 1964. Left (west) half is illuminated at high Sun elevations, 

which enhance albedo (brightness) variations, and right (east) part at lower Sun elevations, 

which enhance relief by casting shadows. Terminator (line between dark and illuminated 

areas) is at long 51 ° E. Extensive ray systems surround craters Copernicus (upper left center) 

and Tycho (near bottom). Mare Nectaris, surrounded by concentric rings, is along terminator at lower right.

Note: Before 1961, lunar directions were stated with reference to their position in the sky 

as viewed from the Earth. Mare Orientate, the "Eastern Sea," is on the left limb of the Moon 

as seen in the Northern Hemisphere, that is, near the east horizon of the Earth. In 1961, the 

"astronautical convention" was adopted in anticipation of manned spaceflight: The direction

from which the Sun rises on the Moon was henceforth called east, as it is on the Earth. The 

definition of north remained the same, but after 1961, more publications began orienting their 

figures with north at the top rather than at the bottom (as it is seen in astronomic telescopes). 

The 0° meridian is in the center of the nearside, and the diametrically opposite longitudal line 

on the farside is 180°. On most maps, longitudes increase both eastward and westward from 

the 0° meridian until meeting at the 180° meridian, which is also the terrestrial convention. 

Some lunar maps use a 360° system of longitudes, increasing eastward from the 0° central 

meridian. In this volume, the Orientate limb is considered the west, photographs are oriented 

with north at the top except as noted, nearside longitudes are less than 90°, and farside 

longitudes are greater than 90°.
1. GENERAL FEATURES

CONTENTS

Surface ... 

Subsurface

Page

3

12

SURFACE

Near full Moon, the naked eye sees a contrast between dark and 

light surfaces (of low and high albedo, respectively) that has been 

fancied as a "man in the Moon" or other configurations (fig. 1.1). In 

1609, Galileo noted that the dark spots are smooth and the brighter 

areas rugged (Whitaker, 1978). These terrain types are still designated by their 17th-century names maria (singular, mare) and terrae 

(singular, terra; commonly known as uplands or highlands; table 1.1). 

The maria constitute about 16 percent, and the terrae 84 percent, of 

the lunar surface. Maria occupy about 30 percent of the lunar nearside, the hemisphere visible from the Earth; spacecraft exploration, 

beginning with the Soviet Luna 3 in 1959 (table 1.2; Barabashov and 

others, 1961; Kopal and Mikhailov, 1962, p. 1-44; Whitaker, 1963), 

showed that they constitute only about 2 percent of the farside (figs. 

1.2-1.4). South of about lat 35° S., however, the proportions are 

reversed; the southern farside is richer in maria than is the southern 

nearside (fig. 1.5).

Most of the maria are approximately circular. The circular maria 

are bordered by annular or arcuate, commonly mountainous terra 

rims. The terra structures in which the circular maria lie are called 

ringed basins or, simply, basins (see chap. 4; table 1.1; Hartmann and

TABLE 1.1. Names of common lunar landforms

[Plural form in parentheses; the singular form is properly used adjectival ly--for example, "mare basalt" rather than "maria basalt."]

[Latin generic terms and proper names of individual features used here are approved by 

the International Astronomical Union (Arthur and others, 1963; Andersson and Whitaker, 

1982). Small or obscure craters have traditionally been designated by letters after 

the names of a nearby, more conspicuous crater (Madler system of nomenclature); following Andersson and Whitaker, these convenient lettered names are used in this volume 

in preference to some new names that were invented for use on the large-scale charts 

made from Apollo orbital data]

Latin name Common name Description

Mare (maria)---- -- Sea (not used in this Dark, smooth plains.

vol ume).

Lacus, palus, sinus -- - Small mare. 

Terra (terrae) 1 -------- Highlands, uplands, Rugged, relatively bright (high albedo)

continents. terrain. 

Mons (montes)---------- Mount, mountain(s)------ High massif(s), generally forming

arcuate ranges. 

Promontorium - Promontory----- ------- Mountains partly enclosed by mare.

Rupes-- Scarp Fault in mare or high arcuate scarp in

terra. 

Dorsum (dorsa)--------- Mare ridge, wrinkle Narrow ridge, mostly in mare.

ridge. 

Rima ( rimae)-------- - Ri 11 e- Narrow, elongate depression (sinuous,

arcuate, or straight). 

Vallis (valles)-------- Valley--------------- - Wide, elongate depression, commonly

consisting of inconspicuous craters. 

Catena (catenae)------- Chain------------ ----- Chain of distinct craters.

_.. -._. Crater- --- - --- Circular or subcircular depression,

generally bounded by a raised rim. 2 

 ._ .. _____ _-__ Basin, ringed basin Large craterlike depression containing

one or more rings in addition to a

rim. 3

1 There is no sharp distinction between an individual "terra" and the "terrae." 

"Terrestrial" refers to the planet Earth.

2 Usage has established the term "ejecta," for the material thrown out of craters, 

as a singular noun, despite its origin as a Latin plural.

3 The terms "mare" and "basin" are commonly confused. Genetically, these two features are only indirectly related (see chap. 2). Rasins are terra structures, not all 

of which contain maria. Efforts to alleviate this serious semantic confusion have led 

to use of such terms as "mountain-bordered mare," "mare basin," "dry mare," "thalassoid" 

(no lava; Lipsky, 1965), and "multiringed basin." Terms for basins that contain the 

word "mare" are unsatisfactory, and even "basin" is misleading because these structures 

are characterized by huge mountainous rings as much as by the excavated depression. 

Nevertheless, it is probably too late to coin a new term. In this volume, the term 

"basin" or "ringed basin" is used for all lunar excavations at least 300 km in diameter; 

a more exact definition has not yet been agreed upon (see chap. 4).

Kuiper, 1962). Most large mountainous rings or arcs that do not 

encompass maria are also parts of basins (figs. 1.4, 1.5). Concentric 

rings characterize all well-exposed basins.

The maria are mostly level and smooth at coarse scales but 

contain local topographic relief. Mare ridges (dorsa) are long intricate 

welts on the mare surfaces (fig. 1.6). Rilles (rimae) are narrow troughs 

or grooves much longer than they are wide; they include genetically 

distinct sinuous, arcuate, and straight varieties. Few ridges but many 

arcuate and straight rilles cut the terrae as well as the maria. Darkmantling materials are as dark as or darker than the maria but 

assume part of the topographic form of the underlying terrain (figs. 

1.6,1.7).

Craters, ranging in size from those visible with binoculars to 

micropits observed on returned samples, are ubiquitous on the lunar 

surface. At coarse scales, they are much more numerous on the terrae 

than on the maria. Most lunar craters have rims elevated above, and 

floors depressed below, the surrounding terrain. Craters smaller than 

about 16 to 21 km in diameter have relatively featureless interiors; 

larger craters are more complex, possessing central peaks, arcuate 

wall terraces, and other interior landforms (see chap. 3). Some crater 

exteriors resemble the adjacent terrain except for a short rim flank; 

others display coarse concentric structure near the rim crest, grooves 

on a lower rim flank, and numerous smaller satellitic craters, which. 

are most conspicuous between about one and three crater radii from 

the rim (fig. 1.6). Bright rays radiate hundreds or thousands of kilometers from some craters (figs. 1.1,1.6).

Most terrae appear at first glance to consist of little except large, 

randomly distributed, overlapping craters. However, regional differences in terra morphology emerge upon further inspection 

(Hackman and Mason, 1961). The central and northern nearside is 

characterized by ridges and grooves radial to the Imbrium basin, 

which contains Mare Imbrium (figs. 1.6-1.8). The terra east of the 

neighboring circular basin, Serenitatis, has a choppier, less regular 

pattern of elevations and depressions (fig. 1.7). Several regions contain particularly dense concentrations of craters that are grouped in 

chains or clusters. Concentric arcs of the Nectaris basin, which 

encompass the small Mare Nectaris, are conspicuous in the southeast 

quadrant (fig. 1.1). An even more conspicuous system of concentric 

rings, radial lineations, and satellitic craters surrounds Mare Orientale on the west limb of the Moon (fig. 1.9). Tkrra plains that are 

smooth and level like the maria, but lighter in color, occupy more 

crater floors and other depressions in and near the Imbrium- and 

Orientale-radial terrains (figs. 1.8, 1.9) than in any other region. 

These textural patterns in the Imbrium, Nectaris, and Orientale 

regions, as described in detail in this volume, play major roles in 

elucidating the geology of the Moon. They are expressions of major 

stratigraphic units and form the basis for interpreting less distinctive 

terrae.

The farside seems, at first, to be even less regularly patterned 

than the nearside. This difference is due mostly to the absence of 

maria and of the extensive circum-Imbrium radial pattern that 

characterize the nearside terrae. Later descriptions show that concentric and radial patterns also characterize the farside (fig. 1.4), 

although they are generally less extensive, less pronounced, and less 

well photographed than those of the nearside. The southern farside, 

which contains most of the farside's maria, also contains the farside's 

most conspicuous concentric rings and other noncrater morphologies 

(fig. 1.5). A major purpose of this volume is to show the basic stratigraphic order that underlies the morphologic features of the terrae.

3
THE GEOLOGIC HISTORY OF THE MOON

FIGURE 1.2. East limb (right edge of lunar disc as seen from the Earth) and adjacent part of 

farside, divided by long 90° E. Mare Crisium (C) is on nearside, and Maria Marginis (M) 

and Smythii (S) partly on farside. Maria also fill such craters as Lomonosov (L; 93 km, 27° 

N., 98°E.)andTsiolkovskiy(T; 180km, 20° S., 129° E.; partly in shadow; compare figs. 

1.3, 1.4). Farside terrane in view, which is otherwise mostly terra, includes craters Fabry 

(Fa; 179km, 43° N., 101°E.), Fleming (Fl; 130km), Hilbert(H; 170km), King(K; 77 

km), and Pasteur (P; 235 km). A large subcircular area of light-colored plains lies between 

Fleming and Lomonosov. Terminator is at long 131 ° E.; left-hand (west) edge of photograph 

is at nearly the same longitude as terminator in figure I.I. Apollo 16 frame M 3021, 

photographed by Apollo 16 mapping camera on Earthbound flight in April 1972.

Note: This volume includes photographs taken with three types of cameras carried in lunar 

orbit by Apollo spacecraft: mapping or metric (M), panoramic (P), and hand-held or 

bracket-mounted Hasselblads (H). All missions carried Hasselblads; Apollos 15 through 17 

carried both mapping and panoramic cameras as well (pi. 2; Masursky and others, 1978)i

Note: Most crater diameters and positions in this volume are from Andersson and 

Whitaker (1982). Basin diameters and most crater diameters used in frequency studies are 

from my measurements. Latitude is given before longitude throughout the volume. Coordinates differ considerably among various maps, especially on the limbs and farside. One 

degree of lunar latitude covers about 30 km.
1. GENERAL FEATURES

FIGURE 1.3. Part of southern farside centered on mare-filled crater Tsiolkovskiy (T). 

Overlaps with fig. 1.2; compare Tsiolkovskiy and Hilbert (H; 170 km, 18° S., 108° E.). 

Other craters: Langemak (L; 102 km; contains small mare patch), Fermi (F; 238 km), Milne 

(M; 262 km; fig. 1.5), Jenner 0; 72 km; fig. 1.5), Pavlov (P; 141 km), Roche (R; 146 km; 

superposed crater Pauli contains mare). PI, small ringed basin Planck (325 km, 60° S., 136°

E.; figs. 1.4, 1.5). Orbiter 3 frame M-121.

Note: Photographs transmitted from unmanned Lunar Orbiter spacecraft in 1966 and 

1967 (table 1.2; Levin and others, 1968; Mutch, 1970) are labeled as follows: "Orbiter" 

followed by mission number (1 5), M (medium resolution) or H (high resolution), and frame 

number. Best Lunar Orbiter coverage of each area is outlined on plate 2.
THE GEOLOGIC HISTORY OF THE MOON

FIGURE 1 A. Equatorial and southern farside. Most of area is heavily cratered, smoother in 

Gagarin (G; 272 km, 20° S., 149° E.) and north of Heaviside (H; 163 km, 11 ° S., 167° 

E.). A distinct linear trough is north of Keeler (K; 169 km). Arcuate massifs below and to left 

(west) of Keeler and Heaviside are parts of Keeler-Heaviside basin (KH). Maria fill 

Tsiolkovskiy (T), Mare Ingenii (MI, in Ingenii basin), Jules Verne (]\f; 134 km), part of

Poincare basin (Po; compare fig. 1.5), and other depressions. Other craters: Chaplygin (C; 

124 km) and Van de Graaff (V, double); other basins: Planck (PI, compare fig. 1.3) and 

Mendeleev (M; 320 km, 6° N., 142° E.). Mosaic of Orbiter 1 frame M-l 15 (left) and 

Orbiter 2 frame M-75 (right).

FIGURE 1.5. Southeast limb, showing concentration of maria and various ringed basins (see 

chap. 4). A, massifs of Australe basin, which enclose Mare Australe, consisting of many 

small mare matches resting in such craters as Jenner (J; compare fig. 1.3) and Lyot (L; 141 

km); An, Antoniadi crater or basin (135 km); C, outer ring of Crisium basin; M, Milne 

(compare fig. 1.3); MS, Mare Smythii in Smythii basin (compare fig. 1.2); PI, Planck basin 

(barely discernible; compare fig. 1.3); Po, Poincare basin, containing several mare patches

(compare fig. 1.4); S, Schrodinger double-ringed basin (320 km, 76° S., 134° E.). 

Indefinite basins (table 4.2) include Amundsen-Ganswindt (AG), Baimer-Kapteyn (BK), 

and Sikorsky-Rittenhouse (SR). Grooves are radial to Schrodinger and to Nectans basin 

(N; basin is outside photograph). Crater Humboldt (H; 207 km, 27° S., 81° E.) contains 

rilles and small mare patches. Orbiter 4 frame M 9.

THE GEOLOGIC HISTORY OF THE MOON

FIGURE 1.6. Telescopic view of circular Mare Imbrium (MI) and Imbrium basin. MF, 

Mare Fngoris; PP, Palus Putredinis; arrow, Apollo 15 landing site. Imbrium-basin rim is 

composed of Monies Alpes (MA I), Monies Apenninus (MAp), Monies Caucasus (MC), 

and lerra occupied by Iridum crater bounding mare fealure Sinus Iridum (SI). Lunar 

stratigraphic scheme of Shoemaker and Hackman (1962) is based on relalions among 

Imbrium basin, planar material of Apennine Bench (AB), Archimedes (left of A), mare 

material, Eraloslhenes (E), and Copernicus (C; 95 km; compare fig. 1.1; salellilic cralers are

visible easl of crater) (see chaps. 2, 7). "Wrinkle ridges" (dorsa) are above A; rilles are below 

PP. Crater pairs (see chap. 3) include Aristillus and Aulolycus (righl of A), Caroline 

Herschel and Heis (CH), Feuillee and Beer (FB), and Helicon and Leverrier (HL). Olher 

fealures: Cassini (Ca), Euler (Eu), Lambert (L), Manilius (M), Plato (PI), Pythagoras 

(Py), and Timocharis (T); many irregular craters north of Mare Frigoris; d, dark-manlled 

lerra surface. Mounl Wilson Observalory pholograph, catalog No. 257.
1. GENERAL FEATURES

FIGURE 1.7. Telescopic view centered on Mare Seremtatis (MS), overlapping area of figure 

1.6. Lineations in Monies Haemus (H) on southern part of Serenitatis rim are radial to 

Monies Apenninus; nonlineate hummocky terrain adjoins Monies Caucasus and Monies 

Alpes. Terrain easl of Mare Serenitalis consisls of irregular cralers and massifs. Olher 

fealures: Ar, Arisloleles (87 km, 50° N., 17° E.); At, Alias (87 km); E, Eudoxus; H,

Hercules; JC, Julius Caesar; M, Manilius; MH, Mare Humboldlinaum; MT, Mare 

Tranquillilalis; P, Posidonius; R, Killer (29 km, 2° N., 19° E.); RH, Rima Hyginus; 

densely cralered norlhern terrain including W. Bond (W; 158 km, 65° N., 4° E.); d, darkmantling material. Arrows, Apollo 15 and 17 landing sites. Mounl Wilson Observatory 

photograph, calalog No. 262.
10 THE GEOLOGIC HISTORY OF THE MOON

/ 'fArMK

( ^\^^m^i

»*BL.? '\;,AV«t((J?J

^:^»^
1. GENERAL FEATURES 11

FIGURE 1.9. Southwest limb (lower left as seen from northern 

hemisphere of the Earth). Arrows indicate long 90 W.; upper 

arrow on equator and lower on south pole. Conspicuous lineations are radial to Mare Orientale and ringed Orientale basin 

(centered at 20° S., 95° W.). Light-colored plains form part of 

terrain outside Orientale radials, for example, in Schiller-Zucchius basin (SZ), in central part of crater Schickard (S; 227 

km, 44° S., 55° W.), and in and near crater Wargentin (W; 84 

km). Other basins: Bailly (B; 300 km, 67° S., 69C W.), 

Grimaldi (G; partly mare-filled), Mendel-Rydberg (MR, 

barely visible), and South Pole-Aitken (SA, mountainous 

massifs); OP, part of Oceanus Procellarum in Procellarum 

basin. Footprint-shaped crater is Schiller (180 km). Orbiter 4 

frame M-180.

FIGURE 1.8. South-central nearside, including parts of "Fra Mauro peninsula" (FM) and 

Mare Nubium (MN), mare-filled crater Pitatus (Pi), fresh crater Tycho (T; 85 km, 43° S., 

11 ° W.; compare fig. 1.1), moderately fresh crater Werner (We), and north-south-trending 

"backbone" of terra, including "chain" of large craters Ptolemaeus (P; 153 km, 9° S., 2° 

W), Alphonsus (Alp), Arzachel (Ar), Purbach (Pu), Regiomontanus (R), and Walter 

(Wa, 140 km, 33 S., 1 ° E). Smooth terra plains lighter than maria fill many craters, including

those in "chain" and Albategnius (Alb), Deslandres (D), Hipparchus (H), Orontius (O), 

Playfair G (PG), and Stofler (S). Rupes Recta (RR, "Straight Wall") and grooves 

("sculpture") in upper right quadrant are radial to Imbnum basin. Apollo 14 landing site is 

just beyond north edge of photograph, above letters FM. Mount Wilson Observatory 

photograph.
12 THE GEOLOGIC HISTORY OF THE MOON 

TABLE 1.2. Spaceflights that provided lunar data

[Dates are beginnings of mission activities at the Moon. Ranger 7 impacted on August 1, 1964 (G.m.t.)]

Mission Date Type of mission Landing location or orbital parameters

Luna 2

Luna 3

Ranger 7

Ranger 8

Ranger 9

Zond 3

Luna 9

Luna 10

Surveyor 1

Lunar Orbiter 1

Luna 11

Luna 12

Lunar Orbiter 2

Luna 13

Lunar Orbiter 3

Surveyor 3

Lunar Orbiter 4

Explorer 35

Lunar Orbiter 5

Surveyor b

Surveyor 6

Surveyor 7

Zond 6

Apollo 8

Apollo 10

Apollo 11

Zond 7

Apollo 12

Apollo 13

Luna 16

Zond 8

Luna 17

Apollo 14

Apollo 15

Luna 19

Luna 20

Apollo 16

Apollo 17

Luna 21

Luna 24

Sept.

Oct.

July

Feb.

Mar.

July

Feb.

Apr.

June

Aug.

Aug.

Oct.

Nov.

Dec.

Feb.

Apr.

May

July

Aug.

Sept.

Nov.

Jan.

Nov.

Dec.

May

July

Aug.

Nov.

Apr.

Sept.

Oct.

Nov.

Feb.

July

Oct.

Feb.

Apr.

Dec.

Jan.

Aug.

1959

1959

1964

1965

1965

1965

1966

1966

1966

1966

1966

1966

1966

1966

1967

1967

1967

1967

1967

1967

1967

1968

1968

1968

1969

1969

1969

1969

1970

1970

1970

1970

1971

1971

1971

1972

1972

1972

1973

1976

Fi rst impact- - 

First unmanned flyby photography------------

First preimpact photography 

Preimpact photography 

Last preimpact photography -

Unmanned flyby photography -

First unmanned landing---- --------------

First unmanned orbital; gamma-ray 

Unmanned landing-- 

First unmanned orbital photography- -

Unmanned orbital photography 

do 

do 

Unmanned 1 and ing -

Unmanned orbital photographyUnmanned landing --- 

Unmanned orbital photography (global) 

Unmanned orbital magnetics 

Unmanned orbital photography 

Unmanned landing, first chemical analysis 

Unmanned 1 and ingFirst unmanned landing in terra 

Unmanned flyby, first returned filmFirst manned orbitalManned orbital------ ------

First manned landing 

Unmanned flyby, returned film-- 

Manned landing -

Manned flyby (aborted landing)- 

First unmanned sample return -

Unmanned flyby, returned filmLunakhod 1, first unmanned rover 

First manned landing in terra - 

Manned landing, first long mission -

Unmanned orbitalUnmanned sample return --

Manned 1 and ing-- - -

Last manned landingLunakhod 2, unmanned roverLast unmanned sample return 

Flank of crater Autolycus, lat 30° N., long 0°.

Farside and east limb; altitude, 65,000 km.

Mare Nubium (Cognitum), lat 10.6° S., long 20.7° W.

Mare Tranquil 1 itatis, lat. 2.6° N., long 24.7° E.

Floor of crater Alphonsus, lat 12.9° S., long 2.4° W.

Farside and west limb; altitude, 9,960 to 11,570 km.

Oceanus Procellarum, lat 7.1° N., long 65.4° W.

Perilune, 350 km.

Mare in crater Flamsteed P, lat 2.5° S., long 43.2° W.

Inclination, 12°; perilunes, 190 and 40 km.

Perilune, 165 km.

Perilune, 100 km.

Inclination, 12°; perilune, 50 km.

Oceanus Procellarum, lat 18.9° N., long 62.1° W.

Inclination, 21°; perilune, 55 km.

Oceanus Procellarum, lat 3.2° S., long 23.4° W.

Inclination, 85°; perilune, 2,705 km.

Perilune, 830 km (returned data until Feb. 1972).

Inclination, 85°; perilunes, 195 and 100 km.

Mare Tranquil 1 itatis, lat 1.4° N., long 23.1° E.

Sinus Medii, lat 0.5° N. , long 1.5° W.

Flank of crater Tycho, lat 40.9° S., long 11.5° W.

Altitude, >3,300 km.

Perilune, 110 km.

Best perilune, 15 km.

Mare Tranquil 1 itatis, lat 0.7° N., 23.4° E.

Southern farside and west limb; perilune, 2,200 km(?).

Oceanus Procellarum, lat 3.2° S., long 23.4° W.

Southern farside and west limb.

Mare Fecunditati s, lat 0.7° S., long 56.3° E.

Altitude, >1,120 km.

Sinus Iridum, lat 38.3° N., long 35.0° W.

Fra Mauro highlands, lat 3.7° S., long 17.5° W.

Apennine-Hadley region, lat 26.1° N., long 3.7° E.

Best perilune, 77 km(?).

Crisium-basin rim, lat 3.5° N., long 56.5° E.

Descartes highlands, lat 9.0° S., long 15.5° E.

Taurus-Littrow Valley, lat 20.2° N., long 30.8° E.

Mare in crater Le Monnier, lat 25.8° N., long 30.5° E.

Mare Crisium, lat 12.7° N. , long 62.2° E.

SUBSURFACE

Spacecraft exploration has provided some information about the 

configuration of materials below the lunar surface. The uppermost 

layers of both the maria and the terrae consist of fragmental material 

called regolith (Shoemaker and others, 1967a, 1968). Regoliths are 

generally thinner than about 5 or 6 m on the maria and thicker than 

that on the terrae (Shoemaker and Morris, 1970; Cooper and others, 

1974). They dominate the lunar scene at closeup scales (fig. 1.10) but 

are not evident in the telescopic and Lunar Orbiter photographs 

shown here, except as they affect albedo (figs. 1.1-1.9). The term "soil" 

is commonly used either as a synonym for regolith or in reference to 

its fine surficial material. Because this volume emphasizes regional 

relations, it does not discuss the regolith in detail.

The underlying bedrock of the maria is basaltic and has a density 

of 3.3 to 3.4 g/cm3 . Mare basalt typically extends hundreds of meters 

below the surface and locally reaches depths of about 5 km (see chap. 

5). In contrast, the terra bedrock is feldspathic (plagioclase-rich) 

material with a lower density estimated at 2.90 to 3.05 g/cm3 (Basaltic 

Volcanism Study Project, 1981, p. 671). Seismic data indicate that this 

terra material forms a crust 45 to 60 km thick in the west-central 

nearside (fig. 1.11; Toksoz and others, 1974; Koyama and Nakamura, 

1979) and 75 km thick in part of the southeastern nearside (beneath 

the Apollo 16 landing site; Nakamura, 1981). Extrapolations of these 

measurements to other regions are based on elevation data, estimates 

of the crustal density, and models of isostatic compensation, all of 

which are subject to further refinement (Solomon, 1978; Thurbur and 

Solomon, 1978). Where measured, average elevations are higher, relative to the Moon's center of mass, on the farside than on the nearside 

(Kaula and others, 1974; Bills and Ferrari, 1977); the crust maybe as 

thick as 120 km under some elevated terra on the farside (fig. 1.12; 

Bills and Ferrari, 1977). Most estimates of the mean crustal thickness 

fall within the range 74 ± 12 km (Kaula and others, 1974; Bills and

Ferrari, 1977; Kaula, 1977; Haines and Metzger, 1980; Basaltic Volcanism Study Project, 1981, p. 671; Taylor, 1982, p. 180-182). A mean 

thickness of 62 km corresponds to about 10 percent of the Moon's 

volume, of 74 km to about 12 percent, and of 86 km to about 14 percent 

(radius, 1,738 km).

Little is known about lunar intracrustal structure, except that 

seismic velocities appear to increase at about 20 to 25 km below the 

surface of southern Oceanus Procellarum (Toksoz and others, 1974). 

This discontinuity may indicate a change in physical state (open 

cracks above; solid, denser material below) or in chemical composition 

(Todd and others, 1973; Herzberg and Baker, 1980). At least the upper 

few kilometers of the terra crust consists of breccia, a rock type 

composed of angular fragments (clasts) set in a finer-grained matrix.

Beneath the terra crust and constituting all or most of the 

remaining lunar volume is the ultramafic lunar mantle (fig. 1.11). Its 

density is close to that of mare basalt and to the mean lunar density, 

3.34 g/cm3. Seismic data suggest that the mantle is fairly uniform at 

least to l,100±100-km depth, although small seismic-velocity 

changes have been modeled (Goins and others, 1979). Seismic shear 

waves are attenuated below 1,100 ±100 km. This central part of the 

Moon may or may not include melted zones and (or) a chemically 

distinct (metallic or sulfide-rich) core (Wiskerchen and Sonett, 1977; 

Goins and others, 1979; Taylor, 1982).

In the past, the mechanically deformable elastic lithosphere 

seems to have coincided with the terra crust, which may be considered the petrologic or chemical lithosphere (see chaps. 6,8). Today, the 

elastic lithosphere must include much of the mantle as well.

These basic facts or assumptions about the lunar subsurface are 

needed as background for later discussions of lunar tectonism and 

petrogenesis and for perspective on the overall constitution of the 

Moon. This volume, however, mostly discusses the three-dimensional 

form of the materials in the upper few kilometers or tens of kilometers 

beneath the terra and mare surfaces.
1. GENERAL FEATURES 13

FIGURE 1.10. Astronaut views of lunar surface.

A. Lunar Module and astronaut at Apollo 11 landing site. Regolith consists of loose, 

footprint-compacted material and a few rocks. Apollo 11 frame H-5931.

B. Wall of Rima Hadley (Hadley Rille) at Apollo 15 landing site, showing the only inplace 

outcrops of lunar strata visited by astronauts, overlain by thin regolith and loose boulders. 

Monies Apennmus, in distance, appear smooth because of cover of loose debris 

(compare rugged appearance of mountains in fig. 1.6). Apollo 15 frame H 12115.

90° W

EARTH

MARE 60°, W 

(BASALT)

HIGH-VELOCITY 

LAYER?,

30° W30° E

60 KM

60° E

CRUST

Shallow

moonquakes 

Negative Vs gradient

Deep 

moonquakes

100 KM

EXPLANATION

n Seismometers 

 Moonquakes

FIGURE 1.11. Two interpretations of major features of lunar crust and mantle.

A. Schematic equatorial slice through entire Moon. All known endogenic moonquakes are 

on nearside. Crust believed to be thicker on farside than on nearside. Longitudes of 

Apollo seismometers are shown. From Coins and others (1979, fig. 6b).

 Depth (km) Zone V p(km/s) (7 

I Crust 0.1-7 0.25

Upper -8 0.25 

mantle

6,000

4,000

Middle 7.5-8 0.29-0.35 < 1,500 

mantle

1,400-1,600

IV Lower 7.5-8 

 mantle

<100

V Core ? 4-6 

B

B. Similar crustal structure but different mantle structure hypothesized by Latham and 

others (1978). V' compressional-wave velocity; Vs, shear-wave velocity; CT, Poisson's 

ratio; Q , quality factor for shear waves.
14 THE GEOLOGIC HISTORY OF THE MOON

45N 45N

300 3(5 330 345 0 15 

LONGITUDE (DEGREES)

45S

45N: 45N

45S

270

LONGITUDE (DEGREES)

B

KHm CONTOUR

FIGURE 1.12. Approximate crustal thicknesses (in kilometers) on nearside (A) and farside (B). From Bills and Ferrari

(1977, fig. 3).

FIGURE 2.1 (OVERLEAF). Geologic map of part of areas covered by figures 1.6 and 1.7. Important units include: Cq, Copernican crater materials; CId, Copernican, Eratostheman, or Imbnan 

dark-mantling materials; EC, Eratosthenian crater materials; Em, Eratosthenian mare materials; Ic t and Ic2, Imbrian crater materials; Im, Imbrian mare materials; Ip, Imbnan plains materials 

(Apennine Bench Formation); lal, Alpes Formation; lap, materials of Monies Apennmus; If, FraMauro Formation; Iplt, Imbrian or pre-Imbnan materials, undivided; plr, pre-Imbnan rugged 

materials. Arrow, Apollo 15 landing site. From Wilhelms and McCauley (1971).
2. THE STRATIGRAPHIC APPROACH

CONTENTS

Page

General principles ................................................. 17

Case histories ..................................................... 17

Craters ........................................................ 17

Mare versus basin .............................................. 19

Terra materials ................................................. 19

Correlation of samples and stratigraphic units ......................... 21

GENERAL PRINCIPLES

Some earlier observers, influenced by experience with terrestrial 

geology, interpreted the Moon's surface in structural terms. Real or 

imaginary alignments of landforms were construed as faults or folds, 

and the present topography was explained as the product of progressive endogenic deformation of an originally simple surface (for 

example, Fielder, 1965). Crater and basin rims were thought to have 

been emplaced gradually along arcuate fissures. Chronologic 

sequences can be partly inferred in these structural models. A small 

crater inside a larger, or one crater rim that cuts across another, are 

signs of relative age in all but the most contrived scenarios. A sharp 

scarp that cuts a mare surface must be younger than the surface.

Our present understanding of lunar geology, however, has 

resulted from interpreting surface relations in terms of stratigraphic 

units (figs. 2.1, 2.2). Building upon centuries of thought that apparently began in 1669 with Nicolaus Steno (Woodford, 1965, p. 2-6), 

stratigraphers studying either the Earth or the Moon treat obaervations in terms of three-dimensional units of material. They do not 

consider a crater or basin as an isolated landform but as the source of 

a deposit. Similarly, a mare is not merely a surface with a certain 

color or smoothness but the upper bound of a stack of three-dimensional material layers. Each crater or mare deposit (1) is stratiform or 

tabular, (2) has a finite, normally varying thickness, (3) is laterally 

continuous over a finite area, (4) rests on other units, and (5) is 

bounded above either by additional units or by a free surface (fig. 

2.25; Shoemaker and Hackman, 1962; Mutch, 1970; Wilhelms, 1970b). 

Even topographically undistinctive terrains are composed of discrete 

rock units (fig. 2.3).

These geometric properties of discrete rock bodies imply that 

each body was formed at some instant or over some finite interval 

during the course of geologic time. Upon deposition, each crater and 

mare unit extended until it pinched out naturally or abutted against 

an obstacle. Interruptions of such depositional patterns as radiality of 

crater ejecta indicate blockage by an obstacle, superposition of a 

younger unit, or transection by a later structure (fig. 2.25). Simply 

put, younger units overlie and thus modify older units. Age relations 

can be detected as far as the units extend. For the Moon, superpositional and transectional relations can generally be seen and interpreted in temporal terms much more quickly and efficiently from a 

photograph than on the surface. These simple observations form the 

basis for any understanding of lunar geologic history.

The detection and mapping of stratigraphic units and sequences 

do not necessarily imply that the genesis of the units is known but 

only that each unit or sequence was formed by a single process or 

related processes. A crater deposit could equally well be composed of 

impact ejecta or extruded volcanic material; nonstratigraphic criteria 

may be needed to distinguish between these origins. Matters of recognition and interpretation of units are kept separate as far as possible 

in photogeologic work (Mutch, 1970; Wilhelms, 1970b, 1972b). A major 

purpose of stratigraphic studies, however, is to help determine the 

origins of the units and of their constituent rocks. The debate about

internal or endogenic origin versus impact or exogenic origins of lunar 

features is a thread running through lunar studies since their beginning.

CASE HISTORIES

Three examples illustrate the application of stratigraphic principles to problems of origin. Before direct exploration began, 

stratigraphic and theoretical studies had formulated pertinent questions and obtained many answers to genetic problems. Then, the 

major remaining problems were solved in general terms by data 

obtained directly from the Moon's surface by nine missions in the 

seven years between the flight of Apollo 11 in 1969 and the Soviet 

unmanned sampling mission Luna 24 in 1976 (table 1.2). The search 

for more specific answers, especially to the third question, the origin 

of terra materials, remains a field of active investigation. Later chapters of this volume further consider these "case histories" in a stratigraphic context.

Craters

The exogenic-endogenic controversy about the origin of craters 

probably occupied more early literature than did any other lunar 

topic (reviewed by Baldwin, 1949, 1963; Firsoff, 1961; Shoemaker, 

1962b). The only major competitor for journal space was the Moon's 

surficial layer, and even that subject was commonly discussed with 

regard to impact-versus-volcanic arguments about crater origin (for 

example, Kopal and Mikhailov, 1962, p. 371-565; Salisbury and 

Glaser, 1964). The debate continued through the era of Ranger, early 

Luna, Surveyor, and Lunar Orbiter exploration (1964-68). Caldera 

origin was favored or entertained for several types of craters and is 

still favored by a few observers (Green, 1971,1976;McCall, 1965,1980). 

However, the impact origin of large fresh craters typified by Copernicus (fig. 1.6), and of most craters that share its principal features, 

had been settled in most minds at the beginning of the space age 

(Baldwin, 1965, p. 137).

Part of the solution was stratigraphic. An origin had to be consistent both with the morphology of typical crater deposits and with the 

spatial distribution of craters superposed on other stratigraphic 

units. The exterior deposits are massive, extensive, and similar in 

lateral morphologic gradation around craters ranging over more than 

five orders of magnitude in diameter. Therefore, the crater-formation 

process (1) released enormous energies and (2) operated similarly at 

all scales. These properties characterize impacts of cosmic projectiles, 

whose approach velocities range from the escape velocity of the Moon 

(2.4 km/s) to about 70 km/s, most typically from 16 to 20 km/s (Gault, 

1974; Wetherill, 1977b). Relative to a planetary target, these velocities 

are hypervelocity, that is, greater than the speed of sound in the 

impacted medium (Baldwin, 1949, 1963; Shoemaker, 1962b). The 

hypervelocity projectiles range in size from dust particles to small 

asteroids. Thus, large primary impacts from space generate almost

17
18 THE GEOLOGIC HISTORY OF THE MOON

unlimited kinetic energies. Projectiles launched from the primary 

crater at lesser velocities (max 2.4 km/s) create morphologically varied secondary craters over great distances. The repetitive map patterns of crater deposits and the detailed morphologies of both primary 

and secondary craters (see chap. 3) match those of experimental 

impact and explosion craters much more closely than those of volcanic 

craters.

Stratigraphic relations, in combination with the properties 

expected of cosmic projectiles, also explain the spatial distribution of 

craters. The number of craters superposed on a given terrane is 

generally proportional to the age of the terrane (Gilbert, 1893; Baldwin, 1949, 1963; Opik, 1960; Shoemaker, 1962a, b; Shoemaker and

others, 1962a). For example, craters are more abundant on the older 

terrae than on the younger maria (figs. 1.6,1.7). Different ages of the 

terranes and not different origins of the craters account for this 

relation. Within a given terrane, small craters are always more abundant than large craters, the sizes and frequencies are systematically 

related, and the most conspicuous craters are randomly scattered. 

This distribution is consistent with the inverse mass-frequency distribution of observable cosmic objects (Baldwin, 1949, 1963; Opik, 

1960; Shoemaker and others, 1962b; Hartmann, 1965a, b). Nonrandom distributions, which are also observed, are equally diagnostic of 

impact origins: Secondary craters are grouped around larger primaries. Apparently nonrandom distributions of large craters, such as the

FIGURE 2.2. Stratigraphic relations of craters Delisle (De; 25 km) and Diophantus (Di; 18 km), mare materials, and Imbrium basin. Massifs of Imbrium protrude 

through mare surface as islands (1). Mare unit (2) is overlain by ejecta of Delisle (3). Another mare unit containing sinuous rille (4) embays Delisle ejecta. Secondary 

craters of Diophantus (5) are superposed on rilled mare unit and Delisle. 

A. Apollo 15 frame M-2076.
2. THE STRATIGRAPHIC APPROACH 19

north-south-trending "chain" in the south-central nearside terra (fig. 

1.8), which were ascribed to major subsurface structures in the early 

literature (for example, Fielder, 1965, p. 56), actually are coincidental 

alignments of primary craters of different ages (Baldwin, 1949, p. 

158 160). When large craters of the same age are plotted together, 

almost all lineaments disappear (pis. 6 11). Afalloff in crater density 

in the terra near the mare borders results from the inverse sizefrequency distribution: The missing craters are mantled by ejecta of 

the largest members of the impact series, the ringed basins (figs. 1.6, 

1.7). Thus, in any one epoch, the distributions of most large lunar 

craters are truly random and would require an internal generating 

process as random as primary impact a prerequisite in conflict with 

the spatial regularities commonly cited by proponents of a volcanic or 

tectonic origin.

The origin of craters having apparently atypical morphologies or 

size-frequency distributions remained to be learned by a combination 

of remotely based analysis and sampling during the 1970's. Although 

no large craters were sampled individually, the overwhelming evidence from returned rock samples is that almost all lunar craters are 

of impact origin. The large degraded craters of the terrae, which 

apparently possess only a truncated rim flank, originally resembled 

fresh craters but have been eroded by impacts and deposition of later 

ejecta (fig. 2.3; chaps. 8, 9). Continued study has uncovered impact or 

modification processes that account for most of the odd landforms 

once ascribed to volcani&m. Most elongate and irregular craters and 

clustered craters were formed by oblique or simultaneous primary or 

secondary impacts (chap. 3). Anomalous uplift of floors of impact 

craters in basins accounts for most remaining departures from the 

typical morphologies of large craters (chap. 6). Circular craters with 

dark halos were formed by impact excavation of dark materials from 

beneath lighter strata (chap. 13). A few small craters, all associated 

with mare or other dark deposits, may be endogenic (chap. 5). 

Impacts, therefore, created most lunar craters and thus are emphasized in this volume.

Mare versus basin

One of the most important products of the historical approach 

was the discovery that a mare and the basin that contains it are 

distinct features. Before the 1960's, these two features were almost 

universally thought to have been formed by the same process, either 

exogenic or endogenic; the terms "mare" and "basin" were equated. 

Even later, one might read that "Mare Imbrium" was created by a 

giant impact or that the "Imbrium basin" yielded samples of basaltic 

lava. There is conclusive stratigraphic evidence, however, that the 

mare materials are younger than the basin materials. Such craters as 

Archimedes that lie inside the Imbrium basin (figs. 1.6, 2.1) must be 

younger than that basin (Baldwin, 1949), except in the unlikely case 

that the basin rim grew up along internal ring fractures. The mare 

materials that fill Archimedes are younger still (Shoemaker and 

Hackman, 1962; Baldwin, 1963; Mutch, 1970; Wilhelms, 1970b). In

stratigraphic terms (Shoemaker and Hackman, 1962), the sequence, 

from oldest to youngest, is: (1) basin material, (2) plains material of 

the Apennine Bench, (3) deposits of Archimedes, and (4) mare material (figs. 2.1, 2.4).

Early work also established the genetic relation of the circumImbrium terrane to the Imbrium basin. Gilbert (1893) and Baldwin 

(1949,1978) perceived the radiality of the "Imbrium sculpture" system of grooves, and Shoemaker and Hackman (1962) added the recognition of craterlike Imbrium ejecta deposits (figs. 1.6-1.8). Mare 

Serenitatis fills another circular basin, which is overlain by the 

Imbrium sculpture or deposits (fig. 1.7). Yet Mare Serenitatis is 

unaffected by the Imbrium deposits, and so considerable time must 

have intervened after the Serenitatis basin formed and before Mare 

Serenitatis filled it (Baldwin, 1949, p. 210-213).

Maria and basins were, therefore, formed by different processes. 

Even before the Apollo missions, most workers had accepted these 

simple stratigraphic observations and knew that the maria are of 

volcanic and the basins of impact origin. Furthermore, the maria were 

identified as basaltic by their dark color and characteristic landforms. 

The basins were known to be exogenic by their similarity to craters 

and by the fact that only the kinetic energies of asteroidal masses 

impacting at cosmic velocities could supply the requisite energies of 

formation.

These conclusions were then confirmed by the first four Apollo 

landings between July 1969 and July 1971 (table 1.2). Apollo 11 

returned rocks with basaltic composition and unmistakable igneous 

textures. Apollo 12 sampled other basalt flows half an aeon younger 

that could not have been generated by the same event as the Apollo 11 

basalt samples. 2- 1 The first nonmare mission, Apollo 14, returned 

entirely different rock complex impact breccia from deposits of 

the Imbrium basin (fig. 2.5A). Landing in the area covered by figure 

2.1, Apollo 15 returned samples whose radiometric ages demonstrate 

a half-aeon age gap between the Imbrium basin and some of the mare 

basalt it contains (fig. 2.5B).

Terra materials

"Origin" of terra materials may mean either the chemical differentiation and igneous evolution that shortly followed the Moon's 

formation (see chap. 8), or the process that emplaced the visible 

landforms and photogeologically observed stratigraphic units. Stratigraphers, and this volume, concentrate on the timing and processes 

of the second, emplacement phase of the rocks' history.

Interpretations of emplacement processes and of relative ages 

are complementary. Recognition that the circum-Imbrium material is 

of impact origin enabled it to be used as a stratigraphic horizon;

z-! One aeon = 109 years, or 1 billion years in American usage. Ages in this volume have been recalculated using 

the radioactive-decay constants recommended in 1977 by the International Union of Geological Sciences (Steiger 

and Jager, 1977). Most of them differ numerically from the ages given in the references cited in this volume.

B. Cutaway view including geologic cross section drawn along bent line in A. Interpretations of age relations depend on interpretations of landforms as expressions of 

three-dimensional, laterally continuous units whose depositional pattern is interrupted only by blockage by older units or superposition of younger units.
20 THE GEOLOGIC HISTORY OF THE MOON

FIGURE 2.3. Vicinity of crater Tycho (T; 85 km), showing relative degradation of lunar craters. According to the principle of uniformitarianism (Albritton, 1967; Mutch, 1970; Wilhelms and 

McCauley, 1971), such other craters as Orontius (Or), Pictet (P), and Saussure (S) once resembled Tycho but have lost textural details with the passage of time. Extent of Tycho deposits and 

secondary craters suggests that similar but now-invisible deposits surround older craters and compose intercrater terrain. Orbiter 4 frame H 119.
2. THE STRATIGRAPHIC APPROACH 21

materials older and younger than the basin could be distinguished 

over much of the nearside outside the limits of the topographic basin. 

Moreover, the absolute ages of the spot samples collected by Apollo 14 

could be extrapolated over the entire mapped extent of the unit (fig. 

2.1). The samples from the Apollo 15 and 17 landing sites (figs. 2.5B,D) 

were also correlated with basins soon after the laboratory analyses 

(see chaps. 9,10).

Some materials outside basin rims, however, were less readily 

interpretable. Two morphologic units sampled by Apollo 16 at the 

fourth major terra sampling site, the Cayley Plains and Descartes 

Mountains (fig. 2.5C), have proved crucial in assessing the general 

origin of the terrae. The history of their interpretation illustrates the 

interaction between photogeologic stratigraphy and its verification 

from actual samples of the mapped units.

The Cayley and other light-colored terra plains cover about 5 

percent of the lunar terra surface and are the most distinctive terra 

landforms after the more craterlike ejecta of fresh basins (Wilhelms 

and McCauley, 1971; Howard and others, 1974). Superpositional relations and crater densities indicate that the plains-forming materials 

are older than the mare materials. Some of the photogeologic properties of the plains are marelike: They are smooth and level and are 

concentrated in depressions. In other ways, the plains-forming materials are like basin materials: They are brighter than the maria and 

may grade from thick in depressions to thin on adjacent more rugged 

terra. Accordingly, the plains-forming materials have been interpreted as both volcanic and impact deposits.

The concentration of the plains near Imbrium and their gradations with the coarse-textured basin ejecta implied lateral continuity 

of the two types of deposits and, thus, an impact-ejecta origin of the 

plains-forming material (Eggleton and Marshall, 1962; Eggleton, 

1964,1965). Similar superpositional relations of apparently isolated 

patches implied that all these patches belong to the same unit. The 

Descartes Mountains resemble coarser parts of the Imbrium ejecta 

blanket (Eggleton and Marshall, 1962).

Later, the same stratigraphic evidence was interpreted differently and helped support a revival of volcanic hypotheses (summarized by Wilhelms, 1970b, Wilhelms and McCauley, 1971, and 

Ulrich and others, 1981). The concentration of plains near basins was 

thought to result from marelike flooding of basin-related depressions, 

and the plains deposits were interpreted as younger material superposed on the blanket. Just as lateral continuity implies restricted 

time and mode of formation, its apparent absence may legitimately be 

interpreted as indicating an origin by various processes over extended 

times. Volcanic interpretations included (1) marelike materials 

brightened by longer exposure to impact cratering, or (2) materials 

more silicic than the basaltic maria. The plains-forming and gradational mantling materials were commonly thought to be facies of 

regional pyroclastic blankets, probably silicic ash-flow tuff, a highly 

fluid material that spreads widely and that partly conforms to the 

substrate (for example, Howard and Masursky, 1968; Cummings, 

1972). The hummocky Descartes materials were interpreted as volcanic on the basis of their close morphologic similarities to certain 

terrestrial landforms (Milton, 1968a; El-Baz and Roosa, 1972; Head 

and Goetz, 1972; Trask and McCauley, 1972). Volcanic interpretations 

prevailed when Apollo landing sites were chosen (Hinners, 1972). 

However, a choice between mechanisms required analysis of actual 

samples.

One of the most significant turning points in the course of lunar 

geologic thinking came in April 1972, when Apollo 16 returned samples of complex terra breccia from typical patches of plains and 

Descartes materials (fig. 2.5C; Howard and others, 1974; Ulrich and 

others, 1981). As a result, volcanic interpretations were replaced by 

impact interpretations. Regional stratigraphic relations turned out 

to be better indicators of origin than did volcanic analogs. The significance of this terra sampling transcends the revised interpretations of 

the sampled units, because the hypothesis of impact origin could also 

be extended to undistinctive, previously uninterpreted terra materials peripheral to craters and basins (fig. 2.3). Like the circumImbrium terrane, such undistinctive terranes are coeval with the 

crater or basin they surround and constitute discrete stratigraphic 

horizons; they did not form piecemeal and do not present a hopelessly 

complex problem for relative dating. More and more basins have been 

identified since 1972, and their deposits have been increasingly recognized as similar both to those of craters and to one another (pi. 3). This 

series of analogous basin deposits constitutes the main stratigraphic 

framework of the lunar terrae.

CORRELATION OF SAMPLES AND 

STRATIGRAPHIC UNITS

Return of materials from the Moon has established the general 

origins of most lunar stratigraphic units. The maria consist of 

basaltic flows and pyroclastic blankets derived by internal melting 

(fig. 2.6). The terrae consist of complex, partly shock-melted breccia 

deposits that were assembled and emplaced by impacts (figs. 2.7,2.8).

However, not all questions of origin and age have been solved by 

sampling. Especially serious is the fact that many returned samples 

have not been definitely identified with particular photogeologic 

units.

One problem is that all samples, except some of mare basalt from 

the Apollo 15 landing site (fig. 1.10B), were collected not from bedrock 

outcrops but from the overlying regoliths. Regoliths are composed of 

highly mixed materials and conceal the underlying bedrock stratigraphy and structure. Sample provenance commonly must be determined from the relative abundances of rock types at various points on 

the surface. This problem of correlating samples with source beds is 

more severe than is usually encountered on the Earth.

Second, not only the regolith but also the bedrock breccia itself is 

recycled from earlier deposits. A given sample thus may have 

acquired its chemical and textural properties during or before the 

impact. For example, the fact that the morphology of the photogeologically visible Cayley and Descartes units is gradational with 

that of the Imbrium-basin deposits does not necessarily indicate that 

the materials of those units are of Imbrium origin, because the morphology may have been imposed when the Imbrium impact reworked 

earlier materials. This situation is partly analogous to that encountered in terrestrial sedimentary conglomerate and partly unique to 

lunar terra breccia. The bedding and matrix fabric of a conglomerate 

are normally formed during sedimentation, whereas the component 

clasts are relicts of an earlier rock. The matrix textures of a lunar 

terra breccia also are generally acquired during ejection, emplacement, and subsequent cooling. The clasts, however, may have been

<o £

8 c

CD 0)

m£

EXPLANATION

Crater material

Crater material

M Mare material 

I". *£ : l Crater material 

MP-~| Plains material 

bi Pvl Basin material 

I* c » I Subcrater breccia

Ic |m ^Jc __ EC ^ ,^Xm*-^___^x<^

^^s^^^^g^ggy^l^^ ,r ^ > K'Xv x,' i i^ ~~ '£ /V( 7 v':\;-, ^N%J -S-1 \ 

X'^rATrVi-7^v* ,^T- ^»»j£?*7~ ,~,P r ~ r, ~^7 > JxT-~'<Tv~C'/^f/x~7 ^ i^^^'y^T^'^'^Vx^'^' ""'^ ^'^^'--''-^]^^l 1 -^'~- "' 

^ ''-- " N i-,Vis>V-TV"" ~ x " x -x - v v ^~ J ' v/ - " w «-/x^"

FIGURE 2.4. Generalized geologic cross section based on area of figure 2.1, showing major lunar time-stratigraphic units.
22 THE GEOLOGIC HISTORY OF THE MOON

formed either by an earlier event or in an earlier stage of the same 

event (see chap. 3); early-formed matrices may be broken up to 

become clasts in the final deposit. Intense mixing and recycling 

characterize the impact process.

Even a relatively uniform crystalline matrix of a terra breccia 

may be hard to date. It is not always clear whether impact-melted 

rocks (fig. 2.8) were heated sufficiently to reequilibrate the isotopes 

used in geochronology. Thus, the laboratory ages of the melt rocks 

may date an impact or an endogenic melting that preceded emplace-

,,- <4-v:^

. ~, ,,,-^ wl&1HL*i' ^SKt .^£?5;^lJ.7:*llW»'mr^l^P*sS

FIGURE 2.5. Settings of four latest and most elaborate Apollo sampling missions (arrows). 

Each frame is an oblique view taken by an Apollo orbital mapping camera; boom of gammaray spectrometer protrudes into two views.

A. Region of Apollo 14 landing site on Fra Mauro peninsula. Large crater in center, to left 

of boom, is Fra Mauro (95 km, 6° S., 17° W.; compare fig. 1.8); Bonpland (60 km, 

left) and Parry (48 km, right) in foreground. View northward. Apollo 16 frame 

M-1419.

B. Region of Apollo 15 landing site in Palus Putredinis near Monies Apenninus. Large 

craters at upper left are Aristillus (55 km, 34 N., 1 ° E.) and Autolycus (39 km, 31 ° N., 

1.5° E.) (compare figs. 1.6, 1.7). View northward. Apollo 15 frame M-1537.

ment of the stratigraphic unit that contains them. Substantial 

petrologic, geochemical, geochronologic, and photogeologic work is 

required to distinguish the times of origin of the constituents of a 

terra breccia.

Lunar breccia poses the additional complexity that recycled preimpact units may have been situated either in the primary target area 

of the new impact or outside that target area. Secondary impacts 

rework exterior target materials and incorporate them into new 

deposits, and the exterior and interior materials may have been 

lithologically similar before the impacts. Some investigators doubt 

whether the large volumes of melt rock found at two landing sites 

(Apollos 14, 16) are relatable to basins centered at great distances 

from these sites. Local origins in craters nearer the landing sites, 

followed by reworking during the basin impacts, are favored by many

c

C. Region of Apollo 16 landing site west of west rim of Nectaris basin (under boom;

compare fig. 1.1). Fresh crater at bottom is Descartes A (16 km, 12° S., 15° E.),

superposed on rim of crater Descartes. View westward. Apollo 16 frame M 566. 

D. Region of Apollo 17 landing site in Taurus-Littrow Valley east of Mare Serenitatis.

Large crater at top is Posidonius (95 km, 32° N., 30° E.; compare fig. 1.7). View

northward. Apollo 17 frame M-939.
2. THE STRATIGRAPHIC APPROACH 23

investigators, though not by me. Much of the uncertainty about the 

preimpact position of the constituent materials results from ignorance of the mechanics of very large impacts. Neither the size of the 

excavated part of basins nor the amount of melt they generate is 

agreed upon. These questions occupy considerable space in this volume.

In summary, the dominance of impact craters and basins on the 

Moon appears to be well established by analyses of the returned 

samples, although only a few individual units have been sampled 

directly. In the evolution of thought toward impact mechanisms, more 

and more layered rock bodies have emerged from anonymity among 

the Moon's seemingly chaotic features to take their place in the lunar 

stratigraphic column. The lithologic characterization and absolute 

ages of many of these units present more difficult problems.

B

FIGURE 2.6. Thin sections of typical mare basalt.

A. Sample 12051. Laths of plagioclase (light gray), surrounded and partly enclosed by

grains of pyroxene (medium gray) (subophitic texture). Plane-polarized light; field of

view, 2.2 mm. 

B. Sample 15538. Laths of plagioclase (long parallel structure) partly enclose pyroxene

grains poikilitically. Crossed polarizers; field of view, 2.2 mm.

Note: The Lunar Receiving Laboratory (LRL) at the U.S. National Aeronautics and 

Space Administration's (NASA) Manned Spacecraft Center (MSC), Houston, Tex., 

assigned numbers to lunar samples upon their arrival from the Moon according to the 

following scheme. The numbers contain five digits, followed by a comma and additional digits 

if the sample has been split. The first digits represent the mission number: 10, Apollo 11; 12, 

Apollo 12; 14, Apollo 14; 15, Apollo 15; 6, Apollo 16; 7, Apollo 17. In some Apollo 15 

and most Apollo 16 and 17 samples, the subsequent digit represents the station from which 

the sample was collected. For Apollo 16: 0, station 10, the region near the landed lunar 

module (LM) and the geophysical instruments (Apollo Lunar Surface Experiments Package 

[ALSEP]); 3, station 13; 7, station 11 (stations 3 and 7 had been designated in premission 

planning but were dropped; Muehlberger and others, I960). Some station numbers include 

intrastation samples as well. The last digit of Apollo 15, 16, and 17 numbers refers to the size 

of the sample: 0, unsieved material; 1,2,3, and 4, increasingly large pieces of sieved material; 

5, 6, 7, and 9, "rocks," that is, samples larger than 1 cm across. The third and fourth digits 

are complexly derived designations for specific samples defined in the lunar-sample catalogs 

prepared by the LRL; for example, the fourth digit of Apollo 16 and 1 7 numbers, if odd, 

refers to parts of large rocks and, if even, to fragments from the soil.

B

FIGURE 2.7. Clast-nch "black and white" breccia (sample 15445) from station 7, Apollo 

15 landing site, on flank of Montes Apenninus.

A. "Mug shot" made when sample first arrived at LRL from the Moon. 

B. Thin section of part of sample (15445,66), showing ragged, chaotically arranged 

plagioclase crystals and other fragments in aphanitic matrix. Plane-polarized light; field 

of view, about 2 mm.

FIGURE 2.8. Crystalline, igneous-appearing texture of impact-melt rock Oames > 1973), 

sample 14310,170 from Apollo 14 landing site. Texture is partly subophitic, like that of 

basalt sample in figure 2.6A, and partly intersertal (fine-grained minerals in interstices of 

larger plagioclase crystals).


FIGURE 3.1 (OVERLEAF). Craters on farside area centered at 20° S., 162° E., in Keeler-Heaviside basin (massifs of ring constitute horizon). Large complex crater in center is Keeler (169 km); 

smaller fresh crater on east (left) wall of Keeler is Plante (38 km); degraded crater in lower left is Stratton (71 km). Small simple craters include overlapping, aligned secondary craters between 

Keeler and Stratton, superposed on other materials. Compare figure 1.4. Apollo 12 frame H-4961.
3. CRATER MATERIALS

CONTENTS

	Page

Introduction ....................................................... 27

General features .................................................. 27

Typical morphology ............................................. 27

Secondary-impact craters ....................................... 29

Atypical craters ................................................. 32

Cratering mechanics ............................................... 33

Introduction .................................................... 33

Shock compression ............................................. 40

Cavity excavation and growth .................................... 41

Ejecta deposition ............................................... 42

Secondary cratering and ground surge ........................... 42

Deformation and nonballistic ejection ............................ 43

Peak and terrace formation ...................................... 43

Impact melting .................................................. 44

Formation times ................................................ 45

Impact breccia .................................................... 45

Terrestrial analogs .............................................. 45

Lunar-terra samples ............................................. 45

Summary of crater-material origins .................................. 47

Introduction .................................................... 47

Rim material ................................................... 47

Wall material ................................................... 48

Peak material .................................................. 49

Floor material .................................................. 50

INTRODUCTION

Stratigraphers have understandably devoted much attention to 

the deposits of craters, the most conspicuous lunar landforms (fig. 

3.1). Much of the upper lunar crust consists of interfingering beds of 

crater material. Crater deposits provide more stratigraphic datum 

horizons for reconstructing lunar geologic history than do any other 

lunar materials, and large fresh craters can be relatively dated over 

extensive areas. Interpretations of basins and terra samples of basin 

deposits depend on knowledge of their smaller relatives, craters. The 

importance of craters and their deposits requires a detailed description here of their appearance and formative processes.

To some extent, crater materials can be mapped and relatively 

dated without knowledge of their origin. They obey the same basic 

laws of sedimentation as does terrestrial sediment, despite their 

radiation from randomly distributed point sources of energy (Mutch, 

1970, p. 164). However, interpretation of the terrane beyond the 

obvious influence of a crater depends largely on knowledge of the 

crater's origin. On the Moon, the effects of an impact crater extend 

much farther than those of an endogenic crater of the same size. 

Craters lacking morphologies diagnostic of origin and age, therefore, 

have a different stratigraphic significance if they are degraded impact 

craters rather than separate genetic types. Thus, the question of 

crater origin occupies much of the next section.

Because impact is now known to be the main crater-generating 

process on the Moon, the rest of this chapter and this volume stresses 

impact craters. The section below entitled "Cratering Mechanics" 

discusses in detail how the distribution and interrelations of cratermaterial facies arose. The brecciated and melted products of impact 

are then described in general terms in preparation for later descriptions of the impact-generated geologic units that have been sampled. 

Finally, I summarize the origins of crater-material units as mapped 

geologically.

GENERAL FEATURES

Typical morphology

Individual primary-impact craters resemble one another more 

than they differ. Borrowing a term from stellar astronomy, Wilhelms 

and McCauley (1971) called the series of morphologically related craters the main sequence. The rims of these craters are nearly circular. 

Their floors lie below the level of the adjacent terrain. Their inner 

walls slope steeply, and the flank outside the raised rim crest slopes 

more gently. Systematic outward gradations of morphology reflect the 

effects of ejection and deposition of a three-dimensional, laterally 

continuous unit of inner-rim material and gradational outer deposits 

of secondary craters. Main-sequence craters are randomly scattered 

over a given terrain in numbers inversely proportional to crater 

diameter.

In general, lunar craters increase in morphologic complexity 

with increasing size (figs. 3.1, 3.2). The size-morphology series is not 

entirely gradational but undergoes a fairly abrupt discontinuity at 

crater diameters of about 16 to 21 km (fig. 3.3; Pike, 1974,1980a, b, c). 

Smaller fresh craters have simple, smooth interior profiles, smooth 

and highly circular rim crests, and depth/diameter (d/D) ratios of 

about l/5 (table 3.1). Their floors are commonly flat or gently sloping 

and are evidently composed of rubbly or fine debris accumulated from 

the walls (fig. 3.2B). Many ejecta blankets of the younger craters of 

this relatively simple type display radial textures; some have subconcentric dunelike forms (fig. 3.2A).

Larger craters are more complex (figs. 3.2C-E, 3.4). In 

unmodified form, they have one or more of the following interior 

features: (1) a broad floor that is generally level but is interrupted by 

various hills and mounds; (2) a centrally disposed hill, peak, or peak 

complex; (3) single or multiple blocks or slices of material slumped 

from the walls; (4) continuous terraces on the wall that represent

27
28 THE GEOLOGIC HISTORY OF THE MOON

FIGURE 3.2. Simple and complex lunar impact craters.

A. Linne(2.5km, 28° N., 12° E.). Interior profile is smooth except for minor rubble; ejecta

is subconcentric and dunelike. Apollo 15 frame P 9353. 

B. Taruntius H (8.5 km, 0.5° N., 50° E.). Profile is smooth except for level floor, composed

of rubble from walls; rim appears smooth, though not favorably illuminated in this

photograph. Apollo 10 frame H-4253. 

C. Arago (26 km, 6° N. 21 ° E.). Large wall terraces, evidently formed by slumping; peak

and wall merge. Apollo 10 frame H-4630.

D

D. Tycho (85 km, 43° S., 11 ° W.). Characterized by very crisp, fresh-appearing topography much more complex than that of Arago. Floor mounds are fissured. Pools of impact 

melt are superposed on terraces and rim flank; radial flow texture of interior and exterior 

melt is also visible. Concentric inner-rim texture grades to radial outer texture. Orbiter 5 

frame M-125.

E. Hausen (167 km, 66° S., 88° W.). Floor is broader relative to diameter than in Tycho; 

peak is relatively smaller, though large absolutely. Possible nnglike pattern of smaller 

peaks is visible. Terraces and hummocky wall masses are conspicuous. Secondary craters 

and herringbone pattern are conspicuous north and southeast of crater. Orbiter 4 frame 

H-193.
3. CRATER MATERIALS 

TABLE 3.1. General crater-identification criteria

[Contributions from Richard J. Pike]

29

Property Primary impact Secondary impact Endogenic

Slope of size-frequency curve (cumulative)- -1.9 for postmare craters <2 km in diameter ------

Circularity 1 - Circular, <15-20 km in diameter; crenate, >20 km in

diameter (0.70<C<0.95; median, 0.82). 

Depth - Much deeper than surrounding terrain where unmodified 

Interior profile Simple, 16-21 km in diameter; complex, >16-21 km in

diameter. 

Rim-flank profile - Rugged near crest, then concave to 2 radii ---------

Rim-flank texture - Concentric, hummocky near crest; radial to about 1

radius. 

Ejecta distribution Mostly symmetrical (but see fig. 3.11A)-------------

Mutual relations - - Interference features (fig. 3.14A) or "pushthrough"

(fig. 3.25) rare. 

Spatial distribution- Random on a geologic unit except for rare pairs or*

triplets (fig. 3.14).

- Mostly <30 km in diameter; proportional to size of source Mostly «20 km in diameter.

- -3.6 to -4.0 About -2 (Greeley and Gault, 1979).

Varies (0.35<C<0.75; median, 0.54) Varies (0.35<C<0.85; median, 0.55).

- Shallow except far from source - - Varies; floor commonly near level of surrounding

terrain. 

Featureless except where filled Varies.

- Inconspicuous except in large craters- - Varies; mostly smooth and low. 

V-shaped or linear intercrater ridges Inconspicuous.

Directed away from source- - - Widespread. 

Interference features common; downrange overlap common Varies.

(fig. 3.4F). 

Concentrated about source in clusters, chains, or loops; Amid dark material, on domes, or aligned on rilles.

commonly radial to basins; on bright rays when young.

 'Circularity (C) is defined as the ratio of the area of the inscribed circle to that of the circumscribed circle (fitted to planimetric outline of rim crest).

wholesale circumferential failure of the rim, as opposed to the blocks 

or slumps or to the minor debris wasting seen in simple craters; and 

(5) a d/D ratio that varies with diameter, from about l/5 for small 

complex craters to about lAo for the largest (fig. 3.3; Pike, 1980c). Their 

rims are scalloped or irregular, though still more or less radially 

symmetrical; some of the scallops are the source areas of the wall 

blocks or terraces (fig. 3.2C).

Exterior features are also better developed in large than in small 

craters. Rim topography adjacent to the crest and out to about half a 

crater radius is elevated, rugged, and commonly concentrically structured (figs. 3.2D, E\ Ejecta is lower and more radially structured 

beyond this rugged collar. Between one and two radii from the crest, 

the radial pattern passes into a zone dominated by negative landforms, the secondary craters. Whereas the interior features of small 

and large craters differ in origin, the exterior features of all craters 

are similar in origin. Following Dence (1964,1965), simple and complex are used here as technical terms for the two size-related morphologic classes.

Secondary-impact craters

Secondary-impact craters differ in most respects from their parents (table 3.1). Sizes are controlled not by the nearly unrestricted 

masses and kinetic energies of cosmic fragments but by the size of the 

primary crater and the 2.4-km/s lunar escape velocity, above which 

the ejected fragments would leave the Moon. Because of the lower 

impact velocities, rim-crest circularity is less commonly developed 

than in hypervelocity primaries (fig. 3.4). However, circular secondaries do form at large distances from their sources (fig. 3.5) and are 

difficult to distinguish from primaries if not clustered. Because the 

ejecta projectiles that form secondaries are larger relative to crater 

size than the hypervelocity cosmic projectiles, irregularities in projectile shape are more manifest in secondaries than in primaries. In 

addition, unbonded debris may create secondary craters (Schultz and 

Mendenhall, 1979). Most interior profiles of secondaries are as smooth

to . cc 4

I

LU

Q 

DC

QC 

O

0.6

I I I I

10 30 

CRATER DIAMETER , IN KILOMETERS

100

FIGURE 3.3. Depth-diameter (d/D) ratios of simple craters (steep slopes, left) and complex 

craters (shallow slopes, right), including 1 36 craters on lunar terrae (uplands) and 203 craters 

on lunar maria; craters 4.2 to 95 km in diameter are plotted. Simple-to-complex transition 

occurs at about 21 -km diameter in terrae and at 16-km diameter in maria. Complex craters 

show greater differences in d/D ratio in the two substrates than do simple craters. Square and 

crosses denote craters with transitional morphologies. From Pike (I980a, fig. 9).

as or smoother than those of small primaries, but shallower (Pike and 

Wilhelms, 1978). Ejected blocks are uncommon, and the exterior 

textures of individual secondaries are also smoother than those of 

primaries. Ejecta of grouped secondaries, however, may be texturally 

complex (figs. 3AC-F, 3.6).

Spatial grouping is the main difference from primaries and is the 

main diagnostic characteristic of secondary-impact craters. Whereas 

primaries are randomly grouped, secondaries are highly concentrated. Secondaries generally occur in linear or curving chains or in 

patches and clusters. Only a few of the farflung projectiles may 

separate enough to form seemingly randomly scattered secondaries.

Secondaries may have revealed more about the cratering process 

than have the primary craters. Early investigators equipped with 

good photographs or observing the Moon visually with telescopes 

were impressed by the myriad small craters that are satellitic to large 

craters of the Copernicus type (fig. 3.4). What remains the most 

convincing set of arguments for the impact origin of both the satellitic 

and the primary craters was assembled at the beginning of the space 

age by Shoemaker (1962b). On the basis of an excellent telescopic 

photograph (fig. 3.4A), he mapped the satellitic craters of Copernicus 

(fig. 3.4B) and successfully accounted for their pattern by cratering 

and ballistic theory. He showed that the chains, loops, and clusters 

were probably excavated by secondary impacts of ejecta derived from 

certain structures in the bedrock struck by the Copernicus primary 

impact. His analysis was supported afterward by field study of ray 

loops formed from an identifiable bed in the target rock of one of the 

fresh Henbury meteorite craters in Australia (Milton and Michel, 

1965), by laboratory experiments (Gault and others, 1968b), and by 

examination of high-velocity to hypervelocity missile-impact craters 

in natural materials (Moore, 1971,1976).

Although primary impact of an object from space followed by 

secondary impact of the resulting ejecta was shown to be consistent 

with the lunar and terrestrial patterns of satellitic craters, other 

mechanisms continued to be invoked. An explosive origin by the 

sudden release of accumulated volcanic gases could theoretically 

explain the patterns except for the enormous energy required, estimated for Copernicus by Shoemaker (1962b, p. 333) at 7.5 x 1021 J. 

Such energies could not accumulate in a planetary crust because the 

weak rocks could not contain them without premature release (Taylor, 

1982, p. 63). Endogenic mechanisms fail abjectly to explain the ray 

pattern of Tycho, part of which extends to the limbs of the Moon (fig. 

3.6) and thus would require enormous internal energies or globalscale fault or fissure systems centered about a point. One advocate of 

such structural systems (Alter, 1963) realized that an impact must 

have formed Tycho but postulated that the secondaries formed endogenically along impact-opened cracks, because the rays are not 

exactly radial.

This offcenter relation of the rays was among the first arguments 

for internal origin to be finally refuted by Lunar Orbiter photography 

in 1967. Shoemaker (1962b, 1964; Shoemaker and Hackman, 1962, p. 

290), Baldwin (1963, p. 355), the "endogenist" Firsoff (1961), and the 

impact-plus-endogeny proponent Alter (1963) had all noted that the 

nonradial rays consist of elements which are individually radial to 

the primary crater and which commonly originate at secondary craters (fig. 3.4A). Orbiter photographs show that the ray elements 

coincide with ridged ejecta of the secondary craters that was cast away 

from the primary crater; the secondary-crater ejecta has a herringbone or bird's-foot pattern (figs. 3AC-F, 3.6). Septa divide many 

crater pairs, as they do several of the Henbury craters (Milton and 

Michel, 1965; Milton, 1968b). In a major advance, the ridges, herringbone pattern, septa, and even domelike features were closely
*

locharis/

Kepler

Pytheas -

Eratosthenes

Copernicus " ' '-.-v' r~

Lansberg

o

Lalande J

Mosting 

N ' - '

FIGURE 3.4. Copernicus (93 km, 10° N., 20° W.), the classic locality for study of 

secondary-impact craters. 

A. Highly detailed telescopic photograph taken in 1929 by F.G. Pease with the 100-inch

telescope of Mount Wilson Observatory. 

B. Map of secondary craters of Copernicus, constructed in 1960 by Shoemaker (1962b) on

basis of photograph in figure 3.4A. 

C. Lunar Orbiter photograph of Copernicus and northeast exterior, showing rays and

subconcentric chains of secondary craters (Rimae Stadius). Orbiter 4 frame H 121. 

D. Copernicus on horizon (arrow), Rima Stadius I (S), and other secondaries superposed on

Mare Imbrium near crater Pytheas (P; 20 km). V-shaped ejecta is downrange from

Copernicus at each cluster of secondaries. View southward. Apollo 17 frame M 2444. 

E. Detail of Rima Stadius I, showing V-shaped herringbone pattern. Orbiter 5 frame

M-144.
3. CRATER MATERIALS 31

imitated in laboratory experiments (fig. 3.7) by Oberbeck and Morrison (1973a, b, 1974). The intersection angles of the lunar V-shaped 

ridges were fully modeled by various spacings and timings of nearly 

simultaneous artificial impacts that caused cones of crater ejecta to 

interact complexly upon collision. Ironically, some of the grossly offcenter satellitic chains that have the classic herringbone pattern 

modeled by Oberbeck and Morrison are the Stadius chains east of 

Copernicus (figs. 3AC-E), which were interpreted by some of the 

most astute proponents of lunar impact before the Orbiter photography as volcanic (Shoemaker, 1962b, p. 302; Shoemaker and Hackman, 

1962, p. 298; Baldwin, 1963, p. 378; Schmitt and others, 1967).

FIGURE 3.5. Group of small circular craters (left center) in crater Gagarin (272 km, 20° S., 

149° E.; compare fig. 1.4) are secondary to a distant crater. Largest crater superposed on 

Gagarin nm is Raspletin (R; 49 km), presumably a primary. Apollo 15 frame M 293.

F. Southeast sector of Copernicus. Conspicuous secondaries with herringbone pattern lie mostly beyond (to right of) one crater radius (white dots). Superposed crater in center is Copernicus H (4.3

km; see chap. 13). Orbiter 5 frame M-147.
32 THE GEOLOGIC HISTORY OF THE MOON

Only the most obdurate endogenists (Green, 1971; McCall, 1980) 

could still believe in the volcanic origin of such craters as Copernicus 

and Tycho or their satellitic craters after 1967. Secondary-impact 

origin was also readily extrapolated to larger-scale associations, for 

example, the 180-km-diameter primary crater Petavius and its retinue of small craters (Hodges, 1973b). For larger craters, volcanic and 

tectonic interpretations continued to appear for several years. A 

secondary-impact origin was recognized for many, but not all, satellitic craters (max 10 km diam) of the 260-km-diameter Iridum crater 

(fig. 3.8; Ulrich, 1969; Scott and Eggleton, 1973). The satellitic craters 

of the Orientale basin (fig. 3.9) were recognized as secondary-impact 

craters by some geologists (Offield, 1971; Wilhelms and McCauley, 

1971) but were interpreted endogenically by others well into the 

period of Apollo exploration (for example, Karlstrom, 1974). Internal 

origins were widely favored for seemingly noncompound craters 

larger than 5 km in diameter whose rims have such irregularities as 

straight segments or reentrants (fig. 3.10C; Wilhelms and McCauley, 

1971).

Several systematic relations support the secondary-impact 

origin for all these satellitic craters, of all sizes. The size ratios and 

spatial relations of secondaries to primaries remain much the same 

around primaries ranging in diameter from 1 km (Oberbeck and 

others, 1974) to more than 1,000 km (ringed basins; Wilhelms, 1976; 

Wilhelms and others, 1978). The ratios of the largest satellitic-crater 

diameters to the primary-crater diameter decrease relatively little, 

from about 0.05 for 100-km-diameter primaries (Shoemaker, 1965, p. 

121), through 0.04 for the Iridum crater (Scott and Eggleton, 1973), to 

about 0.02 for large basins (Wilhelms and others, 1978). Size-frequency distributions of satellitic craters also are remarkably similar 

over a wide size range; cumulative plots of craters satellitic to nuclearexplosion craters, large lunar primary craters (Shoemaker, 1965), and 

basins (Wilhelms and others, 1978) all slope between about 3.6 and 

-4.0. These slopes, which are much steeper than the -1.8 to -2 

typical of primary craters, confirm the visual impression that secondaries in a given cluster are more nearly equal in size than are 

primaries in any given population. Over the entire size range, satellitic craters are concentrated at distances of one to two diameters 

from the primary's center (one to three radii from the rim crest; figs.

FIGURE 3.6. Cluster of crisp-textured secondary craters of Tycho (on east rim of Ptolemaeus 

at 9.2° S., 1.0° E.). Fine herringbone pattern is evident above directional arrows. From 

Lucchitta (1977a). Apollo 16 frame P-4653.

3.4, 3.8-3.10). Even the detailed map patterns of satellitic-crater 

fields of craters and basins are similar; chains, clusters, and loops are 

as characteristic of basin secondaries as they are of crater secondaries 

(Offield, 1971; Stuart-Alexander, 1971; Wilhelms and McCauley, 1971; 

Stuart-Alexander and Tabor, 1972; Scott, 1972b; Hodges, 1973a, b; 

Saunders and Wilhelms, 1974; Oberbeck and others, 1975; Schultz, 

1976b, p. 276; Wilhelms, 1976; Ulrich and others, 1981, pi. 12). Patterns 

of ridges between and distal to clustered basin secondaries resemble 

those of crater secondaries and those formed in laboratory experiments (figs. 3.4, 3.7; Oberbeck and Morrison, 1973a, b, 1974, 1976; 

Oberbeck and others, 1975; Wilhelms, 1976). The only significant 

difference is a greater radiality of many basin-secondary groups; low 

impact velocities and grazing impacts create groovelike crater chains 

near basins (figs. 3.9A, 3.1QA).

The accumulated evidence on distribution and morphology (table 

3.1) leaves little doubt that secondary craters compose a large percentage of lunar craters. They probably outnumber primaries at diameters smaller than 20 km and also occur with diameters of at least 30 

km (Wilhelms and others, 1978).

Atypical craters

Continued research has expanded the types of morphology 

ascribable to impact. Although hypervelocity impacts normally create 

circular craters, impacts at angles less than 10° in weak materials or 

about 30° for certain combinations of target material, projectile material, and velocity may generate noncircular craters (Gault and 

Wedekind, 1978). Elongate craters, such as Messier and Schiller (fig. 

3.11), have been interpreted as volcanic or volcanotectonic. However, 

craters formed by artificial oblique impacts mimic their shapes (figs. 

3.12,3.13; Moore, 1976; Gault and Wedekind, 1978). Ejecta symmetry 

in these experiments was typically bilateral; ejecta was concentrated 

in lateral or downtrajectory directions, as is the Messier ejecta (fig. 

3.1LA). The Messier ejecta possesses such typical impact characteristics as radial ridges, secondary craters, and rays. The position of 

Schiller along a basin ring (fig. 3.1 IB) was considered as supporting 

the endogenic interpretation (Offield, 1971; Schultz, 1976b, p. 20). 

Schiller, however, consists of overlapping elliptical craters that could 

have been created by oblique, nearly simultaneous impact of a fragmented projectile or by a very low-angle impact of the type simulated 

by Gault and Wedekind (top center, fig. 3.12).

Similarities of neighboring craters have long impressed scrutinizers of the lunar surface (for example, Baldwin, 1963, p. 189; see 

discussion by several observers in Hess and others, 1966, p. 308-309). 

Examples of pairs include Sabine and Ritter, Messier and Messier A, 

Heis and Caroline Herschel, Helicon and Le Verrier, and Atlas and 

Hercules (figs. 1.6,1.7,3.1LA, 3.14). Some of these pairs may be accidental Atlas and Hercules probably differ in age but too many pairs 

exist to be entirely coincidental. Although endogeny was commonly 

invoked (for example, De Hon, 1971), diagnostic features show that an 

impact origin is more likely. Straight septa dividing some of these 

pairs (fig. 3.14A) had been considered evidence for a volcanic origin 

until they also were found at Henbury (Milton and Michel, 1965; 

Milton, 1968b) and among secondary-impact craters (Oberbeck, 

1971b). Many of the craters are too large to be secondaries. Pairs of 

large impact craters, such as East and West Clearwater, Quebec, also 

occur on the Earth (for example, Dence, 1964,1965; Oberbeck, 1971b).

Two primary-impact mechanisms have been proposed to explain 

the lunar groupings. First, Sekiguchi (1970) showed that tidal forces 

may break up weak approaching bodies before they impact. Second, 

small bodies may orbit mutually in space. The existence of these 

miniature planetary systems was surmised by Baldwin (1963, p. 21, 

189) and may have been substantiated astronomically (for example, 

Binzel and Van Flandern, 1979).

Another atypical class of craters are those called smooth-rimmed 

(fig. 3.15; Wilhelms and McCauley, 1971). They lack the rough rim 

texture of main-sequence craters and have been considered to be of 

nonimpact origin. Apollo astronauts called them delta-rim craters 

because of their equal exterior and interior slopes another departure from the main sequence, with its shallow outer and steeper inner 

slopes. They were targeted for special attention during the Apollo 

orbital missions because they were widely hypothesized to be calderas (El-Baz and others, 1972; Evans and El-Baz, 1973). Most 

smooth-rimmed craters are about 20 to 40 km across and occur near
3. CRATER MATERIALS 33

the borders of maria, a reasonable site for volcanism. The possible 

significance of smooth- rimmed craters was brought home when 

Lunar Orbiter 4 photographed two craters inside the Orientale basin 

that are, therefore, of the same maximum age (fig. 3.15A; McCauley, 

1968). Maunder is a typical fresh impact crater, complete with high 

and rough inner rim, lower and radially textured outer rim flank, 

deep floor, rugged central peak, wall terraces, and secondary craters. 

The neighboring crater Kopff is opposite in each of these properties; it 

has a smooth "delta" rim, elevated floor, no peak or terraces, and no 

obvious secondary craters.

Although the smooth-rimmed craters are atypical, they also are 

probably of impact origin. Amultivariate analysis of 11 pairs of dimensions of terrestrial and lunar craters shows that they group with 

impact craters (Pike, 1980c). Kopff may have been formed by an 

impact in a soft substrate (Wilhelms and McCauley, 1971; Guest and 

Greeley, 1977, p. 115, 153), or it may be merely a premare "hybrid" 

crater whose rim was smoothed by volcanism and whose floor was 

uplifted before and after mare flooding. Other smooth rims and 

elevated floors are consistent with similar modifications or with a 

secondary-impact origin (fig. 3.15). The more typical crater Maunder 

is simply younger than the smooth-rimmed craters and is not affected 

by either volcanism or floor uplift.

In summary, we are left with impacts as the generators of most 

lunar craters. Chapter 5 describes the relatively uncommon endogenic craters, and chapter 6 the modification of impact craters by floor 

uplift. With these relatively few exceptions and a few anomalous 

groups or individual craters (fig. 3.16), most features of lunar craters 

are compatible with current impact theory.

CRATERING MECHANICS

Introduction

The basic processes that form simple impact craters are now 

relatively well understood after two decades of intensive research on 

laboratory impact craters, explosion craters, and natural terrestrial 

and lunar craters. This section describes an idealized sequence of 

events based on studies of these relatively simple craters. It adds 

interpretations of complex craters and suggests some of the difficulties that confront investigators trying to understand craters and 

ringed basins more fully. The discussion draws heavily on earlier 

summary works by Shoemaker (1960,1962b, 1963), Moore and others

I I I

km

FIGURE 3.7. Craters formed by near-simultaneous impacts in the 

laboratory, compared with morphologically similar secondary clusters 

of Copernicus (from Oberbeck and Morrison, 1973, p. 32-25). In 

each pair of examples, laboratory craters are on left, and lunar craters 

on right. V, experimental impact velocity; S/D, ratio of separation 

between impact points to average crater diameter; 6, impact angle of 

incidence measured from normal to surface; a, angle between crater 

axis of symmetry and flightline or radial line from Copernicus. Projectiles impacted from direction below photographs. Lunar photographs, 

Apollo 15 frame M-1699 (upper left) and Orbiter 4 frame H-121 

(all others) (compare fig. 3.4C).

V = 0.82 km/s 

S/D = 0.26 

6 = 82.5°, a = 0°

Set A

V = 0.74 km/s 

S/D = 0.01 

6 = 60°, a = 0°

SetB

V = 0.77 km/s 

S/D = 0.89 

6 = 75°, a = 35°

SetC

V = 0.77 km/s 

S/D = 0.90 

6 = 75°. a = 69°

SetD
34 THE GEOLOGIC HISTORY OF THE MOON
3. CRATER MATERIALS 35

B

FIGURE 3.9. Diverse Orientale-basin secondary craters north of the basin.

A. Chain (Vallis Bohr, arrow at bottom), clustered circular craters, fissured crater floors (f), 

and herringbone pattern (h, near top), all formed by impact of Orientale ejecta. Head of 

arrow marks one basin diameter (930 km) from basin center (one radius from rim). Thick 

ejecta plains form scarp at p (see chap. 4). Large crater occupying most of photograph 

width is Einstein (170 km, 17° N., 88.5° W.). Orbiter 4 frame H-188. 

B. Domelike interference feature and radial ejecta of crater Struve L (above scale bar; 15 

km, 21° N., 76° W.), 1,350 km north-northeast of Orientale center. Orbiter 4 frame 

H-174.

FIGURE 3.8. Iridum crater (260 km; compare fig. 1.6) enclosing Sinus Iridum (SI). Almost entire terra in scene is blanketed by Iridum deposits or secondary craters; secondaries appear at about 

one radius from rim of Iridum crater. Superposed craters are Bianchini (B; 38 km), Mairan (white M; 40 km, 42° N., 43° W.), and Sharp (S; 40 km); pre-Iridum craters are La Condamine (L; 

37 km, 53° N., 28° W.) and Maupertuis (black M; 46 km). Mosaic of Orbiter 4 frames H-139, H-145, and H-151 (from right to left).
36 THE GEOLOGIC HISTORY OF THE MOON

FIGURE 3.10. Secondary craters of Imbrium and Nectaris basins.

A. Regional view, showing distance of one basin radius (dashes, 580 km) from Imbrium rim (MA, Monies Apenninus); dotted outline, Nectaris rim. Locations of figures 3.1 OB, C, E and 

3.16A, B are outlined; figure 3. IOD lies below area of photograph. Orbiter 4 frame M-108.
3. CRATER MATERIALS 37

B. "Imbrium sculpture," evidently composed of coalescing elliptical craters (sc). Large

crater is Hipparchus (151 km, 6° S., 5° E.). Gamma-ray spectrometer boom protrudes

from right. Apollo 16 frame M-839. 

C. Crater Catharina D (CD; 9 km, 19° S., 21 ° E.) has been considered volcanic (Wilhelms

and McCauley, 1971), but is oriented and situated properly (fig. 3.10A) to be an

Imbnum-basin secondary. Orbiter 4 frame H 84.

20km

D. S-shaped group of Imbrium secondaries looping southward from rim of crater Riccius 

(R; 71 km, 37° S., 27° E.) to crater Nicolai (N; 42 km, 42° S., 26° E.) and back 

toward basin in another loop ending at crater Barocius G (B; 27 km). Arrow indicates 

direction to Nectaris basin; partly filled craters south of (below) Nicolai and west (left) of 

arrow are probably secondaries of Nectaris. Large crater Janssen (190 km; compare fig. 

9.2) is partly visible in lower right. Orbiter 4 frames H-83 (right) and H-88 (left).

E. Complex morphology of Delaunay group of craters, suggesting volcanism or interference 

of several large basin-secondary craters (Holt, 1974). View centered at 22.5° S., 3° E. 

Orbiter 4 frame H-101.
38 THE GEOLOGIC HISTORY OF THE MOON

B

FIGURE 3. 11. Irregular primary-impact craters.

A. Messier (right) and Messier A (2° S., 47° E.). Radial ejecta and rays are north and 

south of Messier, and long double rays to west (down trajectory). Apollo H-frame 

(number unknown).

B. Schiller (S, footprint-shaped crater; 180-km long axis), superposed on ring of doubleringed Schiller-Zucchius basin (compare fig. 1.9); Z, crater Zucchius (64 km, 61 ° S., 

50° W.). Orbiter 4 frame M-155.

IMPACT 

ANGLE

2.12<

SOLID ALUMINUM 

6.35 mm diameter

HOLLOW ALUMINUM 

9.53 mm diameter

PYREX 

6.35 diameter

4.75(

7.33(

FIGURE 3. 12. Craters formed by oblique impacts in laboratory (from Gault and Wedekind, 1978, fig. 4). Spherical projectiles are described at top; target is 

noncohesive quartz sand. Impact velocity, about 1.7 km/s; impact angles, above plane of target surface; trajectories, from left.
3. CRATER MATERIALS 39

\

td

^_

ts/tm

l\ t\I \

r<J / .' x v //

'*

^>

Crater I' 

\ A Crater rim

ts/tm

0 5 10 METERS

*

\/

I/

FIGURE 3. 13. Ray patterns formed by oblique impacts (from right).

A. In laboratory (Gault and Wedekind, 1978, p. 3855). Impact angle, 5° above horizontal. 

B. By ballistic missile at White Sands Missile Range, N. Mex. (Moore, 1976, p. B8). 

Material units: te, thick ejecta; td, thin to discontinuous ejecta; ts/tm, scattered ejecta and 

target material. Arrow indicates missile trajectory; impact angle, about 45°.

FIGURE 3. 14. Probable pairs of primary-impact craters.

A. Bessarion B in Oceanus Procellarum (largest member of pair 12 km perpendicular to

axis of pair; 17° N., 42° W.). Simultaneous impact is indicated by septum. Orbiter 4

frame H-144. 

B. Van de Graaff (230-km long axis, 27° S., 172° E.; compare fig. 1.4). Younger complex

crater is superposed at upper left. View southward. Apollo 17 frame H-22959. 

C. Ritter (left, 29 km) and Sabine (right, 30 km, 1° N., 20° E.), classic lunar "calderas" at

edge of Mare Tranquillitatis near parallel Rimae Hypatia. Ranger 8 frame A 34.
40 THE GEOLOGIC HISTORY OF THE MOON

(1961), Baldwin (1963), Gault and others (1968b), Gault (1974), Oberbeck (1975), Roddy and others (1977), and Melosh (1980). Much of the 

material in the volume edited by Roddy and others (1977) was summarized by Cooper (1977).

Shock compression

An impact crater results from the meeting of an irresistible force 

with an immovable object (Baldwin, 1963, p. 6). A hypervelocity collision generates intense high-pressure shock waves that propagate into 

both the target and the projectile (Shoemaker, 1960,1962b; Melosh, 

1980). As the shock front moves downward and outward into the 

target, masses are set into motion with particle velocities much 

greater than the speed of sound in the various materials. Pressures 

and energies within the shock wave are commonly so great that parts

of both the target and projectile are melted and vaporized around the 

impact zone (Gault and Heitowit, 1963). More significantly for the 

ultimate crater, the shock wave strongly compresses and energizes a 

mass of target material much greater than that of the projectile 

(Shoemaker, 1962b, p. 317; Gault and others, 1968b). Because peak 

pressures in the shock wave are orders of magnitude above the 

strengths of all rock materials, the materials flow hydrodynamically 

(Shoemaker, 1960; Gault and others, 1968b; Dence and others, 1977; 

Roddy, 1977). In homogeneous targets, shock waves propagate outward in approximately spherical fronts (fig. 3.17; Gault and others, 

1968b). The pressures in the shock wave quickly diminish radially 

outward, until eventually the shock wave decays into an elastic wave 

(Shoemaker, 1960). The distance at which the shock wave becomes 

elastic is a function of original kinetic energy, projectile penetration 

depth, duration of contact, rock properties, and deflections due to 

layering and other inhomogeneities (Shoemaker, 1962b, p. 320).

B

FIGURE 3.15. Smooth-rimmed class of craters (Wilhelms and McCauley, 1971).

A. Contrasting craters in Orientale basin: Kopff, right (42 km, 17° S., 89.5° W.), a 

smooth-rimmed crater; Maunder, left (55 km, 15° S., 94° W.), a typical impact crater. 

Orbiter 4 frames H-187 (right) and H-195 (left).

B. Crozier-McClure group on Fecunditatis-basin rim (14° S., 51 ° E; each of three central 

craters, 21 24 km across). Clustering and smooth rims suggested caldera origin 

(Wilhelms and McCauley, 1971), but radial orientation and size range are also consistent

D E

with Imbrium-secondary origin. Orbiter 4 frame H 60. 

C. Gambart (25 km, 1 ° N., 15° W.). Dark rim material is probably volcanic, but current

interpretations favor superposed pyroclastic material rather than volcanic ejecta of crater.

Orbiter 4 frame H-120.

D. Lassell (23 km, 15.5° S., 8° W.) in eastern Mare Nubium. Orbiter 4 frame H-113. 

E. Daniell (29 km, 35° N., 31 ° E.) in Lacus Somniorum. Volcanic origin is suggested, but

not proved, by irregular rim crest, elevated floor, and mare fill. Orbiter 4 frame H 79.
3. CRATER MATERIALS 41

Cavity excavation and growth

Very early in the sequence of events, even before the shock wave 

reaches the projectile's trailing edge, small amounts of both the projectile and target material may be jetted from the sides of the impact 

zone at velocities that may exceed the initial velocity of the projectile 

(fig. 3.17A; Gault and others, 1963,1968b; Kieffer and Simonds, 1980). 

Most ejection, however, takes place at velocities first comparable to 

and then much lower than the initial projectile velocity. The agents of 

excavation in this main cratering-flow stage of ejection and cavity 

growth are rarefactions set up when the shock wave intersects the 

free surfaces and other discontinuities in the target and projectile. 

The result is a sudden decompression. Particles of material are 

deflected from the initially radial motions induced by the shock wave 

into upward and outward trajectories curving back toward the surface 

on the heels of the expanding shock wave (figs. 3.17B-E; Shoemaker, 

1960,1962b, p. 320; Gault and others, 1968b; Cooper, 1977, p. 37-38; 

Grieve and others, 1977, p. 801-804; Kreyenhagen and Schuster, 1977;

FIGURE 3.16. Craters of undetermined origin (see fig. 3. IDA for locations).

A. Miiller group on east of rim of Ptolemaeus; largest crater is Miiller (22 km, 8° S., 2° E.). 

Large craters exhibit morphology, size, and distal overlap typical of Imbrium-secondary 

craters at this radial distance (800 km from Monies Apenninus); "Imbrium sculpture" 

above and below chain is radial to Imbrium. Small chain is also typical of secondaries in 

morphology and overlap but is not radial to Imbrium or any other likely source. Apollo 

16 frame M-1671.

B. Abulfeda chain, parallel to Miiller chain and also site of numerous craters. A, Abulfeda 

(65 km, 14° S., 14° E.). Orbiter 4 frame H-89.

Roddy, 1977, p. 297; Trulio, 1977; Melosh, 1980). Thus, most impactcrater excavation is a response to the disequilibrium suddenly 

induced by the intensely energetic, penetrative shock wave and is not 

the direct result of an expansion like that induced by a true chemical 

explosion (Shoemaker, 1962b, p. 316). Another difference from most 

explosions is that the energy from a hypervelocity impact is released 

along the length of the path of penetration, which may be short or may 

be an elongate, cigar-shaped zone (Jones and Sandford, 1977, p. 1009; 

Gault and Wedekind, 1978). Explosions at moderately shallow depths, 

however, may mimic impact effects (Shoemaker, 1960,1963; Baldwin, 

1963, chap. 7; Roddy, 1968,1977; Oberbeck, 1971a, 1977; Melosh, 1980) 

because their energy may be coupled into the ground similarly 

(Cooper and Sauer, 1977; Knowles and Erode, 1977, p. 874; 

Kreyenhagen and Schuster, 1977; Trulio, 1977).

Because the decompression and not the direct, intense shock 

compression excavates most of the cavity, ejection endures longer 

than the shock compression (Gault and others, 1968b). After most of 

the projectile material has been ejected, the cavity, lined with partly 

molten material, continues to grow behind the advancing shock wave 

by ejection of target material (figs. 3.17.E, F). Though basically orderly, 

at least in simple craters (Shoemaker, 1960; Gault and others, 1968b), 

the ejection process is more subject to vagaries stemming from 

inhomogeneous properties and structures of the target than is the 

shock-compression phase. For example, inhomogeneities are the 

probable cause of the pattern of secondary chains and loops (Shoemaker, 1962b).

Ejection of the target material occurs approximately in the order 

it is enveloped by the shock wave (Stbffler and others, 1975). In 

general, more highly shocked materials high in the target and just 

outside the impact zone leave first at the highest velocities and angles 

(measured above the horizontal). Molten material may be shot into 

high ballistic trajectories. Subsequently, expanding concentric zones 

that include more moderately shocked materials from increasingly 

deep target materials are successively ejected. The ejecta forms an 

upward- and outward-flaring curtain of debris in the shape of an 

inverted lampshade (a frustum). The materials that form this curtain 

are sheared up along the walls of the growing cavity to the cavity lip, 

where they leave at angles parallel to the walls (Gault and others, 

1968b; Oberbeck, 1977, p. 46; Orphal, 1977); the curtain is like an 

extension of the crater wall. The curtain continuously expands outward during crater growth. It apparently always remains thin, at 

least in small craters (fig. 3.17; Oberbeck, 1975; Oberbeck and Morrison, 1976). Typical ejection angles for the main, middle stages of 

experimental crater formation are 40° to 60° above horizontal but also 

depart from these values, depending on such properties of the target 

material as layering and competence (Andrews, 1977; Orphal, 1977; 

Wisotski, 1977).

C. Large group of craters with rectangular outlines, alternatively interpreted as originating 

by volcamsm (Mutch and Saunders, 1972) or secondary impacts. Shape and clustering 

are consistent with source to lower right, but no likely source is in that direction. Group is 

radial to Onentale basin, but Onentale secondaries at this distance are generally sharper 

(compare fig. 4.6). Several other large craters are also clustered and may be basin 

secondaries (Schultz, 1976b, p. 276). Crater under north arrow is Asclepi (43 km, 55° 

S., 25° E.). Orbiter 4 frame H-IOO.
42 THE GEOLOGIC HISTORY OF THE MOON

FIGURE 3.17. Stages in formation of a simple impact crater. Drawing by Donald E. Davis, 

courtesy of the artist. 

A. Initial contact and jetting.

B, C, D. Compressional shock wave propagates outward, and cavity grows by rarefaction 

behind shock wave while projectile is consumed.

E. Cavity continues growth after projectile has been consumed.

F. Maximum crater size.

C. Frustum-shaped curtain of ejected debris continues outward expansion after cavity ceases

growth and overturned ejecta flap comes to rest. 

H. Final crater configuration.

The cavity of a typical small simple crater ceases to expand 

downward when^it has acquired depths of V5 to Vz of the diameter 

(Dence and others, 1977, p. 250-253; Knowles and Erode, 1977, p. 

890-891; Orphal, 1977, p. 909). Shearing flow at the walls may then 

continue to broaden the crater after this maximum depth is reached 

(Orphal, 1977; Piekutowski, 1977; Swift, 1977). Weakly shocked and 

nearly undamaged ejecta derived from near the walls leaves last, 

during the final stages of crater excavation, under relatively low 

stresses and at low velocities and ejection angles (Shoemaker, 1962b, 

p. 335; Oberbeck, 1975; Stoffler and others, 1975; Andrews, 1977; 

Cooper, 1977, p. 25; Orphal, 1977). The size of ejected debris increases 

during cavity growth owing to decreasing shock pressures, decreasing fragmentation of the wallrock, and lower ejection velocities. 

Finally, ejection ceases as the tensile strength of the rock overcomes 

the power of the rarefaction wave to move it.

Ejecta deposition

After the cavity ceases to grow, the inverted frustum-shaped 

ejecta curtain continues to advance outward beyond the cavity (fig. 

3.17G). The curtain continuously decreases in height while expanding 

in diameter, as materials at the base are deposited on the surface from 

the cavity rim outward. Because most ejecta is launched from a 

simple crater at nearly constant exit angles but at decreasing 

velocities, the ejecta front slopes outward at a nearly constant angle 

with the surface, generally 40°-50° (Oberbeck, 1975,1977; Oberbeck 

and Morrison, 1976; Andrews, 1977, p. 1090-1092; Cooper, 1977, p. 

38-39). As a result of this velocity distribution in the curtain, ejecta 

deposition occurs in approximately the reverse order of excavation. 

The first material to be deposited, from the base of the ejecta curtain, 

is the last to have been engulfed by the shock wave and to be sheared 

from the crater walls (figs. 3.17,3.18). That material is lofted or barely 

pushed over the rim at low velocities and soon lands near the crater 

rim. In simple craters, the near-rim material may form a more or less 

coherent overturned flap in which the stratigraphic sequence of 

redeposited target materials is the reverse of their preimpact 

sequence (fig. 3.17H; Shoemaker, 1960, 1962b, 1963; Roberts, 1966; 

Roddy, 1968,1976,1977, p. 201; Stoffler and others, 1975; Moore, 1976; 

Oberbeck and Morrison, 1976). Some of the relief on crater rims is also 

due to structural upthrust and outthrust in a manner dependent on 

target structure and depth of energy release (Shoemaker, 1960,1963; 

Roberts, 1968; Gault and others, 1968b). For a considerable time after 

the cavity stops growing, the curtain continues to move outward and 

to drop material from its lower edge into an ever-expanding but 

thinning ring of deposits.

Secondary cratering and ground surge

With increasing distance from the crater, the ejecta that strikes 

the surface forms secondary craters rather than building up a deposit. 

At some distance, the curtain separates into filaments of debris,

whose impact creates loops and chains of secondary craters and rays. 

Circular secondaries are formed by the last material to impact the 

surface, that which was launched first at the highest velocities and in 

the longest, highest trajectories from sources high in the target near 

the impact zone.

The general picture, then, around craters as well as ringed basins 

is one of outward-thinning deposits of primary ejecta grading into 

increasingly conspicuous secondary craters and their ejecta deposits 

(Shoemaker and Hackman, 1962; Schmitt and others, 1967; Ulrich, 

1969; Guest, 1973). The secondary impacts excavate material of the 

local terrain around the primary crater (Moore and others, 1974; 

Oberbeck and others, 1974,1975,1977; Morrison and Oberbeck, 1975, 

1978; Oberbeck, 1975,1977; Oberbeck and Morrison, 1976; Horz, 1981). 

The amount of local material excavated increases outward. The picture is complicated around many craters and, particularly, around 

basins by the superposition on secondary craters of material that has 

flowed from points closer to the crater or basin (Shoemaker and 

Hackman, 1962, p. 291; Hodges, 1973b; Scott and Eggleton, 1973; 

Moore and others, 1974; Morrison and Oberbeck, 1975; Wilhelms, 

1976,1980). In the picture above of the regular advance of ejecta in a 

narrow curtain, the ejecta that forms the secondaries was still in 

flight when the ground-flow deposits left the crater, but the flying 

ejecta impacts the surface before the ground-flow deposits reach the 

secondary-impact zone.

The question of how the ground flow originates has been much 

debated because of its importance in studies of lunar-sample provenance. Chao (1974,1977) and Chao and Minkin (1977) cited evidence

Tangent above

100

Tangent below

200

Numbers on contours indicate 

the range of ejecta as a 

function of R/V V'

200

FIGURE 3.18. Relation between sources of ejecta and range at which it is deposited (from 

Cooper, 1977, fig. 8). R, range; V, volume of ejecta; charge, 1.7 g of lead azide; target, dry 

sand. Results are similar to those of much larger experiments (Cooper, 1977, p. 25). Zones of 

target material nearest charge were ejected first and to greatest ranges; successively larger and 

more distal zones were deposited nearer rim. 

A. Charge tangential above surface. 

B. Charge tangential below surface, similar to energy release of impacts.
3. CRATER MATERIALS 43

from the 25-km-diameter Ries complex crater or ringed basin, southern Germany, which suggests that the inner ejecta was pushed over 

the rim and moved outward along the surface by gliding and rolling. 

Ground flow may also have emplaced most ejecta of the 1.7-kmdiameter simple crater Lonar, India (Fudali and others, 1980). In 

contrast, V.R. Oberbeck and his coworkers believe that all ejecta 

travels in ballistic trajectories and that surface flow is initiated when 

the ejecta impacts the surface after flight, whereupon it moves outward, together with the abundant excavated local material, under the 

forces of momentum and gravity. The mechanism proposed by Chao 

(1974,1977) would also incorporate local material, because the surface 

movement occurs under confining pressure, but the proportions of 

local to primary material might be less than in the ballistic-impact 

mechanism.

To avoid prejudging the proportions of primary ejecta and local 

material present, use of the term "continuous deposits" is recommended over "ejecta" for the inner ground-covering material (Oberbeck and others, 1974; Oberbeck, 1975). The term "base surge" is 

commonly used for the outward flow (for example, Lindsay, 1976) but 

implies gas- or water-assisted transport, which is unlikely on the 

Moon. Therefore, the terms "ground surge," "debris surge," or "debris 

flow" are preferable (Moore and others, 1974; Morrison and Oberbeck, 

1975; Oberbeck, 1975).

Deformation and nonballistic ejection

Beyond the sphere of shocked material that is launched into 

ballistic flight, other target material is less highly shocked but is 

deformed and sheared. Some of this material is pressed downward 

and outward along the crater floor and walls in curving paths that 

resemble those that precede the ballistic ejection (Dence and others, 

1977; Croft, 1980). Part of this peripheral material may surge over the 

walls and exit the crater at very low velocities without leaving the 

surface.

Additional shock-damaged material is not permanently ejected. 

Some may be lofted above the crater and fall back inside; other 

material is mildly brecciated or fractured in place without significant 

dislocation. Both the fallback and the inplace material form a breccia 

lens in the bottom of the crater that grades downward and outward 

into a zone of fractured rock (Shoemaker, 1960, 1962b, 1963; Chao, 

1977; Dence and others, 1977, p. 250-253).

The volumes of the cavity from which material was ejected and of 

the deformed zone may obey different scaling laws (Dence and others, 

1977, p. 266-268; Croft, 1980, 1981). In general, a representative 

linear dimension of a crater say, the radius or the cube root of the 

volume will scale to the kinetic energy of impact (Shoemaker, 

1962b; Shoemaker and others, 1962a; Baldwin, 1963; Gault and 

Moore, 1965; Gault and others, 1968b; Gault, 1970,1974; Cooper, 1977, 

p. 16-24; Dence and others, 1977, p. 264-271; Roddy and others, 1977, 

p. 1133-1296; Melosh, 1980). This linear dimension is also controlled 

both by gravitational attraction, which inhibits lofting and flight of 

ejecta, and by rock strength, which influences dissipation of shock 

energy. Target strength becomes less important, and gravity more 

important, with increasing impact magnitude; gravity dominates in 

craters larger than about 100 m in diameter (Moore and others, 1963; 

Gault and Moore, 1965; Gault, 1974; Chabai, 1977; Gault and 

Wedekind, 1977; Gaffney, 1978). Smaller craters are created for a 

given impact kinetic energy when gravity dominates than when 

strength dominates. In large craters and basins, the excavation cavities may be "gravity craters," whereas their exterior zones may be 

"strength craters" (Croft, 1980,1981).

Various terms, most of them used in more than one sense, have 

been applied to the excavated and nonexcavated parts of a crater. Both 

"transient crater" and "true crater" have been used to describe the 

combined zones out to the limit of the nonexcavated breccia, before 

modification by slumping. However, the term "transient crater" commonly denotes a stage in the cavity growth. "Apparent crater" has 

been used to denote the depression that is seen excluding the inplace 

breccia; alternatively, it has been defined as that part of the depression which lies below the precrater ground surface. Therefore, the 

usage of each term by a given author must be checked. In this volume, 

I avoid these terms because of their ambiguity and use only "excavation cavity," for the cavity from which material has been removed, 

either permanently or temporarily before falling back.

Peak and terrace formation

The most conspicuous indicator of internal deformation in craters is the presence of central peaks and wall terraces in complex 

craters (Gilbert, 1893; Baldwin, 1949,1963; Dence, 1965,1968; Quaide 

and others, 1965; Howard, 1974; Dence and others, 1977; Pike, 1980a, 

b, c). Diameters of complex craters are generally larger than 16 km in 

the maria and 21 km in the terrae (fig. 3.3; Pike, 1980a, b). Peaks 

increase in size and complexity in proportion to crater size up to 

crater diameters of about 40 or 50 km, but diminish in relative size in 

larger craters (figs. 3.2C-E; Murray, 1980, p. 283). Wall failure first 

appears as scalloping in small and rudimentary complex craters (fig. 

3.2C; Head, 1976b; Settle and Head, 1979). Continuous terraces, 

which probably form by base failure (Grieve and others, 1977; Melosh, 

1977), are observed in larger complex craters. Within the transition 

size range, d/D ratios also decrease markedly, and still larger craters 

have broad shallow floors (figs. 3.2D, E\ Thus, peaks, terraces, and 

relatively shallow floors all seem to be somehow related. However, 

these features do not become evident at exactly the same diameter; for 

example, peaks appear before terraces (Smith and Sanchez, 1973; 

Head, 1976b; Cintala and others, 1977; Smith and Hartnell, 1978; 

Wood and Andersson, 1978; Settle and Head, 1979; Pike, 1980a, b, c).

Interpretations of the simple-to-complex transition are commonly reduced to a choice between "push" and "pull" mechanisms: 

Did the walls of a deep bowl-shaped cavity first collapse centripetally 

into terraces and push up the central peak, or did the peak originate 

by some type of rebound of the subcrater material that pulled the 

walls inward to their collapse? The material trajectories and the 

ultimate geometry would be similar in both cases. Some form of the 

"pull" model seems to be supported by appearance of peaks before 

terraces in the sequence of feature development. The issue largely 

revolves about the question of the original depth and shape of complex 

craters: Did they grow to d/D ratios that are proportional to those of 

simple craters and then collapse in the "modification stage" (Dence, 

1968; Dence and others, 1977; Grieve and others, 1977), or did they 

grow nonproportionally and start out with shallower depths (Croft, 

1978,1980; Settle and Head, 1979; Pike, 1980b)?

A variant of the "pull" model involves a more complex excavation 

cavity, consisting of a deep inner part surrounded by a shallower 

shelf. Only the central part grows proportionally. A key feature is 

subhorizontal inward motion of the subcrater material during crater 

growth (fig. 3.19). Material beneath the central depression is severely 

deformed and sharply uplifted to become the central peak (Roddy, 

1968, 1976, 1977, 1979; Milton and Roddy, 1972; Milton and others, 

1972; Wilshire and others, 1972a; Offield and Pohn, 1979). In contrast, 

strata beneath the shallow floor surrounding the peaks are mildly 

deformed and are neither uplifted nor downdropped substantially. 

Support is withdrawn from the walls, which collapse most completely 

in large craters, where the inward motion is greatest.

Causes of the deformation and of the simple-to-complex transition remain particularly debatable (Quaide and others, 1965; Pike, 

1980a, b). Some phenomenon may relate shallow depths of burst, peak 

formation, and shallow floors. Shallow bursts, which may form peaks 

in very small craters, have been documented by stratigraphic relations and the orientation of shatter cones in such terrestrial craters as 

Grosses Bluff, Australia (Milton and others, 1972) and Flynn Creek, 

Tenn. (Roddy, 1977). Low-density projectiles would release their 

kinetic energy near the surface because of their shallow penetration 

(Roddy, 1968,1977), as would very large projectiles of any density (fig. 

3.20). Large bodies penetrate to shallower depths relative to their 

kinetic energy than do smaller objects, and require more time to be 

consumed (Baldwin, 1963, p. 164-184). Thus, (1) peaks may form 

because of the shallow energy release; (2) more energy is directed 

laterally than in deeper bursts, so that craters are shallow; and (3) 

energy coupling is sustained over longer times, so that more of the 

target material is heated than in rapid energy releases (see next 

section below). Many properties of complex craters and ringed basins 

(see chap. 4) are explainable in this way.

Shallow bursts do not explain all phenomena, however, because 

the diameters in the simple-to-complex transition and the morphologic properties of peaks and floors differ from region to region 

and from planet to planet (Cintala and others, 1977; Pike, 1980a, b). 

Two factors are commonly proposed as modulators or even as initiators of the complex-cratering process and also of basin-ring formation: (1) differing gravitational attraction at the planets' surfaces
44 THE GEOLOGIC HISTORY OF THE MOON

FIGURE 3.19. Model of formation of features of complex crater, according to Pike (1980B, fig. 9). Only strata A, B, and C are exposed in central uplift, where they dip vertically or steeply. 

Stratum F has been stripped from crater, although it forms a slump block (b). Stratum G is composed of ejecta and intracrater breccia. Faults creating terraces (t) pinch out at shallow depths. 

Arcuate dotted lines indicate zones of deformation in center. Modeled after four complex craters on Earth; scale applies to average horizontal and vertical dimensions of those craters.

0.10

0.10- 6.0

0.00

-0.00

O

O

FIGURE 3.20. Relative depth of penetration and effect on crater diameter of large and small 

projectiles (Baldwin, 1963, fig. 29). D, logarithm of diameter (in feet). Depth/diameter 

ratios are typical of lunar craters. ///W/3 , scaled depth of burst; H, depth (in feet); W, 

explosive energy (in pounds of TNT equivalent). Left panel, penetration assumed beneath 

original ground surface to two projectile diameters for all seven sizes of craters; impact velocity 

is set at 10 mi/s; unlikely result of very shallow crater overlying deep burst is obtained for 

largest craters. Right panel, more realistic variable and shallow depths of penetration, 

including surface bursts for two largest sizes. Projectiles flatten during consumption by shock 

wave (Baldwin, 1963, p. 167).

(Hartmann, 1972a; Gault and others, 1975; Pike, 1980a, b), and (2) 

substrate properties, especially layering. Discontinuities between 

layers affect the coupling of impact energy into the target and refract 

and reflect the shock waves (Sabaneyev, 1962; Oberbeck and Quaide, 

1967,1968; Quaide and Oberbeck, 1968; Rinehart, 1975; Head, 1976b; 

Oberbeck, 1975, 1977; Piekutowski, 1977; Hodges and Wilhelms, 

1978). Different stratification may explain the different simple-tocomplex transition diameters between maria and terrae. Gravity or 

substrate properties, however, cannot be the sole factor in complexcrater formation, considering that peaks appear in all sufficiently 

large craters on all impacted moons and planets (Pike, 1980a, b).

In summary, shallow energy releases, strong gravitational 

attraction, and sharp contrasts in target stratification all seem to 

enhance formation of the peaks, shallow floors, and terraces characteristic of complex craters. Different combinations of these factors 

among planets or on a single planet, such as the Moon, yield quantitatively different results. None of these factors, however, may be the 

fundamental cause of complex craters. Changes in the basic physical 

effects of impact above some energy threshold (for example, Melosh, 

1980) ultimately may be found to be more significant.

Impact melting

Estimates of the amount and stratigraphic relations of impact 

melt are important to the interpretations of lunar breccia presented 

later in this volume (chaps. 9,10). Projectile size, impact velocity, and 

density of both the projectile and the target affect the partitioning of 

energy among mechanical excavation, deformation, and impact melting (Baldwin, 1963; Gault and Heitowit, 1963; Gault and others, 

1968b, 1975; Dence, 1971; Ahrens and O'Keefe, 1972; O'Keefe and 

Ahrens, 1977; Kiefferand Simonds, 1980; Melosh, 1980). High-velocity 

impacts of small projectiles generate small amounts of very hot melt,

whereas slow large impacts generate larger amounts of less thoroughly melted material (Rehfuss, 1974; O'Keefe and Ahrens, 1975, 

1977,1978; Grieve and others, 1977, p. 809; Lange and Ahrens, 1979). 

For the same mass and velocity that is, the same kinetic energy 

low-density projectiles generate more melt than do dense projectiles 

(Kieffer and Simonds, 1980). The melt traps much of the heat energy 

of the impact and does not contribute to the excavation process except 

for the melt that may have been lost by early jetting (Shoemaker, 

1962b, p. 316; Gault, 1974). Therefore, in proportion to kinetic energy, 

large impactors melt more material and eject less fragmental material than do small impactors.

Emplacement of impact melt is one of the last processes to run its 

course during a cratering event but one of the first to be initiated. Melt 

originates in the impact zone and is pressed outward along the walls 

of the growing cavity (fig. 3.21). On the Moon, what appear to be flows 

of impact melt are superposed on crater walls and flanks, and pools of 

melt rest in depressions both inside and outside the rim crest (figs. 

3.22, 3.23). These relations indicate that the melt is still mobile after 

the fragmental ejecta leaves the crater (Shoemaker and others, 1968; 

Guest, 1973; Moore and others, 1974; Howard, 1975; Howard and 

Wilshire, 1975; Schultz, 1976b, p. 228-237; Hawke and Head, 1977a). 

Apparently the melt sloshes over the rim along with some of the last 

fragmental ejecta, and locally flows back downward into the cavity 

after terracing ceases (Howard, 1975; Howard and Wilshire, 1975; 

Grieve and others, 1977; Phinney and Simonds, 1977).

 E|ecta

V Projectili

B

C

Glass injected Vesicular 

along fractures glass Glass 

splash 

__^k

Laminated and banded materials 

. (arrested cataclastic flow)

V^ Compressed and 

>y N^ sheared target

Tensile fractures ractures-, / ^ \

I I > \

FIGURE 3.21. Formation and mixing of shock-grade zones in a simple crater, showing four 

stages (Wilshire and Moore, 1974, fig. 11).
3. CRATER MATERIALS 45

Formation times

Impact craters form in less than a minute (Melosh, 1980, p. 80). 

Schmidt (1981) suggested that formational time is proportional to the 

sixth power of the excavated volume.

In some terrestrial craters, the melt contains breccia from the 

peak (Simonds and others, 1976b). This observation confirms that 

peaks form early in the cratering sequence, as suggested above, 

because impact melt solidifies very quickly. In large craters on the 

Canadian shield, the time between impact and melt solidification was 

about 100 s (Onorato and others, 1976, 1978). Final cooling below 

temperatures able to metamorphose the melt and clasts required less 

than 2,000 years (Onorato and others, 1978).

IMPACT BRECCIA

Terrestrial analogs

Studies of impact-generated terrestrial materials, which began 

during the 1960's in conscious preparation for the Apollo landings 

(Chao, 1967; Engelhardt, 1967,1971; French and Short, 1968; French, 

1977, p. 155), have provided essential clues for interpretations of the 

returned samples. Progressive mineralogic and textural changes in 

the target rock have been correlated with intensity of the outwarddecaying shock wave (French and Short, 1968; Dence and others, 

1977; Robertson and Grieve, 1977).

Among the first impact-generated materials to attract attention 

were tektites, small, glassy, aerodynamically shaped objects found in 

strewn fields at several terrestrial localities. They have been thought 

to be lunar in origin (O'Keefe and Cameron, 1962; O'Keefe, 1963) but 

are much closer in composition to the target materials of terrestrial 

craters (King, 1976, chap. 2; Glass, 1982, chap. 6). The tektites of each 

strewn field probably originated during early jetting of very hot 

material from a terrestrial impact (fig. 3.21A; Kieffer and Simonds, 

1980), were carried to various altitudes in and above the atmosphere 

by a fireball, and were dispersed by gases in the fireball and by 

stratospheric winds (Jones and Sandford, 1977).

Lunar studies have depended heavily on analogies to the Ries 

crater in Germany and the numerous craters or ringed basins on the 

Canadian shield. The diameters of the most conspicuous rings of the 

three largest Canadian craters, Manicouagan and West and East

Clearwater, are 65, 32, and 23 km, respectively; however, the exterior 

ejecta has been eroded away and the excavation diameters are uncertain. The Canadian craters have furnished much data about shock 

metamorphism and impact melting (Dence, 1968, 1971; Grieve and 

others, 1974,1977; Grieve, 1975,1978,1980; Floran and Dence, 1976; 

Simonds and others, 1976a, b, 1978a, b; Dence and others, 1977; 

Phinney and Simonds, 1977; Phinney and others, 1977, 1978; 

Robertson and Grieve, 1977; Floran and others, 1978; Grieve and 

Floran, 1978). The larger craters possess massive interior sheets of 

impact-melt rock, about 100 m thick, which texturally resembles 

igneous rock. The 25-km-diameter Ries crater has furnished fewer 

impact-melt rocks (as parts of the deposit called suevite) but abundant unshocked fragmental ejecta (mostly in the Bunte Breccia) 

(Engelhardt, 1967; Dennis, 1971; Chao, 1974,1977; Gall and others, 

1975; Chao and Minkin, 1977; Pohl and others, 1977; Horz and 

Banholzer, 1980). Representative types of terrestrial impact materials from these and other craters have been illustrated and compared 

with lunar impact materials by Stoffler and others (1979).

Lunar-terra samples

Except for tektites, the whole range of shock grades is much 

better represented on the lunar terrae than on the eroded Earth. 

Terra-breccia deposits display complex shock effects, ranging from 

cataclasis of monolithologic rocks, through solid-state metamorphism, to intricate assemblages of cogenetic and foreign clasts 

within clastic and impact-melted matrices (see fig. 2.7, chaps. 8-10, 

and summaries by James, 1977; Phinney and others, 1977; Stoffler 

and others, 1979, 1980; Taylor, 1982, p. 187-201). Some individual 

specimens consist entirely of breccia or entirely of melt rock (fig. 2.8), 

and have been so designated in the literature. However, these rock 

types must have been parts of deposits containing both melted and 

unmelted materials, and so the term "breccia" is used here as a 

general term for lunar impact materials. 3- 1

3 :Only the breccia that forms deposits of the terra bedrock is discussed here. At the time of the first two Apollo 

landings, the only known breccia was that created in regoliths, because the terrae had not yet been visited. The 

regolith breccia (also called soil breccia, microbreccia, glassy breccia, vitric-matrix breccia, and numerous other 

names) constituted an obviously different type of material from the igneous basaltic lavas found at those sites. When 

the much coarser, more massive, more complex, and more extensive bedrock breccia became known, the nomenclature did not fully adapt to the differences. Some petrologic literature does not clearly distinguish between 

regolith breccia, formed by innumerable small impacts near the lunar surface, and bedrock breccia, formed by fewer 

and larger impacts (also called "megaregolith"). In this volume, "breccia" denotes bedrock breccia unless otherwise 

indicated, and "megaregolith" is not used because it obscures the fact that the bedrock breccia forms discrete strata.

 -^^ ' -

FIGURE 3.22. Material subunits of a typical large crater, modeled after craters Theophilus and Madler. Rim materials are three-dimensional deposits, from rim crest outward: rh, hummocky; rr, 

radial; re, cratered (secondary craters). Structural units, not sharply demarcated in subsurface, are floor material (f), peak material (p), and wall material (w). Planar impact-melt rock (m) overlies 

other units. Material of simple crater (c) overlies unit rr. Drawing by Donald E. Davis, courtesy of the artist.
46 THE GEOLOGIC HISTORY OF THE MOON

Some features of breccia deposits can be related to the outward 

decrease in shock intensity during the growth of a crater (fig. 3.21), 

whereas others illustrate the textural and compositional complexities 

that impact cratering induces.

Impact melts that have glassy or crystalline textures most like 

those of igneous basalt presumably were formed near the impact zone 

(for example, ophitic texture, in which pyroxene encloses less abundant plagioclase laths, or subophitic texture, in which pyroxene 

partly encloses a subequal volume of plagioclase laths; compare figs. 

2.6, 2.8). Poikilitic texture in which numerous small, more nearly 

equant plagioclase grains and other debris are enclosed by large 

pyroxene crystals is considered to be a sign that many minute clasts 

were incorporated in the impact melt (Nabelek and others, 1978). Melt 

rocks with these various textures were recovered from most sampling 

sites, most abundantly from the Apollo 17 terra stations (see chap. 9). 

Melt-poor friable fragmental breccia, likely formed in the outer parts 

of excavation cavities, was returned from all the terra landing sites, 

most abundantly by Apollos 14 and 16 (see chaps. 9,10). Some material 

derived from the outer parts of the excavation cavities and ejected late 

in the cratering sequence consists of coherent blocks (fig. 3.2LD).

This zoning is muddied by the turbulent processes that characterize energetic impacts. Melted material of projectile and target that 

is not jetted is pressed into the growing cavity (fig. 3.2IB). As the 

crater grows, this material, formed near the energy-release zone, is 

injected into the less highly shocked but fragmented rock in outer 

zones (figs. 3.2IB-D; Dence, 1971; Wilshire and Moore, 1974; Simonds 

and others, 1976a, b, 1978b; Chao and Minkin, 1977; Grieve and 

others, 1977; James, 1977; Grieve and Floran, 1978; Stoffler and 

others, 1979). Thus, melt rock appears as matrices and dikes among 

less highly shocked fragments (fig. 3.2LD). Unmelted but crushed 

clastic debris may be similarly injected. These processes are illustrated by dimict breccia, consisting of fragmental breccia and melt 

rock showing mutual intrusion and inclusion relations (James, 1977, 

p. 647; 1981; Stoffler and others, 1981).

Such mingling of shock grades can be quite thorough: Rock and 

mineral fragments, shocked to varying degrees, may be engulfed by 

impact melt (fig. 2.7). Moreover, various breccia and melt types are 

interbedded in the crater interiors (Stoffler and others, 1979). This 

bedding varies in sequence and lithology for craters of different sizes 

and different target materials (Stoffler and others, 1979; Stoffler, 

1981), and commonly juxtaposes materials that have experienced a

FIGURE 3.23. Very fresh crater King on farside (77 km, 5° N., 121° E.)

A. Entire crater. "Lobster claw" central peak connects with wall and is overlain by lower 

wall material (enlarged in fig. 3.32). Pool of impact melt outside northwest rim is 

enlarged in figure 3.36. Apollo 16 frame H-19580.

wide range of shock pressures. At any stage in the cratering process, 

the material being ejected may consist of a mixture of complete melt, 

fragment-laden melt, and unmelted fragments.

Shearing of material from the crater walls and transporting it 

outside the crater further mixes the ejecta and streaks it into bands 

or pods. Glass-coated striated fragments evincing shearing during 

ejection or transport are known (Wilshire and Moore, 1974). Some 

samples even appear to be accretionary bombs (Stoeser and others, 

1974; Wood, 1975d; Spudis and Ryder, 1981). Mixing with the substrate 

material increases outward from the crater because of surface transport and increasing secondary-impact velocity of the ejecta. Lunar 

breccia containing clasts of melt formed in the same impact that 

emplaced the breccia has been referred to as "suevite"; clasts of 

impact-melted rock in fragmental deposits analogous to the Ries 

Bunte Breccia predate the deposit (Stoffler and others, 1979, 1980). 

However, distinguishing the two types of deposit is difficult on the 

Moon, where the source crater of a breccia is commonly unknown.

After the deposits come to rest, the intimate mixture of hot or 

superheated melt and warm or cold fragments of varying compositions in all conceivable combinations of grain size and abundance 

leads to further modifications of the deposit (Simonds, 1975; Simonds 

and others, 1976a, b). Melt and hot clasts are differentially quenched 

during the rapid initial cooling of the deposit. Sharp local differences 

in crystallinity, interface texture, and remaining content of fragments may arise. The second, much slower stage of cooling further 

metamorphoses the deposit and underlying materials. The terrestrial metamorphic terms "granoblastic" or "granulitic" (anhedral 

grains of subequal size) are applied to many lunar textures believed 

to have arisen from thermal metamorphism (Warner and others, 

1977; Taylor, 1982, p. 198-200). Thus, the complexities arising in a 

single impact may be so great that they mimic the relations arising 

from multiple events (Simonds, 1975).

The complexities of lunar impact deposits at the scale of the 

outcrop or the hand specimen are not surprising in view of the 

complexity of each individual impact and the fact that each volume of 

material may be similarly reworked by many subsequent impacts. 

Consequently, any attempt to categorize lunar impact materials is 

bound to fall short of reality. The most useful in a series of attempts

B. Northeast rim, showing concentric textures of rim material downdropped along with 

terraces inside rim and locally lying outside rim (lower part of photograph). Gradation 

from concentric to radial pattern of outer rim material indicates increasing radial-flow 

component of outer material. Apollo 16 frame M 1579.
3. CRATER MATERIALS 47

was probably that by Stoffler and others (1980), who compiled a long 

list of earlier names that should greatly aid the reader of the technical 

literature. Stoffler and others (1980) used textures and structures, 

which are signs of the process that formed the rocks, as their main 

basis for classification. Essentially, this classification describes 

matrices as either fragmental, crystallized from a melt, still glassy or 

partially devitrified, or metamorphosed. Clast content can be 

described according to need. From the geologic point of view, this 

classification improves upon those based on chemical composition. 

Many purposes are served by a still simpler breakdown into clast-rich 

and melt-rich (clast-poor) breccia types.

Reconstruction of the history of lunar breccia deposits severely 

challenges the petrologist, geochemist, and geologist. The place and 

process of original formation of the clasts and of the raw materials for 

the melts (chap. 8) must be disentangled from the process by which 

they arrived at the place where they were collected (chaps. 9,10). Each 

clast may yield a different isotopic age, attesting to many endogenic 

and impact events in the sample's history, or only to incomplete 

migration of Ar, Rb, and Sr isotopes (chaps. 9,10). The crater or basin 

in which most of them originated can only be inferred. Unsuspected 

complexities in the process of forming craters and, especially, basins 

may have shaped the materials of the lunar terrae.

SUMMARY OF 

CRATER-MATERIAL ORIGIN

Introduction

Around almost all impact craters, the processes described above 

have created similar patterns of continuous ejecta grading outward to 

secondary craters. The major apparent differences among craters are 

between the interiors of simple and complex craters. Partly understood variations in the target material, such as in structure, cohesion, 

and topographic relief, also affect the details of morphology; and 

variations in projectile velocity, density, and impact angle affect both 

the morphology and the proportions of melted and ejected material. 

Nevertheless, the basic morphologic patterns of impact craters are 

much more alike than those of volcanic craters.

This similarity has greatly aided the reconstruction of lunar 

geologic history. Stratigraphically based lunar geologic mapping has 

become increasingly effective as the analogies between modified and 

unmodified deposits have become clearer. The material subunits in 

and around craters were generally recognized as stratigraphic 

entities when they were mapped geologically in the 1960's (Mutch, 

1970, p. 165-175; Wilhelms, 1970b, p. 40-42), and their mapped limits 

remain generally valid. The interpretations of some units, however, 

have changed (compare Schmitt and others, 1967, and Howard, 1975). 

To enable modern use of these maps, this section updates the interpretations of each class of unit on the basis of the newer research 

summarized above, and suggests what rock types are likely to occur 

in each unit and in the analogous units of basins.

Rim material

A raised rim, formed by uplift of the target and by deposition of 

ejecta, surrounds the crater. The geologic unit "rim material" (a 

member of the formation "crater material") essentially corresponds to 

the continuous ejecta (fig. 3.22). Smooth, clean-looking rim surfaces 

may have been stripped of ejecta by late-stage radial flow (Guest, 

1973). On the near flanks of simple craters, ejecta materials have been 

deposited in inverse order of their original sequence. Original stratigraphic units are more jumbled in complex craters (Horz, 1981). 

Impact melt forms ponds in depressions and coats other parts of the 

rim (figs. 3.23). The deposits also probably incorporate clasts and 

veins of impact-melt rock along with fragmental ejecta. The relative 

amounts of melted and unmelted ejecta depend on factors of projectile 

size, density, and velocity, and on such target properties as density 

and volatile content.

The hummocky facies of rim material (a submember) consists of 

late-emplaced material that was derived from near the crater wall 

and overlies the structurally uplifted and outthrust zone of larger 

craters. Blocks of relatively undeformed material are scattered on the 

rim (figs. 3.2A, 3.23.B). The rim's generally concentric structure 

attests to final emplacement by outthrusting (fig. 3.23). In places,

younger crater rims have breached older rims, and the hummocky 

materials flowed rapidly (fig. 3.24) or sluggishly (fig. 3.25) into the 

gap. Therefore, the deposits were emplaced along the surface. Around 

old craters, the hummocky rim facies, without the topographic subtleties of young craters, is the only distinguishable rim material (fig.

3.26).

Outward from the hummocky rim material, ridges in radial and 

herringbone patterns form a continuous-appearing wreath with 

irregular lobate margins (fig. 3.22). In the well-photographed complex 

craters illustrated here, the radial rim facies (submember) is clearly 

an outward gradation of the hummocky rim material (figs. 3.23, 3.24,

3.27). Surface flow after ejection is indicated by the textures curving 

from concentric to radial patterns (figs. 3.235, 3.24, 3.27A). Some of 

the ejecta apparently was molten (fig. 3.275). The gradation with 

materials at the rim crest and the content of impact melt indicate that 

the ejected ground-flow materials are dominated by primary ejecta. 

The outer part of the radial rim material, however, contains locally 

derived material because substrate material was incorporated either 

by ground flow of the ejecta under confining pressure or during 

ballistic emplacement of the ejecta.

FIGURE 3.24. Lineations of northwestern ejecta of crater Tsiolkovskiy, indicating ground 

surge beyond concentric near-rim deposits. Smaller crater in upper left, probably filled with 

additional Tsiolkovskiy ejecta, is Liitke (39 km). Apollo 17 frame M 2608.

FIGURE 3.25. Concentric ejecta of crater Shirakatsi (51 km, left), pushed through gap in rim 

of older or nearly contemporaneous crater Dobrovol'skiy (38 km). North of Tsiolkovskiy, 

centered at 13° S., 129° E. Apollo 17 frame M-2608.
48 THE GEOLOGIC HISTORY OF THE MOON

At some radius, ballistically generated secondary craters appear. 

Much of the ejecta thrown from secondaries has a herringbone pattern that originated by intersection of cones of their ejecta (figs. 

3.4C-F, 3.6, 3.7). In the ideal case, derivation of the farthest-thrown 

materials from near the impact zone (fig. 3.18) means that the secondaries were formed by highly shocked material from shallow depths in 

the primary target. Thus, whatever primary material is recovered 

from the secondary ejecta will contain impact melt or other high 

shock grades, whereas the local material in the secondary ejecta 

should be less shocked because of the lower kinetic energies of the 

secondary impacts. The proportions are not known, but locally 

derived material probably increases with distance from the primary 

crater. Some geologic maps distinguish a rim-material facies called 

cratered rim material (fig. 3.22) and identify the underlying unit of 

which it is composed (Ulrich, 1969).

The ejecta of young secondary craters commonly appears as rays 

or bright splashes of freshly exposed material (figs. 3.4A, C, 3.6,3.28). 

Most of this ray material was probably excavated from the local 

terrain by secondary impacts. Many investigators of rays have suspected that fine primary ejecta also composes parts of rays (for example, Baldwin, 1963, p. 358; Schultz, 1976a, p. 204; Pieters and others, 

1982). Absolute dating of Copernicus and Aristillus depends on the 

interpretation that the samples recovered on rays hundreds of kilometers from those craters contain primary ejecta (see chap. 13).

Wall material

Wall material is a mixed unit that includes materials of the 

terraces and all the slumps, debris, and impact melt that coat crater 

walls. Walls may expose precrater strata (fig. 3.29). Edges of the 

terraces may also contain precrater target rock, and tops may expose 

downdropped rim material. Depressions in the terrace tops, which 

are most common on the rears of the outward-tilted downdropped 

surfaces, commonly contain pools of flat-surfaced material now recognized as impact melt (figs. 3.2D, 3.23B; Howard, 1975; Howard and

* ;-* ' * ,'£. *

MiJ^

Wilshire, 1975; Hawke and Head, 1977a). Dribbles or cascades of 

additional melt or debris, commonly bounded by levees resulting from 

the flow, connect some terraces and the floor (fig. 3.30; Howard, 1975). 

These melt features, which are seen only on high-resolution photographs, merely contribute to the overall rough appearance of the 

walls as seen on telescopic photographs. Coarse wall hummocks are 

probably slumps (fig. 3.27A).

FIGURE 3.26. Ejecta of moderately fresh (Eratosthenian) crater Werner (70 km, 28° S., 3° 

E., above), superposed on crater Aliacensis (80 km, below), whose subdued walls, raised 

rim wreath, and central peak were probably similar to those of Werner when first formed; 

Ahacensis-radial ejecta is no longer visible. Rim nearest Werner is more highly degraded 

than distal rim. Orbiter 4 frame H-100.

B

FIGURE 3.27. Eastern ejecta of Tsiolkovskiy (180 km).

A. Ground-surge emplacement is indicated by diversion from concentric to radial lineations 

(arrows). Impact melt is visible at upper arrow and in box at bottom. Apollo 15 frame

M-757.

B. Impact melt, partly enclosed in box in figure 3.27A, covers radial ejecta. Apollo 15 

frame P-9580.
3. CRATER MATERIALS 49

FIGURE 3.28. Splashes of bright material dug from beneath darker surficial layers by 

secondary-impact craters of a crater out of scene to lower left. Crater in picture is about 5.5 

km across; near west rim of Gagarin. Apollo 15 frame P-8941.

FIGURE 3.29. Crater Euler (28 km, 23° N., 29° W.), showing wall materials weakly 

coalesced into terraces at bottom and uncoalesced in upper left. Precrater stratigraphic units 

are evident by albedo differences on upper right wall. Apollo 15 frame P-10274.

On many geologic maps, material slowly uncovered or displaced 

downslope on crater walls is mapped as a unit called bright slope 

material (Mutch, 1970, p. 166-169; Wilhelms, 1970b, p. 39). This unit 

was early recognized by geologic mappers and ascribed to renewal of 

exposure of lunar rock by downslope movement (Shoemaker and 

Hackman, 1962, p. 297; Shoemaker, 1965, p. 128); it was generally 

assigned a Copernican age regardless of the age of the underlying 

unit. A substantial content of fresh material has been confirmed by 

spectral studies (see chap. 5), and all high-resolution photographs of 

steep lunar slopes reveal mass-wasted debris deposits (figs. 

3.31 3.33). The time of exposure of materials in the colluvium at the 

bases of massifs at the Apollo 15 and 17 landing sites has been dated 

isotopically at tens or hundreds of millions of years, in contrast to 

several aeons for the source rock. The colluvium and slope debris 

contain material foreign to the source and introduced by impacts. The 

oldest (darkest) slopes probably have accumulated the most exotic 

material (fig. 3.31).

Peak material

The geologic unit "peak material" was probably derived from 

greater depths than any other lunar-crater material because the 

material in peaks formerly lay beneath the excavated part of the 

crater (fig. 3.19). Peaks range considerably in morphology from low 

hills and single centralized pinnacles, through multiple jagged peaks 

clustered around a center, to dispersed smaller arrays (figs. 3.2C-E, 

3.29, 3.32-3.35,4.2). All these forms are consistent with uplift mechanisms. Peaks probably formed during crater excavation by intensive

FIGURE 3.31. Bright streaks of fresh material and darker, longer-stabilized material on wall 

of crater Lalande A (13 km, 6.5° S., 10° W.). Floor is composed of material already 

accumulated by similar downslope movement. Apollo 16 frame P-5400.

FICURE 3.30. Cascade of impact melt (arrow) from terrace to floor on north wall of Copernicus; w, steep part of wall just below rim. Cascade has apparently drained melt pool (right of arrow end)

through groove (right of arrowhead) (Howard, 1975). Arrow is 5 km long. Orbiter 5 frame M-152.
50 THE GEOLOGIC HISTORY OF THE MOON

deformation of the central part of the target. The violent and extreme 

uplift suggested by several lines of evidence given above seems to be 

supported by the observation by Murray (1980) that some peaks seem 

to have toppled over the crater rim (fig. 3.34). Because of the intense 

deformation, peaks are structurally complex. Few, if any, peaks are 

volcanic constructs accumulated after cessation of cratering, contrary 

to the interpretations presented on many geologic maps and in much 

other earlier lunar literature.

Floor material

Some floor materials are coeval with peak materials, some 

slightly younger, and some much younger. Floor materials are 

diverse, and photographs should be consulted during their reinterpretation. Some constitute peaklike uplifts (figs. 3.33,3.35), which 

in many craters form an approximately annular pattern that foreshadows the larger rings of basins (fig. 3.2E; chap. 4). Separation of 

such lumps from central peaks may be more a matter of scale-dependent mapping convention than of genetic significance. Some hummocky floor materials bridge the gap between peak and wall, and may 

represent the toes of inward-displaced wall slivers or parts of the floor 

uplift (figs. 3.2E, 3.23, 3.29; Howard, 1975). Other floor materials 

consist of debris from the walls (figs. 3.31, 3.33).

Impact melt, formerly believed to be volcanic, partly covers many 

crater floors. Fissured floor materials that appeared on Lunar Orbiter 

5 photographs of such young craters as Aristarchus, Tycho, and 

Copernicus (figs. 3.2Z), 3.35) clearly were originally molten, and in the 

spirit of the times, many workers once considered them to be volcanic 

(Offield, 1971; Strom and Fielder, 1971; Pohn, 1972; Schultz, 1976b, p. 

72). Some smooth floor materials (fig. 3.32) and other plains materials 

(figs. 3.2Z), 3.36) retained volcanic interpretations longer than any 

other geologic-map units. The smooth "ponds" are superposed on the 

floor, wall, and rim, and thus were mapped not as crater materials but 

as special postcrater units (Milton, 1968a; Wilhelms and McCauley, 

1971); they were commonly ascribed to volcanic extrusions released or 

localized by the impact (Strom and Fielder, 1971). However, these 

smooth and fissured materials have been identified as impact melt by

detailed observations of the freshest and best photographed craters 

(figs. 3.23, 3.36; Shoemaker and others, 1968; Howard, 1975; Howard 

and Wilshire, 1975; Hawke and Head, 1977a; Gault and Wedekind, 

1978). Impact melt veneers much of the crater interior and rim, and is 

ponded in favorable depressions (fig. 3.36). Fissures that continue 

from the floor over many of the floor uplifts indicate that melt rock 

coats these uplifts as well (fig. 3.35; Howard and Wilshire, 1975). The 

fissures and many irregular pits suggest drainage of impact melt into 

underlying porous breccia (fig. 3.35; Howard, 1975), not volcanism 

(Hartmann, 1968). Some melt flows came to rest later than the peak 

(fig. 3.32), in confirmation of terrestrial observations indicating that 

peak formation is even quicker than melt solidification. The only 

internally generated materials associated with craters are the mare 

materials that flood many craters and, possibly, small amounts of 

other material, such as that in the floor of the endogenic crater 

Hyginus (Pike, 1976).

Impact-melted floor materials probably would provide a chemically representative sample of the target rocks because target materials are extensively homogenized during impact melting (Dence, 

1971; Grieve and others, 1974; Grieve, 1975; Simonds and others, 

1976a, b). Conversely, the material on older crater floors may be 

completely unrelated to the crater that contains them or to the subcrater material (Aliacensis, fig. 3.26). Even craters that lack recognized deposits of mare basalt or ejecta may be filled by debris derived 

from the walls or introduced gradually as fragments by random 

distant impacts. Such slowly accumulated floor materials may yield 

random samples of large parts of the lunar surface.

In summary, geologic mapping of the materials of lunar craters 

has proved to be generally accurate in delineating significant units 

and determining mutual age relations, but the interpretations of 

many mapped units have changed. The absolute-age differences 

detected among crater-material facies by stratigraphic relations have 

shrunk from millions of years to minutes. The materials have moved 

continuously closer to an almost exclusive impact interpretation and 

away from volcanic or hybrid interpretations, under the prodding of 

sample analyses and continued photogeologic, Earth-analog, and 

experimental studies.

FIGURE 3.32. Part of "lobster claw" peak of crater King (compare fig. 3.23A), showing 

superposition of impact melt in gap of peak. Apollo 16 frame P-4998.
3. CRATER MATERIALS 51

FIGURE 3.33. Crater Proclus (28 km, 16° N., 47° E.), showing terrace-free walls, slump from upper wall (S), coarse floor mounds 

probably produced by floor uplift, and fissured impact melt on floor (arrow). Apollo 17 frame P 2265.

FIGURE 3.34. Crater Anaxagoras (51 km, 73° N., 10° W.). Peak material apparently 

toppled over rim at arrow (Murray, 1980). Orbiter 4 frame H-128.
FiGURt 3.35. Impact melt and floor mounds of Copernicus; melt is superposed on mounds at arrows. Fissures indicate shrinkage of melt and, possibly, drainage into subsurface cavities. Orbiter 5

frame H-153.
3. CRATER MATERIALS 53

FicURE 3.36. Lineations and festoons indicating flow of impact melt in exterior pool of crater King (fig. 3.23A). Melt has flowed over rim material and collected in depressions (Howard and

Wilshire, 1975). Apollo 16 frame P-5000.


FIGURE 4.1. (OVERLEAF) Orientale basin and its deposits. Scene centered on farside. Double-ringed basins are Hertzsprung (above) and Apollo (below, with some mare fill). Painting by

Donald E. Davis, courtesy of the artist.
4. BASIN MATERIALS ORIENTALE

CONTENTS

	Page

Introduction ....................................................... 57

Orientale exterior, by John F. McCauley .............................. 66

General features ................................................ 66

Inner facies of the Hevelius Formation ............................. 66

Outer facies of the Hevelius Formation and secondary craters ....... 67

Relation between ground-flow deposits and secondary craters ...... 67

Plains .......................................................... 72

Orientale interior, by John F McCauley .............................. 72

Montes Rook Formation ......................................... 73

Maunder Formation ............................................. 76

Topography and gravity ............................................ 77

Origin of rings .................................................... 77

Introduction .................................................... 77

Genetic models ................................................. 77

Discussion ..................................................... 79

Conclusions .................................................... 81

Summary of basin-material origins .................................. 81

INTRODUCTION

Most simple craters, almost all complex craters, and all ringed 

basins were formed by impacts. Collectively, these lunar features 

constitute a series of impact excavations that have many common 

properties over the entire size range but also show many differences 

(figs. 4.1-4.3).

Appreciation of the importance to lunar geology of the largest 

members of this size-morphology series, the basins, has grown slowly. 

Gilbert (1893) may have been the first to perceive the Imbrium basin 

as the largest example of the series on the nearside. This interpretation was supported by other influential early workers who 

increasingly recognized the distinction between basin and mare 

(Baldwin, 1949, 1963; Kuiper, 1959; Hartmann and Kuiper, 1962; 

Mason and Hackman, 1962; Shoemaker, 1962a, b; Shoemaker and 

Hackman, 1962). Shoemaker and Hackman (1962) pointed out the 

similarity of the Imbrium-basin ejecta to crater ejecta. Several other 

authors enunciated the differences in style from craters, while affirming a basic similarity of origin, by stressing the linearity of the basinradial Imbrium sculpture (Baldwin, 1949, chap. 11; 1978) and the 

concentricity of the ring systems (Baldwin, 1949, p. 37-46; 1963; 

Hartmann and Kuiper, 1962; Hartmann, 1981). A landmark study of 

rectified telescopic photographs by Hartmann and Kuiper (1962) 

documented the ring patterns and regular ring spacings of 12 telescopically visible basins, including Orientale; they also coined the 

term "basin" (Hartmann and Kuiper, 1962, p. 62; Hartmann and 

Wood, 1971, p. 3-5; Hartmann, 1981). McCauley (1964b, 1967a, b, 1968) 

then pointed out similarities of the circum-Orientale terrane to the 

Imbrium ejecta. Titley and Eggleton (1964; Titley, 1967) tentatively 

recognized ejecta of the Humorum basin, and Wilhelms and Trask 

(1965) suggested that the terrane around other basins contains basin 

ejecta and secondary-impact craters. Stuart-Alexander and Howard 

(1970) and Hartmann and Wood (1971) reviewed the photographs 

obtained of the whole Moon in 1966 and 1967 by the five Lunar 

Orbiters and described the properties, distribution, and probable 

origins of many additional basins and their maria.

Although these studies established the existence and the impact 

origin of many large basins, the morphology and extent of the exterior 

deposits and secondary craters were poorly known before Apollo 

exploration began (Mutch, 1970; Wilhelms, 1970b). Volcanic and 

tectonic hypotheses still competed vigorously with impact hypotheses. The first new data that would eventually lead to the current 

impact model were the superb photographs of the Orientale basin 

obtained in 1967 by Lunar Orbiter 4 (figs. 1.9, 4.4). The Orientale 

deposits are less obscured by later material than are those of any 

other large well-photographed basin. The second and vital clue came 

from the discovery of complex impact breccia at every sampling site in 

the terrae, in particular, the Apollo 16 landing site, previously 

thought to be underlain by volcanic materials (see chap. 2; Howard 

and others, 1974). The third new development was the extension of 

impact interpretations to the rest of the terrae, an exercise summarized in this volume. More and more basins have been recognized 

and established as the sources of surrounding deposits and satellitic 

craters (pi. 3; fig. 4.3; tables 4.1-4.3).

FIGURE 4.2 (OVERLEAF). Crater-to-basin transition.

A. Antoniadi (135 km, 70° S., 1 72° W.), with both an inner ring and a small central peak.

Mare material fills depression inside inner ring and partly floods peak. Orbiter 4 frame

M-8. 

B. Compton(162km, 56°N., 105° E.), with conspicuous central peak and low but distinct

inner "peak ring." Cracks indicate floor uplift. Orbiter 5 frame M-181. 

C. Petavius (177 km, 25° S., 60° E.), with massive complex central peak and vague

ringlike configuration of low peaks between central peak and wall. Uplift extending to

middle of wall is indicated by grabens. Slight filling by dark material is evident along

walls. Orbiter 4 frame H-184. 

D. Gauss (177 km, 36° N., 79° E.), with entirely different floor structure from that of

Petavius, although both craters are of same size. Orbiter 4 frame H-165. 

E. Humboldt (207 km, 27° S., 81 ° E.), with central peak and line of peaks. Other rugged

features of original floor are higher in north than in south; cracks indicate uplift of entire

floor. Light-colored plains may be later fill; dark crescents are later mare fill. Orbiter 4

frame M-12. 

F. Landau (221 km, 43° N., 119° W.), with few distinct floor features; irregularly

depressed in northeast. Orbiter 5 frame M-12. 

G. Campbell (225 km, 45° N., 153° E.), with some offcenter rugged relief. An offcenter

depression, as in Landau, is here filled by a small mare patch. Orbiter 5 frame M-103

(rectified). 

H. Clavius (225 km, 58° S., 14° W.), with small offcenter peaks, probably partly covered

by later plains fill. Orbiter 4 frame M-124. 

/. Schwarzschild (235 km, 71° N., 120° E.), with rugged, arcuate, crescent-shaped floor

topography. Orbiter 4 frame M-92. 

/. Milne (262 km, 32° S., 113° E.), with complex single to triple arcs. Milne and

Schwarzschild are the most basinlike craters smaller than 300 km in diameter. Orbiter 3

frame M-121. 

K. Gagarin (272 km, 20° S., 149° E.), with indistinct ring-form structure, though larger

than Milne. Orbiter 1 frame M-l 16.

57
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FIGURE 4.3. Definite lunar basins at least 300 km in diameter, in order of increasing size and 

complexity (see table 4.1 )L

E. Mendeleev.
L. Apollo.

M. Coulomb-Sarton.

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64 THE GEOLOGIC HISTORY OF THE MOON

TABLE 4.1. Definite lunar basins at least 300 km in diameter

[Basins listed in order of increasing size and complexity. Names all in capitals have deposits napped on plate 3. Names are derived, where possible, either from the mare fill of a basin (for example, the Imbrium 

basin from Mare Imbrium) or from a name given by the International Astronomical Union, considering them as craters (for example, Apollo). Otherwise, they are named here from two superposed, unrelated craters:

Name used here (reference^ Earlier designation (reference) 

Sclil Mer-^ucchius (W79)- Basin near Schiller (HKS2J. 

Coulomb-Sarton (L78) Unnamed basin B (HW71). 

Freundlich-Sharonov (S73) 

Mendel-Rydberg (W79) SE. limb basin (HK62) . 

Keeler-Heaviside (S78) 

South Pole-Aitken (S78)--- - Big backside basin (common usage).

Best photographs: Photograph used in figure 4.3 is underlined; each includes one-half basin radius beyond topographic rim, where possible. HK62, telescopic photograph from Hartmann and Kuiper (1962); r, rectified; 

A16, Apollo 16 frame; M10, Manner 10 frame (Mercury-Venus mission that photographed the Moon djring flyby). Others are Lunar Orbiter frames.

Major rings: Diameters are listed for rings that are observed or inferred to extend more than half a basin circumference. Approximate radii are listed for partial segments that may not form complete rings. Additional small segments are plotted on plate 3. Topographic basin rim (also called main nm in descriptions) is underlined, queried where uncertain.

References: 849, Baldwin (1949); B63, Baldwin (1963); G93, Gilbert (1893); HK62, Hartmann and Kuiper (1962); HW71, Hartmann and Wood (1971); L78, Lucchitta (1978); M67, McCauley (1967ai b); S77, Scott and others 

(1977); S78, Stuart-Alexander (1978); SH62, Shoemaker and Hackman (1962); SH70, Stuart-Alexander and Howard (1970); T67, Titley (1967); WE77, Wilhelms and El-Baz (1977); WM71, Wilhelmi and McCauley (1971); W79, 

Wilhelms and others (1979). Chapter numbers refer to the present volume.] ^

Basin Figure Center (lat, long)

FREUNDLICH-SHARONOV 4.3P 18.5° N., 175° E.

Keeler-Heaviside 4.3T 10" S., 162° E.

4.4 

4.5

Best 

photographs

4 H-179 

4 M-193 

4 M-179

5 H-21 

4 M-ft

4 H-160 

4 M-148 

4 M-136

3 M-121 

4 M-8 

4 M-9 

4 M-82 

5 M-65

1 M-115 

1 M-136

5 M-29 

5 H-29 

5 M-25

5 M-65(r) 

2 M-75 

3-M-121 

4-M-82

4 H-168 

4 1-lfil

1 M-38 

1 M-28

5 M-103(r) 

5 M-TZT

5 H-30 

5 M-30

5 '1-5

5 M-?j(r)

5 H-26 

b H-28 

5 M-26 

b M-28(r) 

TM-24

2 M-34(r) 

b M-79

4 H-165 

4 M-23 

4 M- 174 

4 M-152 

4 M-165 

M10 2277

4 H-193 

4 M-193 

4 M-136

Telescopic ~t~VFn 

2 M-7b

4 M-143 

4~H-136 

4 M-136

4 M-9, 11 

2 M-196 

A16 M-3035

4 1-17

4 M-71 

4 M-83 

4 M-84 

HK62(r) 

4 M-52

4 M-9 

4 M"-ll 

4 M-118

4 M-194(r) 

5 M-19 

5 M-22 

See fig. 4.4

Major rings

Diameter Radius 

(km) (km)

300 

T5TJ

320 

T5T

325 

165

325

T7F

330 

W

330 

T50

340 

175

360

430

440 

220

275

445

505

530(?) 

£00

560,,, ...

570 

205 

255 

140

600

600 

600_

155

200

450

740 

420

780(?) 

^^ 270 

170

820(?) 

350 

280 

440 

325

480

330

360

860 

600 

450 

350

880 

275 

95

650 

930 

620 

480 

320

Description

Degraded craterlike rim. Discontinuous inner ring, connected to rim by outward-sloping 

shelf. Deposits obscured by Orientale.

Fresh craterlike rim and wall. Inner ring consists of rugged crags forming nearly complete circle, joined to wall in two places by exposed shelf. Hummocky probable impact 

impact melt on floor. Extensive craterlike deposits, linear secondary chains.

Outer ring craterlike though degraded; bounded inward by steep scarp, outward by gentle 

flank. Inner ring mountainous. No discernible deposits.

Poorly photographed; outer ring apparently craterlike, inner ring partial. No discernible deposits.

Craterlike outer rim with flat-bottomed troughs. Inner ring of discontinuous small 

peaks, connected to wall by outward-sloping shelf. Craterlike deposits and numerous, 

though subdued, secondary craters.

Craterlike rim. Discontinuous, mostly obscured inner ring of narrow ridges. Extensive 

craterlike deposits and secondaries.

Apparently craterlike rim. Continuous but battered lower inner ring. No deposits 

observed.

Craterlike rim. Low and ridgelike inner ring, continuous except where cratered. 

Craterlike deposits and secondary craters west of basin.

Obscured by Orientale; probably craterlike rim. Discontinuous rjgged inner ring. No 

deposits observed.

Craterlike rim locally divided by terracing. Inner ring partly buried but detectable 

for almost a complete circle. Deposits *<ell exposed to west; extensive secondaries.

Crude figure-8 pattern of outer and inner rings. Rim has crude echelon pattern suggestive of large-scale terracing. Inner ring complete for half of periphery; has 

steep inner scarp and broad, gentle outer flank. Deposits and secondaries poorly 

photographed but visibly reduce prehasin crater population 400 to 500 km from rim.

Interrupted obscured main rim. Bifurcations indicate terracing. Inner ring partly 

continuous and bounded by steep inward scarp and gentler outer flank, partly discontinuous and rugged. Concavity or discontinuity inside inner ring indicated by mare 

localization. Mare patches indicate outward slope of interring shelf and possible 

exterior "ings. Deposits dnd secondaries extend about 500 km southward.

Heavily obscured by Bi rkhof f and Lorentz deposits. Rings uncertain; 400-km-di ameter 

ring could he the rim.

Outer ring questionable; consists of discontinuous but large massifs. Inner, 325-kmdiameter, mare-bounding ring, more complete, though battered, may be the nain rim. 

Deposits not directly observed but may breach south rim of Keeler-Heaviside basin.

Craterlike rim in most sectors, irregular in northwest and indistinct in north. Half 

of complete and distinct innermost ring resembles crater nm; rest consists of rugged peaks. Intermediate rings conspicuous, single, and benchlike or ridgelike in 

most sectors; ragged, bifurcated, echelon in northwest; inconspicuous in north, where 

ragged ridges and hills form upland nearly at level of surrounding terrain. Deposits 

extensive, locally 600 km from rim; some secondary-crater lobes 900 km from rim.

Rim complete; seems craterlike but is poorly photographed. Ridges suggest interior 

ring structure, but connection into rings or terraces is uncertain. Distinct radial 

grooves and ridges and mantling of craters south of basin indicate deposits and secondaries extending at least 600 Kin from rim.

Figure-8 shape; each part of main rim about 600 km in diameter, equivalent in total 

area to a circle 700 km in diameter. Bel'kovich (?00 km diam) is either an indepenpendent crater or third part of inner basin. Intermediate ring, fairly distinct in 

south-southwest, consists of hills in southwest, many large hills in southeast, and 

nearly continuous hilly area in north about at level of surrounding terrain. External 

600-km-radius arcs southwest of main rim. Distinct deposits and secondary craters 

(<20 km diam) extend more than 600 km from southeast rim.

Distinct separation of rings in south, but continuity uncertain. Faint deposits in 

southwest.

Ring structure indistinct in west half except in Montes Haemus. Haemus, Taurus-Li ttrow 

massifs, and other margins of mare probably form main rim (Head, 1979b). "Vitruvius 

front" (Head, 1974b) and other irregular elevations form long arcs with 45J-kii radii. 

Mare ridges suggest a buried 42n-km-diameter ring. Basin may be double, with small 

northern and large southern components (Scott, 1972a; Wolfe and Reed, 1976); center 

as listed is that of combined components.

All rings and arcs composed of large but disconnected massifs whose connection as rings 

and relation to excavation boundary are uncertain. Deposits not observed.

820-km-di ameter ring defined by steep notched scarp in south and by discontinuous lower 

ragged, semiradial or semiconcentric ridges or massifs elsewhere. Part of most conspicuous mare-bounding ring (440 km diam) connects with 280-km-radius arcs as sloping 

flank, elsewhere is divided into semidistinct rings with bridges to 820-km-diameter 

ring. 325-km-diameter ring marked by terra shelf and mare bench. Most distinct ring 

separation is in south. Deposits and secondary craters distinct only in southeast.

840-km-diameter ring defined by partial arcs bounded by steep inner scarp and gentle

360-km-di ameter ring continuous but low. Deposits not observed.

Main topographic rim is well-defined scarp on west (Rupes Altai), discontinuous mareembayed mesas on east. Intermediate nng broad and low, bifurcates in west, forms 

discontinuous plateaus in east. Inner ring diverse, low. Possible additional ring 

defined by mare border, 240 km in diameter. Deposits extend 650 to 950 km, typically 

750 kn, from center; secondaries extend 700 to 1,500 km.

Outer ring includes large massifs. Intermediate, poorly defined ring contains massifs, 

locally connected by rugged structure to outer ring. Inner "ring" consists of 

closed series of irregular elevations. No deposits observed.

Most nearly complete concentric ring structure among large basins except in the west. 

Outermost conspicuous ring is Cordillera, bounded inwardly by steep scarp or gentle 

lip and topped by some ridgelike massifs. Semicircular depression is outside this 

ring, and a raised 650-km-radius ring may be east of basm. Extension of Cordillera 

west of basin merges witn extensions of Rook (620 and 480 km diam) rings as rugged,

References

Chapter 9; 

Chapter 10 

Chapter R; 

Chapter 8;

Chapter 9; 

Chapter 8; 

Chapter 8;

Chapter 8; 

Chapter 3; 

Chapter 9; 

Chapter 9;

Chapter S;

Chapter -i;

Chapter 3; 

Chapter 9;

Chapter B; 

Chapter 9;

Chapter 9; 

Chapter 9;

Chapter 8; 

Chapter 9;

Chapter 8; 

Chapter 9;

Chapter 3;

Chapters 4 

M67.

HW71 

; H70 

HK62 

HW71

WE77 

HW71 

HW71

HW71 

HK62

HW71 

HW71

SH70

HW71 

S78. 

SH70

S78. 

HK52

HK62 

B49.

S78. 

HK62

WE77 

B49,

SH70 

, 10;

, W79. 

, HW71 

, W79. 

, W79.

, S78. 

, L78. 

, W79.

, WM71 

, S78. 

, S7S.

, HU71 

, L78.

, HW71

, L78. 

, HW71

, T67,

HK62,

, WE77 

HK62,

, W79.

, S73.

, W79.

WM71.

WM71.

, W79. 

B63,

irregularly structured upland. Other sectors of outer Rook ring consist of large 

massifs with steep inner and gentler outer slopes in a sawtooth pattern. Large section of inner Rook r ing and small section of 320-km-diameter ring also have large 

massifs; otherwise, inner Rook ring consists of smaller massifs or a steplike lip, 

and 320-km-diameter ring is defined by a lip or gradual fallofr. Knobby hills suggest a plateaulike central uplift. Radial structures in inner oasin. Extensive 

ejecta and secondaries.
4. BASIN MATERIALS ORIENTALE 

TABLE 4.1. Definite lunar basins at least 300 km in diameter Continued

65

Basin Figure Center (lat, long)

Best 

photographs

Mdjor

Di ameter 

(km)

rings

Radius Description 

(Km)

References

South Pole-Aitken ---

4.3Z 17.5° N., 58.5° E. 

4.9

4.3AA 33° N., 18° W.

1,060

1.160 570~

450

380

Outer ring consists of tnree differently expressed segments: rimlike scarp in north, 

ridge bordering a trough in northwest, and irregular scarp in southeast. Elevations 

fall off inward from this ring, which also marks start of secondary chains. Ridges 

inside northwest ridge and massifs inside southeast scarp are about 380 and 4bO km 

in radius, respectively, but are not closely concentric. Main mare-bounding nm is 

most conspicuous structure, divisible into 500- and 635-km-diameter rings, the inner 

consisting of large massifs with missing sector in east, and the outer of nigh but 

irregular massifs and other ridges and hills, separated from inner ring by deep irregular depressions. Innermost ring defined by a mare bench and small hills. Deposits 

and secondary craters extensive to east (1,600 kin from center?), secondaries extensive 

to north (900 km) and south (1,100 km).

Center as shown and 1,160-km ring are formed by Montes Carpatus, Apenninus, and Alpes. 

North shore of Mare Frigoris has 900-km radius from same center and is another major 

rim; both parts connect with Montes Caucasus. Lineations in Caucasus radial to same 

center. Montes Recta and Spi tzbergensis and rnare ridges define 670-km-dia,neter ring 

about a more northerly center; other connections are less certain and not given here. 

Lineate ejecta is north of Frigoris and soutn and southeast of Carpatus-Apenninus (Fra 

Mauro Formation); knobby ejecta (Alpes Formation) is closer to r ings in all sectors, 

imbricate on south Apenninus flank (material of Montes Apenninus). Secondary craters 

abundant in all exposed sectors to 1,400-2,700 Kin from center, typically to 2,300 km.

Huge massifs and scarps define single rim averaging 2,200 Km in diameter, or two rings 

about 2,500 and 1,800-2,000 km in diameter.

Chapter 9; HK62, SH70, 

WM71, WE77.

Chapter 10; G93, B49, B63, 

HK62, SH62, WM71.

Chapter 8; S78, W79.

TABLE 4.2. Probable and possible lunar basins at least 300 km in diameter

[Basins naned by convention described in table 4.1. All basins a^e pre-Nectanan (see chap. 8) except possibly Sikorsky-Rittenhouse (Nectarian?). Diameter refers to outer or only observed ring.

References: B53, Ba n dwin '1953); Bfi9, Ba'dwin [ 11691; C74-R1, Cadogan '1974, 19811; HH77, Hawke and Hea-J [1977bl; MAS1, Maxwel 1 and Andre .'19S1); SH70, Stuart-Alexander and Howard (1970); WE77, ' lilhelms and El-Baz 

(1977); WM71, Wilhel'ns and McCauley (1971); W79, WilheVns and others (1979); *I81, Wh-'.aker (1931)]

Center (lat, long)

300

310

'dence and status as basin

Arc.iate radges adjacent to Mpndel-Rydberg; possibly part of that basin-- ------

Craterlike rim and inner hills. Buried by Scnrodinger deposits; possibly 

part of South Pole-Aitken basin.

Craterlike rm and trace D c inner -Tig. Buried by Schrodinger deposits; 

possibly part of So.jtn Pole-Ait<en basin.

Concentration of plains nside and around rugged arcuate hills. Two rings 

 nay be exposed. Probable hasm.

B63, p. 195; named here.

Balmer basm of 1/W1; renamed here.

\fil; similar to Gargantuan basin of C74-81.

Because of the pivotal role played by investigations of Orientale 

both before and after the Apollo 16 mission, Orientale deposits are 

taken as models for others and are thoroughly described in this 

chapter. Similarities to the Orientale deposits appear repeatedly in 

the rest of the terrae and reveal the presence of basin deposits in all 

degrees of preservation from the fresh to the barely discernible.

Although all primary-impact excavations are surrounded by 

basically similar ejecta and secondary craters, their interior morphologies differ sharply from the small to the large end members of 

the impact series (figs. 3.2, 4.2, 4.3). Morphologic complexity 

increases most obviously in the simple-to-complex transition at diameters of about 20 km and again at diameters of about 300 km. Like the 

first, this second transition has exceptions. All unburied lunar impact 

cavities larger than 300 km in diameter display interior rings; 

however, some cavities smaller than 300 km also contain rings. 

Antoniadi (135 km diam) and Compton (162 km diam) each have a

peak centered inside an inner ring. Schwarzschild (235 km diam) has 

a partial ring, and Milne (262 km diam) has several. Several craters 

display ringlike patterns of internal peaks (figs. 3.2#, 4.2C-E). Some 

tectonically modified craters appear to possess a concentric structure 

(see chap. 6). On the other hand, several partly buried circular depressions larger than 300 km in diameter do not display the expected 

second ring.

These deviations muddy the definition of the term "basin." All 

authors agree that cavities larger than 300 km in diameter that have 

concentric rings instead of central peaks are basins, but opinions 

differ about whether this definition should be based entirely on size or 

on the presence of a second ring regardless of size (Wood and Head, 

1976). This chapter concentrates on excavations at least 300 km in 

diameter, and classifies those without observed interior rings as 

indefinite basins (table 4.2). The problem of nomenclature for transitional features smaller than 300 km is not resolved here.
66 THE GEOLOGIC HISTORY OF THE MOON

The origin of rings is among the most important unsolved problems of lunar geology. Orientale is also critical to this problem. Its 

four major rings are the most regularly spaced on the Moon (fig. 4.3). 

The 930-km-diameter Montes Cordillera ring, partly scarplike and 

partly topped by massifs, bounds the basin (figs. 4.1, 4.4). Two jagged 

massive rings, collectively called Montes Rook, have diameters of 480 

and 620 km and lie inside the Cordillera. The fourth, innermost ring, 

320 km in diameter, consists partly of scarps and partly of massifs. A 

low fifth ring, 1,300 km in diameter, outside the Cordillera has also 

been suggested (Hartmann and Kuiper, 1962). Each ring is 1.3 to 1.5 

times larger than the one inside it, an interval that has been generalized as the square root of 2 (Hartmann and Wood, 1971). Although 

they are more regular than they are typical, as I will show below 

(compare figs. 4.4, 4.5), the Orientale rings are generally taken as the 

model for the spacing and formation of other rings in most of the 

literature, including the section below entitled "Origin of Rings."

Central to the controversy is the identity of the boundary of 

excavation. Some investigators, including myself, believe that the 

excavation cavity of basins which corresponds to the excavation cavity 

of simple craters is approximately the present basin itself, that is, the 

depression. The topographic basin rim that bounds most of the basin 

depression thus corresponds to the rim crest of simple craters. The 

topographic rim of Orientale is the Cordillera. In this hypothesis, the 

topographic rim bounds the excavation cavity and marks the inner 

boundary of the most voluminous ejecta. Many other investigators, 

however, believe that the primary, original cavity was smaller than 

the present basin. They identify the boundary of the excavation cavity 

with one of the interior rings or place it at a radial position formerly 

occupied by a raised ring and now by a trough. Stratigraphic questions that could be better answered by knowing the identity of the 

boundary include: (1) the position of the lithologic boundary between 

the ejecta and the interior deposits, (2) the volume of excavated 

material, (3) the ejecta's source depth in the terra crust or mantle, (4) 

the amount of melt generated by the impact, and (5) the position of the 

facies change between primary ejecta and ejecta of the secondary 

craters.

ORIENTALE EXTERIOR

By John F. McCauley

General features

The exterior ejecta deposits of Orientale form nearly concentric 

zones similar to the concentric facies of crater materials. In proportion to cavity size, the extent of morphologically similar facies is the 

same for both basins and large craters (Moore and others, 1974; 

Morrison and Oberbeck, 1975). A continuous deposit extends outward 

an average distance of 450 km, or one basin radius, from Montes 

Cordillera. Beyond this deposit, clustered secondary craters accompanied by their ridgelike and planar ejecta form a more varied terrane of discontinuous deposits. In some sectors, dense concentrations 

of secondaries extend 900 km (two basin radii) from the Cordillera; 

secondaries also extend beyond 1,850 km in at least one raylike string 

(fig. 4.6; Wilhelms and others, 1978, p. 3752), much like secondary 

clusters along the bright rays of young craters. The most conspicuous 

deposits of all types are shown on plate 3, and the continuous and 

discontinuous deposits are distinguished.

The continuous deposits closest to the rim and the discontinuous 

deposits at greater distances are collectively designated the Hevelius 

Formation, the first Stratigraphic unit whose relation to Orientale 

was recognized (McCauley, 1967a, b). This definition was based on 

telescopic observations and was originally applied to a tract centered 

around lat 2° N., long 68° W., that is now considered part of the outer 

facies. Scott and others (1977) later expanded the definition to include 

the continuous, more coarsely textured blanket farther west, which 

had also been detected telescopically (McCauley, 1964b). The Hevelius 

and other Orientale materials, except the plains materials also discussed here, collectively constitute the Orientale Group (McCauley, 

1977; Scott and others, 1977).

Inner facies of the Hevelius Formation

The continuous deposits of the Hevelius Formation form an enormous doughnut-shaped zone, elongate north-south and from 300 to at

least 600 km wide, outside Montes Cordillera (pi. 3; figs. 1.9, 4.1, 

4.4A-G; Scott and others, 1977). This zone is difficult to delineate 

west of Orientale because of the obliquity and poor quality of the 

Lunar Orbiter 5 photographs of that area (pi. 3B; fig. 4.5). However, 

this part of the Hevelius, called by Scott and others (1977) the inner 

facies of the Hevelius Formation, is generally the most distinctive unit 

of the Orientale Group. It displays abundant attributes of a massive 

ejecta deposit that acquired a distinctive textural imprint during flow 

along the rugged prebasin surface.

The inner, continuous Hevelius is texturally characterized by 

elongate ridges and irregular troughs generally radial to Orientale; 

figure 4.4C shows a typical locality. The ridges are curvilinear and 

mutually crosscutting; some ridges have nearly streamlined forms, 

and others appear ropy or filamentary. These landforms clearly 

resulted from surface flow of deposits differentially affected by obstacles. Some of the larger ridges have lobate edges that resemble the 

fronts of viscous flows of lava or debris (fig. 4.4D). Numerous narrow 

linear grooves are also visible, particularly in the lower and smoother 

parts of the deposits. Another characteristic are the myriads of vague 

elongate depressions and chains of depressions, visible only in good 

photographs, that are overlain by ridged and grooved materials (fig. 

4.4C). The shapes, subdued appearance, and distribution pattern of 

these depressions indicate that they are secondary craters, partly or 

completely buried by slightly later flows of ejected deposits. Although 

these buried secondary craters contribute a negative-relief pattern, 

the generally larger positive forms dominate the scene. Their distribution pattern and topography argue for strong centrifugal depositional movements within the materials that make up the inner, 

continuous Hevelius.

The inner facies of the Hevelius Formation generally forms a 

continuous ground cover that deeply buries or subdues large primary 

craters on the pre-Orientale surface (figs. 4.1, 4.4F, G). The depth of 

burial of these craters generally decreases outward from the basin. 

Estimates based on the depth of crater burial indicate that the inner 

facies of the Hevelius decreases in thickness from 3 to 4 km near 

Montes Cordillera to no more than several hundred meters at its 

outer edges (Moore and others, 1974). However, these estimates 

assume that the buried craters had the dID ratios of fresh craters, 

which compose only part of any random population of lunar craters. 

Theoretical models of ejecta thickness (McGetchin and others, 1973) 

are also very uncertain in view of the number of simplifying assumptions that underlie them (Solomon and Head, 1979, p. 1669). Calculations are hampered by uncertain cavity volumes, uncertain degree of 

bulking of the ejected material (Croft, 1978), and the difficulties of 

scaling from one gravity field to another (Gault and others, 1975). 

Basin deposits are visibly irregular and lobate (pi. 3), and are not 

equally thick all around a basin at a given radial distance. Thus, their 

thicknesses are poorly known but are probably less than 4 km.

The well-developed radial or semiradial trends prevail within 

most of the inner facies of the Hevelius except in a sector of about 50° 

of arc east of the basin (figs. 4.4E-G). This anomalous region has little 

radial texture and is part of the concentric facies of Moore and others 

(1974), who noted an outward gradation here and elsewhere from this 

concentric facies to more nearly radial material (top, fig. 4.4F). The 

inner Hevelius in the anomalous area (fig. 4.4F) has stubbier ridges 

and lobes than in more typical localities (figs. 4.4C, D, G). Large 

patches of smooth to rolling weakly lineate terrain are also present, 

some of which show numerous small ridges and troughs concentric 

with the basin. In addition, the depressions formed by buried secondary craters are not so large and do not contribute so significantly to 

the surface texture as in more typical parts of the inner Hevelius (figs. 

4.4C, D, G). The long axes of many of these depressions are also 

concentric rather than radial to the basin. Prebasin craters are not so 

deeply buried as in the more strongly radial parts of the Hevelius 

Formation at comparable distances from the basin. All indications in 

this sector point to ejection of less material at lower velocities than 

elsewhere.

Small patches of what has been described as the fissured facies 

and interpreted as impact melt (Moore and others, 1974) occur within 

the inner facies of the Hevelius. These smooth but locally fractured 

deposits lie in depressions and contrast sharply with the topography 

of their surroundings (fig. 4.4S); they also are somewhat darker than 

the adjacent material. Their patterns of fissures resemble those 

caused by shrinkage of a coherent material. These small deposits 

(areas less than a few tens of kilometers on a side) are widely scat-
4. BASIN MATERIALS ORIENTALE 67

tered, and are included in the inner facies of the Hevelius on the map 

by Scott and others (1977) and in the inner Orientale unit on plate 3. 

They resemble the smaller impact-melt pools on the flanks of craters 

(figs. 3.23, 3.27, 3.36).

In mapping the Orientale region at 1:5,000,000 scale, Scott and 

others (1977) distinguished a transverse facies of the Hevelius Formation, also included here in the inner Orientale zone (pi. 3). This 

transverse facies consists of patches of closely spaced, ropy-looking, 

intertwined ridges and intervening troughs arranged concentrically 

around Orientale. Typical dimensions of these patches are about 30 

by 60 km. They are common on the far sides of crater floors (figs. 

4.4F-H) 300 to 600 km from Montes Cordillera. They seem to be more 

abundant on the east side of the basin, but this apparent asymmetry 

may be due to the poor photographic coverage of the west side (fig. 

4.5). These subparallel features, which have been termed "deceleration dunes" (McCauley, 1968), are good evidence for the surface 

flowage of Orientale ejecta. In the crater Riccioli, transverse ridges 

are banked up on the far side of the crater floor, beyond which a 

transversely textured flow lobe, 100 km long, emanates from a breech 

in the distal crater wall (fig. 4.4H). The influence of topography on 

several patches of less well defined transversely textured materials is 

less obvious (figs. 4.4F-H).

The inner boundary of the Hevelius Formation has been perceived both as sharply demarcated by Montes Cordillera (Baldwin, 

1974a; Wilhelms and others, 1977; Hodges and Wilhelms, 1978; Murray, 1980) and, alternatively, as irregular, gradational, and not everywhere coincident with Montes Cordillera (McCauley, 1977; Scott and 

others, 1977). Much of the Cordillera scarp is draped by the Hevelius 

Formation (figs. 4.4A-F). Locally, however, sharp breaks in texture 

mark the inner contact of the Hevelius along the steeper parts of the 

Cordillera scarp, particularly where this scarp is shadowed on photographs (fig. 4.4E).

Outer facies of the Hevelius Formation and 

secondary craters

The outer facies of the Hevelius Formation (Scott and others, 

1977), like the outer parts of all ejecta blankets, exhibits a greater 

variety of surface textures than the inner facies. The facies boundary 

is gradational but readily detectable on good photographs as the limit 

of the large radial lobate ridges typical of the inner facies of the 

Hevelius (figs. 4.4G, H). This contact is reproducibly mappable within 

a few tens of kilometers on the nearside but is hard to detect on the 

farside. Subdued, elongate, sinuous to irregular secondary-crater 

depressions typical of the inner facies are replaced in the outer facies 

by crisp secondary craters that are not overrun by ejecta. Buried 

prebasin primary craters are also more conspicuous in the outer than 

in the inner facies. Large patches of smooth plains also become 

abundant near the facies boundary and interrupt the radial depositional patterns. The boundary is analogous to that between the continuous inner deposits and the herringbone-patterned fields at one 

crater radius from the rim crest of craters.

The tract near the type area of the Hevelius Formation 

(McCauley, 1967a) contains faintly lineate, thin, and discontinuous 

material that mantles the wall of the crater Hevelius nearest to 

Orientale and grades to smooth plains on the crater floor. The adjacent high ground shows numerous relatively straight grooves and 

ridges, in contrast to the larger and more sinuous topographic forms 

of the continuous Hevelius. Some of these smaller and straighter 

grooves locally resemble mantled grabens but were probably created 

during emplacement of the Hevelius Formation (loc. 2, fig. 4.4H). 

Obvious depositional patterns are subordinate or absent, and this 

"lineated terra material" (Wilshire, 1973) or "grooved facies" (Moore 

and others, 1974) has been considered to be partly of erosional origin. 

However, depositional origins and erosional or secondary-impact origins are difficult to distinguish where the grooves are separated from 

terrain with more diagnostic features (Wilhelms, 1980).

In many places, the outer facies of the Hevelius Formation is 

rolling to hummocky and furrowed, particularly near the crater 

Criiger and eastward to Oceanus Procellarum (fig. 4.4H). This terrane generally shows only vague radial lineations and elongate 

depressions, as well as numerous subdued pits. Although a secondary-impact origin for this hilly and pitted terrain was mentioned 

(McCauley, 1973), a volcanic origin was preferred before the Apollo 16 

results were obtained (Wilhelms and McCauley, 1971; McCauley, 1973;

Wilshire, 1973). Continued studies have shown that the strangetextured terrane is laterally continuous with Hevelius deposits similar to those shown in figure 4.4F (Scott and others, 1977). The terrane 

probably consists of a pre-Orientale surface mantled by Orientale 

ejecta that had less radial momentum or velocity than other parts of 

the blanket.

Although some deposits and grooves are of primary-ejecta origin 

and others of doubtful origin, the outer zone of Orientale materials 

mapped here (pi. 3A) is clearly characterized by the products of 

secondary impacts. Obvious secondary craters large enough to map at 

1:5,000,000 scale were considered by Scott and others (1977) as a 

separate facies of the Hevelius Formation, but secondary craters are 

not mapped separately from deposits here (pis. 3, 6-8). Secondary 

craters of Orientale were described by Offield (1971), Wilhelms and 

McCauley (1971), Moore and others (1974), and Wilhelms (1976) (see 

chap. 3). The Orientale secondaries overlap in chains and clusters 

that are scaled-up equivalents of the secondaries described around 

Copernicus by Shoemaker (1962b). Diameters of individual craters 

are as large as 28 km, but most are from 7 to 20 km (Wilhelms and 

others, 1978), typically about 10 km (figs. 3.9, 4.4C, G, H). Although 

some distant clusters have downrange herringbone patterns (fig. 3.9), 

these patterns are subordinate to straight ridges nearly radial to the 

basin (figs. 4.4G, H). Bowl-shaped to elongate gougelike features are 

typical of the Orientale-secondary field. Their downrange sides show 

wavy or partly braided deposits that partly or completely obscure the 

underlying topography. Many individual craters have well-developed 

rims and appear to be moderately fresh; others are very subdued and 

have been mistaken for prebasin craters (figs. 4.4G, H; Wilhelms, 

1976). Orientale secondaries, including some that extend great distances from the basin (fig. 4.6), are abundant beyond a large lobe of 

the Hevelius Formation southeast of the basin (pi. 3; fig. 4.4G). Secondaries are especially large and abundant north of the basin (fig. 

3.9) and on the poorly photographed west side (fig. 4.5), but are rare in 

the concentrically textured sector east of the basin (figs. 4.4E-G).

The outer facies of the Hevelius Formation, like the outer parts of 

all impact-crater deposits, represents a progressively thinner to more 

discontinuous cover that may locally extend to great distances from 

Orientale without leaving a perceptible imprint on the topography. 

The problem of identifying the outermost limit of ejecta is intrinsic to 

mapping all impact and explosion craters. Here, the outer contact is 

placed where no further linear or Orientale-radial structure can be 

observed on Lunar Orbiter 4 H-frames (pi. 3). Beyond this contact, 

only isolated secondary craters and crater clusters are recognizable 

as related to Orientale (figs. 4.6, 7.7,10.5). Farflung Orientale ejecta 

could have reached points thousands of kilometers from the basin 

(Chao and others, 1975; Moore and others, 1974). It may even be 

concentrated in the furrowed and pitted terrain antipodal to Orientale near Mare Marginis (fig. 4.7; Schultz and Gault, 1975a; Wilhelms 

and El-Baz, 1977).

Relation between ground-flow deposits and 

secondary craters

As discussed above, numerous vague and irregular depressions, 

gouges, and clusters of radially lineate, buried secondaries occur 

within both the outer and inner facies of the Hevelius Formation (figs. 

4.4C, D, F, G). Some chainlike secondary craters buried by the continuous Hevelius extend to the Cordillera ring (fig. 4.4D), as do similar 

chains at the double-ring Schrodinger basin (fig. 1.5). Surficial flow 

textures characteristic of the Hevelius Formation overlie these 

chains; thus, the chains were probably formed by early-arriving clots 

of ejecta traveling in relatively low ballistic trajectories from the 

basin, and were later overrun by a slower-moving ground surge (Morrison and Oberbeck, 1975; Oberbeck and Morrison, 1976). Morrison 

and Oberbeck (1975) suggested that most of the other radial and 

concentric depositional ridges of the inner facies of the Hevelius were 

created by interference and interaction of ejecta, as are the herringbone patterns around craters. Such an origin would imply that 

the ridges contain little primary basin ejecta. This process probably 

did form many of the ridges in the outer parts of the secondary field 

(figs. 3.9, 4.4G, H\ but not the long sinuous ridges and grooves of the 

inner facies of the Hevelius, which were deflected by obstacles in a way 

that indicates surface flowage of thick ejecta (fig. 4.4G). Furthermore, 

the burial of the secondaries indicates that primary ejecta constitutes 

most of the inner facies; any secondary ejecta present there must have
68 THE GEOLOGIC HISTORY OF THE MOON

100°

80° ,

70°

r^f/Jwat'l *VVv'>' !' r'*V I

:^t^^j4;.*

>'.V,:^ t N^l^lt wv 4

illlffe ^"' T^'d?X^*^'^"Ci

r ^*'yi v *'*\ '-'Kn ?;V '''A

^^^Mifia

-10°

-20°

100°

400

I

600

I

800 KILOMETERS

_I

FIGURE 4.4. Orientale basin and periphery.

A. Basin structure. 1, Cordillera ring (930 km); 2, outer Rook ring (620 km); 3, inner Rook ring (480 km); 4, innermost ring (320 km). After U.S. Geological Survey (1978).

B. Central Orientale basin, showing the Hevelius Formation (1), fractured impact melt (2), 

overlap of the Monies Rook Formation onto the Hevelius (3), Cordillera ring overlain by 

Hevelius (4) and Monies Rook (5) Formations, typical knobby Monies Rook Formalion (6), ouler Monies Rook ring massifs (7), inner Rook ring (8, 9a, 9b), massifs of 

320-km-diameler ring (10), fissured domical Maunder Formalion burying part of 320- 

km-diameter ring (II), and level Maunder Formalion cul by grabens (12). Inner Rook 

ring massifs in soulh (9b) are widely separated by ihe Maunder Formalion. Perched 

former higher slands of level part of ihe Maunder indicate subsidence (13). Onenlalebasin materials are exposed from beneath shallow mare near center of basin (14). Circular 

slructure of outer Rook ring (15) is possibly derived from prebasin crater. Crater 

Maunder overlies basin and mare (16). Orbiter 4 frame H 195.
C. Typical continuous inner facies of the Hevelius Formation (bottom half of 

photograph). Many ridges are flow features (1); depressions are buried secondary craters (2). Large secondaries (3) emerge among smooth deposits (4) farther 

from basin. Fissures indicate possible melt deposits (5). Complex landforms were 

generated by multiple near-simultaneous secondary impacts (6). Overlaps with 

figure 3.9A. Orbiter 4 frame H-188.
D. Cordillera ring (1) south of basin, overlain by chain of Orientalesecondary craters, Vallis Bouvard (2). Hevelius Formation mantles 

the secondary chain and extends outward as a thick blanket (3). Lobes 

of nonlmeate (4) or leveed (5) material continue the Hevelius deposit 

outward. Orbiter 4 frame H 186.

E. Monies Rook Formation, decelerated by Cordillera ring into dunehke 

(1) or radially textured (2) deposits. Monies Rook or Hevelius 

Formation, with radial textures, overlies Cordillera scarp (3). 

Maunder Formation is at left (4), inside outer Rook ring (5). Fissured 

ejecta (6) is probable impact melt. Orbiter 4 frame H 181.

F. Radial (1) and concentric (2) textures of the Hevelius Formation directly east of 

Cordillera ring (3). Flow of the Hevelius was blocked by crater Rocca and piled up as 

dunelike mounds (4). Orbiter 4 frame H 173.

C. Transition between inner (1) and outer facies of the Hevelius Formation. Outer facies 

includes secondary craters (2), ridged deposits of secondaries (3), and planar deposits 

(4). Flow of the Hevelius was diverted by now-buried obstacles (5) and blocked by 

crater Inghirami (6). Some secondary craters filled by deposits look deceptively old (7). 

Orbiter 4 frame H-172.

H. Type area of the Hevelius Formation, centered at 1.7° N., 68.5° W. in crater Hevelius 

(1). Linear troughs radial to Orientale are partly filled by deposits (2). Deposits were 

decelerated by rim of crater Riccioli (3), but lobe of textured deposits extends northeastward of rim (4). Lobe of plains deposits inside Grimaldi (5), playout of decelerated 

deposits (6; same as at loc. 4, fig. 4.4F) into planar deposits (7), and pitted terrain (8) 

are also related to Orientale. Postbasin features: grabens (9), crater Criiger (10), and 

other mare-filled depressions (11). Mosaic of Orbiter 4 frames H 168 (bottom) and 

H-169 (top).

72 THE GEOLOGIC HISTORY OF THE MOON

been overwhelmed by the primary ejecta. This question of the proportions of primary versus secondary ejecta in basin deposits is explored 

in discussions of the Apollo 14 and 16 samples in chapter 10.

The photographic evidence, therefore, suggests an early stage of 

secondary cratering followed by a ground-hugging debris surge that 

first inundated and then mixed with the earlier debris. The lastformed Orientale features, however, may also be secondary craters. 

Unusually numerous, crisp circular craters, 10 to 15 km in diameter, 

which at first glance appear to be primary craters younger than 

Orientale, are present in the west-limb region of the Moon (figs. 4.1, 

4.4, 4.5). These may be very late secondary craters excavated by 

projectiles that were ejected from the Orientale basin at high angles

FIGURE 4.5. West and northwest sectors of Orientale basin (O) on farside (compare fig. 

4.1), showing radial pattern of continuous inner facies of the Hevelius Formation (he) and 

secondary craters (sc). CS, Coulomb-Sarton basin; L, Lorentz basin. Orbiter 5 frame 

M-16.

and that almost reached escape velocity (as proposed by Shoemaker 

and others, 1968, p. 40-44, for craters on the rim of Tycho). Circular 

craters, presumably formed by high-velocity impacts, should be 

widely dispersed over the Moon not only from Orientale but also from 

all the older basins (figs. 4.6, 4.7).

Plains

The Hevelius Formation is gradational with the terra plains. 

These plains are abundant in the transition zone between the inner 

and outer facies of the Hevelius and occupy additional large tracts in 

the outer facies. These relations have been much discussed since the 

discovery by Apollo 16 of impact breccia in similar light-colored plains 

in the central lunar highlands. Volcanic origins have been discounted 

in all but a few areas (see chaps. 2,10; Neukum, 1977). The principal 

currently competing hypotheses are an origin as primary ejecta and 

as ejecta from secondary craters. Relations at Orientale provide the 

best photogeologic clues to the plains' origin.

Some nearly planar deposits have a distinguishable relief that is 

an important clue to their origin. Ridges of the inner facies of the 

Hevelius grade into faint rises in the plains of the transition zone 

(figs. 4.4G, H). Flat-topped lobes with planar or faintly textured 

surfaces (mapped as a separate nonlineated facies by Wilhelms and 

others, 1979) seem to have oozed from the more rugged Hevelius (fig. 

4.4D; Eggleton and Schaber, 1972); some of these smooth deposits are 

leveed (loc. 5, fig. 4.4D; Hodges and others, 1973; Moore and others, 

1974). Fluidlike flow of the semiplanar material is indicated. A substantial content of impact-melted rock may be responsible, although 

fluidlike flow of completely clastic debris is not excluded. Some textured deposits are superposed on secondary craters (figs. 4.4D, F). 

Little secondary ejecta can be contained in textured deposits with 

these stratigraphic relations.

By extension, fluidlike emplacement of material segregated from 

the textured ejecta blanket is suggested for the completely planar 

deposits that have ponded in depressions (Eggleton and Schaber, 

1972). This origin is even likely for such deposits as those in the crater 

Wargentin (fig. 1.9), whose apparently thick fill was long considered to 

be evidence for lunar volcanism. The light-colored deposits in Wargentin lie just beyond the zone where the continuous Hevelius grades 

into the plains. Although Wargentin may have been partly filled by 

volcanic materials, the present surface deposit is probably of Orientale ejecta (Schultz, 1976b, p. 92).

Emplacement as secondary-crater ejecta, as hypothesized for all 

light-colored lunar terra plains by Oberbeck and others (1974,1975), 

is likely for plains that lie far from the inner facies of the Hevelius but 

adjacent to clusters of secondary craters (fig. 7.6). The dispute 

between a primary and secondary origin for terra plains, therefore, is 

one not of exclusive genesis but of relative contributions by the two 

mechanisms.

The plains illustrate that the appearances and distributions of 

lunar geologic units may be consistent with more than one genetic 

interpretation. A further example is in the south-central highlands of 

the nearside, where careful mapping showed that plains and mantling deposits become less conspicuous from west to east (Cummings, 

1972; Mutch and Saunders, 1972). The mappers interpreted the differences as due to variations in the thickness of the deposits, which 

they considered to be a function of the duration of volcanism. An 

impact interpretation, more consistent with current understanding 

of the lunar terrae, is that the plains are ejected materials which thin 

away from their source, Orientale.

ORIENTALE INTERIOR

By John F. McCauley

All or almost all the materials of the Hevelius Formation lie 

outside the topographic basin rim of Orientale (pi. 3). Additional 

Orientale deposits occur mostly in the basin interior or on the Cordillera rim close to the basin. Two major units are named the Montes 

Rook Formation and the Maunder Formation (fig. 4.8); these formations complete the Orientale Group (table 4.3; McCauley, 1977; Scott 

and others, 1977). Morphologic and, presumably, genetic analogs 

occur in craters (figs. 3.2Z), 3.32-3.35) but are less conspicuous than 

in Orientale. Their analogs in other basins are less obvious than those
4. BASIN MATERIALS ORIENTALE 

TABLE 4.3. Geologic units of lunar basins

73

[Fm, Formation; n/n, recognized but not named; n/r, not recognized. Most units of other basins are named informally. References: H66, Hackman (1966); H74, 

Head (1974c); M74, Moore and others (1974); S71, Stuart-Alexander (1971); S77, Scott and others (1977); WM71, Wilhelms and McCauley, 1971; W79, Wilhelms 

and others (1979)]

Location Morphology Orientale (references) Imbrium (references) Nectaris (references) Interpretation

Inner basin Wavy, fissured 

Do Level, fissured-- 

Knobby -

- Maunder Fm (S77); corrugated 

facies (H74).

- Maunder Fm (S77) -

n/rShelf and inner 

flank.

Rim and other Massifs 

rings.

Inner flank Concentric, 

wreathlike.

Flank Radial, coarse 

Wavy, fissured---

Apennine Bench Fm(?)

(H66). 

Alpes Fm (WM71); Fra

Mauro Fm, hummocky

(H66). 

Rugged material (WM71) Rugged material (WM71)-

n/r 

n/r 

Hilly material (WM71)-

Flank depressions.

Flank, transition zone. 

Do 

Montes Rook Fm, knobby 

facies (S77); domical 

facies (H74). 

Montes Rook Fm, massif facies

(S77).

Concentric facies (M74) Fra Mauro Fm, hummocky n/r 

(H66); material of 

Montes Apenninus (WM71).

- Hevelius Fm, inner facies Fra Mauro Fm (H66, WM71) Janssen Fm (S71, WM71) 

(S77); radial facies (M74).

- Fissured facies (M74) ---- n/n n/r-------------------

Concentric, 

dune!ike. 

Wide valleysHevelius Fm, transverse

facies (S77). 

n/n (Vallis Bouvard type)-

Hevelius Fm, nonlineate

member (W79). 

n/n 

Transition zone Lobate, smooth

or leveed. 

Do Subcircular

flat-floored

craters.

Transition, Narrow grooves Grooved facies (M74) 

outer zones. 

Do Subcircular

bowl-shaped

craters.

Hilly material (WM71)-

n/n 

n/r 

n/rMaterial of crater chains

and clusters (WM71). 

n/r -

Irregular-crater material 

(WM71).

n/n 

Undivided crater material 

(WM71).

n/n 

Do Radial, fine, 

linear, or 

braided.

Do- Planar, smooth 

Satellitic-crater material 

(WM71, W79); Hevelius Fm, 

secondary-crater facies 

(S77).

Hevelius Fm, outer facies 

S77, W79).

Material of crater clusters and chains (WM71).

Terra material (WM71)- 

- Plains material or facies Plains materialMaterial of crater chains 

and clusters (WM71); 

satellitic crater 

material (W79).

Janssen Fm (S71, WM71) -

Pitted-plains material 

(WM71).

Impact melt, buckled.

Impact melt, fluid, little

deformed. 

Late low-velocity high-angle

ejecta.

Upthrust and covered by

ejecta. 

Primary ejecta pushed over

rim.

Primary ejecta, flowed on

surface. 

Ejected impact melt.

Primary ejecta, decelerated

flow. 

Low-angle secondary impact.

Primary ejecta, fluidlike

flow. 

Secondary impact, filled by

surface flow.

Eroded by secondary impact

or surface flow. 

Secondary impact.

Secondary-impact ejecta.

Primary or secondary ejecta, 

ponded.

of the exterior deposits because of mare flooding and other modifications. Thus, Orientale provides a unique opportunity for study of the 

interior materials of basins.

Montes Rook Formation

The Montes Rook Formation (or the knobby facies of that formation, Scott and others, 1977) is characterized by multitudinous 

smooth, equidimensional to elongate knobs, about 2 to 5 km across, 

set in a matrix of smooth to rolling materials (figs. 4.45, D, E\ In 

places, the formation contains weak lineations, mostly radial but 

locally concentric with the rings (fig. 4.4E); much of it has a ropy or 

imbricate texture. Most of the formation occurs between the Montes 

Cordillera and Montes Rook rings, although several large patches are 

exterior to the Cordillera (fig. 4.4B; Scott and others, 1977). Islands of 

lineate material, either belonging to the Hevelius Formation or 

intrinsic to the Montes Rook Formation, lie within the confines of the 

Montes Rook Formation (fig. 4.4E). Vague circular structures, less 

conspicuous than the large primary craters buried by the Hevelius 

Formation (fig. 4.4B), may be buried by the Montes Rook Formation 

(figs. 4.4A, B-E; Scott and others, 1977).

In places, the Montes Rook Formation appears to embay the 

coarsely lineate Hevelius (fig. 4.4B), and so the Montes Rook is the 

later unit in the basin sequence. Like the Hevelius, the Montes Rook 

material must have left the crater at least locally because it is so 

clearly draped over the Cordillera scarp (fig. 4.45). Thus, it cannot be 

the product of inward slumping of the crater walls along the Cordillera scarp immediately or shortly after final excavation of the 

crater, as proposed by Head (1974c), but is a distinct type of ejecta. The 

Cordillera scarp is draped by ejecta, either of the Hevelius or of the 

Montes Rook Formation, over most of its visible extent (figs. 4.4A-F); 

thus, the scarp predates final ejecta deposition and cannot be the 

product of postcratering collapse as is shown on the hypothetical 

cross sections by Hartmann and Yale (1968), Head (1974c), and 

Howard and others (1974).

The coarse, knobby, weakly lineate texture of the Montes Rook 

Formation indicates a much smaller radial-ejection component than 

that of the strongly lineate Hevelius. The Montes Rook Formation 

probably consists of relatively coherent material ejected in relatively 

low-velocity, high-angle trajectories, possibly in a distinct second 

pulse of ejecta. Steep late-forming plumes of ejecta are observed

inside the main ejecta cone in some experimental craters formed in 

layered target materials (Oberbeck, 1975, 1977; Andrews, 1977; 

Piekutowski, 1977). Blocky materials are lifted steeply from the lower, 

more cohesive layer and dropped onto the crater and onto the ejecta 

already emplaced by the earlier, larger plume. In this sense, they 

constitute an "overturned flap" displaying inverted stratigraphy, as 

on the rims of simple craters (McCauley, 1977); however, the two ejecta 

types do not form coherent layers on basin rims. A similar relation 

between two deposits has been mapped at the Ries crater, Germany 

(see chap. 3), which also formed in a layered target. The moderately to 

intensely shocked suevite, derived from the crystalline basement, 

was lofted and then dropped on unshocked Bunte Breccia, derived 

moments earlier from the overlying sedimentary beds. In both the 

experimental and natural craters, the earlier material exits at low 

angles, and the later at higher angles.

Similar phenomena but a different cause were suggested by 

Murray (1980), who compared the Montes Rook knobs to the coarse 

hummocks of probable impact melt on the floors of fresh complex 

craters (figs. 3.2D, 3.32-3.35) and suggested that they were flipped 

upward and slightly outward as the basin or crater floor rebounded 

sharply during or after the main cratering. These suggestions entail 

quarrying of the Montes Rook Formation from a deeper crustal layer 

than the Hevelius Formation, or from the lunar mantle (McCauley, 

1977; Scott and others, 1977; Hodges and Wilhelms, 1978).

Scott and others (1977) mapped a massif facies of the Montes 

Rook Formation in addition to the knobby facies just described. The 

massifs of Orientale and other basins typically consist of rectilinear to 

equidimensional blocks, ranging from a few kilometers to as much as 

100 km in length (figs. 4.45, D, E\ The smoothness of their surfaces 

(except at very high resolution) indicates slope steepness. Their backslopes are typically gentler than their inward-facing slopes (fig. 4.45). 

Heights are estimated to be as much as 6 km above the plains at their 

base (Head, 1974b; Howard and others, 1974; Moore and others, 1974); 

this relief is greater than that between the Cordillera ring and the 

adjacent shelf. Scott and others (1977) considered the massif facies to 

consist of a stratigraphic ejecta deposit superposed on a structural 

uplift, as for similar massifs of Imbrium (Carr and others, 1971; 

Wilhelms and McCauley, 1971). This interpretation is consistent with 

the photogeologic appearance of fresh massifs and with the stratigraphy of crater rims. The proportions of ejecta and structurally uplifted 

material remain unknown.
74 THE GEOLOGIC HISTORY OF THE MOON

FiGURE 4.6. South-central terra of nearside, including secondary-impact craters of distant basins.

A. Arrow in crater Clavius (C; 225 km, 48° S., 15°W.) points to Orientale basin, centered 2,000 km to northwest; small sharp craters around arrow are probably secondary to Onentale. 

Arrow in Longomontanus (L; 145 km, 49.5° S., 22° W.) points to Imbrium basin, centered 2,600 km due north; large overlapping craters around arrow are secondary to Imbrium 

(Saunders and Wilhelms, 1974). Similarcraterson west and north rim of crater Maginus(M; 163 km, 50°S., 6° W.) also are probably secondaries of Imbrium. Right-hand arrow points to 

Nectans basin, centered 1,450 km away; three craters above arrow and degraded group below end of shaft are possibly secondary to Nectans, although lineations are not strictly radial to 

that basin. Mosaic of Orbiter 4 frames H-l 19, H-124, and H-136 (right to left).
4. BASIN MATERIALS ORIENTALE 75

B. Area just east of figure 4.6A. Arrows point toward Orientale center and touch fresh-appearing craters that are probable Orientale secondaries. B, crater Baco (70 km, 51 ° S., 19° E.). Orbiter

4 frame H-100.
76 THE GEOLOGIC HISTORY OF THE MOON

FIGURE 4.7. Region of Orientale antipode; arrow indicates antipode (20° N., 85° E.). 

Most of terra is pitted in manner apparently characteristic of antipodes of young basins. 

MM, Mare Marginis, with bright swirls superposed; G, bright crater Goddard A (12 km), 

possible source of swirls; J, crater Joliot (143 km); L, mare-flooded fringe of secondary 

craters of Lomonosov. Mare material in quadrant below and right of Goddard A is among 

oldest mare units (see chap. 10). Orbiter 4 frame H 17.

Maunder Formation

The Maunder Formation of the Orientale Group is named after 

the unrelated large impact crater Maunder on the north floor of the 

central Orientale basin (fig. 4.45). This formation, like others of the 

Orientale Group, has been described under various names (table 4.3). 

The Maunder consists of smooth to rolling plains of intermediate 

albedo, many of which are closely fissured or fractured. As far as 

observed, the Maunder lies between Mare Orientale and the outer 

Rook ring and extends the peaks of the inner Rook ring (figs. 4.45, E; 

4.8).

The fractures are a secondary, superposed property of the unit 

and thus not a proper part of the definition of a rock-stratigraphic 

unit (see chap. 7), but they nevertheless characterize it; they led to 

such descriptions as "corrugated" and "crackly." These fractures are 

crudely concentric or radial to the basin (fig. 4.45). Most fractures are 

open, downward-tapering V-shaped gashes, whereas others are 

larger, flat-floored complex grabens. Many fractures radiate from 

local centers, and others are aligned along the axes of broad gentle 

elongate domes, which also characterize the Maunder terrain (fig. 

4.45). The fractures are more abundant inside the inner Rook ring 

than close to the outer Rook ring. Other, larger radial and concentric 

grabens that are more pervasive probably were formed by later basin 

subsidence (pi. 5; chap. 6).

Texturally, the Maunder Formation is a scaled-up equivalent of 

the materials that cover the floors of such fresh unfilled large craters 

as Tycho and Copernicus (Howard, 1975; Howard and Wilshire, 1975). 

Though once interpreted as volcanic (Schultz, 1976b, p. 70 77), these 

floor materials are more convincingly interpreted as impact melts, on 

the basis of a comparison with terrestrial analogs and stratigraphic 

analysis (see chap. 3). The thickness of the Orientale melt is difficult

APPROXIMATE SCALE AT CENTER

0 10OO KILOMETERS

EXPLANATION

Mare materials

Post-Orientale crater materials

= ^^""1

Orientale Group

Pre-Orientale terra

FIGURE 4.8. Geologic map of Orientale basin, showing major units discussed in this volume. 

Units of the Orientale Group, from left to right, in "Explanation" are: Maunder Formation 

(double vertical bars), Monies Rook Formation (check marks, knobby ejecta; dots, massifs), 

Hevelius Formation (crosses, inner facies; dashes, outer facies), transversely structured parts 

of the Hevelius Formation (small areas with wavy lines), and some of the many secondary 

craters (black). From McCauley (1977, fig. I; see also Scott and others, 1977).
4. BASIN MATERIALS ORIENTALE 77

to ascertain; Head (1974c) based an estimate of about 1 km thick on 

extrapolations from early work on terrestrial craters. The draping 

and mantling characteristics of much of the Maunder Formation 

suggest that its material behaved in a fluid manner. Perched 

shorelines locally visible (fig. 4.4B) suggest that it subsided to its 

present position from some higher topographic level. Sinuous rilles 

like those of maria (chap. 5) may mark local flowage of the melt (loc. 13, 

fig. 4.45; Greeley, 1976).

The emplacement of the melt and the formation of the fracture 

pattern are believed to have been nearly contemporaneous with formation of the elongate ridges and domes, which are more commonly 

radial or concentric with the basin than equidimensional (fig. 4.4B). 

Mapping of excavations within experimental explosion craters shows 

that similar structures are compressional anticlines and synclines 

that formed late in the cratering sequence (Roddy, 1976; McCauley, 

1977). If the domes in Orientale also formed late, the Maunder Formation may have been cracked by their uplift. The geometry of the cracks 

is also consistent with draping on preexisting topography (Head, 

1974c), which would be their origin if the folding ended before the melt 

solidified.

TOPOGRAPHY AND GRAVITY

Favorably illuminated photographs (fig. 1.9) show that the elevations of the concentric troughs of the Orientale basin decrease inward 

and that Mare Orientale has the lowest surface. The summits of 

Montes Cordillera and Montes Rook appear to differ less in elevation 

than do the adjacent troughs. Absolute elevations are less certain 

because Orientale is not covered by photogrammetric-quality photographs. Preliminary measurements based on Apollo laser altimetry 

of the basin's north shelf and on profiles obtained from telescopic and 

Lunar Orbiter photographs suggest that Mare Orientale lies 2.5 to 

5.5 km below, and parts of the rings and exterior terrain rise the same 

distance above, the 1,738-km average radius of lunar terra surfaces 

(relative to the center of figure) (Howard and others, 1974; Scott and 

others, 1977; Head and others, 1981; S.S.C. Wu, written commun., 

1982). If the mare-basalt fill is less than 1 km thick (Howard and 

others, 1974), the top of the impact melt in the basin center lies about 3 

to 6 km below the 1,738-km datum and about 6 to 11 km below the 

exterior.

Stairstep structure is also evident in stereoscopic photographs of 

the more heavily modified Crisium basin, which is a less representative basin but which lay more favorably under the Apollo 15 and 17 

ground tracks (fig. 4.9). In the terra rim surrounding Mare Crisium, 

interring troughs decrease in elevation inward toward the mare. The 

mountain summits do not decrease in elevation so much (fig. 4.9); that 

is, the overall relief between a given summit and the adjacent trough 

increases inward, as it does at Orientale. Absolute elevation differences are also similar Mare Crisium lies about 4 km below the 

1,738-km datum (Sjogren and Wollenhaupt, 1976) and 3 to 6 km below 

the basin exterior (fig. 4.9). Relative and absolute elevations of 

another measured basin, Smythii, are remarkably like those of Orientale, despite these basins' substantial difference in relative age 

(Strain and El-Baz, 1979).

Perturbations of spacecraft orbits have revealed some important 

features of the Moon's gravity structure. The terrae are gravitationally bland at short wavelengths (see chap. 6; Bills and Ferrari, 

1977; Ferrari, 1977). The relatively short-wavelength gravity anomalies with the largest amplitudes are positive, and are called mascons 

(for "mass concentrations"; Muller and Sjogren, 1968). The gravity 

anomaly of the Orientale interior is positive, but if the basalt of Mare 

Orientale is removed from the gravity models, a ringlike negative 

anomaly remains (Scott, 1974; Sjogren and Smith, 1976). The gravity 

structure of a given basin as a whole requires two bodies the surficial mare basalt and a raised part of the mantle (fig. 4.10; Wise and 

Yates, 1970; Phillips and others, 1972; Scott, 1974; Bowin and others, 

1975; Sjogren and Smith, 1976). Each basin impact removed an 

unknown but significant thickness of crustal material (Scott, 1974). 

The denser mantle then rose to compensate isostatically for the lost

mass. On the basis of models for Serenitatis (fig. 4.10), the mantle may 

be about 20 to 25 km higher under Orientale than in nonbasin regions 

(Sjogren and Smith, 1976); the same value was found for the Grimaldi 

basin (Phillips and Dvorak, 1981). The negative anomaly, also found in 

many basins and fresh craters (Ferrari, 1977; Dvorak and Phillips, 

1978), suggests that compensation was incomplete (Sjogren and 

Smith, 1976). Further alteration of the gravity structure by the later 

addition of mare basalt resulted in a superisostatic mascon in several 

basins (see chaps. 5,6; Baldwin, 1968; Phillips and others, 1972; Scott, 

1974; Solomon and Head, 1979,1980).

Small positive anomalies displayed by Montes Cordillera 

(Sjogren and Smith, 1976) and Montes Apenninus of the Imbrium 

basin (Ferrari and others, 1978) indicate (1) a loading of the crust by 

material removed from the basin and (2) a lithosphere sufficiently 

strong to support much of this added load.

These gravity models are central to later discussions of basinexcavation depths, mare-basalt localization (chaps. 5, 11, 12), and 

tectonism (chap. 6). The crust-mantle discontinuity (lunar Moho) 

probably resembles a series of domelike swells, whose diameter and 

elevation are approximately proportional to diameters of the overlying basins and which are superposed where basins are superposed. 

The largest known basin, alternatively called Gargantuan (Cadogan, 

1974,1981) or Procellarum (pi. 3; table 4.2; Whitaker, 1981), presumably overlies the largest mantle uplift and the thinnest crust. This 

basin was probably responsible for thinning of the crust from the 

computed average of about 75 km to the observed value of 45 to 60 km 

under southern Oceanus Procellarum (chap. 1). This thinning had 

major consequences for later lunar geologic evolution (see chaps. 5,6, 

8-12).

ORIGIN OF RINGS

Introduction

The origin of basin rings has caused considerable interest and 

controversy because of its importance in impact mechanics and stratigraphy, and because it has proved to be such an obdurate problem 

(see reviews by Hartmann and Kuiper, 1962; Baldwin, 1963; Mutch, 

1970; Hartmann and Wood, 1971; Howard and others, 1974; Moore and 

others, 1974; Brennan, 1976; Wood and Head, 1976; McCauley, 1977; 

Wilhelms and others, 1977; Hodges and Wilhelms, 1978; and Croft, 

1981). The identity of the boundary of excavation is particularly significant. Discussions of basin-forming mechanics are generally based 

on knowledge of cratering mechanics derived from small craters, 

which can be generated experimentally and relatively easily modeled 

(see chap. 3). Complex craters and large terrestrial astroblemes (1) 

appear always to have been shallow and not to owe their complex 

internal structure to the collapse of a deep cavity, (2) contain peaks 

that seem to have formed by compression of a central zone during 

subhorizontal centripetal movement, (3) are terraced as a result of 

this centripetal movement, and (4) are surrounded by deposits and 

concentrated secondary craters whose extent remains approximately 

proportional to crater diameter (chap. 3). Some, but not all, of these 

properties can be extrapolated to basins.

Genetic models

The general increase in complexity with increasing size of impact 

cavities (figs. 4.2, 4.3, 4.11) has suggested to some investigators that 

the peaks and terraces of craters and the rings and shelves of basins 

differ in degree of development because of scale differences. Alternatively, they may have been formed by different processes. Properties of the target material, such as strength and layering, or 

properties of the impactor, such as impact angle, velocity, size, or 

shape, may be relatively more important in the larger events. The 

physical effects of large and small impacts may differ fundamentally 

in ways that are not yet clear.

A. Large-basin models. The main, topographic rim is the boundary of excavation (Baldwin, 1974a; Chao and others, 1975; Oberbeck, 

1975; Wilhelms and others, 1977; Hodges and Wilhelms, 1978). Interior rings could have formed in two alternative ways:
78 THE GEOLOGIC HISTORY OF THE MOON

1. Nested craters. In these models, rings are the rims of subcavities of the main excavation (Oberbeck, 1975; Wilhelms and others, 

1977; Hodges and Wilhelms, 1978). Nesting is observed in small 

craters in the lunar regolith (Oberbeck and Quaide, 1967, 1968; 

Quaide and Oberbeck, 1968; Oberbeck, 1977), and probably in the 

Ries (Chao, 1977; Horz and Ostertag, 1979) and other large terrestrial 

craters or ringed basins (Phinney and Simonds, 1977; Hodges and 

Wilhelms, 1978). The initiator of the dual or multiple cavities in 

experimental and natural terrestrial craters is a discontinuity in the 

target materials between an overlying weak layer and a less easily 

cratered substrate. Separate ejecta plumes may be ejected from the 

nested craters.

2. Floor uplifts. Various mechanisms based on cratering mechanics have been proposed. The morphologic similarity between central 

peaks and rings consisting of peaks (fig. 4.35) suggests that the rings 

are expanded peaks (Hartmann, 1972a; Baldwin, 1974a; Head, 1974c; 

McCauley, 1977; Hodges and Wilhelms, 1978). Centripetal movements, akin to those hypothesized for complex craters and observed in 

large chemical-explosion craters (Jones, 1977; Roddy, 1976,1977), may 

form ringlike structures in the innermost basin (McCauley, 1977).

A more violent, oscillatory mechanism akin to that observed in 

temporary craters formed in liquid has also been suggested. This 

"tsunami" mechanism, induced by the seismic energy of impact, was 

first proposed for external rings (Baldwin, 1949,1972; Van Dorn, 1968, 

1969). Baldwin (1974a) then reapplied it to rings of the cavity interior,

which is more likely to be fluidized by the intense and sustained 

shock pressures intrinsic to impacts of large bodies (Baldwin, 1963, p. 

163-169; 1974a). Dence and Grieve (1979; Grieve, 1980) recently 

revived the tsunami mechanism as an explanation for rings, on the 

basis of the high uplifts of severely weakened subcrater material that 

they inferred in terrestrial craters; and Murray (1980) independently 

proposed it on the basis of several photogeologic observations similar 

to those mentioned in this chapter. The oscillatory model is favored 

here and is described additionally below (fig. 4.12).

B. Small-basin models. An inner ring or the former position of an 

inner ring marks the boundary of basin excavation.

1. Late formation. Passive "megaterracing" is thought to widen 

the topographic basin by inward movement and rotation of a large 

annular tract of exterior terrain after excavation has ceased 

(McCauley, 1968; Mackin, 1969; Hartmann and Yale, 1968; Hartmann 

and Wood, 1971; Dence and Plant, 1972; Short and Forman, 1972; 

McGetchin and others, 1973; Gault, 1974; Head, 1974c, 1977; Howard 

and others, 1974; Moore and others, 1974; Head and others, 1975; 

Schultz and Gault, 1975a; Croft, 1981). According to this hypothesis, 

the rear slip surface of the terrace cuts the rim flank outside the 

excavation boundary and carries that boundary inward. The 

topographic basin rim is the site of the slip at the rear of the "megaterrace," and the migrated excavation rim becomes an inner ring.

2. Contemporaneous formation. A similar geometry results from 

centripetal movement and rotation during the excavation process

>7500m 

6600-7 500m 

5700-6600m 

<5700m 

Postbasin material

50°E

FIGURE 4.9. Topographic map of southern Crisium basin. Arbitrary vertical datum is 1,730,000 m from Moon's center. Elevations of summits (white numbers on black) and of troughs (black 

numbers) are shown in hundreds of meters above this datum. Mare Crisium, 3,800 to 3,900 m; Mare Fecunditatis, 5,000 to 5,200 m; Mare Spumans, 5,500 to 5,700 m; Mare Undarum, 

6,400 to 6,700 m. Postbasin materials include crater and mare deposits (Olson and Wilhelms, 1974). Topography generalized from LM 62 (1st ed., Sept. 1978), U.S. Defense Mapping 

Agency.
4. BASIN MATERIALS ORIENTALE 79

(McCauley, 1977; Scott and others, 1977). This more active process is 

observed in large chemical-explosion craters, where centripetal 

movement opens concentric exterior cracks and depresses an annular 

shelf between the excavation boundary and the remnant topographic 

rim. The rim and false rim are overlain by ejecta.

Discussion

1. The topographic rims of such double-ring basins as Schrodinger so 

closely resemble those of craters in morphology and in their demarcation of the inner boundary of craterlike deposits they must be 

equivalent in origin (Wilhelms and others, 1977; Hodges and 

Wilhelms, 1978). That is, the rims approximate the boundary of 

excavation, except as the excavation cavity was widened by visible 

craterlike terraces. Thus, the 160-km-diameter inner ring is a 

feature of the crater interior and not of the boundary of excavation. 

The inner ring cannot bound the Schrodinger excavation because 

such a cavity would be smaller than many craters, which have no 

second ring outside their rims.

2. Similarly, I believe that the topographic rims of larger basins, such 

as the Cordillera of Orientale, are also the boundaries of excavation. This more controversial conclusion is based on: (a) The 

abrupt demarcation of the Hevelius Formation at the Cordillera; 

(b) the elevated massifs along parts of the Cordillera; (c) the 

absence or ghostlike form4 x of prebasin craters inside the Cordillera, in contrast to their abundance outside the Cordillera 

(Baldwin, 1974a); (d) the rimlike morphology, including steep 

inner and gentler outer slopes and both scalloped and linear 

segments (Murray, 1980); (e) the qualitative similarity of both the 

continuous and discontinuous deposits to crater deposits in overall 

morphology and zoning; and (f) the proportional extent of the 

Orientale ejecta in comparison with crater deposits (Wilhelms and 

others, 1977; Murray, 1980, fig. 7). Montes Cordillera extend as a 

major well-defined ring around three-fourths of the basin and are 

not a minor, external structure. Similarly, the 570-km-diameter 

ring of Hertzsprung and the 600-km-diameter ring of the southern 

part of Humboldtianum are the boundaries of those excavations 

(figs. 4.3O, Q, 4.11; table 4.1). Ejecta extends an average of onebasin

4 lrThe presence of ghostlike circular forms inside the Cordillera (Scott and others, 1977) does not obviate the 

interpretation that the Cordillera is the excavation boundary. If the excavations are shallow, surficial material could 

be stripped away without destroying the deeper structure of prebasin craters.

FIGURE 4.11. Spacing of observed basin rings and ring segments, plotted by radial distance 

from basin center (after table 4.1, with some approximate additions from pi. 3); see table 4.1 

for letter symbols. Length of arc within each sector is approximately proportional to observed 

circumferential extent of ring; width of line is approximately proportional to prominence. 

Homologous rings in different basins are connected by dots. Double-ringed structure persists 

through Korolev (J) or Apollo (L). Innermost and intermediate series begin in Hertzsprung 

(O) or, possibly, in poorly observed Coulomb-Sarton (M). These new homologous series 

expand in larger basins through Orientale (Y). Extensive new structure (not completely 

described here) appears in Crisium (Z) and Imbrium (AA).

300 r250 

200 -

-400 -350 -300 -250 -200 -150 -100 -50 0 50 

KILOMETERS

150 200 250 300 350 400

FIGURE 4.10. Two-body gravity model of Mare Serenitatis (Bowin and others, 1975). A thin plate on surface (mare material) and a mantle uplift each 

have densities of 3.4 g/cm3 (shaded); terra-crustal density, 2.9 g/cm3 . Values in space above mare are observed free-air anomalies from two spacecraft 

orbits (in milligals); curved lines are isoanomalies of computed free-air gravity.
80 THE GEOLOGIC HISTORY OF THE MOON

D

E

FIGURE 4.12. Model of basin formation favored here. Drawing by Donald E. Davis, 

courtesy of the artist. 

A. Early stages of growth, resembling simple crater. Target rock is greatly weakened by

shock wave. 

B. Cavity grows and becomes shallower relative to depth, ejection angle decreases, and

central uplift begins as in complex crater.

C. Central uplift collapses to form a temporary double-ringed basin. 

D. Center rebounds, as in a waterdrop crater. 

E. Wavelike motion enlarges inner ring shown in C and D, new innermost ring forms from

central peak in D, boundary of excavation becomes lower, and ejecta blankets basin walls

and exterior. 

F. Excavation boundary reaches final size, and two inner rings and a central uplift have

formed. 

G. Ejecta continues to move beyond excavation; continued oscillation of interior causes surge

of ejecta from intermediate nng and enlarges innermost ring. 

H. Exterior deposits are mostly in place, two rings inside excavation boundary have frozen,

but innermost ring continues to grow, and center nses once again.

/. Innermost ring has grown to final size, and central uplift relaxes into low broad mound. 

J. Final configuration.

diameter from each of these rings, which are also more rimlike 

than any others in the basin. Although parts of the interior rings 

(265 km diam in Hertzsprung and 300 km diam in Humboldtianum) are rimlike, these rings cannot bound the excavations because 

the similar-size rim crests of Bailly (300 km diam) and 

Schrodinger (320 km diam) are not surrounded by ring systems or 

such extensive ejecta (Wilhelms and others, 1977; Murray, 1980, p. 

273).

3. Continued, passive terracing after the completion of ejecta deposition is precluded by superposition of the Montes Rook Formation 

on the Cordillera (McCauley, 1977). Ring formation, whatever its 

cause, must be an active process driven by the cratering flow.

4. Considerable evidence suggests that basins are shallow excavations (fig. 3.20; Head and others, 1975). First, they presumably 

were never deeper relative to their diameters than were craters. 

Copernicus, which is exactly Vio the diameter of Orientale, now has 

a dID ratio of Vzs. According to the hypothesis presented earlier, 

fresh craters were not much deeper when formed than they are 

now (Pike, 1980a). If Orientale's original dID ratio was Vw, its 

original, transitory depth was 37 km. This value is close to that 

obtained by adding the estimated present depth of 6-11 km (exterior terrain to top of impact melt) to the crustal thinning implied 

by the mantle uplift of 20-25 km. (Its present d/D ratio is between 

Viss and Vs5, a decrease that indicates substantial rebound since 

basin formation; furthermore, the negative gravity anomaly would 

be much greater than observed if no rebound had occurred.) Moreover, these ratios apply only to the central basin inside the 320-kmdiameter ring; the shelves, which occupy much more area, are 

several kilometers shallower. That the cavity-wall slopes were 

gentle during excavation as well is suggested by the massive surface flows of ejecta and the linear, chainlike arrangement of secondary craters close to the rim (fig. 4.4D) both features indicating low-angle ejection. The apparent paucity or absence of mantle 

material in the returned samples (chap. 8) furthermore suggests 

that not even the 3,200-km-diameter Procellarum basin was excavated through the 75-km-thick crust implying an original d/D 

ratio of less than Vis for this largest lunar basin.

5. The major interior rings (for example, Montes Rook) probably 

were formed by some form of active floor uplift. In morphology, 

they resemble central peaks more closely than they do terrace lips. 

Their gentle outer slopes and steep inner slopes resemble those of 

crater rim crests, but the jagged massifs and peaks of their most 

rugged parts more nearly resemble central peaks. The Rook rings 

resemble the inner, rugged ring of Schrodinger much more closely 

than they do the Schrodinger rim.

6. Even if rings and central peaks are both formed by uplifts, the 

uplift process probably differs. Although a transition from craters 

to basins is commonly stressed, the interior structure of most 

craters 200 to 300 km in diameter is poorly developed (fig. 4.2), and 

peaks actually decrease in absolute width in craters larger than 

about 70 km in diameter (fig. 3.2E; Murray, 1980, p. 282). Therefore, central deformation may be less intense in intermediate-size 

impacts than in typical complex craters and typical basins. This 

lull suggests different deformational processes (Hale and Head, 

1979; Murray, 1980).

7. Ring uplift may occur during a wavelike deformation of the cavity 

floor, as follows (fig. 4.12). Highly shocked, effectively fluidized 

material of the cavity rebounds, probably because of general 

involvement of the whole floor in wavelike motion to compensate 

for the cavity A smaller central zone then partly collapses when
4. BASIN MATERIALS ORIENTALE 81

gravity exceeds the inertial force and rock strength. If basin 

growth stops at this point owing to insufficient energy and all 

motions "freeze," the result is a double-ring basin similar to those 

as large as Korolev (440 km) in the size series (figs. 4.3,4.11,4.12C; 

table 4.1). In larger impacts, the collapsed part again rises in the 

center, and so on for as many cycles as are consistent with basin 

size, target strength, subsurface layering, other properties of the 

target, and, possibly, projectile size and velocity. Apollo, 

Hertzsprung, and Humboldtianum are large enough to have 

acquired a partial third ring (fig. 4.12E). Orientale has a fourth 

ring and a central uplift that signals the start of an additional 

cycle. Material may flow inward from outside the crater to supply 

the uplift, as is observed in waterdrop experiments. Sublithospheric flow of an early weak, plastic lunar asthenosphere 

may have enhanced wavelike ring formation (Van Dorn, 1968,1969; 

Murray, 1980; McKinnon, 1981). The uplift freezes from the outside 

in because the inner part is the most highly fluidized. This explanation attractively accounts both for regular structure, like that in 

waterdrop experiments, and for irregular structure, as would be 

expected in heterogeneous and fractured rock materials that do 

not behave entirely like fluids.

8. Terrane outside the excavation is also deformed. Chapter 3 (see 

subsection entitled "Deformation and Nonballistic Ejection") mentions that craters may consist of two parts whose sizes are controlled by different scaling laws: an inner excavation cavity scaled 

to gravity, and an outer unexcavated but deformed zone scaled to 

rock strength (Croft, 1981). The deformed zone expands relative to 

the excavation cavity in proportion to impact energy, and thus may 

be very large for basins. External deformation, which in my model 

forms fewer rings than is believed by most investigators, has 

apparently formed arches concentric with the topographic basin 

rims of Humboldtianum, Orientale, and possibly other basins (pi. 

3; table 4.1). Experimental analogs show how these external arches 

may have formed. Mounds and concentric rings were uplifted 

around experimental craters created by certain combinations of 

depth of burst and layering (Piekutowski, 1977), as well as around 

some large chemical-explosion craters (Baldwin, 1963, p. 122). Seismic energy is transmitted outside the cavity by elastic waves after 

the shock wave decays (Gault, 1974; Schultz and Gault, 1975a). The 

elastic waves are refracted and reflected by discontinuities 

between material layers (Rinehart, 1975; Cooper and Sauer, 1977) 

and might propagate quite far under favorable circumstances 

(McKinnon, 1981). Major topographic features, however, are 

unlikely to have formed by such exterior processes. Although several large rises and troughs in the Sinus Medii region have been 

interpreted as external anticlines and synclines of Imbrium (Baldwin, 1963, p. 322; Wilhelms, 1964), mapping of the Procellarum 

basin suggests, instead, that these and many other rings are part 

of that giant three-ring basin (pi. 3; Whitaker, 1981).

9. The topographic rim, which bounds the ejection, is commonly 

irregular in large basins. For example, Montes Caucasus east of 

Mare Imbrium appear to split into two distinct arcs, one continuing northwestward as Montes Alpes and the other continuing 

northeastward as the north shore of Mare Frigoris (fig. 4.3AA). 

Like the Rook ring of Orientale, the Alpes are the inner bound of 

knobby ejecta of Imbrium (Alpes Formation; see chap. 10). Like the 

Cordillera rim of Orientale, the northern Caucasus and the Frigoris shore are the inner bound of radial secondary chains and surface-flow ejecta. The Imbrium rings may have split in the middle of 

Montes Caucasus because the deformation and excavation were 

blocked by the preexisting Serenitatis massifs in the south but 

were freer to expand in the north (Spudis and Head, 1977) into a 

trough of the Procellarum basin (pi. 3). Spalling, as observed ahead 

of ejection in some explosion craters (Cooper, 1977, p. 40; Knowles 

and Erode, 1977, p. 893; Kreyenhagen and Schuster, 1977, p. 989; 

Maxwell, 1977, p. 1004; Piekutowski, 1977, p. 85-88), may soften up 

parts of the surrounding terrain for later ejection (Cintala and 

others, 1978, p. 3805). Ejection from the outer shelves occurred at 

very shallow angles, more like a surficial stripping than the conical plumes of simple craters.

Conclusions

The clear division of basin structure into rings may have been 

overemphasized, considering the pattern of hills, peaks, and ridges

that lack evident ring structure in several sectors of Hertzsprung, 

Humboldtianum, Humorum, Nectaris, Serenitatis, western Orientale, Crisium, and Imbrium (fig. 4.3). Partly regular and partly random processes of floor uplift seem to be required to explain this 

combination of ringlike and irregular patterns. The mechanism 

favored here is oscillatory uplift of the cavity floor.

The conjecture about ring origin presented here has advanced 

some ideas as partial explanations that have previously been proposed as total explanations. Terracing has slightly widened the rim of 

Schrbdinger and larger rims, such as the west rim of Korolev (fig. 

4.3J), but decreases in magnitude with basin size, rather than 

increasing as predicted by megaterrace models (Murray, 1980, p. 272). 

In the largest basins, terracing creates only minor slump slivers, such 

as those northwest of Montes Apenninus (fig. 10.7; Wilhelms and 

others, 1977; Wilhelms, 1980). The crust-mantle interface or other 

layering may modulate more basic processes by enhancing ejection, 

ring uplift, and external deformation. Such oddities as Antoniadi and 

Compton (figs. 4.2A, B) must have been caused by some such local 

effects as layering (Hodges and Wilhelms, 1978). Oscillatory uplift 

may account both for regularities where enhanced and for irregularities where resisted by rock properties. Knobby ejecta of the 

interior and near-rim exterior may be flipped out by the oscillation. 

Oblique impacts probably produced asymmetric ejecta distributions 

(table 4.4) and may have caused internal asymmetries by distributing 

coupled energy unevenly, because they release their energy not at a 

point but along an elongate sloping zone (Gault and Wedekind, 1978). 

A double or triple impact excavated the Humboldtianum basin (Lucchitta, 1978), and other multiple impacts probably caused the elongation of Moscoviense and Serenitatis. External deformation did occur 

but created only second-order features, not the main topographic rim 

or other major arcs and troughs. The topographic rim is the boundary 

of the excavation cavity but is irregular in many large basins. In 

conclusion, basins are highly complex natural systems whose origin 

cannot be learned by straightforward extrapolation from simple cratering mechanics.

SUMMARY OF 

BASIN-MATERIAL ORIGINS

Orientale and similar, older basins and their deposits were 

formed in the following sequence of events. (1) A relatively deep, 

craterlike cavity developed upon impact of a large cosmic projectile; 

(2) this cavity expanded laterally and ejected massive amounts of 

material at low angles, forming gouges and thick lineate deposits; (3) 

the interior rebounded and ejected blocky melt-rich ejecta at higher 

angles and lower velocities; (4) the rebound collapsed, then oscillated 

further, each time to successively smaller radii, and left a ring at each 

hingeline; (5) the last oscillation left a central uplift (known only in 

Orientale). Ejecta was deposited and secondary craters were formed 

during events 3 through 5; then, (6) impact melt settled into place, 

and (7) distant secondary craters were formed.

The interpretations of the well-exposed, well-photographed Orientale deposits presented in this chapter are the principal foundation 

for interpreting older terra materials. Deposits of 17 basins and the 

262-km-diameter transitional feature Milne (fig. 4.2J) are mapped 

here (pi. 3; table 4.1). The presence of rings establishes the existence of 

11 basins whose deposits are not identified but which presumably are 

present; 17 other basins are also mapped whose existence can be 

disputed (table 4.2). Topographic textures indicative of primary and 

secondary ejecta emplacement are the first to become blurred over 

time. The larger secondary craters survive longer but eventually are 

also obliterated or can no longer be traced to their source basin. Even 

the mountainous rings may be so thickly buried or heavily pitted by 

later cratering that they become unrecognizable. Probably, therefore, 

fewer basins are tabulated (tables 4.1, 4.2) and illustrated here than 

were actually formed.

Deposits outside the topographic basin rim grade from thick to 

thin. When fresh, the deposits closest to the rim were coarsely textured, as is the inner facies of the Hevelius Formation at Orientale. 

These materials probably consist mostly of primary ejecta from the 

basin. At an average distance of one basin radius, secondary craters 

begin to dominate the scene. Secondary craters are visible around 

many basins whose textured ejecta deposits are invisible (see chaps.
82 THE GEOLOGIC HISTORY OF THE MOON

8, 9). The inner and outer zones are here mapped separately only 

where their approximate transition can best be distinguished, 

around the Orientale, Imbrium, and Nectaris basins (pi. 3). The 

outer zone is dominated by secondary ejecta derived from local terrain, but the radius where the balance shifts is uncertain. The outer 

deposits persist for two radii on the average but thin gradually to 

apparent disappearance. Distinctive grooves at the basin antipodes 

may indicate a Moon-wide influence of large basins. Much of the 

barely discernible fine-scale relief of the lunar terrae must have 

originated by basin ejecta deposition and secondary cratering.

Deposits of both the primary and the secondary zones are lobate 

and asymmetric. Salients and raylike stringers of secondary craters 

persist outward for several basin radii in some sectors. "Bow tie" 

patterns suggesting oblique impacts at low angles are particularly 

common (pi. 3; table 4.4). Although irregularities are not so evident 

around old basins, they are inferrable. Therefore, the presence at a 

given point of the deposits of a given basin cannot be predicted, unless 

diagnostic textures or mantling relations are visible. Thicknesses are 

very uncertain.

Several types of deposits are seen only in Orientale or a few other 

young large basins (table 4.3). Abundant plains mark the transition 

from the primary to the secondary ejecta of Orientale, Imbrium, and 

Nectaris but are inconspicuous around smaller or older basins. The 

presence of degraded older plains is, nevertheless, highly probable in 

these positions. Planar and fissured impact melt is visible in Orientale (Maunder Formation), but in no other basin with equal certainty. 

At the large young basins, the presence of the knobby ejecta (Montes 

Rook and Alpes Formations) inside the basin, on the topographic rim, 

and on the earlier, thicker deposit (Hevelius Formation and the Fra 

Mauro Formation around Imbrium) suggests late-stage ejection during floor rebound.

Despite the incompleteness of the lunar record and the shortcomings of the data, a history of lunar basin deposits can be pieced

TABLE 4.4. Basin asymmetries suggestive of impact angle

[Azimuths measured clockwise from north. Basins yrouped into four age categories (see tables

4.1, 9.3)]

Basin

Schrodinger-- 

Imbrium- ------

Humboldtianum 

1 Moscoviense-----

Nectaris-- --

Freundl ichSharonov.

20

75

38

1 Q

24

61

5

26

16

19

Center 

(lat, long)

0 S 95° W

0 S., 134° E.

0 N., 19° W.

0 N 129° W

0 N., 84° E.

° S 157° W

0 N., 147° E.

° S., 34° E.

0 N., 175° E.

Azimuths of 

ejecta lobes

150°-330°

160°

0°-155°

10°

145°

0°-210°

Azimuths of centralbasin elongation

Sector 140° elevated

10°-190°

30°-210°

Interpreted 

impact anyle

From 60° or 240°.

From 70° or 250°

(3?0°-340°?).

or 340°.

From 260°.

From 100°.

From 55° or 235°.

From 105° or

285°.

together from superpositional relations, superposed crater densities, 

and degree of morphologic degradation relative to Orientale. This 

history constitutes most of the history of the lunar terrae. Chapters 7 

through 10 show how knowledge of the mechanisms that emplaced 

the Orientale deposits supports the interpretations and dating of the 

older deposits. Breccia and impact-melted rock ascribable to impact 

processes (chap. 3) were recovered from basin deposits by all the 

terra-sampling missions (pi. 3; fig. 4.13). Apollo 14 returned meltbearing ejecta of Imbrium corresponding to the Hevelius Formation 

(chap. 10), Apollo 15 sampled massifs of Imbrium (chap. 10). Apollo 17 

sampled melt-rich materials from the massifs of the Serenitatis basin 

(chap. 9). Apollo 16 recovered materials of circum-Imbrium plains and 

hilly-and-furrowed materials that are analogous to Orientale deposits but whose origin is controversial (chaps. 9,10).

20IC) Ap 15 0), Ap 16 IN) Ap 14 (I)

Secondary 

ejecta

FIGURE 4. \3. Schematic geologic section of typical lunar basin (based on Orientale), showing locations of Apollo (Ap) and Luna(L) sampling sites relative to basin facies sampled (C, Cnsium; I, 

Imbrium; N, Nectaris; S, Serenitatis). Brackets indicate uncertain positions; alternatives are given for Apollo 16, depending on the basin it sampled (see chaps. 9, 10). Right side is based on 

figure 4.4, left side on figure 4.5. Vertical exaggeration, x20.

FIGURE 5.1 (OVERLEAF). Lobate mare-basalt flows in Mare Imbrium. Apollo 15 frame M-1556.
5. MARE MATERIALS

CONTENTS

Introduction .......................

Mapping properties ...............

Morphology ...................

Blanketing .....................

Albedo ........................

Color ..........................

Orbital chemistry ...............

Other properties ...............

Integration .....................

Thickness .........................

Returned samples ................

Introduction ...................

Composition ..................

Classification ..................

Origin and emplacement .........

INTRODUCTION

The darker half of the Galilean dichotomy, the smooth maria, are 

topographically and geologically simpler than the ringed basins that 

contain them. The maria are the most earthlike and probably the 

most easily understood lunar features (fig. 5.1). They cover about 16 

percent of the lunar surface, or 30 percent of the nearside and 2 

percent of the farside. If younger craters were absent, maria would 

cover about an additional 1 percent of the Moon (dashed contacts, pi. 

4). Most maria flood large basins and their peripheral depressions; 

they also flood craters and small basins that are superposed on larger 

basins. Some mare occurrences have led to the identification of basins 

where none were otherwise suspected. A few small maria that are not 

related to known basins (pi. 4; figs. 5.2, 5.3) may, similarly, indicate 

the presence of undiscovered basins.

Because of their resemblance to terrestrial volcanic terranes, the 

maria were more readily interpreted before being sampled than were 

the terrae. The approximately level mare surfaces were correctly 

taken as signs of fluid, lavalike emplacement (Gilbert, 1893; Urey, 

1952; Baldwin, 1949,1963; Kuiper, 1959; Shoemaker and Hackman, 

1962). Their landforms were perceived as analogous to volcanic landforms on the Earth (Kuiper, 1959; Shoemaker and Hackman, 1962; 

Quaide, 1965). Although nonbasaltic composition (especially ash-flow 

tuff; O'Keefe and Cameron, 1962; Mackin, 1969) and such exotic 

origins as mobile dust (Gold, 1955) and water-laid sediment (Gilvarry, 

1970) were considered, the low albedo and certain diagnostic landforms suggested basic lava to most workers (for example, Baldwin, 

1963, p. 336). Basaltic composition seemed to be substantiated by 

Surveyor 5 and 6 analyses in 1967 (Gault and others, 1968a, p. 287), 

and was established by the Apollo 11 samples in July 1969. Basalt has 

been returned from all the mare sites visited. The rock name "basalt" 

can thus be used as a near-synonym for the less committal term 

"mare materials," even in unsampled regions.5 - 1

Page

85

86

86

89

94

96

97

98

99

99

101

101

101

101

102

s-1 Strictly speaking, the term "maria" (singular and adjectival form, "mare") refers to the surfaces, and "mare 

materials" to the rock-stratigraphic units that underlie the surfaces; however, "mare" and "maria" are commonly 

used for the materials as well.

Stratigraphic relations of the mare-basalt flows are also readily 

understood in principle because they are similar to relations familiar 

on the Earth. The margins of the flows follow the contours of the 

containing basins, as would any later fluidlike deposits. Relative 

crater densities are consistent with the younger age of these flows.

FIGURE 5.2. Isolated mare patch in crater Kohlschiitter (arrow; 53 km, 27° N., 154° E.). 

Many craters of subequal size are probably secondary to Mendeleev basin (see chap. 9), 

outside lower left corner of scene. Boom of Apollo 16 gamma-ray spectrometer protrudes 

into view at right. Apollo 16 frame M 729.

85
86 THE GEOLOGIC HISTORY OF THE MOON

The main obstacle to acceptance of these obvious relations was the 

peculiar confusion between mare and basin (see chap. 2), which 

persisted until the early 1960's and is still not entirely dispelled.

By definition, mare materials (1) form generally smooth and level 

surfaces and (2) are dark. In some areas, mare units can be subdivided by their topographic characteristics. However, most subdivisions are based on optical, spectral, or geochemical properties 

commonly called remote-sensing properties. In this respect, mare 

units differ from impact units, because they are igneous-rock bodies 

whose compositional unity is more obvious than is morphology due to 

emplacement. The recognition of mare subunits has advanced greatly 

since 1972 over its rather primitive status before Apollo exploration 

(Head and others, 1978b). The following sections discuss the properties by which maria and their subunits are recognized and interpreted by remote observations, and the last section of this chapter 

briefly describes and interprets the samples of mare basalt returned 

from the Moon.

MAPPING PROPERTIES

Morphology

Certain areally minor, though genetically significant, morphologies interrupt the flat mare surfaces. The most significant are 

lava-flow fronts. These lobate edges resemble terrestrial lava lobes 

and thus indicate that at least some maria consist of volcanic lavas. 

The first flow fronts to be described and the largest yet known are in 

Mare Imbrium (fig. 5.1; Schaber, 1969). They are visible in good 

telescopic photographs and were shown to coincide with color boundaries, presumed to be some sort of compositional contacts (Kuiper, 

1965, p. 29-32; Whitaker, 1966, p. 79-83). These observations quieted 

exotic hypotheses of mare origin in the minds of most workers. The 

Imbrium lobes range from 10 to more than 35 m in height and bound 

flows as much as 1,200 km long (Schaber, 1973). These great dimensions may indicate that each flow formed by rapid eruption of fluid 

lava from a single large vent (Schaber, 1973). The Imbrium flows are 

probably the only relatively recent large flows still exposed, although 

earlier, now-buried eruptions may have been equally large. Scarps of 

shorter and thinner flows are observed in Oceanus Procellarum, 

Mare Cognitum, Mare Vaporum, and Sinus Medii (Schaber and others, 1976).

Otherwise, lunar mare flows are either too thin or, alternatively, 

too floodlike (Holcomb, 1971; Greeley, 1976) to have formed lobate 

scarps. A pooled, massive style of some mare units is suggested by

marginal terraces (fig. 5.3). Holcomb (1971) pointed out the similarity 

of these "bathtub rings" perched on the terrae around some maria to 

those left by subsidence of Hawaiian lava lakes. The terraces suggest 

withdrawal of ponded lava by way of observed sinuous rilles (fig. 5.4; 

Holcomb, 1971; K.A. Howard, in Carr and others, 1971; Howard and 

others, 1972; Greeley and Spudis, 1978).

Several types of positive landforms, mostly formed by extrusion 

of basalt, are scattered in several, and concentrated in a few, maria 

(pi. 4; Guest and Murray, 1976). Most domes are subtle circular 

features with low profiles, commonly topped by small smooth-rimmed 

craters (figs. 5.5, 5.6); they are visible in telescopic visual observations or photographs made at low Sun illuminations (Arthur, 1962). 

The apparent internal origin of the domes was a major factor in 

endogenic interpretations of the maria, and their low profiles suggested fluid volcanism characteristic of mafic magmas. Low mare 

domes are concentrated in Mare Tranquillitatis (Vitruvius region, fig. 

5.5) and western Mare Insularum (Hortensius and Milichius regions, 

fig. 5.6) but occur in other maria as well. Steeper, rough-surfaced 

features commonly called cones are concentrated in the Marius Hills 

(fig. 5.7; McCauley, 1967a, b, 1968,1969b; Guest, 1971) and are scattered in several maria, for example, Mare Cognitum, Mare Tranquillitatis (Wilhelms, 1972a), and the border of Mare Serenitatis 

(Bryan and Adams, 1973; Scott, 1973). Some cones occur in the terrae 

(fig. 5.8). The difference between the domes and cones presumably 

reflects variations in composition or in the eruptive process, and the 

Marius Hills were a favored site for Apollo missions in hopes that they 

would yield differentiated volcanic-rock types (McCauley, 1969b). The 

morphologies of domes or cones are mimicked by kipukas or steptoes, 

synonyms for protruding terra or mare deposits isolated from the rest 

of their geologic unit by floods of later mare materials (see fig. 6.3). 

The Riimker hills (fig. 5.9) apparently consist of several domes and 

other mare-related units partly burying Imbrium-radial terra 

(Guest, 1971). None of the positive lunar volcanic landforms remotely 

approaches a terrestrial or martian shield volcano or stratovolcano in 

size. Almost all the landforms listed here occur in shallow maria 

(Arthur, 1962, p. 322; Head, 1976a). Some 200 such landforms were 

recently classified by Head and Gifford (1980).

FIGURE 5.3. Lacus Sohtudinis, showing down (d) and up (u) sides of marginal terraces of 

mare. Scene centered at 27° S., 103° E. Apollo 15 frame M-2627.

FIGURE 5.4. "Scablands" north of Anstarchus Plateau. Mare lava probably drained out of 

the region in the sinuous rille (left). Large craters in lower right are possibly secondary to 

Orientale basin. Scene centered at 30° N., 53.5° W. Apollo 15 frame P-349.
5. MARE MATERIALS 87

FIGURE 5.5. Mare domes in northern Mare Tranquillitatis. Numerous elongate pits are on summits of low domes or in flat mare (arrows). Mare units are numbered in order of 

decreasing age. Crater below number 3 (youngest unit) is Maraldi B (7 km). Scene centered at 15° N., 35° E. Apollo 17 frame M-305.

FIGURE 5.6. Detail of mare domes and other endogenic features, including circular domes with elongate central craters (1, 2); irregular domes with circular, nevertheless 

probably endogenic, craters (3); concentrations of similar large, soft-appearing circular craters that could be endogenic (4, 5); sinuous rille (6); and elongate endogenic vent 

or collapse feature (7> Scene centered at 13° N., 30.5° W. Orbiter 5 frame M-164.
88 THE GEOLOGIC HISTORY OF THE MOON

Visible source vents are rare in the maria, because the lavas 

probably cover their sources as do flood-basalt lavas on the Earth 

(Greeley, 1976). Some endogenic craters do appear amidst the maria, 

however, and are fairly numerous along their borders (figs. 5.1QA-D); 

others lie in small patches of dark material in the terrae (figs. 5.8, 

5.WE H). Their internal origin is determined by this association with 

volcanic materials or by their alignment along faults. For example, 

random impact is unlikely to be responsible for the many craters 

exactly centered on what are evidently volcanic domes (figs. 5.5, 5.6). 

The craters or rimless pits along Rima Hyginus are similarly eliminated statistically as impact craters because they are perfectly 

aligned along the rille (fig. 5.WE). Although one leg of Rima Hyginus 

is radial to the Imbrium basin, both the rilles and the craters are too 

young to have been formed directly by the impact (see chap. 6; 

Wilhelms, 1968; Pike, 1976). Like many other endogenic craters (figs. 

5.1QA-D), the rille craters and the main Hyginus crater at the rille 

"elbow" could have formed entirely by collapse without the extrusion

of any material (Pike, 1976). Single craters situated along rilles, 

amidst dark materials, and with elongate or irregular shapes are also 

probably endogenic (figs. 5.WF H). A few craters not situated along 

rilles resemble secondary-impact craters, but endogeny may be 

assumed if the craters are highly irregular and are centered in dark 

deposits (figs. 5.10G, H\ Although the dark ejecta of endogenic craters 

is commonly more extensive relative to the crater size than is the 

ejecta of impact craters (table 3.1), the volcanic ejecta is smoother, and 

the craters are all small (mostly less than 10 km diam). Despite the 

early popularity of a caldera origin for such large craters as the 

smooth-rimmed class (fig. 3.15), large calderas are not expected on the 

Moon; calderas form by collapse or magma withdrawal into shallow 

magma chambers, which probably never existed on the Moon.

Sinuous rilles are another areally minor but visually arresting 

type of mare feature. The Apollo 15 landing was targeted next to a 

large sinuous rille, Rima Hadley (fig. 5.11; chaps. 10,11). Several other 

rilles were considered for landings to determine their formational

1 ' '^>K^ ^ --V* ''?S'-,.*-^-/

 ''w^l*»l'4-n^ . . y/..'./ A'; j|--. :'. . 

f i irV''Wi; 't^; * i'-i* ~. »>* ' i1 ." -.^.ff.a, ';. A.-- i* .

FIGURE 5.7. Manus Hills, a complex of diverse, rough-textured domes or cones surrounded by smoother mare materials. Orblter 5 frame M 211.
5. MARE MATERIALS 89

FIGURE 5.8. Rough cones (black-and-white arrow) situated along linear rille and amidst 

dark-mantling material or mare lava. Smooth-profiled irregular pit (white arrow) further 

indicates basaltic volcanism in this terra setting. Crater rim separates craters Fra Mauro 

(above) and Bonpland (below). Black-and-white arrow is 5 km long; scale varies because of 

obliquity. Apollo 16 frame P-5425.

FIGURE 5.9. Riimker hills (Mons Riimker), a complex of subdued hills in northern Oceanus 

Procellarum (41° N., 58° W.). Imbrium-radial lineations (upper right to lower left) suggest 

an origin as terra covered by thin layer of dark-mantling material.

process and to collect samples of volatile or other exotic erupted 

materials (fig. 5.12). Fieldwork on the Earth and at Rima Hadley has 

shown that these meandering grooves are probably lava channels or 

collapsed lava tubes, which are the sites of the last flows of molten 

magma in a lava unit (Kuiper and others, 1966, p. 199; Oberbeck and 

others, 1969; Greeley, 1971; Murray, 1971; Cruikshank and Wood, 1972; 

Howard and others, 1972; Guest and Greeley, 1977, p. 62-68). Very 

long and narrow rilles are likely to have been channels (always open); 

others, including Hadley, display alternately roofed and unroofed 

segments that indicate an origin as tubes (figs. 5.11,5.13). The sources 

of some lunar rilles are in the terra or at the mare-terra contact (figs. 

5.11, 5.12). Several appear to have erosionally incised the substrate, 

including terra material (fig. 5.12C; Carr, 1974). Sinuous rilles and 

other volcanic landforms of the maria were thoroughly described and 

illustrated by Lowman (1969), Schultz (1976b), and Masursky and 

others (1978).

Blanketing

The property of blanketing or mantling terra topography distinguishes another class of dark, mare-related units, the darkmantling materials. They are as dark as or darker than the maria, 

but their observed relief is partly rugged. The largest and most 

numerous tracts with this seemingly anomalous combination of properties lie around the southern periphery of the Imbrium basin (pi. 4; 

figs. 1.6,1.7, 5.14, 5.15). This position led to their interpretation as a 

dark facies of Imbrium ejecta (Shoemaker and Hackman, 1962; 

Hackman, 1966). Carr (1965a, 1966a, b), however, showed that the 

large dark areas on the southeast slope of Montes Apenninus are not 

composed of Imbrium ejecta. He observed that flat surfaces next to the 

dark ejecta are equally dark, and so the same dark material probably 

covers both the ejecta and the much younger mare material (fig. 5.16). 

Moreover, other patches of mantling material are peripheral to Mare 

Serenitatis (figs. 5.17, 5.18; Carr, 1965a, 1966a, b), Mare Humorum 

(Titley, 1967), Mare Nectaris (Elston, 1972), Mare Crisium (Olson and 

Wilhelms, 1974), Mare Smythii (Wilhelms and El-Baz, 1977), Oceanus 

Procellarum (Moore, 1965,1967; McCauley, 1967a, b), and, in fact, all 

carefully scrutinized maria (pi. 4; Head, 1974a, 1976a; Wilhelms, 

1980). This concentration near maria and the dark color indicate a 

genetic relation of the dark-mantling materials to the maria. The flat 

parts of maria may also contain dark-mantling materials, but these 

are difficult to distinguish from mare lavas.

A more specific origin for the dark-mantling materials is suggested by other physical properties. On both telescopic and spacecraft 

photographs, they appear fine-textured, relatively uncratered, and 

smoother than the underlying topography (figs. 5.15,5.16). No blocks 

are visible in the 2-m resolution of the best Apollo photographs. 

Mantles drape differentially over mountains, low hills, and flat surfaces, and are thickest in depressions and thin or absent on sharp 

hills. The brightness of some hilltops surrounded by the deposits 

indicates nondeposition or shedding of the dark material. Shedding is 

the preferred explanation because some dark deposits lie at higher 

elevations than some bright summits. Also, the action of mass wasting is indicated by talus streaks that descend from the bright outcrops (fig. 5.16). Shedding is consistent with the propensity of unconsolidated material to slide downslope.

The morphologic properties and distribution of the dark-mantling materials suggest that they consist of volcanic fragments, 

pyroclastic material. The Marius Hills (McCauley, 1964a), the 

Aristarchus Plateau, and Montes Harbinger (figs. 5.7, 5.12A; Moore, 

1964b) were originally interpreted as terra features covered by volcanic ash (figs. 5.7, 5.12). Carr (1965a, 1966a, b) noted that the dark 

deposits are situated near likely eruptive fissures in the form of rilles 

in the maria (fig. 5.16). Unlike lavas, few dark-mantling materials 

lack elongate or irregular, commonly rille-straddling craters, gashes, 

or dark cones that are likely source vents (figs. 5.8, 5.WF H, 5.12, 

5.16). Pyroclastic origin was favored for the dark-mantling materials 

in all the early photogeologic work and has been confirmed by sampling data (see chap. 11). The dark halos of many small endogenic 

craters, such as those in Alphonsus (fig. 5.10F), are also dark-mantling deposits. Head and Wilson (1979), however, suggested that the 

eruptive mechanisms differ: The extensive, regional deposits are 

thought to be strombolian (continuous eruptions), and the smaller 

halos vulcanian (intermittent eruptions).
D

FIGURE 5.10. Endogenic craters.

A. Craters aligned along graben on south margin of Mare 

Serenitatis (18° N., 26° E.; compare fig. 11.13) Apollo 

17 frame P-2321.

B. Aratus CA (24.5° N., 11° E.), compared by Greeley 

(1973) to a volcanic crater in Idaho. Stereoscopic pair of 

Apollo 15 frames P-9350 (right) and P-9355 (left).

C. Probable collapse craters in southern Mare Imbrium; collapse of the deep, well-defined crater (5 km long) broke 

mare surface, whereas collapse of the gentler crater probably did not. Apollo 1 7 frame P-3125.

D. "D-caldera" (Ina), an irregular crater (3 km long) with 

peculiar floor texture. In Lacus Felicitatis (19° N., 5° E.), 

perched on Monies Apenninus flank, 300 to 400 m above 

nearest large mare (Vaporum; Strain and El-Baz, I960). 

Apollo 17 frame M-1672.

E. Rima Hyginus; largest crater is Hyginus (9 km, 8° N., 6° 

E.} Orbiter 4 frame H-102.

R Dark-haloed craters along fissures in floor of crater 

Alphonsus (compare fig. 1.8; Carr, 1969; McCauley, 

I969a). Arrow (5 km long) indicates Ranger 9 impact 

point (12.9° S., 2.4° W.). Apollo 16 frame P-4656.

G. Irregular crater, fissure, and dark-mantling material (center) in Monies Riphaeus (6° S., 27° W.). Apollo 16 frame 

P-5453.

H. Crater next to Rima Bode II (13° N., 4° W.), amidst 

dark-mantling material (patch d3, fig. 5.14). Orbiter 5 

frame H-122.
i; ,*' -Y *% ; .V . rv 5

/': '?'' V '*";- ?',, *l .'">;

FIGURE 5.11. Rima Hadley (Hadley rille) at Apollo 15 landing site (arrow). Rille hugs and 

possibly erodes terra at A, and is bridged at B and C by basalt. D, dark-mantled "North 

Complex," surrounded by level mare. Ledges are visible at E. StG, St. George Crater (2.25 

km} Apollo 15 frame H-11720.

FIGURE 5.12. Aristarchus Plateau-Monies Harbinger complex of sinuous

rilles.

A. Monies Harbinger and crater Prinz (P). Large fresh crater is Aristarchus 

(A; 40 km, 24° N., 47° W.) Arrow indicates rille in C. View southward. Apollo 15 frame M-2606.

B. Aristarchus Plateau. Aristarchus is at left edge; H, Herodotus (35 km, 

23° N., 50° W.). Arrow indicates rille in D. View southward. Apollo 15 

frame M-2611.

D

C. Terra ridge in Monies Harbinger (arrow in A), eroded by rille. Apollo 15

frame P-321. 

D. Small meandering rille within largest lunar sinuous rille, Vallis Schroteri

(arrow in B), which probably formed in massive lava flow(s)i Apollo 15

frame P-341.
92 THE GEOLOGIC HISTORY OF THE MOON

FIGURE 5.13. Sinuous complex, consisting of ridge and elliptical and curved pits, that may 

have originated as a lava tube whose surficial crust collapsed intermittently along its length 

(Oberbeck and others, 1969). Crater in upper right corner is Gruithuisen K (6 km, 35° N., 

43° W.). Obiter 5 frame M-182.

 i*

,(G>

o

tl

FIGURE 5.15. Ray-crossed patch of dark-mantling material south of Copernicus (dl) and 

large patch south of Sinus Aestuum (d2; compare fig. 5.14). Sharp-nmmed (delta-rimmed) 

crater at lower left is Gambart (G; 25 km, I ° N., 15° W.; compare fig. 3.15C). Telescopic 

photograph.

FIGURE 5.14. Telescopic full-Moon photograph of Maria Imbrium (upper left) and Serenitatis (right). Bright crater Copernicus (C) is at left edge. Dark patches are south of Copernicus (dl), on 

border of Sinus Aestuum (d2, d3), on border of Mare Vaporum (d4), and along margin of Mare Serenitatis (d5, d6). Serenitatis is also bordered by other dark mare units. U.S. Naval 

Observatory, Flagstaff, Ariz., photograph.
5. MARE MATERIALS 93

FIGURE 5.16. Dark-manthng materials on southwest rim of Serenitatis basin (d5, fig. 5.14). 

Decrease in number of exposed bright terra peaks from lower left to upper right indicates 

increase in thickness of dark-mantling material toward the mare. Talus of dark-mantling 

material (t) covering brighter slope indicates shedding of superposed dark deposit. Blackand-white arrow indicates flooding of dark-mantling material by mare lava. Arrowheads

denote endogenic craters. Lacus Odii (LO), which is elevated substantially (approx 1.5 km) 

above Mare Serenitatis, may consist of pooled dark-mantling material (Wilhelms, 1980). 

Arcuate rilles are Rimae Sulpicius Callus; S, crater Sulpicius Callus (12 km, 20° N., 12° 

E.). Stereoscopic pair of Apollo 17 frames M-2701 (right) and M-2703 (left).

FIGURE 5.17. Eastern Mare Serenitatis and adjacent terrain. Dark-mantling material (d6; 

compare fig. 5.14) is in region of Apollo 17 landing site. Central Serenitatis is depressed 

below darker border zone; grabens (see chap. 6) are in border zone above crater Plinius (P; 

43 km, 15 N., 24 E.). Elevation of Mare Tranquillitatis (bottom) is approximately the same 

as that of Serenitatis border zone. Large crater Posidonius (95 km) along mare-terra contact 

near top of photograph has an uplifted floor (see chap. 6). Telescopic photograph.
94 THE GEOLOGIC HISTORY OF THE MOON

Albedo

Because of the contrast between exposures of feldspathic terra 

breccia and mafic mare basalt, the maria can be delineated even by 

unsophisticated observations (figs. 1.1,5.14). For example, Luna 3 was 

able to reveal the important fact that the farside is deficient in maria, 

despite poor camera resolution (Barabashov and others, 1961; 

Whitaker, 1963).

This observation illustrates that the mare-terra contrast shows 

up best when the Sun illumination is high; shadows become very rare, 

and true brightness differences are observed, at Sun angles higher 

than about 45°. The highest possible Sun angle, exactly normal to the 

surface, is called zero phase; the reflectance at zero phase is known as 

normal albedo. Zero phase is never observed from the Earth because 

the Moon would be eclipsed, and so quantities described by terms 

such as reflectance, reflectivity, brightness, or, simply, albedo are 

used more or less interchangeably as convenient approximations to 

true normal albedo. Mare albedos generally range from 0.07 to 0.10; 

that is, the maria reflect 7 to 10 percent of the incident visible light 

(Pohn and others, 1970). With minor exceptions, the total range of 

lunar visible albedos is about 0.07 to 0.24.

Considerable effort has been applied to quantifying albedo 

because of its value in distinguishing geologic units (Pohn and others, 

1970). For geologic mapping of intramare units, however, high-Sun 

photographs with good spatial resolution are required, even if not 

accurately calibrated from region to region, because they show contacts between albedo units. Mare units have been extensively mapped 

since preparation of the geologic map of Mare Serenitatis (fig. 5.18; 

Carr, 1966a), where the boundaries are relatively unobscured by rays 

and the contrast between the dark border and the brighter interior is 

evident. Four Serenitatis albedo units became standards for mapping 

other areas (Wilhelms, 1970b, p. F30). Although the causes of these 

albedo differences were not known, the fact that albedo correlates to 

some extent with mare units was clear.

A leading hypothesis was that albedo is related to age. Several 

lines of evidence suggested that dark units are young and light units 

old. Except with respect to large craters (Dodd and others, 1963), the 

inward-sloping Serenitatis border mare appeared to contain fewer 

craters than the central mare. The dark-mantling materials of the 

border are superposed on some mare materials and appeared to be 

young. Dark-mantling or mare materials are superposed on lighter 

planar materials in other areas, for example, in Alphonsus, Mare

South

SCALE 1.1 000000

LAMBERT CONFORMAL PROJECTION

STANDARD PARALLELS 21°20' AND 42°4O'

\ ( \ [ \ v r v. y \ ~\

FIGURE 5. 1 8. Geologic map of Mare Serenitatis and margin (from Carr, 1 966a). Dark mare (Ipm3, Ipm4) and dark-mantling materials (Emd, Els) of border zone are mapped in dark tones. 

High-numbered units were thought to be young; they actually are young in northeast and northwest, but are old in south. Unit boundaries remain valid.
5. MARE MATERIALS 95

Imbrium, Palus Putredinis, Julius Caesar, and Sinus Medii (figs. 1.6, 

5.WE, F, 10.35, 11.13). Moreover, crater rays are generally brighter 

than the maria they cross, and so cratering was assumed to be an 

"instant aging" process that brightened materials. Correlation of 

albedo units thus became a means of correlating mare ages from area 

to area (Wilhelms, 1970b), and albedo units were shown as age units 

on most lunar geologic maps.

Albedo is now known to be more a function of composition than of 

age, as suspected by some geologic mappers (Trask and Titley, 1966). 

The dark basalt of the Mare Serenitatis border that was sampled by 

Apollo 17 yielded some of the oldest absolute ages yet obtained on 

exposed mare materials (see chap. 11). Albedo is a function of exposure 

to the space environment and of the alteration processes that create 

regoliths (soils) from bedrock. It is controlled by both the amount and 

the composition of glasses developed from crystalline material by 

small impacts (Adams and McCord, 1970,1971,1973; Nash and Conel, 

1973). These glasses are darker than their parent crystalline materials. The most significant glasses in the darkening process are those 

that bind other regolith fragments into the complex aggregates 

known as agglutinates. Agglutinates accumulate in the regolith as a 

function of exposure age. Eventually, the amount accumulated

reaches a steady state in which new micrometeorite impacts destroy 

as much agglutinitic material as they create. The albedo is then at an 

equilibrium value, and the regoliths are termed mature (Charette 

and others, 1974). The observable properties of mature regoliths are 

relatively homogeneous over the extent of a given bedrock unit.

In mature regoliths, the composition of the agglutinates strongly 

affects the albedo; regoliths highest in iron and titanium are the 

darkest (Adams and McCord, 1970,1973; Pieters, 1978, p. 2839-2840). 

Because all undisturbed mare regoliths are believed to be mature, the 

darkest regoliths form from the lavas that are highest in Fe and Ti. 

Higher contents of these elements also explain why the maria are 

darker than the terrae. The darkening agents may be Fe-Ti-rich 

glasses, Fe and Ti oxides in such opaque minerals as ilmenite that are 

contained in the agglutinates (Pieters, 1978), and (or) metallic Fe 

reduced from oxides by impacts and solar-wind sputtering (Hapke 

and others, 1975). Although Fe may be the greater contributor to total 

darkening, the differences in albedo among mare units may be due 

more to differences in Ti content, because Ti content varies more 

among mare basalt units than does Fe content (Pieters, 1978, p. 2840). 

In other words, all mare regoliths are dark because their glasses 

contain Fe and Ti, but the darkest contain the most Ti.

le Monni* C * 

& (430 j \ fi- * \ \

t_

30"

28°
96 THE GEOLOGIC HISTORY OF THE MOON

This correlation between darkness and Fe-Ti content applies 

only to mature regoliths. Fresh exposures of mare materials may not 

have been exposed long enough to accumulate impact glass in steadystate amounts. Such exposures occur in two main settings. First, new 

glass is shed, and nonglassy material continuously exposed, on steep 

slopes; the fresh material is brighter than the less disturbed material 

(figs. 3.28, 3.31, 5.16). In general, lunar slopes are brighter than 

adjacent more nearly level terrain composed of the same geologic unit 

(Shoemaker and Hackman, 1962). Second, relatively recent impacts 

exhume fresh material of mare lavas and redeposit it as bright rays. 

On bright, fresh lunar surfaces, more crystalline material than glass 

is exposed (Adams and McCord, 1970, 1971, 1973; Conel and Nash, 

1970; Charette and others, 1974). Therefore, the formation of a crater 

ray in a given unit is not necessarily one of "instant aging" but of 

"instant rejuvenation." The correlation between youth and darkness 

demonstrated by dark-on-light superpositions is valid where the 

younger materials are richer in Fe and Ti, as in Alphonsus and the 

other examples given above, but not for all juxtapositions of dark and 

light.

In summary, geologic maps prepared before the mid-1970's correctly show relative ray and slope ages, and the mare units generally 

are correctly delineated. The mare units, however, should be interpreted as compositional, not age, units except where youth and composition correlate.

Color

Color differences are subtle on the Moon. The spectra of all areas 

are similar, and increase in reflectance toward longer wavelengths; 

that is, all lunar colors are red. Differences do exist, however, and can 

be detected with instrumented Earth-based telescopes and enhanced 

by various photographic, instrumental, and computerized imageprocessing techniques. Differences also appear conspicuous to the 

trained astronaut eye (Evans and El-Baz, 1973; Lucchitta and 

Schmitt, 1974). Color differences have proved to be valuable discriminators of lunar mare and dark-mantling units and, since the returned 

samples were analyzed, have helped extrapolate known compositions 

to unsampled units (table 5.1).

Reflectance spectra over wavelengths of 0.3 to 2.5 |xm (ultraviolet 

to near-infrared) have been intensively studied during the space age 

for information about lunar compositions obtainable by remote 

means (fig. 5.19; Adams and McCord, 1970,1973; McCord and others, 

1972a, b, 1976; Adams and others, 1974; Charette and others, 1974; 

Pieters and others, 1974,1975,1980,1982; Johnson and others, 1975, 

1977; Head and others, 1978a, b; Pieters, 1977, 1978; Basaltic Volcanism Study Project, 1981, chap. 2). Areas 10 to 20 km across are 

measured photometrically. Absorption features due to mineral constituents of the surface material may be detected near wavelengths of 

0.4,1.0,1.2, and 2.0 |xm. These features are weak for agglutinate-rich 

regoliths and stronger for crystalline rocks (fig. 5.19C). To compare 

spectral differences, spectra are commonly divided by a standard 

spectrum of a relatively uniform mare surface in central Mare 

Serenitatis (designated MS-2; spectral class mISP, table 5.1). Small 

departures from this curve are then exaggerated in graphs, which are 

normalized to unity at 0.57 |xm to eliminate albedo effects and to 

emphasize color differences (fig. 5.19). Most lunar spectra obtained 

before 1978 cover the spectral range 0.3 to 1.1 |xm (figs. 5.19A, B; see 

references and historical background in Pieters, 1978, and Moore and 

others, 1980b). Spectra in the 1.0- and 2.0-|xm bands have been 

obtained more recently (figs. 5.19C, D; McCord and others, 1981; 

Pieters and others, 1982).

For mapping and discriminating units, the most useful colorenhancement displays are images that show the spatial variations of 

simple spectral parameters (Whitaker, 1966, 1972b; Johnson and 

others, 1975,1977; McCord and others, 1976). Black-and-white photographs prepared by combining a photographic negative taken at 

one wavelength with a positive taken at another provide much data at 

good spatial resolution (fig. 5.20; Whitaker, 1966,1972b). Two major 

mare subdivisions, relatively red and relatively blue, can be defined. 

For example, Mare Tranquillitatis and the border of Mare Serenitatis 

are relatively blue, and central Serenitatis is relatively red. The 

young Imbrium flows are relatively blue, and most of the rest of Mare 

Imbrium is relatively red. Several intermediate colors can also be 

resolved.

TABLE 5.1. Remotely sensed properties of mare units

[Spectral class from Pieters (1978) and Basaltic Volcanism Study Project (1981, chap. 2). 

Ultraviolet/visible ratio: H, high (probably more than 5 weight percent TiOg); h, medium 

high; m, medium; L, low (probably less than 2 weight percent TiOg). Albedo: D, dark (max 

0.08); I, intermediate (0.03-0.09); B, bright (higher than appro* 0.09); 1-um absorption 

band: S, strong; G, gentle; W, weak. 2-ym absorption band: P, prominent; A, attenuated; 

-, undetermined.

Color (fig. 5.20): b, blue; i, intermediate; r, red.

Age (see chaps. 11-13): II, earliest Late Imbrian; 12, middle late Imbrian; 13, latest Late 

Imbrian; E-I, near the Imbrian-Eratosthenian boundary; E, Eratosthenian; C?, Copernican 

(spectral class uncertain).

Interpretation based on Basaltic Volcanism Study Project (1981, p. 452-456)]

Symbol 

(pi. 4)

Spectral 

class Mare Color Samples Age Interpretation

HDWA

HDSA

Tranquil!itatis, 

Serenitatis 

border.

Procel1 arum-- 

hDSA Imbrium, N. Procel 1 arum.

hDSP Humorum, S. Procell arum.

b Apollos

11, 17.

b 

b 

b

hDGhDWA

mISP

Nub i urn--- - -

Fecunditatis, 

Serenitatis 

border.

Central Serenitatis, W. Procel - 

larum.

Luna 16

12

E, C?

E

E-I

13? 

E, 13

mIG- Widespread- i, r Apollo 12 E-I

9

10

11

12

13

mBGLISP

LIGLBGLBSP

Nectaris, S.

Fecunditatis. 

Crisium, Roris,

Frigoris. 

E. Imbrium, W.

Serenitatis. 

Somniorum, Mortis,

Imbrium border,

N. Procellarum,

Nubium. 

Frigoris, Imbrium,

Nubium.

Apollo 16? 

Luna 24 

Apollo 15

12

11-13, 

E.

13

11-13

I2-E

Highest Ti.

High Ti, radioactivity, 

Fe?, olivine?

Similar to HDSA but 

probably lower Ti.

Similar in Ti, Fe, and 

radioactivity to HDSA 

and hDSA, but possibly 

richer in pyroxene.

Medium Ti.

Medium Ti, high Al.

Low to medium Ti. Possibly richer in Fe in 

Serenitatis than in 

Procellarum.

Class of generally lowto medium-Ti basalt.

Low Ti, high Al.

Very low Ti, high Fe.

Low Ti, high radioactivity.

Very low Ti, probably 

high Al.

Low Ti, high in Fe or 

pyroxene, high radioactivity.

Such multispectral images contain considerable compositional 

information. To a first approximation, the excess of ultraviolet or 

near-ultraviolet color over the midvisual that is, blueness is a 

measure of the content of both FeO and TiO2 in agglutinates within 

mature regoliths (Pieters, 1978, p. 2835). In figure 5.20 (made from 

photographs filtered at 0.37 and 0.61 |xm), dark areas are bluer and 

thus richer in these elements. The reason for the correlation (Pieters, 

1978, p. 2835-2841) is somewhat complex. As discussed, Fe and Ti 

darken mature regoliths, and the dark materials in the regolith 

suppress the intensities of individual absorption bands. The more Fe 

and Ti in pure glass, the more ultraviolet light is absorbed, that is, the 

redder the spectrum. The opposite can be observed, however, for many 

Fe- and Ti-rich mature regoliths, because dark regoliths containing 

well-developed agglutinates also strongly absorb in the visible and 

near-infrared. The resulting greatly reduced spectral contrast causes 

the apparent blue color of dark mature mare regoliths (Pieters, 1978).

Low Fe and Ti contents, however, do not necessarily cause the 

opposite effect. Red or intermediate-color maria are lower in Ti than 

relatively blue maria, but their Ti content is not constrained (Pieters, 

1978, p. 2842-2843). The role of Fe in the red maria is unclear. The 

strength of an absorption feature near 1 |xm is affected both by the Fe 

content of the regolith glass and by the proportion and composition of 

the mafic minerals present. This interrelation has not yet been fully 

examined (Adams, 1975; Pieters, 1978, p. 2840-2841).

The dark-mantling materials have distinctive reflectance spectra (fig. 5.19D; Pieters and others, 1973; Adams and others, 1974; 

Charette and others, 1974). Mantles that appear black on the surface 

or bluish-gray to an orbiting astronaut are spectrally blue as well, and 

those that appear orange on the surface or reddish-brown to tan from 

orbit are spectrally red. These colors are due to both the composition 

and the crystallinity of the pyroclastic dark-mantling materials. 

Some mantled regions differ spectrally from any sampled materials.

The bright interiors of fresh, young craters on both maria and 

terrae differ spectrally from the surrounding mature regoliths 

because their regoliths contain less agglutinitic material and more 

unaltered mineral fragments (fig. 5.19C). These properties affect the 

general properties of the spectral continuum and allow detection of 

specific near-infrared absorption features (Charette and others, 1974; 

Pieters, 1977). Visual wavelengths and, thus, the color-difference 

image in figure 5.20 do not fully reveal the spectral properties of fresh 

craters; these properties are better revealed by spectral images in the 

visible and near-infrared because of the conspicuous 1.0- or 2.0-|xm 

absorption bands of pyroxene (fig. 5.19C; McCord and others, 1976).
5. MARE MATERIALS 97

Orbital chemistry

Two instruments flown in lunar orbit on the Apollo 15 and 16 

missions have also permitted extrapolations of sample compositions 

to unsampled parts of the Moon. The gamma-ray spectrometer and X-- 

ray-fluorescence spectrometer Adler and Trombka, 1977; Basaltic 

Volcanism Study Project, 1981, chap. 2) are currently the only means 

of estimating compositions on the farside (pi. 2). Whereas the reflec-

«i 1.00

I8 1.00

1.00

1.00

1.00

1.00

1.00

0.96

O H1 APOLLO 11

A H2FLAMSTEED

* hi LUNA 16

 h2 GS-3

D hSMHO

o ml APOLLO 12

	HDWA- 

/MS2 HDSA - 

/MS2 hDW- 

/MS2 hDSA 

/MS2 hDSP 

/MS2 mlGm1' NECTARIS 3 /MS2 mBGMARE SOILS

00~o8ooo'

 

1.00

1.00

1.00

1.00

1.00 

0.96

0.92 

0.88

i iii i i r i i i i i | 

- m4 MS-2 /MS2 MISP I 

- m4 MH-LI (305) /MS2 MIS- - 

-0 11'APOLLO 15 /MS2 LIG- - 

- A CRISIUM 4 /MS2 LIS- - 

- L1 SOMIN 3 /MS2 LBG- - 

o L4 PLATO 7 /MS2 LBS- -

d

- ^f£***f4§*$*«f$ ^ ^i-

- $|I 6O^O "-

: -^ ;

; J ' "%;

\ ,,. *..*«*««**** :

_ t o 0 $-

- 0

- o 

I ^° IMAGING FILTERS '_

Ii 

 

--Uv5 -- Vis' H.^ -NEAR IR -

ill II 1 1_

Me

J

Ayy B W

0.3 0.5 0.7 0.9 1.1 

WAVELENGTH, *m

0.3 0.5 0.7 0.9 1.1 

WAVELENGTH, M m

B

MS2/CONT 9/16/81

MSR/CONT 9/16/81 RP2

RPOLLO 16/CONTJNUUM (.73,1.6)

DESCRRTES 2/CONTJNUUM 11/23/80

Highlands

0.80 1.00 1.20 l.UO 1.60 1.80 2.00 2.20 2.110

WRVELENGTH (MICRONS)

tance spectra help determine mineral and glass contents, the two 

orbital instruments determine elemental compositions independently of mineralogy.

The compositions determined are those of only the surficial 

layers, the uppermost tens of centimeters in the case of the gammaray data and the uppermost tens of micrometers for the X-ray data. 

Investigators are nevertheless confident that most surficial regolith 

materials represent underlying materials because the sharp boundUJ 

CJ

z 

en ~_ t < 

o

UJ

ui:

UJ

cv:

do

o

UJ

CO v '

a flPOLLO 16

b flPOLLO 17

c MflRE SERENITflTIS 2

d LITTROW 1

- i.oo

1.00

- 1.00

-- 1.00

0.80 1.00 1.20 1.40 1.60 1.80 2.00 2.20 2.40

WflVELENGTH (MICRONS)

D

FIGURE 5.19. Typical relative-reflectance spectra of lunar soils (surficial fine materials of

regoliths) obtained by Earth-based telescopes. Courtesy of C.M. Pieters.

A. Spectra of mare soils between 0.3 and 1.1 |xm, relative to an area in central Mare 

Serenitatis (MS2 or MS-2) and scaled to unity at 0.57 (Jtm. Designations to left of area 

names are from classification of Pieters and McCord (1976), and those to right (see table 

5.1) are from Pieters (1978). After Pieters (1978, fig. 3).

B. Additional mare spectra of same type. Slopes of plots emphasize differences of spectra 

from that of comparison area MS 2 (horizontal). Bandpasses of filters used in various 

spectral-imaging studies are shown at bottom (see Pieters, 1978, fig. 3).

C. Near-infrared spectra of areas 5-10 km across, scaled to unity at 1.02 (Jim and divided 

by a straight-line continuum. For a given terrane (highland or mare), spectra of soils (s) 

and interiors of fresh craters (c) are qualitatively similar, but spectral contrast of soil 

spectra is reduced by alteration products in the soils. Absorption bands of highlands are 

centered near 0.90 (jtm, indicating that Ca-poor pyroxene is the dominant mafic mineral; 

absorption bands of maria are near 0.98 (jtm, indicating that Ca-rich pyroxene is the 

dominant mafic mineral.

D. Spectra in the visible and near-infrared, scaled to unity at 1.02 (Jim: a, terra soil around 

Apollo 16 landing site, calibrated with returned soil sample 62231; b, Apollo 17 

landing site, a Ti-rich basaltic soil; c, standard area MS2, probably a low-Ti basalt; d, 

pyroclastic dark-mantling material near the Apollo 1 7 landing site.
98 THE GEOLOGIC HISTORY OF THE MOON

aries between mare and terra and among mare units, as seen on the 

color-difference images (fig. 5.20), have persisted since the maria 

formed (Kuiper, 1965, p. 13). Small fragments of nonlocal material are 

present in all regoliths, though not in sufficient quantity to mask the 

local material.

The gamma-ray spectrometer measures the natural radioactivity of K, Th, and U, as well as gamma rays emitted from other 

elements during nuclear reactions induced by high-energy cosmicray bombardment (Metzger and others, 1973, 1977; Adler and 

Trombka, 1977; Reedy, 1978; Pasaltic Volcanism Study Project, 1981, 

chap. 2). The most useful results are obtained for concentrations of 

Th, K, Fe, Mg, and Ti (Bielefeld and others, 1976; Metzger and others, 

1977; Haines and others, 1978; Metzger and Parker, 1979; Davis, 1980; 

Haines and Metzger, 1980). The Ti values obtained by the gamma-ray 

and spectral techniques agree moderately well (Basaltic Volcanism 

Study Project, 1981, p. 453). The gamma-ray spectrometers were 

limited by low spatial resolution, from 2° by 2° (3,600 km2) to 10° by 10° 

(90,000 km2) of the lunar surface; they analyzed about 22 percent of 

the total surface. On the basis of this areal sample, the maria are 

generally more radioactive than the terrae, and the western nearside 

maria more radioactive than the eastern (Soderblom and others, 

1977).

The X-ray spectrometer detects the elements Al, Mg, and Si from 

the intensity of secondary X-ray fluorescence induced by solar X-rays 

(Adler and others, 1972; Adler and Trombka, 1977). Thus, the data are 

limited to illuminated parts of the lunar surface, and interpretations 

of the data must compensate for variations in solar activity. In all, 

about 10 percent of the lunar surface was analyzed. The Mg/Al ratio is 

most sensitive to terra-mare differences because it is substantially 

greater for mare materials, whereas Si content differs little between 

the maria and the terrae. Highly magnesian and highly aluminous 

mare and dark-mantling materials are also detectable by this ratio, 

other factors being equal (Andre and others, 1977, 1979a; Schonfeld 

and Bielefeld, 1978; Conca and Hubbard, 1979; Hubbard, 1979). Suitable data analysis has resolved areas as small as about 1° by 1° (900 

km2). At these resolutions, ejecta of individual craters and, therefore, 

the subsurface materials exhumed and redeposited in the ejecta can 

be analyzed (Andre and others, 1975,1978,1979b).

Other properties

Two other regions of the electromagnetic spectrum, accessible 

from the Earth or with the aid of a spacecraft, have potential value in 

interpretations of the maria. The first region, in the thermal infrared,

FIGURE 5.20. Color-difference image, produced by subtracting photograph taken at 0.37 |xm from one taken at 0.61 jun. Longer wavelengths (red) appear

brighter and shorter wavelengths (blue) darker. Courtesy of E. A. Whitaker.
5. MARE MATERIALS 99

measures temperatures during a total eclipse of the Moon (Saari and 

others, 1966; Shorthill and Saari, 1969; Shorthill, 1973) or from orbit 

(Mendell and Low, 1975; Schultz and Mendell, 1978). High eclipse 

temperatures are valuable indicators of fresh blocky surfaces and 

thus of youthful craters, which may be undetectable in low-resolution 

photographs (see chap. 13; Winter, 1970; Thompson and others, 1974). 

For mare units, there seems to be a correlation between high eclipse 

temperatures and the youth of a unit (Moore and others, 1980b, p. 

B65). The knowledge gained from craters suggests that this correlation results from the coarse, blocky regoliths that thinly cover young 

mare units.

Lunar applications of radar have been intensively studied. 

Earth-based signals at 3.8 and 70 cm have yielded the most detail, 

including radar maps of the whole nearside disc (Zisk and others, 

1974; Thompson, 1974,1979; Moore and others, 1980b, p. B52-B61). 

Many factors affect radar echoes (Moore and others, 1980b, p. B53). 

For geologic purposes, the echoes ideally would measure depth of 

penetration and surface roughness at scales proportional to wavelength larger than 70 cm for 70-cm radar and larger than 3.8 cm for 

3.8-cm radar. Chemical composition, dielectric constant, electromagnetic absorption, fine-scale roughness, and regional tilt may all additionally affect the echoes. Despite the difficulty of interpretation, 

radar in combination with other remote-sensing techniques is a 

useful measure of crater youth: Depolarized radar echoes from craters decrease with decreasing concentrations of blocks visible in 

photographs and with decreasing infrared eclipse temperatures 

(Thompson and others, 1974,1980). Radar data, infrared eclipse temperatures, and distinctive color spectra together help characterize 

mare units (Schaber and others, 1975; Moore and others, 1980b) and 

reveal the presence of block-free dark-mantling materials that might 

otherwise be taken for lava (Zisk and others, 1977).

During the Apollo 14, 15, and 16 missions, radar signals were 

relayed to and from the Moon by orbiting satellites (Moore and others, 

1980b, p. B34-B41). Though technically highly successful, this 

"bistatic" radar experiment generated little new scientific information beyond the fact that bistatic-radar roughness correlates with the 

visual appearance of a terrain. Young mare units are rougher to radar 

than older mare units.

Integration

The examples already given show that remote-sensing techniques are most valuable when used in concert. Correlations among 

data sets are readily visualized when converted to a common format 

and displayed as false-color images showing any desired weighted 

combination of the data sets (Eliason and Soderblom, 1977; La Jolla 

Consortium, 1977; Soderblom and others, 1977). Geologic mapping 

requires images with good spatial resolution, including low-Sun 

images that enhance topographic detail, high-Sun photographs that 

show brightness differences (figs. 1.1, 5.14), and multispectral images 

(fig. 5.20). To characterize chemical and physical properties not determinable from such images, lower-resolution data from other measurements are located within each resulting map unit and then 

extrapolated to the entire area it covers.

Generalizations based on integration of the data, to be discussed 

further in chapters 11 through 13, include:

1. The western maria in the overflown strips are more highly radioactive than the eastern.

2. Ti-rich, spectrally blue lavas form at least half of Mare Tranquillitatis, the adjoining border of Serenitatis, the lobate flows of 

Mare Imbrium, central Mare Humorum, and much of Oceanus 

Procellarum. Most of these lavas are in maria that overlie the 

central basin and middle trough of the Procellarum basin (pi. 4). 

The Tranquillitatis-Serenitatis units, which include the Apollo 11 

and 17 landing sites, belong to a different high-Ti spectral class 

(HDWA) than do the western units (HDSA, hDSA, hDSP).

3. Class hDWA occurs at several mare margins and fills most of Mare 

Vaporum. On the color-difference photographs, its color appears to 

be intermediate between red and blue and has been called 

"orange" (Basaltic Volcanism Study Project, 1981, p. 237). A patch 

at the Fecunditatis margin includes the Luna 16 landing site.

4. Red spectral classes, thought to be lowest in Ti content, are concentrated in Lacus Somniorum, Mare Frigoris, Sinus Roris, and 

northern Mare Imbrium; that is, they occur in diverse settings 

over the north half of the Procellarum basin. Apollo 15 landed

within one of these classes (LIG-). Red or orange class LISP also 

occurs in Mare Crisium, where it probably underlies the Luna 24 

landing site.

5. Red spectral class mISP forms extensive areas in central 

Serenitatis and the outer trough of western Procellarum. These 

two occurrences may differ in composition (table 5.1); they were 

not sampled.

6. Orange or (partly) red spectral class mIG- covers much of the area 

west of the central meridian and south of the strong blue-red 

association in Mare Imbrium. Apollo 12 landed within this belt. 

This class also forms much of Mare Crisium and Mare Fecunditatis. These concentrations may have little significance, however, 

because class mIG - is a group of low- to medium-Ti mare units 

that may differ compositionally (Pieters, 1978).

7. One interpretation of the orbital geochemistry is that Maria 

Smythii, Fecunditatis, and Crisium are all Al-rich (Conca and 

Hubbard, 1979; Hubbard, 1979). Orange spectral classes mBG-, 

characteristic of Mare Nectaris, and LBG-, common among the 

northern very red units, also are probably Al-rich (C.M. Pieters, 

oral commun., 1982).

THICKNESS

Recognition that the lunar mare surfaces are underlain by threedimensional deposits led early to measurements of their thickness. 

Marshall (1961) estimated thicknesses in Oceanus Procellarum from 

the degree of obliteration of flooded craters, whose preburial profile 

was assumed to equal that of fresh craters of the same diameter 

(Baldwin, 1949, p. 128-138). Baldwin (1970), Eggleton and others 

(1974), and De Hon (1974) refined this method on the basis of better 

data for the profiles and rim dimensions. Mare-basalt thicknesses, 

including thicknesses of individual flows, also subtly affect the sizefrequency distributions of craters (Eggleton and others, 1974; Neukum and others, 1975a). R.A. De Hon and his coworkers systematically extended these measurement techniques to most of the nearside 

maria (fig. 5.21). They found that the eastern mare materials average 

200 to 400 m in thickness (De Hon and Waskom, 1976) and the 

western closer to 400 m (De Hon, 1979). Lenses thicker than 1,200 or 

1,500 m also occur; more accurate values are hard to determine 

because few craters are visible in such thick mare sections (De Hon, 

1979).

The accuracy of other mare-thickness determinations is uncertain. On the one hand, thicknesses will be overestimated if the craters 

on which they are based were highly degraded or filled before the 

mare flooding (Hbrz, 1978), or if they have been shallowed by floor 

uplift (as is, in fact, observed in basins; see chap. 6). Horz (1978) 

estimated that De Hon's figures are a factor of 2 too high. On the other 

hand, thicknesses will be underestimated if the craters are perched 

on older mare materials. Head (1979a) pointed out that the flooding of 

the present central highlands of the nearside, which must resemble 

the floors of some of the older basins, would require more mare 

material than believed by either De Hon (1979) or Horz (1981).

A few geophysical determinations of thickness are available. The 

gravity data are consistent with thicknesses of 2 to 4 km in the 

centers of mascon maria (Sjogren and others, 1974; Solomon and 

Head, 1979, 1980), the largest of which are Imbrium, Serenitatis, 

Crisium, Humorum, Smythii, Nectaris, and Orientale (fig. 5.22; table 

6.1; Solomon and Head, 1980). Large nonmascon maria include Tranquillitatis, Fecunditatis, and most of Procellarum. The 20- to 25-kmdeep seismic discontinuity under Oceanus Procellarum was first 

interpreted as the base of the basalt section (Toksoz and others, 1974) 

but is now thought to represent a physical or chemical discontinuity 

within the terra crust (see chaps. 1, 8). Most of the mare fill of 

northern Oceanus Procellarum is thinner than 500 m. However, in 

one place in the north and in several places in the south, it reaches 

thicknesses greater than 1,000 m where superposed basins or large 

craters deepen the original Procellarum-basin floor (De Hon, 1979).

A radar sounder flown on the Apollo 17 mission directly measured depths to discontinuities in two mascon maria, Serenitatis and 

Crisium (Peeples and others, 1978). If certain assumptions about the 

dielectric constant of the overflown materials are correct, one horizon 

in Mare Serenitatis lies about 1 km beneath the surface, and another 

1.6 km deep in the west to 2.0 km deep in the east (Peeples and others, 

1978). One horizon apparently within the basalt section of Mare
100 THE GEOLOGIC HISTORY OF THE MOON

Crisium lies 0.8 to 1.0 km below the surface of the shelf that separates 

the conspicuous ridge system and the basin rim (figs. 4.9, 5.23; 

Maxwell and Phillips, 1978; Peeples and others, 1978). What appears 

to be the same interface (Maxwell and Phillips, 1978) is 1.4 km deep 

inside the shelf.

Mare thicknesses have also been inferred from the geochemical 

experiments (Andre and others, 1979b). The craters Peirce (19 km) 

and Picard (23 km) penetrated the 1.4-km-thick layer in Crisium and 

excavated additional basalt whose Mg-rich chemistry is evident in the 

X-ray data (fig. 5.235; Andre and others, 1978,1979a, b). These craters 

may have excavated terra materials from beneath the mare (Head and 

others, 1978a). On the basis of assumed excavation depths, the total 

mare thickness in Crisium may be 2.4 to 3.4 km (Maxwell and Phillips, 1978).

Progressive decrease in the number of visible premare craters, 

the gravity data, and the radar-sounder profiles indicate that the 

mare materials are thickest in the mare centers (Baldwin, 1970). This 

relation reflects the topography of the host basin: Each concentric 

shelf of a basin is lower than the surrounding one, and the center is 

the lowest of all (chap. 4; figs. 4.9, 5.16, 5.17, 5.23). Elevations of the 

mare surfaces also mimic this step structure to some degree. For 

example, Mare Crisium lies about 2 km lower than peripheral Mare 

Spumans and 3 km lower than Mare Undarum (fig. 4.9). However, the 

mare surface drops in elevation less than the basin floor because of the 

inward mare thickening.

90 W

--30 N

ISOPACH MAP OF WESTERN MARE BASALT

-30 S

50V700 EDO

KILOMETERS

300 20C IOO/\M 2i5 Q 2j5 50/^0 200 300 <IH 500 EDO 700^50° 

~ ~" '

 /

20° 30° 40° 50° 60 7O° 8O° 9O° 1OO

 I I I I I i I i -3d

FIGURE 5.21. Thickness of mare materials on western (A; De Hon, 1979) and eastern (B; 

De Hon and Waskom, 1976) nearside.

FIGURE 5.22. Lunar gravity anomalies (from Muller and Sjogren, 1968).

5 -

B

FIGURE 5.23. Thickness of mare-basalt flows in Mare Crisium.

A. Western Mare Crisium and part of Crisium basin, showing difference in elevation 

(arrows) between mare shelf and depressed center. P, crater Peirce (19 km, 18° N., 

53.5° E.). Dashed lines indicate basin-concentric troughs, as interpreted by Wilhelms 

(1970b). Telescopic photograph.

B. Diagrammatic cross section of Mare Crisium, based on radar-sounder profile (Peeples 

and others, 1978). Arrow denotes shelf indicated by arrows in A. Crater P (Peirce; also 

Picard) penetrates younger mare unit and excavates older mare unit (Mg-rich; Andre 

and others, 1978) and, possibly, terra underlying the mare section (Head and others, 

I978a). Scales are diagrammatic, exaggerated vertically.
5. MARE MATERIALS 101

In summary, enough mare thicknesses have been measured to 

indicate that mare basalt is volumetrically minor. Few depths are 

likely to exceed 4 km; the thickest sections are in the centers of the 

circular maria. Thicknesses on mascon-mare shelves and nonmascon 

maria are no more than about 0.5 km. Likely lower and upper limits 

for the average mare-basalt thickness on the whole Moon are probably 0.2 and 1.0 km, respectively. Thus, the upper 75 km of the Moon 

would consist of 0.24 percent mare basalt if the average thickness is 

1.0 km, and of only 0.047 percent if it is 0.2 km (based on coverage of 17 

percent of the total lunar area of 38 x 106 km2 ). Subsurface dikes and 

sills would add an unknown volume to these small values.

RETURNED SAMPLES

Introduction

The samples of basalt returned from the lunar maria present 

fewer problems than do those of breccia from the terrae, because most 

lavas acquire their chemical and physical properties when their magmas originate and when they are emplaced, and not from an earlier 

geologic unit. Most rock-size basalt samples are not shock-melted or 

metamorphosed. Their textures originated when the unit from which 

they were collected solidified. Their radiometric ages almost unambiguously date the magma's crystallization and, therefore, the unit's 

emplacement. No doubt remains that the sampled basalt flows originated as magmas melted within the Moon.

All nine sampling missions returned mare basalt, in greatly 

varying amounts. Apollos 11 and 12 landed on maria (table 1.2) and 

returned 16 and 36 rock-size (min 1 cm) basalt samples, respectively 

The Apollo 11 material established the important fact that the maria 

are ancient in absolute age, though young stratigraphically. Subsequent studies showed that they were derived by partial melting from 

an already differentiated mantle source. Apollo 12 collected the 

youngest sample suite (see chap. 12) and proved that maria differ in 

age and composition from region to region. The remaining missions 

found a wide scatter in age and composition (see chap. 11), and partly 

clarified and partly complicated the picture of lunar-basalt 

petrogenesis. Apollos 15 and 17 were multipurpose missions that 

landed on mare surfaces and returned large amounts of basalt. Lunas 

16 and 24 collected small cores of mare regolith. Even Apollos 14 and 

16 and Luna 20, which landed on the terrae, obtained small bits of 

mare basalt. All the landing sites were near the margins of maria (pi. 

4) or in the terrae (pi. 3). Therefore, only relatively thin basalt sections, and no central mascon sections, were sampled.

Composition

Lunar-mare rocks are designated "basalt" by the same general 

criteria used to define terrestrial basalt (Hubbard and Gast, 1971): 

They are dark-colored, mafic, fine-grained extrusive or shallow-intrusive rocks. Their textures are familiar from terrestrial basalt, except 

that they are fresher because they have not been chemically 

weathered (fig. 2.6). Clinopyroxene and plagioclase together compose 

75 to 90 percent of most mare basalt; pyroxene is the more abundant 

mineral (Papike and others, 1976). Some compositional groups contain as much as 20 percent olivine and 24 percent optically opaque 

minerals (Fe-Ti oxides, of which ilmenite is the most abundant; 

Papike and others, 1976, table 3).

In elemental abundances, lunar-mare and terrestrial basalt differ substantially except in Ca content (10-11 weight percent CaO) 

(table 5.2; Hubbard and Gast, 1971; Taylor, 1975, 1982; Papike and 

others, 1976; Basaltic Volcanism Study Project, 1981, chap. 1). Lunarmare basalt (1) contains no detectable H2O; (2) is very low in alkalis 

(less than 1 weight percent), especially Na2O; (3) is generally high in 

TiO2 low lunar and average terrestrial contents are about the 

same; (4) is low in A12O3 (8-15 weight percent) and SiO2 (mostly 

39-49 weight percent); maximum lunar and average terrestrial contents of these oxides are about the same; (5) is very high in FeO (17-22 

weight percent) higher than terrestrial basalt; and (6) is generally 

richer in MgO (7-18 weight percent) the lowest lunar and average 

terrestrial contents are about the same. An additional important 

difference is in the extreme degree of reduction: Lunar-mare basalt 

contains essentially no Fe3 + ; most Fe occurs as Fe2 + ; and a minor

amount of native Fe, formed by reduction at or near the surface, is 

always present (Basaltic Volcanism Study Project, 1981, chaps. 1, 3; 

Tkylor, 1982, p. 292-294, 312).

Mare basalt also differs chemically from terra melt rocks of 

impact or volcanic origin. It is higher in FeO and MgO contents and 

lower in A12O3 and CaO contents and Al2O3/CaO ratio. Its rare-earthelement patterns, as normalized to those of chondritic meteorites, 

generally have negative europium (Eu) anomalies, in contrast to the 

positive Eu anomalies of most terra materials (see chap. 8). Lunar 

basalt is very low in siderophile elements, elements that are concentrated in iron meteorites (see chap. 8). It differs further from impact 

melts in its lack of included debris. Lunar-mare basalt and terra 

materials thus, the whole Moon are very poor in volatiles (Wetherill, 1971; Tkylor, 1975; 1982, p. 300-307).

Pyroclastic glasses, which constitute the dark-mantling materials, are partly similar and partly dissimilar to the mare lavas in 

bulk composition. Orange and black glasses from the Apollo 17 landing site are rich in Ti, but differ in trace elements and Mg from the 

high-Ti lavas in the same region (Heiken and others, 1974). Apollo 15 

collected both Ti-poor green glasses and Ti-rich red glasses, neither of 

which are identical to the Ti-poor Apollo 15 lavas (Delano, 1979,1980). 

Smaller amounts of other pyroclastic glasses are mixed in other 

sampled regoliths (Wood, 1975a).

Classification

Sampled mare basalt is generally classified on the basis of majorelement content (table 5.2). Some authors distinguish sharply 

between high-Ti and low-Ti groups because the largest collections are 

very distinct in this element; Apollos 11 and 17 returned exceptionally 

Ti-rich basalt, whereas Apollos 12 and 15 returned low-Ti basalt (table 

5.2; Papike and others, 1976; Taylor, 1975). The Ti gap is occupied by a 

few samples from the Apollo 12 landing site and small fragments from 

the Apollo 14 terra landing site (in breccia sample 14063, 7.3 weight 

percent TiO2 ; Ridley, 1975). Small fragments of a very low titanium 

(VLT) group were found later in the Apollo 17 and Luna 24 regoliths 

(Papike and Vaniman, 1978). High-Al, feldspathic basalt is recognized 

as a distinct category, although it also was recovered in only small 

amounts (Ridley, 1975; Taylor, 1975,1982; Taylor and Jakes, 1977); it 

characterizes collections from the eastern maria (see chap. 11) and 

premare basalt of the mare type (chap. 9). Small rare glass droplets 

and fragments found in all regoliths may represent additional 

magma types (Wood, 1975a; Binder and others, 1980). The continuum 

of spectral classes (Pieters, 1978) suggests that the gaps in Ti content 

and in other discriminative compositions are filled by unsampled 

basalt types (Papike and Vaniman, 1978). Only about a third of the 

observed number of spectral classes may have been sampled (Pieters, 

1978).

There is now good agreement about which groups are significant, 

but less agreement about the classifying nomenclature. Most designations include the sampling mission. The scheme used here (table 

5.2; Basaltic Volcanism Study Project, 1981, sec. 1.2.9; Taylor, 1982, p.

TABLE 5.2. Classification of mare-basalt samples used in this volume

[After Basaltic Volcanism Study Project (1981, sec. 1.2.9.) and Taylor (1982, table 6.1). All

values in weight percent]

Group Distinguishing 

chemistry

Apollo 11 high-K high Ti basalt High K, high Ti 

Apollo 12 ilmenite basalt Intermediate Ti 

Apollo 12 pigeonite basalt Low Ti, high Si 

Apollo 12 olivine basalt Low Ti, high Fe, Mg

Luna 24 very low Ti basalt Very low Ti, high Al

Apollo 12 feldspathic basalt High Al , low Ti 

Apollo 16 feldspathic basalt High Al, low to 

intermediate Ti.

A1 2 0 3

8-10

Ti02

5-9

<1.5

K2 0

>0.3

.03-. 11

<0.04

MgO

7-10

10-18

10-11
102 THE GEOLOGIC HISTORY OF THE MOON

282-284) lists the main chemical properties that distinguish 

sampled groups. Most group names also include the name of a characteristic mineral observed optically (a modal mineral); other classifications refer to a mineral that would theoretically crystallize from a 

melt of the rock's composition (norm). For example, the terms 

"pigeonite (low-Ca clinopyroxene) basalt" (table 5.2), "Si-rich basalt," 

and "quartz-normative basalt" have all been applied to the same 

Apollo 12 or Apollo 15 samples.'The Apollo 17 basalt suite is subdivided by trace-element contents (see chap. 11; Rhodes and others, 

1976; Warner and others, 1979). Although textural terms are also used 

to designate some groups (chap. 11), classifications based on texture 

cut across those based on sampling site and chemistry (Warner, 1971). 

Textural variety within a basalt suite from a given site reflects differences in the crystallization and cooling histories of the lavas more 

than differences in composition (Lofgren and others, 1975). Some 

compositional differences within a given site have resulted from 

minor fractionation after emplacement, though fewer than in many 

terrestrial lavas (Taylor, 1982, p. 334).

ORIGIN AND 

EMPLACEMENT

The mare-basalt magmas originated by partial melting of ultramafic mantle material (mostly or entirely olivine and pyroxene) 

(Basaltic Volcanism Study Project, 1981, chaps. 1, 3, 4, 9). The Eu 

anomaly and other trace-element data demonstrate that the mantle 

does not consist of primitive, undifferentiated lunar material but 

segregated from the bulk Moon soon after it formed 4.55 aeons ago 

(see chap. 8; reviews by Taylor, 1975,1982).

As discussed in chapters 12 and 13, known basalt flows were 

extruded until 3.5 aeons after the Moon formed. Thus, a substantial 

time gap separated the global differentiation that formed the mantle 

source from the partial melting that led to basalt extrusion. Because 

formation of mare basalt was a two-stage process, a source of heat in 

addition to that remaining from planetary accretion was apparently 

necessary to initiate the second-stage melting (Taylor, 1975, 1982). 

The amounts of radioactive K, Th, and U believed to have been 

present in the source regions were sufficient to generate enough heat 

to melt the very small observed amounts of basalt (Taylor, 1982, p. 

310). These radioactive elements are most abundant in the Ti-rich 

basalt.

This partial melting is generally thought to have taken place 

between 60 and 500 km below the surface (Taylor and Jakes, 1974; 

Green and others, 1975; Kesson and Lindsley, 1976; Delano, 1979, 

1980; Taylor, 1982, p. 325). More precise estimates are controversial, 

especially concerning high-Ti basalt. Various schemes have associated basalt composition with source composition, but, again, the data 

are equivocal and the interpretations diverse. For example, high-Ti 

basalt may be derived from clinopyroxene-rich zones, VLT magmas 

from olivine-orthopyroxene zones, and high-Al basalt from 

clinopyroxene-plagioclase zones (see summary and references by Taylor, 1982, p. 330). The mantle may be crudely layered; one suggestion 

is that the low-Ti source lies 150 to 250 km below the surface, the highTi source 100 to 150 km deep, and the high-Al source 60 to 100 km 

deep, immediately beneath the crust (Taylor and Jakes, 1974). Alternatively, the high-Ti source may be lower than the low-Ti source 

(Nyquist and others, 1977). Delano (1979, 1980) suggested that the 

pyroclastic glasses originated at the greatest depths, possibly 400 to 

500 km below the surface, in an environment where carbon-oxygen 

gases may have formed (Sato, 1979).

The presence of glass and the vesicularity of many lavas are the 

only major indications of volatile activity on the Moon. The fact that 

their eruptions were driven by volatiles may explain why patches of 

dark-mantling materials appear where no lavas are known, as in such 

isolated craters as Petavius and Humboldt (fig. 4.2).

Whether or not layering exists in the mantle, the variety of 

sampled and unsampled compositions indicates considerable vertical 

and horizontal heterogeneity in the basalt-source zones. In particular, the compositional variety of basalt erupted at a given site 

simultaneously or within a very short time suggests that the sources 

were compositionally heterogeneous over short distances (see chaps. 

11,12; Green and others, 1975; Rhodes and others, 1976,1977; Nyquist 

and others, 1979a; Taylor, 1982, p. 301, 320-321).

After the magmas were generated, most were modified in the 

original magma chambers, on the way to the surface, or on the surface 

(Basaltic Volcanism Study Project, 1981, p. 399-408, 498-513, 

577-591). On the Moon as on the Earth (and probably all planets), 

early-formed crystals fractionated from the liquid magma, mantle or 

crustal materials were assimilated, and magmas from separate 

chambers mingled. Only a few magmas may have been primary in the 

sense that they remained unchanged after melting of the mantle 

source.

This volume can best contribute to the multifaceted question of 

mare-basalt origin by examining the geologic setting and emplacement histories of the observed mare units. Therefore, I consider why 

maria are concentrated in basins, especially in particular basins.

Basic hydrostatic principles suggest that extrusion of basaltic 

magmas is favored by (1) depressed surface elevations, (2) thin crusts 

and lithospheres, (3) low magma densities, and (4) deep melting sites 

(Solomon, 1975; Wilson and Head, 1981). Conditions 1 and 2 are most 

favorable in basins. A basin impact removes several kilometers or tens 

of kilometers of crustal material. The resulting depression is greatly 

reduced by an uplift of the mantle to compensate for the lost mass 

(chap. 4) but will partly survive if the lithosphere is strong. Addition 

of the ejecta also thickens the crust near the basin by a few kilometers. The depression not only collects whatever magmas are 

extruded but also contains most of the extrusion vents (fig. 5.24). Even 

if magmas melt at the same depth and ascend with equal freedom 

everywhere on the Moon, more would be extruded in the depressions 

than on the adjacent terra. A rising column of magma that would be 

erupted in a basin may only be intruded as dikes and sills outside the 

basin.

Basins may also affect extrusion in more active ways. A basin 

impact generates heat that may aid the melting and weaken the 

lithospheric barrier to magma ascent (Solomon and Head, 1980, p. 

134). Subbasin structures, such as fractured or brecciated zones similar to those beneath craters, may facilitate magma movement and 

divert some magmas to extrusion sites high on the basin periphery 

(figs. 5.10D, 5.16; G, fig. 5.24).

Superposition of two or more basins or large craters enhances the 

likelihood of mare extrusion at a given spot. The first or the only 

extrusions will occur in the superposed excavations (Head, 1976a).

FIGURE 5.24. Diagrammatic cross section of typical basin 1,000 km in diameter; no vertical 

exaggeration. Crust is about 75 km thick outside basin (see chap. 1) and is estimated at 55 km 

thick beneath basin center; deformed zone (above dotted line) is estimated at 25 km thick, and 

mantle uplift at 20 km. Magma-ascent paths, 60 km high above mantle surface, are indicated 

diagrammatically in nine places (A-I) as systems of dikes and sills. At A, lava was extruded 

onto floor of a crater superposed on the basin, and spilled beyond crater rim; at B, lava was 

extruded directly onto basin floor; extrusions A and B combined to flood basin center. At C,

lava was extruded onto trough but did not overtop trough edges. At D, crust was too thick to 

allow an equal column of magma to be extruded. At E, floor of superposed crater decreased 

magma-ascent path by amount necessary to allow extrusion; crater interior was flooded. At F, 

contact of deformed zone and less severely deformed crust afforded path for lateral migration 

of lava, which was extruded high on basin flank at G. Crust at H and I outside basin was too 

thick to allow extrusion, even beneath a crater (I) equal in depth to the flooded craters inside 

basin.
5. MARE MATERIALS 103

This association is clearly illustrated on the farside, where such 

craters as Leibnitz, Lyot, and von Karman and such small basins as 

Apollo, Planck, and Poincare, all of which are superposed on the 

South Pole-Aitken and Australe basins, contain most of the farside's 

maria (pi. 4; fig. 5.25). In my opinion, similar superpositions of the 

still-larger Imbrium, Serenitatis, Tranquillitatis, Nubium, and other 

basins on the 3,200-km-diameter Procellarum basin cause the familiar concentrations of nearside maria.

Factors other than basin distribution have also been thought to 

affect mare distribution. Differential crustal thickness is commonly 

cited as the explanation of the 15-fold hemispheric dichotomy in mare 

distribution (pi. 4); the farside crust is thought to be thicker than the 

nearside crust (chap. 1). The Procellarum basin, however, may cause 

the difference. All the Moon's major maria except Crisium, Fecunditatis, Humboldtianum, Nectaris, Orientale, and Smythii lie in 

basins superposed on Procellarum (fig. 5.26). Of these maria, only 

Crisium contains a thick section of mare, and Mare Crisium is much 

smaller (500 km) than would be expected from the Moon's fourth 

largest basin (1,060 km diam; chap. 9; table 4.1).

The relative importance of magma density and melting depth is 

difficult to evaluate because of the many unknowns. High-Ti magmas 

are the densest, and high-Al magmas the least dense (Solomon, 1975). 

In the scheme of Taylor and Jakes (1974), low-Ti magmas originate at 

greatest, and high-Al magmas at shallowest, depths. Thus, density 

and depth may cancel each other as factors favoring extrusion. Basincontrolled crustal thickness may play a greater role in affecting the 

composition of extruded lavas. High-Al magmas preferentially fill 

basins formed in thick crusts (Fecunditatis, Nectaris, Smythii), 

whereas low-Ti magmas preferentially form over thin crusts (in the 

Procellarum basin, especially where Serenitatis and Imbrium are 

superposed on that giant basin). These factors are further evaluated 

in chapters 11 and 12 after the historical dimension has been added.

FIGURE 5.25. Eastern nearside, east limb, and part of farside of Moon, showing spotty 

distribution of mare patches in sparsely flooded Australe basin (lower right edge), peripheral 

"lakes" concentric with Crisium basin, and other maria outside main mare concentration. 

Apollo 11 frame H-6665, taken during return to Earth after first lunar landing.

B

FIGURE 5.26. Procellarum basin as drawn by Whitaker (1981). Courtesy of E.A.

Whitaker.

A. Telescopic full-Moon photograph of nearside, showing center (P) and parts of three rings 

of Procellarum basin. I, center of Imbrium basin. Maria are concentrated inside largest 

ring of Procellarum basin. In much early work, an asymmetry in basin distribution was 

thought to cause the hemispheric dichotomy in mare distribution, but basins are distributed randomly (pi. 3); Procellarum may account for most of the dichotomy. Right 

side of photograph overlaps with left side of figure 5.25. From Whitaker (1981, fig. 2).

B. Observed scarps and terra margins concentric with basin (large dots), complete rings that 

best fit the topography (small dots), mare ridges (lines; see chap. 6), and major 

landmarks. From Whitaker (1981, fig. I).

6.STRUCTURE

m

' ; * 

. */.
FIGURE 6.1 (OVERLEAF). Lee-Lincoln scarp west of Apollo 17 landing site, partly covered by light-colored landslide of debris from South Massif (bottom). Scene centered at 20.25° N.,

30.75° E., covers 7.5 km north-south. Apollo 17 frame P-2309.
6. STRUCTURE

CONTENTS

Page

Introduction ...................................................... 107

Structures in maria and basins ..................................... 107

Mare ridges ................................................... 107

Arcuate rilles .................................................. Ill

Origin by mare subsidence ..................................... 112

Crater-floor fractures ............................................. 113

Straight rilles ..................................................... 113

General features ............................................... 113

Imbrium sculpture .............................................. 113

Other systems ................................................. 115

Interpretation lithospheric thickness ............................... 115

INTRODUCTION

Tectonic structures, like endogenic materials, play a smaller 

role in lunar geology than was once widely thought. Miscellaneous 

aligned landforms were interpreted as fault-controlled during much 

of the telescopically based mapping of the 1960's. Geologists commonly mapped faults without the objectivity that proved so valuable 

in mapping material units. Basin rings were generally mapped as 

fault-bounded, and lineament trends were commonly assigned to the 

"lunar grid," a hypothetical system of conjugate fractures (Fielder, 

1961, 1965; Strom, 1964; Mutch, 1970, p. 247-250; Casella, 1976). 

High-resolution photographs, however, have substantiated the view 

of some astute early observers that faults are rare on the Moon 

(Arthur, 1962, p. 321; Baldwin, 1963, chap. 22). Except for the faults 

described in this chapter and for relatively minor gravitational faulting along some rings, I believe that most supposed faults are coincidental alignments of unrelated features. The lunar grid is probably 

an artifact of the plotting methods (Wise, 1982).

The present lunar surface, therefore, has not been extensively 

reshaped by endogenic forces, as has the Earth, but generally 

retains the pattern imposed by cratering and deposition of stratigraphic units. The dominance of lunar geology by basins and maria 

extends to most of the tectonic structures that do exist. Most of these 

structures are inside basins, and many transect both the basin and 

the mare (pi. 5). Some structures at mare margins that are related to 

dark pyroclastic materials and irregular craters have already been 

illustrated (figs. 5.10, 5.16). This chapter more thoroughly describes 

and interprets basin-related mare ridges, arcuate rilles, and craterfloor fractures.

Some straight rilles of fault origin also cut surfaces outside the 

most conspicuous mare-filled basins. These rilles have not yet been 

satisfactorily explained. Global deformation is commonly invoked, 

but this chapter presents alternative, basin-related interpretations 

for these structures as well. In particular, the 3,200-km-diameter 

Procellarum basin (Whitaker, 1981), whose importance to terra 

topography, lithospheric thickness, and mare localization has 

already been mentioned (chaps. 4, 5; pis. 3, 4; fig. 5.26), is suggested 

as a major control on the straight rilles as well as the ridges and 

arcuate faults.

STRUCTURES IN 

MARIA AND BASINS

More ridges

Deformation inside mare-filled basins is most manifest in two 

classes of structures mare ridges and arcuate grabens. Mare ridges, 

also known as wrinkle ridges and now formally named dorsa, are 

linear positive features that occur in most maria (pi. 5). Most mare 

ridges are on the nearside because most maria are there, but even 

small isolated farside maria have ridges (fig. 5.2). Also, some ridges 

extend from mare margins into the terrae, for example, the "LeeLincoln" scarp at the Apollo 17 landing site (fig. 6.1) and some basinradial ridges south of Mare Humorum (fig. 9.24B; Saunders and 

Wilhelms, 1974). The ridges generally occur in systems that are 

subconcentric and subradial to maria (pi. 5; figs. 5.17, 5.23A, 6.2). 

Systems in Oceanus Procellarum (fig. 6.3), which appear to be linear 

and parallel (Scott, 1974), are concentric with the Procellarum basin 

(fig. 5.26£; Whitaker, 1981).

In detail, mare ridges consist of linear segments that merge, 

overlap, or are arranged in echelon (figs. 6.2, 6.3); many extend as 

complex, branching echelon systems for hundreds of kilometers. 

Ridges generally have two distinct morphologic parts a broad arch 

and a narrow superposed spine (Strom, 1971). Some arches are composed of overlapping domes. They range from several kilometers to as 

much as 10 km in width and commonly are about 100 m high; in Mare 

Serenitatis, some arches reach a height of 350 m (fig. 6.2; 

Muehlberger, 1974). Such large arches appear on all photographs, but 

many others appear only on photographs taken at low-Sun illuminations (less than 15°) and are not necessarily disclosed by high photographic resolutions. Many arches are asymmetric, bounded on one 

side by scarps or monoclinal bends that offset the mare surface by 50 

to 100 m (Lucchitta, 1976). In places, the asymmetry shifts sides along 

the length of the ridge. Some scarps resemble flow lobes in shape. The 

sharp spinelike ridges commonly form shorter, narrower, tortuously 

braided networks along the crests and flanks of the broader arches, 

and may occupy 25 to 60 percent of the total arch width (figs. 6.2, 6.3; 

Strom, 1971). The spines are commonly 100 m high and about 200 m 

wide (Lucchitta, 1976). Some spines extend from the arches onto the

107
108 THE GEOLOGIC HISTORY OF THE MOON

FIGURE 6.2. Border zone between Mare Serenitatis and Serenitatis-basin rim, showing structures concentric with mare. View northward.

A. West and southwest border. Rugged terra is Monies Apenninus. Grabens concentric with and Iruncaled by mare are Rimae Sulpicius Callus (compare fig. 5.16). Apollo 17 frame M 952.
6. STRUCTURE 109

fi. East border. Twosetsofgrabens diverge upward from near bottom of photograph; one is closely concentric with the mare (a), and the other diverges eastward (b). L, LeMonmer(61 km, 27° 

N., 31° E.). Dark-mantling material at bottom (d) is truncated by lighter central mare. Arrow, Apollo 17 landing site (compare fig. 5.17). Apollo 17 frame M-940.
110 THE GEOLOGIC HISTORY OF THE MOON

adjacent flat mare surfaces; others are solitary. In some places, both 

the arches and the narrow spines are superposed on gentle linear 

rises, as much as 500 m high and 25 km wide, that can be detected 

only on low-Sun images or by topographic data (Lucchitta, 1976).

Two main schools of interpretation of mare ridges have emerged. 

The volcanic school favors intrusion and extrusion of lava, controlled 

by global or basin-related tectonic patterns (Fielder, 1965; Quaide, 

1965; Hartmann and Wood, 1971; Strom, 1971; Scott, 1974), or a variant autointrusion of lavas into fractures (Hodges, 1973c). The 

tectonic school, which was founded by Baldwin (1963, p. 380-382; 

1968), favors structural deformation of solid materials and is supported by the studies discussed in this chapter. Most investigators 

agree that many circular wrinkle ridges were formed by compaction 

over crater rims (fig. 6.4). Other, larger circular patterns of mare 

ridges are ascribed to settling of mare sections over buried basin 

rings; such ridges are, in fact, the only basis for locating the inner 

rings of Imbrium, Serenitatis, Crisium, and other deeply filled basins 

(table 4.1; Hartmann and Kuiper, 1962; Hartmann and Wood, 1971; 

Wilhelms and McCauley, 1971; Maxwell and others, 1975; Brennan, 

1976).

Many exposed parts of crater rims, basin rings, and other buried 

terrae assume domelike landforms (figs. 6.5A, B). These landforms 

and the associated ridges have been interpreted as ringlike volcanic 

complexes. The moldinglike accumulations of material at the bases of 

the terra remnants resemble volcanic flows (O'Keefe and others, 1967; 

Strom, 1971), and their light color implies silicic composition. 

However, the "moldings" are common along contacts between mare 

surfaces and terra slopes, including those of obvious impact features 

(figs. 6.1, 6.5), where they consist of debris accumulated from the 

slopes (Milton, 1967; Offield, 1972). Therefore, these accumulations 

are stratigraphically younger than the ridges and the mare materials 

at the bases of the terra islands, whereas their source bedrock is older 

than the maria.

Evidence for dislocations of nonmare units (figs. 6.6, 6.7) has led 

to elaboration of the tectonic hypothesis. The complex ridge morphology suggests crumpling of the surface under compressional 

stress. Howard and Muehlberger (1973) suggested that the compression created thrust faults along intramare gliding horizons. Bryan 

(1973) and Maxwell and others (1975) explained the compression as a

FIGURE 6.4. Mare ridges, formed by subsidence of mare basalt over rim of crater Lambert 

R (bottom; 55 km, 24° N., 20.5° W.). Unburied crater with shadowed interior above center 

is Lambert (30 km). Apollo 15 frame M-1010.

7^^mm^&'

FIGURE 6.3. Complex mare ridges and flooded terra islands in southern Oceanus Procellarum near crater Letronne (left edge of left frame; Colton and others, 1972). Some ridges and kipukas or 

steptoes of elevated, presumably thin maria (t) mark subsurface extensions of crater rims and other terra. Photograph covers area of lat 7.0° 13.5° S., long 32.5° 41.5° W. Object at center 

right of both frames is boom of Apollo 16 gamma-ray spectrometer. Stereoscopic pair of Apollo 16 frames M 2837 (right) and M 2839 (left).
6. STRUCTURE 111

buckling of the surface caused by settling of the maria into a reduced 

area; thick sections of mare basalt settle most, and the stresses are 

concentrated where the basalt thins (Maxwell and others, 1975). 

Lucchitta (1976, 1977b) also favored settling and suggested that the 

ridges lie along fault systems that include both normal and reverse 

faults, which are manifested by compression at the surface (fig. 6.8). 

In these hypotheses, accepted here, most ridges originate basically by 

compression resulting from vertical tectonism, although some are 

probably of volcanic origin (fig. 5.13).

Arcuote rilles

Arcuate rilles (rimae) are flat-floored, steep-walled troughs (figs. 

6.9, 6.10, 6.11). Each trough consists of a floor a few kilometers wide 

and linear subparallel facing scarps 50 to 250 m high (Golombek, 

1979). The troughs are grabens created by extension, like many similar grabens on the Earth (Baldwin, 1949, p. 197-199; 1963; Quaide, 

1965; McGill, 1971; Golombek, 1979). Like the ridges, most arcuate 

grabens occur in parallel, crosscutting, or echelon sets; some occur

FIGURE 6.5. Moldmglike accumulations of debris at bases of partly inundated impact 

features. 

A. Light-colored domelike features along mare ridge are terra islands that partly deflect the

ridges. Apollo 16 frame H-19244. 

B. Part of crater Flamsteed P ("Flamsteed Ring") concentric with mare ridge (arrow).

Obiter 3 frame M-181. 

C. Interior of crater Maraldi. Apollo 17 frame P-2302.
112 THE GEOLOGIC HISTORY OF THE MOON

singly. Most grabens are concentric with circular maria and the 

containing basins (pi. 5; figs. 5.16, 5.17, 6.9-6.11). Moreover, they lie 

within the topographic-basin rim as that rim is interpreted in chapter 

4 (except for some subconcentric grabens east of Serenitatis, which 

may be controlled by the middle Procellarum-basin rim; set b, fig. 

6.2B). Grabens cut both the mare and the basin material (figs. 6.9, 

6.10). With minor exceptions, they occur only in basins superposed on 

the Procellarum basin, particularly Imbrium, Serenitatis, Tranquillitatis, and Humorum (pi. 5), and are rare or absent in the large 

basins Crisium, Nectaris, and Smythii and on the farside.

Origin by mare subsidence

The recurring association of mare ridges and arcuate grabens 

suggests a genetic connection (Baldwin, 1963,1968). Their geometric 

relation is similar in many basins: Subradial ridges innermost, arcuate rilles outermost, and concentric ridges in between. Discovery of

FIGURE 6.6. Crater Bonpland D (6 km, 10° S., 18° W.) in Mare Cognitum, offset along 

mare ridge. Apollo 16 frame P 5429.

FIGURE 6.7. Ridges continuous between mare and terra and appearing to thrust material 

from bottom over material at top. Scene centered at 7° S., 29° W., in Monies Riphaeus. 

View northward. Apollo 16 frame P-5452.

the mascons (chap. 5; Muller and Sjogren, 1968) led to the currently 

prevailing interpretation, which substantiates Baldwin's (1963, p. 

380 382) concept. The mare-basalt masses constitute the superisostatic loads in the mascon basins (Baldwin, 1968; Wise and Yates, 

1970; Phillips and others, 1972; Bowin and others, 1975; Sjogren and 

Smith, 1976; Melosh, 1978; Solomon and Head, 1979, 1980). These 

superisostatic loads are generally smaller than they would be if no 

isostatic compensation had taken place since mare emplacement 

began (Solomon and Head, 1980, p. 136); therefore, the basalt must 

have subsided. This subsidence stretched the peripheral parts of the 

mare and the basin floor, creating the grabens, and compressed the 

center, creating the ridges (fig. 6.12; Baldwin, 1968; Melosh, 1978; 

Solomon and Head, 1979,1980).

FIGURE 6.8. Interpretation of mare ridges by Lucchitta (1976), based on models by A. R. 

Sanford. Thrust relations on surface originate by vertical displacements along steep faults in 

subsurface. Normal faulting occurs under tension in relatively raised block.

FIGURE 6.9. Arcuate rilles concentric with Mare Humorum and Humorum basin, between 

southeastern Mare Humorum (left) and westward extension of Mare Nubium (right). 

Craters include Campanus (C; 48 km, 28° S., 29° W.) and more highly degraded Mercator 

(M; 47 km). Orbiter 4 frame H-132.
6. STRUCTURE 113

CRATER-FLOOR FRACTURES

Fractures in the floor materials of craters are among the commonest lunar tectonic structures (figs. 6.13, 6.14). They are concentrated in and near maria and basins without maria (pi. 5). Fracturedfloor craters larger than about 150 km also occur in nonbasin settings, 

for example, Petavius and Humboldt (figs. 4.2, 9.5, 9.6). These two 

craters are gravitationally neutral; they have neither the negative 

anomalies of fresh craters nor the positive anomalies due to mare 

basalt (Dvorak and Phillips, 1978).

The fractured floors are higher relative to the crater rims and 

surrounding terrain than are the floors of typical impact craters 

(Pike, 1971). Many of the floors, however, occur in craters with other 

morphologies diagnostic of impact origin such as central peaks, as 

in Gassendi and Posidonius (figs. 5.17,6.13A), and the whole range of 

exterior impact phenomena including rays, as at Taruntius (fig. 

6.13B). Uplift of impact-crater floors is the evident explanation. The 

presence of fractured floors in Petavius and Humboldt (fig. 4.2) suggests that originally negative gravity anomalies resulting from mass 

loss during impact excavation were erased by the uplifts. Vitello on 

the Mare Humorum border (fig. 6.13C) possesses a freshly fractured 

floor, a mantle of dark material, and an enhancement in thermalinfrared wavelengths during eclipse (Shorthill and Saari, 1969) all 

features pointing to a caldera origin or, at least, to eruptions from the 

floor fractures (Saunders and Wilhelms, 1974). However, these features are also consistent with floor uplift (fig. 6.15), which opened 

fissures fresh enough to expose blocks that appear "warm" during 

eclipse, and with mantling by the dark material a common phenomenon along other mare borders as well (Titley, 1967).

The uplift interpretation also suggests an explanation for a wide 

range of other lunar landforms (Brennan, 1975; Schultz, 1976a). The 

shallow tilted floors and arcuate structures of 15 craters larger than 

20 km in diameter in the Smythii basin (fig. 6.13E; Wilhelms and ElBaz, 1977) suggest cylinderlike uplifts that broke free where the wall 

and floor meet (fig. 6.15). Especially great uplift probably formed the 

irregular, knobby elevated interior of Gaudibert, along the Mare 

Nectaris margin (fig. 6.13F; Brennan, 1975), which had been thought 

to be volcanic (Elston, 1972). Even small craters with such elevated 

interior structures as nested rims or rings of bulbous material (figs. 

6.13Z), G) may be modified by tectonic uplift rather than by volcanism, 

as is commonly believed (for example, Schultz, 1976b, p. 12-15). Twin 

impacts followed by uplift account for the nearly identical neighboring shallow-floored craters Sabine and Ritter along the Tranquillitatis border (fig. 3.14C), which superficially resemble terrestrial 

calderas more than they do impact craters (Morris and Wilhelms, 

1967; DeHon, 1971).

Although many fractures and unusual landforms have been created by uplifts of impact-crater floors, a few fractured-floor craters 

have a less well organized fracture pattern that suggests shrinkage of 

a cohesive material (fig. 6.13H). Most such craters lie near but outside 

basins, and the fractured material may be impact melt ejected from 

the basins (Moore and others, 1974; Wilhelms and others, 1979), 

unless it is otherwise-unknown terra volcanic material (Schultz, 

1976b, p. 68; Stuart-Alexander, 1978).

Igneous intrusions have been suggested as the cause of floor 

uplifts because the uplifts are concentrated near maria and are 

commonly overlain by dark mantles or small maria (Brennan, 1975; 

Schultz, 1976a). Intrusions seem reasonable, considering the affected 

craters' setting. All that is required, however, is isostatic leveling 

resulting from viscous relaxation of the substrate (Danes, 1965; Hall 

and others, 1981). The required plasticity is probably due to the 

weaker lithosphere beneath basins and large craters.

STRAIGHT RULES

General features

More puzzling than the arcuate grabens and floor fractures are 

rilles that are here called straight by contrast with the other, curved 

or complexly shaped fractures. Although long stretches are straight 

(pi. 5; figs. 6.16-6.21), most straight rilles curve or sharply inflect as a 

whole. In detailed shape of profile and in their echelon patterns, 

straight and arcuate rilles are similar (fig. 6.16). The straight rilles 

are also grabens (Quaide, 1965). Some rare single fault scarps not part 

of a rille are also considered here.

Imbrium sculpture

Faulting initiated by the Imbrium impact was the preferred 

explanation during the 1960's for the "Imbrium sculpture" system of 

radial grooves and ridges (Kuiper, 1959; Shoemaker and Hackman, 

1962; Hartmann, 1963, 1964b; Strom, 1964; Wilhelms, 1970b, p. 15; 

Wilhelms and McCauley, 1971). A coalescence of elliptical craters to 

form the grooves is evident on Apollo photographs (fig. 3.10B), 

however, confirming the secondary-ejecta origin proposed by Gilbert 

(1893) and Baldwin (1949, 1963). Nevertheless, some true faults are 

radial or subradial to Imbrium; those cutting the Apennine Bench, 

for example, probably formed in response to adjustments of Mare 

Imbrium and the Imbrium basin (fig. 6.10).

Vallis Alpes (the Alpine Valley, fig. 6.17) lacks the smoothly scalloped matching walls and raised lips characteristic of low-angle secondary chains. Some of its walls match but are jagged, steep, and

FIGURE 6.10. Apennine Bench, showing grabens concentric with Imbrium basin cutting basin, bench, and mare material (lower right). Radial grabens and vent for dark-mantling material (arrow) 

are also present (Imbrium center is outside upper left corner of photograph). Large rugged mass is Monies Archimedes, overlain and surrounded by planar Apennine Bench Formation (see chap. 

10) and by secondary craters of Archimedes, centered north of view. Mare at upper right is Palus Putredinis, west of Apollo 15 landing site. Pair of craters at left edge of bench are Feuillee (left; 9 

km) and Beer (9.5 km). Mosaic of Apollo 15 frames M-589 through M-594 (from right to left).
114 THE GEOLOGIC HISTORY OF THE MOON

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 '^ 

t rsrr-i-^rj. - ^- ' ,- =»-*--, =- frf -~-^ . ' .- f" 1 ^, rr*-; r: '^Xi-t-^-jf jf*fc-\«'-~:^'-' * % * "- ^~» - " ' =" »«:,, y »y. ,

 - . , fl ? ' ' *' 1 fi ' i:T> ' ' * *' * ~ * -'y?O ^- '" N;^V..'.<r -K. 'wi: &&-, . >w-v * ~ - -; i"

FIGURE 6.11. Arcuate rilles in border zone of Mare Tranquillitatis (5° N., 21 ° E.) south of crater Arago, about 90 km east of mare edge. Intricate structure shown here is probably also common in

other, less well photographed rille systems. Orbiter 2 frame M 43.
6. STRUCTURE 115

composed of linear segments. The valley is probably a graben formed 

shortly after the Imbrium impact, as was once proposed for all sculpture. Its radiality is consistent with an origin by isostatic doming of 

the subbasin mantle (see chap. 4).

Such radial or subradial faults as the Straight Wall (Rupes Recta, 

fig. 1.8), one leg of Rima Hyginus (figs. 5.1QE, 6.16A), the Cauchy set 

(fig. 6.18), and several parallel grabens in southwestern Mare Fecunditatis (fig. 6.19) are more puzzling because they cut mare materials. 

They cannot be direct products of the Imbrium impact. If they are 

related to Imbrium, they may have formed by some sort of rejuvenation of radial fractures (Wilhelms, 1972a; Holt, 1974; Mason and 

others, 1976).

Other systems

The geometric relation of other straight rilles to basins is less 

apparent. Many grabens in the central and east-central equatorial 

zone of the nearside are concentrated along trends between 20° and 

30° north of west (pi. 5). Most of these grabens belong to a set that 

extends 1,600 km east-southeastward from Rima Bode II (fig. 5.1QH), 

through the eastern leg of Rima Hyginus and Rimae Ariadaeus (fig. 

6.16), to the Imbrium-radial rilles in Mare Fecunditatis (figs. 6.19, 

9.7). The only basin to which all these grabens are geometrically 

related is Procellarum, to whose center (23° N., 15° W.; Whitaker, 1981) 

they are radial or subradial. Their trend is paralleled in other areas 

by the Cauchy set of structures (also subradial to Procellarum; fig. 

6.18), the short Miiller crater chain (fig. 3.16A), and the 300-km-long 

Abulfeda crater chain (fig. 3.16B). The origin of these crater chains, 

which are radial to no known basin or large crater, is unclear (see 

chap. 3). Some grabens near Sinus Medii have a complementary trend 

of 20°-30° north of east (pi. 5; fig. 10.28).

At the east end of Sinus Medii is the complex Triesnecker system 

of grabens (pi. 5; fig. 6.20). This system as a whole is oriented northsouth but includes many other trends, which mutual transection 

relations show to be contemporaneous (Wilhelms, 1968). East-west 

extension is the apparent cause.

The Moon's most complex system of (exposed) grabens occupies 

the southwest shore of Oceanus Procellarum from about lat 5° N. to 

lat 22° S. (fig. 6.21). The prevailing trends are nearly concentric with 

the Procellarum shoreline or radial to the general region of both the 

Imbrium- and Procellarum-basin centers. Other straight rilles lie 

along the northwest Procellarum shore, which also contains the 

largest lunar concentration of fractured-floor craters (pi. 5; fig. 6.14).

INTERPRETATION 

LITHOSPHERIC THICKNESS

Many, if not most, of the structures described in this chapter are 

ascribable to thinning of the elastic lithosphere beneath basins. The 

subsidence of basin-filling mare basalt is aided by a thin lithosphere 

and inhibited by a thick lithosphere (Melosh, 1978; Solomon and 

Head, 1980). Solomon and Head (1980) suggested that a lithospheric 

thickness of about 75 km has prevented graben formation in the 

Crisium, Nectaris, and Smythii basins. The crust in the region of the 

seismic stations on the southwest nearside is 45 to 60 km thick (chap. 

1). Because arcuate grabens have formed in this seismically explored 

region, such crustal thicknesses apparently facilitate mare subsidence and graben formation. Because the estimates of lithospheric 

and crustal thickness are similar in magnitude, the feldspathic crust 

and the elastic lithosphere were probably equivalent at the time of 

graben formation. The mantle constituted the more plastic 

asthenosphere.

Chapter 4 suggests that the Procellarum-basin impact exerted a 

major control over crustal (lithospheric) thickness (fig. 6.22). Consequently, arcuate grabens are restricted almost entirely to the interior of that basin (pi. 5). Even the old shallow basin Tranquillitatis 

and the thin basalt of Mare Tranquillitatis, which lie inside Procellarum, are cut by marginal grabens (fig. 6.11). Lithospheric thickness related to Procellarum also may be reflected in the gravity 

structure. Outside Procellarum, even such relatively thin maria as 

Nectaris, Orientale, and Smythii preserve mascons. Despite their 

thinness, these maria lie at low elevations (more than 3.5 km below 

the average lunar sphere, 1,738 km in radius; fig. 6.23; table 6.1).

Thus, the thick lithosphere apparently hindered isostatic uplift 

of basins, extrusion of mare basalt, sinking of mascons, and formation 

of grabens. Each condition was opposite inside Procellarum. Oceanus 

Procellarum and Mare Tranquillitatis, though areally large, have 

only small, local gravity highs (one part of southern Procellarum and 

the mare-ridge feature Lamont in Tranquillitatis; figs .11.1,11.9; Scott, 

1974). The superisostatic loads per unit area in Imbrium and 

Serenitatis are smaller than would be expected from their thick 

mare-basalt sections (table 6.1; Solomon and Head, 1980). Apparently, 

the thin lithosphere inside Procellarum abetted early isostatic uplift 

of such basins as Tranquillitatis and permitted later isostatic sinking 

in response to loading even by thin basalt. The concentration of maria 

in Procellarum and in the superposed basins is another result of the 

thinner lithosphere. Although mascons are not known on the farside, 

the depth of the giant South Pole-Aitken basin (5-7 km; StuartAlexander, 1978), the paucity of farside maria except in the basins 

superposed on South Pole-Aitken (pi. 4), and the absence of grabens 

all suggest farside lithospheric thicknesses at least as great as those 

on the non-Procellarum nearside.

The factors that localize the straight rilles are less clear but may 

be further effects of the Procellarum basin. The complex system along 

the west Procellarum shore comprises the grabens most likely to be 

related to this giant basin. Widespread extension of the Procellarum 

margin is suggested by the trends and extent of the fracture system, 

which are consistent with a broad regional uplift of the mantle under 

the basin.

More speculatively, the many straight rilles on the central and 

east-central nearside may have a similar origin. Mantle uplift 

beneath the Procellarum basin may explain the 1,600-km-long system 

and such "Imbrium-radial" faults as Rupes Recta, the grabens cutting the Fra Mauro peninsula (fig. 5.8), and the Cauchy set. Continuation of mild uplift into the time of mare-basalt extrusion would 

account for the otherwise-puzzling transection of the maria by these 

faults. The Triesnecker system, whose overall orientation diverges 

from others in its vicinity, may have been localized by Sinus Medii, 

whose southeast boundary is a Procellarum ring (pi. 4; fig. 5.26). 

Sinus Medii may have been isostatically uplifted because of viscous 

relaxation of the weak lithosphere, in the manner of a fractured crater 

floor. Tides, relaxation of a tidal bulge (Melosh, 1977), or some other 

global effect would seem to be attractive alternative causes of the 

deformation of this near-Earth zone. To my knowledge, however, the 

observed distribution of lunar faults is inconsistent with any of the 

global mechanisms that have been thus far proposed. The Triesnecker system and the systems of long straight grabens, therefore, 

provide additional examples of basin-related alternatives to global 

origins for lunar structures.

A progressive thickening over time of the lunar lithosphere in all 

regions (Howard, 1970) is evidenced by increasing resistance to deformation (Solomon and Head, 1980). The arcuate and straight grabens 

cut terra units and old mare units (chap. 11), whereas the ridges 

deform both old and young mare units (Lucchitta and Watkins, 1978). 

Thus, the graben opening was the first deformation to cease. Mare 

extrusion ceased somewhat later, when the elastic lithosphere 

exceeded 100 km in thickness (Solomon and Head, 1979). Presumably, 

most previous conduits for ascent of mare magmas were shut off. 

Wrinkling of the central mare surfaces and crater-floor uplift were the 

longest enduring modifications. All these changes were presumably 

the result of general global cooling. Decreasing plasticity of part of the 

mantle added thickness to the lithosphere, which once consisted only 

of the terra crust. Although the mascons indicate that the superisostatic loads still exist, continued subsidence is unlikely because 

endogenic moonquakes do not all correlate with maria (Nakamura 

and others, 1979; Solomon and Head, 1979). If any asthenosphere 

remains today, it is too deep to affect surface deformation. The only 

recent deformation that has been suggested since analysis of the 

Apollo results is minor thrust faulting in the terrae (Binder, 1982).

Younger mare basalt

- 'f''-"'-vOlder mare basalt

FIGURE 6.12. Diagrammatic cross section showing stresses responsible for graben opening 

and ridge formation during subsidence of mare basalt in a double-ringed basin. Short vertical 

lines cutting older mare unit and underlying basin material denote grabens.
116 THE GEOLOGIC HISTORY OF THE MOON

FIGURE 6.13. Raised fractured floors of craters.

A. Gassendi (110 km, 17.5° S., 40° W.), on north border of Mare Humorum. Floor is 

almost as high as mare, but peaks indicate impact origin. View southward. Apollo 16 

frame H-19295.

B. Taruntius (56 km, 5.5° N., 46.5° E.), on northwest border of Mare Fecunditatis, a 

typical fresh impact crater except for raised floor. Or biter 1 frame M 31.

C. Vitello (42 km, 30.5° S., 37.5° W.), on south margin of Mare Humorum, with darkmantled rim and very fresh fractures on uplifted floor. Orbiter 4 frame H 136.

D. Pitatus (97 km, 30°S., 13.5°W.), on south border of Mare Nubium. Concentric rilles 

result from uplift of crater floor and later mare fill. Small bull's-eye crater at left may have 

similar cause. Chain of craters at right are secondary-impact craters of Imbrium basin. 

Line of white dots in lower left is photoprocessing artifact. Orbiter 4 frame H 119.

E. Group of double-ringed craters, 30 to 35 km in diameter, in Mare Smythii. Cylinderlike 

uplift of floor is indicated (fig. 6.15). Orbiter 4 frame H 17.
6. STRUCTURE 117

H

G.

F. Gaudibert (34 km, 11 ° S., 38° E.) and two smaller craters on margin of Mare Nectaris. 

Bulbous landforms are possible evidence for viscous extrusive volcanism (Elston, 1972) 

but were probably created by extensive uplift (Brennan, 1975). Dark-mantling material 

surrounds teardrop crater along fracture in lower right crater. Orbiter 4 frame H 72. 

Liouville DA (11 km, 46.5°, 52° W.), on border of Sinus Roris, possibly containing 

volcanic extrusions (Scott and Eggleton, 1973) but probably another uplift phenomenon. Sinuous rille Rima Sharp I is at right. Orbiter 4 frame H-163.

H. Tamm (38 km, 4.5° S., 146.5° E., foreground) and Van den Bos (32 km), 225 km 

southeast of Mendeleev-basin rim, filled by fissured, viscous-appearing material possibly 

emplaced as impact melt of Mendeleev. View southward. Apollo 10 frame H-4966.

TABLE 6.1. Physical properties of major nearside basins containing large maria

[D, diameter (km); basin D's from table 4.1.

Average elevation: Elevation of central mare inside obvious sloping margins, in km below 1,738-km datum (S.S.C. Wu, written

commun., 1983).

Mascon mass: excess mass in units of 1020 g (compiled by Solomon and Head, 1980, table 2). 

Mascon load: superisostatic load in center of mare, in units of 10 8 dyne/cm 2 (calculated by Solomon and Head, 1980, table 2,

using different mare diameters than those listed here; they took Smythii as 340 km, Grimaldi as 100 km, and Nectaris as 400 km

in diameter)]

Basin

Imbrium (with 

Frigoris) . 

Imbrium (without 

Frigoris).

Nor"t~ar*ic__ _

Tranquill itatis 

Humbol dti anum 

£rima1 Hi __________

Basin D

1,500 

1,160

1,060 

930 

860 

840 

820 

775 

740 

690 

690 

600 (each) 

/nn

Mare D

1,450 

960

420 

280 

300 

200 

350 

600 

600 

600 

600 

120 

i dn

Ratio

.83

.40 

.30 

.35 

.24 

.43 

.77 

.81 

.87 

.87 

.20 

TI;

Outer D 

of graben 

zone

1,160

<420 

500 

320 

<200 

750 

710 

850 

600 

620 

<120 

ifin

Ratio 

to basin

1.00

<.40 

.54 

.37 

<.24 

.91 

.92 

1.15 

.87 

.90 

<.20

R/l

D of 1-km 

isopach

350

210 

150(?) 

150 

0 

125 

150 

300(?) 

150

o(?)

0(?) 

r\f->\

Ratio 

to basin

.30

.20 

.16(7) 

.17 

0 

.15 

.19 

.41 

.22

o(?)

0(?) 

r\n\

Ratio 

to mare

.35

.50 

.54 

.50 

0 

.36 

.25 

.50 

.25

o(?) 

o(?)

r\lf\

Average 

elevation

2.5-3.5

4.0-4.5 

4.0 

>3.5 

4.0-4.5

7

1.5->2.5 

3.0-4.0 

2.0-3.0 

>3.0 

? 

t n * K

Mascon 

mass

23.0

9.4 

3.3 

9.0 

5.0 

5.2 

0 

14.0

o(?)

0

o(?)

7fi

Mascon 

load

.9

1.3 

.9 

1.3 

1.4 

1.4 

0 

1.3 

0(?) 

0

o(?)

1 /I
118 THE GEOLOGIC HISTORY OF THE MOON

!fei»m--V""

*

FIGURE 6.14. Numerous fractured-floor craters along west border of Oceanus Procellarum 

between lat 28° N. and 57° N. R, crater Rontgen (126 km, 33° N., 91 ° W.), superposed on 

Lorentz basin. Linear structure at bottom is radial to Imbrium basin. Orbiter 4 frame 

H-183.

FIGURE 6.15. Crater floor uplifted (arrows) from original position (dashes) lower than

surrounding terrain.

FIGURE 6.16. Straight rilles Hyginus (A) and Ariadaeus (fi). Rima Ariadaeus continues 

westward in B as rille indicated by arrow in A. Largest crater in B is Silberschlag (13 km, 6° 

N., 6.5° E.). Surface offsets and echelon patterns are obvious. Apollo 10 frames H-4648 

(A) and H-4645 (B).
6. STRUCTURE 119

FIGURE 6.17. Vallis Alpes (Alpine Valley), 200 km long and oriented radially to Imbrium 

basin. Knobby material is Alpes Formation (see chaps. 4, 10). 

A.Regional setting, centered 47° N., 0°. Orbiter 4 frame H-115. 

B. Detail. View southwestward. Orbiter 5 frame H 102.

FIGURE 6.18. Mare Tranquillitatis, faulted by Cauchy rille (Rima Cauchy, right) and 

complex scarp (Rupes Cauchy, left), oriented radially to Imbrium basin (beyond top of 

photograph). Fresh crater between the two structures is Cauchy (12 km, 10° N., 39° E.). 

Part of crater Taruntius (compare fig. 6.13B) is visible in lower right corner. Apollo 11 frame 

H-6231.

FIGURE 6.19. Grabens cutting Mare Fecunditatis, crater Goclenius (longest dimension, 72 

km; 10° S., 45° E.), and mare fill of Goclenius. View southwestward. Apollo 8 frame 

H-2225.
120 THE GEOLOGIC HISTORY OF THE MOON

FIGURE 6.20. Triesnecker system of grabens east of crater Triesnecker (left center; 26 km, 

4° N., 3.5° E.). Part of northwest branch of Rima Hyginus is at right, oriented northwest 

radially to Imbrium basin. Apollo 10 frame H-4816.

FIGURE 6.21. Grabens along west margin of Oceanus Procellarum south of crater Hevelius 

(106 km, 2° N., 67° W., cut by top edge of photograph). Circular mare below Hevelius fills 

Grimaldi basin (compare fig. 4.4f/). Orbiter 4 frame M-161.

FIGURE 6.22. Diagrammatic cross section of half of Procellarum basin, showing interpreted 

crustal thinning from 75 km outside basin to 25 km in basin center, caused by excavation and 

mantle uplift; dashed line denotes preimpact position of mantle-crust interface (lunar Moho). 

Ring radii (Whitaker, 1981) are shown above lunar surface; basin rim is at 1,600 km. Dotted 

lines denote additional excavation and crustal thinning by later basins; basin in Procellarum 

center penetrates crust to mantle, but same-size basin in intermediate and outer Procellarum 

troughs bottoms in crust. Vertical and horizontal dimensions and curvature to true scale. 

Surface relief and relative relief of concentric rings and shelves are barely detectable.

GAGARIN

180°-

FECUNDITATIS

N. TRANQUILLITATIS

. NECTARIS (TONGUE) 

SERENITATISN%A \ \

PUTREDINIS

0 » EARTH

NUBIUM 

COGNITUM

HERTZSPRUNG

90° W

FIGURE 6.23. Elevations of lunar maria compared with those of the terrae, most of which 

define a sphere 1 , 738 km in radius about the center of gravity (Sjogren and Wollenhaupt, 

1 976). Derived from laser altimetry and photogrammetry obtained during Apollo 1 5 and 1 6 

missions (Kaula and others, 1973, 1974; Sjogren and Wollenhaupt, 1976). Smythii, 

Crisium, Serenitatis, Imbrium (S. Imbrium), and Grimaldi vary in elevation, although all 

have large mascons. Elevations of maria in Aitken, Tsiolkovskiy, and Van de Graaff 

probably reflect elevations of crater floors; Aitken is on rim of South Pole- Aitken basin, and 

van de Graaff well inside. Chaplygin, Gagarin, Heaviside, Hertzsprung, Keeler, Mendeleev, and Pasteur are filled by terra plains, not maria. Polar plot relative to 1,738-kmradius sphere about lunar center of gravity.
pouad

uassuep UOUBLUJO-J 

OIUBO|OA uj PUB 

BJ-J ojnB|/\j UOUBUJJO-J lunuqiui) (uiSBq 

qoodg 

AjJBg ueuqiuj 

S9U9S

31AJI1 UNO 

UNO

OIHdVdOllVHIS 

31MI1

oiHdvuoiiv«is->iood snisin 

S]9V
FIGURE 7.1 (OVERLEAF). Diagrammatic lunar stratigraphic column illustrating relations among three classes of stratigraphic units (youngest at top). Three named basin materials are the denning 

rock-stratigraphic units of the oldest time-stratigraphic units. Individual lunar rock-stratigraphic units (mare or crater materials; table 7.2) formed within very brief time intervals, but groups of 

rock-stratigraphic units may cross time-stratigraphic boundaries.
7. RELATIVE AGES

CONTENTS

Page

Introduction ...................................................... 123

Stratigraphic nomenclature ........................................ 123

Superpositions ................................................... 125

Mare-crater relations .......................................... 125

Crater-crater relations .......................................... 127

Basin-crater relations .......................................... 127

Mapping conventions .......................................... 127

Crater dating .................................................... 129

General principles ............................................. 129

Size-frequency relations ........................................ 129

Morphology of large craters .................................... 129

Morphology of small craters, by Newell J. Trask .................. 131

DL method .................................................... 133

Summary ........................................................ 133

INTRODUCTION

The goals of both terrestrial and lunar stratigraphy are to integrate geologic units into a stratigraphic column applicable over the 

whole planet and to calibrate this column with absolute ages. The 

first step in reconstructing the relative stratigraphy is the identification of the material units of craters, basins, and maria (see chaps. 

2-5); stratigraphic relations to neighboring units are generally recognized in the course of recognizing and defining a unit. The next 

step is to relate as many units as possible to areally extensive units. 

Units contacting one of these key stratigraphic horizons are dated 

directly by geometric relations, and others are dated means of the 

statistics and morphologies of superposed craters. In present practice, laterally extensive deposits of the Nectaris, Imbrium, and Orientale basins divide the time-stratigraphic column into four major 

sequences, from oldest to youngest: pre-Nectarian, Nectarian, 

Lower Imbrian, and Upper Imbrian through Copernican (fig. 7.1;

TABLE 7.1. Time-stratigraphic nomenclature in use on lunar geologic maps of the U.S.

Geological Survey

[Diagrammatic; no absolute-age spans implied. Upper and Lower series of the Imbrian 

System are informal and were not recognized on all geologic maps from 1970 to 

1979]

1959-63 

(Shoemaker and 

Hackman, 1962)

Copernican

Eratosthenian

Procellarian

Imbrian

1963-70 

(Shoemaker, 

1964; McCauley 

1967b)

Copernican

Eratosthenian

Archimedian

Apenninian

1970-75 

(Wilhelms, 

1970b)

Copernican

Eratosthenian

1975-79 

(Stuart-Alexander 

and Wilhelms, 

1975)

Copernican

Eratosthenian

Nectarian

Pre-Nectarian

This volume

Copernican

Eratosthenian

Upper 

Imbrian

Lower 

Imbrian

Nectarian

Pre-Nectarian

table 7.1). The first three of these sequences, which are older than 

the visible mare materials, are also dominated internally by the 

deposits of basins. The fourth (youngest) sequence consists of mare 

and crater materials. This chapter explains the general methods of 

stratigraphic analysis that are employed in the next six chapters 

and on plates 6 through 11 to subdivide and calibrate the lunar 

stratigraphic column.

STRATIGRAPHIC 

NOMENCLATURE

The task of reconstructing the lunar stratigraphic column has 

benefited from application of the threefold code of stratigraphic 

nomenclature developed in North America for terrestrial geology 

(American Commission on Stratigraphic Nomenclature, 1970). This 

code is just as applicable to the Moon as to the Earth (Mutch, 1970, 

chap. 5; Wilhelms, 1970b).

Complete understanding of the stratigraphic nomenclature 

used in this volume requires some familiarity with the three major 

types of stratigraphic units recognized by the code (fig. 7.1). Rockstratigraphic (rock) units are the observed, physical units that can 

be identified and mapped objectively. Ideally, they are defined by 

physical properties intrinsic to their emplacement process and are 

described additionally by other reproducibly observable properties. 

The basic rock-stratigraphic unit is the formation; formations can be 

combined into groups and divided into members. Most of the foregoing discussions in this volume have concerned descriptions and 

interpretations of rock-stratigraphic units.

Time-stratigraphic (time-rock) units include all the rock-stratigraphic units emplaced on a planet within a given time interval. The 

basic time-stratigraphic unit is the system, which can be divided into 

series (and finer subdivisions in terrestrial geology). Time-stratigraphic units are defined on the basis of specific rock-stratigraphic 

units. Time-stratigraphic units do not overlap, and the upper boundary of one unit is the lower boundary of the next.

The time intervals corresponding to the time-stratigraphic units 

are the third type of unit. They are defined by the corresponding timestratigraphic units and are not physical units. Periods correspond to 

systems, and epochs to series.

Lunar time-stratigraphic nomenclature has changed somewhat 

since the concept of lunar stratigraphy was introduced by Shoemaker 

and Hackman (1962). The first two maps published by the U.S. Geological Survey at a scale of 1:1,000,000 (Hackman, 1962; Marshall,

123
124 THE GEOLOGIC HISTORY OF THE MOON

1963) used the original scheme, and the remaining 42 maps of this 

series used modified versions (fig. 7.2A; table 7.1). Subsequent mapping, which covers the Moon at a 1:5,000,000 scale (fig. 7.2B), has 

incorporated additional refinements (table 7.1).

The purpose of distinguishing rock-stratigraphic and time-stratigraphic units is illustrated by changes in concept and definition of 

the Imbrian System (Imbriarc refers to the system, and Imbriwm to 

the basin). Initially, Shoemaker and Hackman (1962, p. 293-294) 

defined the Imbrian System as equivalent to the "immense sheet" of 

material around Mare Imbrium now given such rock-stratigraphic 

names as the Fra Mauro Formation (table 4.3). They named the mare 

material the "Procellarian System", after the largest expanse of mare 

material, Oceanus Procellarum. The concept of the Imbrian Period 

was quickly extended to cover a longer timespan than that necessary 

to emplace the sheet (Shoemaker and Hackman, 1962, p. 298-299). 

The crater Archimedes, for example, was recognized as younger than 

the Imbrium basin but older than the Procellarian System (fig. 1.6). 

After Shoemaker and Hackman's (1962) report was prepared but 

before it was formally published, the Imbrian System was divided 

into the Apenninian Series, equated with the "regional material of

the Imbrian System" (the sheet), and the younger Archimedian 

Series, equated with the deposits of Archimedes and other craters of 

similar stratigraphic position (Shoemaker and others, 1962b).

One problem with these classifications was that they equated 

time-stratigraphic and rock-stratigraphic units. The definition of the 

Procellarian System by Shoemaker and Hackman (1962) was founded 

on the belief, based on telescopic crater counts and consistent stratigraphic relations, that the mare materials were formed within a 

short time interval (Shoemaker and Hackman, 1962, p. 299; Shoemaker and others, 1962a; Baldwin, 1963, p. 309). However, continued 

crater counts (Dodd and others, 1963; Hartmann, 1967) and stratigraphic observations (Carr, 1966a, b; McCauley, 1967a, b; Wilhelms, 

1968,1970b) showed that they vary substantially in age. It is now clear 

that sections of mare basalt cross time-stratigraphic boundaries; 

mare and crater materials interfinger well into the Eratosthenian 

System (fig. 7.1; table 7.2; McCauley, 1967b, p. 436; Wilhelms, 1970b, p. 

F36; 1980). This extraterrestrial example vividly illustrates the 

rationale behind the threefold stratigraphic code, which sprang from 

discoveries that many terrestrial lithologic units are time-transgressive. The rapid alternation of volcanic and impact deposits could

-_ . . ,.--

-.'XL-. - ' -' **/ £ ' 'J&. ~^3^*\'--*fc* '^^

s 

A

B

FIGURE 7.2. Index maps of lunar geologic mapping.

A. Nearside, showing coverage by 1:1,000,000-scale series. Number above quadrangle name refers to 

base chart (LAC series) by the U.S. Air Forces Aeronautical Chart and Information Center 

(ACIC). I-number refers to published map in U.S. Geological Survey Miscellaneous Investigations 

Series.

B. Nearside and farside coverage by geologic maps at 1:5,000,000 scale: 1 703, Wilhelms and 

McCauley (1971; same coverage as in fig. 7.2A); 1-948, Wilhelms and El-Baz (1977); 1-1034, 

Scott and others (1977); 1-1047, Stuart-Alexander (1978); 1-1062, Lucchitta (1978); I-l 162, 

Wilhelms and others (1979).
7. RELATIVE AGES 125

TABLE 7.2. Selected stratigraphic units in southern Mare Imbrium and part of northern Oceanus Procellarum

[Stated superpositional relations refer to units in actual contact. After Wilhelms (1980)]

Feature Materials in area

Oldest

superposed

unit

Youngest

subjacent

unit

AristarchusKepler 

Pytheas----- 

Copernicus--

Diophantus--

Mare unit 6- 

Mare unit 5- 

Delisle 

Euler 

Mare unit 4 

Timocharis - 

EratosthenesLambert------

Mare unit 3 

Krieger 

Mare unit 2 

Mare unit 1 

Aristarchus

plateau. 

Prinz---------

Imbrium basinVery fresh secondary craters None------

do Aristarchus 

Fresh ejecta and secondary craters Kepler - 

Large clusters of secondary craters Pytheas -

along major rays.

Fresh ejecta, small fresh secondary Copernicuscraters; no rays.

Flow-textured lavas (DL=140-195 m) Diophantus 

Flow-textured^ lavas (DL=170-215 m) Mare unit 6- 

Fresh ejecta 'and secondary craters; no Mare unit 5-

rays.

Fresh ejecta and secondary craters Mare unit 5- 

Flow-textured lavas (D, =210-245 m) Euler - 

Fresh ejecta and secondary craters Mare unit 4- 

Extensive, subdued secondary craters Timocharis 

Slightly subdued ejecta and secondary Mare unit 4-

craters.

Smooth lavas (DL^255 m) Mare unit 4- 

Asymmetric ejecta and secondary craters Mare unit 3- 

Smooth lavas (DL=240-385 m) Krieger 

Smooth lavas (DL=270-385 m) Mare unit 2- 

Dark-mantling material (DMM) Mare unit 1-

Subdued ejecta and secondary craters DMM ------

Isolated islands- Prinz------

Kepler. 

Pytheas. 

Copernicus. 

Diophantus.

 Mare unit 6.

Mare unit 5. 

Delisle. 

Mare unit 4.

Mare unit 4. 

Timocharis. 

Eratosthenes. 

Mare unit 3. 

Mare unit 2.

Krieger. 

Mare unit 2. 

Mare unit 1. 

DMM. 

Prinz.

Imbrium basin. 

None of above.

not have been expressed by Shoemaker and Hackman's nomenclature 

without misuse of the time-stratigraphic concept (McCauley, 1967b; 

Mutch, 1970; Wilhelms, 1970b).

One possible modification would have been to redefine the Procellarian System to include crater deposits that interfinger with mare 

materials. The course that was taken, at a meeting of lunar stratigraphers held in Flagstaff, Ariz., in November 1963, was to drop the 

name "Procellarian System" and to designate mare materials by the 

rock-stratigraphic name "Procellarum Group" (Hackman, 1964; 

McCauley, 1967b; Wilhelms, 1970b). The upper boundary of the timestratigraphic Imbrian System was to be defined by some part of the 

rock-stratigraphic Procellarum Group. Other rock-stratigraphic 

names were to be devised for Eratosthenian and Copernican mare 

materials when they were discovered. The contemporaneous crater 

materials were mapped as material units assigned to the appropriate 

system or series. Later, even the name "Procellarum Group" was 

dropped in favor of the informal rock-stratigraphic or lunar-material 

name "mare material," because "Procellarum Group" retained a timestratigraphic connotation (Wilhelms, 197Gb). 7- 1

The 1963 and 1970 stratigraphic schemes were used throughout 

most of the l:l,000,000-scale mapping program (table 7.1). The 

Imbrian System included all materials from the base of the Fra 

Mauro Formation stratigraphically up to unspecified mare units 

between the premare crater Archimedes and the postmare crater 

Eratosthenes (figs. 1.6,2.1; Wilhelms, 1970b, p. 23). Thus, the Imbrian 

System included all basin, crater, mare, and other materials 

emplaced from the time the Imbrium basin formed until the time its 

flooding was mostly complete. The amount of fill excluded from the 

Imbrian System was not specified in this working definition.

Chapters 10 and 11 refine the definition of the Imbrian System 

and define two series, Lower and Upper Imbrian, separated by the 

Orientale-basin materials. An earlier major change in the stratigraphic scheme divided the pre-Imbrian of Shoemaker and 

Hackman (1962) into the pre-Nectarian and Nectarian Systems 

(table 7.1; Stuart-Alexander and Wilhelms, 1975). Although the definition of the Imbrian-Eratosthenian and Eratosthenian-Copernican 

boundaries remains imprecise, no further serious obstacles have 

appeared in the lunar stratigraphic nomenclatural scheme.

7-1The equation of the Imbrian System or the Apenninian Series with rocks interpreted as Imbrium ejecta 

would have been less of a practical problem because the ejecta has not proved to be time-transgressive, as are the 

mare materials. However, assumptions about the time significance of rock units so commonly prove to be erroneous 

in terrestrial and planetary geology that the threefold code is best adhered to from the beginning of the study of any 

new planet (Mutch, 1970).

Several other formal rock-stratigraphic units were introduced during the mid-1960's as antidotes to the 

potential confusion of time and rock units and as means of clarifying local stratigraphic relations (Wilhelms, 1970b). 

Some of these names are still used, but present practice emphasizes Moon-wide similarities of agiven class of unit by 

use of such informal names as "mare materials" and "crater materials." Such units may be assigned to timestratigraphic units or subdivided by physical properties as required. Formal names are now normally given only (1) 

to units that are difficult to describe simply and that form important regional deposits, such as the Hevelius or Alpes 

Formations; or (2) to separate patches of similar-appearing deposits that are being distinguished for some reason, 

such as the Apennine Bench and Cayley Formations. Whitford-Stark and Head (1980), however, resuggested the 

practice of formally naming mare units.

SUPERPOSITIONS

Mare-crater relations

Superpositions are recognized on photographs by topographic 

contrasts that are not explicable by facies changes within one unit. 

Rugged landforms may be muted by smoother blankets, and smooth 

topography may be crosscut by rugged texture. Many crater deposits 

and mare flow units illustrate such relations. Some superpositional 

relations are obvious (fig. 7.3); others are established by subtle flooding of secondary craters by thin mare flows (fig. 7.4). The difference 

between a slightly flooded secondary crater and an unmodified secondary superposed on other lavas may be detectable only on the best 

photographs.

Secondary craters greatly extend the area over which the primary crater can be relatively dated. Secondaries along conspicuous 

rays serve best for long-distance correlations because the rays visible 

on favorably illuminated and processed photographs point out the 

source (fig. 7.4). For example, the clustered craters at the Ranger 7 

impact site were tentatively identified as secondary to both Copernicus and Tycho despite their 600- and 1,000-km distances, respectively, from those primary craters (U.S. Geological Survey, 1971). 

Tycho, Copernicus, Kepler, and other young craters have been identified as the sources of projectiles that reached every potential Apollo 

landing site mapped in the equatorial belt, hundreds or even thousands of kilometers from the source craters (Carr and Titley, 1969; 

Carr, 1970; Grolier, 1970a, b; Cummings, 1971; Pohn, 1971; Rowan, 

1971a). In principle, secondary craters can also be traced to nonrayed 

craters by analysis of the flight directions of the ejecta projectiles: 

Large secondary craters generally occur in a cluster uprange toward 

the primary, and herringbone patterns, radial grooves and ridges, 

and smooth deposits occur downrange (figs. 3.4, 3.6-3.8, 7.4). Stratigraphic relations both at the primary crater and at the outlying 

clusters can thus be used to date units relative to this dispersed 

stratigraphic datum.

&#/&*£&#.

FIGURE 7.3. Crater Plato (101 km, 52° N., 9° W.X between Mare Imbrium (bottom) and 

Mare Frigoris (top). Mare materials fill Plato and truncate its south rim deposits (arrows), 

thus are younger than Plato. Fully developed rim deposits and secondary craters of Plato are 

preserved where superposed on terra (part of Monies Alpes); thus, Plato is younger than 

Monies Alpes. Secondary craters of Plato are superposed on secondary craters of Indum 

crater (arrow). Or biter 4 frame H 127.
-

SM^^fe!-Slitef^': : :r-^^ ^f'frn'- ,v G

Ifc, H-A^iis* iS- ' ' ^.fi, - *"'* - - '** >4.

FIGURE 7.4. Part of Mare Imbrium, including crater Euler (E; 28 km, 23° N., 29° W.) 

and inundated massifs of Imbrium basin (bottom, fig. 7.4A). White arrows indicate same 

scene in both frames; left-hand arrow points to crater Aristarchus, 650 km to west, and lower 

arrow to Copernicus, centered 450 km to south-southeast. Alignment of secondary-crater 

groups indicates these primary craters as sources of the excavating projectiles. 

A. Ejecta and secondary craters of Euler, subtly flooded by mare material (black-and-white

arrowheads). Apollo 17 frame M-2291. 

B. Aristarchus secondaries superposed on Copernicus secondaries (black-and-white

arrowX indicating that Aristarchus is the younger crater. Apollo 17 frame M-3093.
7. RELATIVE AGES 127

The ages of 20 mare and crater units in southern Mare Imbrium 

and northern Oceanus Procellarum are bracketed by the telltale 

marks of secondary craters and mare embayments (fig. 7.4; table 7.2; 

Wilhelms, 1980). This region is favored both by photographs of the 

necessary quality and by relatively extensive mare flows that can be 

used to correlate units over hundreds of kilometers (Schaber, 1973; 

Wilhelms, 1980). More commonly, mare materials mappable as single 

flow units extend only a few tens of kilometers (Schaber and others, 

1976). Small units are generally dated by methods based on small 

superposed craters, as described below in the section entitled "Crater 

Dating."

Crater-crater relations

Superpositions among crater materials establish local stratigraphic sequences in all parts of the Moon. How near craters must be 

to allow relative dating depends on the freshness of their deposits and 

secondary craters. Rayed and other fresh craters can be dated relative 

to other craters as they can to mare units, that is, if the secondary 

craters happen to contact the other crater. Several pairs of rayed and 

nonrayed craters listed in table 7.2 are relatively dated by the chance 

impacts of younger secondaries upon older. Degraded craters lacking 

obvious secondaries can be relatively dated if one rim happens to cut 

out part of another, the so-called "cookie cutter" obliteration.

Figure 3.26 illustrates the more common case between these 

extremes, in which the radial ejecta of one crater (Werner) is superposed on an older crater lacking ejecta texture (Aliacensis). The 

Aliacensis rim is swamped within about one Werner radius from the 

Werner rim and is visibly marked to about one additional radius. If 

the radial textures of Werner were not visible, the degradation of the 

Aliacensis rim could nevertheless be detected one radius, but not two 

radii, from the Werner rim. These relations provide two guidelines 

used during preparation of the paleogeologic maps in this volume 

(pis. 6 11): (1) Subdual of a crater's rim within one radius of another 

rim indicates that the first crater is the older, and (2) ages are 

determinable at greater distances where textural evidence is available. Each case, however, is individual.

Basin-crater relations

Similar observations help date the extensive deposits of ringed 

basins relative to one another and to other units. Massif-lined rings of 

some basins truncate those of other basins (fig. 1.7), and the continuous deposits of younger basins swamp older basins and craters (fig. 

7.5). Secondary-impact craters of basins are so common and extensive 

on the Moon that units more than 1,000 km from the basin rim can 

potentially be related by superpositions (figs. 7.6, 7.7; Wilhelms, 

1976). The source of morphologic features distant from basins is even 

easier to determine than is the source of crater secondaries because 

more outer lineaments of basins are radial. Many superpositions are 

detectable on the basis of radial structure even in highly degraded 

terranes. The possibility that, around old basins, the linear or 

chainlike pattern of basin secondaries may be indistinguishable from 

the radial-flow lineations does not matter in establishing superpositional relations.

Near a basin rim, craters younger than a basin appear much 

fresher, of course, than those older. Farther from the rim, this difference becomes increasingly less apparent as prebasin craters are 

increasingly less affected by the basin deposits. Morphologic differences between buried and superposed craters caused by basin 

ejecta are generally detectable even around old basins out to one 

basin radius from the rim (the average range of the continuous ejecta) 

and commonly out to one diameter (the average range of abundant 

secondary craters). Basin deposits, however, are lobate and asymmetric, and so each case must be considered individually. Southeast of 

Orientale, subtle grooves and ridges radial to that basin can be used 

to distinguish pre- from post-Orientale units beyond 1,300 km from 

the Orientale rim (Wilhelms and others, 1979). Distant craters are 

easily datable if fortuitously struck by large projectiles (fig. 7.6). 

Orientale ejecta in "rays" may have excavated secondary craters as far 

away as the central and southeastern terra (figs. 4.6, 7.7; Wilhelms 

and others, 1978; Wilhelms, 1980). On the other hand, dating of some 

craters only a few hundred kilometers from the rim may be ambiguous. For example, the swale-and-hummock topography of the zone 

near the crater Criiger could be interpreted either as superposed but

poorly developed Orientale-secondary craters or as well developed 

Orientale secondaries blanketed by Criiger ejecta (fig. 7.8). The 

weaker primary textures around older basins cause even more ambiguities (fig. 7.9). Some older craters may be scored or cratered in one 

sector but unaffected in another (figs. 7.6, 7.9). Therefore, determination of superpositional relations depends on knowledge of the whole 

geologic setting of the units in question.

Mapping conventions

Assignment of circumbasin materials to time-stratigraphic units 

older than, contemporaneous with, or younger than a given basin 

depends partly on interpretations of the topography, as is again well 

illustrated by the history of changes in mapping convention around 

the Imbrium basin.

During most geologic mapping at 1:1,000,000 scale, the Imbrium 

sculpture was thought to originate by impact-induced faulting (see 

chap. 6). Mappers assigned units clearly cut by the sculpture to the 

pre-Imbrian, and units superposed on the grooves to the Imbrian or 

younger systems (fig. 7.7). The grooved terrain was mapped as 

Imbrian where the grooves and adjacent swales were interpreted as 

intrinsic to the Fra Mauro Formation and thus contemporaneous 

with Imbrium (for example, Howard and Masursky, 1968; Wilhelms, 

1968).

The newer interpretation of the rimmed outer grooves as secondary-impact chains changes the convention in detail though not in 

principle. Units impacted by the chains of secondaries are still 

mapped as pre-Imbrian, that is, either Nectarian or pre-Nectarian. A 

difference from earlier mapping conventions, however, results from 

expansion of the terrane interpreted as secondary to Imbrium. Most 

terrane mapped in the past as pre-Imbrian lineate material (fig. 7.7; 

Wilhelms and McCauley, 1971) is here mapped as Imbrium-related to 

emphasize the age of the surficial morphology (pis. 3, 8). Lunar 

Orbiter photographs have revealed similar radially lineate terrane 

over a vast area (see chap. 10). Lunar Orbiter and Apollo photographs, 

coupled with the Apollo 16 results, have also led to changed age 

assignments for such features as "volcanic domes," reinterpreted as 

parts of the Imbrium-secondary regime and thus lowermost Imbrian 

instead of Imbrian or younger (see chap. 10).

FIGURE 7.5. Terrain southeast of Imbrium basin, lying just beyond lower right corner of 

photograph. Circular mare in top half of photograph is Mare Vaporum, about 220 km 

across. Linear- and knob-textured terrae are parts of Imbrium-basm ejecta. Circularity of 

mare suggests that it fills a Vaporum basin or crater, which was deeply buried by Imbrium 

ejecta before mare flooding. Several small craters (arrows) were superposed on Imbrium 

material before mare flooding. Graben near left-hand crater (above and parallel to arrow) is 

also post-Imbnum and premare (graben is part of set possibly not related to basins; see chap. 

6). Rima Hyginus is in background (H). View southerly. Apollo 17 frame M 1674.
a **«..¥»«

S-. -V

FIGURE 7.6. Terrain southeast of Orientale basin, heavily affected by its ejecta. Large circular 

crater is Schickard (227 km, 44° S., 55° W.), whose rim is scored by linear ejecta of Orientale 

secondaries (s). Barrage of secondaries with linear or braided ejecta destroyed part of rim of 

crater Schiller C (SC; 49 km) but left other rim segments untouched; Schiller C is 1,475 km from 

Orientale center and 1,000 km from rim (Monies Cordillera). Orbiter 4 frame H-160.
7. RELATIVE AGES 129

In lunar geologic mapping, the age of an ejecta unit applies to the 

time it was emplaced, which is not necessarily the same as the isotopic 

ages of the contained materials. 72 The materials contained in 

Imbrium-secondary ejecta may consist partly of projectile material 

from the Imbrium basin, they may be identical to materials of the 

impacted and redistributed pre-Imbrian terrane, or they may contain 

intimate mixtures of the primary and redistributed "local" material. 

As illustrated by the Apollo 16 materials, the provenance of the 

constituent materials is commonly difficult to determine either photogeologically or petrographically (see chaps. 2, 9, 10). Mapping of 

large tracts as related to Imbrium (pi. 8) stresses the interpreted 

basin age and origin of the morphology over the age and provenance of 

the contained materials. Similar conventions are used here for all 

basin peripheries to show the full extent of basin-related deposits 

(pis. 3, 6-8,12).

CRATER DATING

General principles

This account has shown how the Moon's extensive datum horizons, as well as less extensive strata, can potentially be ranked 

stratigraphically by detection of superpositional relations among 

their depositional features. These conceptually simple methods must 

be supplemented on much of the lunar surface by chronologic interpretations of small craters superposed on other units. Craters are 

used to correlate isolated units in the same way as fossils are used on 

the Earth (Mutch, 1970; McGill, 1977). Craters are thus (1) stratigraphic units in their own right and (2) time counters for other units, 

including other, larger crater units.

Which of these stratigraphic roles a given crater plays depends 

on the purpose of the investigation and on the size of the crater 

relative to the steady-state size for the cratered surface. The continuous bombardment of cosmic projectiles creates craters of all sizes, 

from micropits to basins. Within any given time interval, far more 

small than large excavations are produced. A surface becomes saturated with the smallest craters, so that new impacts destroy as many 

such craters as they add. Craters of these sizes are thus in a steadystate or equilibrium condition. The steady-state sizes for a surface of a 

given age are limited by a computable crater diameter, conventionally 

designated Cs (Moore, 1964a; Shoemaker, 1965; Trask, 1966; Shoemaker and others, 1967a, b; 1969; Gault, 1970; Marcus, 1970; Morris 

and Shoemaker, 1970; Soderblom, 1970; Soderblom and Lebofsky, 

1972). Craters smaller than Cs are in all states of aging, from the very 

fresh to the totally flattened. Craters larger than Cs range from very 

fresh to only partly destroyed. Cs increases over time (fig. 7.1QB; table 

7.3).

The first of two dating methods described here, size-frequency 

counts, is based on craters larger than Cs. I describe it only briefly 

because it is familiar and has long been widely used (Opik, 1960; 

Shoemaker and others, 1962a; Baldwin, 1963,1964). The second set of 

methods is based on erosional morphology created mostly by small 

craters acting in the steady state. The visible craters they have 

degraded are assumed to have been morphologically similar when

7 2On all but the most recent geologic maps (Wilhelms and others, 1979), massif materials were exempted from 

the practice of assigning ages according to time of emplacement. Massifs were assigned the age of the component 

prebasin rock to stress the interpretation that they are structural uplifts of relatively cohesive materials not 

severely broken up and redistributed by the impact. For example, Imbrium massifs were mapped as pre-Imbrian 

massif material, and Nectarian basin massifs and rugged material as Nectarian or pre-Nectarian (Wilhelms and 

McCauley, 1971; Wilhelms and El-Baz, 1977). In this volume, massifs are assigned the age of the basin of which they 

are parts.

FIGURE 7.7. Pre- and post-Imbnum craters, distinguishable over extensive area. PreImbrian age of Flammarion (F; 75 km), Ptolemaeus (Pt; 153 km), and Alphonsus (Al; 1 19 

km) is demonstrated by superposition of linear "Imbrium sculpture" (Gilbert, 1893); preImbrian age of Purbach (Pu; 118 km), Regiomontanus (R; 124 km), and La Caille (L; 68 

km) is demonstrated by superposition of more nearly circular secondary craters. Plains filling 

the craters also postdate sculpture and secondary craters, possibly by a very short time. 

Regiomontanus is 1,400 km from Imbrium rim (Monies Apenninus) and 2,100 km from 

center. Craters Herschel (H; 41 km), Arzachel (Ar; 97 km), Thebit (T; 57 km), and Werner 

(W; 70 km) lack Imbrium effects and thus are post-Imbrium (Imbrian and Eratosthenian). 

Arrows indicate possible secondary craters of Orientale basin, centered 2,600 km in 

direction of arrow shafts (south of west). Mosting (M; 25 km, 0.5° S., 6° W.) and the craters 

superposed on rim of Thebit (Thebit A and Thebit L) are Copernican. Orbiter 4 frame 

H-108.

formed, within a given size range. Surfaces are assumed to have been 

worn down uniformly over the whole Moon in proportion to their time 

of exposure. The small eroding impacts were both primary and secondary; the secondaries formed in proportion to the number of primaries (Moore, 1964a; Shoemaker, 1965; Gault, 1970; Soderblom, 

1970; Neukum and others, 1975a). This gradual erosion by small 

impacts seldom creates novel landforms but quantitatively alters the 

original crater morphology. Modifications that, in contrast, took place 

rapidly or episodically can be recognized by identifying superposed 

ejecta blankets, mare filling, or such anomalies as tectonically 

uplifted floors. In practice, the morphologically based methods are 

divided into those applicable to craters either larger or smaller than 

about 3 km in diameter.

Size-frequency relations

The use of statistical studies of craters as a dating tool is based on 

the assumption that primary impacts crater a surface in proportion to 

its time of exposure. The analyst must, therefore, (1) count only 

craters larger than Cs; (2) distinguish primary, secondary, and endogenic craters; and (3) eliminate buried craters from the counts (fig. 

7.10; Neukum and others, 1975a,b; Young, 1977; Wilhelms and others, 

1978; Basaltic Volcanism Study Project, 1981, chap. 8). For most stratigraphic purposes, craters should be counted on individual geologic 

units and not over whole "provinces." Where the geology is too complex or the data inadequate to attain this ideal, as is commonly the 

case (Hartmann, 1972c, p. 54), other methods must be applied. Thus, 

stratigraphic analysis should accompany any use of craters in determining ages.

Figure 7.10 summarizes the properties of size-frequency curves 

likely to be encountered in lunar work. The inverse relation between 

the masses and frequencies of impacting primary bodies produces 

inverse size-frequency relations of craters (fig. 7.10A). Continued 

bombardment by the same population of bodies moves the production 

curves higher on the graph, but the slopes remain the same (fig. 

7.WB). The slope of the steady-state distribution also remains the 

same, always shallower than the production curves (fig. 7.10.B). 

"Rollovers" of the end of a curve representing small diameters result 

either from inability to observe small craters or from mutual obliteration of craters in and near the steady state (fig. 7.10C). Subsequent 

burial removes the smaller members of a crater population and thus 

truncates the production curve (fig. 7.10D); this truncation is theoretically abrupt but, in fact, may resemble the effects illustrated in 

figure 7.1QD. Resumed cratering may have the same production curve 

as did preburial cratering, but the offset in the curve remains visible 

because craters of the sizes that were buried remain relatively 

depleted (fig. 7. WE). Finally, craters may not all have the same production function (figs. 7.10F, G); cumulative plots of craters larger 

than a few kilometers across are commonly observed to slope at 1.8 

(Basaltic Volcanism Study Project, 1981, ch. 8), but bends appear at 

smaller diameters. Plots of secondary craters differ considerably in 

slope from those of primaries (fig. 7.1 OF; Shoemaker, 1965; Wilhelms 

and others, 1978), and the addition of substantial numbers of secondaries to crater populations may cause the bends observed in many 

crater-frequency curves below diameters of about 2 km (fig. 7.10G).

Additional aspects of and problems with these methods are discussed in later chapters of this volume (see reviews by Baldwin, 1963, 

1964; Hartmann, 1964a, 1972c; Chapman and Haefner, 1967; Gault, 

1970; Greeley and Gault, 1970; Neukum and others, 1975a, b; Young, 

1975; McGill, 1977; Wilhelms and others, 1978; Crater Analysis Techniques Working Group, 1979; Basaltic Volcanism Study Project, 1981, 

chap. 8).

Morphology of large craters

To some extent, the overall morphology of large craters indicates 

relative age and stratigraphic correlation. Baldwin (1949, p. 134; 1963, 

p. 191) may have been the first to show that older lunar craters are 

systematically shallower than younger craters of the same size. 

Crater classifications have been developed independently on this 

basis (Arthur and others, 1963; Baldwin, 1963, chap. 9; Wood and 

Andersson, 1978).

The morphologic system most used in the U.S. Geological Survey's lunar mapping program was based less on overall morphology 

than on the morphologies of parts of craters (fig. 7.11). Crater-subunit
130 THE GEOLOGIC HISTORY OF THE MOON

TABLE 7.3. Stratigraphic criteria for lunar time-stratigraphic units

[Cs , limitiny diameter of the steady state; all values approximate, estimated from the data of Boyce and Johnson (1977) and Moore and others (1980b, table 1), from rollovers in the 

curves in fiyure 7.16, and from the approximate formula D L =1.7CS (Moore and others, 1980b, p. B8).

D|_, diameter of largest crater theoretically eroded to 1° interior slopes (see text); n/a, not applicable because method is invalid for the stated system or series (see chap. ID).

Subunits: Approximate diameters of craters datable by evaluation of their material-subunit morphologies.

Crater frequencies (cumulative); see chapters 8 through 13 for the basis of the stated values; n/a, not applicable because craters of this size are not useful for age determinations on the system or series.

Stratigraphic superpositions: X, Stratigraphic relations determinable; radii, distance from the rim crest within which the relations are generally detenninable; n/a, not applicable because no units are known.

Stated values of C s , D|_, and crater frequencies are the ranges for the entire system or series; they overlap for adjacent systems or series because of uncertainties, and depend 

partly on the substrate (in parentheses)]

System or series

Eratosthenian System- - 

Upper Imbrian SeriesLower Imbrian SeriesPre-Nectarian system 

C s (m)

? 

<100 (mare) 

80-30D (mare)

32D-860 (basin) 

8DD-4.000? (basin) 

>4,ODD? (basin)

D L (m)

<165 (mare), 

£?OD (crater)

145-250 (mare) 

23D-550 (mare)

n/a 

n/a 

n/a

Subunits 

(crater 

diameter, 

in km)

>3 

>3 

>3

>5 

>2D 

>20-30

Crater frequency 

(number per km 2 )

>1 km

<7.5xlO'4 (mare), 

<1.0xlD~ 3 (crater)

7.5X10" 1* to 

 v.2.5xl(T 3 (mare)

^ 2.5xlO~ 3 (mare) to 

-X-2.2X1D- 2

-x.2.2-4.8xlD' 2 (basin) 

n/a 

n/a

>20 km

n/a 

n/a 

2.8xlO~5

1.8-3. 3xl(T5 

2.3-8.8xlD~5 

>7.0xlO' 5

Crater-mare

X 

X 

X

n/a 

n/a 

n/a

Stratigraphic superpositions

Crater-crater

^2 radii for large cratersCrater-basin

n/a 

n/a 

n/a

X 

X 

X

morphology was extensively applied as an age criterion before basin 

deposits and their secondary-crater fields were completely mapped, 

and is still used in areas lacking good stratigraphic-datum planes. 

The method, developed by Pohn and Offield (1970), was the initial 

means for extending the Imbrium-region stratigraphy to the rest of 

the Moon (Offield and Pohn, 1970). Correlations are best made by 

comparing the morphologies of isolated craters with those of craters 

previously dated relative to one of the regional datum planes. The

methods work best for primary-impact craters (for which they were 

intended) and less securely for secondaries, which vary greatly in 

morphology around a given source and do not initially resemble 

primary craters, except where situated far from the source.

Pohn and Offield (1970) first divided craters into three size 

classes based on the planimetric shapes of the rims of fresh craters of 

each class: More than 45 km diameter, round with distinct rim crenulations; 20 to 45 km, polygonal; 8 to 20 km, round (fig. 7.11; Mutch,

fy^$.KW .;

' *; * ^ :'« «** " ' » * s -

Ift.f^^l^^^,-^^;

» ''^\^^v : ^^wl^j^t'^"1r

t-^mim^

FIGURE 7.8. Large mare-filled crater Criiger (46 km, 17 S., 67 W.), 370 km from Monies Cordillera. Dunelike topography to southwest in large crater Darwin is decelerated surface-flow ejecta

of Orientale (see chap. 4), but Stratigraphic relation of other hummocky ejecta to Criiger is unclear. Orbiter 4 frame H 168.
7. RELATIVE AGES 131

1970, p. 272). The first two classes are complex; the third, simple. 

Because of these differences in the initial morphologies, a slightly 

different morphologic continuum is definable for each size class. 

Craters from 1 to 8 km in diameter can also be ranked by morphologic 

criteria (Offield and Pohn, 1970, p. C164). Craters in each class are 

ranked in order of age by an arbitrary decimal scale, 0.0 (oldest) to 7.0 

(youngest). The seven criteria for estimating age illustrated in figure 

7.11 still appear to be valid, except that interior radial channels do not 

occur even in many old craters.

This classification was widely used and generally successful. The 

sequence of eight ringed basins and the relative ages of several mare 

provinces derived by Offield and Pohn (1970) agree with those given in 

this volume. Plains units were found to be only slightly older than old 

mare units a correct result that was misinterpreted on some geologic maps to mean that a terra type of volcanic plains is coeval with 

the mare type (for example, Rowan, 1971b). Offield and Pohn (1970) 

believed the light plains to be volcanic, but their age ranks showed the 

plains to be consistent with Orientale and Imbrium provenances. 

They also stated that "Numerous craters not obviously related to the 

Orientale and Imbrium basins are determined to be of the same age 

as the basins and are tentatively identified as basin secondaries" 

(Offield and Pohn, 1970, p. C168) an important result subsequently 

extended to craters that "look" older and younger than the basins 

(Wilhelms, 1976; Oberbeck and others, 1977; Wilhelms and others, 

1978).

This volume also uses morphology as an age criterion where 

necessary. The next six chapters give the morphologic characteristics 

and a list of typical craters for each age. No attempt is made, however, 

to establish type craters or formulas for recognizing crater ages. Such 

an attempt would mislead more than help, in view of the many 

variables of photographic resolution, illumination, and, above all, 

geologic setting, that affect apparent crater morphology. Instead, the 

type craters are described and illustrated in their geologic context 

throughout the volume.

FIGURE 7.9. Rim of pre-Nectarian crater Sacrobosco (S; 98 km, 24° S., 17° E.), lineate 

radially to Nectaris basin; basin rim (Rupes Altai, RA) is only 100 km northeast. 

Unlabeled west rim of Sacrobosco is less conspicuously affected by Nectaris material. Most 

smaller craters are probably secondaries of Imbrium and Nectaris basins. Orbiter 4 frame 

H-89.

Morphology of small craters

By Newell J. Trask (modified from Trask, 1971)

To fill the need for stratigraphic correlations of units of interest to 

early Apollo landings, Trask (1969,1971) devised a system based on 

the morphology of small craters (less than 3 km diam). These ubiquitous features were and remain necessary tools for dating mare 

units because of the uneven distribution of stratigraphically discrete 

and extensive mare flows or crater deposits at closeup scales. The 

craters on the mare surfaces range in morphology from sharp to 

indistinct. The sharpest craters have well-defined rim crests and 

bright, blocky ejecta blankets. Less sharp craters have rounded rim 

crests, little in the way of an ejecta blanket, and few surrounding 

blocks. The most indistinct craters are very shallow depressions and 

lack a well-defined rim crest. Figure 7.12 shows examples of this 

continuum of crater types for various crater diameters. Where craters 

of the same approximate size are superposed, the sharper are superposed on the less distinct, indicating that the sharper craters are the 

younger. These observations suggest that most new craters are sharp 

and distinct and become progressively more subdued and indistinct 

over time; in other words, craters in the lower frames of figure 7.12 

once looked like those in the upper frames. The sharpest craters (top, 

fig. 7.12) have all the characteristics of impact craters, including rays 

and ray loops. Thus, nearly all craters in the 0.05- to 3-km diameter 

range were probably formed by impacts. Older surfaces should, therefore, have more craters than younger surfaces and should also have 

more older-appearing craters than younger surfaces. Such a relation 

between the number of craters on the surface and the number of oldappearing, subdued craters is, in fact, observed.

Observation of these relations on the entire suite of Lunar 

Orbiter photographs led Trask (1969,1971) to the following working 

hypotheses. (1) When first formed, craters appear sharp and fresh. (2) 

All craters undergo modification to less sharp and more subdued 

forms by erosion of the rim crest and infilling of the floors. (3) Smaller 

craters disappear sooner than larger ones. (4) There exists for any 

surface a critical crater diameter (Csl, Cs2, Cs3, fig. 7.10B), which 

increases over time for a given surface; craters smaller than Cs that 

are as old as the surface have had time to be completely destroyed, 

whereas those larger have not.

These hypotheses, in conjunction with the assumption that the 

rate of crater degradation is approximately the same everywhere on 

the maria, formed the basis for assigning geologic ages to craters in 

geologic mapping of potential Apollo landing sites (Trask, 1969,1971). 

By concentrating on the oldest craters, surfaces could also be dated. 

Figures 7.12 and 7.13 illustrate this system of assigning ages. The 

horizontal lines in both figures represent equal points in time, or 

isochrons. Accordingly, a gentle depression 50 m in diameter, for 

example, is of the same age as a 1-km-diameter crater that is only 

moderately subdued; and a 50-m-diameter moderately subdued 

crater is of the same age (Copernican) as a 1-km-diameter crater with 

well developed rays. The degree of crater degradation implied by this 

assumed equivalence seems reasonable in the light of theoretical 

studies on the extent and mechanism of crater subdual by the impact 

of small particles (Soderblom, 1970). The number, spacing, and names 

of the time intervals on the vertical axes of figures 7.12 and 7.13 were 

chosen to be consistent with the stratigraphic scheme used in 

regional mapping. In that scheme, the base of the Imbrian System is 

defined by the Fra Mauro Formation, upon which gentle depressions 

as much as 1 km in diameter are superposed; therefore, 1-km-diameter gentle depressions are placed at the base of the Imbrian System in 

the scheme of figures 7.12 and 7.13. Similarly, the dividing line 

between the Eratosthenian and Copernican Systems generally is 

placed between craters with bright halos and craters without bright 

halos in the diameter range 5 to 10 km in both small- and large-scale 

mapping, although individual cases may depart from this rule. The 

six numbered subdivisions of the Copernican System were chosen 

solely for convenience in the detailed site mapping and are not used at 

regional scales. The approach and assumptions are similar to those 

applied in the system for determining relative ages of craters larger 

than 8 km in diameter that was described above.

Trask (1969,1971) examined high-resolution Lunar Orbiter 1, 2, 

3, and 5 photographs of parts of the maria to determine whether 

systematic differences exist in the relative ages of the oldest superposed craters as determined by this scheme. The isochrons labeled 

"Eratosthenian mare" and "Imbrian mare" in figure 7.13 refer to two
132 THE GEOLOGIC HISTORY OF THE MOON

mare surfaces with different crater populations and different apparent relative ages. Craters 700 m in diameter on the Imbrian mare 

range in morphology from indistinct gentle depressions, barely mappable as craters, to fresh excavations; craters of the same size on the 

Eratosthenian mare range from strongly subdued but still quite 

distinct to fresh. Thus, craters of the same size may be older on the 

Imbrian than on the Eratosthenian mare, as expected (fig. 7.14). The 

Apollo 12 landing site (fig. 7.14B) falls near the EratosthenianImbrian boundary in this system. Still younger mare surfaces not 

visited have gentle, indistinct depressions no smaller than 300 m in 

diameter (fig. 7.14A). The most highly degraded craters 10 m in 

diameter on Copernican 6 surfaces would be bright and blocky. In 

summary, a search for the largest craters barely visible on a surface 

can be used to date surfaces.

A difficulty with this system and all others based on morphology 

is the likelihood that craters of different origins may have different 

initial morphologies (Trask, 1969). Secondary-impact craters near 

their sources will begin their existence with the subdued rim crests 

and shallow depths characteristic of older primary craters of their 

size. Attempts were made in mapping of the Apollo landing sites to 

distinguish secondary craters by their clustering, alignment, and 

elongation; this system was applied to the secondaries cautiously and 

with allowance for the probable initial differences.

Another difficulty is that the physical characteristics of the materials in which craters form influence their sizes and morphologies. 

Craters in such low-cohesion materials as regolith, pyroclastic glass, 

and fragmental terra breccia will appear more subdued even when 

first formed than will craters in basaltic bedrock (see chaps. 12, 13;

DIAMETER DIAMETER DIAMETER

DIAMETER

10" 

10' 

10 6

DIAMETER

IO3 

10 2 

10' 

10°

10-'

10~2

io-3 

10-* 

io-5 

io-6

\

\- \

I \ I

10 3

10 !

10°

N 10-

io-5

IO 2 IO 3 IO 4 IO 5 IO 6

CRATER DIAMETER, IN METERS

io-2 10-' 10° io 1 io2

CRATER DIAMETER (D), IN KILOMETERS

FIGURE 7.10. Idealized cumulative crater-frequency curves representing different geologic 

and observational conditions.

A. Crater-production functions with indicated slopes. 

B. Distributions of three different ages but same production function. Limiting diameters of

steady state for all three ages (Cs] , Cs2 , C53 ) define a curve sloping more gently than

production functions; if production slope is 3, slope of steady-state distribution is 2

(Trask, 1966). 

C. Gradual rollover for frequencies of small craters, caused either by (1) inability to observe

small craters with available photographic resolution or by (2) removal of small craters by

mutual interactions of primary- and secondary-crater populations. 

D. Production function (slope, 2), truncated by burial of smallest and shallowest craters.

Rollover is gradual because diameter cutoff is not precise under natural conditions. 

E. Frequency of D (dashed), altered by renewed cratenng. 

F. Idealized plots of secondary craters (steeper slope, about 3.6) and primary craters

(slope, about 2). After Shoemaker (1965). 

G. "Calibration distribution" of Neukum and others (1975a), showing different slopes in

three diameter ranges. Steep slope in middle regime probably results from mixtures of

secondaries and primaries, as in F.
7. RELATIVE AGES 133

Lucchitta and Sanchez, 1975). Moreover, craters formed in low-cohesion material degrade more quickly than those in cohesive bedrock. 

These effects are complex in craters formed in interlayered rock 

types, such as regolith and basalt (Young and others, 1974; Young, 

1975, 1977; Schultz and others, 1977; Chapman and others, 1979). 

Other postformational modifications, such as downslope movement 

dependent on angle of slope, will also cause variations in morphologies and frequencies among contemporaneous craters of the 

same size (Trask, 1966,1969). Despite the need to consider so many 

variables, the general validity of this dating system in any one area is 

supported by the superpositions of craters assigned young ages on 

craters assigned older ages, and by the densities of small superposed 

craters.

DL method

A quantitative refinement of the above model has been extensively applied by the U.S. Geological Survey in dating geologic units 

too small to be dated reliably by crater counts (Soderblom and Boyce, 

1972; Boyce and Dial, 1973, 1975; Schaber, 1973; Boyce and others, 

1974,1975; Boyce, 1976; Moore and others, 1980b, p. 3-18). Craters 200 

to 2,000 m in diameter are generally examined. The method is based 

on a theoretical model of crater erosion operating in the steady state 

(Soderblom, 1970; Soderblom and Lebofsky, 1972). As in the method of 

Trask described above, the largest crater almost destroyed on a surface is sought. However, because a search will not necessarily disclose 

that crater, the crater's size is calculated. On a given surface, craters 

somewhat larger than Cs will be less nearly obliterated than the 

craters of diameter Cs. The hypothetical larger craters selected are 

those whose walls slope at 1°; their diameters are DL (in meters). As a 

rule of thumb, DL = 1.7CS .

Two related techniques are used to determine DL values. The 

technique originally devised (Soderblom and Lebofsky, 1972) first 

determined the largest crater whose slope is eroded to the Sunillumination angle of the photograph being examined. In practice,

Polygonality

Satellitic

craters

Radial 

fades

Sharp 

rim crest

Radial 

channels Terraces

Class I (>45 km)

Class II (20-45 km)

vy

Class III (1-20 km)

FIGURE 7.11. Changes in diagnostic crater features with age (Pohn and Offield, 1970). 

Numbers 0 through 7 denote seven morphologic stages corresponding to relative ages, from 

oldest to youngest. Width of symbol indicates relative degree of development of each feature at 

each stage.

this crater (diameter, Ds) is bracketed by finding the largest 

unshadowed crater and the smallest diameters at which all craters 

are clearly shadowed. Ds is then converted mathematically to DL (DL 

is thus a theoretical diameter). In the second technique (Boyce and 

Dial, 1975), which can be applied to photographs of lower quality than 

the first, the crater (diameter, Dm ) that is shadowed to the midpoint of 

its floor is bracketed. The upper bound on Dm is the diameter for 

which all craters of that size and larger are shadowed to the midpoint; 

its lower bound is the largest diameter of those craters that are 

shadowed halfway or slightly less. Dm is also converted mathematically to DL.

DL values of geologic units are quoted extensively in this volume. 

The method appears to have yielded correct results for mare units 

(Wilhelms, 1980) but less secure results for light-plains units (see 

chap. 10). Its successful application depends on several conditions: (1) 

Photographic quality sufficient to distinguish true shadows and true 

diameters; (2) Sun-illumination angles between 8° and 20° above 

horizontal (Soderblom and Lebofsky technique), or between 8° and 

30° (Boyce and Dial technique); (3) a sample area of at least 100 km2 

for young surfaces (Sun angle, 10°) and 1,800 km2 for terra plains (Sun 

angle, 20°; Boyce and others, 1975; Moore and others, 1980b, fig. 7); (4) 

the same initial crater shapes; (5) crater-wall slopes of from 8° to 25°, 

the interval over which the underlying theoretical model, assuming 

erosion by small impacts, applies; (6) craters exposed, not blanketed 

by later deposits; (7) craters having a slope of - 3 on cumulative sizefrequency plots when formed; (8) craters formed on level surfaces 

(detectable on the photographs); and (9) inclusion of craters superposed on only one geologic unit. Condition 7, which may not apply to 

primary craters larger than 2 km in diameter (fig. 7.10G), probably 

explains why DL and size-frequency values are inconsistent for light 

plains and other units characterized by large superposed craters 

(chap. 10; fig. 7.15;Neukumandothers, 1975a). Condition 9 may be the 

one most frequently violated, because measurements have commonly 

been made at the center of rectangles laid along a grid rather than 

within the confines of units mapped geologically. Thus, small units 

may have been missed.

Most of the pitfalls of the DL method, or any other attempt to 

determine the ages of lunar and planetary materials, can be avoided 

by commonsense and examination of a surface from the viewpoint of 

the sequence of events through which it was shaped.

SUMMARY

Lunar stratigraphic units are relatively dated by a combination 

of superpositional criteria and methods based on superposed craters. 

Pairs of adjacent crater, basin, or mare units can normally be dated by 

straightforward observations of superpositions and transections 

based on the simple concept that the younger units modify the older. 

In favorable cases, secondary-impact craters extend the areas in 

which these observations can be applied.

Craters play a dual role as individual stratigraphic units treated 

by these relations and as time counters for other units' ages. Crater 

dating rests on the assumption, implied by the law of superposition, 

that the craters overlying a given stratigraphic unit have size-frequency distributions and morphologies typical of stratigraphic intervals younger than the unit. Old units support more craters than do 

younger units. Topographically sharp craters are superposed on all 

units, but large subdued craters only on old units (unless the subdual 

results from burial by still younger deposits).

The properties of the inverse size-frequency relation result in a 

hierarchy of dating methods (table 7.3). Direct observation of craters 

smaller than Cs is not useful in dating because fewer craters are 

observed than were formed. However, observation of the cumulative 

effect of these craters on morphologies is useful in dating. Craters 

slightly larger than Cs are dated by the method of Trask described 

above, or by the DL method. The determined ages bracket the ages of 

small planar units on which the craters formed. Still larger craters 

(more than 3 km diam) are dated by assessments of the degradation of 

their subunits. Ages of whatever units underlie or overlie these craters can then be bracketed. Sufficiently extensive units can be dated 

not only by individual craters but also by statistical evaluation of the 

superposed craters larger than Cs. In general, because Cs increases 

over time, increasingly larger craters must be depended on for dating 

of increasingly older units.
134 THE GEOLOGIC HISTORY OF THE MOON

Crater diameter

1-3 KILOMETERS

0 1 KILOMETER

400 METERS-1 KILOMETER 200-400 METERS 100-200 METERS

 ^^St *i£**-- sA*'-- ->

50-100 METERS

FIGURE 7.12. Craters smaller than 3 km in diameter, arrayed by age and size. Sun-illumination angles are mostly 19.5° above horizontal; a few are from 

18.2° to 23°. Craters of each size decrease in sharpness with age. From Trask (1971, fig. 2).

LU 

O

< |^4 

LU »

F 5 

<

_i

LU 

CC

Eratosthenian mare

Imbrian mare

10m 50m 100m 500m 1km 5km 10km 

CRATER DIAMETER (RIM CREST TO RIM CREST)

FIGURE 7.13. Assumed relations between sizes, properties, and relative ages of craters from 

0.01 to 10 km in diameter. Categories are intergradational; horizontal lines are isochrons 

indicating crater populations on Eratosthenian and Imbrian mare materials. From Trask 

(1971, fig. 3).
7. RELATIVE AGES 135

Thus, there is a substantial practical difference in dating the six 

divisions of the lunar stratigraphic column considered here (fig. 7.16; 

table 7.3):

1. Copernican units are dated by the frequencies and morphologies of 

craters so small as to be visible only on the best photographs.

2. Eratosthenian crater units photographed at high resolution can 

also be dated by size-frequency counts. Eratosthenian mare units 

and some impact-melt pools can be dated by the DL method.

3. Late Imbrian mare units are readily accessible to the DL method.

4. Only a few crater units of Late or Early Imbrian age are large 

enough to support enough craters larger than Cs for statistically 

valid size-frequency counts. These craters are dated mostly by the

degradational morphology of their own subunits.

5. The Early Imbrian Orientale and Imbrium basin materials are 

extensive enough, young enough, and well enough photographed 

to be dated by craters over a large diameter range (fig. 7.16).

6. The steady-state size increases for older units (fig. 7.16; table 7.3), 

and only craters larger than 20 km in diameter are thought here to 

be valid as counters for the whole Nectarian and pre-Nectarian 

age range. Morphologies of individual Nectarian and pre-Nectarian craters larger than about 20 km in diameter supplement 

the size-frequency counts. 

The next six chapters show in detail how the lunar stratigraphic

column is constructed by these methods.

f *ty».v,ryi , , « ,,; ^.rv".-* - ,-, * *^a»;-*s n^f%^Ti , »'-! %' ;,> : ,%*"it T'^v 's *5?.^f^

iL^*'^" ^X-"" ' '»v * '** " '^ :?"^ ";* r:^ -.v ».^5^|j

.;*V.r>

> * ,-; "» -\ \t ' " » * ;\ - J "*., ", ;»';^i'-vi:»^.^ V'.;-,.^.i;^,^v ... ' = ; *,^..- .'- % ^t^'^

' ' V'-rH"-lj!:v * *^" ., '.*- " -A "* '*' ''" k*j^«*'^

5 KILOMETERS

5 KILOMETERS 5 KILOMETERS 

I

D

FIGURE 7. 1 4. Surfaces of four ages, showing differences in crater populations. Sun angles are all from 2 1 . 3° to 22.8° above horizontal. From Trask 

(1971, fig. 4).

A. Young Eratosthenian mare in southern Oceanus Procellarum, characterized by small sharp craters and much level intercrater terrain. 

B. Older Eratosthenian mare near Apollo 1 2 landing site; larger and more subdued craters are present. 

C. Imbrian mare near Apollo 1 1 landing site, characterized by high density of subdued craters and little noncratered terrain. Arrow indicates low

swell discussed by Trask ( 1 971).

D. Fra Mauro Formation near type area. Saturation of surface by large subdued craters suggests that many more craters formed than are now 

visible.
136 THE GEOLOGIC HISTORY OF THE MOON

1000

100

10

1 ' ' "I ' ' ' ' I' 'I ' ' ' ' ' ' "I ' ' ' ' ' ' "I 

Lunar light plains regions

1. West of Apollo 15 LS 5. Parry floor 10. Ptolemaeus floor

2. Apollo 16 LS 7. Bonpland floor 11. Albategnius floor

3. North of Apollo 16 LS 8. Guericke floor 12. Meton floor

4. Fra Mauro floor 1 5. Mendeleev floor

16 Other regions

6. Apollo 14 LS 20. Apollo 12 LS9. Mare west of Archimedes 21. Copernicus

13. Mare Tranquillitatis 22. Mosting

14. Apollo 11 LS 23. Dawes

16. Mare near Brayley 24. Tycho

17. Mare Serenitatis (light int.) 25. Eratosthenes

18. Oceanus Procellarum 26. Aristarchus

19. Apollo 15 LS (mare)

Orientate 

event

Imbrium 

event

10- 10- 10- 10-

CUMULATIVE CRATER FREQUENCY PER KM 2 FOR D=ONE KM

FIGURE 7.15. Comparison of DL values and cumulative crater frequencies, normalized to 

crater diameters of I km (Neukum, 1977). For many plains whose DL value is 550 m, DL 

values and crater frequencies do not correlate. "Event" refers to basin materials; Imbrium 

"event" counts were made on Montes Apenninus (Neukum and others, 1975a). LS, landing 

site.

10C

10-

o

2 io-2

10'

10-

io- 5 -

10-

0.1 0.2 0.5 10 20 50 100 200 500 1000

CRATER DIAMETER, IN KILOMETERS

FIGURE 7. 1 6. Cumulative size-frequency distributions of craters in each time-stratigraphic 

system. Curves are solid within the diameter range where frequencies are most accurately 

determined, and long-dashed where extrapolated. Short-dashed curves represent average 

frequencies of primary-impact craters of only the age indicated: pNc, pre-Nectarian; Nc, 

Nectarian; Ic, Imbrian; CEc, Copernican and Eratosthenian, undivided (after Wilhelms 

and others, 1 978). Other curves include all primary craters accumulated on the units. Where 

photographic data permit, successively smaller craters can be counted on successively younger 

units (many CEc and Ic craters smaller than 4.5 km in diameter are visible, but average 

frequencies were not determined). Numbers of smallest plotted Nc, pNc, Nectaris, and AlKhwarizmi/King craters are diminished relative to larger craters of same age (compare fig. 

7. 1 OC). Rollover in curve indicating this deficiency is sharper and extends to larger diameters 

in pNc than Nc (rollover is less pronounced for Al-Khwarizmi/King than for pNc because 

former curve is plot of all accumulated craters, including CEc and Ic, which have relatively 

linear distributions). Al-Khwarizmi/King curve is steep at largest plotted diameters because 

of small statistical sample (see supplementary table). Imbrium-basin material is only lunar unit 

for which statistically valid frequencies in entire diameter range are available. Counts on 

Copernicus from Neukum and Konig (1 976); counts on Montes Apenninus and on Fra 

Mauro Formation (combined here) from Neukum and others (I975a, b); Imbrium- and 

Nectaris-basin counts of craters at least 20 km in diameter from present study. Data and 

system boundaries further discussed in chapters 8 through 1 3 .
a ,> « *2b**^S». * * %

 S^Sr,

^ fcv-
FIGURE 8.1 (OVERLEAF). Predominantly pre-Nectarian terrane on northern farside, including densely cratered tracts and smoother areas covered by basin deposits and secondary craters.

Orbiter 5 frame M-53.
8. PRE-NECTARIAN SYSTEM

CONTENTS

Introduction ......................

Crustal petrology and geochemistry

Composition ..................

Igneous or impact origin? ......

Early crustal petrogenesis .........

Magma ocean ................

Multiple plutons ...............

Crustal zonation ...............

Basin and crater materials .........

Introduction ...................

Recognition criteria ............

Basins ........................

Craters .......................

Volcanic rocks ...................

Chronology ......................

Page

139

139

139

140

142

142

142

143

143

143

145

145

147

156

156

INTRODUCTION

The lunar pre-Nectarian system8- 1 provides our clearest look at 

the early Solar System. The barrage of cosmic bodies that must have 

affected all moons and planets is richly recorded by pre-Nectarian 

basins and large craters, which are packed more densely than those 

of any other age (pi. 6; fig. 8.1). Their deposits cover a triangular area 

of the southern nearside and large parts of the farside. Smaller 

craters presumably formed in proportional abundance but were 

degraded by other craters (see chap. 7). Fine distinctions among 

facies of basin and crater deposits, which are observed in rayed craters and the Orientale basin, have similarly been blurred but can 

commonly be detected by comparisons with the younger analogs. 

Any maria that may have formed must have been even more completely devastated by the rain of impacts. Faults and folds dating 

from the pre-Nectarian are unknown, although they may have 

formed then as they did later. Therefore, efforts to understand early 

lunar and Solar-System history depend largely on interpreting the 

generally subdued and modified deposits of pre-Nectarian basins 

and craters.

Another look into the distant past is provided by the samples 

returned by the Apollo and Luna spacecraft (pi. 3). When fully evaluated, data from the sample analyses should show much about the 

composition of Solar-System material, the bulk composition of the 

Moon, and the processes by which a body of the Moon's size and 

composition differentiates. A major problem confronting analysts is 

distinguishing the effects of these early chemical and petrologic 

processes from those of the later impacts which formed the deposits 

that contain the samples. By definition, all materials formed since 

the origin of the Moon until the formation of the Nectaris basin are 

pre-Nectarian (Stuart-Alexander and Wilhelms, 1975). Accordingly, 

both the stratigraphic units that are still visible and the plutonic 

rocks of the early terra crust that were originally created by igneous 

processes are pre-Nectarian.

It is important to stress this distinction between the internal 

origin and impact emplacement of lunar terra rocks. The various 

groups of scientists that have investigated these two phases in the 

rocks' history have not always communicated effectively. Most 

petrologists and geochemists have concentrated on the endogenic 

phases, and their interest in the origin and chemistry of the Moon 

and of its early crust is reflected in the most commonly used lunar 

petrologic nomenclature. For example, names derived from ter8- 1 1b indicate that the pre-Nectarian system is not a formal stratigraphic name, the term is not capitalized 

(see introduction to the section in this chapter entitled "Basin and Crater Materials").

restrial igneous petrology have been applied, inappropriately, to 

rocks melted and emplaced by impacts (Smith, 1974, p. 233; Irving, 

1975, p. 364; Prinz and Keil, 1977). Stratigraphers are more interested in the times and processes of assembly and deposition of the 

breccia beds. These events in the rocks' histories are emphasized 

here and in the breccia-classification scheme discussed above in 

chapter 3 (Stoffler and others, 1979, 1980). The early crust is not 

described comprehensively here (see Taylor, 1982) but is treated 

mainly from the standpoint of its role as raw material for later 

stratigraphic units.

Unfortunately for those interested in the early Moon, no rocks 

of the early crust are preserved in outcrop, and only small fragments 

believed to be relatively unaltered samples of that crust have been 

recovered from the breccia deposits formed and emplaced by later 

impacts. Still more seriously for attempts to decipher the early 

Moon, no deposits identifiable with a specific pre-Nectarian stratigraphic unit were sampled. All the terra landings took place on 

Nectarian and Imbrian deposits because operational factors prevented Apollo and Luna from landing on the farside and southern 

nearside tracts where pre-Nectarian materials are exposed (pi. 6). 

Therefore, the photogeologically decipherable record and the sample 

record for the pre-Nectarian are connected much more tenuously 

than for later time-stratigraphic systems. Correlation of pre-Nectarian absolute and relative time scales and determination of the 

rates of impact and volcanic activity are difficult problems confronting stratigraphic and laboratory analysis.

CRUSTAL PETROLOGY AND 

GEOCHEMISTRY

Composition

The lunar terrae are composed of a crust of material that differentiated from the bulk Moon early in lunar history. The crustal 

composition was determined during early pre-Nectarian time by 

processes that are basically igneous, and was later mixed and otherwise modified by impacts. This mixing aids attempts to determine the 

average crustal composition, although it obscures whatever geochemical provinces may once have existed. Data from unmanned 

Surveyor analyses (Gault and others, 1968a; Turkevich and others, 

1969; Phinney and others, 1970; Turkevich, 1971), returned-sample 

analyses, and orbital geochemical sensing (see chap. 5) currently 

combine to indicate an average terra composition rich in Al (25-29 

weight percent A12O3 ) and Ca (approx 14 weight percent CaO) (Taylor, 

1975,1982; Basaltic Volcanism Study Project, 1981, p. 666-678).

139
140 THE GEOLOGIC HISTORY OF THE MOON

This chemistry is reflected in the low density of the crust 

(2.90-3.05 g/cm3 ; Haines and Metzger, 1980; Basaltic Volcanism 

Study Project, 1981, p. 671) and in the fact that highly calcic plagioclase (mostly greater than 90 percent anorthite; Smith and Steele, 

1975) is the most abundant sampled mineral. The next most abundant minerals in terra samples are low-Ca pyroxenes (mostly 

orthopyroxene, with subordinate pigeonite), olivine, and augite (highCa clinopyroxene). Aluminous or chrome spinel, ilmenite, silica minerals, troilite, and Fe metal are common accessories, and many other 

trace minerals are present (Smith, 1974; Frondel, 1975; Smith and 

Steele, 1975; Taylor, 1975, table 4.2; Cadogan, 1981, chap. 3). Lunar 

rocks vary far less in mode (optically visible minerals) and in mineral 

chemistry than do terrestrial rocks.

A series of rocks commonly called the ANT suite composes most 

of the returned terra material, where 'ANT" refers to the plutonicrock types anorthosite (min 90 percent plagioclase), norite (plagioclase and orthopyroxene), and froctolite (plagioclase and olivine; 

table 8.1; Prinz and Keil, 1977; Stoffler and others, 1980). The plagioclase-rich composition of the terrae was first recognized in studies 

of small fragments in the Apollo 11 mare regolith (Wood, 1970; Wood 

and others, 1970), and the ANT suite was defined on the basis of small 

fragments in the Luna 16 and 20 regolith samples (Keil and others, 

1972; Prinz and others, 1973; Taylor, G. J., and others, 1973; Prinz and 

Keil, 1977). Because the importance of impact mixing and melting was 

not fully appreciated during these early analyses, compositional 

types were commonly delineated incorrectly on the basis of mixed, 

brecciated, fine-grained, or even glassy rocks as if they were plutonic 

igneous rocks containing large optically identifiable crystals. The 

name "anorthosite" has been especially misused for materials that 

are not anorthosites in either the compositional or textural sense.

In ANT terminology, the "average" terra rock is an anorthositic 

norite or a noritic anorthosite (approx 70 percent normative plagioclase; Taylor, 1975,1982). Unfortunately for the nonspecialist, this 

composition is also referred to as "highland basalt." This term was 

coined during study of regolith glasses (Reid and others, 1972a) and is 

one of several terms containing the word "basalt" that, especially in 

literature of the early 1970's, refer not to crystalline extrusive rocks 

but to magmas whose existence was predicted on the basis of glassy 

fragments (Apollo Soil Survey, 1971,1974; Reid and others, 1972a, b; 

Prinz and others, 1973). Early workers commonly stated that "highland basalt" represents the most important primary magma of the 

terrae.

ANT rocks occur at all terra-sampling sites, though greatly varying in amount from the small fragments returned by Apollos 11,12, 

and 14 and Lunas 16 and 20 to many kilograms from the Apollo 15,16, 

and 17 landing sites. Earthlike granite and granodiorite are absent, 

but rare fragments of "granite" and "quartz monzodiorite" have been 

found (Ryder and others, 1975a; Ryder, 1976); these materials consist 

dominantly of quartz and potassium feldspar and contain varying 

amounts of relatively Fe-rich pyroxene and olivine and of relatively 

sodic plagioclase.

The most abundant type of highly differentiated lunar terra 

material is known, confusingly, by many partial synonyms, of which 

KREEP and KREEP-rich rock are the most nearly appropriate (table 

8.1). The major-element composition of this material is not well 

defined, and the term refers primarily to an assemblage of minor and 

trace elements present in characteristic proportions (Meyer, 1977; 

Warren and Wasson, 1979c). The catchy acronym "KREEP" (coined by 

Hubbard and others, 1971, p. 343) refers to potassium (K), rare-earth 

elements (REE), and phosphorus (P); other diagnostic trace elements 

are Ba, Rb, Th, U, and Zr. These distinctive elements share the 

properties of large ionic radius (they are large-ion lithophiles [LIL]) 

and exclusion from the major mineral phases ("incompatible"). 

Although the name "KREEP basalt" commonly appears (Hubbard 

and Gast, 1971; Hubbard and others, 1971; Meyer and others, 1971), 

very few KREEP-rich samples have igneous textures. The term 

"KREEP" is properly used in a chemical, not a lithologic, sense for 

these diversely textured rocks. Major normative minerals of KREEPrich materials are similar to those of the ANT rocks, except that 

plagioclase is more sodic (mostly less than An88) and orthopyroxene 

and pigeonite are richer in Fe. Thus, both the minor- and majormineral phases indicate advanced differentiation according to the 

Bowen (1928) reaction series. The incompatible elements are the 

longest to survive in residual liquids of magmas from which crystals 

have separated, and are the first to appear in new partial melts.

TABLE 8.1. Major-oxide compositions of major types of lunar terra

material

[All values in weight percent. Others, mostly Ti02 , 

Na 2 0, K2 0, and Cr2 0 3 (Taylor, 1982, table 5.5); still 

other trace elements characterize KREEP (Warren and 

Wasson, 1979c). After Taylor (1975, p. 233-253; 

1982, p. 201-233, tables 5.5, 8.2; S.R. Taylor, oral 

commun., 1982); data on high-K KREEP from Reid and 

others (1972b).

ANT: Compositions generally range from anorthosite (A, 

left) to norite (N, right) or, for MgO, to troctolite. Gabbro is synonymous with norite, except that 

high-Ca pyroxene characterizes gabbro and low-Ca 

pyroxene characterizes norite (Taylor, G.J., and 

others, 1973). In the terrae, low-Ca pyroxene is 

more abundant than high-Ca pyroxene.

KREEP: Compositions range from low-K (left) to high-K 

(right). KREEP-rich materials are also called KREEP 

basalt, Fra Mauro basalt (Reid and others, 1972b), 

nonmare basalt (Hubbard and Gast, 1971), high-Al 

basalt, alkali high-Al basalt (Prinz and others, 

1973), norite (Wood, 1972a), KREEP norite, and many 

other names and acronyms (Taylor, 1975, p. 215, 272, 

table 5.8). Low-, medium-, and high-K KREEP-rich 

materials are commonly referred to as low-, medium-, 

and high-K Fra Mauro basalt (LKFM, MKFM, and HKFM, 

respectively).

Material 

Si02 

AloOo 

C. "J

CaO 

FeO 

MgO 

ANT (A-N)

44-50

O C 1C

19-10

1-11

1-15

<-7

KREEP 

(low-high K)

47-53

OO 1C

12-10

10

11-6

<3 5

Average terra

45 0

24.6

1C O

6.6

6

O

<1.2

A common classification recognizes three categories of KREEPrich materials that are similar in major-element bulk composition but 

differ in the proportions of the distinctive minor elements, 

exemplified by potassium: high K (more than 1 weight percent K2O), 

medium K (approx 0.5 weight percent K2O), and low K (approx 0.1 

weight percent K2O)(Taylor, 1975, p. 234-237). Also, the Fe/Mg ratio 

and SiO2 content increase as the K content increases. Low-K KREEP 

is the most abundant category and, in fact, may be the second most 

abundant component of the terra crust after the ANT suite (approx 

20 percent; Taylor, 1975, p. 234, 252). Low-K KREEP occurs in glasses 

or fragment-laden impact melts in some abundance at all landing 

sites except the mare-dominated Apollo 11 site (James, 1980). It was 

first suggested as a magma type on the basis of glassy fragments 

(Reid and others, 1972b). As a legacy of the large Apollo 14 collection, 

these three categories of material are commonly called high-K, 

medium-K, or low-K Fra Mauro basalt (HKFM, MKFM, LKFM) 

(Apollo Soil Survey, 1971; Reid and others, 1972b). However, this 

volume avoids the term "basalt" for nonvolcanic rocks and restricts 

the name "Fra Mauro" to the rock-stratigraphic unit the Fra Mauro 

Formation.

Igneous or impact origin?

A petrologically and stratigraphically important question is how 

many terra rocks have compositions that result from igneous processes and how many have compositions that were assembled by 

impacts. Until about the time of the Third Lunar Science Conference 

in 1972, most investigators seemed to favor an origin of certain Apollo 

14 samples with KREEP-rich compositions and textures indicating 

crystallization from melts as endogenic volcanic basalt ("Fra Mauro 

basalt"; fig. 2.8). The current consensus, however, is that almost all 

terra melt rocks with basaltlike textures originated as impact melts 

(Dence and Plant, 1972; Green and others, 1972; Taylor and others, 

1972; James, 1973; Irving, 1975,1977; McKay and others, 1978,1979). 

The main clues are (1) trace elements indicating contamination by 

meteorites and (2) small-scale textural and compositional heterogeneity, especially among clasts.

Although impact-created rocks dominate the terrae, some igneous rocks from the early pre-Nectarian crust were recovered (fig. 8.2; 

Warren and Wasson, 1977,1979b, 1980a, b; Norman and Ryder, 1979). 

A diligent search has been conducted for samples of these "pristine" 

rocks, which have also been referred to as meteorite-free, endogenic,
8. PRE-NECTARIAN SYSTEM 141

endogenous, indigenous, primitive, and plutonic (see references in 

Warren and Wasson, 1980a, p. 81-82). The following are the main 

criteria for pristinity that currently are most widely accepted.

1. A low content of siderophile elements, many of which are characteristic of iron meteorites and chondrites (for example, Hertogen 

and others, 1977; Morgan and others, 1977). The highly siderophilic 

elements Au, Ir, Os, Pd, and Re probably accompanied metallic Fe 

during early differentiation of Solar-System material and thus 

were concentrated in planetary cores and small Fe-rich bodies. A 

few investigators believe that the silicates of the Moon are intrinsically richer in these elements than is generally believed and that 

pristine rocks low in siderophiles are impact-melt rocks from 

which siderophiles have been removed by extraction of metal 

(Delano and Ringwood, 1978; Wanke and others, 1978). Most investigators, however, believe that high concentrations of these elements indicate meteoritic contamination, although their absence 

does not necessarily prove pristinity (Irving, 1975, 1977; Warren 

and Wasson, 1977; Anders, 1978; Norman and Ryder, 1979, p. 

533-636; Taylor, 1982, p. 222).

2. Igneous, especially cumulus, textures. Impact melts rarely cool 

slowly enough to develop cumulus textures, which are conspicuous 

in layered intrusions on the Earth, where such factors as gravity, 

liquid density, and fluid currents have produced quasi-sedimentary features (Jackson, 1971; Jackson and others, 1975; Raedeke 

and McCallum, 1980). Thus, lunar cumulus textures indicate pristinity (fig. 8.2D). Such textures furthermore suggest that the rock's 

composition arose by igneous processes. In contrast, metamorphic 

textures commonly form in deposits of impact melt and breccia 

whose compositions arose by mechanical mixing of preexisting

FIGURE 8.2. "Pristine" terra samples.

A. Cataclastic anorthosite (sample 60025), as photographed in Lunar Receiving Laboratory after arrival from Apollo 16 landing site. Little texture is evident. Scale cube is 1 cm 

on a side in this and all other "mug shots."

B. Thin section of sample 60025. Most optically observable crystals (modes) are of 

plagioclase. Texture is seriate. Crossed polarizers; field of view, about 2 mm.

rocks. Volcanic textures may form either in true volcanic rocks or 

in impact melts, and thus are ambiguous when not supplemented 

by other clues, such as siderophile content.

3. Grain size larger than 3 mm, indicative of slow cooling in a plutonic environment. In many pristine samples, however, the original igneous textures and the coarse grain size have been obscured 

by subsequent shock-induced cataclasis.

It seems clear today that the bulk compositions of most terra 

rocks are determined by impact-induced mixing. Plots of bulk majorelement compositions of lunar terra rocks reveal that individual 

endogenic cumulates and other pristine samples vary much more in 

composition than do obviously impact-generated polymict breccias 

and regoliths containing siderophiles (Warren and Wasson, 1977; 

1979b, p. 601; 1980b; Norman and Ryder, 1979, p. 554; Taylor, 1982, p. 

202). The "anorthositic norite" or "highland basalt" average terra 

composition, granulitic impactite (Warner and others, 1977; Bickel 

and Warner, 1978), very high alumina "basalt" (Dowty and others, 

1973,1974b; Hubbard and others, 1974), andlow-KKREEPare among 

the compositional types that fall in the middle of such plots away from 

known pristine rocks. Low-K KREEP is found mostly (Ryder and 

others, 1977) or entirely in siderophile-bearing fragment-laden melts

C. Coarse-grained granulitic troctolite (sample 76535) from North Massif, Apollo 17 

landing site. Major modes are plagioclase (58 percent) and olivine (37 percent), with 

accessory orthopyroxene (4 percent) (Dymek and others, 1975).

D. Thin section of pristine pink spinel troctolite clast (sample 67435,14). Original igneous 

cumulate texture is largely preserved (Prinz and others, 1973). Euhedral to subhedral 

olivine (gray) is enclosed poikilitically by plagioclase (white). Plane-polarized light; field 

of view, about 2 mm.
142 THE GEOLOGIC HISTORY OF THE MOON

(Meyer, 1977; Hess and others, 1977; McKay and Weill, 1977; Reid and 

others, 1977; Shih, 1977; Warren and Wasson, 1979b). This extensive 

impact mixing is to be expected from the geologic style of the Moon 

because each new impact mixes and melts rocks, most of which 

already are complex breccia deposits with compositions far removed 

from those of their endogenic origins. New hybrid compositions are 

formed by each event, and continued impacting creates a hybrid 

average of the original pristine compositions. Mixing models attempt 

to reconstruct the observed compositions of regoliths and bedrockbreccia deposits from a set of end members believed to represent 

original magma types (for example, Hubbard and Gast, 1971; Schonfeld and Meyer, 1972; see recent critiques by Ryder, 1979, and Taylor, 

1982, p. 201-205).

In summary, the composition of the lunar terra crust was derived 

by mixing pristine ANT compositions with smaller amounts of highly 

differentiated KREEP-rich material. Although a few intact fragments 

of the early crustal rocks were recovered, the textures and structures 

of almost all returned terra samples were created by impacts. Many of 

the samples assembled by impacts retain the early endogenic compositions, whereas others have been compositionally as well as texturally altered during the impact reworking. Problems still persist in 

identifying the sources of specific samples and determining the relative amounts of pristine and impact-mixed materials. However, 

enough clues about the early crust have been assembled from this 

obscure record to reconstruct the following partial plutonic history of 

the early crust.

EARLY CRUSTAL PETROGENESIS

Magma ocean

Analysis of the small fragments from the Apollo 11 regolith led to 

what is still the leading model for the origin of the early crust (Smith 

and others, 1970; Wood and others, 1970). The terra crust and the 

mantle source of mare-basalt magmas are widely thought to have 

separated from a very large volume of magma popularly called a 

magma ocean (Walker and others, 1975; Wood, 1975b). The first hint of 

an early widespread differentiation of primitive materials came from 

REE abundances. Europium is overabundant in "average" terra rocks 

and underabundant in mare basalt relative to chondritic REE abundances (for example, Taylor, G.J., and others, 1973; Smith, 1974; 

Taylor, 1975, p. 154-159). This relation is consistent with an efficient 

mechanical separation of the ultramafic mare-source material, as 

cumulates, from a magma from which large amounts of plagioclase 

had already crystallized and extracted Eu relative to the other REE's. 

Plagioclase has large lattice sites that accept the large divalent 

(reduced) ion that only Eu among the REE's possesses.

"Oceanic" depths of 200 to 400 km are commonly favored, 

although lesser or greater depths are not excluded (half the lunar 

volume lies above, and half below, 360-km depth) (Walker and others, 

1975; Hubbard and Minear, 1975; Herbert and others, 1977; Taylor, 

1978; Longhi, 1980; Binder, 1982). Because of its lesser density, the 

plagioclase-rich terra-crustal material presumably floated in the 

magma ocean, whereas ultramafic cumulus crystals sank to form the 

upper mantle that became the mare source (for example, Smith and 

others, 1970; Wood and others, 1970; Taylor and Jakes, 1974; Taylor, 

1975,1978; Wood, 1975b, c; Herbert and others, 1977; Ryder and Wood, 

1977; Longhi and Boudreau, 1979; Warren and Wasson, 1979a, 1980a). 

The divide between the sunken and floated materials is generally 

assumed to be the seismic discontinuity at the base of the crust. If the 

ocean was 360 km deep and the crust is 75 km thick, the crustal 

material constituted about a quarter of the original oceanic volume.

The origin of KREEP in general and its role in "magma 

oceanography" in particular have generated a very large literature 

because of its persistent appearance in lunar samples and its exotic 

highly differentiated chemistry (see reviews by Taylor, 1975, 1982; 

Irving, 1977; Meyer, 1977; Warren and Wasson, 1979c; James, 1980; 

Vaniman and Papike, 1980). Its origin has been ascribed to partial 

melting of ANT material or, as commonly in more recent literature, to 

a late stage in fractionation of the magma ocean. KREEP has been 

considered to be derived from a minor liquid accumulated in dispersed niches throughout the magma ocean or the crust, or, alternatively, concentrated in certain regions. The question of its origin is 

difficult to resolve because because most or all analyzed KREEP is a 

chemical constituent incorporated into other rocks, rather than a 

distinct rock type itself.

Multiple plutons

Study of pristine rocks, many of which are thought to be relics of 

the primordial crustal differentiation, has shown that the magmaocean model outlined above is too simple. Many of the pristine crustal 

rocks, especially those from the Apollo 17 landing site, are cumulate 

norite and troctolite that apparently formed by settling in their 

parent magmas (Norman and Ryder, 1979, p. 538). Data from the 

orbital geochemical instruments, as interpreted by mixing models, 

indicate that the topmost part of the "average" crust can contain no 

more than about 60 percent anorthosite (Haskin and Korotev, 1981) 

and that compositions range laterally from highly feldspathic to 

highly mafic (Spudis and Hawke, 1981). Moreover, the orbital instruments covered only a small percentage of the Moon and are subject to 

calibration problems (see chap. 5). Therefore, mafic materials may be 

more abundant than once thought.

Compositional subtleties of the mafic and anorthositic ANT 

rocks suggest that more than one magma system may have built the 

terra crust. For example, if the Mg/(Mg + Fe) ratios in mafic minerals 

(olivine, pyroxene) in the pristine rocks are plotted against the 

Ca/(Ca + Na + K) ratios in plagioclase, two distinct trends are definable (fig. 8.3; Steele and Smith, 1973; Taylor, G.J., and others, 1973; 

Warner and others, 1976; Warren and Wasson, 1977,1979a, b, 1980a; 

Wood, 1977; Binder, 1980; James, 1980; Longhi, 1980; Raedeke and 

McCallum, 1980; Taylor, 1982, p. 248-251). One trend consists of 

highly feldspathic rocks (anorthosite and troctolitic anorthosite) that 

contain relatively calcic, uniform plagioclase and relatively Fe-rich, 

compositionally varying mafic minerals; these rocks are termed ferroan anorthosite because of the Fe-rich mafic minerals (Dowty and 

others, 1974a). The other trend consists of more mafic rocks (troctolite, spinel-bearing troctolite, norite, gabbronorite, and a few ultramafic rocks) that have highly varying plagioclase and relatively magnesian, compositionally varying mafic minerals; these rocks are 

referred to as Mg-rich plutonic rocks or the Mg suite. In an individual 

rock of the Mg suite, Mg-rich mafic minerals are associated with Carich plagioclase, and Fe-rich mafic minerals with more sodic plagioclase. These associations are what is to be expected in typical 

cotectic igneous differentiation (Bowen, 1928). The anorthosite was at 

first thought to be more anomalous (McCallum and others, 1975)

95

90

85

so

01 Opl 

a Very probably or definitely pristine

_ v T Probably pristine

o

LL

< 75

65

60

55

50 -

Mg-rich trend

85 90 95

An CONTENT OF PLAGIOCLASE, IN MOLE PERCENT

100

FIGURE 8.3. Ferroan-anorthosite and Mg-rich trends in ANT rocks, separated by a 

distinct hiatus. Examples shown are for olivine (ol) and orthopyroxene (opx) in pristine rocks 

(Warren and Wasson, I980a>
8. PRE-NECTARIAN SYSTEM 143

because Fe-rich mafic minerals are associated with Ca-rich plagioclase (Wood, 1975b; 1977, p. 48). However, the same type of mineralogic variation has been discovered in anorthosite of the Stillwater 

layered igneous complex (Raedeke and McCallum, 1980). No simple 

magma system is likely to have generated both trends simultaneously 

(Wood, 1975b; Longhi and Boudreau, 1979; Warren and Wasson, 

1979a, b, 1980a, b; Raedeke and McCallum, 1980).

To accommodate the intricacies and the global scale that both 

seem required, several authors have proposed "separate pluton" models, in which layered Mg-suite plutons (Jackson and others, 1975; 

Raedeke and McCallum, 1980) intruded from below into an alreadyfloated crust of ferroan anorthosite (Binder, 1980; James, 1980; Warren and Wasson, 1980a). Alternatively, the early pre-Nectarian crust 

may never have been an "ocean" but only a system of magma "pods" or 

"ponds" (Wetherill, 1975b, 1976; Simonds, 1979). Complex petrologic 

schemes dating from the accretionary period (Simonds, 1979; Smith, 

1982) or operating in an early convecting magma ocean (Longhi, 1980) 

have also been proposed. Impacts must have bombarded the primitive 

Moon catastrophically; they would have pierced any solid crust that 

formed on an "ocean" and have stirred and mixed the liquid, influencing the early geochemistry and petrogenesis (Wetherill, 1975b; Wood, 

1975b; Hartmann, 1980; Minear, 1980; Simonds, 1979; Warren and 

Wasson, 1980a; Taylor, 1982, p. 248-251). Although the magma-ocean 

concept is still favored by most petrologists and geochemists, it is 

becoming increasingly modified and complex.

Crustal zonation

Possible layering or, more likely, crude vertical zonation of the 

crustal material has been studied by a combination of seismic, 

petrologic, and impact investigations. Seismic data indicate a discontinuity in the crust at about 20- or 25-km depth (Toksoz, 1974; Toksbz 

and others, 1974; Goins and others, 1979) that is commonly ascribed to 

an interface between low-density, relatively porous crustal rock and 

denser sealed material of the same composition (Todd and others, 

1973; Bills and Ferrari, 1977; Haines and Metzger, 1980). Alternatively, or additionally, this seismic feature could represent a compositional discontinuity between more feldspathic material above 

and more mafic material below (Herzberg and Baker, 1980; Warren 

and Wasson, 1980a). A distinct KREEP-rich layer has also been suggested, either between layers of anorthosite and more mafic material 

(Ryder and Wood, 1977) or at the base of the crust (Warren and 

Wasson, 1979c).

Because basins probe the subsurface, the surface distribution of 

ejecta compositions should help in reconstructing any layering or 

zonation that exists. The large size of basins might suggest excavation 

as deep as the mantle. Mineral assemblages in some spinel-bearing 

Mg-suite rocks have been cited as requiring pressures available only 

in the deep crust or mantle of the Moon (Herzberg, 1978; Herzberg 

and Baker, 1980). However, alternative phase-equilibria schemes 

allow an origin as shallow as 25 km, well within the crust (Warren and 

Wasson, 1979a; Herzberg and Baker, 1980). Although most terrestrial 

dunite is probably of mantle origin, at least one Mg-suite lunar dunite 

probably originated in the crust, judging from its association with, 

and mineralogic similarity to, troctolite sample 76535, whose cooling 

and reequilibration depths were probably in the range 10-30 km 

(Gooley and others, 1974; Dymek and others, 1975). Orbital geochemistry suggests crustal composition even for Imbrium-basin 

ejecta, which I believe was derived from the cavity of the Moon's 

largest known basin, Procellarum (fig. 8.4). Therefore, little, if any, 

mantle material was brought to the surface in impact ejecta, and 

shallow basin excavations restricted to the crust are implied.

Crustal thickness and surface composition seem to be related; 

the more aluminous the material, the thicker the crust (Schonfeld, 

1977; Haines and Metzger, 1980). That is, the Pratt (unequal density) 

model of isostasy seems to be more nearly applicable to the terrae 

than is the Airy model (Solomon, 1978). The Airy (equal density) 

model was assumed in arriving at the commonly stated conclusion 

that the farside crust is thicker than the nearside crust (Kaula and 

others, 1974). Most investigators have nevertheless concluded that 

the farside crust is thicker (see chap. 1), even under Pratt or mixed 

isostatic models (Haines and Metzger, 1980).

This difference in crustal composition and thickness, however, 

may not be hemispheric in extent. The orbital geochemistry (Haines

and Metzger, 1980) suggests that the crust is both thinner and more 

magnesian beneath the two largest lunar basins, Procellarum and 

South Pole-Aitken, than in other regions where data are available. 

Seismic data (see chap. 1) suggest that the crust is 50 or 60 km thick 

under southern Oceanus Procellarum and about 75 km thick under 

the highly feldspathic Apollo 16 landing site, which is near the Procellarum rim (pi. 3). Crustal thicknesses are similar on the southern 

nearside and western (long 90°-180° E.) farside (fig. 1.12; Bills and 

Ferrari, 1977). The X-ray spectrometer detected higher Al/Mg ratios 

on the farside and on parts of the eastern and southeastern nearside 

than in the Imbrium-Procellarum region (Adler and others, 1972, 

1973; Hubbard and others, 1978). No data known to me contradict the 

conclusions that the thick aluminous crust is restricted to the rims 

and exteriors of the two giant basins and that the thinner, more 

magnesian crust is concentrated in basin interiors. Additionally, the 

gamma-ray spectrometer detected the high radioactivity characteristic of KREEP in the Imbrium-central Procellarum region 

(Metzger and others, 1973,1977), where I would expect the thinnest 

part of the lunar crust.

The sample data support the conclusion that the giant basins 

control the crustal thickness and, thus, the compositional zone that 

was accessible to later impacts (fig. 8.4). The greatest depths were 

reached by the Imbrium basin, whose ejecta (Fra Mauro Formation) is 

correspondingly the lunar deposit richest in KREEP (see chap. 10). 

Mg-suite rocks and low-K KREEP, probably derived from the same or 

shallower depths, apparently compose most of the Imbrium (Apollo 

15; chap. 10) and Serenitatis (Apollo 17; chap. 9) rims. The Apollo 16 

and Luna 20 landing sites are the richest in feldspathic ANT materials (chap. 9), as befits their derivation from the thick crust of the 

outer Procellarum shelf or exterior terrain.

This picture, though internally consistent, is subject to revision 

or rejection on the basis of better data. Lateral zonation of primordial 

origin also is possible. Only four major sample suites from four 

landing sites have been collected from the terrae, and each suite 

added complicating detail to the lunar petrogenetic picture. On the 

basis of the complexities generally uncovered when geologic systems 

are examined closely, fine-scale local intrusions only crudely 

organized into petrologic provinces might be expected. The pattern 

that seems to reflect global layers exposed by large basins might be 

fortuitous.

BASIN AND CRATER AAATERIALS

Introduction

The severely deformed, chemically and texturally mixed, siderophile-rich materials found in the Nectarian and Imbrian deposits 

attest to a high rate of impacts during pre-Nectarian time as well. 

More specific information on the characteristics of individual preNectarian units is sparse because none were sampled directly and 

because the chemical and petrologic criteria for recognizing such 

units are ambiguous. Attempts have been made to identify individual 

impacting projectiles with specific basins by means of distinctive 

siderophile signatures (Hertogen and others, 1977; Morgan and others, 1977), but the success of these attempts is limited by knowledge of 

the prebasin stratigraphy. The study of pre-Nectarian impact deposits thus requires stratigraphic analysis.

Originally, pre-Nectarian and Nectarian were combined as preImbrian (Shoemaker and Hackman, 1962). This designation sufficed 

for most old units mapped at 1:1,000,000 scale because the mapping 

was confined to the central nearside, which is dominated by the 

Imbrium basin (pi. 3). If they were subdivided at all, pre-Imbrian 

craters were numbered according to the morphologic criteria established by Pohn and Offield (1970).

The desirability of stratigraphically dividing the pre-Imbrian 

emerged during mapping of the nearside regions farthest from 

Imbrium. While mapping south of Mare Nectaris, Stuart-Alexander 

(1971) recognized a lineate surface as the expression of Nectaris-basin 

deposits and secondary-impact craters similar to those of the Orientale and Imbrium basins (fig. 8.5); she named the deposit the "Janssen Formation" (Stuart-Alexander, 1971). At first, the Janssen was 

used to subdivide pre-Imbrian crater materials into numbered preand post-Nectaris categories descendant from Pohn and Offield's
IMBRIUM < 

Present surface ? f '

Post-Procellarum surface

KREEP-rich basalt

SERENITATIS CRISIUM -3 

A'

FIGURE 8.4. Hypothetical crustal layers intersected by impact basins.

A. Index map based on plate 3, showing basins from which samples were obtained. PI, P2, 

P3, rings of Procellarum basin (dotted where inside younger basins; inner rings dashed). 

Shading, Mg-rich terrane after formation of Procellarum and before formation of other 

basins.

B. Diagrammatic geologic fence diagram, suggesting crustal and mantle configuration after 

Procellarum impact, subsequent uplift, and partial fill by hypothetical KREEP-rich 

volcanic basalt (black). Three gradational crustal layers are intersected by Procellarumbasin surface: upper Al-rich material is exposed outside basin and on outermost shelf

(between PI and P2); Mg-rich material is exposed on intermediate shelf (between P2 

and P3) and in central basin; and lower-crustal KREEP-rich material is near surface (or 

may be exposed) near center of basin. Nectaris and Crisium basins were formed mostly in 

Al-rich terrane (Apollo 16 and Luna 20 samples, respectively); Serenitatis basin was 

formed in Mg-rich terrane (Apollo 17 samples); Imbrium basin was formed in Mg-rich 

terrane (Apollo 15 samples) and also ejected large amounts of KREEP-rich material 

and (or) KREEP-rich basalt (Apollo 14 samples). Arrows indicate paths of ejecta. 

Curvature of surface trace and of the Moon is ignored; vertical exaggeration is arbitrary.
8. PRE-NECTARIAN SYSTEM 145

numeric designations (Stuart-Alexander, 1971; Wilhelms and 

McCauley, 1971; Scott, 1972b). In 1975, the pre-Imbrian was formally 

divided into the pre-Nectarian and Nectarian (Stuart-Alexander and 

Wilhelms, 1975). This division is shown on all subsequently published 

geologic maps, where appropriate, and is used here. The utility of this 

subdivision was proved as the limb, farside, and polar regions were 

mapped further and large numbers of pre-Nectarian and Nectarian 

units were identified (Scott and others, 1977; Wilhelms and El-Baz, 

1977; Lucchitta, 1978; Stuart-Alexander, 1978; Wilhelms and others, 

1979). The term "pre-Imbrian," however, is still used where more 

convenient, as in places on the nearside where the finer distinction 

cannot easily be made. Because neither "pre-Nectarian" nor "preImbrian" has a formally defined base, such terms as "system" and 

"period" used in conjunction with them are informal and not capitalized.

Recognition criteria

A start in establishing a pre-Nectarian age is to examine terrain 

covered by the deposits and secondary craters of the Nectaris basin 

(fig. 8.5). Deposits overlain by these materials are, by definition, preNectarian. Pre-Nectarian basins in direct contact with Nectarisbasin materials include Australe, Fecunditatis, Mutus-Vlacq, Tranquillitatis, and Werner-Airy (pi. 6; table 8.2). Craters lightly nicked by 

Nectaris secondaries but not deeply buried by deposits are type 

examples of pre-Nectarian craters (figs. 8.5C, D\

Units not in contact with Nectaris-basin materials can potentially be dated by comparing the size-frequency distributions of 

superposed primary craters with those of craters superposed on Nectaris (fig. 8.6; table 8.2). Size-frequency curves lying entirely above 

that for Nectaris demonstrate a pre-Nectarian age. One curve for a 

basin here considered to be pre-Nectarian (Apollo) crosses below the 

Nectaris curve, but only at the statistically weak end representing 

the frequencies of large craters.

The differences among the many curves demonstrate the fact 

that the lunar terrae do not have some single characteristic sizefrequency distribution but are composed of basin blankets, each of 

which has reset the visible crater population of the surfaces it covers. 

The slopes of the cumulative crater-frequency curves decrease and 

become more highly curved from the youngest to the oldest surfaces 

(fig. 8.6). Some of the crater-frequency curves approach a slope of 1 

for craters as large as 60 or 70 km in diameter, a slope substantially 

lower than that of 2 or 1.8 which is commonly believed to represent the production frequency for younger craters (for example, 

Basaltic Volcanism Study Project, 1981, chap. 8). Although the -1 

slope has not been modeled in studies of crater saturation (Woronow, 

1977,1978), it nevertheless probably results from mutual obliteration 

of pre-Nectarian and Nectarian craters.

Several factors in addition to this destruction of craters limit the 

counts to craters at least 20 km in diameter. Meaningful counts are 

restricted to (1) superposed craters and (2) primary craters. A superposed crater useful in dating the basin, and a subjacent crater whose 

inclusion in the counts would be erroneous, may differ only slightly in 

appearance. Numbers of secondary craters generally do not indicate 

time of exposure because such craters are generated in bursts when 

their parent primaries form. The superposed-buried and primarysecondary distinctions can be made only in well exposed, well photographed terrains. Some counts of pre-Imbrian craters are inevitably restricted to small areas because of burial by later deposits; the 

buried craters are visible but are useless for dating (fig. 8.7). Photographic quality on the farside is inadequate for interpreting craters 

smaller than 20 km in diameter on many basins, especially at latitudes poleward of about 40°.

The results of these natural and artificial shortcomings are indistinct separations of the curves for old basins and statistically poor 

curves for the smallest and the youngest (least cratered) units. No 

individual craters are dated here by crater densities except the 265- 

km-diameter crater-basin transitional feature Milne (figs. 4.2J, 8.11). 

For most pre-Nectarian craters, the small superposed craters that are 

required for statistically significant counts have been erased by 

mutual obliteration and degradation on slopes. Several counts of 

superposed craters are listed in table 8.2 as "poor," meaning that they 

contain few total craters or lack small craters because of burial. These 

counts are not plotted in figure 8.6 except for the marginally useful 

one for the Coulomb-Sarton basin.

Pre-Nectarian surfaces may be heavily or only moderately cratered, depending on age and happenstance position relative to younger secondary clusters. Although densities of superposed secondary 

craters are almost useless for dating, clustered secondaries are very 

useful for establishing relative age if they can be traced to their source 

basins (fig. 8.7A). Once a basin is established as pre-Nectarian, superposition of its secondary clusters on another basin or crater unit 

establishes the pre-Nectarian age of that unit as well. Where available, superpositions prevail over crater counts as age criteria. Secondary clusters can also be used to date small units, such as those of 

individual pre-Imbrian craters that are inaccessible to size-frequency 

studies. Numerous stratigraphic relations between pre-Nectarian 

and younger basins and between pre-Nectarian basins have been 

established by such superpositional relations (table 8.2).

Basins

This volume divides 30 pre-Nectarian basins into 9 age groups 

(table 8.2). Each group is headed by one basin whose relative age 

seems to be well established by crater densities or superpositional 

relations. Additional basins are tentatively placed in the groups on 

weaker grounds.

The first group consists of the giant South Pole-Aitken and 

Procellarum basins. South Pole-Aitken is characterized by saturation 

of superposed impact structures as large as basin size (fig. 8.8; table 

8.2). The history of its discovery illustrates the need for mapping large 

areas. When Hartmann and Kuiper (1962) realized that basin rings 

are prevalent on the Moon, they predicted the presence of a large 

farside basin on the basis of telescopic photographs showing huge 

mountains near the Moon's south pole. This prediction went 

unnoticed when Soviet Zond photographic altimetry and Apollo laser 

altimetry revealed a large depression as much as 5 to 7 km below the 

average farside surface (see review by Stuart-Alexander, 1978). Some 

massifs had been photographed by Apollo 8, but their connection with 

a basin was not immediately perceived (fig. 8.9; Wilhelms and others, 

1969). Wilhelms and El-Baz (1977) mapped other massifs without 

knowing of their connection with the basin. Several observers then 

independently discovered that mountains over a huge area of the 

southern farside and southernmost nearside form a giant ring (pis. 3, 

6; fig. 8.8; Howard and others, 1974; Schultz, 1976b, p. 306; StuartAlexander, 1978; E.A. Whitaker, written commun., 1976). Now, there 

is little doubt about the existence of this "Big Backside Basin," 

although the number and exact diameters of the rings are uncertain 

(Wilhelms and others, 1979).

Regional studies also suggest the presence of the Procellarum 

(Whitaker, 1981) or "Gargantuan" (Cadogan, 1974,1981) basin. Some 

of the profound effects of this 3,200-km-diameter basin (pi. 3; fig. 8.10; 

Whitaker, 1981) on the later geologic evolution of the Moon have 

already been described. (1) Its excavation thinned the crust considerably, possibly by 50 km, and the mantle rose in isostatic compensation 

(figs. 6.22, 8.4). (2) As a result, the terra surface was lower, and the 

elastic lithosphere thinner or weaker, during mare extrusion and 

structural deformation in the Procellarum region than elsewhere on 

the Moon. (3) Their relation to Procellarum accounts for numerous 

concentric ridges and troughs of the terrae, which, like the mountains 

now known to be parts of South Pole-Aitken, would be anomalous if 

not part of a basin (fig. 8.10). (4) Redistribution of crustal materials by 

the Procellarum impact may be an important factor in the distribution of anorthositic, Mg-suite, and KREEP-rich terra materials (fig. 

8.4). (5) More speculatively, Procellarum ejecta and, possibly, ejecta of 

the South Pole-Aitken basin may account for the suggested thickening of the farside crust (Cadogan, 1974; proposed previously in somewhat different form by Wood, 1973). (6) The Procellarum basin is the 

reason for the existence and circular shape of Oceanus Procellarum 

(pi. 4; fig. 5.26), which otherwise would have had to form by a special 

"continental" type of mare-basalt volcanism otherwise unknown on 

the Moon. (7) Other maria, contained in basins superposed on Procellarum, are more numerous, larger, and deeper than maria in 

equally large or larger basins outside Procellarum (pi. 4). (8) Isostatic 

compensation has gone further inside Procellarum than on the 

thicker lithosphere outside. (9) More uplifted crater floors lie in and 

near Procellarum than elsewhere (pi. 5). (10) Arcuate rilles, opened by 

mare subsidence, are abundant inside but rare outside the basin (pi. 

5); straight rilles may also be controlled by this largest lunar feature 

(see chap. 6). Additional effects of Procellarum on mare volcanism are 

described in chapters 11 through 13.
146 THE GEOLOGIC HISTORY OF THE MOON

Age group 2 (table 8.2) rather arbitrarily consists of eight basins 

whose existence is uncertain; their collection into a single group is 

based on the assumption that if these basins were younger (and if 

they exist), they would be more obvious. The -1 slope of the 

cumulative size-frequency distribution of the numerous craters 

counted on Al-Khwarizmi/King (figs. 8.6, 8.11) characterizes that 

ancient basin. If this basin does not exist, the -1 slope characterizes 

other ancient lunar terrane, possibly South Pole-Aitken deposits, 

which probably extend into this region. Two of the age-group 2 basins, 

Flamsteed-Billy and Insularum, were identified on the basis of circular patterns of terra islands in the western maria (pi. 3; fig. 8.10; 

Wilhelms and McCauley, 1971). This criterion for basin identification 

arose from the discovery that seemingly isolated elevations generally 

are parts of basin rings.

The basins of age groups 3 and 4 are also so heavily degraded that 

they have lost most traces of ejecta, and superposed and prebasin 

craters are difficult to distinguish in their peripheries. Nevertheless, 

they are still likely to be true basins. Saturation by large craters 

appears to affect the size-frequency plots significantly; the slopes of 

some curves for craters believed to be superposed approach 1. 

Because of this saturation or deep burial by later basin or mare 

deposits, ranking of basins by age within one of these older groups is 

not considered feasible, even where crater frequencies appear to 

differ substantially (for example, Lomonosov-Fleming versus MutusVlacq).

Identification of Mutus-Vlacq and Lomonosov-Fleming as basins 

is another result of mapping with the impact model in mind. Extensive tracts of "pitted plains" of pre-Imbrian age are concentrated 

southwest of the Nectaris basin (pi. 7; Wilhelms and McCauley, 1971; 

Scott, 1972b). These plains are centered within a partial circle of 

massifs and ridges (fig. 8.12; Wilhelms and others, 1979). This concentration of plains suggests the presence of a basin, Mutus-Vlacq, 

regardless of the plains' origin (see chap. 9). Some of the Mutus-Vlacq 

massifs had been identified in earlier mapping but were considered to 

be "volcanic domes" (Wilhelms and McCauley, 1971; Mutch and Saunders, 1972) before the Apollo 16 results downgraded the significance of 

terra volcanism. Similarly, Lomonosov-Fleming was identified by 

mapping mounds, ridges, and scarps that form a circle 620 km across 

and that encompass extensive light-colored plains (fig. 8.13; Wilhelms 

and El-Baz, 1977). Although no interior ring is observed in either of 

these basins, the term "ringed basin" is appropriate because the 

presence of another ring or rings beneath the plains fill is predictable 

by analogy with better exposed basins.

Although the indistinct Nubium, Fecunditatis, and Tranquillitatis basins cannot be accurately dated by superposed craters, 

they are subjectively ranked with the other age-group 3 basins on the 

supposition that they would be equally conspicuous if not flooded by 

mare materials. Fecunditatis probably has at least two exposed and 

one buried ring and thus has the most elaborate visible ring structure 

of the poorly observed basins listed in table 4.2. The 680-km-diameter 

ring, overlain by the young fractured-floor crater Taruntius 

(Wilhelms, 1972b), is the most conspicuous, but another ring, about 

980 km in diameter, is probably the main topographic rim.

Age group 4 consists of three well-established basins. The diameters, numbers of rings, and age relations of Keeler-Heaviside and 

Ingenii, however, are somewhat uncertain because several nearby 

massifs may belong to the South Pole-Aitken basin (figs. 1.4, 8.7).

Although the two basins assigned to group 5 are poorly photographed or obscured by younger deposits, their deposits apparently 

have diminished the crater densities within a basin radius or diameter of the rim, as expected around basins. Smythii contains a distinct 

mare (fig. 8.11) and a mascon (table 6.1) but has only indistinct rings. 

Coulomb-Sarton (figs. 4.3M, 8.14) was discovered shortly after photography of the farside by the Lunar Orbiter spacecraft and was 

designated "unnamed basin B" (Hartmann and Wood, 1971). The rings 

are vague, and the identity of the topographic rim uncertain (table 

4.1). Wood and Head (1976) withdrew the basin from the list of lunar 

basins. Massifs and abundant interior plains do exist, however, and 

the basin is here considered to be definitely established (Lucchitta, 

1978).

FIGURE 8.5. Boundary between pre-Nectarian units and Nectarian System southeast of Nectaris basin (Stuart-Alexander, 1971; Stuart-Alexander and Wilhelms, 1975; Wilhelms and El-Baz, 

1977; Wilhelms and others, 1979). 

A. Regional view showing concentrations of textured deposits and secondary craters of Nectaris basin (hachured black outlines) and areas of B through D (white outlines). Orbiter 4 frame

M-52.
8. PRE-NECTARIAN SYSTEM 147

fi. Poorly exposed pre-Nectarian deposits. Part of Vallis Rheita (VR), a secondary chain 

of Nectaris basin, overlain by Nectarian craters Rheita (R; 70 km, 37° S., 42° E.) and 

Young D (YD; 46 km). Other Nectaris-basin secondary chains (Nsc) are overlain by 

Nectarian primary crater Stemheil (S; 67 km). Crater Metius (M; 88 km) also is 

probably Nectarian because it interrupts pattern of Nectaris-basin deposits. Fabricius 

(F; 78 km) is much younger, probably Eratosthenian. Orbiter 4 frame H 71.

C. Moderately well exposed pre-Nectarian craters, including Vega (V; 76 km, 45° S., 63° 

E.), Fraunhofer E (FE; 42 km), Fraunhofer J (FJ; 63 km), Reimarus R (RR; 35 km), 

and Peirescius (P; 62 km), overlain by circular and subcircular Nectaris-basin secondary 

craters (Nsc). Vega and Peirescius are good type examples of pre-Nectarian crater 

morphology because large rim segments are preserved. Orbiter 4 frame H 52.

The well-dated basins of the four youngest age groups (table 8.2) 

have observable exterior deposits and secondary craters where well 

exposed and well photographed. The craterlike rims of Birkhoff and 

Lorentz are surrounded by ejecta and secondary craters to distances 

of at least one basin diameter (fig. 8.14). Secondary craters of Apollo 

are visible in the moderately well photographed zone south of the 

basin (fig. 8.15; Wilhelms and others, 1979), although they are covered 

by younger materials in the north and northeast. Apollo, Birkhoff, 

and the large crater or small basin Milne were considered to be 

Nectarian because of their evident ejecta (Wilhelms and El-Baz, 1977; 

Lucchitta, 1978; Wilhelms and others, 1979); the crater counts, 

however, indicate a pre-Nectarian age (as proposed for Apollo by 

Stuart-Alexander, 1978). Grimaldi ejecta appears to have smoothed 

its periphery, although it is too heavily obscured by Orientale deposits 

to demonstrate textured ejecta (figs. 4.37, 4.477; McCauley, 1973). The 

Planck and Schiller-Zucchius basins similarly lack textured ejecta 

but have smoothed the surrounding terrain and reduced the crater 

densities.

Freundlich-Sharonov is an especially instructive basin. The 

basin rings and interior are poorly photographed (fig. 4.3P), and their 

pattern emerged only upon mapping (Stuart-Alexander, 1978). 

Stuart-Alexander (1978) noted that superposed craters are relatively 

sparse in the surrounding terrain. A good photograph of the south 

basin periphery reveals relatively uncratered terrain lineate radially 

to the basin (fig. 8.7A). Analogy with Orientale indicates that the 

lineations are not tectonic or volcanic but were formed by the deposition of Freundlich-Sharonov ejecta. The basin would probably not 

have been discovered or interpreted as a typical impact basin without 

knowledge of other basins.

The progress of basin discovery suggests that additional preNectarian basins will be found. Moreover, many basins probably 

formed before the South Pole-Aitken and Procellarum basins but are 

now completely obliterated. Basin deposits probably covered all parts 

of the Moon many times over in the manner displayed by the visible 

basins.

Craters

In the absence of reliable crater counts, morphologic comparisons are used to date stratigraphically isolated craters. These

 : 1 \.*Hir. X 1 >rm

 ' &

D. Well-exposed pre-Nectarian terrane, lightly pitted by Nectaris-basin secondary craters, 

affording additional typical pre-Nectarian morphologies. Subdued pits are superposed 

on pre-Nectarian craters Nearch (N; 76 km, 58.5° S., 39° E), Rosenberger (R; 96 

km), Rosenberger C (RC; 47 km), Rosenberger B (RB; 33 km), and probably Vlacq 

(V; 89 km). Well-defined part of Nectaris-basin secondary field is in upper right (Nsc). 

Orbiter 4 frame H-83.
148 THE GEOLOGIC HISTORY OF THE MOON

TABLE 8.2. Pre-Nectarian basins

[Age yroups in order of increasing age, each group arbitrarily Headed by the largest basin of toe yroup. Diameters from table 4.1. Parenthetical Qasins, 

existence not definitely established; parenthetical areas, buried; parenthetical densities, poor sample. Superposed craters are at least 20 km in diameter 

unless otherwise noted. Underlying and overlying basins determined from directly observed superpositional relations; n.d., no relations detected]

Aye 

group

9

8

7

6

5

4

3

2

1

Superposed craters

uiasnctcr ........... ... 

3aSln (Km) Apea Density 

Number ,,,,£ I 2 \ (craters/ 

(l ° km } 106 Km 2)

Apollo 505 57 0.480 119 

Grimaldi 430 15 (0.154) (97)

Freundlicn- 600 81 .629 129 

Sharonov.

BirKhoff 330 50 .401 127

(>21 Km)

Planck 325 9 .082 (110)

Schiller- 325 16 (0.143) (112) 

Zucchius.

(Arnundsen- 355 11 (0.102) (108) 

Ganswindt) .

Lorentz 360 33 .208 159

Coulomb-Sarton 530 43 (0.296) (145)

Keeler-Heaviside-- 780 69 .371 186 

Poincare 340 32 (0.168, (190)

partly)

Inyenii 650 37 .228 162 

(>22 km)

Lomonosov- 620 63 .356 177 

Fleming. (>23 km)

(Mutus-Vlacq) 700 80 .336 225

Tranquillitatis 800 

Australe 880 129 (0.608, (>212)

partly)

(Al-Khwarizmi/ 590 63 .320 197 

King)

(Pinyre-Hausen) 300 --- 

(Balmer-Kapteyn)-- 550

(Flamsteed- 570 

Billy).

(Grissom-White) 600

(Tsiolkovskiy- 700 

Stark) .

South Pole- 2,500 

Ait Ken.

(Procellarum)----- 3,200 - 

Underlying basin

Maine Figure

Lorentz, Coulomb- 8.14 

Sarton.

South Pole-Aitken(?)-- 1.5

8.7

n.d 

South Pole-Aitken(?) 1.5 

Tranquillitatis -- 

n.d 

n d- -------------

n.d 

n.d ------------------

Overlying oasin

Name

Oriental e, Hertzsprung, Korolev.

Orientale, Imbrium----

Orientale, Irnbrium, 

Hertzsprung, Bi rkhoff, Lorentz.

Freund lich- Sharonov--- 

Schrodinger, PlanckKeel er-Heaviside(?) -

Imbrium, Crisium, 

Humboldtianum.

Crisium-- -- -

Orientale, MendelRydberg.

Orientale, Imbrium -

Crisium, Humboldtianum.

Hertzsprung, Apollo---

Orientale, MendelRydberg. 

Schrodinger, Planck, 

Poincare, Australe, 

( Sikorsky-Rittenhouse), (AmundsenGanswindt) . 

Hertzsprung, Korolev, 

Apol lo.

Keeler-Heaviside, 

Ingenii .

Most nearside 8 

basins.

Fi gjre

9.21 

1.9

8.7

9.21

8.14 

.14, 9.25

1.5

1.9, 7.6 

9.24

1.5

10.2 

8.14

1.5 

8.14

1.4 

8.7

1.5 

8.7

8.13

.8, 8.10 

8.10

8.12

11.1

.7, 5.17

.2, 11.1

1.5 

.5, 9.5, 

9.6

9.5 

9.2

8.11 

1.9

9.2 

1.5 

8.10

9.4

8.10 

8.15

1.3 

1.4

1.9 

1.5

9.21

9.6 

8.7

8.15 

.10, 10.2

Remarks

 

Small craters 

buried.

Oo. 

Deeply buried.

Small craters 

buried.

Small craters 

buried.

Small craters 

obscured.

Do. 

Deeply buried.

Deeply buried.

Complex burial 

pattern.

Deeply buried.
8. PRE-NECTARIAN SYSTEM 149

craters are compared with typical pre-Nectarian craters in the circum-Nectaris zone. Very few craters smaller than 20 km in diameter 

have been recognized as unequivocally pre-Nectarian because of 

severe degradation and the ambiguous significance of morphology 

(Wilhelms and others, 1978). In general, pre-Nectarian craters larger 

than 20 km in diameter that are degraded by the accumulated impact 

flux and not by catastrophic deposition of later deposits have lost finescale textures but still retain the overall shape of fresh craters (figs. 

8.5B-D, 8.7C, 8.11, 8.14B, C; table 8.3). A short, sloping rim flank is 

still visible, but not the fine radial texture of the outer flanks that is 

displayed by younger craters; pre-Nectarian radial textures are visible only around basins, where they correspond to chains of secondary 

craters or, possibly, to large ridges of primary ejecta (fig. 8.7). The 

general morphology, but little sharp detail, of the crater rim-wall 

terraces is visible; the terraces are commonly amalgamated, and 

their edges cut by subradial notches. Floors are shallower, and rims 

more rounded, than in younger craters. Despite the degradation, 

many complete and well-defined rim segments are preserved. 

Because of this preservation of the overall morphology of the main 

impact sequence, even in older craters, good practice is to compare 

morphologies not only of the whole crater but also of small areas of the 

same size on each crater.

Like crater densities, morphologies of craters can be used as 

dating tools only if their geologic setting is also considered. Later 

deposits may degrade craters differentially (fig. 8.5). Moreover, typical "pre-Nectarian" morphologies may result from burial of a Nectarian crater by Nectarian or Imbrian basin deposits. For example, 

the subdued craters shown in figure 9.21 that are buried by deposits of 

the Nectarian Hertzsprung basin could be either Nectarian or preNectarian. Much sharper, unburied craters could be pre-Nectarian 

(figs. 8.7C, 8.14B, C).

200

150 O = Coulomb-barton 

x = Ingenii

AK = AI-Khwarizmi/King 

FS = Freundlich-Sharonov 

KH = Keeler-Heaviside 

LF = Lomonosov-Fleming 

pNc = Average pre-Nectarian 

crater frequency

30 40 50 60 70 80 90100 

CRATER DIAMETER, IN KILOMETERS

150

FIGURE 8.6. Cumulative size-frequency distributions of craters at least 20 km in diameter 

superposed on pre-Nectarian basins. Materials of Nectarian and Imbrian Systems are 

indicated. Curve pNc (from Wilhelms and others, 1978) includes only pre-Nectarian 

craters, whereas each basin curve includes all postbasin primary craters of sizes indicated; 

curve pNc is average for best exposed pre-Nectarian terranes. Error bars represent square 

root of plotted point, added to and subtracted from value of point (compare table 8.2). Data 

tabulated in supplementary table.

TABLE 8.3. Representative pre-Nectarian craters

[Cross rules divide table into craters smaller than 60 km, 60 to 119 km, and at least 

120 km in diameter; interiors only are mapped in plate 6. F-S, Freundlich-Sharonov 

basin; N?, possibly Nectarian]

Crater Diameter 

(km)

Daedalus U 

Con'olis L- -----

Rosenberger B 

Reimarus R -----

Daedalus R ---------

Fraunhofer E- ------

Krusenstern --------

Rosenberger C ------

Blanchinus -------

Pei rescius 

Apianus ------------

Fraunhofer J -----

Barbier ------------

Riccius ------------

Li cetus ------------

Nearch __________

Vega ---------------

Birkhoff X 

Mutus _______

Esnau It-Pel terie 

Cyrano _ __-_

Rabbi Levi ------

Zagut 

Vlacq 

Barrow -------------

Playfair G ----- -

Fra Mauro 

Rosenberger ------

Schlesinger --------

Sacrobosco ---------

Pur-bach ------------

Goldschmidt --------

Drontius 

Messala ------------

Regi omontanus 

Furnerius 

StOfler 

Szilard ----------

Paschen ---------

Stebbins ---------

Fowler -------------

Curie --------------

Hi rayama -------

Lyot 

Joliot -

Hipparchus ---------

Ptolemaeus ---------

W. Bond 

Richardson ---------

Heaviside ----------

Magi nus ----------

Einstein -----------

Tsander ------------

Fabry 

Von Karman ---------

Janssen ----------

Galois ---------

Landau -------------

Campbel 1 

Von Karman M -------

Schickard 

Deslandres 

Pasteur ------------

Lei bni tz -----------

Fermi 

Milne --------------

Gagarin 

30

32

33

35

41

42

47

47

61

62

63

63

67

71

75

76

76

77

78

79

81

81

84

89

93

94

95

96

97

98

118

120

122

124

124

125

126

127

133

135

136

139

139

141

143

151

153

158

161

163

163

170

171

179

179

190

207

221

225

225

227

234

235

236

238

262

272

Center 

(lat) (long)

4° S.

2° S.

52° S.

48° S.

8° S.

43° S.

26° S.

52° S.

25° S.

47° S.

27° S.

42° S.

24° S.

37° S.

47° S.

59° S.

45° S.

62° N.

64° S.

48° N.

21° S.

35° S.

32° S.

53° S.

71° N.

24° S.

6° S.

55° S.

47° N.

24° S.

26° S.

73° N.

40° S.

39° N.

28° S.

36° S.

41° S.

34° N.

14° S.

65° N.

43° N.

23° S.

6° S.

50° S.

26° N.

6° S.

9° S.

65° N.

31° N.

11° S.

50° S.

17° N.

6° N.

43° N.

44° S.

45° S.

14° S.

43° N.

45° N.

47° S.

44° S.

33° S.

11° S.

38° S.

20° S.

31° S.

20° S.

175° E.

173° E.

46° E.

64° E.

75° E.

62° E.

6° E.

42° E.

31° E.

68° E.

8° E.

64° E.

158° E.

27° E.

7° E.

39° E.

63° E.

150° W.

30° E.

141° W.

158° E.

24° E.

22° E.

39° E.

8° E.

7° E.

17° W.

43° E.

139° W.

17° E.

2° W.

3° W.

4° W.

60° E.

1° W.

60° E.

6° E.

106° E.

140° W.

143° W.

145° W.

92° E.

94° E.

84° E.

93° E.

5° E.

2° W.

4° E.

100° E.

167° E.

6° W.

89° W.

149° W.

101° E.

176° E.

42° E.

152° W.

119° W.

153° E.

176° E.

55° W.

5° W.

105° E.

179° E.

123° E.

113° E.

149° E.

Figure

8.7

8.7

8.5

8.5

8.7

8.5

9.27

8.5

9.27

8.5

9.27

8.5

8.7

3.10D, 11.5

10.31

8.5

8.5

8.14

8.12

8.14

8.7

11.5

11.5

8.5

10.15

1.8, 9.27

2.5A, 10.19

8.5

8.14

7.9

1.8, 7.7

10.15

1.8, 2.3

9.3

1.8, 7.7

9.5

1.8, 10.31

8.13

9.21

8.14

8.14

8.11, 10.46

8.11

1.5, 10.46

4.7, 9.4, 9.28

1.8, 3.10, 10.27

1.8, 7.7, 10.28

1.7, 10.6, 10.15

8.13

1.4, 8.7

4.6A

3.9A

9.21

1.2

8.15

3.100, 9.1, 9.2

9.21

4.2F, 8.14

4.2G, 12.9

8.15

1.9, 7.6

1.8

1.2, 8.11

8.15

1.3

1.3, 4.2J, 8.11

1.4, 3.5, 4.2K

Remarks

Older than F-S.

Younger than F-S.

Pitted by Nectaris

basin.

Older than F-S.

Pitted by Nectaris

basi n.

Mantled by Nectaris

basin.

Mantled by Nectaris

basin.

Typical .

__ _

Pitted by Nectaris

basin.

Relatively fresh.

N?

Typical .

Do.

___

N?

Relatively fresh.

N?

 

N?; covered by the

Fra Mauro and

Cayley Formations.

Typical .

Do.

N?

Typical .

___

Typical .

N?

 

 

 

 

 

 

 

 

 

Typical .

_ 

__-

Covered by the

Janssen Formation.

_ 

Typical .

Do.

__ _

Ring arcs.

Typical .

Despite the problems, the results of dating basins by superpositional relations, crater counts, and assessments of superposed crater 

morphology are generally consistent. No known craters with preNectarian morphology are clearly superposed on Nectarian basins. 

This general internal consistency is the outcome of a series of iterations in which the morphologic criteria for pre-Nectarian age were 

repeatedly reassessed. Anomalous rebound of craters inside basins 

and anomalous degradation by deposits of younger craters were 

detected by geologic mapping. Thus, some assignments of crater age 

in this volume differ from those shown on earlier geologic maps, 

including those by Wilhelms and El-Baz (1977) and Wilhelms and 

others (1979). Pre-Nectarian craters are subdivided on the paleogeologic map (pi. 6) as exposed (solid black), buried (solid rim-crest 

outline), or possibly Nectarian (dashed rim-crest outline).

A total of 1,208 pre-Nectarian craters larger than 30 km in 

diameter, 456 of which are unburied by pre-Nectarian or younger 

deposits, are mapped here (pi. 6). They vary greatly in areal density 

because the units that underlie them vary in age and because the 

thickness of overlying deposits affects the number of buried craters
150 THE GEOLOGIC HISTORY OF THE MOON

FIGURE 8.7. Stratigraphic relations between pre-Nectarian units on farside.

A. Radia! hneations (arrows) of Freundlich-Sharonov basin (FS), superposed on Keeler-Heaviside basin (named for unrelated craters Keeler [K; 169 km, 10° S., 162° E., probably Lower 

Imbnan]andHeaviside[H; 163 km, pre-Nectarian]; compare figs. 4.3P, S); areas of B and C are outlined. Orbiter 2 frames M 75 (lower left) and M 34 (upper right); overlaps figure 

1.4.
8. PRE-NECTARIAN SYSTEM 151

that can still be seen. A density of 77 unburied pre-Nectarian craters 

per million square kilometers (422 craters min 30 km diam in 

5.45 x 106 km2) was determined for terrain of the oldest basins (age 

groups 1-5, table 8.2) and other heavily cratered terrain. The younger 

pre-Nectarian basins (age groups 6 9) display a much lower average 

density of 10 such craters (34 in 3.24 x 106 km2 ). Extrapolation of the 

larger value to the whole Moon (area, 38 x 106 km2 ) suggests that a 

minimum of 2,940 craters larger than 30 km in diameter formed

- 'jf' \

' '

B. Pre-Nectarian terrane pitted by probable secondary craters of Freundlich-Sharonov 

basin (sc) and other, smaller secondaries. Craters older than Freundlich-Sharonov 

include Daedalus R (DR; 41 km) and Daedalus U (DU; 30 km); Coriolis L (CL; 32 

km) may be younger. A much younger crater chain secondary to a distant impact crater or 

basin, probably Orientate (centered 2,800 km to east), is also superposed on Coriolis L 

and Daedalus U. Crater Daedalus S (DS; 20 km) is also young, probably Imbrian. 

Orbiter 2 frame H-33.

during pre-Nectarian time. Assuming that their deposits were as 

extensive relative to crater size as are those of young craters, the 

deposits covered about 110 percent of the Moon's surface area. Based 

on the Nectarian frequency (see chap. 9), which is less affected by 

saturation, the actual number of pre-Nectarian craters of this size 

that formed while the 28 pre-Nectarian basins of groups 2 through 9 

were forming was closer to 3,400. A still greater, unknown number 

formed earlier.

C. Craters Cyrano (81 km) and Barbier (67 km), similar to pre-Nectarian craters in figure 

8.5 and of a morphologic type absent in B, are beyond range of Freundlich-Sharonov 

deposits; thus, Cyrano and Barbier are probably pre-Nectarian and possibly older than 

Freundlich-Sharonov basin, despite moderately sharp topography. Orbiter 2 frame 

H-75.
152 THE GEOLOGIC HISTOKY OF THE MOON

FIGURE 8.8. Pre-Nectarian South Pole-Aitken basin, centered around lat 56° S., long 180°. Airbrush drawing by Donald E. Davis, courtesy of the artist.

FIGURE 8.9. South Pole-Aitken massifs photographed by Apollo 8. Arrow indicates crater 

Mohorovicic A (20 km, 16° S., 163° W.) atop one arcuate massif; other massifs along lat 22° 

S occupy horizon. At time of Apollo 8 mission, relation of the large mountains to a basin had 

not yet been recognized. View southeastward. Apollo 8 frame H 2319.
8. PRE-NECTARIAN SYSTEM 153

FIGURE 8.10. Part of southern Oceanus Procellarum region, including Nectarian Humorum basin (H), pre-Nectarian Insularum (In) and Nubium (N) basins, and three rings of pre-Nectanan 

Procellarum basin (Whitaker, 1981); central Procellarum ring is expressed here mostly by mare ridges (compare fig. 5.26). Shallowness of mare is indicated by many islands belonging to basins 

and craters that are superposed on Procellarum basin (pi. 3). Orbiter 4 frame M 125.
FIGURE 8.11. Large section of western farside, centered on craters Pasteur (P; 

235 km, 12° S., 105° E., Pre-Nectanan) and Hilbert (H; I 70 km, 18° S., 

108° E., Nectarian), both filled by cratered plains. White dots, indefinite preNectanan basin Al-Khwarizmi/King (Wilhelms and El-Baz, 1977), named for 

craters Al-Khwarizmi (A; 65 km, 7° N., 107° E., Nectarian) and King (K; 

77 km, 5° N., 121 ° E., Copernican); most conspicuous part of rim is southwest

of letter A. MS, Mare Smythii; s and black dashes, rim of pre-Nectarian 

Smythii basin. Lacus Solitudinis (LS) is not related to known basins. PreNectarian crater or small basin Milne (M; 265 km, 31° S., 113° E.) has 

basinlike ring structure but is omitted from table 4.1 because of its size. Orbiter 2 

frame M-196.
8. PRE-NECTARIAN SYSTEM 155

FIGURE 8.12. Pre-Nectarian Mutus-Vlacq and Werner-Airy basins.

A. South-central nearside and area around south pole including Imbrian Schrodinger basin 

on farside (s; 320 km, 76° S., 134° E.). Black dashes, rims of Mutus-Vlacq (below) 

and Werner-Airy (indistinct, above). Arrowheads, prominent massifs of Mutus-Vlacq; 

those indicated by larger arrowheads are also shown in B. Black p, Nectarian plains 

concentrated in Mutus-Vlacq basin. N, Nectaris-basin rim (Rupes Altai). PreNectarian craters include Ptolemaeus (white p; 153 km, 9° S., 2° W.) and Vlacq, (v; 

89 km, 53.5° S., 39° E.); Nectarian craters include Maurolycus (m; 114 km, 42° S., 

14° E.) and Airy (a; 37 km, 18° S., 5.5° E., possibly Imbrium-basin secondary crater); 

Werner (w; 70 km, 28° S., 3.5° E.) is Eratosthenian. Orbiter 4 frame M-95.

B. Detail of Mutus-Vlacq massifs. Large arrows indicate same massifs as large arrowheads 

in A. Small arrow indicates hill possibly belonging to an inner ring. M, Mutus (78 km, 

63.5° S., 30° E., pre-Nectarian). Orbiter 4 frame H-82.
156 THE GEOLOGIC HISTORY OF THE MOON

VOLCANIC ROCKS

The discovery of ancient volcanic rocks on the Moon would be of 

great interest from the standpoints of both petrogenesis and stratigraphy. The unforewarned geologist might think that such rocks have 

been found in abundance because the terms "highland basalt," 

"KREEP basalt," and numerous other "basalts" are so common in the 

literature. However, no pristine samples of terra composition are 

known to be both pre-Nectarian in age and volcanic in origin; volcanic 

KREEP basalt of Early Imbrian age is the oldest currently recognized 

terra-type volcanic material (see chap. 10). Fragments of mare basalt 

have been recovered from breccia deposits emplaced in the Nectarian 

and Imbrian Periods, and a few have been radiometrically dated 

(table 9.5; chaps. 9,10). Whether they are pre-Nectarian depends on 

the absolute age of the Nectaris basin (see chap. 9); none are of 

definite pre-Nectarian age.

Moreover, no plains or other terrane of pre-Nectarian age sufficiently resemble maria to be definitely established as volcanic. Plains 

in such basins as Mutus-Vlacq (fig. 8.12) may consist of pre-Nectarian 

basalt underlying Nectarian ejecta, but this is uncertain (see chap. 9). 

All pre-Nectarian interbasin terrane is ascribable to basin deposits 

and is here mapped accordingly (pi. 6).

Despite the absence of firm evidence for pre-Nectarian volcanism, mare and KREEP basalt may have been extruded during the 

period. The extensive early plutonism suggests accompanying volcanism, and some of the recovered mare-basalt clasts are, at least, 

very close to pre-Nectarian in age. Another line of indirect evidence 

for volcanism is the concentration of KREEP revealed by gamma-ray 

spectrometers in the Imbrium-Procellarum region. The Procellarum 

basin may have unroofed deep-seated KREEP-rich material that was 

then excavated by further impacts; alternatively, however, KREEP 

basalt may actually have been extruded volcanically into the basin 

(fig. 8.4B; Wood, 1972b; Cadogan, 1974, 1981). These pre-Nectarian 

extrusions could have been recycled by later impacts and mare extrusions, so that the KREEP chemistry appears to be concentrated today

FIGURE 8.13. Pre-Nectarian Lomonosov-Fleming basin (dashes outline rim), named for 

unrelated craters Fleming (F; 130 km, 15° N., 109.5° E., Nectarian) and mare-filled 

Lomonosov (L; 93 km, 27.5° N., 98° E., Lower Imbrian). Pre-Nectarian craters include 

Richardson (R; 161 km) and Szilard (S; 127 km)i Nectarian and Imbrian plains occupy 

much of basin interior, especially between arrows and east of Fleming. Groovelike crater 

chains (arrows; above, Catena Artamonov; below, Catena Dziewulski) are radial to and, 

probably, secondary to the Imbnum basin, centered 3,000 km away on opposite side of 

Moon. Apollo 12 frame H-8296.

in deposits of Archimedes, Aristarchus, Copernicus, and other craters 

in the region. Continued sampling might well disclose pre-Nectarian 

volcanic rocks here and elsewhere on the Moon.

CHRONOLOGY

The earliest pre-Nectarian event was the origin of the Moon. 

Before the space age, the Moon's age was inferred from U-Th-Pb ages 

of the Earth and meteorites to be 4.55 ± 0.07 aeons, assuming that all 

these bodies formed contemporaneously (Patterson, 1956). Absolute 

dating of lunar rocks based on radiometric ages and assumed initial 

isotopic ratios (model ages) is consistent with 4.55 aeons ago as the 

approximate time of lunar origin (Tatsumoto and others, 1973,1977; 

Tera and Wasserburg, 1974; Nyquist, 1977; Tilton and Chen, 1979).

Some recovered clasts of ANT rock yield very old isotopic ages 

(table 8.4). Ferroan anorthosite is difficult to date because it contains 

very little Rb, Nd, and Sm; its low initial 87Sr/86Sr ratios are consistent with, but do not require, very primitive ages (Papanastassiou 

and Wasserburg, 1975, p. 1483; Tatsumoto and others, 1977). More 

precise ages appearing to be crystallization ages have been obtained 

for pristine Mg-suite rocks (table 8.4). They range as old as 4.51 ± 0.07 

aeons for inclusions in olivine of troctolite sample 76535 from the 

Apollo 17 landing site (Papanastassiou and Wasserburg, 1976). The 

analysts were uneasy about the similarity of this age to that of the 

Moon, but the age may, indeed, have been set by igneous melting in 

the Moon shortly after its accretion (Carlson and Lugmair, 1981b). A 

later event in the same rock's history at about 4.26 aeons ago is 

recorded by Sm-Nd mineral isochrons and 40Ar-39Ar plateaus. If the 

4.51-aeon age is an artifact, this later age may date the crystallization 

of the troctolite, or it may date the end of a long cooling episode during 

which isotopic reequilibration was possible (Gooley and others, 1974; 

Lugmair and others, 1976; Carlson and Lugmair, 1981b).

Interpretation of such internal isochrons as dates of igneous 

melting, along with evidence for very early geochemical fractionations (Rb-Sr, Pb-U, Sm-Nd), leads to the conclusion that the differentiation which formed the lunar crust began during or immediately 

after accretion (for example, Papanastassiou and Wasserburg, 1971a, 

1975, 1976; Tera and Wasserburg, 1974; Lugmair and others, 1975; 

Nunes and others, 1975; Herbert and others, 1977; Nyquist, 1977; 

Tatsumoto and others, 1977; Lugmair and Carlson, 1978; Carlson and 

Lugmair, 1981a, b). How long this differentiation lasted and the extent 

of igneous activity after crustal formation but before extrusion of the 

visible mare basalt are unanswered questions. The geologically 

important date of when the crust solidifed enough to support basinand-crater topography probably is between 4.4 and 4.2 aeons ago; the

TABLE 8.4. Radiometric ages of samples of pristine pre-Nectarian lunar rocks

[Sample 15455, Apollo 15 (Ryder and Bower, 1976); numbers beginning with 6, Apollo 16 (111rich and otners, 

1981); numbers beginning with 7, Apollo 17 (Wolfe and others, 1981). Numbers after commas represent splits 

of sample. Mineralogy: cpx, clinopyroxene; ol, olivine; opx, orthopyroxene; pi, plagioclase.

Ayes calculated using decay constants recommended by the International Union of Geological Sciences (IUGS) 

(Steiyer and Jager, 1977).

References: B75, Boyard and others (1975); CL81a, Carlson and Luymair (1981a); CLSlb, Carlson and Lugmair 

(1981b); C75, Compston and others (1975); HS75, Husain and Schaefer (1975); HW75, Huneke and Wasserbury 

(1975); J77a, Jessberyer and others (1977a); L75, Leich and others (1975); L76, Lugmair and otners (1976); 

N76, Nakamura and others (1976); N79b, Nyquist and others (1979b); NSla, Nyquist and others (igSla); PW75, 

Papanastassiou and Wasserbury (1975); PW76, Papanastassiou and Wasserbury (1976); S74, Stettler and others 

(1974)]

,. , Source rock and litnology 

"'P (minerals in volume percent)

Aye 

(aeons)

Reference

15455

67435

67667

72255

72417

732bS, 

27,45

Norite clast (60-65 pi, 35-40 

opx) from black-and-white 

breccia.

Tiny plagioclase clasts from 

354-y polymict breccia.

7.9-g rake sample of Iherzolite 

(58 ol, 21 pi, 15 opx, 5 cpx).

Plagioclase from 2-cm "Civet 

Cat" norite clast from boulder 

1, station 2.

Dunite (93 ol); part of clast 

72415-7241R from boulder 3, 

station 2.

Granulated 0.9-y norite clast 

(53 pi, 40 opx) from 394-g 

polymict breccia 73255.

156-g rake sample of troctolite 

(37-60 ol, 35-58 pi, 4-5 opx).

4.48+0.12 Rb-Sr

4.35±0.05 Ar-Ar

4.18+0.07 Sm-Nd

3.93+0.03 Ar-Ar 

4.08+n.05 Rr-Sr

4.45+0.10 Rb-Sr

N79b Isotopic systems disturbed (CLRlb).

CL81a Crystallization or impact-excavation 

age.

L75 Possibly as old as 4.36+0.13 aeons 

C75 (C75).

PW75 Imprecise whole-rock isochrons.

4.23+O.U5 Sm-Nd CLBlb Crystallization age; not disturbed.

4.16+0.04 Ar-Ar

77215, Norite clas 

37 boulder.

78236 Shocked norite clast (opx=pl) 

from station 8 boulder.

4.26+0.02 

4.26+0.06 

4.27+0.08 

4.51±0.07

3.98±0.04 

4.33±0.04 

4.37±0.07

4.34+0.05 

4.38+0.02

4.39 

4.43+0.05

Ar-Ar 

Sm-Nd 

Ar-Ar 

Rb-Sr

Ar-Ar 

Rb-Sr 

Sm-Nd

Sm-Nd 

Rb-Sr 

Ar-Ar 

Sm-Nd

HW75 Crystallization or impact-excavation

age.

HS75 Do. 

L76 Do. 

B75 Do. 

PW76 Crystallization age or artifact.

S74 Crystallization age at least 4.33

N76 aeons (CLSlb).

N76

CLBlb Crystallization age is probably 

NSla close to 4.34 aeons (CLSlb). 

N81a 

NSla
8. PRE-NECTARIAN SYSTEM 157

older date is more common in the literature (for example, Tera and 

Wasserburg, 1974). Taylor (1982, p. 246), however, suggested that 

crystallization proceeded rapidly until 4.3 aeons ago and that intense 

impact bombardment kept the crust thoroughly mixed and chemically requilibrated until about 4.2 aeons ago. The picture of igneous 

petrogenesis that is emerging furthermore suggests that the differentiation was complex. Local zones may have solidified early, especially 

near the surface, and local hotspots probably persisted until late. 

Early-formed crust was probably intruded by later plutons, dikes, 

and sills. This complex, late-ending history (Nunes and others, 1975; 

NakamuraandTatsumoto, 1977, p. 2312; Carlson and Lugmair, 1981b) 

seems more likely than the simpler scenarios that appear in much of 

the literature.

Some pre-Nectarian clasts record probable postcrystallization 

impacts. The 4.26-aeon event in the history of troctolite sample 76535 

may have been an impact that ended the cooling (Papanastassiou and 

Wasserburg, 1976; Carlson and Lugmair, 1981b). Dunite sample 

72415 72418 was shocked at least once, probably during the preNectarian, before incorporation into the Apollo 17 melt-rich breccia 

deposits (Dytnek and others, 1975). The more abundant Apollo 17 

norite samples were also shocked before final emplacement 

(Nakamura and Tatsumoto, 1977; see summary by James, 1980). The 

shock events may have been the formation of one of the pre-Nectarian 

basins in the Serenitatis target area, such as Tranquillitatis (pi. 6). 

Similarly complex, extended histories have been detected in other 

primitive lunar rocks (table 8.4). The later events indicate that the 

isochrons have been disturbed, and so the pristine ages are not 

accurate except as minima.

The impact processing manifest in samples from younger deposits generally supports the high early impact rate indicated by the 

abundance of observed pre-Nectarian basins and craters. For example, widespread thermal metamorphism at shallow depths, probably 

caused by accumulation of hot ejecta blankets, is suggested as the 

origin of the granulitic textures of many impact-mixed rocks (Warner 

and others, 1977; Bickel and Warner, 1978; James, 1980). However, 

more specific information about pre-Nectarian chronology applicable 

to the stratigraphic record has not been forthcoming from the rocks. 

Attempts to calibrate the impact flux by means of the amounts of 

siderophile elements or impact-melt rock are hindered by uncertainties about projectile-to-ejecta ratio or energy partitioning between 

impact melt and excavation (for example, Lange and Ahrens, 1979).

Another potential method of establishing the emplacement ages 

of pre-Nectarian units is the classic one of calibrating crater frequencies with absolute ages. However, the relative ages of ancient surfaces 

are difficult to determine precisely by crater frequencies because they 

are saturated to large crater diameters (Baldwin, 1969; Hartmann, 

1975); all the frequency curves will have similar intercepts at 20-km 

diameter and will differ only by subtle differences in the steady-state 

diameters. In figure 8.16, an envelope of curves is fitted to (1) the 

approximately known relative ages and well known absolute ages of 

the Imbrian and Eratosthenian mare-basalt samples (see chaps. 11, 

12), (2) the well known relative and absolute ages (3.85 aeons) of the 

Imbrium basin (chap. 10), (3) the well known relative age and the 

estimated absolute age (3.92 aeons) of the Nectaris basin (chap. 9), 

and (4) the highest crater frequency found in this study, that within 

the limits of the possible Al-Khwarizmi/King basin. The boxes for the 

Imbrium, Nectaris, and Al-Khwarizmi/King surfaces are bounded 

below by the actual observed cumulative frequencies of all craters 

larger than 20 km in diameter and above by projections to 20-km 

diameter from the parts of the curves representing craters large 

enough to have escaped steady-state obliteration (fig. 8.6). Proper 

comparison of the boxes requires comparing the same edge of each. 

This extrapolation suggests that the Al-Khwarizmi/King crater population is between 3.97 and 4.03 aeons old and that, accordingly, the

basin itself, part of pre-Nectarian group 2, is equally old or somewhat 

older.

A steep exponential decline like that plotted in figure 8.16 is only 

one of many possible models of the early impact rate (Hartmann, 

1965a, 1966,1972c; Baldwin, 1971; Shoemaker, 1972,1981; Wetherill, 

1975a, 1977a, b, 1981; Basaltic Volcanism Study Project, 1981, chap. 8). 

Possibly, the pre-Nectarian impact rate was closer to the Nectarian 

rate (19,000 craters and 157 basins per aeon; see chap. 9) than is 

generally assumed. In this case and if, additionally, the pre-Nectarian crater size-frequency distribution was originally the same as 

the Nectarian (small craters not yet destroyed), 3,400 craters from 30 

to 300 km in diameter, as well as the 28 basins of pre-Nectarian 

groups 2 through 9, formed from 4.1 to 3.92 aeons ago. A similar figure 

of 4.06 aeons for the age of the oldest group-2 basins is derived by 

assuming a constant rate of surface blanketing by basin deposits 

(Nectarian basin deposits one basin diameter in radius theoretically 

cover 41 percent of the Moon, and group 2-9 pre-Nectarian deposits 

cover 84 percent). Thus, the models of steeply declining and constant 

impact rates bracket the age of the oldest group-2 basins and accompanying craters between about 4.1 and 4.0 aeons. The younger age is 

more likely if the early impact rate declined as steeply as shown in 

figure 8.16, and may even be too old if the impacts declined still more 

steeply. On the other hand, 4.1 aeons may be too young if the Nectarian Period lasted longer than 0.07 aeon. I tentatively adopt 4.1 

aeons as the age of the oldest group-2 basins because of this intermediate position among the likely estimates.

The two giant basins of group 1, Procellarum and South PoleAitken, probably formed before all the other visible basins and all the 

visible craters. Because these basins are probably saturated by 

impact structures of basin size, their ages are very difficult to estimate without more calibration points or better models of the impact 

rate than are available. The two models of the impact rate discussed 

here imply different scenarios in the early era when these basins 

formed. If the impacts declined steeply all through pre-Nectarian 

time, far more basins formed than are observed. This scenario is, in 

fact, consistent with the extreme impact processing observed in the 

lunar crustal rocks and the relative paucity of "pristine" samples. If 

extrapolated to the early pre-Nectarian, however, such steep curves 

also imply that an unrealistically great mass of material impacted the 

Moon since its accretion (R.B. Baldwin, oral commun., 1983). This 

difficulty is avoided if basins formed throughout pre-Nectarian and 

Nectarian time at an approximately constant rate. About 14 basins in 

addition to the two giant basins would then have formed from 4.2 to 

4.1 aeons ago still enough for considerable impact processing of the 

crust.

In summary, the many uncertainties preclude accurate estimates of any pre-Nectarian event after the origin of the Moon 4.55 

aeons ago. The many great impacts that must have preceded 4.2 aeons 

may have struck either a solid or a mushy crust; the geologic record 

does not specify which, but frequencies of craters and basins are, at 

least, consistent with the 4.2-aeon date of crustal solidification given 

earlier. Procellarum, South Pole-Aitken, and at least 14 now-obliterated basins formed between crustal solidification and the oldest of 

the 28 pre-Nectarian basins of groups 2 through 9. The group-2 

basins began about forming 4.1 aeons ago if the Nectarian Period 

lasted 0.07 aeon (from 3.92 to 3.85 aeons ago) and if the impact rate 

was approximately constant during pre-Nectarian and Nectarian 

time. These conservative estimates of 4.2, 4.1, and 3.92 aeons for the 

major milestones of the pre-Nectarian period are consistent with 

available data but may be considerably altered by better data on the 

age of the Nectaris basin or the early impact rate. At least, preNectarian basins and craters clearly formed in such abundance that 

they repeatedly destroyed earlier features and reworked earlier 

crustal materials.
' V V >* w» ^'V'£v 1 y^*)* r » ! * , \\> '«, '/ - *%r, -y^v. M , 

t .^ ':i.^>V J ^^^4 <f ;y, f
8. PRE-NECTARIAN SYSTEM 159

B. Detail of Coulomb-Sarton basin and superposed craters, including Coulomb (C) and Sarton (S). Orbiter 5 frame H 24. C. Craters superposed on Birkhoff include Stebbms (S; 135 

km, 64.°5 N., 142.5° W.) and Birkhoff X (BX; 77 

km, 62° N., 149.5° W.), probably pre-Nectarian; and 

Carnot (C; 52.5° N., 144° W.), fresher appearing than 

the others, probably Nectarian. Numerous pits on floor 

of Birkhoff are probably secondary craters of Stebbms, 

Carnot, and other large nearby craters. Birkhoff deposits 

or secondary craters overlie pre-Nectarian craters Fowler 

(F), Schlesmger (Sch; 97 km, 47.5° N., 138.5° W.), 

and probably Esnault-Pelterie (E-P; 79 km, 47.5° N., 

141.5° W.); sharper rim of E-P suggests, alternatively, 

degradation only by Carnot deposits and superposition on 

Birkhoff deposits. Secondary craters of Orientale and 

(or) Hertzsprung basins are conspicuous in Fowler (compare A). Orbiter 5 frame H-29.

FIGURE 8.14. Stratigraphic relations of pre-Nectarian basins Birkhoff (59° N., 145.5° 

W.), Coulomb-Sarton (530? km, 52° N., 123°W.), and Lorentz (360 km, 34.5° N., 97° 

W.) on northeastern farside.

A. Widely spaced dashes outline basin rims. Interior of Coulomb-Sarton basin (C-S) is 

pitted by small subequal subdued craters probably secondary to Birkhoff and Lorentz, 

and is overlain by radial chains secondary to Nectarian Hertzsprung basin (Nh), 

centered in direction of arrow 1,350 km south of crater Fowler (136 km, 43° N., 145° 

W., pre-Nectarian); dashed line encloses additional secondaries and other outer deposits 

of Hertzsprung. Craters superposed on C-S basin include Landau (221 km, 43° N, 

119° W., pre-Nectarian), Coulomb (89 km, 54.5° N., 114.5° W., Nectarian), and 

Sarton (69 km, 49.5° N., 121° W., Nectarian). Overlap of old craters by Lorentz 

ejecta is visible at X. Sharp small craters in lower left near crater Fersman (143 km, 18° 

N., 125° W., pre-Nectarian) are secondary to the Imbrian Orientale basin (lo; centered 

in direction of arrow) and to Hertzsprung. Kovalevskaya (111 km) is probably Upper 

Imbrian. Orbiter 5 frame M-8.
160 THE GEOLOGIC HISTORY OF THE MOON

10~3

FIGURE 8.15. Region west and southwest of Apollo basin, including Apollo-secondary 

craters (right of dashed line; most evident examples at letters as); black-and-white arrows 

point toward basin center. Center of indefinite basin Grissom-White (GW) is also indicated, 

but basin could be part of Apollo's peripheral structure. Pre-Nectarian craters include 

Leibnitz (L; 236 km), Von Karman (V; 179 km, 44° S., 176° E.) and Von Karman M 

(VM; 225 km). Oppenheimer (O; 206 km, 35° S., 166° W.) is believed to be Nectarian, 

though poorly photographed. Orbiter 5 frame M-43, rectified at University of Arizona's 

Lunar and Planetary Laboratory, Tucson, Ariz.

10

"as

Si q°

I- CC

< W a: a10

10-

Al- Khwarizmi / King

craters ~\ | mbrjum basin 

20-km craters J

10'

10

Is£*

LU

CC O

£ W

< DC

CC LU 

LU ^

!<°

10

10-

4.5 4.0 3.5 3.0 2.5 

ABSOLUTE AGE (AEONS BEFORE PRESENT)

FIGURE 8.16. Cratenng rate during pre-Nectarian, Nectarian, and part of Imbrian time. 

Boxes represent observed crater densities (height) and absolute ages (width) given in later 

chapters; crosses represent my estimates for most likely values. Craters larger than 20 km in 

diameter (left vertical axis) were counted on Imbrium-basin and older units, and small craters 

(right vertical axis) on Imbrium-basin and younger units; Imbrium basin (fig. 10.4; table 7.3) 

provides calibration between cumulative frequencies of craters at 20- and I-km diameters. 

Straight segments of curves drawn through corners of boxes enclose likely envelope of 

cratering frequencies as a function of time (dashed); heavy solid line is drawn through centers 

of crosses. Crater densities, but not absolute ages, are available for Al-Khwarizmi/King.
'6 NVIHV133N W31SAS 
FIGURE 9.1 (OVERLEAF). Deposits and secondary craters of Nectaris basin southeast of basin rim (upper left). View centered on Vallis Rheita (VR), a Nectaris-basm secondary-crater chain, 

and crater Rheita, a superposed Nectarian primary-impact crater (R; 70 km, 37° S., 47° E.). Pre-Nectanan terrane is at lower left; largest crater is Janssen (j; 190 km, 45° S., 42° E.). 

Compare figures 8.15 and 9.2. Orbiter 4 frame M-52.
9. NECTARIAN SYSTEM

CONTENTS

	Page

Introduction ...................................................... 163

Nectaris basin ................................................... 163

General features .............................................. 163

Nectaris component in Apollo 16 samples ........................ 164

Absolute age .................................................. 168

Crisium basin .................................................... 170

General features ............................................... 170

Luna 20 samples ............................................... 171

Serenitatis basin .................................................. 171

General features ............................................... 171

Relative age ................................................... 173

Petrology of Apollo 17 samples ................................. 173

Emplacement process .......................................... 174

Absolute age .................................................. 177

Other basin and crater materials ................................... 178

Basins ........................................................ 178

Craters ....................................................... 180

Volcanic rocks ................................................... 190

Chronology ...................................................... 190

INTRODUCTION

The Nectarian System includes all materials of the Nectaris 

basin and all younger lunar materials emplaced before the deposition of Imbrium-basin materials (Stuart-Alexander and Wilhelms, 

1975). Thus, the Nectarian Period comprises the timespan between 

the Nectaris and Imbrium impacts. This system contains fewer rockstratigraphic units than does the pre-Nectarian but is nevertheless 

widespread on the Moon (pi. 7). Nectarian basin materials cover a 

large area on the nearside and farside centered on the east limb 

(figs. 9.1-9.6). Other concentrations are on the farside around the 

Orientale deposits. Nectarian craters, formed at a declining but still 

frequent rate (fig. 8.16), are scattered randomly wherever younger 

deposits are absent or thin. Nectarian plains are more evident than 

are pre-Nectarian; photogeologic and sample data to be discussed 

here suggest that some Nectarian plains may be partly volcanic in 

origin.

Samples were collected from as many as three Nectarian basins. 

The Apollo 17 mission collected large amounts of material of the 

Serenitatis basin, which is probably Nectarian. Small fragments 

returned by Luna 20 from a spot on the Crisium-basin flank probably were emplaced in the ejecta of that Nectarian basin. Apollo 16, 

arguably, returned materials of the Nectaris basin. As always, the 

term "Nectarian" in this context refers to the photogeologically 

observed stratigraphy (see chaps. 2, 7, 8). Most of the materials 

contained in Nectarian units originated endogenically during the 

pre-Nectarian, and, additionally, had been extensively reworked by 

pre-Nectarian impacts before emplacement during the Nectarian 

Period. In general, the samples can be identified with specific Nectarian stratigraphic units better than with the pre-Nectarian units 

in which they once resided. The possible Nectaris-basin provenance 

of the important and controversial Apollo 16 sample suite is discussed in this chapter, and alternative interpretations are weighed 

in chapter 10.

NECTARIS BASIN

General features

Deposits of the Nectaris basin clearly separate pre-Nectarian 

and younger Nectarian deposits over a wide area (pis. 3, 7; figs. 8.5, 

9.1, 9.2, 9.5, 9.6). At 860 km in diameter, the Nectaris basin is only 70 

km smaller than the Orientale basin (table 4.1), and probably 

resembled Orientale before modification by later basin and crater 

deposits. The west and south sectors have a moderately well defined 

ring structure spaced like the Orientale rings (fig. 9.2A). In fact, 

Nectaris was chosen as the type basin by Hartmann and Kuiper 

(1962).

The landscape south and southeast of the Nectaris basin is a 

degraded equivalent of the Hevelius Formation and the secondary 

craters south and southeast of Orientale (figs. 4.4, 9.1; Stuart-Alexander, 1971; Wilhelms and McCauley, 1971; Stuart-Alexander and 

Tabor, 1972; Stuart-Alexander and Wilhelms, 1975; Wilhelms and ElBaz, 1977). Thick mantling of pre-Nectarian terrane is evident from 

the paucity of large craters in comparison with the region farther 

south (figs. 8.5, 9.1). Like the circum-Orientale terrane, the Janssen 

Formation (fig. 9.2; Stuart-Alexander, 1971) and much other circumNectaris terrane are strongly lineate radially to the basin center. 

Vallis Rheita (figs. 9.1, 9.2) and similar troughs (fig. 9.5; Hodges, 

1973b), consist of subtle overlapping or tangential elliptical craters 

and evidently were formed by secondary impact, as was Vallis 

Bouvard on the south Orientale-basin flank (fig. 4.4D). Subcircular 

satellitic craters of Nectaris (figs. 9.1, 9.25) also resemble the secondary craters of Orientale in shape, cluster geometry, and distance from 

the source basin, although they appear less crisp. At Nectaris, the 

primary-secondary transitional zone that is reproducibly mappable 

east of Orientale and is also detectable at Imbrium (chap. 4) has been 

blurred by "aging" and secondary cratering from Imbrium (fig. 3.10A); 

nevertheless, an attempt has been made here to identify the contact

163
164 THE GEOLOGIC HISTORY OF THE MOON

between the two Nectaris fades (pis. 3, 7).

The north and east sectors of Nectaris have an overall arcuate 

map pattern, but individual rings are inconspicuous (fig. 9.7). These 

sectors resemble western Orientale in this respect (fig. 4.5), although 

the Nectaris uplifts are more flat-topped than jagged. The northern 

and eastern Nectaris ejecta is largely buried by younger deposits (pi. 

7). Any secondary chains that were formed in these sectors have been 

buried by younger deposits of the maria and of the post-Nectaris 

Crisium, Serenitatis, and Imbrium basins (pis. 3, 7; figs. 9.3-9.6).

More light-colored plains are concentrated inside the Nectaris 

basin than in other settings except the Orientale- and Imbrium-basin 

peripheries (fig. 10.34). This setting suggests an analogy with the 

Maunder Formation of Orientale, that basin's interior impact melt. 

Hummocky material on the northern Nectaris terra is similar in 

topography and relative position to the Montes Rook Formation and 

may be the Nectaris equivalent (fig. 9.7; Elston, 1972; Wilhelms, 

1972c). However, this hummocky material lies on the distal sides of 

Imbrium-secondary craters and thus may consist of their decelerated 

ejecta (Wilhelms, 1980, p. 45).

Light-colored plains are also widespread south of the basin (pi. 

7). Wilhelms and McCauley (1971) and Scott (1972b) interpreted the 

large patches of cratered or "pitted" plains southwest of the basin (fig. 

9.25) as volcanic in origin by analogy with the younger, Imbrian 

plains also thought to be volcanic before the Apollo 16 analyses. The 

subequal "pits" and larger, also subequal, flat-floored depressions 

were thought to be volcanic because of their grouping. The superposition of the pitted plains and smoother Imbrian plains on the larger 

craters was interpreted as analogous to the flooding of craters by 

mare materials. An alternative explanation, however, seems to be 

required by the Apollo 16 results. The plains grade with the lineate 

Janssen Formation, as do circum-Orientale or circum-Imbrium 

plains with the lineate ejecta of those basins. Therefore, the stratigraphic relations here can be interpreted by the impact model of 

lunar-terra formation. The large grouped craters are probably Nectaris secondaries, and most of the smaller craters or "pits" are probably Imbrium secondaries. The plains were likely emplaced by one of 

the ejecta-deposition mechanisms whereby a surge of debris arrives 

at a site after the secondaries form. This concentration of plains may 

result from catchment of the Nectaris ejecta by the Mutus-Vlacq 

basin (fig. 8.12; Wilhelms and others, 1979). In addition, a marelike 

volcanic deposit, extruded into the Mutus-Vlacq basin during preNectarian time, may underlie the plains and contribute to their level 

topography (see section below entitled "Volcanic Rocks").

The terrain west of Nectaris differs considerably from that in the 

other sectors. Rupes Altai (the Altai scarp) constitutes a continuous, 

well-defined topographic rim (fig. 9.2A). Despite its proximity to this 

rim, the western region contains little evidence of Nectaris ejecta. 

Thick, strongly lineate ejecta is visible at the same distance from the 

southwest and southeast rim sectors. Part of the reason for the 

difference is that the western Nectaris deposits were modified by 

Imbrium projectiles when the strong Imbrium-radial pattern of craters and grooves was formed (figs. 3.10, 9.2A). Additionally, Nectaris 

ejecta may have been deposited less thickly here than in the south.

Nectaris component in Apollo 16 samples

The thickness of the western Nectaris ejecta is important in the 

interpretation of samples collected by the Apollo 16 mission. The 

astronauts landed 60 km west of the Kant Plateau, a broad part of the 

Nectaris-basin rim crest about 100 km wide (pi. 3; figs. 2.5C, 9.8). At 

Orientale, thick ejecta is visible in the most nearly analogous position, 160 km east of the Cordillera scarp, even though this point is in

FIGURE 9.2. Materials of Nectaris basin.

A. Regional view showing Mare Nectaris and Nectaris basin (upper right quadrant) and

area of B (outlined). Letter R is below crater Rheita and right of Vallis Rheita. Orbiter

4 frame M-83. 

B. Detail of Janssen Formation (Nj), plains deposits (Np), Nectaris-basin secondary

craters (Nnsc), and Imbrium-basin secondary craters (lisc). Nj, type area of the Janssen

Formation (Stuart-Alexander, 1971), which partly buries pre-Nectarian crater Janssen;

arrow is radial to center of Nectaris basin. Craters Nicolai (42 km, 42° S., 26° E.) and

Pitiscus (82 km, 50° S., 30° E.) are Nectarian. Nj and Np thickly to thinly bury Nnsc;

lisc overlies Nectarian units (Scott, I972b). Orbiter 4 frame H 83.
9. NECTARIAN SYSTEM 165

the relatively "excluded" zone of ejecta (pi. 3; fig. 4.4F). Thus, Nectaris 

ejecta probably either underlies or composes the surface materials at 

the Apollo 16 landing site. However, the search for Nectaris materials 

has been inconclusive.

Two photogeologic units were sampled by Apollo 16. The first is 

the Cayley Formation, which forms the Cayley plains (figs. 9.8, 9.9). 

This formation was emplaced during the Imbrian Period (chap. 10); 

however, it may consist of Nectaris materials reworked during its 

emplacement. The Cayley was sampled in the central and southern 

parts of the region (all stations except 11 and 13, fig. 9.9). Station 10, 

near the landed lunar module (LM) and the geophysical instruments 

(Apollo Lunar Surface Experiments Package [ALSEP]), is probably 

the most intensely examined part of the Moon (Muehlberger and 

others, 1980; Schaber, 1981). Nevertheless, no outcrops or large boulders of the Cayley were found. A regolith 10 to 15 m thick (Cooper and 

others, 1974; Freeman, 1981) obscures the internal stratigraphy of the 

Cayley. The Cayley may contain coherent, lithologically uniform 

strata or lenses (Hodges and Muehlberger, 1981), or may consist of 

loosely aggregated poorly sorted clastic debris. Some of the mechanisms described in chapter 10 could have emplaced a Cayley Formation consisting of poorly sorted Nectaris debris during the Imbrian 

Period.

Regolith samples show that impact-melted materials are common in the Cayley Formation (James and Horz, 1981). Their sources 

and ages, however, are major problems. A tight compositional cluster 

(approx 29-31 weight percent A12O3 ) and a looser cluster (17-21 

weight percent A12O3 ) suggest to Ryder (1981) that two groups of melt 

rocks came from two distinct melt sheets. The first group is richer in 

Al than the average terra (25-26 weight percent A12O3 ; Taylor, 1982) 

and is mostly subophitic to intergranular in texture, like much volcanic basalt but resulting from thorough impact melting (lithologic 

type SM, compositional type V, table 9.1). The low-Al samples are the 

richest in KREEP, are laden with fragmental material, and are typically poikilitic in texture (lithologic type PM, compositional type K). 

The poikilitic texture results from the relatively mafic composition 

and from the presence of many clasts, which form nuclei for crystallization of the melt (Lofgren, 1977; Nabelek and others, 1978). In 

addition, many melt rocks have A12O3 contents intermediate between 

the two clusters (lithologic type IM, compositional type I); their 

compositions are commonly described inappropriately as "very high 

alumina (VHA) basalt" (Hubbard and others, 1973,1974).

The largest samples of the Cayley Formation were excavated by 

South Ray Crater, 680 m wide and 130 m deep, which is 2 million years 

old (table 13.2). Materials probably ejected from South Ray were 

sampled at stations 4 and 8 (Reed, 1981; Sanchez, 1981; Ulrich and 

Reed, 1981; Ulrich and others, 1981b). Among other materials is a 

distinctive dimict (dilithologic) breccia that apparently characterizes 

the Cayley Formation (James and Horz, 1981). These "black and 

white" rocks consist of cataclastic anorthosite and fine-grained 

impact-melt rock of "VHA" composition. The dimict breccia seems to 

be analogous to terrestrial dike-breccia deposits formed in the outer 

parts of impact craters (fig. 3.21). During cratering, melt intruded into 

cataclastic feldspathic material in the walls of a growing crater or 

basin and then was quenched, fragmented, and intruded by other 

clastic feldspathic material (James, 1977,1981). Fragments of similar 

anorthosite and melt rock that also abound in the regolith probably 

originated in dimict breccia.

The melt rocks of the Cayley Formation may have been emplaced 

while hot as beds or pods (Hodges and Muehlberger, 1981), or they may 

have been emplaced cold in isolated boulders that South Ray and 

other craters happened to penetrate. If individual impacts generate 

uniform melt sheets, one melt type, at most, is likely to have been 

emplaced hot when the Cayley was emplaced (Ryder, 1981). Thus, 

reworking of an already complex deposit is suggested. Nonetheless, 

large single impacts can generate great lithologic and compositional 

complexities (see chap. 3). It is still unclear how many melt sheets are 

represented by these melt rocks.

The other unit sampled by Apollo 16 is material of the Descartes 

Mountains (fig. 9.9; Milton and Hodges, 1972). For convenience, this 

unit is here called the Descartes Formation, although it has not been 

defined formally. The Descartes is divided into two gradational facies. 

The northern or Smoky Mountain facies is grooved radially to the 

Imbrium basin. Grooves in the southern or Stone Mountain facies 

curve away from this Imbrium-radial direction into a more nearly 

transverse trend approximately where the formation contacts the 

Kant Plateau (figs. 2.5C, 9.9). This embankment of the Descartes

Formation against the Nectaris-basin rim has been considered as 

evidence for a partial Nectaris origin of the topography (Head, 1972; 

Wilhelms, 1972c). Thus, both the Imbrium and Nectaris basins may 

have contributed to the morphology of the Descartes Formation; the 

same is true for the materials of which it is composed.

Large sample collections were made at stations 11 and 13 on 

Smoky Mountain from ejecta of the 1-km-wide, 50-million-year-old 

(table 13.2) North Ray Crater (fig. 9.10; Stoffler and others, 1981; 

Ulrich and others, 1981b). Station 11 is on the crater rim, and station 13 

is about one crater diameter from the rim (fig. 9.9). North Ray penetrated a 50-m-high mass-wasted ridge that is part of Smoky Mountain 

(300-400 m high overall). The 230-m depth of North Ray probably 

ensures that the Smoky Mountain material was excavated (Ulrich, 

1973; Ulrich and others, 1981b). The station 11 and 13 samples are 

mostly of fragmental feldspathic breccia whose fragmental melt-poor 

matrix is composed of material similar to that of the larger clasts (fig. 

9.10). About half the clasts are of cataclastic anorthosite (Minkin and 

others, 1977; James, 1981), a quarter of granulitic breccia, and a 

quarter of feldspathic fragment-laden impact melt (lithologic types 

CA, G, and FM, respectively, table 9.1). The granulitic breccia is 

thermally metamorphosed (Warner and others, 1977; James, 1980, 

1981). The relation of the FM-type melts to the breccia deposit in 

which they are found is not known. They may have been cogenetic 

with the fragmental breccia (as in suevite at the Ries; see chap. 3; 

Stoffler and others, 1979) or may predate the deposition of the breccia. 

Determining which is the case here would be an important step in 

interpretation of the Descartes Formation, especially considering the 

rarity of impact melts in the breccia matrices; such melts are generally the most valuable for dating impact deposits.

Small amounts of the Stone Mountain facies may have been 

sampled at stations 4 and 5, where feldspathic material like that at 

stations 11 and 13 was recovered. Contamination by Cayley material

Note: Figures 9.3 through 9.6 constitute a series of overlapping oblique photographs taken 

by Lunar Orbiter 4 of areas covered poorly from lower altitudes earlier in the mission (pi. 2); 

Sun illumination is opposite from that in other nearside photographs.

FIGURE 9.3. (p. 166, left) Part of northeast quadrant of nearside, including Mare Crisium 

and Crisium basin (bottom; centered at 17.5° N., 58.5° E.). Major Nectarian craters are 

Cleomedes (Cl; 126 km) and Endymion (E; 125 km; pi. 7; table 9.4). Arrows indicate 

secondary-crater chains radial to Crisium (C) and Imbrium (I) basins; terra south of dashed 

line is covered by deposits of Crisium basin, except where covered by younger deposits of 

Cleomedes, Geminus (G; 86 km, Eratosthenian), and Taruntius (T; 56 km, Copernican). 

A, Atlas (87 km, Lower Imbrian); H, Hercules (69 km, Eratosthenian); M, Messala (124 

km, pre-Nectarian, partly covered by Crisium deposits). Orbiter 4 frame H 191.

FIGURE 9.4. (p. 166, right) Part of northeast quadrant near east limb, including two 

distinct rings of Humboldtianum basin (300- and 600-km diameter, centered at 61 ° N., 84° 

E.) between Mare Humboldtianum (H) and crater Endymion (E; compare fig. 9.3). 

Arrows indicate Humboldtianum-radial or -subradial directions; arrow I is also radial to 

Imbrium (fig. 9.3), and arrow below crater Zeno (Z) is also radial to Crisium. Zeno is 

superposed on Humboldtianum basin and overlain by Crisium-basin secondary-crater cluster 

(compare fig. 4.3Q). Deposits and clustered craters between Gauss (G; 177 km, a 

Nectarian crater with fractured floor) and Joliot (J; 143 km, pre-Nectarian) are probably 

outer materials of Humboldtianum basin; deposits in lower third of photograph may be related 

to Crisium (C). Basin deposits underlie Nectarian craters Gauss, Berosus (B; 74 km), and 

Neper (N; 137 km, 9 N., 85 E.), Lower Imbrian crater Hahn (Ha; 84 km), and mare 

basalt (M; Mare Marginis) (compare fig. 4.7) Orbiter 4 frame H-165.

FIGURE 9.5. (p. 167, left) Terrain south of Mare Crisium (C) and east of Nectaris basin. 

Black-and-white arrows are radial to basin centers and denote large secondary-crater chains 

of Crisium basin (c) and Nectaris basin (N; lower chain is part of Vallis Rheita); g, graben. 

Outer deposits of Crisium (Nco) and Nectaris (Nno) include secondary craters and terramantling deposits. Luna 20 landing site (L20) is on Crisium-basin rim (compare fig. 11.16). 

Nj, Janssen Formation, partly covers pre-Nectarian crater Furnerius (F; 125 km, 36° S., 

60° E.). Np, plains of Nectarian age; INp, deposit in pre-Nectarian Balmer-Kapteyn 

basin, probably consisting of Nectarian plains overlain by Imbrian plains, y, deposits 

younger than Nectarian. Primary and secondary deposits of floor-fractured Lower Imbrian 

crater Petavius (P; 177 km) are superposed on plains, Crisium-secondary chain (c; Vallis 

Palitzsch), and crater Adams (A; 66 km); Adams and Pontecoulant (Po; 91 km, 59° S., 66° 

E.), both Nectarian, overlie basin deposits. L, Langrenus (132 km), an Eratosthenian crater 

superposed on Crisium deposits and Mare Fecunditatis. Orbiter 4 frame H 184.

FIGURE 9.6. (p. 167, right) Eastward continuation of geology shown in figure 9.5. 

Southward extent of Crisium-basin deposits (Nco) is obscured by large, floor-fractured, 

Upper Imbrian crater Humboldt (H; 207 km, 27° S., 81 ° E., unrelated to Humboldtianum 

basin); clustered craters between H and arrow N (indicating Nectaris center) may be 

Crisium-basin secondaries. C, direction to center of Crisium; Nno, outer deposits of Nectaris 

basin. Po, Pontecoulant (91 km, 59° S., 66° E.); m, part of Mare Australe. Orbiter 4 

frame H-178.
 ^'V*

 ^r'r

 *&&&& JEW**- «,»r V'&»V «*$3r^<ife:
-f^-rt^c A** %4^ 

. r -'j-i*i-^ :
168 THE GEOLOGIC HISTORY OF THE MOON

(or a similarity of Cayley and Descartes material), however, is likely at 

those stations, considering that many large samples consist of dimict 

breccia with a 2-million-year exposure age (Sanchez, 1981).

In summary, the Cayley and Descartes Formations appear to 

differ in several respects in addition to morphology (James, 1981; 

James and Horz, 1981; Stoffler and others, 1981). (1) The North Ray 

Crater samples (probable Descartes) are more aluminous than the 

South Ray Crater samples (Cayley). (2) The Cayley appears to contain 

more impact-melt rock, whereas the Descartes is characterized by 

friable fragmental feldspathic breccia. (3) More of the Cayley impactmelt rocks are of low-K KREEP and "VHA" composition. (4) The 

Cayley apparently has a higher magnetism as measured on the 

surface, a possible indication of buried melt-rich breccia (Strangway 

and others, 1973). Finally, (5), the Cayley is characterized by sharper 

craters, more blocks ejected onto the surface from craters, and a 

thinner regolith (Ulrich, 1973), properties consistent with the Cayley 

being a more cohesive unit than the Descartes. A possible explanation 

for these lithologic and topographic differences is that the Descartes 

Formation consists mostly of Nectaris ejecta and the Cayley of meltrich Imbrium ejecta.

This complex region is still not well understood. Neither the 

Cayley nor the Descartes Formation can be said to be well characterized lithologically or stratigraphically. Apart from a boulder composed of lithified regolith (sample 62195), the only boulders sampled 

are those from North and South Ray Craters. The observed differences may not be qualitative but may reflect relative proportions of 

the same constituents (O.B. James and M.M. Lindstrom, oral commun., 1983). Analyses of the returned samples have stressed their 

highly feldspathic composition, and only now are beginning to 

address problems of the origin and emplacement age of the breccia 

deposits.

.^4 .- .^;..-.j.: ^ ,.,-p^.^ ^; i * 4*w'^

"i -^ . r ---V^'';

A"'^'

u%

FIGURE 9.7. North rim of Nectaris basin. Hilly terrain (h) is possible Nectaris equivalent 

of Monies Rook and Alpes Formations. Topographic freshness, however, suggests emplacement by Imbrium-basin secondary craters (arrow is radial to Imbrium). Imbrium-radial 

grabens are in lower third of photograph (see chap. 6). View centered at 3° S., 37.5° E. 

Orbiter 4 frame H-72.

TABLE 9.1. 40Ar-3t>Ar ages of Apollo 16 samples

[After James (1981).

Sample numbers (Ulrich and others, 1981): 60-, station 10, central part of area 

near Lunar Module and Apollo Lunar Science Experiment Package (ALSEP); 61-, station 1, central area on rim of 300-m-diameter Flag Crater; 615-, chips from rake 

sample 61500; 63-, station 13, on ejecta of North Ray Crater, about 800 m from 

rim: 635-, 2- to 4-mm fines from rake sample 63500; 64-, station 4, on Stone 

Mountain; 65-, station 5, on Stone Mountain; 66-, station 6, near base of Stone 

Mountain; 67-, station 11, on southeast rim of North Ray Crater; 676-, 2- to 4- 

mm fines from rake sample 67600; 677-, 2- to 4-mm fines from rake sample 67700; 

68-, station 8, on ejecta of South Ray Crater, 3.3 km northeast of rim; 685-, 

fragment from rake sample 68500; identical to DM (see below). Digits and letters after commas refer to splits of sample.

Lithologies: CA, cataclastic (mechanically disrupted) anorthosite, found as 

clasts in feldspathic fragmental breccia, as one of the major lithologies in 

dimict breccia, and as monolithologic fragments in the regolith; some recrystallized. DC, anorthosite in dimict breccia. DM, impact-melt rock in dimict 

breccia; texturally diverse; mostly fine grained, compositionally uniform meltrock component of the breccia. FM, feldspathic fragment-laden impact-melt rock; 

diverse texture, nonpoikilitic. G, granulitic breccia (granulitic impactite of 

Warner and others, 1977); probably recrystallized or partially melted after 

heating of earlier breccia (Warner and others, 1977; James, 1980, 1981); compositionally heterogeneous: granulitic (granoblastic) texture consists of small 

equant crystals lacking wel1-developed crystalline forms. IM, intergranular 

impact-melt rock; fine texture: intergranular texture consists of pyroxene 

filling voids between plagioclase crystals. PM, poikilitic impact-melt rock: 

poikilitic texture consists of large single grains (commonly of pyroxene) enclosing many small crystals (commonly of plagioclase). SM, subophitic impactmelt rock; fine texture, grading to intergranular; glassy residuum also present: subophitic texture consists of pyroxene filling voids and partly surrounding plagioclase crystals.

A1 2 0 3 contents from Ryder and Norman (1980) and James (1981).

Surface unit is unit in which sample was emplaced at surface and from which it was 

collected (Ulrich and others, 1981): Ca, Cinco a crater: local, in regolith 

developed from underlying unit; NR, rim or other ejecta of North Ray Crater; SR, 

ejecta of South Ray Crater.

Source unit is photogeologic unit in which sample was emplaced near the landing 

site: C, Cayley Formation; D, Descartes Formation.

Ages recalculated from original values, using International Union of Geological 

Sciences (IUGS) decay constants (Steiger and JSger, 1977), except for those 

determined by Schaeffer and Husain (1973, 1974), Schaeffer and others (1976), 

and Schaeffer and Schaeffer (1977); reevaluation of the monitor used by these 

analysts nearly compensates for the changes in decay constants (0. A. Schaeffer, 

in James, 1981, p. 210).

References: H73, Huneke and others (1973); J74, Jessberger and others (1974); 

J77a, Jessberger and others (1977a); K73b, Kirsten and others (1973b); M78, 

Maurer and others (1978); P75, Phinney and others (1975); S73, Stettler and 

others (1973); S76, Schaeffer and others (1976); SH73, Schaeffer and Hussain 

(1973); SH74, Schaeffer and Husain (1974); SS77, Schaeffer and Schaeffer (1977); 

T73, Turner and others (1973); TC75, Turner and Cadogan (1975)]

Sample Lithology AloOo 

(wt pel)

Surface 

unit

Source 

unit

Age 

(aeons) References

60015

68415

68415

68415

67603

65015

63503

65055

67703

67703,

67016

67603,

65055

64536

60315.

67703,

67703,

64535,

65015

63503,

63503,

68416,

63503,

68503,

63503,

67603,

60315,

65785,

67075

68503,

66043,

61503.

61503,

68415,

63503,

67603,

67703,

63503,

63503,

67703,

60025,

,22

,49

,10

,50

,1-ME

,61

,17-LN

,12B

,14-NF

,14-NG

,100

,1-MQ

,12A

,3

6,1

14-NI

14-ND

7

17-LE

17-LC

34

17-LK

16,33

,17-LI

,1-MD

,19

13

16,1 

2,5

1,11

1,18

10

17-LH

1-MM

14-NL

17-LL

17-LF

14-NN

86

CA 

SM 

SM 

SM 

PM 

PM? 

IM 

SM

G

G

?

G

SM 

DC 

PM 

PM 

SM 

DC 

PM 

G 

G

SM 

IM 

IM 

G

FM 

PM 

SM 

CA? 

PM 

PM 

PM 

IM

Plagioclase 

FM 

FM 

FM 

FM 

FM 

FM 

DC?

35± 

25-30

27 

25-30

27

20±

23 

?

29

28

30±

28

17+

25

24

?

20+ 

33

25 

25-30

26

23

31

32 

<20 

20-25 

>30

23

25-30 

33

31

32

33 

33 

32 

35+

SR

SR or local 

SR or local 

SR or local

NR 

?

NR

SR?

NR

NR

NR

NR

SR?

SR or Ca 

SR or local

NR

NR 

SR or Ca

?

NR

NR

SR

NR 

Local

NR

NR

SR or local 

Local

NR

Local 

Local

SR? 

NR 

NR 

NR 

NR 

NR 

NR 

SR

C 

C 

C 

C 

D 

?

D

C?

D

D

D

D

C?

C

C

D

D

C

?

b

D 

C 

D 

C 

D 

D 

C 

D 

D 

C 

D?

3.50±0.05 

3.74±0.04 

3.80±0.04 

3.80±0.06 

3.83±0.05 

3.86+0.04 

3.87+0.04 

3.89±0.02 

3.89+0.04 

3.89+0.04 

3.89±0.05 

3.90+0.06 

3.90+0.02 

3.91+0.01 

3.91+0.02 

3.91+0.04 

3.91+0.04 

3.92+0.02

3.92

3.92+0.04 

3.93+0.04 

3.94+0.05 

3.94+0.05 

3.95+0.06 

3.96±0.04 

3.96±0.05 

3.97±0.03 

3.97+0.02 

3.98±0.05 

4.00+0.05 

4.01+0.05 

4.02±0.02 

4.02±0.01 

4.03±0.04 

4.09±0.07 

4.10+0.05 

4.12+0.05 

4.13+0.04 

4.13+0.1 

4.14+0.04 

4.17-4.21

SH74, P75.

S73.

H73.

K73b.

M78.

K73b.

M78.

J77a.

M78.

M78.

TC75.

M78.

J77a.

J77a.

S76.

M78.

M78.

J77a.

J74.

M78.

M78.

K73b.

M78.

SH73.

M78.

M78.

K73b.

SS77.

T73.

SH73.

SH73.

SS77.

S76.

H73.

M78.

M78.

M78.

M78.

M78.

M78.

SH74.

Absolute age

Absolute ages determined by the 40Ar-39Ar method on materials 

obtained from the Cayley and Descartes Formations do not clearly 

indicate a difference in provenance consistent with the suggested 

lithologic differences. Pre-Imbrian ages (older than 3.85 aeons) are 

more common than Imbrium-basin ages. The following interpretation suggests an age for the Nectaris basin (James, 1981).
9. NECTARIAN SYSTEM 169

The dated material that is most critical with respect to a Nectaris-basin age was derived from the feldspathic fragmental breccia 

deposits penetrated by North Ray Crater (table 9.1). Coarse fines (2-4 

mm) of melt type FM collected by rake from the North Ray regolith 

yielded a cluster of six dates between 4.09 and 4.14 aeons and one 

younger date of 3.96 aeons (Maurer and others, 1978). Because the 

young fragment (67603,1-MD) is compositionally and texturally 

identical to the apparently older clasts, 3.96 aeons is a maximum age 

for the FM-type materials (James, 1981). The older ages were probably 

inherited from an earlier event and were not reset when the feldspathic fragmental breccia deposits were assembled. Their close clustering suggests an actual event at about 4.1 aeons. Wetherill (1981) 

suggested that the older event was the Nectaris impact, whereas 

Maurer and others (1978) and James (1981) proposed that it was a preNectarian crater impact(s).

The other chronologically significant type of coarse fines derived 

from the feldspathic fragmental breccia deposits is granulitic breccia. 

The ages of six fragments range from 3.89 to 3.96 aeons, a range that 

also includes the age of the youngest FM-type melt (Maurer and 

others, 1978). This range may reflect analytical inaccuracies and may 

record a single metamorphic episode generated by a major source of

near-surface heat, probably a hot basin-ejecta blanket (Warner and 

others, 1977). The geologic setting of the samples suggests the Nectaris ejecta as the source of this heat. The average of the age cluster, 

3.92 aeons, or the youngest age, 3.89 aeons, which is also the age of a 

dark clast and a plagioclase separate from sample 67016 (Turner and 

Cadogan, 1975), may date the Nectaris impact (Jessberger and others, 

1974; Turner, 1977; Maurer and others, 1978; James, 1981).

The wide range of isotopic ages presents obvious problems for 

determining the emplacement ages of the Cayley and Descartes Formations. An age occupies almost every 10-million-year interval from 

4.14 to 3.86 aeons ago, and four ages are younger (table 9.1). The 

inevitably long list of uncertainties includes: (1) analytical errors; (2) 

incomplete reequilibration of argon isotopes by impacts; (3) addition 

or removal of argon by later events; (4) varying degrees of equilibration of clasts and matrices in the same sample, which are commonly 

analyzed together; (5) failure to collect thoroughly reset materials 

that do date an event; (6) uncertainty whether the recorded events are 

the crystallization of plutonic rocks, metamorphism (granulitic breccia), the main emplacing impact, or minor later impacts; and (7) 

uncertainty whether the photogeologic unit underlying the sampling 

stations was actually sampled (especially questionable for the

FIGURE 9.8. Geologic map of part of nearside (Wilhelms and McCauley, 1971), including Apollo 11 (see chap. 11) and Apollo 16 landing sites. Im, mare basalt of Imbrian age (Apollo 11); Ip, 

plains material of Imbrian age (Cayley Formation; Apollo 16); Ihf, hilly and furrowed material of Imbrian age (Descartes Formation; Apollo 16). Arcs indicate basin rings as interpreted at time of 

map's preparation.
170 THE GEOLOGIC HISTORY OF THE MOON

Descartes Formation). The problems are illustrated by sample 60315 

from the Cayley, which yielded ages of 3.91 to 3.97 aeons (table 9.1). 

Therefore, the dating may not be accurate; the ages listed in table 9.1 

may date no real event at all.

These ages, however, are at least consistent with formation of the 

Nectaris basin at about 3.92 ± 0.03 aeons, and this date is tentatively 

accepted in this volume (fig. 8.16). The basin is unlikely to be younger 

than that age range because no samples of granulitic breccia or typeFM melt rock recovered from the Descartes Formation have yielded 

younger ages (James, 1981). It is unlikely to be older than that age 

range because some type-FM melts were reset during that time 

interval, because metamorphism of granulitic breccia requires a 

major heat source not available after the Nectaris impact, and 

because the impact-rate curve (fig. 8.16) would have to bend to fit both

the available Imbrian ages and an older age for Nectaris (see chap. 8). 

The ages older than 3.92 ±0.03 aeons (table 9.1) date pre-Nectarian 

events, if they date any discrete events. The few younger ages may 

have been introduced or partly reset during the Imbrium impact by 

the mechanisms described in chapter 10.

CRISIUM BASIN

General features

Crisium has already been referred to in connection with basin 

topography, ring origin, mare-basalt thickness, and structural modification because quantitative data were obtained by overflights of

 y "f \my: *-* *.

: % W^' : ' *"*' 

% *

FIGURE 9.9. Region of Apollo landing 16 site. Facies of the Descartes Formation in the Smoky Mountains is furrowed radially to 

Imbrium basin; facies in Stone Mountain is furrowed transversely to Imbrium and radially to Nectaris basin. Lines, astronaut 

traverses; numbered circles, sampling stations. Stereoscopic pair of Apollo 16 frames P-4558 (right) and P-4563 (left).
9. NECTARIAN SYSTEM 171

Apollo 15 and 17 instruments (pi. 2; figs. 4.9, 5.23). As in many large 

basins, the identity of its topographic rim is unclear. The massive 

mare-bounding rim, about 635 km in diameter, would seem to be the 

logical choice as the topographic rim and, therefore, the boundary of 

excavation. It probably does bound a deep cavity. Incomplete but 

locally conspicuous arcs, however, form a near-circle, 1,060 km in 

diameter, whose massifs are as high as or higher than those of the 

635-km-diameter ring (figs. 4.3Z, 4.9, 9.11). These arcs are locally 

rimlike (raised crest, steep inner slope, gentler outer flank) and mark 

the inner boundary of conspicuous secondary chains (figs. 9.3, 9.5, 

9.6). Thus, the outer zone may have been partly excavated. A large 

basin excavation is consistent with the many topographic complexities within the 1,060-km periphery, and that diameter is favored here 

as the excavation diameter of Crisium.

Another fact favoring a large excavation is the probable great 

extent of the ejecta and secondary craters. Crisium-secondary craters 

extend as far as 900 km northward and 1,100 km southward of the 

basin center (pi. 7; figs. 9.3-9.6). One group of secondaries seems to 

be superposed on the crater Zeno, which, in turn, is superposed on the 

Humboldtianum basin (fig. 9.4). The extent of the continuous ejecta is 

less clear. Hummocky terrain occupies some sectors close to the basin, 

as does the somewhat similar, though topographically sharper, 

Montes Rook Formation of Orientale (fig. 9.11; Scott and Pohn, 1972; 

Scott and others, 1972; Head, 1974b; Wilhelms, 1973,1980). Crisium 

ejecta probably covers terrain in the poorly photographed zone east of 

Crisium, which is smoother than the adjacent more heavily cratered 

pre-Nectarian terrane (fig. 9.4). An extent of Crisium ejecta as great 

as 1,600 km east of the Crisium center is suggested by a concentration 

of radially arranged craters 10 to 20 km in diameter.

B

FIGURE 9,10. Feldspathic fragmental breccia (sample 67015) collected from rim of North 

Ray Crater. 

A. Whole sample. Melt-rich clasts are dark; light material includes feldspathic matrix and

clasts. Sides of cube (letter N), 1 cm. Courtesy of Lunar and Planetary Institute,

Houston, Tex. 

B. Thin section of sample 67015,88. Senate-textured matrix consists of cataclastic plagioclase and minor pyroxene. Clasts include glassy-matrix breccia (g), clast-laden,

crystalline impact-melt breccia (m), and monomineralic plagioclase (p). Plane-polarized

light; field of view, 1.5 mm. Courtesy of P.D. Spudis,

The observed pattern of Crisium ejecta is three-pronged, rather 

like that of the crater Messier (fig. 3.1 LA). A further similarity with 

Messier is the elliptical outline of Crisium, elongate approximately 

east-west. The east sector of the mare-bounding ring is almost missing (fig. 4.3Z). An oblique impact from the west, possibly by a fragmented body, is suggested as the cause of these asymmetries (table 

4.4).

Luna 20 samples

The Luna 20 mission returned a small core of regolith material 

from a part of the Crisium flank that is outside the mare-bounding 

rings and that is here considered to lie inside a buried part of the main 

rim (pis. 3, 7; fig. 9.5). Most of the fragments are of aluminous ANT 

rock not unlike the Apollo 16 suite; some are of low-K KREEP composition (Prinz and others, 1973; Taylor, G.J., and others, 1973). They 

have been dated at 3.84 ±0.04 aeons (Podosek and others, 1973), an 

age consistent either with a late Nectarian relative age of Crisium or 

with an Imbrian age of superposed Imbrium-basin or Imbrian-crater 

material. In view of the small size of the sample and the difficulties 

encountered in dating the larger Apollo 16 and 17 sample suites, no 

firm conclusions should be based on this age despite its apparent 

analytical precision.

SERENITATIS BASIN

General features

Although Mare Serenitatis is one of the Moon's most conspicuous 

circular maria (fig. 1.7), the Serenitatis basin is indistinct and was not 

even listed as a basin by Hartmann and Kuiper (1962). Its west sector 

has been nearly destroyed by the Imbrium basin, and its north rim is 

covered by hummocky Imbrium ejecta (Alpes Formation; figs. 1.7, 

10.12). The map pattern of the north rim and the presence of two 

mascons suggest that Serenitatis consists of a smaller northern basin 

and a larger southern basin (Baldwin, 1963, p. 320; Scott, 1972a, 1974; 

Reed and Wolfe, 1975; Wolfe and Reed, 1976; Head, 1979c). Although 

two independent impacts have been proposed, the considerable evidence for multiple simultaneous crater impacts (see chap. 3) and 

basin impacts (Humboldtianum, Moscoviense, possibly Crisium) suggests such an origin here as well. Montes Haemus, the south 

topographic rim of Serenitatis or "South Serenitatis," is conspicuous, 

though covered by lineate and smooth Imbrium ejecta (figs. 1.7,9.12).

The east sector is of most interest because part of it was sampled 

by Apollo 17. The terra east of Mare Serenitatis consists of angular 

massifs that are partly concentric and partly radial to the mare (fig. 

9.13). A general ring form is visible, but no single ring stands out. The 

largest massifs surround the Taurus-Littrow valley, in which Apollo 

17 landed. The vague ring containing these massifs resembles the 

outer Rook ring of Orientale (Head, 1974b, 1979c; Reed and Wolfe, 

1975; Solomon and Head, 1979). Although parts of an indistinct ring 

east of the Littrow massifs and ending near the crater Vitrivius (fig. 

9.3) are higher than the massifs (Head, 1979c), this "Vitruvius front" 

is not sharply delineated like Montes Cordillera of Orientale. Therefore, eastern Serenitatis resembles the irregularly structured western Orientale more closely than it does the regularly ringed eastern 

Orientale (fig. 4.1). As is common for large basins, the east boundary 

of the excavation cavity of Serenitatis is thus poorly defined.

One possibility is that the Vitruvius front is the excavation 

boundary; small amounts of material may have been ejected from the 

zone between the Littrow and Vitruvius rings, as proposed above for 

Crisium and in chapter 4 for other large basins. The most straightforward interpretation, however, is that the excavation boundary passes 

along the east edge of Mare Serenitatis through the largest massifs of 

the Littrow ring (figs. 5.17, 9.12, 9.13). I tentatively favor this interpretation despite the similarity of the Littrow ring to the outer Rook 

ring, which I believe lies inside the basin excavation. A diameter of 

about 740 km 190 km smaller than Orientale fits the largest 

Littrow and Haemus massifs. This identification of the rim is 

strengthened by Whitaker's (1981) interpretation of the Vitruvius ring 

as a part of the middle ring of the Procellarum basin (pi. 3; fig. 9.11) 

an interpretation that also explains why the terra east of Serenitatis 

is so irregular and why no relation of rings to basins is obvious.
iS-W&Wtrf'V fj,-' - *. HI-- fSi 'Ifsfl v ' . . f"/(Bi- w

*^ \:^J 1 *>, n '.4 .> 4»,^r '
9. NECTARIAN SYSTEM 173

Relative age

Serenitatis was long considered to be an old basin because of its 

highly degraded appearance (Stuart-Alexander and Howard, 1970; 

Head, 1974b). Numerous degraded craters in the terra east of the 

mare were thought to be superposed on the basin (figs. 1.7,5.17, 9.11, 

9.14). The interpretation of these craters and the position of the rim 

are critical to the important question of the relative age of the 

Serenitatis basin. The clustering and nearly subequal size of some of 

the craters suggest an origin as secondary craters of the Imbrium 

basin (for example, Littrow; figs. 9.11,9.13; Wilhelms, 1980). Almost all 

the other craters resemble the degraded prebasin craters visible close 

to the rims of the Orientale and Nectaris basins (figs. 4.4, 7.9). 

Therefore, these craters are older than Serenitatis, provided that the 

Serenitatis rim lies west of the craters near the edge of the mare. Only 

three primary craters larger than 20 km in diameter visible east of 

Serenitatis are clearly younger than Serenitatis and older than 

Imbrium (Kirchhoff, Maraldi, Romer A; figs. 9.11, 9.14). Le Monnier 

(61 km, 27° N., 31° E.; fig. 6.25) also seems likely to lie in this

FIGURE 9.11. Mare Cnsium and circum-Crisium terrain. Primary-impact craters common 

to both scenes include Macrobius (M; 64 km, 21 ° N., 46° E., Lower Imbrian) and Proclus 

(P; 28 km, late Copermcan; compare fig. 3. 33). Trough concentric with Crisium is west (left) 

of Macrobius. Hummocky material, possibly ejected from Crisium or Serenitatis basin, is 

between Proclus and Mare Tranquillitatis (MT).

A. Nectarian primary craters include Maraldi (Ma; 40 km, 19° N., 35° E.) and Tisserand 

(T; 37 km); Romer (R; 40 km) is Copernican. Craters labeled "c" are believed to be 

secondary-impact craters of Serenitatis or Imbrium basin because of compound shape, 

alignment with basin radials, varying, diverse morphology of ridges, grooves, and 

mounds, and concentration in size range 10-20 km. Deposits of Imbrian age (Ip, Itl) 

and fine-scale morphology indicated by other symbols (1, It, pi, t), possibly related to 

Imbrium basin, are superposed on rugged features of Crisium and Serenitatis basins 

(Wilhelms, 1980). SA, Sinus Amoris; MT, Mare Tranquillitatis; m, other mare 

patches. Rugged terra along north-south midline of photograph is probably part of 

Procellarum basin (pi. 3). From Wilhelms (1980, fig. 19). Apollo 17 frames M-195 

(right) and M-302 (left).

B. Crisium-basm massifs, continuous in west and discontinuous (but poorly illuminated) in 

east. Eratosthenian craters Peirce (Pe; 19 km, 18° N., 53.5° E.) and Picard (Pi; 23 

km, 15° N., 55° E.) penetrate mare (see chap. 5). Mosaic of Apollo 17 frames 

M-274, M-278, M-281, M-286, M-289, M-293, and M-294 (from right to 

left).

intermediate stratigraphic position, but it is only a "ghost" crater 

whose age is uncertain.

Several additional lines of evidence are consistent with a relatively young age of the Serenitatis basin. (1) The near-destruction of 

the west sector by the Imbrium basin, a clear indicator of relative age, 

has no significance for absolute age; the two basins could be almost 

contemporaneous. (2) Serenitatis secondaries may be superposed on 

Crisium ejecta and the Crisium rim (figs. 9.14, 9.15; Wilhelms, 1976). 

(3) No Crisium ejecta is texturally evident near the Taurus-Littrow 

massifs (figs. 9.11, 9.13); if Serenitatis is not the younger of the two 

basins, the Crisium pattern of ejecta must have been highly asymmetric (which, in fact, is possible, as shown by the analogy with 

Messier). (4) The substantial thickness of the basalt in Mare 

Serenitatis indicates little filling or isostatic rebound since formation 

of the Serenitatis basin. (5) Apollo 17 astronaut R.E. Evans described 

the massifs of Serenitatis as much fresher and more boulder-strewn 

than those of Crisium (Evans and El-Baz, 1973, p. 16); although 

morphology has not proved to be a very precise age indicator for basin 

ages, at least this evidence helps counter reliance on Crisium's apparently fresh overall morphology as an indicator of its youth.

In summary, the roughness of the terra east of Serenitatis is due 

to Procellarum and Serenitatis massifs, Imbrium secondaries, 

Imbrium deposits (fig. 9.11), Serenitatis secondaries, .and preSerenitatis craters. None of these features establishes an old age for 

the basin. A relatively young Nectarian age for Serenitatis is favored 

here.

Petrology of Apollo 17 samples

The Apollo 17 landing site (figs. 9.16, 9.17) was chosen mainly to 

provide samples of terra outside strong influence of the Imbrium 

basin and in a known geologic context (Scott and others, 1972; Hinners, 1973). Petrologists, geochemists, and geochronologists held out 

the hope that primitive terra would be sampled. Geologists wanted 

information on basins other than Imbrium, on basin-formation processes in general, and, as always, on dates of specific events. The mare 

and dark-mantling materials were also to be sampled (see chaps. 5, 

11), but Serenitatis-basin materials were the main objective.
174 THE GEOLOGIC HISTORY OF THE MOON

Most of the sampled material is fragment-laden impact melt with 

a wide variety of textures but a narrow range of chemical compositions. Matrices are rich in low-K KREEP (Ryder and Wood, 1977; 

Winzer and others, 1977). Their prevalent textures are the types 

generally called poikilitic, ophitic, or subophitic, which indicate igneous-like crystallization from hot melts (fig. 9.18; Simonds and others, 

1974). Other textures are aphanitic, crystallized from rapidly cooling 

melts (fig. 9.19; James and others, 1978). Most lithic clasts are of 

granulitic breccia, but clasts of Mg-suite cumulates are unusually 

abundant (for example, norite sample 77215 from a 2-m clast, 156-g 

troctolite sample 76535, and dunite sample 72415-72418; table 8.4). 

Many of the Mg-suite cumulates probably came from a single layered 

pluton (Dymek and others, 1975; Jackson and others, 1975; James, 

1980). Multiple samples were collected from five melt-rock boulders, 

three of which were sampled at station 2 on a colluvial apron at the 

base of South Massif, and one each at stations 6 and 7 on basal aprons 

of North Massif (Dymek and others, 1976). All but one of the boulders 

left telltale tracks indicating their source ledges, which crop out twothirds of the way up the massifs (fig. 9.20). Smaller samples of fragment-laden melt rock were taken from station 3, on a light-colored 

landslide derived from South Massif (figs. 6.1, 9.13, 9.16, 9.19). Some

material of an additional terrain, the "sculptured hills," may have 

been collected at station 8. The sculptured hills have a hummocky 

surface resembling the Montes Rook Formation at Orientate or the 

Alpes Formation of Imbrium. The amounts collected are too small, 

however, to assure that they fairly represent that unit (Wolfe and 

others, 1981).

Emplacement process

The fragment-laden melt rocks have generally been interpreted 

as part of the Serenitatis-basin ejecta (Wood, 1975d; Winzer and 

others, 1977). This interpretation is consistent with the similarity of 

the massifs to those of Montes Rook, many of which are surrounded 

by the probably melt-rich Maunder and Montes Rook Formations (see 

chap. 4). The ledges presumably compose either the massifs proper or 

a thick superposed deposit. Superposition of Serenitatis ejecta is 

consistent with photogeologic observations here and at Orientale and 

Imbrium (figs. 4.4E, 10.7). However, the abundance of Imbrium-radial 

material in the vicinity is consistent with superposition of Imbrium 

ejecta on the Serenitatis massifs (fig. 9.13; Wilhelms, 1980).

E R E\N I T A

FIGURE 9.12. Part of geologic map of nearside ( Wilhelms and McCauley, 1971), including Apollo 17 landing site (arrow). Site is on unifcCId, dark-mantling material of Imbrian to Copernican 

age, shown by mission results to be Imbrian (see chap. 11). Adjacent massifs (dark color, plr, pre-Imbrian rugged material) are parts of Serenitatis-basin ring, interpreted here as main basin ring 

(table 4.1), which continues as Montes Haemus near left edge of map. Crisium-basin massifs (same color and symbol) are at right edge of map (compare fig. 9.11). Possible rings of Tranquillitatis 

basin are also shown. Ages of mare units (Em, Im) have been partly revised since publication of map (see chap. 11). Lower left corner of figure adjoins upper right corner of figure 9.8.
9. NECTARIAN SYSTEM 175

FIGURE 9.13. Apollo 17 landing site (lower arrowhead: compare figs. 2.5D, 6.2B). South 

Massif is below and left of arrowhead; North Massif is above arrowhead, d, dark-mantling 

material. Units radial to, and believed related to, Imbnum basin (It, Itl, It, pi; fig. 9.1 \A; 

Wilhelms, 1980) are superposed on much of rugged terrain but not on massifs. Lineations 

parallel to white line are not obviously related to any basin. L, Littrow, possible Imbnumbasin secondary crater. Ma, Nectarian crater Maraldi (compare fig. 9.11 A); V, Upper 

Imbrian crater Vitruvius. Ragged line of massifs extending between Vitruvius and upper 

arrowhead is "Vitruvius front" (Head, I974b)i From Wilhelms (I960, fig. 18); overlaps 

with figure 9.11 A. Stereoscopic pair of Apollo 17 frames M-444 (right) and M-446 

(left)

FIGURE 9.14. Terrain northeast of Serenitatis (arrow S), northwest of Crisium (arrow C), 

and south-southwest of Humboldtianum (arrow H). Craters above (north of) black line are 

probable Serenitatis-basin secondaries, as indicated by radiality of clusters and associated 

linear deposits. Lineate terrain south of black line is not closely radial to any basin but is 

subradial to Imbrium and Crisium. Humboldtianum-basin secondary origin of irregular chain 

or compound craters around arrow H (including Berzelius, B) is suggested by approximate 

radiality. Kirchhoff (K; 25 km, 30° N., 39° E.) and Romer A (R; 35 km, 28° N., 37° E.) 

are probably the only post-Seremtatis, pre-Imbnan primaries in view. From Wilhelms 

(1976, fig. 8). Orbiter 4 frame H-74.

FIGURE 9.15. Southwest rim of Crisium basin. Large crater groups aligned along arrow are 

probably secondaries of a basin; they are 1,000 km and 2,000 km from centers of Serenitatis 

and Imbrium basins, respectively. Degraded morphology favors a Serenitatis source; superposition then shows that Crisium is the older basin. P, Proclus (28 km; compare fig. 9.11). 

Apollo 11 frame H-6230.
176 THE GEOLOGIC HISTORY OF THE MOON

FIGURE 9.16. Premission geologic map of Apollo 17 landing site by B. K. Lucchitta (in Scott and others, 1972). Landing took place 

within proposed circle. Sampled units, in order of decreasing age: pltm, pre-Imbrian terra-massif material; Iplh, Imbrian or preImbnan hilly material ("sculptured hills"); Ips, Imbrian smooth plains material (subfloor basalt; see chap. 11); EC, Eratosthenian 

crater material; Cce, craters now believed secondary to Tycho (chap. 13); Cb, bright Copernican material (light mantle or landslide); 

dh, dark-halo craters (now known to be Copernican impact craters; see chap. 13).

FIGURE 9.17. Astronaut traverses and locations of sampling stations at Apollo 17 landing site.
9. NECTARIAN SYSTEM 177

The general chemical uniformity of the sample suite has led most 

of the authors cited above to conclude that the Apollo 17 samples were 

emplaced in a single impact, whereas others cite differences in chemistry, clast content, or apparent assembly process as evidence for 

multiple impacts (Ryder and others, 1975b; Chao and others, 1976; 

Minkin and others, 1978; Spudis and Ryder, 1981). The argument that 

impact melting and mixing of facies homogenize target materials of 

terrestrial craters (Grieve and others, 1974, 1977; Simonds and others, 1976a, b; Phinney and Simonds, 1977; Grieve and Floran, 1978) is 

commonly used to defend the multiple origins of lunar sample suites

FIGURE 9.18. Vesicular poikilitic impact melt (sample 76215) from station 6 boulder, 

North Massif, Apollo 17 landing site. Probably part of Serenitatis melt-rich ejecta. 

A. Hand specimen. Sides of cube, 1 cm. NASA photograph.

B. Thin section of sample 76215,55. Poikilitic texture is well developed in pyroxene 

oikocryst (bottom center); incipient subophitic texture is around vesicles (black). Crossed 

polarizers; field of view, about 2 mm. Courtesy of RD. Spudis.

FIGURE 9.19. Slab cut through center of aphanitic, clast-laden impact melt (sample 73255) 

from station 3, Apollo 17 landing site (James and others, 1978). Vesicular rind and 

nonvesicular interior suggest origin as melt bomb.

evaluated as "different." Extensive mixing does occur, but not total 

homogenization. Only the zone near the projectile is totally melted; in 

impacts of large, relatively slow projectiles, much additional material 

is incompletely melted and partly preserves the inhomogeneities of 

the target (see chap. 3\ Furthermore, segregation of constituents 

according to density or volatility might create diversity during impact 

ejection and deposition. Therefore, the observed diversity is consistent with a Serenitatis origin for all rock-size samples collected by 

Apollo 17.

Absolute age

A better understanding of the geology and the availability of 

suitable materials have led to a clearer picture of the Apollo 17 

absolute ages than for the Apollo 16 samples. The Apollo 17 ages also 

scatter considerably (table 9.2; Jessberger and others, 1977b, p. 2578; 

1978; James and others, 1978; Wolfe and others, 1981) because of 

incomplete outgassing and isotopic reequilibration, as is usual in 

terra breccia deposits. In fact, dozens of sophisticated Ar-Ar analyses 

by stepwise heating of the same breccia sample (73215) yielded ages 

spread over 300 million years (Jessberger and others, 1977b). The 

laser-melting technique of determining the Ar isotopes resolved 

apparent age differences between parts of the same sample only 0.1 

mm apart (Miiller and others, 1977). However, the same careful

TABLE 9.2. 40Ar-39Ar ages of Apollo 17 terra samples

Stations (Wolfe and others, 1981): 2, at contact between light mantle (landslide) and 

north base of South Massif; 3, on light mantle near lower slope of Lee-Lincoln scarp; 

6, on south slope of North Massif; 7, near base of North Massif, 0.5 km east of station 6.

Samples (Wolfe and others, 1981): All are from boulders or are impact-melt samples 

larger than 1 cm. 722-, station 2, boulder 1; 724-, station 2, boulder 3; 732-, station 3, rocks collected from nm of 10-m-diameter crater; 760-, station 6, subboulder 

5; 762-, station 6, subboulder 1; 763-, station 6, subboulder 2; 77-, station 7; numbers 770-, 771-, and 772- each refer to a separate sample from a 3-m boulder.

Lithologies: Samples 72255, 73215, and 73255 are light gray, others are "blue-gray" or 

"green-gray"; the various colors are commonly thought to derive from different source 

beds.

Ages: Calculated using International Union of Geological Sciences (IUGS) decay constants (Steiger and Jager, 1977).

References: CT76, Cadogan and Turner (1976); E79, Eichhorn and others (1979); H78, 

Huneke (1978); L75, Leich and others (1975); M77, Mliller and others (1977); J77b, 

Jessberger and others (1977b); J78, Jessberger and others (1978); P75, Phinney and 

others (1975); S79, Staudacher and others (1979); S78, Stettler and others (1978); 

TC75, Turner and Cadogan (1975)

Station Sample Lithology Age 

(aeons) References

2 72255 Bulk fragment-laden melt rock 3.95+0.03 L75.

72435 do 3.86±0.04 H78.

3 73215 Felsite clast 3.86±0.04 J77b, J78.

Melt-derived groundmass 3.91±0.03 M77, J77b.

Bulk fragment-laden melt rock 3.92±0.04 J77b.

Clasts of fragment-laden melt rock 3.96-4.17±0.04 J77b.

73235 Bulk fragment-laden melt rock 3.90-o;°8 TC75.

do - 3.92±0.04 P75.

73255 do 3.86±0.04 E79.

do 3.88+0.03 J78.

do 3.89-3.92+0.02 S79.

Felsite clast 3.89+0.03 S79.

73275 Bulk fragment-laden melt rock 3.90+0.05 TC75.

6 76015 Poikilitic melt rock 3.87+0.04 CT76.

Mineral concentrates 3.86+0~.04 to CT76.

3.90+0.06

76215 Poikilitic melt rock 3.88+0.04 CT76.

76235 Clasts 3.87-3.89+0.06 CT76.

76255 do 3.96±0.04 CT76.

76275 Melt rock 3.96+0.04 CT76.

76295 do 3.89-3.90±0.04 CT76.

76315 Clasts 3.91-3.92±0.04 CT76.

Melt rock 3.92+0.04 CT76.

7 77075 Melt-rock dike in large norite clast 3.91+0.04 S78.

77115 Bulk fragment-laden melt rock 3.84±0.03 S78.

77135 Poikilitic melt rock 3.83±0.04 S78.

77215 Plagioclase from norite clast 3.92+0.03 S78.
178 THE GEOLOGIC HISTORY OF THE MOON

application of sophisticated Ar-Ar techniques that disclosed the problems also probably disclosed the youngest materials that have been 

completely reequilibrated the aphanitic groundmass of a felsite 

clast from station 3 (sample 73255; Eichhorn and others, 1979; 

Staudacher and others, 1979). The age of this material is 3.87 ±0.04 

aeons.

The applicability of these precise measurements to Serenitatis 

has been questioned because of the uncertain provenance of the beststudied samples, which were collected at station 3 from float on the 

landslide and not from a known source bed, as were most of the 

boulders. This provenance might suggest superposition on the massifs of a post-Serenitatis deposit, possibly from Imbrium. Sample 

73255 resembles a bomb that could have arrived as an isolated ballistic projectile from Imbrium or elsewhere (Spudis and Ryder, 1981). 

Spudis and Ryder (1981) cited the homogenization model as showing 

that the differences between the station 3 samples and the boulder 

samples are too great for both groups to have been assembled by the 

same impact.

The consortium responsible for studying the station 3 materials, 

however, considered a genetic relation of those materials to the boulder samples to be established by similar compositions (James, 1976; 

James and others, 1978). Moreover, one boulder sample from station 2, 

72435, is coeval with the youngest age from sample 73255,3.86 ± 0.04 

aeons (Huneke, 1978). Other boulder samples yield older ages, but so 

do parts of the aphanatic samples, as noted. Therefore, an age of 3.86 

or 3.87 aeons is accepted here as that of the Serenitatis basin. This age 

is consistent with tentative evidence for the relative youth of 

Serenitatis within the Nectarian System.

OTHER

BASIN AND CRATER 

MATERIALS

Basins

Twelve basins, including Nectaris and Serenitatis, are considered to be of probable Nectarian age (pi. 7; table 9.3). Direct 

superpositional relations establish many age relations between preNectarian, Nectarian, and Imbrian basins and among the Nectarian 

basins. An especially clear sequence is Apollo (pre-Nectarian)- 

Korolev (early Nectarian)-Hertzsprung (late Nectarian)-Orientale 

(Imbrian) (figs. 9.2LA, C). Two, possibly three, age groups of Nectarian 

basins are recognized (table 9.3). Within a single group, the statistical 

error ranges of the counts of superposed craters overlap (fig. 9.22).

The older group consists of the Nectaris, Mendel-Rydberg, 

Moscoviense, and Korolev basins. The Nectaris deposits, which constitute the base of the Nectarian System, are very extensive and the 

superposed craters correspondingly numerous. The density of 79 

craters at least 20 km in diameter per million square kilometers (fig. 

9.22) can thus be taken as expressing the base of the system in sizefrequency terms.

The deep burial by Orientale deposits of the Mendel-Rydberg 

basin has reduced the crater population on that basin (figs. 4.1, 4.3R); 

because the curve lies close to that for Nectaris, the Nectarian age of

**' -", --Source: of statfoh G'boulder'trlirk

 ' ; ' "

FIGURE 9.20. Boulder tracks on North Massif (from Wolfe and others, 1981, fig. 145).
9. NECTARIAN SYSTEM 

TABLE 9.3. Nectarian basins

[Basins are ranked in order of increasing age; ranking is by stated density of craters or by criteria given in text; ranking of groups is 

more certain than ranking within a group. Diameters are from table 4.1. Parenthetical basins, existence not definitely established; 

parenthetical areas, buried; parenthetical densities, poor sample. Superposed craters are at least 20 km in diameter. Underlying and 

overlying basins determined from directly observed superpositional relations; n.d., no relations detected]

179

Age 

group

°f sr Superposed craters

^ Density

Number / , ,,'fi" 7"?\ (craters/ .. r . (10b km z ) ms k 2 1 Figure

Superposed basin

Name Figure

2? Bailly2? (SikorskyRittenhouse),

2 Hertzsprung 

300

570

(0.096)

51 .883

(31)

58

n.d------------------

South Pole-Aitken---

Korolev 9.21 

Birkhoff 8.14, 

Apollo 8.15, 

Lorentz 8.14 

Coulomb-Sarton 8.14, 

(Grissom-White) 8.15 

South Pole-Aitken 9.21

Orientale 1.9. 

SchrOdinger 1.5.

Orientale

9.25

9.21

9.25

2? Serenitatis2 Crisium740

1,060

.108

45 .843

2 Humorum----- 

2 Humboldtianum820 24 .428 

avg 700 32 .515

(83)

53

56

62

Crisium 9.11, 9.14, Imbrium9.15. 

Humboldtianum 9.14

8.14, 9.21, 

9.25.

1.7, 9.11.

Humboldtianum - 9.4

Smythii 1.5

Lomonosov-Fleming - - 

Fecunditatis 9.5

Tranquill itatis 9.11

Marginis 9.4

Nubium 8.10

Moscoviense 9.23

Imbrium 9.11. 

Serenitatis 9.11, 9.14, 

9.15.

Orientale, Imbrium 9.23.

Imbrium- 9.4, 10.6.

Serenitatis 9.14.

Crisium 9.4.

Freundl ichSharonov, 

KeelerHeaviside.

1 Korolev 440 88 1.113 79 FreundlichSharonov, 

Apollo, South 

Pole-Aitken.

Sharonov. 

1? Mendel -Rydberg 630 18 (0.247) (73) (Pingre-Hausen) ,

South PoleAitken.

1 Nectaris 860 102 1.286 79 Mutus-Vlacq 

Tranquil litatis 

9.21 Orientale, 9.21. 

Hertzsprung.

8.12, 9.2 Imbrium 3.10, 9.2.

9.2 

9.2, 11.1 

1.5, 9.6 

8.12

Mendel-Rydberg is doubtful. A small patch of Mendel-Rydberg ejecta 

apparently is exposed southwest of the basin, near the crater Fizeau 

(pi. 7; Wilhelms and others, 1979).

Moscoviense and Korolev ejecta or secondary craters blanket preNectarian craters over distances averaging a basin diameter and 

extending, in salients, more than two basin diameters (pis. 3, 7; figs. 

9.21C, 9.23). Large tracts can thus be dated relative to these basins. 

The southwestern ejecta of Korolev is distinctly lineate (fig. 9.2LD). 

Radial alignment of some fresh secondary chains inside Korolev 

suggests an Orientale origin, and so the adjacent plains may also be of 

Orientale origin (fig. 9.215). The crater densities on Korolev and 

Nectaris are similar, but they overlap only where the Nectaris curve is 

kinked. Although more craters per unit area are superposed on 

Moscoviense than on Nectaris, most of the surplus craters are small. 

The curves as a whole thus indicate that both the Korolev and Moscoviense basins are younger than Nectaris.

The definite basins of the younger group are Humboldtianum, 

Humorum, Crisium, and Hertzsprung. Bailly and Serenitatis are 

probably members of this group, although they cannot be dated reliably by superposed craters. Humboldtianum ejecta covers much of 

the northern part of the east lunar limb (fig. 9.4). Conspicuous radial 

lineations and secondary craters extend 600 km south of the 

topographic rim, which is a clearly defined ring 300 km in radius in 

this sector. An exterior arc 600 km in radius is covered by the ejecta 

(Wilhelms and El-Baz, 1977). Humboldtianum is a double or triple 

basin whose parts form a figure-8 (pi. 3; fig. 9.4; Lucchitta, 1978). A 

ring 150 km in radius is conspicuous in both parts of this figure-8; a 

poorly defined ring between the topographic rim and the 150-kmradius ring is visible in the southern part. The northern part of the 

figure-8 consists of irregular uplifts that do not clearly define rings

(Lucchitta, 1978). Crater counts (fig. 9.22) are consistent with the age 

relation demonstrated by superposition of Crisium secondaries on 

Humboldtianum (fig. 9.4).

The Humorum basin is also irregular in most sectors (figs. 4.3U, 

9.24). The most complete ring, 440 km in diameter, bounds Mare 

Humorum (table 4.1). A conspicuous and rimlike scarp almost twice 

as large (410 km radius) and resembling Montes Cordillera lies south 

of this mare-bounding ring (fig. 9.24S). This scarp is the topographic 

rim in its sector and so is here considered to be the boundary of 

excavation, despite lesser expression in other sectors. Although an 

820-km diameter is, furthermore, consistent with the considerable 

extent of conspicuous deposits and secondary craters southwest of the 

basin (fig. 9.245), this large size cannot be confirmed because 

Humorum deposits are buried in other sectors. Sufficient numbers of 

superposed craters are exposed to date Humorum as part of the 

younger group but not to establish its exact place in the group. 

Humorum and Crisium appear to be nearly contemporaneous.

The fresh-appearing Hertzsprung helps clarify the geology of the 

other basins (fig. 9.25). The secondary craters extend as far as 1,300 

km north of the 570-km-diameter ring. This relation and the nearcompleteness and rimlike morphology indicate that the 570-kmdiameter ring bounds the excavation (see chap. 4). The interior rings 

probably formed by floor deformation by the oscillatory-uplift mechanism (stage E, fig. 4.12). Hertzsprung may be the youngest Nectarian 

basin. Although the ends of the cumulative size-frequency curves 

representing small craters lie above those for Crisium and Humorum, 

the large-crater end lies below (fig. 9.22). The large-diameter end is 

considered to be more reliable because some small craters were probably missed in the counts of poorly photographed parts of Crisium 

and Humboldtianum.
180 THE GEOLOGIC HISTORY OF THE MOON

The craterlike Mendeleev basin is surrounded by lineate ejecta 

and secondary craters in the well photographed southern peripheral 

zone (fig. 9.26A). Because of this morphologic freshness, it is mapped 

as late Nectarian on plate 7. Fairly distinct crater counts suggest, 

however, that Mendeleev may be intermediate in age between the two 

main groups.

Several factors limit the further resolution of Nectarian age 

differences by crater counts. The small Bailly basin, the questionable 

Sikorsky-Rittenhouse basin, and the extensively buried Serenitatis 

basin present insufficient exposed surface to be dated by superposed 

craters. Even for the well-exposed large basins, crater counts are 

statistically no more definitive than for pre-Nectarian basins because 

relatively few craters are available for dating the Nectarian surfaces. 

The statistics of some of these counts could be improved by adding 

craters smaller than the 20-km-diameter cutoff used to prepare figure 9.22, but the necessary primary-secondary and superposed-buried distinctions cannot be made for other, poorly photographed 

basins.

The inability of these criteria to distinguish between some basins 

may reflect natural groupings. Wetherill (1977a) suggested that large 

impacts might cluster in time because related projectiles derived by 

breakup of a single large body from the outer Solar System are 

consumed by collisions with the Moon and terrestrial planets in 

relatively brief spurts (millions or tens of millions of years). The 

substantial gap in crater densities between the youngest large Nectarian basin (probably Hertzsprung; fig. 9.22) and the oldest Imbrian 

basin (Imbrium) is strikingly suggestive of such clustering of large 

impacts.

Craters

Many craters can be identified as Nectarian (table 9.4), and some 

of these craters can be dated within the system relative to the Nectarian basins (pi. 7). Craters superposed on Nectaris-basin materials 

but nicked by Imbrium-secondary craters provide type examples of 

Nectarian crater morphology (figs. 8.5, 9.1, 9.2, 9.27). Nectarian craters larger than 20 km in diameter have generally deeper floors and 

rougher primary topography of rim and walls (as opposed to superposed topography) than do pre-Nectarian craters. Central peaks also 

are more commonly exposed in Nectarian craters, although they may 

be submerged if subjected to deposits from a nearby crater or basin 

(figs. 9.2IB, 9.27). Some of the largest Nectarian craters have radial 

ejecta textures and secondary craters (fig. 8.11). Apparently, typical 

pre-Nectarian crater morphologies are nowhere superposed on Nectarian surfaces, except possibly for some marginal craters that 

appear to be superposed on outer parts of the Korolev and Moscoviense basins, here considered to be early Nectarian (fig. 9.21D), and 

for a few disputed nearside craters that could be pre-Nectarian (pi. 7). 

Although identifiable Nectarian craters smaller than 20 km in diameter are more numerous than identifiable pre-Nectarian craters of 

this size (Wilhelms and others, 1978), their morphologies vary too 

greatly to be useful for dating isolated craters. Their profiles range 

from cone-shaped to flat-floored, and their numbers appear to dwindle greatly at diameters of about 5 km because of mutual obliteration 

near the steady state (table 7.3).

TABLE 9.4. Representative Nectarian craters

[Cross rules divide diameter ranges mapped differently in plate 7: smaller than 30 km, unmapped; 30 to 119 km, interiors mapped; 120 km and larger, exterior deposits mapped where 

exposed. L, long axis of elongate crater; LI?, possibly Lower Imbrian; pN?, possibly preNectarian]

Crater Diameter 

(km)

Center 

(lat) (long) Figure Remarks

Kirchhoff 25 30° N. 39° E.

RBmer A 

Ti sserand---- 

Maraldi -- 

Crookes D---- 

Young D -- 

Mercator----- 

Playfair -- 

Schiller C~- 

KohlschUtterTimi ryazev--- 

Epigenes -- 

Korolev M - 

Congreve -- 

Congreve U---

35

37

40

41

46

47

48

49

53

53

55

58

58

59

21° N. 

19° N. 

10° S. 

44° S. 

29° S. 

24° S. 

55° S. 

14° N.

6° S. 

68° N.

9° S. 0° 

0.5° N.

37° E.

48° E.

35° E.

163° W.

52° E.

26° W.

8° E.

49° W.

154° E.

147° W.

5° W.

157° W.

167° W.

171° W.

9.14

9.11

9.11,

9.21

8.5

6.9

9.27

7.6

5.2

9.21

10.15

9.21

9.21

9.21

K13

Typical.

Post-Hertzsprung, LI?

Relatively old. 

Relatively young.

Mechnikov----

Lebidinskiy-- 

Sechenov-----

Abulfeda-----

Al -KhwarizmiAdams -- 

SteinheilRheita -- 

Berosus--- - 

Al iacensis--- 

Zhukovskiy--- 

piti scus-----

Mersenius-- 

Metius-------

Rocca-- ----

Inghi rami----

PontecoulantKibal 'chich-- 

Pitatus------

Hevelius-----

Carnot --- 

Al phonsus----

60 11° S. 149° W. 9.21

62 8° N. 164° W. 9.21

62 7° S. 143° W. 9.21

65 14° S. 14° E. 3.16B

65 7° N. 106° E. 8.11

66 32° S. 68° E. 9.5

67 49° S. 47° E. 8.5

70 37° S. 47° E. 8.5, 9.1

74 34° N. 70° E. 9.4

80 31° S. 5° E. 3.26, 9.27

81 8° N. 167° W. 9.21

82 50° S. 31° E. 9.2

84 22° S. 49° W. 9.24A

88 40° S. 43° E. 8.5

90 13° S. 73° W. 4.4F

91 48° S. 69° W. 4.4G, 10.42

91 59° S. 66° E. 9.5, 9.6

93 3° N. 147° W. 9.21

97 30° S. 14° W. 1.8, 6.13D, 10.30

106 2° N. 67° W. 4.4H, 6.21

117 52° N. 144° W. 8.14

119 15° S. 2° W. 1.8, 5.10F, 7.7, 10.27

Pre-Hertzsprung. 

Rel atively young. 

Pre-Hertzsprung. 

Typical. 

Do.

Ty p i c a 1. 

Do.

Typical. 

Relatively old. 

Typical.

Do.

Do. 

LI?, pN?

Typical.

Post-Hertzsprung, LI?

Atypical.

Faulted.

pN?

Chaplygin-----

Endymion------

Cleomedes-- - 

Fleming-------

Albategnius--- 

Neper- ------

LongomontanusHilbertGauss - - 

Cl avius-------

Van de GraaffSchwarzschild124

125

126

130

136

137

145

170

177

225

234(L)

235

6° S.

54° N.

28° N.

15° N.

11° S.

9° N.

50° S.

18° S.

36° N.

58° S.

27° S.

71° N.

150° E.

57° E.

56° E.

110° E.

4° E.

85° E.

22° W.

108° E.

79° E.

14° W.

172° E.

120° E.

.4

1.4 

9.3, 9

9.3

1.2, 8.13, 9 

, 10.27-10.29

9.4 

4.6A

1.2, 1.3, 8.11 

1.2, 4.2D 

4.2H, 4

28

10.36

9.4 

6A

1.4, 3.14B, 10.32 

4.21, 10.6

Typical. 

Mare fill. 

Floor fractured. 

Typical.

Typical.

Floor fractured. 

Typical. 

Double, pN? 

Ring arcs.

Plate 7 shows a total of 560 unburied Nectarian craters larger 

than 30 km in diameter within a total area, covered by both preNectarian and Nectarian deposits, of 19.53 x 106 km2 . As expected 

from stratigraphic relations, the crater density is higher on the preNectarian terrane than on the Nectarian basin deposits (35 versus 22 

craters/106 km2, respectively). The density on the pre-Nectarian terrane is the maximum Nectarian density; extrapolating it to the whole 

Moon (area, 38 x 106 km2 ) indicates that a total of 1,330 craters at least 

30 km in diameter were formed during the Nectarian Period. If the 

number of Nectarian basins is 12, 111 such craters were formed 

between each basin impact; if there are 10 basins (excluding MendelRydberg and Sikorsky-Rittenhouse), 133 such craters per basin were 

formed.
9. NECTARIAN SYSTEM 181

FIGURE 9.21. Stratigraphic relations of lower Nectarian Korolev basin.

A. Korolev (K), massifs of ancient pre-Nectarian South Pole-Aitken basin (black arrows indicate massifs in fig. 8.9), northern interior of upper pre-Nectanan Apollo basin (A; letter athwart 

Nectarian crater Barringer, 69 km, 28° S., 150° W.), and deposits of upper Nectarian Hertzsprung basin (H; centered in direction of white arrows). Massifs near antipode of Seremtatis 

basin (S) are grooved, as near Orientale and Imbrium antipodes (see chap. 10); however, grooves may be secondary craters of Korolev, Hertzsprung, and Apollo. Areas of B through E are 

outlined; black-and-white arrow indicates chain shown in B and is subradial to Orientale basin. Orbiter 1 frame M 28.
. Detail inside Korolev basin, showing probable Onentale-basin secondary chain and fine Orientale-radial lineations. Large crater is Korolev M (58 km, 9° S., 

| 57° \X/ ? Nectarian), filled by Imbrian plains deposits. Orbiter 1 frame H-38.

C. Enlargement of part of A, showing stratigraphic relations of Korolev and Hertzsprung 

basins and adjacent craters. Degradation of pre-Nectarian craters such as Galois (G; 

207 km), Paschen (P; 133 km), and Tsander (Ts; 171 km) indicates burial by both 

Korolev and Hertzsprung deposits. Such Lower Nectarian craters as Mechnikov (M; 

60 km) and Sechenov (S; 62 km) are too sharp to be older than nearby Korolev but are 

clearly overlain by Hertzsprung-basin secondary craters (trending northeast-southwest). 

Such fresh-appearing upper Nectarian craters as Kibal'chich (K; 93 km) and Timiryazev (Ti; 53 km) are superposed on Hertzsprung deposits. Sharp Onentale-basin 

secondaries (lose) cross all other units.

D. Korolev basin and part of periphery. Line separates inner Korolev ejecta (Nki; clearly 

radial southwest of basin) from outer deposits containing many secondary craters (Nko). 

Fresh-appearing Nectarian craters superposed on Korolev include Congreve U (CU; 

59 km, 0.5° N., 171 ° W.) and Lebedinskiy (L; 62 km); sharper Engelhardt (E; 43 km) 

may be Lower Imbrian. Congreve (C; 58 km) and Zhukovskiy (Zh; 81 km, 8° N., 167° 

W.) are degraded Nectarian craters superposed on Korolev deposits but have morphologies and crater densities resembling those of pre-Nectarian craters. Sharply 

textured Doppler (D; 100 km, 13° S., 160° W.) is probably Lower Imbrian. Most 

plains inside Korolev are Nectarian (Np), but some lightly cratered patches are Imbrian 

(Ip), possibly of Orientale origin. Np and Ip may overlie basalt flows. Obiter 1 frame 

M-38.
E. Crookes D (center; 41 km), overlain by deposits of Crookes (49 km, IO°S., I65°W., truncated by left edge). Superposition of Crookes Don Korolev demonstrates 

a Nectanan age; however, subdual because of partial burial, especially of rim sector nearest Crookes, would otherwise suggest a pre-Nectanan age. Freshness of 

Crookes indicates a Copermcan age. Orbiter 1 frame H 30.
ft) O C 3 a. 13 °-« "S 

a 3 3J3 /*"N ^

CUMULATIVE NUMBER OF CRATERS PER 106 km2

o 

|

.

a 2. _ 2 H § 3^ =

.

°

q -0 5 '

e - M . ^

3

s<

00

05
9. NECTARIAN SYSTEM 187

FIGURE 9.24. Upper Nectarian Humorum basin. MH, Mare Humorum.

A. Hilly and furrowed terrain (hf) is probably a composite of Humorum basin ejecta, 

Imbrium-basin secondary pits and lineations, and, possibly, thin Orientale deposits; 

arrow (I) is radial to Imbrium. Nectarian crater Mersenius (M; 84 km, 22° S., 49° W.) 

is superposed on Humorum basin but is marked by Imbrium-radial lineations. 

Humorum-basin massifs are conspicuous south of (below) Mersenius. Orbiter 4 frame 

H-149.

 ':T'*-J <!'.; «'*« *"" ',/ " ~

Terrain south-southeast of Humorum, centered in direction of arrow. Arrow marks 

probable secondary-impact craters of Humorum, including a subradial valley resembling 

Vallis Bouvard (compare fig. 4.4D). Major scarps of Humorum are at 1 and 2. H, 

Hainzel A (53 km, 40° S., 34° W., Eratosthenian); V, Vitello (42 km; see chap. 6). 

Orbiter4frameH-136.
188 THE GEOLOGIC HISTORY OF THE MOON

FIGURE 9.25. Upper Nectarian Hertzsprung basin (H). Inner (Nhi) and outer (Nho) 

deposits are superposed on pre-Nectarian basin Birkhoff (B) and overlain by Orientalebasin secondaries (lose) (compare figs. 4.5, 8.14). Orbiter 5 frame M-28.

-

FIGURE 9.26. Craterlike Nectarian basin Mendeleev, filled with Nectarian plains.

A. Part of radial ejecta and secondary field south of basin (between rim and lower edge of 

picture). Prebasin, mostly pre-Nectarian craters diminish in number and expression near 

rim. Orbiter 1 frame M 136.

B. Interior plains. Crater chain (Catena Mendeleev) is radial to, and undoubtedly secondary to, crater Tsiolkovskiy, 750 km south in direction of chain. Apollo 16 frame 

M-875.
FIGURE 9.27. South-central nearside (centered at 27° S., 5° E.). Apianus (Ap; 63 km) is 

lightly scored by lineations radial to Nectans basin (Rupes Altai is 400 km in direction of 

arrow) and thus is probably pre-Nectarian. Playfair G (PG; 94 km), Krusenstern (K; 47 

km), and Blanchmus (B; 61 km) are heavily mantled by probable Nectans deposits and thus 

are pre-Nectarian; their rims are also nicked by Imbrium-basm secondaries (sc). Intercrater

terrain probably consists mostly of Nectaris-basin deposits. Crater Aliacensis (Al; 80 km) 

would appear moderately fresh if it were not degraded by deposits of Werner (W; 70 km, 

Eratosthenian; compare fig. 3.26); Imbrium-basin secondary craters (sc) superposed on rim 

constrain age of Aliacensis to Nectarian. Playfair (P; 48 km) is also moderately fresh and 

probably Nectarian. Orbiter 4 frame H 101.
190 THE GEOLOGIC HISTORY OF THE MOON

VOLCANIC ROCKS

Volcanism began on the Moon before the emplacement of the 

visible maria. This conclusion was reached even before any samples 

were examined, based on the reasoning that generation of basaltic 

magmas cannot have started coincidentally with the Orientale 

impact. Only post-Orientale maria can have remained exposed 

because all earlier maria were obscured by impacts (Hartmann and 

Wood, 1971; Shoemaker, 1972; Hartmann, 1973; Ryder and Taylor, 

1976; Basaltic Volcanism Study Project, 1981, chap. 8). The high preOrientale impact rate must have kept lunar surfaces so thoroughly 

churned that extensive basalt flows did not long remain exposed, and 

Orientale and each earlier large basin impact must have covered most 

of the Moon with debris. Nor did the impacts trigger the volcanism; 

Orientale would be filled more deeply if it had triggered its own 

volcanic fill.

This well-founded inference is confirmed by physical evidence. 

Mare basalt has been found as fragments in bedrock breccia deposits 

(table 9.5; Ryder and Taylor, 1976; Ryder and Spudis, 1980). The 

Apollo 14 collection contains the most numerous of these samples, 

possibly because it contains material ejected from the already-flooded 

Procellarum and Insularum basins. This collection includes four 

large and several small samples of high-Al basalt, identified by cotectic compositions and the absence of siderophiles as endogenic, and by 

other chemical characteristics as related to mare basalt (Ridley, 1975; 

Taylor, 1975, p. 136; 1982, p. 300). Dated samples of high-Al basalt 

include the 251-g sample 14053 and 45-g sample 14072, found as 

individual rocks, and clasts from breccia sample 14321 (table 9.5). 

Additional high-Al basalt was found by Apollo 14 as small clasts in 

breccia samples (for example, fragments in sample 14063 with intermediate Ti contents; Ridley, 1975) and as fragments in the regolith 

(Ryder and Spudis, 1980). Apollo 17 fragment-laden impact-melt 

rocks also contain scattered small fragments of high-Al mare basalt 

(James and McGee, 1980). One sampled boulder contains clasts of 

basalt that may be intermediate between KREEP and mare basalt 

(sample 72275; Ryder and Spudis, 1980). The fact that relatively few 

samples of volcanic basalt have been discovered in terra breccia 

deposits is not surprising in view of the small volume of the crust 

composed of mare basalt today (see chap. 5). A major impact on any 

part of today's Moon would generate breccia much richer in terra 

material than in mare basalt.

The pre-Imbrian basalt fragments have been dated at from 3.85 

to 3.99 aeons (combined range of ages obtained by Ar-Ar and Rb-Sr 

methods, table 9.5; Ryder and Spudis, 1980). If the Nectaris basin 

formed 3.92 aeons ago, these ages span the boundary between the 

Nectarian and pre-Nectarian; most are Nectarian.

TABLE 9.5. Absolute ages of some pre-Imbrian volcanic clasts

[After Ryder and Spudis (1980).

Stations C2, C 1 , and Cl are on rim flank of Cone Crater at the Apollo 14 landing 

site; 2(1), station 2, boulder 1, Apollo 17 landing site.

Ages are calculated using International Union of Geological Sciences (IDGS) radioactive decay constants (Steiger and Jager, 1977), except the age determined 

by Husain and others (see headnote, table 9.1). Ar-Ar ages are ^Ar- Ar plateau ages.

References: C72, Compston and others (1972); C75, Compston and others (1975); 

H72, Husain and others (1972); M73, Mark and others (1973); M75, Mark and 

others (1975); PW71, Papanastassiou and Wasserburg (1971b); S73, Stettler and 

others (1973); T71, Turner and others (1971; see also Turner, 1977); Y72, York 

and others (1972)]

Sample Station Lithology Age Method Reference

14053

14072

Clasts in

14321

Clasts in

72275

C2 251-g high-Al mare-type

basalt.

C 1 45-g high-Al mare-type

basalt.

Cl High-Al mare-type basalt

clasts (sample 14321

is diverse 9-kg breccia) .

2(1) Basalt intermediate between mare and KREEP

types.

3.88±0.04

3.88±0.04

3.89±0.05

3.92±0.08

3.91+0.09

3.91±0.14

3.93±0.12

3.98+0.05

3.85±0.05

3.87+0.04

3.88±0.05

3.93+0.05

3.93+0.12

3.96+0.06

3.96+0.08

3.87+0.04

3.93+0.04

Ar-Ar

Rb-Sr

Ar-Ar

Ar-Ar

Rb-Sr

Rb-Sr

Rb-Sr

Ar-Ar

Ar-Ar

Rb-Sr

Ar-Ar

Ar-Ar

Rb-Sr

Rb-Sr

Rb-Sr

Rb-Sr

Rb-Sr

S73

PW71

T71

H72

C72

M75

M73

Y72

Y72

PW71

Y72

Y72

M73

M73

C72

PW71

C75

The extent of Nectarian and pre-Nectarian volcanism is unclear. 

If the heat source was entirely radiogenic (Taylor, 1982), the ancient 

basalt units were less extensive than the visible maria, because 

radioactive U, Th, and K are not abundant in the Moon. If, however, 

the Moon retained considerable primordial heat during its first aeon 

(Hubbard and Gast, 1971; Wood, 1972b; Toksoz and Johnston, 1974; 

Hubbard and Minear, 1975; Ryder and Taylor, 1976; Solomon and 

Head, 1979,1980), only a small rise in temperature might have been 

sufficient to melt considerable magma.

Photogeologic evidence is consistent with considerable premare 

volcanism. Scattered patches of Nectarian plains have been identified 

in many parts of the terrae on the basis of crater densities (pi. 7; figs. 

9.265, 9.28, 9.29; Neukum and others, 1975b; Wilhelms and El-Baz, 

1977; Stuart-Alexander, 1978; Lucchitta, 1978; Wilhelms and others, 

1979). Some Nectarian plains, including the concentration in the 

Mutus-Vlacq basin already discussed (figs. 8.12, 9.2B), surround Nectarian basins and are stratigraphically related to the lineate deposits 

in the same way the circum-Orientale plains are related to the 

Hevelius Formation. That is, they are younger than the ejecta and 

secondary craters, but by an unknown amount. If they were part of 

the ejecta-deposition process, they are momentarily younger; if they 

are volcanic, they may be substantially younger.

Many of the plains may, in fact, have originated by a combination 

of volcanism and ejecta deposition. Eggleton and Marshall (1962) 

ascribed the planar light-colored surfaces south of the type area of the 

Fra Mauro Formation (figs. 10.18,10.19) and elsewhere in the central 

lunar highlands to covering of premare surfaces by Imbrium ejecta. 

Wilhelms and El-Baz (1977) noted that the Nectarian and Imbrian 

plains in the Lomonosov-Fleming basin (fig. 9.28) are too smooth and 

extensive to be easily explained as basin ejecta and suggested that 

they were volcanically emplaced. Schultz and Spudis (1979) showed 

that the Lomonosov-Fleming plains are spotted by many dark-haloed 

craters that probably excavated dark material from beneath a lightercolored cover. Therefore, the plains may consist of basaltic materials 

overlain by a thin cover of Crisium, Imbrium, or Orientale ejecta. 

Orientale ejecta may have covered basaltic plains in the SchillerZucchius basin near Orientale (Hartmann, 1966; Schultz and Spudis, 

1979), which have diffuse patches darker than most terra plains (fig. 

7.6; Offield, 1971). Numerous other plains, mostly inside basins of preNectarian (Mutus-Vlacq) and Nectarian (Hertzsprung, Korolev) age, 

may also consist partly of premare basalt (pi. 7). Original emplacement in the style of the currently visible mare-basalt flows is indicated.

Geochemical sensing suggests that the dark halos contain both 

mare and KREEP basalt (Schultz and Spudis, 1979; Hawke and 

Spudis, 1980). Dark-haloed craters superposed on light-colored plains 

of Nectarian and Imbrian age (Wilhelms and El-Baz, 1977) within the 

possible Balmer-Kapteyn basin (table 4.2) display the telltale high 

radioactivity of KREEP (Haines and others, 1978), an observation 

suggesting impact exhumation of Nectarian or pre-Nectarian 

KREEP basalt (Hawke and Spudis, 1980). Mare basalt was probably 

exhumed from beneath the plains north of Taruntius and in many 

other localities (Schultz and Spudis, 1979; Hawke and Spudis, 1980). 

Therefore, basalt with a range of compositions may have been 

extruded far back into lunar history. The extent of this volcanism, 

however, remains uncertain.

CHRONOLOGY

If the age of 3.92 aeons favored here for the Nectaris basin is 

correct, then the Nectarian Period began 3.92 aeons ago and ended 

3.85 aeons ago when the Imbrium basin was formed (see chap. 10).

One of the most frequently cited hypotheses in the lunar literature is that a "cataclysm" formed most or all basins within an 

interval of only 0.1 aeon (Tera and others, 1974). This cataclysm is 

generally conceived as a resumption of heavy bombardment after a 

lull following the heavy accretionary bombardment. The main basis 

for this hypothesis is that radiometric ages of terra rocks are concentrated around 3.8 to 3.9 aeons. Serenitatis looks like a relatively old 

basin; yet it has an absolute age that falls within this interval. The 

following data, presented here and in chapter 8, are consistent with 

the contrary view that the cratering rate declined relatively steadily 

(Baldwin, 1974b; Hartmann, 1975). (1) Numerous ages do predate 4.0 

aeons (table 8.4; Baldwin, 1974b). (2) A concentration of late dates 

would be the natural consequence of a steadily declining cratering 

rate because of repeated impact reworking (Hartmann, 1973,1975).
9. NECTARIAN SYSTEM 191

(3) No pre-Nectarian basins were sampled (fig. 9.30). (4) A smooth 

cratering curve (fig. 8.16) is fixed by a 3.92-aeon age for Nectaris, a 

3.85-aeon age for Imbrium, and the ages of mare-basalt samples; this 

curve is consistent with a relatively young Nectarian relative age 

(table 9.3) and a 3.87-aeon absolute age for the Serenitatis basin.

"Minicataclysms" forming some of the basin age groups are possible, 

but no cataclysm in the sense generally meant is recorded in lunar 

rocks. Instead, the record of early lunar impacts reveals a continuous 

barrage of large projectiles that only during the Nectarian Period 

began to approach less than cataclysmic proportions.

'" *./:- . >^ ^^^^^^^fk^^W

FIGURE 9.28. Lomonosov-Fleming basin (compare fig. 8.13). F, Nectarian crater Fleming (130 km, 15° N., 110° E.). Dark halo (arrowhead) indicates basalt buried beneath light-colored 

Nectarian and Imbrian plains of probable basin origin. Mare fills crater Lomonosov (L; 93 km, 27° N., 98° E.), and Lomonosov ejecta is apparently covered by older mare in crater Joliot (under 

letter J); thus, Lomonosov is Lower Imbrian (compare fig. 4.7). Mare at left edge is Mare Marginis. GB, Giordano Bruno (22 km, 36° N., 103° E.), one of the Moon's youngest craters of its 

size. Apollo 16 frame M-3004.

FIGURE 9.29. Crater size-frequency curves by Neukum (1977), showing ages of plains in 

craters Gagarin, Pasteur, and Hilbert relative to Imbrium-basin deposits ("event"). Older, 

probably Nectarian frequencies are a factor of at least 7 higher than the Imbrium-basin 

frequency. Containing craters are pre-Nectarian (Gagarin and Pasteur) and Nectarian 

(Hilbert) (compare figs. 3.5, 8.1 I).

 Gagarin Floor 

D Pasteur Floor 

A Hilbert Floor

10" 10' 10 

CRATER DIAMETER D(km)
192 THE GEOLOGIC HISTORY OF THE MOON

FIGURE 9.30. Geologic cross sections based on figure 8.4A, showing deposits of the Nectarian basins Nectaris (n), Humorum (h), Crisium (c), and Serenitatis (s). Triangles, subbasin breccia 

(overlain by impact melt, not shown). Materials of the later Imbrium basin are derived partly from Serenitatis deposits and partly from pre-Nectarian terrane. Depression created by Procellarum 

basin (see fig. SAB) is still evident, though reduced by superposition of other pre-Nectarian deposits. Dotted lines, present surface.
10. LOWER IMBRIAN SERIES

4
FIGURE 10.1 (overleaf). Units of the Imbrian System where first recognized (compare figs. 1.6, 2.1). View centers on massive, rugged Monies Apenninus, part of the Imbrium-basin rim. 

Imbrium-basin ejecta grades from coarsely hummocky forms on Apennine flank to smoother drumlinlike or whaleback forms near crater Ukert (near bottom; 23 km, 8° N., 1.5° E., probably 

Lower Imbnan). Archimedes (at top; 83 km, 30° N., 4° W.), which is superposed on rugged Monies Archimedes and plains of the Apennine Bench (Hackman, 1966X is flooded by mare 

basalt. Massive slumping from Monies Apenninus is indicated by parallel slructure and oullines of ridges norlhwesl of Apennine fronl. Type area of lop of Imbrian System is in bay of Mare 

Imbrium between the Apennine Bench and crater Eraloslhenes (iruncaled by lefl edge of pholograph). Orbiler 4 frame M-102.
10. LOWER IMBRIAN SERIES

CONTENTS

Page

Introduction ...................................................... 195

Definition ........................................................ 196

Interior and near-rim Imbrium-basin deposits ....................... 196

Introduction ................................................... 196

Apennine Bench Formation Apollo 15 samples .................. 197

Montes Apenninus massifs Apollo 15 samples .................. 198

Material of Montes Apenninus .................................. 202

Alpes Formation ............................................... 203

Fissured plains ................................................ 203

Fra Mauro Formation ............................................. 204

Regional relations ............................................. 204

Setting of the Apollo 14 landing site ............................. 204

Sample provenance ............................................ 204

Petrology ..................................................... 205

Emplacement process .......................................... 211

Absolute age .................................................. 212

Conclusions ................................................... 212

Imbrium-basin secondary craters .................................. 213

Light-plains materials ............................................. 215

Cayley and Descartes Formations .................................. 216

Introduction ................................................... 216

Imbrium-basin origin ........................................... 217

"Local" origin ................................................. 219

Conclusions ................................................... 220

Other basin and crater materials ................................... 221

Recognition criteria ............................................ 221

Frequency .................................................... 222

Other terra materials ............................................. 222

Mare materials ................................................... 223

Chronology ...................................................... 224

INTRODUCTION

The extensively exposed and relatively well sampled Lower 

Imbrian Series is the key to the ages and origins of most materials of 

the lunar terrae. Deposits of the Lower Imbrian Orientale and 

Imbrium basins constitute the Moon's most extensive laterally continuous stratigraphic horizons (pi. 8). The Imbrium basin dominates 

the surface visible from Earth and was the first basin to be treated 

stratigraphically (see chaps. 2, 4; fig. 10.1). Deposits of the Orientale 

basin dominate the terrae on both sides of long 90 W. and serve as 

models for other basin materials, as previous discussions have made 

clear. Thus, the Lower Imbrian Series dominates a large part of the 

lunar surface, despite the small number of events that emplaced its 

rock units (pi. 8). Both groups of basin deposits make excellent stratigraphic markers for dividing the lunar stratigraphic column (table 

7.1) because of their great extent and clear expression.

Because of their early discovery and prominence, the Imbriumbasin deposits attracted more sampling missions than did any other 

class of terra units. The curtailed Apollo 13 and the successful Apollo 

14 nights were sent to a point on the Fra Mauro Formation, the most 

distinctive Imbrium-basin unit. Exploration of the Imbrium rim at 

Montes Apenninus was the main objective of the Apollo 15 mission. 

The Cayley and Descartes Formations, though considered to be volcanic when selected as targets for the Apollo 16 landing (see chap. 2),

are Lower Imbrian photogeologic units emplaced by impact and may 

also contain Imbrium-basin material.

In terms of total volume of returned material, however, rocks 

emplaced during the Early Imbrian Epoch are less well represented 

than Nectarian and pre-Nectarian materials. Apollos 14 and 15 

together returned much less terra material than did Apollo 17. 

Arguably, the abundant Apollo 16 collection may include more material processed in the Nectarian than in the Imbrian Period (see chap. 

9). Effects of igneous crustal processes active in pre-Nectarian time 

have profoundly affected all sample collections (see chap. 8).

The problem of determining which compositions and textures of 

the Apollo 14, 15, and 16 samples were created during the Early 

Imbrian events that emplaced them and which were created earlier 

occupies considerable space in this chapter. The first half of the 

chapter is devoted to the Imbrium basin and to the samples returned 

from its deposits by the Apollo 14 and 15 missions. I develop the 

viewpoint that the Apollo 14 materials owe their properties to their 

emplacement in the Fra Mauro Formation, as originally proposed 

but not currently so widely believed. A similar analysis for the more 

controversial Cayley and Descartes Formations sampled by Apollo 16 

is less conclusive. Resolution of the remaining questions about the 

Apollo 14 through 16 rocks would significantly aid in dating large 

areas on the absolute and relative time scales and in correlating 

compositions with sources in the crust. Nonsampled Lower Imbrian 

units, including the Orientale basin that already has been thoroughly described, are given less attention here.

195
196 THE GEOLOGIC HISTORY OF THE MOON

DEFINITION

Deposits of the Imbrium basin and the radial system of grooves 

called "Imbrian(um) sculpture" have long served as the main reference for dividing materials into Imbrian and pre-Imbrian on the 

nearside (see chap. 7; Gilbert, 1893; Shoemaker and Hackman, 1962). 

In August 1967, photography of the Orientale basin at the end of the 

Lunar Orbiter 4 mission added a major stratigraphic horizon to the 

Imbrian System as it had previously been defined. Superpositional 

relations (fig. 10.2), morphologic freshness (figs. 4.4,10.3), and sizefrequency distributions (fig. 10.4) all indicate that Orientale is younger than Imbrium. Orientale is older than the nearby mare surface, 

which would have been swamped by the basin ejecta if it predated 

Orientale (figs. 10.3,10.5). The Orientale deposits constitute excellent 

stratigraphic markers because of their vigorous topographic expression and clear relations with other units over a wide area (pis. 3, 8). 

They probably would have been used as a stratigraphic datum to 

define a major time-stratigraphic boundary, had they been in a better 

position for study during the telescopic phase of the geologic-mapping

FIGURE 10.2. Superpositions of secondary craters of the Orientale basin on those of the 

Imbrium basin.

A. Regional view of northwest nearside limb, including north pole (black-and-white arrow). 

Orientale basin is at bottom. Black arrow parallels Imbrium sculpture (northeastsouthwest) and marks superposition indicated in B. White L, pre-Nectarian Lorentz 

basin (360 km, 34° N., 97° W.); P, Eratosthenian crater Pythagoras (130 km, 64° N., 

63° W.); black L, Copernican crater Lichtenberg (20 km, 32° N., 68° W.), embayed 

by Copernican mare basalt (see chap. 13). Numerous craters with uplifted floors are near 

terra-mare boundary (see chap. 6). Area of figure 10.3 is outlined. Orbiter 4 frame 

M-189.

program. This key horizon has, in fact, been used on some geologic 

maps to divide informal lower and upper series of the Imbrian System 

(for example, McCauley, 1973).

The two series of the Imbrian System are here defined formally. 

The Lower Imbrian Series includes all materials from the bottom of 

the Fra Mauro Formation to the top of the Hevelius Formation (fig. 

7.1). The type area of the base of the Lower Imbrian Series is the type 

area of the Fra Mauro Formation, lat 0°-2° S., long 16°-17.5° W. (figs. 

10.1,10.19; Wilhelms, 1970b). The type area of the top of the series is 

the type area of the Hevelius Formation, near lat 2° N., long 68° W. (fig. 

4.4H; McCauley, 1967a). A reference area for the top of the series is 

bounded by lat 2° S.-2° N., long 75°-80° W. (fig. 10.3). The Upper 

Imbrian Series includes all materials above the top of the Hevelius 

Formation that is, younger than the Orientale basin to the top of 

the Imbrian System as defined in chapter 11. The Early Imbrian 

Epoch thus began at the moment of the Imbrium impact and ended 

when the Hevelius Formation came to rest. The Late Imbrian Epoch 

began when the Hevelius came to rest and ended when the mare 

materials that define the top of the Imbrian System solidified.

INTERIOR AND NEAR-RIM 

IMBRIUM-BASIN DEPOSITS

Introduction

Imbrium is the Moon's largest well-preserved basin (table 4.1) but 

is not entirely understood. Its basaltic fill, Mare Imbrium, obscures 

an area larger than the entire Orientale topographic basin (table 6.1). 

Interior deposits and large parts of interior rings are visible only at 

the Apennine Bench (figs. 1.6, 2.1, 4.3AA, 10.6). Concentric mare 

ridges help somewhat in locating buried rings (Hartmann and 

Kuiper, 1962; Brennan, 1976). The southwest, west, and northwest 

sectors of the topographic rim are also buried by basalt flows,

B. Detailed view of Onentale-secondary craters, showing hernngbone-patterned ejecta 

interrupting Imbrium-radial chains (arrow at bottom). Orbiter 4 frame H 189.
10. LOWER IMBRIAN SERIES 197

*. V 'tf\ATic * '«.* /.«*«

»ri3*&:ft

*> r^^%*

' .-A. r».-J»* *1& vS tv^H*-1 '

Ht'Wffi^&Ji'to * I f* K.^^ftV ^^i"* .' V^.MC ^ W*L; ' >:

«* : W-i '^^7^' 1 * ''= &' r* ' ^1 .1 *y>,'/^^F'' f!^*-^^C^

5^?..^ -»4f vj j3*i

FIGURE 10.3. Orientale-secondary craters and plains deposits, partly flooded by mare

basalt (arrows); old mare unit is at lower arrow (see fig. 10.2 A for location). Reference area

' for typical exposure of inner facies of the Hevelius Formation (2° N.-2° S., 75-80° W.) is

enclosed by corner markers. Secondary craters of Krafft (K; 51 km, 17° N., 73° W.) and

Cardanus (C; 50 km) are superposed on Orientale deposits and old mare, but are flooded on

pyroclastic materials, and deposits of such large craters as Iridum 

(figs. 3.8, 5.12). The connections between rings in the north and 

northeast are partly buried (pi. 8). Therefore, the identity not only of 

the boundary of excavation but even of the topographic rim is uncertain in the half of the basin periphery from the southwest to the 

northeast sectors. In the east and south, Montes Apenninus and 

Carpatus (figs. 10.6,10.7) constitute the topographic rim and, in my 

opinion, the excavation boundary (see chap. 4). These ranges are 

equivalent to Montes Cordillera in position and morphology 

(Wilhelms and McCauley, 1971; M'Gonigle and Schleicher, 1972; Brennan, 1976) and in origin (Baldwin, 1974a; Wilhelms and others, 1977; 

Hodges and Wilhelms, 1978).

The best exposures of Imbrium-basin materials are on Montes 

Apenninus, in large tracts southeast of that range, and north of Mare 

Frigoris. This section describes the Imbrium materials from the 

inside out and interprets them by comparison with their Orientale 

equivalents.

Apennine Bench Formation Apollo 15 

samples

Hackman (1964,1966) named the level or wavy materials between 

Montes Archimedes and Montes Apenninus the 'Apennine Bench 

Formation" (figs. 2.1,6.10,10.1,10.7). Page (1970) extended this name 

to other, similar plains inside the Imbrium-basin rim that underlie 

deposits of Imbrian craters and the mare (figs. 2.4, 10.8). Overlap of

10

10'

10-

10-

io- 5 _l___I I

0.1 1.0 10 100 

CRATER DIAMETER, IN KILOMETERS

FIGURE 10.4. Cumulative size-frequency distributions of craters superposed on major 

Imbrian units. Counts of craters larger than 20 km in diameter on deposits of the Imbrium and 

Orientale basins are from this study; Nectarian and pre-Nectarian fields are added for 

comparison (figs. 8.6, 9.22). Other curves are plotted from data of Neukum and others 

(1975a); data points for the Apennines are from two counts by Neukum and others (1975a), 

which have a common point at 3 km. Curve for Imbrian craters (Ic) is from Wilhelms and 

others (1978).

west by younger mare; similar ages and sizes of these two craters suggest twin impacts (see 

chap. 3). Ejecta of Copernican crater Olbers A (OA; 43 km, 8° N., 78° W.) is superposed 

on Orientale deposits and partly covered by small patch of very young mare. Olbers (O; 75 

km, probably Nectarian) and Riccioli (R; 146 km, probably pre-Nectarian) are buried by 

Orientale deposits. Orbiter 4 frame H-l 74.
198 THE GEOLOGIC HISTORY OF THE MOON

the Apennine Bench Formation on the rugged Imbrium-basin 

topography shows that the plains are at least slightly younger than 

the basin (figs. 10.7,10.8).

As for other terra-plains deposits, interpretations of origin have 

shifted back and forth between volcanic and impact. E.M. Shoemaker 

and R.J. Hackman (oral commun., 1963) proposed that the Apennine 

Bench Formation consists of Imbrium-basin ejecta or impact melt. 

The impact-melt interpretation is supported by position inside the 

basin and by morphologies indicating subsidence of a fluidlike material (fig. 10.7; Wilhelms, 1980). The Apennine Bench Formation and 

the Maunder Formation of Orientale are similar in these respects 

(table 4.3). Volcanic interpretations, which are also consistent with 

the fluidlike appearance, were favored in the era when lunar lightcolored plains in general were commonly thought to be volcanic 

(Hackman, 1966; Wilhelms, 1970b; Carr and others, 1971; Wilhelms 

and McCauley, 1971).

A volcanic origin is supported by discovery among the Apollo 15 

samples of fragments of possible terra-type volcanic rock. The Apennine Bench Formation probably underlies the mare basalt in the 

region of the landing site (figs. 10.7, 10.8; Carr and others, 1971), 

although the formation is not exposed within range of the astronaut 

traverses (fig. 10.9). The distribution of KREEP-rich fragments in the 

regolith (Basu and McKay, 1979) is consistent with derivation from 

beneath the mare (Spudis, 1978b). Their Th-rich medium-K KREEP 

chemistry matches the composition detected by the orbiting gammaray spectrometer over Apennine Bench exposures west of the landing 

site (Spudis, 1978a; Hawke and Head, 1978b; Metzger and others, 

1979; Clark and Hawke, 1981). Because the fragments have igneous 

textures and no meteoritic siderophile elements, many investigators 

have concluded that they are true endogenic igneous rocks (fig. 10.10; 

Irving, 1975,1977; Dowty and others, 1976; Ryder, 1976; Meyer, 1977; 

Basaltic Volcanism Study Project, 1981, p. 274-278). A further argument for a volcanic origin is that the samples were probably derived 

from a source more radioactive than any now present beneath the

FIGURE 10.5. Partly buried crater group (above arrow) in Oceanus Procellarum (2° S., 

40.5° W.), identified as secondary to Orientale basin by ejecta orientation subradial to 

Orientale (centered 1,700 km to west in direction of arrow) and subequal size (about 5 km; 

compare crater F, Flamsteed FB, 5 km). Orbiter 4 frame H-143.

landing site (as shown by the Apollo 15 heat-flow experiment; Langseth and others, 1976); an Imbrium melt sheet would incorporate 

material similar to the present subbasin material (Spudis, 1978a). An 

episode of KREEP volcanism in the Early Imbrian Epoch is indicated, 

possibly triggered by the Imbrium impact (Spudis, 1978a).

The volcanic interpretation is still not universally accepted. Taylor (1982, p. 214-216) pointed out that the absence of siderophiles and 

fragments does not prove an igneous origin and suggested that 

Montes Apenninus are as rich in KREEP as is the Apennine Bench. 

Some interpretations of the orbital data, however, favor a higher Th 

content for the bench (Metzger and others, 1979).

Three geochronologic methods have yielded ages for the mediumK KREEP-rich samples 15382 and 15386 averaging 3.85 aeons (table 

10.1). This average age probably dates the emplacement of the 

KREEP-rich basalt and constrains the age of the Imbrium basin to 

3.85 aeons or older, regardless of whether the KREEP-rich liquids 

were melted by impact or endogenic heat.

Montes Apenninus massifs Apollo 15 

samples

Steep rugged massifs that are among the Moon's highest mountains form the crest of Montes Apenninus; they probably were 

uplifted structurally and overlain by ejecta (figs. 10.1,10.7,10.8,10.11; 

Carr and others, 1971; Wilhelms, 1980). Large slump masses lie as far 

as 60 km northwest of embayments in the mountains, which were 

evidently the sources of the slumps (figs. 6.10,10.7,10.8; Wilhelms and 

others, 1977; Wilhelms, 1980). Parts of the Apennine Bench Formation 

overlie some of these slumps (Spudis, 1978a; Wilhelms, 1980). Most 

investigators currently believe that Montes Apenninus are an exterior scarp of an Imbrium basin smaller than its topographic rim. 

Others, including me (see section in chap. 4 entitled "Origin of 

Rings"), regard an origin as the excavation boundary as much more 

likely because the Apennines are a massive elevated range of mountains, the largest on the lunar nearside, and because they mark a 

sharp discontinuity in the textures of basin deposits (figs. 10.1,10.6, 

10.7).

The identity of the boundary of excavation in other sectors is less 

clear. The Apennines and Montes Carpatus, another obvious part of 

the topographic-basin rim, define a circle 1,200 km in diameter with 

Montes Alpes and an oval 1,500 km long with the north shore of Mare 

Frigoris (pi. 3; figs. 4.3AA, 10.6). The Apenninus-Carpatus ring apparently divides in Montes Caucasus (figs. 10.6, 10.12). The southern 

Caucasus continue as Montes Alpes, and the northern Caucasus as 

the Frigoris shore (Wilhelms and McCauley, 1971). In the interpretation given in chapter 4, the large oval bounds the Imbrium cavity. The 

hypothesis that the outer parts of large basins are differentially 

excavated may explain the divergence; the Frigoris sector was excavated to a larger diameter than was the Carpatus sector, which was 

blocked by the Serenitatis rim. Lineations in the northern Caucasus 

(figs. 10.6,10.12)point to an impact point at the center of the CarpatusAlpes circle at lat 33° N., long 18° W. (table 4.1).

The origin of Montes Archimedes, the rugged elevation south of 

the crater Archimedes, is ambiguous (fig. 10.1). It lies at the same 

distance from the center as does Montes Alpes (pi. 3; Wilhelms and 

McCauley, 1971) but also resembles the slumps derived from the 

Apennine scarp. Thus, it may be either part of an Imbrium-basin ring 

or a slump.

One of the major objectives of the Apollo 15 mission in July 1971 

was to collect samples of Montes Apenninus. According to impact 

models based on simple craters (see chap. 3), this basin-rim material 

might have come from deep within the Moon. Primitive materials or, 

at least, pre-Imbrian materials might have been uplifted and exposed

FIGURE 10.6. Imbrium basin, bounded by Montes Alpes (Al), Caucasus (white C), and 

Apenninus (Ap). Imbrium-radial crater chains (black-and-white arrows) north of Mare 

Frigoris (F) are superposed on deposits of Nectarian Humboldtianum basin (H; 61° N., 

84° E.; Lucchitta, 1978) and on many craters, including pre-Nectarian W. Bond 

(W; 158km) and Barrow (B; 93 km; compare fig. 10.15). Humboldtianum deposits cover probable Nectarian craters Nansen (N; 122 km, 81° N., 95° E.) 

and Schwarzschild (S; 235 km, 71° N., 120° E.), and are overlain by Nectarian crater 

Endymion (E; 125 km, 54° N., 57° E.), Lower Imbrian crater Compton (black C; 162 km, 

56° N., 105° E.; compare fig. 9.4), and Copernican crater Hayn (87 km, north of H). 

Fewer small prebasin craters are visible near both basins than along upper left edge of 

photograph. White arrow, north pole. Orbiter 4 frame M-l 16.
« :.*,

 CT 4 '»** :"-* *l 3*^L V5 (V ^ « 
200 THE GEOLOGIC HISTORY OF THE MOON

TABLE 10.1. Absolute ages oflmbrium-basin materials

Stations: Apollo 14 (see fig. 10.21): Traverse EVA 1 was within 150 m of the 

Lunar Module (LM); stations Cl, C 1 , and C2 are on the rim flank of Cone Crater 

along traverse EVA 2; F, 400 m east of LM on EVA 2; G, 250 m east of LM on EVA 

2. Apollo 15 (see fig. 10.9): Station 7 is at Spur Crater along the front of 

Montes Apenninus.

Rock types: CMR, crystalline melt rock of uncertain origin (volcanic basalt or 

clast-poor or -rich impact-melt rock); CPM, clast-poor impact-melt rock; FMB, 

Fra Mauro breccia (clast-rich impact-melt rock); xln, crystalline. Lithologic 

descriptions in much of the geochronologic literature are inadequate for categorizing samples further; for example, the term "basalt" is commonly used for 

impact-melt rocks (see chap. 8).

Ages calculated using International Union of Geological Sciences (IUGS) radioactive-decay constants (Steiger and Jager, 1977), except those determined by 

Husain and others (1972) (see headnote, table 9.1). Most Ar-Ar ages are 40Ar- 

 "Ar plateau ages. Age determinations with error ranges of ±0.10 aeon or 

greater and those obviously outside the likely range (for example, those of 

Upper Imbrian mare-basalt grains) are omitted.

References: AD74, Alexander and Davis (1974); AK74, Alexander and Kahl (1974); 

C72, Compston and others (1972); CL79, Carlson and Lugmair (1979); H72, Husain 

and others (1972); K72, Kirsten and others (1972); M72, Murthy and others 

(1972); M74, Mark and others (1974); M79, McKay and others (1979); N74, Nyquist and others (1974); N75, Nyquist and others (1975); PW71, Papanastassiou 

and Wasserburg (1971b); PW76, Papanastassiou and Wasserburg (1976); S73, 

Stettler and others (1973); T71, Turner and others (1971); T72, Turner and 

others (1972); T73, Turner and others (1973; summarized in Turner, 1977); 

WP71, Wasserburg and Papanastassiou (1971); Y72, York and others (1972)]

Sample Station Rock type Age Method Reference

Apollo 14 young group

14001,7,3

14066 

,21,1.02 

,21,1.01 

.04,

EVA 1 2-4 mm CMR-

,21,2. 

,21,2,

14073

01

14150 

,7,3 

,7,2

14161 

,34,2 

,34,6 

,34,5 

,34,4

14167 

,6,1 

,6,7 

,6,3 

,8,1 

,9

14257,12,1

G 

G

EVA 1 

EVA 1

3.81+0.03 

3.84+0.04

510-g FMB

Whole rock 3.84+0.03 

Xln matrix 3.86+0.03 

White xln clast 1 3.88+0.03 

Breccia clast 1 3.90+0.02

10-g CPM (plagioclase 3.80+0.04 

separates). 3.82+0.05

4-10-mm fines

CPM 3.82+0.02 

CPM 3.83+0.01

2-4-mm fines

CMR 3.80+0.08 

CPM 3.84±0.04 

CMR1 3.87+0.08 

CMR 1 3.90+0.06

2-4-mm fines

CMR 3.79+0.03 

CMR 3.82+0.03 

CMR 3.83+0.05 

CMR 1 3.88+0.04 

CMR 1 3.91-3.95

EVA 1 2-4-mm fine 3.87+0.06

Rb-Sr 

Ar-Ar

Ar-Ar 

Ar-Ar 

Ar-Ar 

Ar-Ar

Rb-Sr 

Ar-Ar

Rb-Sr 

Rb-Sr

Ar-Ar 

Ar-Ar 

Ar-Ar 

Ar-Ar

Ar-Ar 

Ar-Ar 

Ar-Ar 

Ar-Ar 

Ar-Ar

Ar-Ar

PW71 

T71

AD74 

AD74 

AD74 

AD74

PW71 

T72

M79 

M79

K72 

K72 

K72 

K72

H72 

H72 

H72 

T71 

Y72

H72

Sample Station Rock type Age Method Reference

14270,1-7

14276

14303, 13, R5.1

14310

EVA 1 26-g FMB 3.83+0.05

EVA 1 13-g CPM 3.80±0.04

EVA 1 Clast from 898-g FMB 3.85+0.04

G 3.4-kg CPM 3.78+0.03

3.79+0.04

3.81±0.05

3.82+0.06

3.83+0.04

3.85+0.04

3.85+0.05

3.85+0.06

3.86+0.03

rAI Ar rAI Ar

Rb-Sr

Ar-Ar

Ar-Ar

Rb-Sr

Ar-Ar

Ar-Ar

Ar-Ar

Rb-Sr

Ar-Ar

Rb-Sr

Rb-Sr

rAIN Ak74/ *r

WP71 nr / J.

K72

UI79 M/ c.

PW71

T72

S73

T71

C72

Y72

M72

M74

Apollo 14 old group (pre-Imbrian)

14053 

14072 

14321

C2 C 1 

Cl

High-Al basalt - 

do -

3 QQ Q Q9

3

Q r- n r\ fSee table 9.5. 

nn _ __ _ ___ __

Do 

Apollo 15 group

15382

1 £ OQ£

15455

7 

7

7

7

KREEP-basalt fragment 

do 

do.? 

Impact melt of breccia

3.82+0.02 

3.84+0.05 

3.85+0.05

3.85+0.08

3.86±0.04 

3.83+0.04

3.86+0.04

Rb-Sr 

Ar-Ar 

Ar-Ar

Sm-Nd

Rb-Sr 

Rb-Sr

Ar-Ar

PW76 

S73 

T73

n 70

N75 

N74

AK74

matrix.

1 Probably older than the Fra Mauro Formation.

Massif

WEST Mare Slump

Alpes Formation and material 

of Montes Apenninus

Apennine Bench Formation

EAST

FIGURE 10.8. Diagrammatic geologic cross section of region of Apollo 15 landing site, 

based on figure 10.7. Indicated history: (1) uplift of massifs and ejection of the Alpes 

Formation and "material of Montes Apenninus," (2) slumping of large blocks from Apennine 

front, (3) emplacement of the Apennine Bench Formation (either impact melt or volcanic 

basalt), and (4) emplacement of mare basalt.

FIGURE 10.7. Region of Apollo 15 landing site (arrow). Mare basalt of Palus Putredinis (PP) is overlain by secondary craters and rays of Autolycus, centered north of area. Sinuous rille is Rima 

Hadley. lab, Apennine Bench Formation; s, slumps from Apennine front; C, Conon (22 km; compare fig. 10.15). Stereoscopic pair of Apollo 1 5 frames M-41 3 (right) and M-415 (left).
10. LOWER IMBRIAN SERIES 201

in the Imbrium-basin massifs. Materials of the Serenitatis basin are 

likely to be included because that basin's west rim is cut off by the 

Apennines and thus lay in the target area of the Imbrium impact (pi. 

3; fig. 9.30; Wilhelms and McCauley, 1971; Spudis and Head, 1977).

Fewer samples than expected were obtained from the massifs. 

Most that were collected are small and of uncertain geologic context 

because outcrops are covered by thick colluvium (fig. 10.11). The 

largest samples are two "black and white rocks" (15445,15455), totaling 1.17 kg in mass and representing a 1-m boulder sampled at station 

7 (figs. 10.9,10.13; Swann and others, 1972). Dark, aphanitic impactmelt rock of low-K KREEP composition encloses light-colored clasts of 

mafic ANT rock (Reid and others, 1977; Ryder and Bower, 1977). The 

ANT rock was crushed or fragmented in one or several episodes 

(James, 1977). Some of the clasts are composed of ancient, apparently

"pristine" noritic and troctolitic rock (table 8.4; Nyquist and others, 

1979b). The impact-melt rock is similar in composition to impact-melt 

rock from the Taurus-Littrow massifs of the Serenitatis basin, though 

richer in Mg (Taylor, S.R., and others, 1973). This similarity may 

support the photogeologic observation that Serenitatis material was 

part of the Imbrium-basin impact target. Alternatively, it may indicate that low-K KREEP is a common crustal constituent (Ryder and 

Wood, 1977).

The lithologic similarity to the Serenitatis melt-rich deposit has 

led to suggestions that the black-and-white rocks are part of a meltrich deposit of the Imbrium basin (Ryder and Bower, 1977; Ryder and 

Wood, 1977). The 40Ar-39Ar gas-retention age of 3.86 + 0.04 aeons 

(table 10.1; Alexander and Kahl, 1974) probably dates the assembly of 

the breccia deposit and, therefore, the Imbrium impact.

FIGURE 10.9. Area of Apollo 15 astronaut traverses. Numbered circles are sampling stations. Prepared by U.S. Defense Mapping Agency; base from Apollo 15 frames P-9370 and

P-9377.
202 THE GEOLOGIC HISTORY OF THE MOON

Material ofMontes Apenninus

Four textural types of Imbrium-basin ejecta lie on the southeast 

Apennine flank (fig. 10.14; table 4.3). One type has been given the 

informal name "material of Montes Apenninus" (Wilhelms and 

McCauley, 1971), not to be confused with the massif material. This 

unit is characterized by closely spaced imbricate slivers concentric 

with the Apennine crest (figs. 10.1,10.7,10.14). Its position near the 

Apennine crest and the coherent morphology of the blocks are consistent either with structural deformation along the excavation rim or 

with low-velocity lofting or pushing of late ejecta over the rim 

(Wilhelms, 1980). The unit resembles the concentric ejecta facies of 

the Orientale basin in morphology, relative position, and areal extent 

(figs. 4.4£, F; Wilhelms, 1980).

FIGURE 10.10. Thin section of crystalline, probably volcanic KREEP-rich sample 

15386,8, showing basaltic intersertal to subophitic texture (Dowty and others, 1976). 

Fragment lacks clastic inclusions and meteoritic siderophile elements (Irving, 1975, 1977; 

Dowty and others, 1976). Crossed polarizers; field of view, about 2 mm.

FIGURE 10.11. Part of Monies Apenninus, showing paucity of outcrops and blocks on 

slope, and colluvial apron at base (a). Apollo 15 frame P 9804.

FIGURE 10.12. Monies Caucasus, showing lineations oblique lo Irend of mounlains and 

radial lo cenler of Imbrium basin. Hummocky Alpes Formalion is above craler Cassim (C; 

57 km, 40° N., 5°E., probably Lower Imbrian) and wesl of Eudoxus (E; 67, 44° N., 16° 

E., probably Copernican), and in part of Monies Apenninus al bollom. A, Arisloteles (87 

km, 50° N., 17° E., probably Eraloslhenian). Orbiler 4 frame H 103.
10. LOWER IMBRIAN SERIES 203

Alpes Formation

Another, more extensive hummocky Imbrium-basin deposit is 

the Alpes Formation (Page, 1970). In contrast with the material of 

Montes Apenninus, the Alpes Formation consists of randomly oriented knobs and larger tracts of smoother material (figs. 10.1,10.12, 

10.14 10.16). This morphology closely resembles that of the Montes 

Rook Formation of Orientale (figs. 4.45, E\ The two units are also

similar in spatial relation to massifs that are considered to be parts of 

major rings (Montes Alpes and the outer Montes Rook ring), which lie 

inside the topographic rims of their respective basins. The Alpes 

Formation, however, extends even farther from the Imbrium rim as 

far as 600 km in the east than does the Montes Rook Formation 

from Montes Cordillera (pi. 8; fig. 4.8). Both units were probably 

emplaced as low- velocity, late-stage ejecta from the inner parts of the 

basins. The Alpes Formation is the more extensive because the 

Imbrium basin is larger than Orientale.

Fissured plains

A third type of material on the Apennine flank forms level or 

wavy pools resting in depressions in the other units (fig. 10.14; 

Wilhelms, 1980). Some of these pools, as well as smooth veneers on 

rougher material, are cut by fissures of the types that suggest 

shrinkage of a formerly molten material (see chap. 3). This Apennine 

material is thus likely to be impact-melt rock like that occupying 

comparable positions on the flanks of craters and the Orientale basin 

(fig. 4.4B).

FIGURE 10.13. "Black-and-white" breccia sample 15455 from rim of Spur Crater, station 

7, Apollo 15 landing site. Black component is impact melt, probably formed by Imbrium 

impact (Ryder and Bower, 1977). White component is pre-Nectarian anorthositic nonte 

(table 8.4). 

A. Hand specimen after arrival from the Moon.

B. Thin section of 15455,29, showing chaotically distributed ANT clasts in melt matrix. 

Plane-polarized light; field of view, about I mm.

By*!

FIGURE 10.14. Types of Imbrium ejecta on flank of Montes Apenninus. ap, northwest-facing Apennine scarp. Concentric "material of Montes Apenninus" is closest to scarp; farther out are 

knobby Alpes Formation (al), radially grooved terrain (g), smooth terrain (s), transverse fissures (f), and level ponds of probable impact melt (m). Arrows, irregular craters in mare, including "D 

caldera" (arrowhead in right frame; chap. 5; compare fig. 5.1OD); d, dark-mantling material. C, fresh crater Conon (22 km; compare fig. 10.7). Stereoscopic pair of Apollo 17 frames M-l 821 

(right) and M-1823 (left). From Wilhelms (I960, fig. 13).
204 THE GEOLOGIC HISTORY OF THE MOON

FRA MAURO FORMATION

Regional relations

A fourth unit on the Apennine flank, better developed farther 

out, is the Fra Mauro Formation (figs. 10.15-10.20; Eggleton, 1964, 

1965; Wilhelms, 1970b). The Fra Mauro Formation extends 600 to 800 

km from the rim of Imbrium in all exposed sectors (pis. 3, 8), and 

grades outward from coarsely hummocky to smooth and wavy (figs. 

10.1, 10.14-10.16). The Fra Mauro Formation obliterates almost all 

prebasin craters to distances of 350 to 600 km from the Apennine 

crest. Thus, the Imbrium-basin deposits resemble exterior deposits of 

craters, whose continuous radial facies also extend beyond a coarsetextured zone out to 1.0 to 1.5 radii from the crater rim (chap. 3; Moore 

and others, 1974; Oberbeck and others, 1974). This similarity to 

craters was the chief basis for concluding that the Fra Mauro Formation is primary ejecta from the Imbrium basin (Shoemaker and 

Hackman, 1962).

The Fra Mauro Formation also resembles the Hevelius Formation of Orientale (fig. 4.4), except that the Fra Mauro includes more 

equidimensional hummocks and fewer distinct ridges. The continuous facies of the Hevelius extends 300 to 600 km from the rim of 

Montes Cordillera, 465 km in radius, which is the probable boundary 

of excavation of Orientale (see chap. 4). These similarities in morphology and extent relative to the diameters of the sources indicate 

that the Fra Mauro, Hevelius, and crater deposits form a related class 

of impact deposits. The inner and much of the outer facies of the 

Hevelius Formation include thick deposits that appear to have flowed 

along the surface (fig. 4.4), and the evident thickness and blockage by 

obstacles indicate that the Fra Mauro is also a major ground-flow 

deposit (Morris and Wilhelms, 1967; Wilhelms, 1980). Grooves and 

ridges indicative of the flow are also present (figs. 10.15 10.17,10.19, 

10.20). Small hills that have been interpreted as volcanic (fig. 10.17; 

Wilhelms, 1968; Wilhelms and McCauley, 1971) may, instead, be 

"deceleration dunes" formed like those of the transverse facies of the 

Hevelius Formation (figs. 4.4G, H\

Setting of the Apollo 14 landing site

The Apollo 14 landing site (figs. 10.18-10.21; lat 3.7° S., long 17.5° 

W) was selected to obtain samples of primary ejecta from the 

Imbrium basin (Eggleton and Offield, 1970). Ridges and grooves 

radial to Imbrium and typical of the Fra Mauro Formation characterize the area around the landing site and most of the northern part 

of the north-south-elongate terra on which it is located (figs. 10.19, 

10.20). The site is 550 km south of Montes Carpatus, which constitutes the south topographic rim of Imbrium and the boundary of 

the excavation cavity (fig. 10.18). Taking the Imbrium impact point as 

lat 33° N., long 18° W, the excavation cavity has a radius of about 540 

km in this sector. Secondary craters are buried near the landing site 

(fig. 10.20; Eggleton and Offield, 1970) and are invisible farther north. 

Despite appearances that suggest dominance by primary ejecta at 

one-radius distances, the currently most hotly debated issue about 

the Apollo 14 rocks is the proportion of primary ejecta to nonbasin 

material in the sample collection.

Sample provenance

The samples returned by Apollo 14 were collected on two traverses outside the Lunar Module (LM) in February 1971 (fig. 10.21; 

Swann and others, 1971, 1977; Sutton and others, 1972). The first

FIGURE 10.15. Several facies of Imbrium ejecta and an extensive stratigraphic section of 

older and younger units. Knobby, nonlineated Alpes Formation (bottom) is transected by 

Vallis Alpes (VA). All pre-Imbrian craters north of Mare Frigoris (F; compare fig. 10.6) 

are transected by Imbrium sculpture and filled by Lower Imbrian plains deposits. W. Bond 

(W; 158 km, 65° N., 4° E.) is also filled by hummocky and braided Fra Mauro Formation. 

W. Bond, Barrow (B; 93 km), and Goldschmidt (G; 120 km) are probably pre-Nectarian. 

Epigenes (E; 55 km) has a higher wall, despite its smaller size, and so it is probably 

Nectarian (not covered earlier by Humboldtianum deposits). Protagoras (P; 22 km) and 

Timaeus (T; 33 km) are not affected by Imbrium basin but are peripherally flooded by 

Imbrian mare basalt, and so they are Imbrian; Protagoras is degraded, possibly Lower 

Imbrian; Timaeus is fresh, Upper Imbrian. Archytas (A; 32 km) is fresh and superposed on 

mare, and so is probably Eratosthenian. Anaxagoras (An; 51 km, compare fig. 3.34) is very 

fresh and bright-rayed (see W. Bond floor), and so it is Copernican. Orbiter 4 frame 

H-116.
10. LOWER IMBRIAN SERIES 205

traverse (extravehicular activity [EVA]), west of the landing site, was 

relatively short but resulted in the collection of many fragments from 

the regolith and some large rocks. Most of the hand specimens of 

coherent rocks were collected on the second EVA, whose main goal 

was the ejecta of Cone Crater, 1.1 km northwest of the landing site 

(figs. 10.21 10.24). Cone is a Copernican crater, 25 million years old 

(table 13.2), 370 m across, and at least 75 m deep (Swann and others, 

1977, pis. 1, 2). Interpretation of the rocks' ultimate provenance 

depends on whether Cone Crater sampled the Fra Mauro Formation, 

the underlying bedrock, or only the regolith built on the Fra Mauro 

Formation. Estimates of regolith thickness at the landing site range 

from 5-12 m, predicted from observations of small craters (Eggleton 

and Offield, 1970), through 8.5 m, inferred from seismic-refraction 

measurements (Cooper and others, 1974), to about 35 m, estimated 

theoretically on the basis of superposed craters (Moore and others, 

1980b, p. 13 14). Because Cone Crater lies on a ridge crest, the 

regolith that it penetrated is unlikely to be thicker than even the 

largest of these values. Thus, Cone Crater probably penetrated the 

regolith and excavated the bedrock unit, which is identified as a 19- to 

76-m-thick layer with a seismic velocity of 299 m/s (fig. 10.23; Kovach 

and Watkins, 1972; Chao, 1973).

Interpretations of what bedrock material the Cone impact excavated from beneath the regolith depend partly on interpretations of 

the ridge. The ridge is 50 to 100 m high (Sutton and others, 1972; 

Swann and others, 1977, pis. 1,2). Hawke and Head (1977b) proposed 

that it is underlain by a large pre-Imbrian crater, in which case Cone 

might have excavated some of the crater material. Eggleton and 

Offield (1970), however, interpreted the ridge as one of the group that 

is characteristic of the Fra Mauro Formation and estimated a Fra 

Mauro thickness of 100 to 200 m on the basis of the ridge relief. In this 

interpretation, Cone probably did not penetrate beneath the Fra 

Mauro Formation. Even if the ridge is underlain by a pre-Imbrian 

crater, a substantial thickness of the Fra Mauro must overlie it (fig. 

10.23; Chao, 1973). Because the bulk of a crater's ejecta comes from the 

upper part of the target material (Stdffler and others, 1974), the hand 

specimens of coherent rock from Cone's ejecta are likely to consist of 

Fra Mauro bedrock (Swann and others, 1977). Many of the smaller 

fragments could have been derived from regolith that overlay the Fra 

Mauro before the Cone impact; these fragments would also consist 

mostly of Fra Mauro material, recycled by impacts. However, they 

probably would also include bits of pre-Imbrian material dug from 

beneath the Fra Mauro Formation by craters older and larger than 

Cone (Head and Hawke, 1975; Swann and others, 1977, p. 28).

Samples important for interpreting and dating the Fra Mauro 

Formation were also collected, on both EVA's, from relatively smooth 

terrain near the LM and at stations A, B, E, F, and G (fig. 10.21). Before 

the mission, this smooth unit was interpreted alternatively as a facies 

of the Fra Mauro Formation or as a superposed volcanic deposit 

(Eggleton and Offield, 1970). It is clearly not volcanic and appears to 

be sufficiently similar in composition and lithology to the Cone Crater 

material as to be thought part of the Fra Mauro Formation (Swann 

and others, 1977). However, the relations among the parts of the 

landing-site region are somewhat unclear because of limited data 

gathered in the field.

Petrology

The coherent rocks from the Cone Crater ejecta exhibit flow 

banding, large-scale stratification, and strong thermal effects, features that suggested to early analysts an origin in a vigorously

FIGURE 10.16. Outward gradation of Imbrium deposits. H, Rima Hyginus (compare figs. 

5.10£, 6.16A). Knobby ejecta (Alpes Formation, lal) grades to smoother Fra Mauro 

Formation (If) in and near Monies Haemus. Surface flow of Imbrium ejecta apparently was 

deflected by Imbrium-secondary chain Boscovich (B; Boscovich Formation was interpreted 

as volcanic by Morris and Wilhelms, 1967} Terrain south of Hyginus, consisting of preImbrian craters (for example, Lade [L; 56 km]; compare fig. 10.17), is scored by Imbrium 

lineations and is mapped either as pre-Imbrian lineated terrain (Wilhelms and McCauley, 

1971) or as part of Imbrium zone (this volume). Arrow, possible Orientale-basin secondary 

chains; Orientale is centered 3,100 km to west in direction of arrow, d, dark-mantling 

deposits near Sulpicius Callus (compare fig. 1.7) and north of Rima Hyginus. Fresh craters 

include Manilius (M; 39 km, 14.5° N., 9° E., probably latest Imbrian), Agrippa (A; 44 

km, younger than mare and probably Eratosthenian), and Godin (G; 35 km, younger, more 

sparsely cratered, brighter, and with a higher reading in thermal infrared than Agrippa; 

probably Copernican). Overlaps with figures 10.1, 10.35, and 10.38. Orbiter 4 frame 

H-97.
206 THE GEOLOGIC HISTORY OF THE MOON

FIGURE 10.17. Lineated terrain at top (north) is Fra Mauro Formation, probably superposed on Imbrium-secondary chains. Crater Lade (bottom; 56 km, I ° S., 10° E.) is filled with hummocky

deposit, probably consisting of Fra Mauro material that flowed over rim and decelerated. Apollo 11 frame H 4552.
10. LOWER IMBRIAN SERIES 207

FIGURE 10.18. Part of geologic map of lunar nearside (Wilhelms and McCauley, 1971), including Apollo 12 and 14 landing sites (arrows). If, Fra 

Mauro Formation; Em, Eratosthenian mare material (Apollo 12; see chap. 12). Montes Carpatus, 550 km north of Apollo 14 landing site, 

constitutes south rim of Imbrium basin. Ring passing near landing sites is part of Insularum basin in one interpretation, but may be ring of 

Procellarum basin (see chap. 8; pis. 3 7).
FIGURE 10.19. Type area of the Fra Mauro Formation. Arrow, Apollo 14 landing site. Fra Mauro Formation, Imbrium sculpture, and plains deposits overlie pre-Nectanan or 

Nectarian craters Fra Mauro (F; 95 km), Bonpland (B; 60 km), and Parry (P; 48 km); Fra Mauro Formation is overlain by Imbrian crater Gambart (G; 25 km, 1 ° N., 15° W.). 

Orbiter4frameH-113.
FIGURE 10.20. Vicinity of Apollo 14 landing site, including Fra Mauro Formation (F) in sharp contact with mare material (M) at triangle. The Fra Mauro Formation is 

strongly lineated radially to Imbrium basin (compare fig. 10.19). Cone Crater is barely visible at tip of arrow. P, buried pre-Imbrian crater. Apollo 12 frame

H-7597.

FIGURE 10.21. Apollo 14 astronaut traverses. Pennant, landing point; ALSEP, geophysical instruments; 

lettered and numbered circles, sampling stations. Numerical scale refers to original map. Airbrush 

drawing by U.S. Geological Survey, published by U.S. Defense Mapping Agency for NASA.

 N * . ? ^ 

. ' ' ' £ ;<
210 THE GEOLOGIC HISTORY OF THE MOON

emplaced, hot ejecta blanket (fig. 10.22; Engelhardt and others, 1972; 

Quaide and Wrigley, 1972; Wilshire and Jackson, 1972). Heat-induced 

textures were first interpreted as the products of metamorphism or 

annealing (Quaide and Wrigley, 1972; Warner, 1972; Wilshire and 

Jackson, 1972) but are now widely considered to be the result of

crystallization of impact melts (Engelhardt and others, 1972; Ryder 

and Bower, 1976; Lofgren, 1977). The impact-melt rock is crystallinematrix breccia that grades from clast-free (fig. 2.8) to heavily laden 

with fragments (fig. 10.24). The most common types of this gradational series consist of more than 15 percent clasts in gray, coherent,

FIGURE 10.22. Large boulders near rim of Cone Crater, containing light and dark materials. From Swann and others (1977, fig. 44).

WSW ENE

EXPLANATION

Pre-Fra Mauro 

material

500 METERS

FIGURE 10.23. Schematic geologic cross section of Apollo 14 landing site (from Chao, 1973). Section is drawn just north of landing site in west to Cone 

Crater in east. Thicknesses of geologic units are based mostly on active-seismic data (Kovach and Watkins, 1972). Compressional-wave velocities of some 

units are indicated. Elevations in vertical scale are based on reference elevation chosen to correspond to mean radius of Moon, 1,738 km.

FIGURE 10.24. Large (8.9 kg) Fra Mauro breccia sample 14321, "Big Bertha," from rim 

of Cone Crater, Apollo 14 landing site: a polymict breccia displaying multiple episodes of 

brecciation (Duncan and others, 1975; Grieve and others, 1975). 

A. Entire hand specimen, showing complex relations of clasts and matrix.

B. Thin section of sample 14321,197, showing abundant mineral and lithic clasts set in 

dark, fine-grained matrix. Light-colored, coarsely crystalline materials along right and 

bottom edges are clasts of high-Al mare basalt (table 9.5). Nonpolarized light; field of 

view, about 1.5 cm.
10. LOWER IMBRIAN SERIES 211

melt-derived matrices (Chao and others, 1972; Engelhardt and others, 1972; Wilshire and Jackson, 1972; Chao, 1973; Simonds and 

others, 1977). These clast-laden breccia samples were collected from 

both Cone Crater and the smooth terrain; because they characterize 

the landing-site region, they have been termed Fra Mauro breccias 

(Chao and others, 1972). Light-colored friable breccia samples ("white 

rocks") are also common but are probably clasts previously included 

in the crystalline-matrix breccia (Chao and others, 1972). Samples of 

almost fragment-free impact melt (Green and others, 1972; Taylor 

and others, 1972; James, 1973; McKay and others, 1978,1979) were 

collected from the regolith on the smooth terrain; the largest are 

samples 14073,14276, and 14310 (fig. 2.8). At one time, these samples 

were widely considered to be volcanic (see chap. 8).

The samples are typically rich in KREEP. Most of the crystallinematrix breccia seems to be relatively uniform in composition; 12 of the 

16 samples analyzed by Simonds and others (1977) cluster together on 

a plot of MgO content versus A12O3 content as does "vitric matrix" 

breccia, consisting of consolidated regolith fragments; not further 

discussed here). Clast-free rock samples 14276 and 14310 contain less 

MgO and more A12O3 than the cluster. McKay and others (1979) found 

that the chemical and isotopic differences among the KREEP-rich 

impact-melt samples are consistent with an origin in a single melt 

sheet.

Emplacement process

A debate about emplacement mechanism persists even among 

those who agree that most samples are from the Fra Mauro Formation and that the Fra Mauro is somehow related to the Imbrium basin 

(fig. 10.25). One view is that the rocks that compose the Fra Mauro are 

parts of the primary ejecta of the basin (Eggleton and Offield, 1970; 

Chao and others, 1972; Dence and Plant, 1972; Engelhardt and others, 1972; Quaide and Wrigley, 1972; Wilshire and Jackson, 1972; 

Chao, 1973). That is, they were parts of the target material that 

traveled at least 550 km to their site of deposition (hypothesis b, fig. 

10.25). This view is consistent with the generally uniform composition 

of the Fra Mauro breccia samples. Either thorough mixing of target 

materials or derivation from a compositionally uniform source, such 

as a KREEP-rich crustal layer or a KREEP-basalt section in the 

Procellarum basin, could have produced this uniformity. The primary-impact mechanism also allows for variations in composition 

and shock grades because of the general disorder of the basin-forming 

process. The differing compositions and shock grades of the friable, 

light-colored materials and the Fra Mauro breccia are readily comprehended in this model as the results of formation in different parts 

of the shock-engulfed target zone, followed by mixing during ejection 

or emplacement (Engelhardt and others, 1972). That is, some of the 

already heterogeneous bedrock-breccia, regolith, and volcanic deposits in the target zone were little altered by the Imbrium impact, 

whereas others were more highly shocked, melted, and homogenized.

The opposing view is that the crystalline-matrix breccia originated at or near the deposition site (Oberbeck and others, 1974; 

Morrison and Oberbeck, 1975; Head and Hawke, 1975; Oberbeck, 

1975; Stoffler and others, 1976; Hawke and Head, 1977b, 1978a, b; 

Simonds and others, 1977). The proponents of this "local" hypothesis

FIGURE 10.25. Geologic cross section illustrating secondary-ejecta (a) and primary-ejecta 

(b) hypotheses for emplacement of the Fra Mauro Formation, drawn along north-south line in 

figure 10.18. In most secondary-ejecta models, Imbrium cavity is bounded by an inner ring, 

ejection trajectories are steep (dashed arc), and transition between primary and secondary 

deposits is close to rim (dashed sawteeth beneath right-hand letter a). In primary-ejecta model 

preferred here, basin rim is at Monies Carpatus, ejection trajectories are shallow (solid arc), 

and transition from primary to secondary deposits occurs south of Apollo 14 landing site 

(solid sawteeth beneath right-hand letter b). Secondary craters form in both models (diagrammatically indicated by notches in substrate): in hypothesis a, their ejecta forms sampled part of 

the Fra Mauro Formation; in hypothesis b, they are deeply buried by primary ejecta at 

Apollo 14 landing site.

are troubled by the diversity of textures and compositions of the 

breccia samples and by the great distance of the Apollo 14 landing site 

from the Imbrium basin. They believe that the terrestrial analogs 

(Chao and others, 1972; Dence and Plant, 1972; Simonds and others, 

1977) show that the abundant melt found in the breccia deposits could 

not have been transported so far. In the local-derivation models, the 

Fra Mauro Formation (and analogous deposits of Orientale and 

smaller craters) consists of mixtures of the primary ejecta with material derived from the local terrane by secondary impacts (hypothesis 

a, fig. 10.25). The radial ridges of the Fra Mauro that are interpreted 

as flows of primary ejecta in the first model are considered in the 

local-derivation model to be analogous to the herringbone texture 

visible among secondary craters of Copernicus-type craters (fig. 3.4), 

or to surges of debris that continued to flow outward after their 

generation by secondary impacts. Such coalesced debris flows may 

have erased the original secondary-ejecta pattern. On the basis of the 

upper equation in figure 10.26, Morrison and Oberbeck (1975) calculated that only 15 to 20 percent of the Fra Mauro Formation at the 

Apollo 14 landing site should consist of primary ejecta. Head and 

Hawke (1975) pointed out craters and basins whose melted and fragmental deposits are reasonable sources of the local materials. They 

considered the KREEP-rich compositions to have been derived from 

volcanic extrusions in the vicinity of the landing site (Hawke and 

Head, 1978b). They stated that mineralogically recorded shock pressures in the narrow range 20-30 GPa should be of Imbrium-basin 

origin and that less highly shocked and, possibly, more highly 

shocked material should be of local origin (Hawke and Head, 1978a). 

Both the primary-ejecta model and the secondary-impact, localmaterial model probably contain some truth. Great diversity in shock 

grade and other lithologic attributes is surely possible at a given point 

in the ejecta deposit of a large impact. On the other hand, the pervasiveness of secondary impacts at great ranges from basins cannot 

be denied, and incorporation of the impacted local material is to be 

expected The issue is thus one of proportion. I believe that primary 

ejecta is the dominant component and that the proponents of secondary emplacement of the entire Fra Mauro Formation have overstated 

their case, for several reasons (Wilhelms and others, 1980):

= 46.4

= 30.5

sin20

+ sin 9

2 tan 60.6

sin 29 . .2

2 tan R

 » sin 8

60.6

-0401 1-134-

D cos 9

-060 I 20 0

D cos 9

10 D =

0 = 75

10 100 1000 

RANGE (km)

FIGURE 10.26. Mass ratio of primary versus local material as a function of range from 

source of incoming ejecta. Dashed lines are based on upper equation, which assumes scaling 

of D = KE ~ 3 - 4 (Oberbeck and others, 1975); solid lines are based on lower equation, which 

assumes scaling of D = KE~ 3 6 (H.J. Moore, written commun., 1980). (X, ratio of mass 

ejected from secondary crater to mass of ejecta projectile that created crater; 0, angle of 

impact (equal to angle of ejection), measured from local surface normal; D, diameter of 

secondary crater (in kilometers); R, range or distance along Moon's surface from point of 

ejection within primary crater or basin to secondary crater (in kilometers). Courtesy of H.J. 

Moore.
212 THE GEOLOGIC HISTORY OF THE MOON

1. The topography of the landing-site region is dominated by ridges 

formed by flow of a thick deposit that has obscured the secondary 

craters. This flow originated either at the basin rim or at secondary-impact sites much closer to the basin than its present resting 

place. Either circumstance would change the "range" input in the 

equation (fig. 10.26).

2. The equation in figure 10.26 is model-dependent in other ways. The 

assumed ejection angle, size of the Imbrium basin, point of derivation within the basin, size of secondary craters (which are difficult 

to see and measure near the landing site), and diameter-energy 

scaling laws are all questionable. An apparently small change in 

the exponent of the energy in the scaling relation from 3.4 to 

- 3.6 changes the equation to the lower form in figure 10.26 and 

greatly affects the results of calculations. Moreover, the equation is 

based on single-body impacts (Morrison and Oberbeck, 1975, p. 

2527), whereas most basin ejecta at great ranges is likely to be 

comminuted (Schultz and Mendell, 1978). Many small low-density 

fragments impacting over an extended time at low velocities will 

build up a deposit of their own material rather than excavate local 

materials greater than their mass (Chao and others, 1975; Morgan 

and others, 1977, p. 672).

3. Photogeologic relations indicate the presence of abundant melt in 

basin-ejecta deposits, (a) Material ponded in depressions in the 

inner Orientale and Imbrium ejecta displays "mud crack" patterns 

of fine fissures, probably resulting from shrinkage of a coherent 

material (figs. 4.4E, 10.14; Moore and others, 1974). These fissured 

ponds resemble the external impact-melt pools of large craters, 

which increase in abundance and distance from the rim with 

increasing crater size and occur as far as two crater radii from 

some rims (chap. 3; Howard and Wilshire, 1975; Hawke and Head, 

1977a). (b) Crater ejecta is coated by additional melt in the form of 

veneers that are gradational with the ponded melt (fig. 3.36; 

Howard, 1975; Howard and Wilshire, 1975; Hawke and Head, 

1977a). (c) Large tongues of nontextured material apparently have 

segregated from the radially lineate Hevelius Formation and have 

formed distant lobes (fig. 4.4D). The lobate, commonly leveed morphology of these segregates suggests emplacement as viscous fluids, for which melt is a reasonable, though not exclusive, 

explanation (Eggleton and Schaber, 1972; Moore and others, 1974). 

Finally, (d) pools in the floors of many secondary craters appear to 

have once been molten, an appearance suggesting considerable 

ranges for ejected melt (figs. 3.9A, 4.4C; Moore and others, 1974; 

Schultz and Mendenhall, 1979).

4. Scaling laws suggest that large, basin-forming impacts generate 

more melt than small, crater-size impacts on either the Moon or 

the Earth (chap. 3; Rehfuss, 1974; O'Keefe and Ahrens, 1977; 

Lange and Ahrens, 1979; Melosh, 1980, p. 77). An increase in the 

ratio of projectile mass to crater size is predicted with increasing 

magnitude of impact (Baldwin, 1963, chap. 8; O'Keefe and Ahrens, 

1975; Gaffney, 1978). The impact energy of large bodies is coupled 

into the target longer, so that the shock energy is distributed over a 

larger volume of rock. In contrast, small fast bodies strike and 

rapidly release their kinetic energy into small, intensely heated 

volumes of rock. The result is that large impacts generate larger 

volumes of less highly heated melt than do small impacts. As a 

corollary, melt of the large impacts would presumably be less 

thoroughly homogenized than that of small impacts.

5. Finally, the target may have been hot already, so that any impact 

would create more melt than would an impact into a cold target 

(Wood, 1975d, p. 515). Evidence of heat at crustal depths is contained in the coarse grain size of Apollonian metamorphic rocks 

(Stewart, 1975) and the textures of granulitic impactites, ascribed 

to accumulated heat from hot ejecta blankets (Warner and others, 

1977). Radiogenic heat may have melted KREEP-rich magmas 

before (see chap. 9) and after (Apennine Bench Formation) the 

Imbrium impact, and liquid magma may even have been ejected 

(Schultz and Mendenhall, 1979).

These points are offered not as proof of a primary-ejecta origin 

but as reasonable alternatives to the local-origin models currently 

widely accepted. The data of Simonds and others (1976b, 1977) also 

seem to be consistent with a primary-ejecta origin, although those 

authors favored the local-origin model. Simonds and others (1976a, b) 

suggested that the Fra Mauro crystalline-matrix breccia formed as 

superheated impact melt plus cold fragmental debris and that the

differences in clast content are solely responsible for the textural 

differences. As discussed above, the compositional data (Simonds and 

others, 1977; McKay and others, 1979) also seem to be more consistent 

with a limited range of source compositions than with the diverse 

"local" target cited by many authors. The main support for the localderivation model seems to be (1) terrestrial-analog craters whose 

ejecta has not been identified (Canadian-shield craters), and (2) the 

upper equation in figure 10.26 (Morrison and Oberbeck, 1975), which 

is highly model-dependent and is qualitatively contradicted by the 

absence of conspicuous Imbrium-secondary craters in the appropriate uprange positions to have been the sources of the Fra Mauro 

Formation.

Absolute age

Two age clusters have been identified in the returned samples by 

both the 40Ar-39Ar (Turner, 1977) and Rb-Sr (Papanastassiou and 

Wasserburg, 1971b) methods, which agree well for these rocks (table 

10.1). The younger cluster was obtained mainly for some clast-poor 

impact-melt rock samples (14073,14276,14310), for a few samples of 

Fra Mauro breccia, and for small fragments from mixed regolith 

samples. The measured ages of the clast-free rocks average 3.82 

aeons; of the clast-free fragments, 3.83 aeons; and of other melt-rock 

samples of uncertain origin, including some that may belong to the 

older cluster, 3.84 aeons (table 10.1). If the dated samples are of 

Imbrium-basin impact-melt rock, as I believe, the basin formed 3.82 

to 3.84 aeons ago, with an uncertainty of only a few tens of millions of 

years. However, the uncertain field relations between samples from 

the smooth terrain and from Cone Crater do not preclude the possibility that some samples are derived from melts of a post-Imbrium 

crater.

The best dated older group of Apollo 14 samples (extremes and 

stated analytical errors larger than 0.1 aeon eliminated) yields ages of 

3.87 to 3.96 aeons (tables 9.5, 10.1). All three large samples (14053, 

14072, and clasts of 14321) came from the Cone Crater flank, and all 

are of high-Al mare basalt (fig. 10.24; table 9.5; Grieve and others, 

1975; Ridley, 1975; Taylor, 1975; Simonds and others, 1977; Ryder and 

Spudis, 1980). They differ further from the younger group in their 

lower initial Sr-isotope ratios (Papanastassiou and Wasserburg, 

1971b). Although error bars overlap, these rocks form a distinct age 

group from the younger rocks. Clasts of breccia sample 14066 and 

some fragments from regolith samples 14161 and 14167, all collected 

within 400 m of the LM, may also belong to the older cluster 

(lithologies uncertain; listed with the younger group in table 10.1). 

The older clasts and regolith fragments appear to be fragments of 

pre-Imbrian rock, including volcanic basalt, incorporated into the 

Imbrium-ejecta breccia before its emplacement near the future landing site (chap. 9; Mark and others, 1975; Wetherill, 1977b). This 

basalt-in-breccia relation is consistent with either the primary- or the 

secondary-ejecta model; beds of basalt could have existed before the 

Imbrium impact either in the Imbrium target area (Wilshire and 

Jackson, 1972) or in the vicinity of the Apollo 14 landing site (Hawke 

and Head, 1978b). In either model, the age of the Imbrium basin is 

constrained as equal to (Compston and others, 1972) or less than 

(Wetherill, 1977b) 3.87 aeons.

Conclusions

The following working hypotheses (fig. 10.25) are consistent with 

all data from the samples known to me and with Moon-wide relations. 

(1) The crystalline-matrix clast-laden Fra Mauro breccia excavated 

from Cone Crater is sufficiently uniform structurally and compositionally to indicate derivation from a single major rock unit. (2) That 

unit is part of the photogeologically observable ridgy unit, the Fra 

Mauro Formation, which is composed mostly of primary ejecta of the 

Imbrium basin. (3) Part of the target material of the Imbrium impact 

consisted of mare basalt. (4) Ages of basalt clasts in the Fra Mauro 

breccia deposits constrain the age of the formation to 3.87 aeons or 

younger. (5) Ages of clast-poor impact-melt rocks in other parts of the 

landing-site region constrain the age of the formation to 3.82 aeons or 

older. (6) These results are consistent with the Apollo 15 ages (table 

10.1). (7) The likely time of the Imbrium-basin impact is narrowed to 

3.845 ±0.03 aeons ago by these combined results.
10. LOWER IMBRIAN SERIES 213

IMBRIUM-BASIN 

SECONDARY CRATERS

Like the Hevelius Formation at Orientale, the Fra Mauro Formation grades outward from a thick and continuous coarsely textured 

blanket, through plains deposits, to a wide zone of secondary craters. 

Southeast of the Imbrium basin, the transition from plains to secondaries occurs about at Sinus Medii, 450 to 750 km from the Apenninus 

crest or 1,150 to 1,450 km from the basin center (figs. 10.16, 10.27, 

10.28). Due south, the Fra Mauro Formation and mare materials have 

mostly buried the secondary craters, some of whose rims protrude 

around the margins of the Fra Mauro peninsula (fig. 10.19). In the 

north, this transition is also poorly defined because so many rugged 

linear chains of the Vallis Bouvard type (fig. 4.4D) are visible through 

the Fra Mauro continuous ejecta (figs. 10.6,10.15). The transition is 

estimated to occur about 300 to 400 km north of Mare Frigoris or 1,000 

to 1,100 km from the basin center (pi. 3). In all observed sectors, the 

transition occurs farther from the basin than for Orientale, as 

expected for a topographically larger basin.

The most striking characteristic of the zone of secondary craters 

is the "Imbrium sculpture" (figs. 10.15, 10.27-10.29). Most of this 

sculpture, once thought to be fault-controlled (see chap. 6), is now 

thought to consist of chains of elliptical secondary-impact craters 

(compare figs. 10.27,10.29). Although some of the finer outer sculpture seems to have formed by surface flow (fig. 10.17), as in the Fra~ 

Mauro closer to the basin and in the Hevelius Formation, the origins 

of individual grooves and ridges cannot be determined so readily as 

around Orientale (Wilhelms, 1980).

Abundant secondary-impact craters appear beyond the 

groovelike chains at distances of about 1,700 km from the basin center 

(figs. 10.30, 10.31). Like the sculpture, the more nearly circular secondaries were once considered to be volcanic or primary-impact craters. One basis for this interpretation was the morphologic diversity 

within a given cluster (figs. 10.28,10.29). However, individual satellitic craters of Orientale and of large craters also differ from one 

another (figs. 3.9, 4.4). Good photographs show that a range in morphologies from very sharp to almost obliterated results from differential burial of a coeval cluster by deposits that flowed from nearer the 

basin (figs. 4.4G, H; Oberbeck, 1975; Wilhelms, 1976; Wilhelms and

ww

»

FIGURE 10.27. Imbrium sculpture in south-central nearside highlands south of Sinus Medii (SM). Craters cut by sculpture include Lade (L; compare figs. 10.16, 10.17), Flammarion (F; 75 km), 

Hipparchus (H; 151 km), Ptolemaeus (P; 153 km), Albategnius (Alb; 136 km, compare figs. 10.29, 10.36), and Alphonsus (Alp; 119 km). Planar fill apparently transects sculpture. 

Subcircular secondary craters of Imbrium are visible below latitude of Ptolemaeus. Catalina Observatory photograph No. 1907.
214 THE GEOLOGIC HISTORY OF THE MOON

others, 1978, p. 3749). The septa between craters and "domes" superposed on their intersections (figs. 3.9B, 4.4C, 10.29 10.31), which 

were once cited as evidence of volcanism, resemble forms created in 

similar positions by the interaction of ejecta in laboratory experiments (fig. 3.7). Only the gross form of the domes and intercrater 

ridges is observed in most of the Imbrium-secondary field. Study of 

Orientale analogs and the laboratory experiments were required 

before the secondary-impact origin of these domes and ridges could be 

supported.

An Imbrium-basin origin of furrows or pits outside the main 

concentration of secondaries (figs. 10.32,10.33) is also likely. A large 

tract west of Mare Humorum (fig. 9.24A), which contains more rugged, hilly topographic elements than most terrain at the same distance from Imbrium (2,000 km from the basin center), was 

interpreted as Humorum-basin ejecta (Titley, 1967) (part of the 

Vitello Formation) or as volcanic (McCauley, 1973; Wilshire, 1973). 

Secondary-impact origin emerged as a reasonable alternative during 

mapping of the whole circum-Imbrium zone. This tract is well within 

the reach of Imbrium ejecta projectiles and contains numerous softappearing, north-northeast-trending Imbrium-radial gougelike 

depressions as long as 30 km. It also adjoins, and is faintly gradational with, pitted terrain related to Orientale (fig. 4.4#). The tract is 

likely to be a composite unit consisting of Humorum ejecta scoured by 

Imbrium-basin secondaries and then mantled by thin Orientale 

ejecta; it is included in the Imbrium-related zone on plates 3 and 8. A 

somewhat similar terrane, previously thought to be volcanic 

(Wilhelms, 1972a; Olson and Wilhelms, 1974), is superposed on the 

south Crisium-basin rim flank, including the vicinity of the Luna 20 

landing site (figs. 10.33,11.16). Imbrium-secondary origin is suggested 

by radiality of the bands of pits and furrows.

Imbrium-basin effects persist farther outward. The first clustered craters that were identified as secondary to the basin (Scott, 

1972b) occur near the craters Riccius and Maurolycus, 1,900 km from 

Montes Apenninus or 2,600 km from the basin center (fig. 3.10Z)). 

Scattered secondary craters of Imbrium occur beyond 3,000 km from 

the basin center and, possibly, in all parts of the lunar terrae. Probable secondaries have been identified at least 3,300 km from the center

FIGURE 10.29. Terrain east of crater Albategnius (truncated by left edge of photograph; 

compare figs. 10.27, 10.28). Imbrium sculpture is generally linear but consists of elliptical 

craters. Domelike intersections of craters and other small hills (arrows) were created by 

multiple secondary impacts (see chap. 3; compare fig. 3.9B), not volcanism. Apollo 16 

frame M-392.

FIGURE 10.28. Strip from Mare Vaporum (top) southward to lat 23.5 S., showing 

transition from smooth, continuous Fra Mauro deposits (compare fig. 10.16) to highly 

diverse secondary craters below latitude of Ptolemaeus and Albategnius (compare figs. 

10.27, 10.36). Mosaic of Orbiter 4 frames H-101 (bottom) and H-102 (top); overlaps 

figures 7.7 and 10.16.
10. LOWER IMBRIAN SERIES 215

on the southeast nearside (fig. 9.2; Stuart-Alexander, 1971; Wilhelms 

and others, 1978, 1979). Linear chains on the farside are so closely 

radial to the basin as to indicate Imbrium origin (fig. 8.13; Wilhelms 

and El-Baz, 1977).

Finally, peculiar grooves at the antipode of Imbrium on the 

farside (figs. 10.32,10.33; Stuart-Alexander, 1978) have been ascribed 

to convergence of Imbrium-basin ejecta (Moore and others, 1974) or to 

convergence of seismic waves from the Imbrium impact shock 

(Schultz and Gault, 1975a, b). Similar terrain north of Mare Marginis 

antipodal to Orientale was noted in chapter 4 (fig. 4.7). All grooved 

terrain may, in fact, be related to basins (fig. 10.33). The Serenitatis 

and Crisium antipodes are similarly grooved (fig. 9.2LA; Hood and 

others, 1979). Strangely, all four of these antipodal grooved regions 

are also characterized by anomalous concentrations of bright swirls 

and by high remanent magnetism as observed from orbit (Hood and 

others, 1979). The cause-and-effect relations of these peculiar coincidences of old topography (the grooves), very young features superposed on maria and lacking intrinsic relief (the swirls), and 

magnetism have yet to be determined. Hood and others (1979) tentatively suggested that the basin impact magnetized the antipodal 

terrain, either by secondary impact or seismicity, so as to deflect the 

solar wind during the rest of geologic time and permit the maintenance of high albedo in the otherwise-fragile swirls.

LIGHT-PLAINS MATERIALS

Around both the Orientale and Imbrium basins, light-colored 

plains are concentrated in the transition zones between the continuous deposits and the secondary craters (figs. 4.4, 10.34, 10.35). The 

clearly visible Orientale relations and the fact that the samples of 

presumably typical plains material returned by Apollo 16 were 

emplaced by impact have combined to shift interpretations from a 

volcanic to an impact origin for terra plains in general (see chaps. 2, 

4). The superposition of plains on the "Imbrium sculpture" (figs. 

10.27,10.35, 10.36), which was once considered prime evidence for a 

volcanic origin (Howard and Masursky, 1968; Milton, 1968a; 

Wilhelms, 1968,1970b), has been shown to be consistent with impact 

emplacement on the basis of similar relations at Orientale. Plains 

occupying depressions thousands of kilometers from Imbrium may 

also be of impact and not volcanic origin because they generally occur 

downrange from secondary-crater clusters (fig. 4.6). Although the 

appearance and crater densities of many plains remain equally consistent with volcanic and impact origins (Neukum and others, 1975a; 

Neukum, 1977), and although volcanic substrates may underlie the 

surficial ejecta in some plains (chap. 9; Eggleton and Marshall, 1962; 

Schultz and Spudis, 1979), impacts evidently formed most of the 

visible plains.

FIGURE 10.30. Chains of Imbnum-secondary craters; arrows point to center of Imbrium 

basin, 2,000 km north. Left arrow is in crater Pitatus (97 km, 30°S., 14°W.). Along chain 

is superposed on crater rim; septa separate many crater pairs of the chain. Chain below right 

arrow shows interference "domes." Mosaic of Orbiter 4 frames H 113 (right) and H 112 

(left). From Wilhelms (1976, fig. 5}
216 THE GEOLOGIC HISTORY OF THE MOON

More patches of well-defined plains are Early Imbrian than of 

any other single age. Burial of the Apennine Bench Formation inside 

the Imbrium-basin rim by the deposits of the craters Archimedes and 

Cassini suggests an Early Imbrian age for this plains unit, and the 

radiometric dates point to an earliest Early Imbrian age. An earliest 

Early Imbrian age is suggested for the plains north of Mare Frigoris 

(fig. 10.15) by textural gradation with the Fra Mauro Formation, 

crater size-frequency values (figs. 7.15,10.37), and visual inspection of 

crater densities. The partly planar Maunder Formation inside the 

Orientale basin almost certainly consists of impact melt contemporaneous with that basin (see chap. 4) and thus is latest Early Imbrian.

DL values of 1,000 to 1,200 m have been identified with the plains 

north of Frigoris and, therefore, with the Imbrium basin (fig. 7.15; 

Boyce and others, 1974; Moore and others, 1980b). Similar values, 

however, have also been obtained for other plains (Boyce and others, 

1974) that are clearly older than the basin (figs. 7.15,8.11, 9.26, 9.28; 

Neukum, 1977; Wilhelms and El-Baz, 1977; Stuart-Alexander, 1978; 

Lucchitta, 1978). The inability of DL values to discriminate among old 

deposits (fig. 7.15) may be due to several factors: (1) Such material 

properties as low cohesion of thick regoliths, and (2) the large sizes of 

superposed craters obviously present (fig. 9.26). For these large craters, (a) the production function may be less than the 2 cumulative 

assumed in DL models (Soderblom, 1970; Soderblom and Lebofsky, 

1972), and (b) initial crater shapes and wall slopes may differ from 

those of the small craters measured during DL determinations. Stratigraphic relations and size-frequency counts are preferred here toD^ 

measurements as indicators of terra-plains age.

The relative age of the Cayley Formation is important to interpretations of the returned materials. On Lunar Orbiter 4 photographs, the Cayley generally seems to have flatter surfaces and 

fewer, sharper craters than the earliest Early Imbrian plains north of 

Mare Frigoris that are clearly contemporaneous with the Imbrium 

basin. (The Cayley seems to be less planar at higher resolution [fig. 

10.38] and on the surface as seen by the astronauts.) DL values of 475 

to 640 m (fig. 7.15; Soderblom and Boyce, 1972; Boyce and others, 

1974), some crater-density studies (fig. 10.37; Boyce and others, 1974; 

Neukum and others, 1975b), and visual inspection also suggest a clear 

temporal hiatus between the plains north of Mare Frigoris and the 

latest Early Imbrian plains at Orientale. The Cayley and circumOrientale plains look alike on Lunar Orbiter 4 photographs. Some 

crater counts and DL measurements also show them to be contemporaneous with each other and with the Hevelius and Maunder Formations (Soderblom and Boyce, 1972; Boyce and others, 1974). Therefore, 

an Orientale age has been suggested for the Cayley plains in the 

central nearside. Suggested mechanisms for creating Orientale-age 

surfaces on plains that are spatially related to Imbrium include

FIGURE 10.31. Short chains of Imbrium secondary craters near craters Stofler (S; 126 km) 

and Licetus (L; 75 km, 47° S., 7° E.), 2,600 km from center of Imbrium basin. Orbiter 4 

frame H-107.

mantling of older, Imbrium-basin-age (and Nectarian) plains by farthrown Orientale ejecta (Hodges and others, 1973; Chao and others, 

1975) and degradation induced by seismic waves from the Orientale 

impact (Schultz and Gault, 1975a). Either mechanism could theoretically have rejuvenated the cratered surface on the older material. 

A dissenting view, based on different evaluations of the same 

data, is that the Cayley Formation is contemporaneous with the 

Imbrium basin. Other crater counts (fig. 10.37; Neukum and others, 

1975a, b; Neukum, 1977), the inability of DL determinations to date 

old plains, the plains' concentration around Imbrium, and the 

improbability that Orientale ejecta was sufficiently thick to cover all 

Cayley patches lead me to support this interpretation.

CAYLEY AND DESCARTES 

FORMATIONS

Introduction

The problem of determining whether the materials composing a 

geologic unit of impact origin were formed before or during the unit's 

emplacement is especially acute for the two units sampled by Apollo 

16. Chapter 9 showed that the distribution and absolute ages of many 

Apollo 16 samples, particularly those from the Descartes Formation, 

are consistent with a Nectaris-basin origin. The overall mountainous 

topography of the Descartes could also be an expression of the Nectaris ejecta. As discussed in the previous section, however, the distribution and planar morphology of the Cayley Formation originated 

in the Imbrian Period. The paucity of overlying primary-impact craters and the topographic sharpness of the furrows (fig. 9.9) suggest 

that at least the surficial part of the Descartes Formation is also of 

Imbrian age. Because few lunar terra surfaces are shaped without 

addition of material, the presence of some Imbrian material in the 

Descartes is expected.

This section discusses the events in the units' histories that 

might have taken place during the Imbrian Period. Three possible 

ramifications of an Imbrium-basin origin are explored first: that the

FIGURE 10.32. Approximate antipode of Imbrium basin (arrow; 38° S., 160° E.). 

Grooves score rims of old craters in vicinity of antipode (Stuart-Alexander, 1978). Bright 

swirls above arrow are in Mare Ingenii. Figure-8 crater at right edge is Van de Graaff (234 

km long, probably Nectarian). Orbiter 2 frame M-75.
10. LOWER IMBRIAN SERIES 217

units consist (1) mostly of Imbrium-basin ejecta, (2) mostly of Nectaris-basin ejecta reworked during the Imbrian Period, or (3) of basin 

ejecta of different provenance (fig. 10.39). Then follows a consideration 

of "local" origins as (1) reworked pre-Imbrian crater materials or (2) 

post-Imbrian materials origins that I consider subordinate in 

importance to a basin origin at the Apollo 16 landing site.

Imbrium-basin origin

Spatial distribution, relative age, and comparison with the Orientale-basin geology, as already discussed, suggest that the agent 

which created the planar surface of the Cayley Formation was the 

Imbrium impact. Similarly, the furrows of the Descartes Formation

Hilly and pitted Hilly and furrowed Furrowed and pitted Grooved

FIGURE 10.33. Fourtypesof lunar terrain with small negative features (after Howard and others, 1974, fig. 13). Original caption states that terrains are not clearly related to basins, but text offers 

basin-related, nonvolcanic interpretations. Antipodal effects may explain patches 1MB (Imbrium) and OR (Orientale) (compare figures 4.7 and 10.32). Pitted and furrowed deposit at long 90° 

E. (in Smythii basin) may be related to Crisium basin.

NEARSIDE 

N

FIGURE 10.34. Terra plains (Howard and others, 1974, fig. 12). Most nearside and eastern farside plains surround Orientale and Imbrium basins except in a one-basin-radius zone occupied by 

continuous basin ejecta; plains occur as far as one-quarter of Moon's circumference from source basin. Other large patches are associated with Nectaris and other Nectarian basins and have been 

reassigned a Nectarian or Nectarian-Imbrian age since this map was prepared (Wilhelms and El-Baz, 1977; Wilhelms and others, 1979).
218 THE GEOLOGIC HISTORY OF THE MOON

seem to be somehow related to the Imbrium basin (Eggleton and 

Marshall, 1962); the grooves of the northern (Smoky Mountain) facies 

of the Descartes are closely radial to Imbrium.

These relations do not specify whether primary- or secondaryimpact mechanisms were more important in the units' emplacement. 

Whereas gradation with the Hevelius Formation indicates that some 

of the Cayley-like Orientale plains deposits consist of pooled primary 

basin ejecta that segregated from the Hevelius (fig. 4.4D; Eggleton 

and Schaber, 1972), similar plains are distal to Orientale-secondary 

craters (fig. 4.4G). The Descartes Formation, though exceptionally 

conspicuous, is similar in general morphology to other circumImbrium hilly-and-furrowed units already interpreted as part of the 

secondary-crater retinue of the Imbrium basin (fig. 10.33). Many 

grooves and scarps of the southern (Stone Mountain) facies, however, 

are transverse to the Imbrium-radial direction (fig. 9.9) and are less 

like secondary craters than like the dunelike deposits of Orientale 

that originated by surface flow of basin ejecta (figs. 4.4F H). Therefore, the major issue being debated, as was the case with the Apollo 14 

materials, is the proportion of Imbrium ejecta to "local" material 

excavated by the Imbrium secondaries.

Diversion by the Kant Plateau of material flowing from Imbrium 

has been suggested as the cause of the transverse trend of the Stone 

Mountain facies (Hodges, 1972; Hodges and others, 1973; Moore and 

others, 1974; Hodges and Muehlberger, 1981). Hodges and Muehlberger proposed that a long trough radial to Imbrium (fig. 10.35) 

facilitated the flow. Flow elsewhere in the region is indicated by a 

thick deposit, evidently from Imbrium, in the crater Andel M west of 

the Apollo 16 landing site (Moore and others, 1974). The Apollo 16 

landing site and many plains-dunes associations around Orientale lie 

two-thirds of a basin diameter from the rims of their respective 

basins.

Most investigators, however, doubt that much primary ejecta 

could have reached the Apollo 16 landing site all the way from the 

Imbrium rim (Montes Apenninus) by surface flow. This viewpoint is 

supported by the abundant evidence for Imbrium secondaries in the 

region. Many pits and deeply etched grooves that seem clearly to be 

secondary to Imbrium are adjacent to the two sampled units (fig. 

10.35). Clustering of a group of craters on the Kant Plateau immediately east of the landing site suggests an origin as Imbrium secondaries. Alignment of the nearly equant pair Dollond B and Dollond C 

northwest of the landing site is also typical of secondary fields, and 

these 32- to 37-km-diameter craters are similar in size to Imbrium 

secondaries in other regions. Any Imbrium material at the site may, 

therefore, have arrived in impacting Imbrium projectiles (Wilhelms, 

1972c) or in a debris flow initiated by secondary impacts of the projectiles (Morrison and Oberbeck, 1975; Oberbeck, 1975; Oberbeck and 

others, 1975).

As for the Apollo 14 landing-site region, an intermediate view is 

more likely than an all-primary or all-secondary origin. A flow initiated by secondary impacts would contain mixed primary and secondary ejecta of Imbrium and could create the same morphologies and 

regional relations as a flow of purely primary material. Morrison and 

Oberbeck (1975) calculated that a deposit dislodged at the Apollo 16 

landing site by Imbrium projectiles would contain only 13 to 18 percent of Imbrium-primary ejecta. At the landing site itself, however, 

whatever secondaries were formed are buried by the Cayley and 

Descartes deposits. Therefore, these deposits either are primary (like 

those at the Apollo 14 site) or originated from secondary craters 

northwest of the Apollo 16 site. Much more Imbrium material than 13 

to 18 percent would be present if the deposit originated to the northwest (fig. 10.26).

FIGURE 10.35. Strip of eastern nearside terra, 900 km long, from Montes Haemus (H) to 

crater Descartes (D; 48 km, 12° S., 16° E.). Montes Haemus is 1,000 km from center of 

Imbrium basin, and Descartes 1,800 km. North half of photographed area is distinctly 

inundated by Imbrium-basin deposits (If, Fra Mauro Formation); formation has accumulated 

against southeast wall of Julius Caesar GO The Cayley Formation in its type area (lea; 

compare fig. 10.38) is apparently gradational with the Fra Mauro Formation; other patches, 

as at Apollo 16 landing site (arrow), seem more isolated, yet are situated near Imbrium-radial 

lineations, as at Dollond B (DB; 37 km) and Dollond C (DC; 32 km). Hilly and furrowed 

material of the Descartes Mountains (around and south of Apollo 16 landing site to bottom of 

photograph), though also seemingly isolated, forms part of Imbrium-radial system north of 

Apollo 16 site and is similar to hilly material farther north (Ih). Mosaic of Orbiter 4 frames 

H-89 (below) and H-90 (above).
10. LOWER IMBRIAN SERIES 219

Any of the debris-flow mechanisms could emplace the Cayley and 

the Descartes Formations either simultaneously or sequentially (fig. 

10.39). Simultaneous emplacement (fig. 10.39A) seems to be consistent with the Orientale relations because the Orientale plains and 

dunelike deposits appear to be gradational (figs. 4.4F-H). If simultaneously emplaced, the differing morphologies might arise either 

because (1) a higher melt content (chap. 9) caused the Cayley deposits 

to pool in depressions (Hodges and Muehlberger, 1981; Ulrich and 

Reed, 1981) or because (2) the Descartes material, but not the Cayley 

material, reached and was crumpled into dunelike topography by the 

Kant Plateau. Alternatively, an embankment of ejecta may have been 

emplaced by the Nectaris basin, followed at the beginning of the 

Imbrian Period by secondary impacts forming the Descartes furrows 

and a debris flow forming the Cayley (fig. 10.39B). The measured 

absolute ages and the apparent lithologic differences between the 

Cayley and the Descartes (see chap. 9; table 9.1) are consistent with 

this sequence.

Among the absolute ages of the sampled impact-melt fragments 

are some that are distinctly younger than the 3.92 ± 0.03-aeon age 

tentatively identified here with the Nectaris impact (table 9.1). A few 

of these ages might indicate an Imbrium-basin origin of materials in 

the units or, at least, in the melt-rich Cayley. Although the methodologic problems are severe (see chap. 9), these ages are at least 

consistent with an Imbrian age of the younger impact-melt rocks.

"Local" origin

In the Imbrium debris-flow models, strata of the Cayley and 

Descartes Formations would possess considerable lateral continuity, 

though also considerable lithologic diversity. In "local origin" models, 

in contrast, the units formed piecemeal and were derived from various relatively small deposits quite diverse in lithology, absolute age, 

and composition. Sample diversity is commonly cited as requiring 

such diverse "local" sources.

Local origin of the Cayley and of plains in general is commonly 

envisioned as dislodgement of debris from crater walls and redeposition in adjacent depressions (Morrison and Oberbeck, 1975; Oberbeck 

and others, 1974,1975,1977; Oberbeck, 1975; Oberbeck and Morrison, 

1976). An early suggestion of this type was that the plains consist of 

accumulated secondary ejecta of craters (Oberbeck and others, 1974). 

The concentration of plains around Imbrium in many additional 

localities indicates, however, that if secondary impacts emplaced the 

Cayley, they were from Imbrium.

FIGURE 10.36. Plains materials superposed on Imbrium-secondary craters, probably consisting of debris dislodged and ejected by secondary impacts (Oberbeck and others, 1974) 

superposed on north (near-Imbrium) wall of crater Albategnius (136 km). Visibility of 

secondary chains buried by plains on floor is enhanced by Sun elevation, only 5° above 

horizontal (from right). Neukum (1977) suggested that plains postdate Imbrium and Orientale basins and are volcanic. Apollo 16 frame M 449.

One of the most widely favored hypotheses has been that the 

deposits originated as material of pre-Imbrian "local" craters (Head, 

1974d; Oberbeck and others, 1974,1975). Head (1974d) suggested that 

an "unnamed crater B" dominates the landing site's geology and 

petrology (fig. 10.40). The Cayley is supposed to consist of melt rock 

from the fallback on this crater's floor, and melt-poor rocks are supposed to have come from the less highly shocked rim materials of 

unnamed crater B and other old craters. Head (1974d) thought that 

unnamed crater B transects the Stone Mountain facies (of Nectarisbasin origin) and is transected by the Smoky Mountain facies. 

Unnamed crater B, however, is a vague circular feature or coalescence 

of subcircular features that lacks the continuity, terraces, and other 

hallmarks of a crater young enough to affect the surface geology; it 

interrupts no furrows. If it exists at all, it is too deeply buried by both 

Descartes facies and the Cayley to have furnished much material to 

the bedrock-breccia deposits of either unit. Furthermore, no other 

craters older than the Cayley and the Descartes Formations are 

sufficiently near to have furnished much material to either bedrock 

deposit (pis. 6, 7). If "local" material constitutes these formations, it is 

Nectaris-basin primary ejecta.

Although a "local" origin in the sense generally meant seems to 

be refuted by the photogeologic relations, several processes that 

might be defined as local probably did contribute to the observed 

diversity of the samples. First, secondary impacts of Imbrium ejecta 

dislodged material, mostly Nectaris-basin primary ejecta but possibly including some crater material, that lay northwest of the landing 

site. Second, the reworked Nectaris-basin ejecta itself consists of 

materials that were derived from diverse small-scale deposits of 

craters and, possibly, of volcanic flows that were in the Nectaris target 

area before the Nectaris impact. Any later impact would have incorporated, shocked, melted, and reworked such earlier impact and volcanic materials. Most of the chemical, chronologic, and lithologic 

diversity of the Apollo 16 samples probably results from this preNectarian diversity, which is a feature of all impact deposits.

The preceding discussion pertains to the bedrock deposits. All 

hypotheses acknowledge that several local processes contributed to

10-

5

O

J

X.

LU

cc

a

LU 

DC

DC 

0

10- 2

io-3

0

10

I I I I I I I I

Light plains 

Nearside 

T Apollo 16 landing site

 Ptolemaeus

 Albategnius 

A Meton 

Farside 

O Medeleev 

O Van de Graaff

10° 10 1

CRATER DIAMETER, IN KILOMETERS

102

FIGURE 10.37. Size-frequency distributions of craters superposed on terra plains of several 

ages (Neukum and others, 1975b; compare fig. 11.2). Curves represent, from oldest to 

youngest: plains in Mendeleev (compare fig. 9.29); Imbrium basin and plains at Apollo 16 

landing site, indistinguishable; plains in Ptolemaeus; Orientale basin, similar to Ptolemaeus 

plains; plains in Albategnius, considerably younger than Orientale basin; and mare materials 

in Van de Graaff.
220 THE GEOLOGIC HISTORY OF THE MOON

the diversity of the overlying relatively thin regolith at the landing 

site:

1. Large craters probably pierced the Cayley Formation and introduced some of the underlying material into the regolith. Seismic 

data indicate that the Cayley is about 200 m thick (Hodges and 

others, 1973; Cooper and others, 1974). Impacts on the order of 

North Ray Crater in size (1 km) could penetrate this thickness. 

Many subdued craters of this size are visible and probably "contaminated" the regolith during the Imbrian Period (fig. 9.9). The 

introduced material probably was mostly Nectaris ejecta but may 

also have included some crater materials.

2. The regolith developed on each formation contains some fragments derived from the other formation by post-Imbrium impacts. 

Origin as Orientale-basin ejecta seemed to be supported by three 

post-Imbrium Ar-Ar ages (table 9.1) and a precise Rb-Sr age of 

3.76 ±0.01 aeons (Papanastassiou and Wasserburg, 1972b) for 

Cayley sample 68415 (Chao and others, 1975). An alternative origin 

more in keeping with the regional relations is that this rock 

(samples 68415, 68416) was derived from impact melt of an 

Imbrian crater that formed nearby a few tens of millions of years 

after the Cayley and Descartes Formations. This origin is supported by the fact that its composition is similar to that of the 

station 11 regolith, which accumulated on the Descartes Formation 

(P.D. Spudis, oral commun., 1983).

3. Like all lunar regoliths, the Apollo 16 regolith undoubtedly contains some material derived by small post-Imbrium impacts from 

other nearby units. For example, fragments of mare basalt were 

probably introduced into Theophilus ejecta (see chap. 11; Delano, 

1975).

Conclusions

1. The relation of surface morphologies to the Imbrium basin suggests that the Cayley Formation and at least part of the Descartes 

Formation are also genetically related to that basin (fig. 10.39).

2. I favor origin of the Cayley by some sort of Imbrium debris-flow 

emplacement, and origin of the Descartes by modification of Nectaris-basin deposits by the same debris flow or by Imbrium-secondary impacts (fig. 10.39).

3. Secondary impacts are more likely to have contributed to the 

deposits than at the Apollo 14 landing site because secondary 

craters are more evident and primary-ejecta features less evident 

than at that site. However, I consider neither the relative role of 

primary and secondary ejecta of the Imbrium basin nor the origin 

of the overall mountainous topography of the Descartes Formation 

to be established.

4. The absolute ages of certain samples support a Nectaris-basin 

origin for much of the Descartes material, and other absolute ages 

suggest addition of some Imbrium material (table 9.1). However, 

the relative amounts of Imbrium and Nectaris ejecta in the 

Descartes and Cayley Formations are unknown.

5. Pre-Imbrian "local" craters contributed little material to the site 

except as parts of the Nectaris ejecta, as minor parts of the 

Imbrium debris flow, and as regolith fragments.

6. Post-Imbrium craters melted some local materials.

7. Although this list of unsolved problems is unsatisfyingly long, at 

least the dominant role of basin deposits over "local" crater deposits appears to be well established.

FIGURE 10.38. Type area of the Cayley Formation (Morris and Wilhelms, 1967).

A. Regional view. Type area (3.5-4.5° N., 15.5-16.25° E.)ispartly within and partly below rectangle outlining B. Crater Cayley (C; 14kmX moderately fresh, with lightly cratered ejecta, 

is probably Eratosthenian. Dionysius (D; 18 km), very fresh, is Copernican. More highly degraded and shallower craters above, below, and left of crater Cayley are probably Imbrian. 

Intercrater septum of pair near right edge of photograph indicates simultaneous (primary or secondary) impact (see chap. 3). Part of Rima Ariadaeus is at top (see chap. 6). Lineated, preImbrian craters truncated by left edge of photograph are also truncated by right edge of figure 10.16. Orbiter 4 frame H-90.
10. LOWER IMBRIAN SERIES 221

OTHER BASIN AND CRATER 

MATERIALS

Recognition criteria

The units discussed heretofore in this chapter are parts of extensive deposits of the Imbrium basin, and the Orientale deposits 

described in chapter 4 constitute another widespread group of contemporaneous units. Like all other lunar time-stratigraphic units, 

the Lower Imbrian Series also contains randomly scattered primarycrater materials (table 10.2).

The regional units establish Imbrian ages of the scattered units 

over a large area of the Moon. Imbrian craters are sandwiched 

between deposits of the Imbrium basin and Upper Imbrian mare 

materials. Examples inside the Imbrium basin include Archimedes, 

Plato, the Iridum crater, and Cassini. Outlying craters, such as Lansberg, whose ejecta is superposed on Imbrium deposits but embayed 

by mare materials, are also Imbrian (fig. 10.41).

Furthermore, stratigraphic relations divide some craters into 

Lower and Upper Imbrian. Craters superposed on Imbrium deposits 

and overlain by Orientale deposits are Lower Imbrian (fig. 10.2). 

Other fresh-appearing craters buried by Orientale ejecta or secondaries are also likely to be Lower Imbrian (fig. 10.42), as are craters 

overlain by relatively old Upper Imbrian crater materials (fig. 10.43).

Counts of craters of the large sizes required to date pre-Nectarian 

and Nectarian surfaces (see chap. 7) are statistically valid for deposits

TABLE 10.2. Possible Lower Imbrian craters

[Cross rules divide diameter ranges mapped differently in plate 8: smaller than 30 km, unmapped; 30 to 59 km, interiors mapped; 60 km and larger, exterior deposits mapped where 

exposed. UI?, possibly Upper Imbrian]

* «« Di (Ser

Ukert 

Inghirami A 

La Condamine 

22

23

34 

37 

46 

46 

48 

57

64 

93 

97 

119 

162 

169 

177

Center 

(lat) (long)

56° N. 

8° N.

45° S. 

53° N. 

17° S. 

50° N. 

28° S. 

40° N.

21° N. 

27° N. 

18° S. 

11° S. 

56° N. 

10° S. 

25° S.

7° E. 

2° E.

65° VJ. 

28° W. 

67° VJ. 

27° VJ. 

28° VJ. 

5° E.

46° E. 

98° E. 

2° VJ. 

42° VJ. 

105° E. 

162° E. 

60° E.

Fi gu re

10.15 

10.1

10.42 

3.8, 10.43 

4.4H, 7.R 

3.8, 10.43 

6.9 

1.6, 10.12

9.11 

1.2, 4.7, 8.13, 9.28 

1.8, 7.7 

6.3, 12. 3C 

4.2B, 9.4, 9.23, 10.6 

1.4, 3.1, 8.7 

4.2C, 9.5

Remarks

...

Covered by Orientale. 

Covered by Iridum.

Covered by Iridum. 

UI?

Typical ; mare fill . 

I/I? 

UI?, half-buried. 

UI?, interior ring.

Floor and rim faulted.

of two Imbrian basins and marginally definitive for a third (fig. 10.4). I 

determined cumulative frequencies of 28 and 22 craters larger than 

20 km in diameter per million square kilometers on Imbrium-basin 

and Orientale-basin deposits, respectively. The large statistical samples suggest that these are good values for the lowermost Imbrian and

B. Area outlined in A. The Cayley Formation is more densely cratered than a mare surface and is saturated by small subdued craters, as expected on a Lower Imbrian surface. Many secondary 

clusters are also superposed large in upper right, small and linear elsewhere (probable source, Dionysius). Relief is enhanced by low Sun illumination, 11 ° above horizon (from right). AB, 

crater Anadaeus B (8 km), whose flat floor is visible in A but is obscured here; Imbrian age is confirmed by crater density similar to that on the Cayley Formation. Orbiter 2 frame M 61.
222 THE GEOLOGIC HISTORY OF THE MOON

the boundary between the two Imbrian series, respectively. Sparse 

crater counts indicate that the Schrbdinger basin is also Imbrian (fig. 

10.4; density, 20 such craters). Schrbdinger is the topographically 

freshest lunar basin of its size and thus appears to be relatively young 

(fig. 10.44). It had been considered Nectarian because of its moderately large density of small craters, although no definite Nectarian 

craters are known to be superposed (Wilhelms and others, 1979). If it 

is Imbrian, superposition of some probable Orientale secondaries (fig. 

10.44) restricts its age to Early Imbrian.

Otherwise, counts of craters smaller than 20 km in diameter are 

required to date Imbrian surfaces not in contact with definite Imbrian 

deposits. Figure 10.4 plots the counts by Neukum and others (1975a) 

of craters as small as 1.4 km in diameter on Montes Apenninus, which 

fall on the projection of the curve determined here for craters larger 

than 20 km in diameter. Crater frequencies for the plains at the 

Apollo 16 landing site nearly coincide with those for the Apennines. 

The counts by Neukum and others (1975a) and myself indicate that 

the Orientale and Imbrium curves are approximately parallel for 

both small and large superposed craters.

Craters that are thought to be too fresh to be Nectarian and that 

are overlain by relatively old mare materials are likely to be Early 

Imbrian for example, Lomonosov (fig. 9.28). Stratigraphically isolated Early Imbrian craters are difficult to distinguish from Nectarian and Late Imbrian craters, and so the compilation here (pi. 8; 

table 10.2) is based mostly on qualitative assessments of morphology 

and some sparse size-frequency counts.

The distinction between small craters of the Lower and Upper 

Imbrian Series is still less clear. Craters smaller than about 20 km in 

diameter in the terrae are hard to date by morphology (Wilhelms and 

others, 1978). Debris from the walls has filled Imbrian craters smaller 

than about 8 km in diameter more deeply than their younger counterparts, and commonly has created shallow bowl-shaped or flat-floored 

interiors from originally steeper more nearly conical profiles. 

However, detection of this distinction requires good photographic 

resolution and Sun illumination (figs. 10.38, 10.41). Fresh Imbrianage primary craters may be indistinguishable from sharp Orientaleor Imbrium-basin secondary craters of similar sizes. Small flatfloored or otherwise subdued Imbrian primary craters may resemble 

some Nectarian primaries or partly buried secondaries of one of the 

Imbrian basins. Imbrian craters disappear at diameters of about 1 km 

on the Fra Mauro Formation (fig. 7.14; Trask, 1971).

FIGURE 10.39. Diagrammatic geologic cross sections drawn across region of Apollo 16

landing site, illustrating alternative interpretations.

A. The Descartes and Cayley Formations were both emplaced in an Imbrium debris surge.

B. The Descartes was emplaced as Nectaris ejecta and furrowed by impacts of Imbrium 

ejecta. The Cayley was emplaced in an Imbrium debris surge. The Cayley would consist 

mostly of Imbrium ejecta if the surge originated near Imbrium, and mostly of Nectaris 

ejecta if the surge originated nearer the Apollo 16 landing site.

Frequency

Despite the difficulties, Lower Imbrian craters are here distinguished from Nectarian and Upper Imbrian craters as well as 

possible and are mapped in plate 8. In a farside area of 14.2 x 106 km2 

that excludes the Orientale deposits which overlie all other materials 

of the series, 73 unburied craters at least 30 km in diameter are 

interpreted as Lower Imbrian. For this density (5.1 craters/106 km2), 

195 craters of this size were formed over the whole Moon.

OTHER TERRA MATERIALS

The discussions of the Orientale and Imbrium basins have 

offered basin-impact alternatives (Howard and others, 1974) to the 

volcanic origins previously considered or assumed for many terra 

landforms. However, volcanic origins for some terra plains (Neukum, 

1977), as well as for a few isolated domelike structures, are still being 

entertained.

The "domes" have a volcanolike morphology in the form of summit or flank pits and unusually red spectra (fig. 10.45; Wilhelms and 

McCauley, 1971; McCauley, 1973; Scott and Eggleton, 1973; Malin, 

1974; Head and McCord, 1978; Basaltic Volcanism Study Project, 1981, 

p. 762-763). The most densely concentrated are the Mairan and 

Gruithuisen domes near the Iridum crater, between Mare Imbrium 

and northern Oceanus Procellarum or Sinus Roris (Scott and Eggleton, 1973). They may represent rare differentiated volcanic materials akin to KREEP-rich material or certain still more highly

DOLLOND C

UNNAMED CRATER A

DOLLOND B

UNNAMED 

CRATER D KANT

CRATER C

^-DESCARTES

KM

FIGURE 10.40. Vicinity of Apollo 16 landing site. 

A. Telescopic photograph.

B. "Unnamed" craters believed by Head (1974d) to constitute area of A and to have 

contributed materials to Apollo 16 samples.
10. LOWER IMBRIAN SERIES 223

differentiated "granitic" or "quartz-monzodioritic" compositions that 

appear in small amounts in lunar samples (Ryder and others, 1975a; 

Ryder, 1976). The craters atop some of these domes argue for volcanic 

origins (fig. 10.45A). Experience with other seemingly volcanic landforms, however, recommends a search for impact origins. Some 

"domes" may result from blanketing of Imbrium massifs and Iridumsecondary craters by Iridum ejecta (P.D. Spudis, oral commun., 1980).

MARE MATERIALS

Whether any Lower Imbrian mare materials are exposed is 

uncertain. None could remain exposed within reach of the Orientale 

ejecta. Judging from the mapped extent of the Orientale deposits and 

secondary craters on the west-limb and farside terrae, the nearside 

maria probably would have been heavily mantled or cratered if their

surface units predated Orientale. Orientale and other large basins 

might even cover the whole Moon with debris (Moore and others, 1974; 

Chao and others, 1975). However, some outlying points maybe spared 

a cover because outer impact ejecta is concentrated in raylike stringers (pi. 3; fig. 4.6).

Certain heavily cratered maria on the east limb may be of preOrientale, Early Imbrian age. One example is the patchily distributed planar material in the Australe basin (fig. 10.46). In albedo, 

these plains are intermediate between typical mare and typical circumbasin terra plains, and the superposed crater densities are similar to those of Orientale (J.M. Diaz and A.B. Watkins, written 

commun., 1976). These marelike materials are older than the Imbrian 

craters Humboldt and Jenner (Wilhelms and El-Baz, 1977). Darkhaloed craters are evident on the marelike plains (Schultz and 

Spudis, 1979). A thin dusting of mare basalt by Orientale ejecta may 

explain the plains' albedo. Other densely cratered materials with

FIGURE 10.41. Imbrian craters on periphery of Imbrium basin. Lansberg (L; 39 km, 0.3° S., 26.5° W.) is superposed on hummocky Imbrium ejecta (h) and embayed by mare materials (arrows). 

Interior and periphery of Lansberg C (LC; 17 km) are flooded more deeply. Rayed, probably Copernican, but somewhat degraded crater Lansberg B (LB; 9 km) is superposed on the mare. 

Craters in 1-km size range are probably Imbrian (I), Eratostheman (E), and Copernican (C). Apollo 12 frame H-8095.
224 THE GEOLOGIC HISTORY OF THE MOON

intermediate albedos lie in Mare Marginis (fig. 4.7) and at the distal 

margin of Mare Frigoris (Lucchitta, 1978).

Nevertheless, none of these criteria prove an Early Imbrian age. 

Even the obviously Late Imbrian materials near Orientale (fig. 10.3) 

are quite heavily cratered. All exposed mare materials may, therefore, 

be Upper Imbrian or younger and are so mapped here (pis. 9-11). 

They can be locally subdivided in considerable detail, as shown in the 

next two chapters.

CHRONOLOGY

Radiometric ages of Imbrium-basin materials are clustered 

within a short interval. The oldest are those of an older cluster of ages 

at the Apollo 14 landing site (3.87-3.96 aeons), but these ages pertain 

to pre-Imbrian rocks. The Apollo 15 KREEP basalt, 3.85 ± 0.04 aeons 

old, probably was derived from a deposit (Apennine Bench Formation) 

that is superposed on other basin materials. This basalt may be the 

best dated of the Apollo 14 and 15 terra materials (G. W. Lugmair, oral 

commun., 1982); it is contemporaneous with or younger than the 

Imbrium basin, depending on its origin. Black-and-white breccia 

believed to have been collected from the Imbrium massifs has been 

dated at 3.86 ± 0.04 aeons. A younger cluster of ages from the Fra

FIGURE 10.42. Terrain centered 600 km southeast of Orientale-basin rim (Monies Cordillera). Largest craler is Inghirami (91 km, 48° S., 69° W.); olher probable primary craler 

older lhan Orienlale deposils is Inghirami A (I; 34 km). Relalively young ages for bolh 

craters are suggested by moderately sharp lexlural elemenls, especially smaller Inghirami A. 

If ihey are Imbrian and nol Neclarian, ihese craters musl be Early Imbrian because ihey are 

buried by Orienlale deposils. Texlured Hevelius Formalion and plains inlergrade. 

Arrowheads indicate embaymenl of Orienlale-secondary craters by plains; most or all of 

clustered craters in lower righl corner are Orienlale secondaries. Orbiler 4 frame H-l 72.

Mauro Formation at the Apollo 14 landing site falls mostly between 

3.82 and 3.84 aeons. The time of the Imbrium impact seems to be well 

constrained at from 3.82 to 3.87 aeons ago; the average and wellrepresented date of 3.85 aeons ago is tentatively adopted here.

One group of basalt samples from the Apollo 11 landing site 

yielded ages as old as 3.84 ±0.03 aeons (table 11.1; Guggisberg and 

others, 1979). The source flows of these samples are not exposed (see 

chap. 11); they are probably Lower Imbrian.

No other Lower Imbrian units have been dated isotopically. The 

interpretation of ages reported for the Cayley Formation at the Apollo 

16 landing site is highly uncertain, except that the Cayley materials 

are no more than 3.92 aeons old. Materials of the Descartes Formation, which may have been emplaced 3.92 ±0.03 aeons ago, were 

modified and probably augmented by Imbrium-basin materials 3.85 

aeons ago.

The upper age boundary of the series is limited by the ages of the 

oldest nearly exposed mare basalt. This material was recovered from 

the Apollo 17 landing site and is described in the next chapter; it is 

3.72 aeons old. Therefore, the Orientale impact that ended the Early 

Imbrian Epoch occurred some time between 3.85 and 3.72 aeons ago. 

Relative crater frequencies (figs. 8.16,10.4) suggest that it occurred 

about 3.8 aeons ago. Thus, the Early Imbrian Epoch spanned about 

50 million years.

FIGURE 10.43. Terra between Mare Frigoris (F) and Mare Imbrium, underlain by 

deposits of Imbrium basin. Lower Imbrian craters La Condamine (L; 37 km, 53° N., 28° 

W.) and Maupertuis (M; 46 km) are swamped by deposits of Iridum crater (Ulrich, 1969), 

which grade outward to conspicuous secondary craters (sc) and to light plains deposits (white 

p). Secondary craters of crater Plato (black P), centered outside right edge of picture, are 

superposed on Iridum-secondary craters. Iridum crater (centered outside lower left corner of 

photograph) is slightly younger than Orientale (fig. 10.4) but older than Plato, and so is 

probably early Late Imbrian. The Straight Range (Monies Recti, R) is part of Imbrium ring 

overlain by Iridum deposits. Orbiter 4 frame H 139.
10. LOWER IMBRIAN SERIES 225

FIGURE 10.44. Schrodinger basin (320 km). Arrow indicates probable secondary craters of Orientale, centered 2,150 km in direction of arrow; overlap and orientation of cluster are 

consistent with Orientale origin (Wilhelms and others, 1979). Large chain to left of arrow is secondary to Schrodinger. (White and black-and-white spots are blemishes.) Orbiter 5 

frame H 21.

FIGURE 10.45. Domical terra landforms of possible volcanic origin. 

A. Gruithuisen domes (36° N., 40° W.). Orbiter 5 frame M-182.

B. Hansteen alpha (12° S., 50° W.). B, crater Billy; H, crater Hansteen. Orbiter 4 

frame H-149.
226 THE GEOLOGIC HISTORY OF THE MOON

FIGURE 10.46. Stratigraphic relations of mare and crater materials in Mare Australe. 

A. Deposits of Jenner 0; 72 km, 42° S., 96° W.) and Humboldt (H; 207 km, 27° S., 81 ° 

E.), both Upper Imbrian, are superposed on mare basalt (white-on-black arrowheads), 

whereas a younger mare unit encroaches on Humboldt secondaries (white arrowheads). 

C, pre-Nectarian crater Curie (1 39 km, 23° S., 92° E.); L, pre-Nectarian crater Lyot. 

Obiter 4 frame M-9.

B. Two mare units in Lyot (141 km, 50° S., 84° W.). Younger unit encroaches on older (white arrowheads). Secondary craters of crater Humboldt are superposed on older unit, including at 

shaft end of black-and-white arrow, which points to Humboldt, 10° W. of N. Older unit may be Lower Imbrian; younger unit is Upper Imbrian. Orbiter 4 frame H-9.
11. UPPER IMBRIAN SERIES

»

*

f
FIGURE 11.1 (OVERLEAF). Region centered on Apollo 11 landing site in Mare Tranquillitatis (arrow). Large crater at bottom is Theophilus (100 km). Rings of Procellarum basin pass under 

mare-ridge feature Lament (ring and radial pattern above arrow) and through rugged terra at right. Telescopic photograph by the Pic du Midi Observatory, France.
11. UPPER IMBRIAN SERIES

CONTENTS

Page

Introduction ...................................................... 229

Definition ........................................................ 229

Crater materials .................................................. 231

Recognition criteria ............................................ 231

Frequency .................................................... 232

Mare materials general stratigraphy and distribution .............. 232

Mare Tranquillitatis ............................................... 235

General stratigraphy ........................................... 235

Apollo 11 samples ............................................. 235

Summary and interpretation .................................... 238

Mare Nectaris ................................................... 238

General stratigraphy ........................................... 238

Apollo 16 samples ............................................. 238

Summary and interpretation .................................... 238

Mare Serenitatis .................................................. 238

General stratigraphy ........................................... 238

Apollo 17 samples ............................................. 239

Summary and interpretation .................................... 241

Mare Fecunditatis ................................................ 241

General stratigraphy ........................................... 241

Luna 16 samples ............................................... 241

Summary and interpretation .................................... 241

Mare Crisium .................................................... 241

General stratigraphy ........................................... 241

Luna 24 samples .............................................. 241

Summary and interpretation .................................... 242

Mare Imbrium .................................................... 243

General stratigraphy ........................................... 243

Apollo 15 samples Palus Putredinis ............................ 243

Summary and interpretation .................................... 244

Other maria inside the Procellarum basin ........................... 244

Near-limb maria outside the Procellarum basin ..................... 245

Farside maria .................................................... 245

Chronology ...................................................... 245

INTRODUCTION

The Upper Imbrian Series contains extensive mare materials, 

subordinate crater materials, and no basins at all (pi. 9; fig. 11.1). This 

geologic style contrasts sharply with the exposed record of the earlier, 

impact-dominated era (see chaps. 8-10; pis. 6-8), which ended with 

the Orientale impact. Only on the farside and only if exterior deposits 

are considered (as shown for craters larger than 60 km in diameter on 

pi. 9) do Upper Imbrian crater deposits exceed Upper Imbrian maria 

in area.

Two-thirds of the nearside mare surface is Upper Imbrian. 

Because of this extent and because the Imbrium impact struck the 

nearside, beginning the Imbrian Period, the Imbrian System dominates the nearside (pis. 8, 9). Absence of an Imbrian cover leaves the 

pre-Nectarian and Nectarian impact materials, which doubtless 

formed randomly on the whole Moon, as the dominant deposits on the 

farside (pis. 6, 7).

Upper Imbrian mare units furnished more samples than any 

other class of lunar geologic unit, because of their accessibility to 

landings and their prominence on the nearside. Samples of mare 

basalt were returned from all landing sites (see chap. 5; table 1.2), and

seven of these sites furnished data on absolute ages and composition 

that are useful in regional studies like the present one. Only one of 

these seven sampling missions, Apollo 12, returned basalt that is here 

regarded as younger than Imbrian (see chap. 12).

The emphasis in this chapter corresponds to the relative importance of the two main types of lunar materials in the Upper Imbrian 

Series. Craters are described briefly in the first section after the 

series definition. The rest of the chapter describes the general stratigraphy of each major mare or mare cluster and describes in detail the 

absolute ages and stratigraphy for those maria from which basaltic 

samples were collected. I show that the volume, sequence, and composition of basaltic eruptions were fundamentally predetermined in 

the foregoing era, when basin impacts thinned the crust, uplifted the 

mantle, and weakened the lithosphere.

DEFINITION

The bottom of the Upper Imbrian Series is the top of the Orientale-basin deposits (see chap. 10). Frequencies of 2.2 x 10 ~ 2 craters at 

least 1 km in diameter per square kilometer and DL values of 550 m,

229
230 THE GEOLOGIC HISTORY OF THE MOON

therefore, are approximate maximums for the Upper Imbrian Series 

(figs. 10.37,11.2; tables 7.3,11.1). To my knowledge, the largest frequencies actually determined on Upper Imbrian materials are about 

2xlO~ 2 craters of this size per square kilometer, for old, probably 

unsampled units near the Apollo 11 landing site (Neukum and others, 

1975b; compiled from data of Greeley and Gault, 1970). The large error 

bar for this determination (table 11.1) overlaps the crater frequency for 

the older Orientale basin and demonstrates the imprecision of such 

counts for small mare units. The largest DL value measured on a 

mare unit is 390 m, for a spot in Mare Nectaris (Boyce and others, 

1975).

No stratigraphic unit comparable to the Orientale deposits 

defines the top of the Upper Imbrian Series and the Imbrian System, 

because the individual mare and crater units that formed after Orientale do not cover a sufficient area to be globally useful as stratigraphic-datum horizons. Therefore, the boundary is defined on the 

basis of DL values.

o

2.

LU

r> 

O 

LU

cr

LU

cr 

o

LU

o

10,-1

10"'

icr

10

-5

T Monies Apenninus (total)

 Monies Apenninus (15-413) 

D Mare Orientale

 Mare Serenitatis

o Apollo 16 LS Light Plains

A Apollo 15 LS

A Mendeleev (floor)

 Mare Crisium

- A

i i i 1 i 111

10rl

100 101

CRATER DIAMETER D (km)

FIGURE 11.2. Crater frequencies of units of Upper Imbrian and other ages.

A. Distinct curves for three time-stratigraphic units: (1) Nectarian System plains on floor 

of Mendeleev; (2) Lower Imbrian Series Montes Apenninus flank (two curves), light 

plains at Apollo 16 landing site (LS), and "Mare Orientale," actually Orientale basin 

(compare fig. 10.37); (3) Upper Imbrian Series eastern Mare Crisium, Palus 

Putredinis near Apollo 15 landing site (both sides of Rima Hadley), and two-thirds of 

central Mare Serenitatis. Upper Imbrian units are near upper boundary of the series. 

From Neukum and others (I975a, fig. 6).

TABLE 11.1. Properties of the Upper Imbrian Series and adjacent time-stratigraphic units

[References: B75, Boyce and others (1975); B76, Boyce (1976); M80b, Moore and others (l<580b); N75a, 

Neukum and others (1975a); N75b, Neukum and others (1975b); SL72, Soderblom and Lebofsky (1972).

Crater frequencies: Cumulative, determined by normalization, using a calibration curve and commonly based 

on craters smaller than 1 km in diameter (see fig. 7.10G; Neukum and others, 1975b).

Ages: Radiometric ages obtained from samples collected from the listed rock-stratigraphic units (see tables 

10.1, 11.3, 12.4)]

Rock-stratigraphic Crater frequency 

Time- unit (larger than

unit Mare Formation P er S 1" ar* 

group kilometer)

Eratosthenian Apollo 12 mare 2.4^'^xlO" 1 N75b 

System.

D L

(m)

210+20 

210+30 

215+45

Reference , " *

B76 

SL72 

SL72

3.16

Upper Imbrian 

Series.

Apollo 15 mare 2.6+^xlO' 3 N75a, N75b 270+15 B76 

Eastern Mare Approx 2.6xlO" 3 N75a 

Crisium.

Central Mare 

Serenitatis.

3.9x10" N75a

Luna 16 mare - - 

Apollo 11 mare 3 - 4li''

225-320 B75

270+30 B76

320+40 

330+40 

335+55

SL72 

SL72 

B76

3.40

3.57

Lower Imbrian 

Series.

2 Apollo 17 mare 

the Apollo 11 

landing site.

 Orientale basin 2.2xlO' 2

N75b

N75a, N75b

360+30 B76 

365+30 M80b

550+50 Chapter 10

3.72

ZT 10 i

-1

o-z.

LU

LU

cr

cr 

LU

cr 

o

io 2

TO'3

O

D

B

Apollo 12

Apollo 14 (all craters)

Apollo 14 (preexisting

craters eliminated) 

Apollo 15 

Apollo 16 

Apollo 17

= 10- 11 D- 2

10,-1 10U 101 

CRATER DIAMETER D (km)

B. Upper Imbrian Apollo 15 mare in comparison with older and younger units. Apollo 17 

mare (sparse data) is also Upper Imbrian. Offset in Apollo 12 curve shows that data for 

two units are combined an older Upper Imbrian unit not sampled and an Eratosthenian unit that probably was sampled (see chap. 12); Eratosthenian unit is barely younger 

than Apollo 15 mare. Lower Imbrian units (see chap. 10) are Apollo 16 Cayley plains 

and Fra Mauro Formation in area of Apollo 14 landing site; one Apollo 14 count ("all 

craters") is hybrid of superposed and buried craters. Curve NE is equilibrium (steady 

state) curve of Trask (1966); Apollo 16 plot shows rollover near steady state at about 

800-m crater diameter (table 7.4); craters too small to count would show rollover for 

younger units. From Neukum and others (1975b, fig. 7).
11. UPPER IMBRIAN SERIES 231

jc^V-"-'^ " **"""" *~ -^£fr *** «" "* ^ffPii*^

- -

 w^"5 - f, m/

 < ^^?^*^ ^ik Jf- "~ S*

^rjb^j - * 3 R R Si*';% 

- r ",- *:S' * * « - __:

" "*r *£ %&f, " '£ : ,:,: .r. -.^-«E ~"?*~^ ^ - *%? t

&&? i - **£ * r. " " -*f.. £^s -- L -..,* "'- . i-*j»:; _ -v-!rrr- » _..... = ;- * - ^-

Boyce and others (1975) determined that the mare units between 

Eratosthenes and Archimedes (figs. 1.6,2.1,10.1,11.3) haveD^ values 

from 190 to 340 m, including error ranges; midpoints range from 220 

to 290 m. Wilhelms (1980) observed that units here and in the rest of 

southern Mare Imbrium which have DL values less than 230 m match 

Eratosthenian units as traditionally mapped, and that units withD^ 

values greater than 250 m match Imbrian units. Therefore, he proposed these values as working criteria for the Imbrian-Eratosthenian 

boundary. Units withDL values from 230 to 250 m were to be assigned 

on the weight of other evidence. This definition, though not formally 

rigorous, has served conveniently to divide a maximum number of 

units into Imbrian and Eratosthenian age groups that harmonize 

with existing mapping. This working criterion of DL = 240 ± 10 m for 

the upper boundary of the Imbrian System and of the Upper Imbrian 

Series is used here.

The characteristic size-frequency distribution for the upper 

boundary of the system and the series is less well defined than those 

for Imbrium- and Orientale-basin deposits because of the statistical 

unreliability of counts for the small units typical of maria. The similar values determined by Neukum and others (1975a, b) for central 

Mare Serenitatis, the region of the Apollo 15 landing site, and eastern 

Mare Crisium are probably representative of units near the upper 

boundary (fig. 11.2; table 11.1). Midpoints of these determinations 

range about from 2.6xlO~ 3 to 3.9xlO~ 3 craters at least 1 km in 

diameter per square kilometer.

Thus, the total range of observed Upper Imbrian cumulative 

frequencies of such craters per square kilometer is from about 

2.2 x 10~ 2 for the Orientale basin to 2.6 x 10~ 3 for these young mare 

units (figs. 10.4,11.2; tables 7.3,11.1). Figures 7.15 and 12.8 and tables 

7.3 and 11.1 give some equivalent!^ and crater-frequency values. As 

discussed in chapters 12 and 13, the values on terra-breccia and marebasalt units are not directly comparable. Interpolations between 

values for breccia deposits and mare lavas are hazardous because the 

correction factor for the Imbrian System has not yet been determined. 

Thus, some of the units considered here to be Upper Imbrian may be 

Eratosthenian, and vice versa.

CRATER MATERIALS

Recognition criteria

Some large craters can readily be assigned to the Upper Imbrian 

Series by stratigraphic criteria (table 11.2). An example is Schliiter, 

which is superposed on Orientale deposits but slightly flooded by 

mare material (fig. 11.4).

TABLE 11.2. Representative Upper Imbrian craters

[Cross rules divide diameter ranges mapped differently in plate 9: smaller than 30 km, unmapped; 30 to 59 km, interiors mapped; 60 km and larger, exterior deposits mapped where 

exposed. E?, possibly Eratosthenian; LI?, possibly Lower Imbrian]

Diameter Center 

(km) (Tat) (long)

St. George ----- 

Bonpland D 

Ariadaeus B - 

Lansberg C ----- 

Daedalus S ----- 

Krieger ------

Vitruvius ------ 

Timaeus ----- 

Bianchini ------ 

Lansberg ------- 

Mairan ------ 

Sharp 

Stiborius ------

Krafft 

Thebit 

Archimedes ----- 

Atlas 

Piccolomini - 

SchlOter 

Posidonius 

Plato 

Tsiolkovskiy 

Humboldt 

2.3 

6 

8 

17 

20 

22

30 

33 

38 

39 

40 

40 

44 

50 

51 

53 

57

72

83 

87 

88 

89 

95 

101 

135 

180 

207 

?fin

25.9° N. 

10° S. 

5° N. 

2° S. 

7° S. 

29° N.

18° N. 

63° N. 

49° N. 0° 

42° N. 

46° N. 

34° S. 

13° N. 

17° N. 

32° S. 

22° S.

42° S.

30° N. 

47° N. 

30° S. 

6° S. 

32° N. 

52° N. 

70° S. 

20° S. 

27° S. 

aa° N

3.4° 

18° 

15° 

29° 

173° 

46°

31° 

0.5° 

34° 

27° 

43° 

40° 

32° 

72° 

73° 

25° 

4°

96°

4° 

44° 

32° 

83° 

30° 

9° 

172° 

129° 

81° 

11°

E. 

U. 

E. 

U. 

E. 

W.

E. 

W. 

U. 

U. 

U. 

W. 

E. 

U. 

W. 

E. 

U.

E.

U. 

E. 

E. 

U. 

E. 

W. 

U. 

E. 

E. 

u

Fi gu re

5.11 

6.6 

10.38 

10.41 

8.7 

11.6

9.13 

10.15 

3.8 

10.41, 13.10 

3.8 

3.8 

11.5, 11.7 

10.3, 11.8 

10.3, 11.8 

11.5 

7.7

1.3, 1.5, 10.46

1.6, 6.10, 10.1 

1.7, 9.3 

11.7 

11.4 

1.7, 5.17, 2.5D 

1.6, 7.3 

1.5, 4.2A 

1.3, 1.4, 3.24, 3.27 

1.5, 4.2E, 9.6, 10.46 

i e, ^ R in /n

Remarks

Apollo 15. 

Faulted. 

Typical . 

Flooded. 

LI? 

Asymmetric (see 

table 7.2).

LI?

LI? 

E?

Typical . 

Twin of Cardanus.

Sandwiched by mare 

units. 

Mare fill. 

Floor uplifted. 

LI? 

Typical . 

LI?; floor uplifted. 

Mare fill. 

LI?; interior ring. 

Mare fill. 

Floor uplifted.

Ma If hnr-ioH
232 THE GEOLOGIC HISTORY OF THE MOON

Other crater materials can be compared with the well-dated ones 

by means of crater densities where photographs are adequate and 

where exposed surfaces are large enough to underlie a statistically 

valid number of craters. These conditions are rarely met for craters of 

the Imbrian System. D.B. Snyder (written commun., 1980) determined that the Iridum crater (fig. 3.8) is Upper Imbrian and that 

Petavius (table 10.2) is Lower Imbrian. Most crater units, however, are 

too small or too poorly photographed to permit counts of a statistically 

significant number of superposed craters larger than the steady-state 

diameters (approx 300-800 m for the Imbrian System; see table 7.3).

Qualitative estimates of the age significance of erosional morphology are, therefore, used here more than stratigraphic or sizefrequency criteria to determine the ages of craters larger than 2 or 3 

km in diameter within the Imbrian System. Morphology reflects 

erosion by small impacts, many of whose individual craters are now 

obliterated because they are in a steady state. The radial ejecta and 

secondary-crater groups of many large craters considered to be Late 

Imbrian are conspicuous, though slightly blurred and visibly cratered; examples are Tsiolkovskiy (figs. 3.24, 3.27; Guest and Murray, 

1969) and Lindenau (fig. 11.5). Krieger, embayed by Late Imbrian mare 

materials, is a very fresh Late Imbrian crater (fig. 11.6). In general, 

however, assignments to the Lower or Upper Imbrian Series are less 

secure than those to a system. For example, such outlying craters as 

Piccolomini that are more subdued than Eratosthenian craters and 

are superposed on Imbrium-basin secondaries must be Imbrian (fig. 

11.7), but whether they are Lower or Upper Imbrian is less clear.

Frequency

Plate 9 shows a total of 162 unburied Upper Imbrian craters 

larger than 30 km in diameter within an area of 35.3 x 106 km2 that 

excludes younger deposits from the total lunar area of 38 x 106 km2 . 

Extrapolation of this density (4.6 craters per million square kilometers) suggests that 174 such craters formed over the whole Moon

FIGURE 11.4. Crater Schliiter (89 km, 6° S., 83° W.), superposed on Hevelius Formation 

at Montes Cordillera (scarp below and left of Schliiter; compare fig. 4.4). Late Imbrian age 

of Schliiter is indicated by its stratigraphic position between the Hevelius Formation and mare 

basalt; thus, the morphology and density of small superposed craters are useful for comparison with other areas (fig. 11.5). Orbiter 4 frame H 181.

during the Late Imbrian Epoch. These values are similar to those for 

the Lower Imbrian Series (5.1 such craters per million square kilometers or 195 such craters on the whole Moon; see chap. 10). As shown 

at the end of this chapter, however, the formation rates differed 

greatly because the two epochs were of unequal duration. The 

cumulative frequency of Lower and Upper Imbrian craters at least 5 

km in diameter superposed on deposits of Imbrium, Orientale, and 

the Iridum crater thus, of different maximum ages is 10 ~ 4 per 

square kilometer (figs. 7.16,10.4; Wilhelms and others, 1978).

MARE MATERIALS GENERAL 

STRATIGRAPHY AND DISTRIBUTION

Imbrian mare materials are divided here (pi. 9; figs. 11.8, 11.9) 

into three groups: old (unit 1), intermediate-age (unit 2), and young 

(unit 3). Finer subdivision is possible in well-photographed areas but 

not in all maria.

Exposures of the old units are the least extensive; they cover only 

0.24 x 106 km2 . They occupy elevated terra tracts in the outer Procellarum-basin shelf and outside this giant pre-Nectarian basin, 

places where their elevation or presence in a generally mare-poor 

area protected them fromlater inundation (figs. 10.3,10.46,11.8). They 

undoubtedly extend beneath the younger units but their extent is 

unknown. The western Procellarum remnants of the old group are 

spectrally blue, as are those on the east shelf of Mare Nubium (not 

distinguished on pi. 4). The more extensive exposures inside and near 

Oceanus Procellarum in the northeast quadrant of the nearside are 

spectrally red. Some units of the old group were discussed in chapter 

10 because they may be Lower Imbrian (fig. 10.46); however, they are 

mapped here with other Upper Imbrian units because that is their

FIGURE 11.5. Craters Lindenau (L; 53 km, 32° S., 25° E.) and Stiborius (S; 44 km, 34 

S., 32 E.), probably Late Imbrian, judging by their similarity to Schliiter (fig. 11.4). 

Rothmann (Ro; 42 km, 31° S., 28° E.), generally similar though less cratered and slightly 

sharper despite its smaller size than Lindenau, is probably Eratosthenian. Pre-Nectarian 

craters buried by Nectaris ejecta include Zagut (Z; 84 km), Rabbi Levi (RL; 81 km), and 

Riccius (Ri; 71 km); sc, Imbrium-secondary craters superposed on Nectaris ejecta. Orbiter 

4 frame H-83.
11. UPPER IMBRIAN SERIES 233

more likely age. DL values of the old group are larger than 390 m, 

values that correspond to ages of 3.75 to 3.80 aeons (Boyce, 1976). 

Size-frequency counts are unavailable.

The intermediate-age and young groups occupy much more area, 

about 1.19xl06 and 2.67 xlO6 km2, respectively (pi. 9). DL values 

range from 290 to 390 m for the intermediate group and from 240 to 

290 m for the young group. These values correpond to ages of 3.50 to 

3.75 and 3.20 to 3.50 aeons, respectively (Boyce, 1976). Crater frequencies do not accurately discriminate these groups because the 

individual units are not sufficiently extensive for statistically valid 

counting of craters. Table 11.1 lists some typical frequencies for craters 

at least 1 km in diameter (Neukum and others, 1975a, b). An approximate dividing frequency is 3 x 10 ~ 3 such craters per square kilometer, but this value is not intended to be definitive for group assignment. Both groups occur in all nearside maria, although one group 

may not occupy enough area to be shown on plate 9. The intermediateage group is abundant in Maria Fecunditatis, Nectaris, Australe, and 

other maria outside Procellarum, and on shelves and other elevated 

terrain of many maria. The young group dominates the surfaces of 

most maria. The intermediate-age and young groups are subequal in 

Mare Tranquillitatis. On the farside, photographic quality permits 

subbivision of mare units (Wilbur, 1978; Walker and El-Baz, 1982) 

only where they are covered by Lunar Orbiter H-frames (parts of 

Maria Australe and Orientale and in the crater Tsiolkovskiy) or by 

Apollo orbital photographs (most of the area north of lat 30° S. and 

west of long 180°; see pi. 2).

The original extent of the three groups is difficult to estimate. 

Assuming that Imbrian mare units underlie all exposed Eratosthenian and Copernican mare units and all the crater deposits that are 

superposed on maria (pis. 4, 10, 11), the total area of Imbrian mare 

units is 6.17 x 106 km2. If each group covers buried terrain in proportion to its exposure, the total unburied and buried extent of each is: 

old group, 6.17 x 106 km2 ; intermediate-age group, 5.18 x 106 km2 ; and 

young group, 4.01 x 106 km2. These calculations probably are reasonably accurate for the young group but less accurate for the intermediate-age group and quite uncertain for the old group, which may not 

have been extruded everywhere that maria are now exposed.

FIGURE 11.7 Crater Piccolomini (P; 88 km, 30° S., 32° £.), superposed on Rupes Altai 

(under north arrow). Piccolomini deposits are superposed on Imbrium-secondary craters (sc), 

as indicated by white-on-black arrows. Imbrium-basin rim (Monies Apenninus) is 1,800 km 

away in direction of white arrows. Craters Rothmann (R) and Stiborius (S) are also visible in 

figure 11.5. Orbiter 4 frame H-76.
234 THE GEOLOGIC HISTORY OF THE MOON

Dark materials that mantle terra topography are also Late 

Imbrian, judging from the absence of Orientale effects on their surfaces (figs. 5.12,5.15-5.17,6.2,11.6). Most known dark-mantlingmaterials occur along the margins of maria, especially those near the 

center of the Procellarum basin (pi. 4). Carr (1965b; 1966a, b) observed 

superpositions of some dark-mantling deposits on mare surfaces 

(figs. 5.16 5.18) and, accordingly, surmised a volcanic origin (see

chap. 5) and an Imbrian or Eratosthenian age. The smoothness and 

low crater densities of the deposit of dark mantle at the Apollo 17 

landing site led later workers to conclude a Copernican age (Scott and 

Pohn, 1972; Scott and others, 1972; Greeley and Gault, 1973). Embayment relations in several areas showed impingement by mare material on dark-mantled hills (fig. 6.25; Scott and others, 1972), but this 

evidence was not accepted as definitive of an old age. Subsequent

FIGURE 11.8. Three mare units near craters Krafft (K; 51 km, 17° N., 73° W.) and Cardanus (C; 50 km, 13° N., 72° W.; compare fig. 10.3). Old unit (1) underlying Cardanus ejecta is sharply 

truncated by units of uppermost Upper Imbrian group (3) and Eratosthenian System (E; pi. 10). Patches too small to show on plate 9 may belong to middle Upper Imbrian group (2). Orbiter 4 

frame H-169.
11. UPPER IMBRIAN SERIES 235

results, however, have confirmed the validity of the photogeologic 

observations of embayments; post-Imbrian dark-mantling deposits 

are rare.

The following six sections describe the sampled Upper Imbrian 

maria in order of decreasing absolute age of their oldest samples. The 

Apollo 11 suite contains basalt of both the old and intermediate-age 

Late Imbrian groups, as well as some Early Imbrian basalt. The next 

two suites, from the Apollo 16 (one fragment) and Apollo 17 landing 

sites, fall entirely within the intermediate-age group. The last three 

Imbrian suites are in the young group. The rest of this chapter 

discusses unsampled maria according to their position relative to the 

Procellarum basin, which, I believe, has controlled mare distribution 

(pis. 4, 9-11) by thinning the elastic lithosphere and terra crust as a 

prelude to the additional thinning inflicted by the immediate mare 

host basins.

MARE TRANQUILLITATIS

General stratigraphy

Mare Tranquillitatis occupies a pre-Nectarian basin, about 800 

km in diameter, with irregular margins and without definite multiple-ring structure (fig. 11.1). The irregular topography in and near this 

basin results from the intersection of the Tranquillitatis, Nectaris, 

Crisium, Fecunditatis, and Serenitatis basins with two throughgoing 

rings of the Procellarum basin (pi. 3; fig. 11.1). The apparent doublebasin structure of Tranquillitatis (Stuart-Alexander and Howard,

1970) results from impingement of the older and younger rings on the 

outline of Tranquillitatis.

Mare Tranquillitatis contains northern and southern belts of the 

intermediate-age group of Upper Imbrian basalt, separated by a belt 

of the young age group (pi. 9). Most of the central belt and the western 

parts of all three belts are of a blue spectral type that is more 

extensive here than in any other mare (pi. 4; table 5.1). The eastern 

parts of the northern and southern belts, however, are spectrally red. 

Sinus Amoris (fig. 9.11) contains thin highly cratered material that is 

bright and spectrally red. The southeast sector of Mare Tranquillitatis contains similar spectrally red or mottled material, whose 

wavy topography and degraded superposed craters lend a terralike 

appearance (Carr, 1970; Wilhelms, 1970a; Wilhelms and McCauley,

1971), but which result from a very thin cover of mare basalt on 

elevated basin structure (Wilhelms, 1972a). Depths of buried craters 

also indicate that the eastern Tranquillitatis basalt section is thinner 

than the western (fig. 5.2LB; De Hon and Waskom, 1976). These 

differences correlate with a position in the Procellarum basin: thick 

and blue closest to the center, thin and red nearer the edge (pi. 4). 

Units of both age groups, as well as the west basin margin, are cut by 

arcuate rilles (figs. 6.11,11.1).

The unique mare-ridge feature Lamont (fig. 11.1) has elicited 

considerable comment and has been diversely interpreted. It consists 

of both concentric and radial ridges (Morris and Wilhelms, 1967), is 

very conspicuous at low Sun elevations, has a small mascon (Scott, 

1974; Dvorak and Phillips, 1979), and lies athwart a ring of the

Procellarum basin (pi. 4; fig. 11.9; Whitaker, 1981). I suggest that a 

premare crater perched on this ring created the circular part of the 

ridge pattern, contains sufficient basalt to appear as a mascon, and 

remained elevated while subsidence crumpled the surrounding 

basalt flows into a radial ridge pattern.

Apollo 11 samples

The first rocks returned from the Moon are among the oldest 

mare-basalt samples. Apollo 11 landed in the southern intermediateage belt on the oldest of three geologic units mapped near the landing 

site (figs. 9.8, 11.1, 11.10; Grolier, 1970a, b). Most of the samples 

obtained presumably came from this unit or from still older buried 

units. A crater frequency between those for the Fra Mauro Formation 

and the Apollo 15 mare (table 11.1) and aDL value of 335 ± 55 m (Boyce, 

1976) place the surface unit in the intermediate-age group mapped 

here (pi. 9). A nearby, older unit that was also dated by crater frequencies (table 11.1) has not been definitely correlated with any samples.

The landing point of the Lunar Module Eagle was about 400 m 

west of the 180-m-wide, 30-m-deep crater "West" between blocky rays 

of that Copernican crater (figs. ll.lOfi, 11.11; Shoemaker and others, 

1969a; Schmitt and others, 1970). Thickness of the regolith is estimated at 3 to 6 m (fig. 11.12; Shoemaker and others, 1970). Beaty and 

Albee (1978, p. 431-432; 1980) suggested that almost all the samples 

were derived from the ejecta of West crater and thus came from no 

more than 30 m deep (fig. 11.12). Because most of the samples were 

exposed to the space environment for about 0.1 aeon (Eberhardt and 

others, 1970; Guggisberg and others, 1979), West is probably about 0.1 

aeon old. At least one sample (10050) has an exposure age of about 0.5 

aeon and thus was probably excavated by a crater older than West, 

possibly from a basalt unit distant from the landing site (Beaty and 

Albee, 1978, p. 431-432; 1980).

The returned basalt is characterized by its high Ti content. This 

property calibrated the telescopic spectral data (McCord, 1969) and 

showed that blue spectra indicate abundant Ti (Adams and McCord, 

1970; McCord and Johnson, 1970). This Ti is contained in several 

modal minerals, mostly ilmenite (10-17 volume percent; James and 

Jackson, 1970).

Two types of basalt were recognized by megascopic differences 

during the first examination of the rock-size samples (larger than 1 

cm) and were referred to as types A and B (U.S. National Aeronautics 

and Space Administration, 1969, p. 123-126; Schmitt and others, 

1970). These two basalt types have proved to be distinct in composition, microscopic texture, and initial Sr-isotopic ratios as well. They 

also differ in crystallization age and display the largest spread in ages 

from any mare-sampling site.

Low-K, high-Ti. Most basalt that was called type B in the 

original classification appears medium-grained and vuggy in hand 

specimen and ophitic or subophitic in thin section (compare fig. 2.6A). 

It is commonly referred to as Apollo 11 ophitic ilmenite basalt (James 

and Jackson, 1970) or Apollo 11 low-K basalt. Of the 10 rock samples of 

this compositional group, 7 have been well dated at from 3.66 to 3.85 

aeons by the Ar-Ar method (table 11.3; Guggisberg and others, 1979). 

Although the Rb-Sr method is less reliable for low-K rocks, the ages

NNW SSE

FIGURE 11.9. Inferred geologic cross section in southern Tranquiilitatis basin, Sinus Asperitatis, and Nectaris basin, drawn approximately from top to bottom of figure 11.1, with a southward 

extension. Units, from oldest to youngest: pN, pre-Nectarian basin, crater, and volcanic deposits, undivided; Nn, Nectaris-basin primary ejecta; Ns and Nc, primary and secondary ejecta of 

Serenitatis and Crisium basins; If, Fra Mauro Formation (Imbrium-basin ejecta); Im|, Im2 , and Im3 , Upper Imbrian mare basalt; EC, Theophilus ejecta. Lamont is drawn as a buried 

Nectarian crater superposed on a ring of Procellarum basin. An inner ring of Nectaris basin passes under Theophilus, thinning the basalt units. Vertical scale approximate; Theophilus 

measurements from Pike (1980b).
236 THE GEOLOGIC HISTORY OF THE MOON

obtained by the Lunatic Asj'lum of the California Institute of Technology agree approximately with the Ar-Ar ages, as does one Sm-Nd 

age (table 11.3; Papanastassiou and others, 1977). Textures, modal 

mineralogy, and compositions differ sufficiently within the sample 

group to indicate at least three separate flows (fig. 11.12; Beaty and 

Albee, 1978). Ar-Ar ages suggest at least four flows (table 11.3; Geiss 

and others, 1977; Guggisberg and others, 1979), or three flows and an 

ejecta blanket containing sample 10050 (Beaty and Albee, 1980). 

Averages of the best Ar-Ar age determinations for each of these four 

units, performed on plagioclase separates where available, are 3.84, 

3.79, 3.70, and 3.67 aeons.

0-6'N

23.6'E

FIGURE I 1.10. Geologic maps of region of Apollo I I landing site.

A. General region (Grolier, I970b). Units, from oldest to youngest: Im,, older Imbrian 

mare material, surrounded by younger mare unit (unsymbolized here); both units are 

parts of intermediate-age Upper Imbrian unit as defined in this chapter; EC, Eratosthenian crater material; Ccd, Copernican "dimple crater" material; Crct, ray-cluster (secondary) craters of Theophilus. Small craters not separately mapped are indicated by 

letters (E) and numbers (1, 2, 5) corresponding to scheme of N.J. Trask (see chap. 7; 

figs. 7.12 7.14). Barbed arrows and dashed lines indicate minor mare topography.

An early conclusion that sample 10003 is older than the average 

for the compositional group (Turner, 1970) has been confirmed and 

extended to samples 10029 and 10062 (table 11.3). The 3.84 ±0.03- 

aeon average Ar-Ar age obtained for plagioclase separates from samples 10003 and 10029 is close to the 3.85±0.03-aeon age of the 

Imbrium basin, and the Rb-Sr and Sm-Nd ages of 10062 are older 

than 3.85 aeons. A pre-Imbrian age is unlikely, however, because 

Imbrium ejecta is evidently thicker on the adjacent terra than the 

total 30-m-thick basalt section from which the samples are believed to 

have been derived (Beaty and Albee, 1980). Also, the error ranges for 

all age determinations on the older basalt samples and the Imbrium

0°40'

0°38'

23°26' 23°28'

1 KILOMETER

B. Detailed map (Grolier, I970a). Cc5 , young Copernican crater "West" (5 in A); heavy 

dashed line delimits greatest concentration of ejected blocks. EC and low-numbered Cc, 

older craters; Cch, material of irregular crater chain (probably secondary); Im, same as 

in A.

4*

FUG »

^TRAVERSES

Approximate Sorfe 1 250

g

FIGURE 11.11. Traverses by Apollo 11 astronauts relative to Lunar 

Module and Little West Crater (seen as depression tangential to right 

edge of black circle over landing site in figure 11.1 OB).
11. UPPER IMBRIAN SERIES 237

TABLE 11.3. Radiometric ages of large Upper Imbrian basalt samples and pyroclastic glasses

sten and o hers (1973a) 

Nyq ist an others (117

Other method : Sm, Sm-Nd

Group

Apollo 15

green

glass.

Apollo 15

ol i vine

basalt.

Apollo 15

pigeonite

Luna 24

VLT

basalt.

ic basalt.

Apollo 17

orangeand-black

glass.

Apollo 17

low-K,

high-Ti

basalt.

Apollo 16

feldspathic basalt.

Apollo 11 

high-K,

Apollo 11 

low-K,

KH74, Klrsten and 

); N75, Nyquist an

U. U-Pb-Th]

Sample

15086

15426

15016

15385

15555

15607

15633

15659

15668

15678

15683

15058

15065

15075

15076

15085

15117

15499

15682

24077 13

24077,63

24170

L24A

L16-B1

74001

74220

70017

70035

70135

70215

70255

71055

74255

75035

75055

75075

79155

66043,2,

17

10017

10022

10024

10031

10032 

10057

10069

10071

10072

10044

10047

10058

10020

10050

10062

10003

i nn?a

Horn (1974) 

d others (19

Type

___

___

(r)

(r)

(r)

(r)

(r)

(r)

(r)

(r)

(r)

OB

OB

U

U

U

U

r

A

U

 

U

rA

A

n

A

A

A

D1

Q1

Q1

?

B2?

B2

R?

L75, Lugmair an 

75); N76. Nyquis

Station

1

7

3

7

9A

9A

9A

9A

9A

9A

8

1

1

1

1

2

4

9A

4

4

LM

1 M

LM

LM

i &j

1

4

5

5

5

 

9

6

i M

LM

i it

i M

LM i M

i ki

LM

LM

i ki

1 M

LM

LMi kj

LM

LM

1 M

n and others (1975);

d others (1975); L7B,

fer (1977); T73, Tats

Rb-Sr age 

(aeons)

3.22±0.05

3.23±0.08

3.25±0.04

3.2U0.04

3.26+0.08

3.33+0.04

3.28+0.04

3.37+0.07

3.35±0.18 

3.4±0.2

3.60+0.18

3.65+0.11

3.74+0.06

3.67+0.09

3.56+0.09

3

/~ o i n i o

3.75+0.06

3.73+0.14

3.69±0.06

3.74±0.06

3.75+0.10

3.76±0.12

3.51+0.05

3.53+0.07

3.55+0.04

3.60±0.06 

3.60+0.02

3.56+0.05

3.63±0.11

3.63±0.11

3.55+0.11

3

c C_L r\ oo

3.92+0.11

3.76+0.08

Lunatic Asylum (1178); 

=e also Nyquist, 1977);

Reference

E73

M72

PW73

PW73

PW73

PW73

r W / C.O.

BA73

N75

N74

E73

N75

T74

MC76

N76

M n c.

T74

MC76

T73

N75

P70

P70

P70

P70

P77

073

P70

M71 

P70

P77

PW75

M71, Murthy and others

Ar-Ar age 

(aeons)

 3.25±0.06

 3.34+0.06

3.34+0.08

. 3.35±0.05

3.22±0.03

3.24±0.06

- 3.27±0.03

 3.27±0.06

 3.23±0.12

- 3.22+0.05

 3.30+0.04

- 3.10+0.05

 3.34+0.05

 3.32+0.03

 3.32+0.03

 3.41±0.20

3.31+0.04

 3.30±0.08

 3.33+0.21

 3.26+0.06

 3.30±0.04

3

0 o.n OH

 3.61+0.12

3.41±0.04

 3.65+0.02

 3.50±0.05

 3.60+0.04

 3.66+0.03 

- 3.66+0.06

3

C C Q C C

3.69+0.07

 3.79+0.04

 3.79+0.02

3.71+0.05

3.71+0.05 

3.73+0.04

 3.77±0.05 

3.67±0.06

3.69+0.02

- 3.75±0.04

- 3.74+0.05

 3.54+0.06 

- 3.54+0.09

- 3.55±0.08 

 3.53+0.06

3.47+0.06

3.48±0.05

3.57±0.05

3.66±0.04 

- 3.68±0.05

- 3.69±0.03 

3.66+0.04

- 3.72+0.04

 3.70+0.03

3.77+0.06

 3.79+0.04

3.85+0.03 

 3.86+0.07

. T Q-un m

'f ShortyVatertVi E'«st .f ^*ll?* £t of U 9. 2 ^northeast

(1971); H72, Murthy and others (1972); HC76, Murthy and Coscio (1976); N74, 

rs (l<)79a); P70, Papanastassiou and others (1170); P72, Podosek and others

the International IKuon of Geological Sciences ^( loss) (Steiger and Jager,

Reference °^eonsT Method Reference

H74 

PH73 

DO 1 

Hs74 

A074 3.3 U TW75

Hs74 

P72 (p) 

Y72 

Hs74 

Hs74 

Hs74 

Hs74 

Hs74 

Hs74 

Hs74 

Hs74 

573 

Hs74 

C7Q ______ _ _

C7 Q

L78 3.30±0.05 Sm 078

OQ1 _____________

SA78 

LJ7 O

H73 3.48±0.03 U TW76

H78 3.63 U T73

UC71 _________ ___

E73a 

 3.76+0.06 Sm N79a

I/U7/1 ________ ___

Tu73 

U71 ________

J75 3.70+0.07 Sm L75

U7 C

V LJ7 A

CU7 °, ----------------

G79 

Tu70 

G79 

G79 

573 (p) 

Tu70 3.57±0.03 Sm P77

G79 (p) 

G79 

Tu70 

S74, G79 

G79 

G79 

G79 (p) - - -

Tu70 3.88+0.06 Sm P77

G79 (p) 

574 (p) -- - 

Tu70 

R7Q !n\ 
238 THE GEOLOGIC HISTORY OF THE MOON

basin overlap and are large for the Rb-Sr and Sm-Nd ages. The 

3.84±0.03-aeon average Ar-Ar age of the plagioclase separates is 

probably close to the actual age of the flow (barring systematic errors 

in all Ar-Ar ages). Extrusion of the low-K, high-Ti basalt probably 

began soon after the Imbrium impact (fig. 11.9). A pre-Orientale 

(Early Imbrian) age is possible because Orientale ejecta may well be 

absent or very thin here.

High-K, high-Ti. The other type of basalt (originally, type A) is 

richer in K2O, as well as in TiO2 , Na2O, CaO, and A12O3 , than the lowK group (Papike and others, 1976). Grain sizes are finer than those of 

the low-K group, and textures are commonly intersertal (unoriented 

crystal grains separated by small amounts of glass or crystallized 

glass); accordingly, the group is also called Apollo 11 intersertal 

ilmenite basalt (James and Jackson, 1970). For this sample group, RbSr ages are more reliable than Ar-Ar ages (Stettler and others, 1974; 

Papanastassiou and others, 1977). Six rocks have been dated by this 

method (table 11.3). Rb-Sr ages from any single laboratory cluster 

tightly (for example, Papanastassiou and others, 1970, 1977). This 

consistency and the data on textures and bulk compositions are 

consistent with derivation of all the rocks of this type from a single 

flow (Papike and others, 1976; Beaty and Albee, 1978, 1980); only 

minor fractionation near the surface (Papike and others, 1976) or in 

shallow magma chambers (Beaty and Albee, 1978) is indicated. Most 

Ar-Ar ages are younger than the Rb-Sr ages, probably because argon 

was lost after crystallization (Turner, 1977). The average of the best 

Rb-Sr ages and the Sm-Nd age is 3.57 aeons.

Exposure ages (duration of exposure to cosmic rays) of the two 

compositional groups differ and have been interpreted as indicating 

that the high-K flow is the uppermost sampled flow and overlies the 

low-K flows (Eberhardt and others, 1970; Beaty and Albee, 1980), in 

accord with the absolute ages (fig. 11.12). Origin of the two basalt types 

from different magma reservoirs is indicated by the different initial 

Sr-isotopic ratios (Papanastassiou and others, 1977). I tentatively 

correlate the frequency of3.4xlO~ 3 craters at least 1 km in diameter 

per square kilometer and the DL value of 335 m listed in table 11.1 

(both with large error bars) with the high-K flow.

Summary and interpretation

If the source section of all the samples is, indeed, only 30 m thick, 

an incredibly low extrusion rate of one flow averaging 6 m in thickness 

every 54 million years is indicated. Other deposits presumably underlie the sampled flows because the mare surface here is relatively 

unbroken by islands belonging to buried craters or other terra (fig. 

11.1). Nonetheless, a thick section of Upper Imbrian flows is unlikely, 

considering the very old age, within the Imbrian System, of some of 

the lavas that were sampled. Therefore, Lower Imbrian, Nectarian, 

and pre-Nectarian deposits must be present in the subsurface (fig. 

11.9). These deposits include thick ejecta of the nearby Nectaris basin, 

thinner deposits from the Serenitatis and Imbrium basins, and possible deposits of the Crisium and Orientale basins. Pre-Nectarian,

40 -

WEST CRATER

Total volume of ejected

bedrock = 2.7x10 5 m 3

Ejecta composition: 

60 percent A 

10 percent B1 

10 percent B2 

20 percent B3

Type D basalt =?

VERTICAL EXAGGERATION X3

I_________I________

100 100

DISTANCE, IN METERS

FIGURE I 1.12. Inferred basalt stratification penetrated by West Crater and yielding the 

Apollo 11 samples (after Beaty and Albee, 1980). A, high-K, high-Ti basalt; B1, B3, and 

B2, flows of low-K, high-Ti basalt (Beaty and Albee, 1978) B1 may be ejecta blanket. No 

definite B3 basalt is listed in table 11.3; samples 10050 and 10062 may be part of that flow 

(Beaty and Albee, 1978). Type D, possible additional basalt type not dated or discussed 

here.

Nectarian, and Lower Imbrian basalt flows may be interbedded with 

this ejecta sequence. Regoliths formed by smaller craters (not shown 

in fig. 11.9) separate all the deposits in proportion to the timespan 

each was exposed to the surface.

Basalt extrusion was also relatively sluggish in the rest of Mare 

Tranquillitatis during the Late Imbrian Epoch. Although interbedded basalt flows and ejecta may reach a thickness of 1.25 km in one 

spot, most of the total section is thin (fig. 5.21B; De Hon and Waskom, 

1976). Two somewhat indistinct color and age units are present.

In summary, the most recent history of Mare Tranquillitatis is 

characterized by thin and slowly extruded, yet areally extensive, 

lavas of intermediate and late Late Imbrian age. The western flows 

are thicker and richer in Ti than the eastern, except that some 

eastern flows of the young Late Imbrian group are also Ti-rich.

MARE NECTARIS

General stratigraphy

Mare Nectaris differs markedly from Mare Tranquillitatis. 

Although the Nectaris basin is much younger and slightly larger than 

the Tranquillitatis basin (860 versus 800 km diam), mare volcanism 

was even more sluggish. Nectaris is much the smaller mare (table 6.1). 

Little or no basalt of the young Late Imbrian group is present (pi. 9; 

fig. 11.9; Boyce, 1976; Whitford-Stark, 1981). The Nectaris basalt has 

subsided enough to open only a few arcuate grabens on the west 

margin of the mare (pi. 5). The spectral class (mBG-, table 5.1) indicates intermediate values of the spectrally measured properties. 

Orbital geochemical data are not available except in the narrow neck, 

Sinus Asperitatis, that connects Nectaris with Tranquillitatis (figs. 

11.1,11.9), and there the surface is contaminated by Theophilus ejecta. 

The Asperitatis basalt section is as thick as or thicker possibly 1.5 

km than those in central Nectaris and Tranquillitatis (fig. 5.21B; 

De Hon and Waskom, 1976). It even coincides with a small mascon 

(Scott, 1974).

Apollo 16 samples

Although it did not land near a mare and was not designed to 

sample mare materials, Apollo 16 returned a few fragments of mare 

basalt. They were presumably thrown to the sampling site in the 

ejecta of an impact crater, probably Theophilus, in which case their 

source is Mare Nectaris (fig. 11.1; Delano, 1975). The fragments 

include both Ti- and Al-rich types. Ar-Ar analysis of one Ti-rich 

sample yielded an age of 3.74 aeons, consistent with the photogeologically old age of most of Mare Nectaris (table 11.3; Schaefer 

and Husain, 1973).

Summary and interpretation

The differences in age and position relative to the Procellarum 

basin between the Nectaris and Tranquillitatis basins may account 

for the differences between Maria Nectaris and Tranquillitatis. First, 

the fact that Mare Tranquillitatis is extensive, yet relatively thin 

(table 6.1; fig. 5.21B; De Hon and Waskom, 1976), probably results 

from premare leveling of the basin floor by thick Nectaris-basin and 

other ejecta deposits and, possibly, by pre-Imbrian basalt (fig. 11.9). 

Second, the thick lithosphere outside Procellarum may have prevented sufficient subsidence to erase the Nectaris mascon, to open 

large grabens, or to provide much space for basalt extrusion. Each 

condition was opposite in Tranquillitatis, where the weak lithosphere 

allowed extensive basalt extrusion and subsidence. This hypothetical 

significance of lithospheric differences is supported by the relatively 

extensive mare development in Sinus Asperitatis, which is outside 

the centers of Nectaris and Tranquillitatis but inside the rim of the 

Procellarum basin (pi. 4).

MARE SERENITATIS

General stratigraphy

Whereas Serenitatis is among the Moon's least well understood 

basins (see chap. 9), Mare Serenitatis displays clear stratigraphic and 

structural relations that have made it a model for other maria (Carr,
11. UPPER IMBRIAN SERIES 239

J2

00

a E en

O

- 

.« s

I*_*

<3 B

.

II

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1966a; Howard and others, 1973; Thompson and others, 1973; 

Muehlberger, 1974; Head, 1979c; Solomon and Head, 1979). The 

southeast margin contains a continuation of the intermediate-age, 

dark, spectrally blue unit that characterizes northwestern Mare 

Tranquillitatis. This border zone is faulted by conspicuous arcuate 

grabens (pis. 4,5,9; figs. 5.17,6.2,11.13). The southwest-marginal unit 

is younger (pi. 9; Howard and others, 1973), of a similar but less blue 

spectral class (hDWA, pi. 4), and unfaulted. The division between the 

two units falls along the middle ring of the Procellarum basin (pi. 4). 

The center of Mare Serenitatis is occupied by younger, relatively 

bright material of spectral class mISP (table 5.1), which, under the 

designation "MS-2," is the standard for normalizing color spectra 

(fig. 5.19). This red or "orange" material appears to be spectrally 

uniform even though it consists of numerous small flows (not mapped 

here). The central units are not cut by grabens but contain large mare 

ridges (see chap. 6; Muehlberger, 1974; Maxwell and others, 1975; 

Solomon and Head, 1979).

In early work, the dark units were thought to be the younger 

(Carr, 1966a, b), partly because the central mare seemed in telescopic 

photographs to have more craters than the border mare. Along the 

east and west margins, this relation has been confirmed; the dark 

border units there (spectral class hDWA, pi. 4) are younger than the 

light (pi. 10). Stereoscopic spacecraft photographs have shown, 

however, that the central mare is younger than the southeast- and 

southwest-border units; it sharply abuts most of the raised border 

zone and transects the rilles (figs. 5.16,6.2,11.13; Howard and others, 

1973). The apparent abundance of craters in the central mare results 

from their greater brightness and sharpness, which properties are 

caused by excavation of fresh rock through the thinner regoliths on 

the younger lavas (Young, 1975); the older border mare has more 

large, subdued craters.

Serenitatis is a massive mare. The 1-km isopach encloses a large 

part of the mare (fig. 5.215; table 6.1), and the radar-sounder and 

mascon data suggest depths of more than 2 km (see chap. 5). The large 

massive load of basalt depressed the basin and the earlier Ti-rich 

basalt and opened the many arcuate grabens (chap. 6; fig. 6.2; 

Muehlberger, 1974; Solomon and Head, 1979, 1980). Serenitatis, 

therefore, differs considerably from the maria already described. Its 

basalt units are strongly deformed, spectrally distinct, sharply divergent in age, thick, and extensive relative to basin size (figs. 11.13,11.14).

Apollo 17 samples

The largest suite of lunar mare basalt was returned by Apollo 17, 

despite the emphasis of this last Apollo mission on the Serenitatisbasin massifs (see chap. 9). The landing site was on the dark floor of a 

marginal embayment of Mare Serenitatis called the Taurus-Littrow 

Valley (figs. 9.13,11.13,11.15; Wolfe and others, 1981). The "subfloor 

basalt" may be as thick as 1.4 km (Cooper and others, 1974); the 

returned samples were excavated by craters from as much as 100 m 

belowthe surface (fig. 11.15; Wolfe and others, 1981). They are as rich as 

or richer in TiO2 than the Apollo 11 basalt samples and are generally 

similar petrologically. Two main types were at first recognized on the 

basis of major-element composition (Papike and others, 1976): (1) a 

dominant mafic Mg-rich olivine basalt especially rich in Ti (max 14 

weight percent TiO2), and (2) a less abundant, coarse-grained basalt 

very similar to the Apollo 11 low-K, high-Ti variety. The second type 

was obtained from the rim of the largest sampled crater, Camelot, 650 

m in diameter (sta. 5), and thus may have lain at the greatest depth of 

all the sampled basalt possibly deeper than 100 m (Wolfe and others, 1981, p. 109-114,204-205).

Subsequent examination of a larger sample suite by Rhodes and 

others (1976) and Warner and others (1979) disclosed that the gaps in 

major-element composition are occupied and led to a fourfold classification into types A, B, C, and an unclassified type U that is based 

more on trace elements. These authors suggested that the coarsegrained, low-K, relatively low-Ti basalt from the Camelot rim was not 

derived from a separate magma but is a highly fractionated form of 

the type A basalt. Other type A materials are scattered over the 

collection area. Type B basalt was collected mostly in the east half of 

the collection area, especially in ejecta of the 600-m-diameter crater 

Steno (sta. 1). These rocks may have lain somewhat above the source of 

the Camelot samples (Wolfe and others, 1981, p. 35, 205). Type C 

basalt, the rarest and most distinctive, was collected only from the 

rim of the 110-m-diameter crater Shorty (sta. 4) and thus probably
240 THE GEOLOGIC HISTORY OF THE MOON

was derived from the shallowest depths (fig. 11.15). The mostly coarsegrained type U rocks came from all parts of the collection area and are 

the most numerous of the rock-size samples that were dated (table 

11.3).

Rhodes and others (1976) suggested that each of the three basalt 

types A, B, and C differs sufficiently from the others to indicate 

extrusion as a separate lava flow. Variation within each type and, 

probably, among some or all of the type U rocks arose by fractional 

crystallization of one of the magmas on or near the surface.

The scatter in the absolute ages is so great (table 11.3) that no firm 

correlations with the A-B-C stratigraphic sequence or composition 

have been proposed. The Ar-Ar ages range from 3.67 to 3.79 aeons, 

and the Rb-Sr ages from 3.56 to 3.76 aeons; two Sm-Nd ages are 3.70 

and 3.76 aeons (table 11.3). The average of the ages obtained by all 

methods having stated error ranges less than 0.1 aeon is 3.72 aeons. 

The absence of obvious correlations among age, composition, and 

sampling locality suggests that all the Apollo 17 basalt was extruded 

within a short time. The age of 3.72 aeons is, therefore, tentatively 

accepted here as the emplacement age of the entire sequence, more 

than 100 m thick. Partial melting of a heterogeneous source in a small 

zone of the mantle is indicated by the small ranges in age and 

composition (Rhodes and others, 1976).

A major Apollo 17 mission objective was to sample the darkmantling materials, which were believed to be young (Scott and 

others, 1972; Hinners, 1973). They were recovered in the form of 

orange and black glass droplets ("beads") with a mean diameter of 

0.04 mm. This material, which composes from a few percent to 20

percent of the regolith, was recovered in greater concentrations from 

the rim of the 19-million-year-old (Eugster and others, 1977) crater 

Shorty (Muehlberger and others, 1973; Heiken and others, 1974; 

Wolfe and others, 1975,1981). Shorty's ejecta appears dark from orbit 

and black and orange on the surface. Orange droplets compose a bed, 

about 25 cm thick, overlying a bed of black droplets. The color difference of these droplets, whose compositions are the same, is due to 

more nearly complete crystallization of the black (Haggerty, 1974; 

Heiken and others, 1974). The devitrified black beads have a blue 

spectrum, and the glassy orange beads a red spectrum (Adams and 

others, 1974; Pieters and others, 1974). Presumably, the degree of 

devitrification accounts for similar spectra elsewhere, such as the red 

spectra (and orange-brown visual appearance) of the dark-mantling 

materials on the Aristachus Plateau (Zisk and others, 1977).

The telescopic spectrum of "pure" dark-mantling materials north 

of the Apollo 17 landing site is similar to the laboratory spectrum of 

the droplets from Shorty but is only partly like that of droplets from 

the regolith around the landing site (Adams and others, 1974). The 

regolith there probably contains a mixture of true dark-mantling 

materials, consisting of orange and black glass derived from a sheet 

originally at least 1.5 m thick, and of Ti-rich fragments derived from 

the subfloor basalt. The total thickness of the weakly bonded glassbead deposit plus the regolith is 5 to 28 m, and the mean thickness 14 

m (fig. 11.15; Wolfe and others, 1975,1981).

A pyroclastic rather than an impact origin is indicated by the 

chemistry and absence of fragments in the glasses and by the distribution of dark-mantling deposits at the margins of several addiA'"

OCEANUS PROCELLARUM

P1 (P2) (I) (P3) 

Im CEm

MARE IMBRIUM 

CEm

° MARE SERENITATIS

< I (P3) < P2

L CEm

FIGURE 11.14. Geologic cross section, drawn along line A-A'" infigure8.4, showing relative depths of mare basalt and inferred shape of crust-mantle boundary; vertical scale arbitrary. PI, P2, 

and P3, rings of Procellarum basin (pi. 3); I, rim of Imbrium basin; symbols in parentheses, buried rings; Im, Imbrian mare basalt; CEm, Copernican and Eratosthenian mare basalt (pis. 10, 

11). Superposed mantle uplifts resulting from superposed basin impacts are indicated diagrammatically; total uplift is thought to depend on crustal thickness remaining after formation of 

Procellarum basin and on size of superposed basins. Numbers 1 through 7 indicate pathways from mantle source of mare materials to surface at time of eruptions, from shortest to longest.

SHORTY

CAMELOT ALSEP

SAN LUISREY

SHERLOCK

4500m

4400 m

4500m

4400m

EXPLANATION 

Ejecta of Shorty and Van Serg craters

 [JTMI

Light mantle Cluster ejecta

Older regolith of the valley floor

Volcanic ash

Subfloor basalt

2 km

FIGURE 11.15. Geologic cross section, drawn approximately east-west in area of figure 9.17, showing regolith, mare, and glass-droplet ("volcanic ash") deposits in Taurus-Littrow Valley. Light mantle, 

landslide from South Massif; cluster ejecta, ejecta of "Central Cluster" craters believed to be secondary to Tycho; ALSEP, Apollo Lunar Surface Experiments Package (mostly geophysical instruments). 

Vertical elevations above arbitrary datum. From Wolfe and others (1981, fig. 248).
11. UPPER IMBRIAN SERIES 241

tional maria (Heiken and others, 1974). These glasses apparently 

were formed during an early "lava fountain" or "fire fountain" stage of 

mare eruption. At Taurus-Littrow, they were deposited on the subfloor basalt before most of the regolith formed (Head, 1974a; Heiken 

and others, 1974). The pyroclastic material was probably derived from 

a different magma than was the sampled basalt (Hubbard and Rhodes, 1976).

Although sampling confirmed the pyroclastic origin of the darkmantling materials, it refuted estimates of a young age. Determinations of absolute age by different laboratories range from 3.50 to 3.66 

aeons (table 11.3); the average of the single determination for sample 

74001 and the three best determinations for sample 74220 is 3.64 

aeons, and this age is adopted here. All these determinations are 

consistent with superposition of the orange glass deposit on the 

subfloor basalt. The low crater density that led to estimates of a young 

age is due to initially subdued shapes and rapid degradation of 

impact craters in the noncohesive deposits (Lucchitta and Sanchez, 

1975).

Summary and interpretation

The extrusional history of the Apollo 17 basalt was very different 

from that of its compositional relative, the Apollo 11 basalt. In the 

Taurus-Littrow embayment, a massive flow or related flows, more 

than 100 m thick, formed in a time so short that age differences are 

geochronologically unresolvable. Fountains driven by volatile material spewed forth liquid droplets that accumulated as deposits of 

glass. More than 100 million years before and after the lavas and 

pyroclastic materials poured out here, a few meters of basalt were 

being added to the margin of Mare Tranquillitatis. Then, after both 

margins were covered by the high-Ti basalt flows, the center of 

Serenitatis was filled by voluminous low-Ti basalt flows that 

depressed and faulted the older Ti-rich basalt and the basin. The 

spectral and stratigraphic uniformity of the many small flows of lowTi basalt suggests high eruption rates.

Although Maria Serenitatis and Tranquillitatis are similar in 

area, the volume of Serenitatis is by far the greater second only to 

that of Mare Imbrium (fig. 5.21; table 6.1). The voluminous volcanism 

is due to the relatively young age of the Serenitatis basin (see chap. 9), 

to its premare filling by only one major deposit (Imbrium ejecta), and, 

probably, to its superposition on a relatively weak part of the 

lithosphere relatively near the center of the Procellarum basin (figs. 

11.9,11.14).

MARE FECUNDITATIS

General stratigraphy

The poorly photographed Mare Fecunditatis occupies a pre-Nectarian basin, about 990 km in diameter, outside the Procellarum 

basin. The mare is shallow, spectrally mottled and bland, relatively 

aluminous in X-ray readings (Adler and others, 1972; Hubbard, 1979), 

and apparently middle to late Late Imbrian in age (pi. 9). Like Mare 

Tranquillitatis, Mare Fecunditatis occupies a large fraction of the 

basin diameter but is relatively thin (table 6.1). Its west edge is faulted 

by arcuate grabens.

Selenographically, Mare Fecunditatis includes a subcircular 

southward extension of the main mare. This second mare, about 250 

km in diameter, may occupy an (unmapped) large crater or small 

basin superposed on the 990-km-diameter Fecunditatis basin. Alternatively, the Fecunditatis impact may have been double.

Luna 16 samples

The Luna 16 unmanned sample-return mission (September 

1970) yielded 101 g of small regolith fragments from Mare Fecunditatis (lat 0.7° S, long 56.3° E.; fig. 11.16; McCauley and Scott, 1972). 

The dominant basalt is a feldspathic, high-Al type distinguished by a 

high FeO/MgO ratio from other mare basalt, KREEP-rich material, 

and high-Al mare-basalt clasts from the Apollo 14 breccia (Kurat and 

others, 1976). The Rb-Sr and Ar-Ar ages agree fairly well at an 

average of 3.38 aeons, but the Ar-Ar age of 3.41 ± 0.02 aeons has much 

the smaller error range (table 11.3); a weighted average of 3.40 aeons 

is adopted here. The sampling site is on a dark unit whose spectral

class (hDW-) indicates a higher Ti content than most of Mare Fecunditatis (pi. 4). The Al-rich sample composition is consistent with the 

orbital detection of relatively high Al/Si ratios for the overflown part 

of the mare (Hubbard, 1979).

Summary and interpretation

The aluminous composition of the basaltic magmas may have 

favored their hydrostatic ascent and extrusion through the thick 

crust in the Fecunditatis region, because they probably are the marebasalt magmas closest in density to the terra crust (2.8-2.9 g/cm3 ; 

Solomon, 1975). Possibly, Al-rich magmas could be extruded where 

denser magmas would remain beneath the surface.

Of all the maria considered here that lie outside Procellarum, 

Mare Fecunditatis occupies the largest percentage of its basin's area 

(table 6.1). A possible explanation is that the younger Nectaris and 

Crisium basins "assisted" the volcanism. The grabens in western 

Fecunditatis lie within the rim of the Nectaris basin (pi. 5). Northern 

Mare Fecunditatis, including the compositionally anomalous (relatively Ti-rich) Luna 16 material, lies within the proposed 1,060-kmdiameter rim of the Crisium basin (pi. 9; fig. 11.16). Therefore, the 

lithospheric thinning required for mare extrusion and structural 

deformation is achieved here by superposition of basins other than 

Procellarum.

MARE CRISIUM

General stratigraphy

Mare Crisium occupies an elliptical area, 420 km north-south by 

560 km east-west (pi. 4). The average diameter of this mare is less 

than half that of the containing Crisium basin in my interpretation 

(1,060km; see chap. 9; tables 4.1,6.1). Crisium is probably between the 

Serenitatis and Nectaris basins in age, closer to Serenitatis (table 

9.3). Mare Crisium has properties in common with both Maria 

Serenitatis and Nectaris. It has a mascon, as do both of these maria. 

The basalt section is about as thick as that of Serenitatis and thicker 

than that of Nectaris (fig. 5.2LB; De Hon and Waskom, 1976). In 

respect to its areal extent relative to the basin diameter, however, it 

more closely resembles Nectaris (if the Crisium basin is indeed as 

large as I suggest). Tectonically, Crisium is more like Nectaris than 

Serenitatis because no arcuate grabens have opened since the visible 

mare formed (pi. 5).

Mare Spumans, Mare Undarum, and other "lakes" surround the 

main Mare Crisium at considerably higher elevation (fig. 4.9). In 

chapter 4, this relation was interpreted as analogous to the steplike 

structure of the Orientale basin and its small superposed maria. The 

"lakes" are as high as or higher than the adjacent Mare Fecunditatis 

(fig. 6.23).

Investigative groups led by J.M. Boyce (Boyce and Johnson, 1977; 

Boyce and others, 1977) and J.W. Head (Head and others, 1978a) 

mapped several age units, but these two groups did not agree on the 

units' distribution. Photographs suitable for age determinations are 

available only in the southern part of the mare under the Apollo 17 

orbits. To me, Crisium seems to be among the stratigraphically most 

uniform maria, and I map it here entirely as part of the youngest of 

the three Late Imbrian age categories (pi. 9). The "lakes" are older.

Head and others (1978a) and Pieters (1978) mapped several color 

units in Mare Crisium (pi. 4). An Mg-rich unit similar to the Luna 16 

basalt is exposed at the southeast margin (hDWA, pi. 4); this may also 

be the unit exposed in the ejecta of the crater Picard (fig. 5.23; Andre 

and others, 1978; Head and others, 1978a). The unit sampled by Luna 

24 is the type area of spectral unit LISP. Unit hDWA is moderately 

rich in Ti, but the other spectral units are relatively poor in Ti and 

rich in Fe (table 5.1). Hubbard (1979) interpreted data obtained by the 

orbiting X-ray spectrometer to mean that Mare Crisium is more 

aluminous than the maria farther west, about equal to Mare Fecunditatis, and less aluminous than only Mare Smythii among the major 

analyzed maria.

Luna 24 samples

The most recently (August 1976) collected lunar sample is a 160- 

cm-long drill core obtained from the regolith of Mare Crisium by the
242 THE GEOLOGIC HISTORY OF THE MOON

unmanned Soviet spacecraft Luna 24 (lat 12.7° N., long 62.2° E.; fig. 

9.3; Butler and Morrison, 1977; Florensky and others, 1977). Most of 

the basalt fragments are very low in Ti (less than 1.5 weight percent 

TiO2), a basalt type previously found only in small amounts in the 

Apollo 17 drill core (Vaniman and Papike, 1977). The Luna 24 VLT 

"ferrobasalt" is richer in FeO than the Apollo 17 VLT basalt and is also 

highly feldspathic (12.5-14 weight percent A12O3 ; Papike and Vaniman, 1978). In addition, the Luna 24 core contains an MgO-rich 

olivine basalt and, possibly, several additional types or subtypes (for 

example, see Ma and others, 1978). The spectra indicate the presence 

of additional units of Mare Crisium that were not sampled. Discovery 

of the VLT basalt caused a rethinking of previous classification 

schemes (Papike and Vaniman, 1978). In general, current thinking 

favors a continuum of lunar-basalt types rather than the pigeonhole 

classifications that emerged from earlier missions.

Despite the small sample size, accurate absolute ages of the 

underlying bedrock apparently were obtained from the regolith fragments (table 11.3; Lunatic Asylum, 1978). Identical Sm-Nd and 

40Ar-39Ar ages fix the crystallization age of the Luna 24 VLT ferrobasalt at 3.30 aeons (Lunatic Asylum, 1978). The Lunatic Asylum

(1978) did not regard Rb-Sr ages as reliable for these basalt samples, 

and so none are given here. An Ar-Ar age for a different sample that is 

0.3 aeon older (table 11.1; Stettler and Albarede, 1978), and the presence of the buried Mg-rich basalt (Andre and others, 1978; Head and 

others, 1978a), suggest an extended history of basalt extrusion in 

Mare Crisium.

Summary and interpretation

The properties of Mare Crisium are consistent with excavation of 

a large, young, deep basin in a previously thick crust. I speculate that 

relatively thick crustal material favored aluminous mare volcanism 

and resistance to deformation, as in Mare Fecunditatis. Unlike 

Fecunditatis, however, the Crisium basin was filled by thick basalt 

because it had not been filled previously by thick basin ejecta (fig. 

11.17). An older and smaller basin filled by thick earlier deposits would 

resemble Fecunditatis. A smaller basin of the same (Nectarian) age as 

Crisium would contain a small mare like Nectaris if it were in the 

same position as Crisium, or a large mare like Serenitatis if it were 

inside the Procellarum basin.

FIGURE 11.16. Mare Fecunditatis and south Crisium-basin rim, showing Luna 16 and Luna 20 (see chap. 9) landing sites. Orbiter I frame M 33.
11. UPPER IMBRIAN SERIES 243

MARE IMBRIUM

General stratigraphy

The center of Mare Imbrium contains the Moon's thickest known 

mare-basalt section; the 1-km isopach extends to the largest diameter 

of all lunar maria and probably encloses several kilometers of basalt 

in the center (De Hon, 1979). Subsidence has opened large concentric 

grabens on the Apennine Bench and, in lesser degree, elsewhere (pi. 

5; figs. 6.10,10.7). Deformation was undoubtedly great in other, nowburied sectors, as is shown by Orientale, where most grabens would 

be buried if filled by a larger mare. Mare Imbrium has a large mascon, 

though not so large (table 6.1) as would be expected from the Moon's 

third largest basin (after Procellarum and South Pole-Aitken). The 

thin lithosphere may have allowed a more complete isostatic adjustment here than at Crisium.

The mare units of Imbrian age in Mare Imbrium are mostly 

spectrally red and bright (pi. 4; fig. 5.20; Head and others, 1978b). As 

in central Serenitatis, these units are more complex than can be 

shown in plate 9 and consist of a patchwork of units ranging in age 

from early to late Late Imbrian (DL = 240-385 m; Boyce and others, 

1975; Boyce, 1976; Wilhelms, 1980). Eratosthenian mare materials in 

Mare Imbrium are dark and relatively blue (see chap. 12). The spectrally red and blue types contrast sharply with each other but are 

each relatively uniform internally.

The Mare Frigoris-Sinus Roris trough is concentric with both the 

Imbrium and Procellarum basins (pis. 3, 4). It contains generally 

bright and thin mare materials of different spectral classes from most 

of Oceanus Procellarum but similar to some of the older units in the 

adjacent part of northern Mare Imbrium (pi. 4).

Apollo 15 samples Palus Putredinis

Apollo 15, the first complex Apollo mission, was able to explore a 

wide area of both mare and terra systematically (see chap. 10). The 

peculiar sinuous rille Rima Hadley (figs. 5.11,10.7,10.9) helped attract 

attention to the landing site, which lies in the marginal mare of Mare 

Imbrium called Palus Putredinis. At one time, the rille or its source 

crater was thought to be a possible source of rare lunar volatile 

materials, although this hope faded with analyses of the volatile-free 

Apollo 11 basalt samples. The spectrally red mare of Palus Putredinis 

was a bonus to the more significant mission objective, materials of the 

Imbrium-basin rim (chap. 10; McCord and others, 1972b). This color 

indicates a low Ti content (Pieters, 1978), and the returned basalt 

samples contain much less of this element (1.5-2.0 weight percent 

TiO2 ) than any other mare suite except that of Luna 24. The Palus 

Putredinis mare differs somewhat in remotely sensed properties from 

the spectrally red part of Mare Imbrium proper (Head and others, 

1978b). Before the mission, the mare's age was correctly determined 

by the Traskmethod (see chap. 7; figs. 7.12-7.14) as near the ImbrianEratosthenian boundary (Carr and others, 1971). As was true for 

Apollo 17, more rock-size mare-basalt samples than terra materials 

were returned because the basalt was more accessible.

A twofold classification of the samples into an older, dominant 

type A and a younger, less abundant type B that was proposed upon 

preliminary examination of the field relations, modal mineralogy, and 

textures of the returned samples (Swann and others, 1972) has withstood further testing (for example, Rhodes and Hubbard, 1973; 

Lofgren and others, 1975). Type A is now generally referred to as 

quartz-normative or pigeonite basalt, and type B as olivine-normative or olivine basalt (table 5.2). Two conspicuous types of 

pyroclastic materials are also among the returned samples.

Pigeonite basalt. The basalt type richer in silica is characterized by abundant modal pyroxene, mostly pigeonite with augite 

mantles. This basalt composes most of the rock-size samples collected 

by Apollo 15. It is exposed in boulders that are nearly in place at the 

rille lip (sta. 9A, fig. 10.9) and in ejecta of the major craters Elbow (400 

m diam, max 80 m deep; sta. 1) and Dune (460 m diam, max 90 m deep; 

sta. 4). Textures are mostly porphyritic, and the groundmass of one 

subgroup is megascopically glassy (vitrophyric). The various textures 

are compatible with derivation from a single flow (Lofgren and others, 

1975). Judging from the field relations, this flow constitutes most or 

all of the 60-m-thick section of layered lavas exposed in the walls of 

Hadley rille (figs. 1.105,11.18; Howard and others, 1972; Swann and 

others, 1972; Lofgren and others, 1975). The Apennine Bench Formation, which may be the source of small KREEP-rich fragments in the 

regolith (see chap. 10), may underlie this flow or a still older marebasalt flow (fig. 11.18).

The best Ar-Ar ages of rocks and of smaller samples collected by a 

rake cluster between 3.30 and 3.32 aeons (table 11.3); Rb-Sr ages 

range from 3.21 to 3.37 aeons. Despite this spread in Rb-Sr ages, 

petrologic and chemical relations among the samples indicate moderate crystal fractionation of a single magma on or near the surface 

(Rhodes and Hubbard, 1973). The average crystallization age of the 

samples dated by both methods is 3.30 aeons.

Olivine basalt. Rock-size olivine basalt was collected only from 

station 9A at the edge of Rima Hadley and from station 3 at Rhysling 

Crater (Swann and others, 1972). The near-surface position of the 

olivine basalt suggested by this distribution is corroborated by its 

greater abundance in smaller fragments, at all collection stations, 

than the pigeonite basalt (Swann and others, 1972; Rhodes and Hubbard, 1973; Lofgren and others, 1975). The petrologic and compositional data are consistent with derivation from a single flow different

0 100 200 300 METERS

NO VERTICAL EXAGGERATION

FIGURE 11.18. Geologic cross section of mare materials at Apollo 15 landing site, drawn 

northwest-southeast on figure 10.7. Rima Hadley profile is from Howard and others (1972). 

Thicknesses of basalt units and presence of older basalt and Apennine Bench Formation are 

hypothetical.

MARE CRISIUM MARE FECUNDITATIS

0

w 1

DC 

LU 9

1- *

04 

* 5

6

O <£>

CM *-

as to

c. c 

 Mn -3 lisc Im 3 _i Nc

. . ^^^^^^smf^^mi^fa^1 ^ , _ , .. __**=_ Jg8"""M6i!Ji!!W¥ / / /

_si Ns, Nv? pN

Nn Imi,lm2

FIGURE 11.17. Inferred geologic cross section in Crisium and Fecunditatis basins, drawn along line A" '-Fin figure 8.4, with an extension to lat 5° S. Units, from oldest to youngest: pN, pre-Nectarian basin, 

crater, and volcanic deposits, undivided; Nn, Nectaris-basin ejecta; Nc, Crisium-basin ejecta (inner knobby and outer lineated types); Ns, Serenitatis-basin ejecta; Nv?, possible Nectarian volcanic deposits; 

lisc, deposits of Imbrium-secondary craters; lrri|, Im2 , and Im3 , Imbrian mare units.
244 THE GEOLOGIC HISTORY OF THE MOON

in parentage from the thicker pigeonite basalt flow (Rhodes and 

Hubbard, 1973).

The 9.6-kg porphyritic olivine-basalt boulder-size sample 15555 

("Great Scott") from station 9A has been dated by three geochronologic methods at 3.25 ± 0.05 aeons (table 11.3). One other rock 

(sample 15016) and numerous walnut-size rake samples from many 

collection stations average 3.27 aeons in age (Husain, 1974); the 

average by all methods is 3.26 aeons. This younger age is consistent 

with stratigraphic superposition on the pigeonite basalt but may not 

be significant in view of the spread in age determinations for each 

flow; the 40-million-year age difference between flows is just at the 

limit of geochronologic discriminability.

Pyroclastic green glass. Apollo 15 also collected more 

pyroclastic material than did any other mission except Apollo 17. 

Green-glass droplets attracted the astronauts' attention and were 

collected from the regolith at two main stations (table 11.3). Emplacement by both impact and volcanism has been considered (see summary by Delano, 1979). Delano (1979) suggested that compositional 

data that were once seen as supporting evidence for an impact origin 

are more consistent with a volcanic origin. He proposed a derivation 

from heterogeneous sources rich in volatile material at 400 ± 50-km 

depth.

Although some older ages have been reported in the literature for 

the green glass, ages close to those of the lavas appear to be more 

likely (table 11.3). A dark-mantling deposit observed photogeologically near the landing site (Carr and others, 1971) could be the source of 

the green-glass samples (Carr and Meyer, 1974). However, the green 

glass, which is not dark in hand specimen, may be from a mantling 

deposit not detected either photogeologically or spectrally.

Pyroclastic red glass. Droplets of red glass were also recovered 

by Apollo 15. Because this glass was found at only one station (6), it 

may be exotic to the landing site (Quaide, 1973). Delano (1980) suggested that its parent magma formed at the great depth of 480 km.

Summary and interpretation

Mare Imbrium is similar to western Mare Serenitatis in general 

stratigraphic sequence. Old, spectrally blue border basalt typical of 

eastern Serenitatis is not exposed. The main basin-filling sequence in 

both maria is thick, Ti-poor, relatively uniform spectrally, but complex in internal stratigraphy. Both sequences occasioned substantial 

basin subsidence, and both were followed by lesser amounts of spectrally blue Eratosthenian basalt erupted from the basin margins 

more in Imbrium than in Serenitatis. These similarities probably 

result from superposition of the two relatively young basins on the 

thin lithosphere and crust in the central Procellarum basin (figs. 8.4, 

11.14).

OTHER MARIA INSIDE 

THE PROCELLARUM BASIN

The outermost shelf of Oceanus Procellarum contains a large 

expanse of basalt of the youngest Late Imbrian age group (pi. 9). Most 

of this tract belongs to the same spectral class as central Serenitatis 

(mISP, pi. 4) but may be poorer in Fe (Basaltic Volcanism Study 

Project, 1981, p. 455). This young Late Imbrian basalt transects older 

Late Imbrian mare units and is overlain in places by Eratosthenian 

units (fig. 11.8). At the north end of the mare-filled part of the shelf and 

in the adjacent middle trough of the Procellarum basin, intermediateage Upper Imbrian material of red spectral class LEG- overlaps the 

very old, spectrally blue units (class hDWA; Luna 16 type) perched on 

the mare margin. Class LEG- (unit 11, pi. 4) is also typical of Lacus 

Somniorum, in the eastern hemisphere, where intermediate-age and 

old Late Imbrian mare units overlie the same shelf and trough of the 

Procellarum basin (pis. 4, 9). Spectrally blue Eratosthenian units 

occupy the middle trough in much of the western nearside and locally 

extend onto the outer shelf (see chap. 12; pi. 10; fig. 12.11). The persistence of similar concentric compositional patterns in so much of 

the Procellarum basin, in otherwise different geologic settings, is 

remarkable.

The middle Procellarum trough is also the site of the Moon's three 

largest complexes of nonplanar volcanic deposits and unusual landforms (see chap. 5; pi. 4). The northernmost complex is the Riimker 

Hills (Mons Riimker), which consists of domes and dark-mantling 

deposits (fig. 5.9). Intermediate is the Aristarchus Plateau, which 

contains the Moon's largest and most numerous sinuous rilles and 

thick dark-mantling materials (fig. 5.12; Zisk and others, 1977). Farther south are the Marius Hills, characterized by diverse cones, 

domes, and sinuous rilles (fig. 5.7; McCauley, 1967a, 1968, 1969b). 

Although the morphologic freshness of the landforms in these three 

complexes suggests an Eratosthenian or even a Copernican age 

(Moore, 1965, 1967; McCauley, 1967a, b; Wilhelms and McCauley, 

1971), they are partly flooded by late Late Imbrian lavas and so are 

mostly or entirely Imbrian in age (Zisk and others, 1977). The middle 

Procellarum trough apparently was conducive not only to late volcanism but also to extrusion of diverse magma types.

Oceanus Procellarum and other western-hemisphere maria east 

of long 40 W. have complex, mottled patterns of unit distribution that 

cannot be completely portrayed here (see Pieters and others, 1980; 

Whitford-Stark and Head, 1980; Wilhelms, 1980). Much of this region 

is covered by spectral class mIG- (Apollo 12 type), which represents a 

class of related compositional units (see chap. 5; Pieters, 1978). This 

complexity probably results from superposition on the larger Procellarum-basin structures of several old basins, such as FlamsteedBilly, Insularum, Nubium (pi. 3), and the small basin or large crater 

that contains Mare Cognitum. The lithospheric and mantle structure 

resulting from these superpositions is complex and presumably influences a complex pattern of mare extrusion.

Oceanus Procellarum is generally stated to be a nonmascon 

mare, but, in fact, parts of its outer troughs contain small linear 

positive anomalies (Scott, 1974). Although Scott (1974) ascribed these 

anomalies to intrusion of dikes beneath linear mare ridges (fig. 6.2), I 

suggest that they are due to relatively thick sections of basalt in the 

troughs. They alternate with parallel negative anomalies aligned 

along the buried rings of the Procellarum basin (compare pi. 3 with 

fig. 2 of Scott, 1974).

The Procellarum pattern of mare distribution seems to persist 

into Mare Humorum, which has the third large mascon inside the 

Procellarum basin (besides Serenitatis and Imbrium; see table 6.1). 

Mare Humorum occupies the complex Nectarian Humorum basin 

(see chap. 9; figs. 4.3C7, 9.24), which is superposed on the outer 

Procellarum shelf and periphery. The central mare is the most extensive, but maria are also abundant in parts of the north basin periphery that overlie the Procellarum basin (pi. 4). The presence of 

numerous conspicuous grabens (figs. 6.9,9.24) indicates considerable 

subsidence. Mare Humorum contains mare units of diverse color (pi. 

4) and age classes (Pieters and others, 1975; Boyce, 1976; Boyce and 

Johnson, 1978). The western part of the basin contains spectral class 

mISP, and the eastern part mIG-, as in Oceanus Procellarum. This 

relatively red unit is covered by a young (late Imbrian and (or) 

Eratosthenian), moderately blue unit originating outside central 

Mare Humorum. The southern contact between the mare and the 

basin is draped by a dark-mantling deposit called the Doppelmayer 

Formation (Titley, 1967).

The central highlands of the Moon apparently were a barrier to 

mare extrusion (pis. 4, 9; fig. 1.8). Between the Insularum-Nubium 

and Tranquillitatis-Asperitatis mare zones, hardly any maria occupy 

the zone between the Procellarum-basin rim and middle ring. Furthermore, the maria inside the middle ring are much smaller than 

elsewhere in the same radial position from the Procellarum center; 

they apparently required the additional lithospheric penetration 

afforded by such small basins or large craters as Vaporum, Aestuum 

(Wilhelms, 1968), and, possibly, a circular Sinus Medii structure (De 

Hon, 1979) not mapped here. The terra borders of these maria are rich 

in dark-mantling deposits (Wilhelms, 1968; Wilhelms and McCauley, 

1971). Penetration by the Mutus-Vlacq and Werner-Airy basins was 

apparently insufficient to allow Late Imbrian volcanism, although 

pre-Nectarian basalt may be buried by ejecta plains in the MutusVlacq basin (see chaps. 8, 9), and a small dark deposit in the crater 

Airy may be volcanic in origin (unless it is ejected Imbrium-basin 

impact melt). Thus, the south-central terra "backbone" may be a firstorder crustal feature. A northward extension seems to block mare 

extrusion in the central part of the outer Procellarum trough (pi. 4). 

These major terra structures may have arisen during the "magma 

ocean" stage, or they may be remnants spared from excavation by 

giant undiscovered basins still older than Procellarum.
11. UPPER IMBRIAN SERIES 245

NEAR-LIMB AAARIA OUTSIDE 

THE PROCELLARUM BASIN

Mare Orientale and Mare Smythii are similar in many respects. 

They are thin mascon maria (Solomon and Head, 1980), and their 

basins are of about the same size (table 4.1). They also are equally 

deep (Sjogren and Wollenhaupt, 1976; Strain and El-Baz, 1979) a 

surprising fact at first glance, considering the difference in age of 

their basins (Imbrian and pre-Nectarian, respectively). Both basins 

contain floor-rebound craters (pi. 5); Smythii has more because it is 

the older basin and thus contains more craters of all types. The ringand-trough structure is similar in the two basins (Strain and El-Baz, 

1979). Smythii lacks arcuate grabens, and those of Orientale, though 

conspicuous and numerous, are confined to a smaller area of the basin 

than those in more deeply filled maria (pi. 5; table 6.1). The paucity of 

grabens is expected from their position on a thick lithosphere, which 

resisted vertical movements in both basins.

As in other maria superposed on thick terra crusts, the Smythii 

basalt appears to be highly aluminous (Conca and Hubbard, 1979; 

Hubbard, 1979), although this appearance may result from mixing of 

the thin mare with terra material (P.D. Spudis, oral commun., 1982). 

Like Crisium, Smythii is low in radioactivity (Metzger and others,

1977). Mare Orientale was not analyzed.

The two maria apparently had somewhat different patterns of 

basalt emplacement. Smythii contains late Imbrian units (pi. 9) 

thinly covered by Eratosthenian units (pi. 10; Boyce and Johnson,

1978). At Orientale, extrusion ended during the middle Late Imbrian 

in the central mare, which is comparable to all of Mare Smythii, but 

migrated to the inner and then to the outer crescentic maria in the 

concentric troughs (Lacus Veris and Autumni, respectively; Greeley, 

1976).

Small, thin, old maria occupy some of the terrain between the 

Humboldtianum and Procellarum basins in the northeast quadrant 

of the nearside (pi. 9; fig. 9.3). A very old, still-undelineated basin may 

exist in this region.

Mare Australe probably contains a more nearly equal amount of 

all three Imbrian mare units than any other studied mare (pi. 9). 

Most of the maria occupy craters superposed on the Australe basin, 

and much less basalt probably would have erupted in this shallow 

weakly defined basin if these craters were absent. Mare ridges, but 

not grabens, are present (pi. 5; Wilhelms and others, 1979).

Mare Marginis is another near-limb mare that contains old (fig. 

4.7) as well as young (Boyce and Johnson, 1978) mare units. Some of 

the basalt is KREEP-rich (Hawke and Spudis, 1980).

FARSIDE MARIA

Most maria that lie entirely on the farside are poorly characterized because of absence of telescopic spectra, paucity of orbital 

geochemical readings, and poor photographic coverage (pi. 2). No 

arcuate grabens are known, and they are probably absent (pi. 5). I 

know of only five farside mare patches that have been dated by crater 

counts (Wilbur, 1978; Walker and El-Baz, 1982), all of which indicate a 

middle Late Imbrian age (pi. 9).

The farside's most extensive maria, those in the South PoleAitken basin, are similar to those of Australe in that they occupy 

parts of small superposed basins (pis. 4, 9) and large craters. Thus, 

South Pole-Aitken played a lithosphere-thinning role on the farside 

comparable to that of the Procellarum basin on the nearside. The fact 

that maria are sparser in South Pole-Aitken than in Procellarum is 

probably due to: (1) the smaller size of South Pole-Aitken, 2,500 

versus 3,200 km; (2) the absence of such large superposed basins as 

Imbrium and Serenitatis; and, possibly, (3) a first-order hemispheric 

dichotomy in crustal thickness, such as is commonly hypothesized 

(see chaps. 1, 6, 8).

Other farside maria are scattered in the few other large basins, in 

a few large craters, or in such seemingly isolated and unexplained 

spots as the crater Kohlschutter and Lacus Solitudinis (pi. 4; figs. 5.2, 

5.3). The Moscoviense basin contains the largest mare patch outside 

the South Pole-Aitken basin; photographic coverage is inadequate for 

dating the mare, but it appears to be old. The mare's relatively large 

area may result from superposition of two basins (D.E. Davis, oral 

commun., 1982), only one of which is mapped here. The conspicuous 

mare in crater Tsiolkovskiy probably results from the considerable

depth of this large Late Imbrian crater, aided by its superposition on 

the Tsiolkovskiy-Stark basin. The mare, which is covered by one of 

the best Lunar Orbiter photographs of the farside (an Orbiter 3 Hframe; pi. 2), is Imbrian in age. DL values of 255 m (Boyce and others, 

1975) suggest a late Late Imbrian age, whereas crater frequencies 

(Wilbur, 1978; Walker and El-Baz, 1982) suggest a middle Late 

Imbrian age. The basalt may be rich in Ti (Metzger and others, 1977; 

Wilbur, 1978). Finally, pre-Imbrian or Early Imbrian basaltic plains 

are likely to be hidden beneath the smooth, light-colored plains 

fillings of such basins as Hertzsprung and Korolev (see chap. 9).

CHRONOLOGY

Late Imbrian mare basalt is the most accurately dated lunar 

material. The Late Imbrian Epoch began immediately after the Orientale impact, which is estimated from superposed crater frequencies 

to have occurred 3.8 aeons ago, between the Imbrium impact 3.85 

aeons ago and the extrusion of the Apollo 17 mare basalt 3.72 aeons 

ago (see chap. 10). The youngest sampled Imbrian mare basalt, from 

the Apollo 15 landing site, formed 3.26 aeons ago. Therefore, the Late 

Imbrian Epoch lasted at least 0.54 aeon and probably a little longer, 

allowing for extrusion of some unsampled units after 3.26 aeons (for 

example, in Maria Serenitatis and Crisium; table 11.1). An approximate date of 3.2 aeons ago, between the Apollo 15 and Apollo 12 dates, 

is adopted here for the end of the Late Imbrian Epoch and the 

Imbrian Period. Therefore, the Late Imbrian Epoch lasted about 0.60 

aeons and the Imbrian Period about 0.65 aeons.

Exposed basalt units of the oldest of the three Late Imbrian age 

groups mapped here are all on the outer shelf of the Procellarum 

basin or outside that basin. Their age is estimated at from 3.75 to 3.80 

aeons; the exposed flows were not sampled, but the 3.79-aeon Ar-Ar 

age of one sample from a buried unit at the Apollo 11 landing site 

(10062; table 11.1) lies within this age range.

Exposed flows of the intermediate-age group are represented by 

the Apollo 17 Ti-rich basalt (3.72 aeons) and the Apollo 11 high-K, Tirich basalt (3.57 aeons). The age range of this group suggested by 

these absolute ages, as calibrated with the relative stratigraphy, is 

3.50 to 3.75 aeons. In addition, therefore, an Apollo 16 fragment 

probably from Mare Nectaris (3.74 aeons), two low-K, high-Ti Apollo 

11 units (3.67 and 3.70 aeons), and the Apollo 17 pyroclastic deposit 

(3.64 aeons) also probably originated in middle Late Imbrian time. 

Orbital geochemistry indicates that compositions of this age are very 

diverse. Many units are deformed by grabens.

The youngest Late Imbrian age group, 3.20 to 3.50 aeons old, is 

represented by the Luna 16 feldspathic basalt (3.40 aeons), the Luna 

24 VLT basalt (3.30 aeons), the Apollo 15 green glass (3.30 aeons), and 

the Apollo 15 low-Ti olivine basalt and pigeonite basalt (3.26 and 3.30 

aeons, respectively). Few units of this group are cut by grabens, but 

most are deformed by mare ridges. Spectrally red or "orange" basalt 

units with medium to low Ti contents characterize the most extensive 

exposures of this age group, that is, those in Mare Serenitatis, Mare 

Imbrium, and western Oceanus Procellarum. Basalt of this composition may be the Moon's most voluminous basaltic type.

Dark-mantling materials lie at the margins of most maria but 

predominate around maria superposed on the center and intermediate trough of the Procellarum basin. Many of these dark-mantling 

deposits are older than the adjacent lavas. Spewing of magmas driven 

by volatile materials probably was important in the early filling of 

basins (Head, 1974a). At 3.64 aeons, the middle Late Imbrian Apollo 

17 orange and black glass is somewhat younger than the subjacent 

Apollo 17 lavas (3.72 aeons). The 3.3-aeon age of the Apollo 15 green 

glass shows that pyroclastic materials were also erupted late in the 

Late Imbrian Epoch.

In the Imbrian Period, the cratering rate changed from the intensive early bombardment (fig. 8.16) to a much-reduced rate that 

characterizes the rest of lunar history. This rapid falloff in impact rate 

is recorded by successively lower crater densities for successively 

younger time-stratigraphic units (pis. 7-11; figs. 7.16,8.6,9.22,10.4). 

During the total Imbrian Period, 9.7 craters larger than 30 km in 

diameter formed per million square kilometers (pis. 8, 9), in contrast 

to about 35 such craters in the Nectarian Period and more than 77 in 

pre-Nectarian time. If the Orientale basin formed 3.8 aeons ago and 

the Late Imbrian Epoch ended 3.2 aeons ago, formation of the 195 

Lower Imbrian craters of this size on the whole Moon required 0.05 

aeons, and that of the 174 Upper Imbrian craters required 0.6 aeons.
246 THE GEOLOGIC HISTORY OF THE MOON

These values translate into rates of 3,900 and 290 impacts per aeon, 

respectively. These rates are transitional between the 19,000 impacts 

per aeon of the Nectarian Period (1,330 craters in 0.07 aeon) and the 

much lower Eratosthenian and Copernican rates of 42 impacts per

aeon (133 craters in 3.2 aeons). Errors in the durations of the periods 

and epochs proposed here would change these values but would not 

alter the conclusion that the rate of impacts decreased sharply 

between 3.85 and 3.2 aeons ago.
fsfipil&Sf^ «J *^,~*&yt jm* *ti£**ifc «-' 1 * i- -J * ^^ M..«?tofc4V^%: -
FIGURE 12.1 (OVERLEAF). Eratosthenian crater and mare materials. L, Lambert (30 km), superposed on Imbrian mare units (small white-on-black arrowheads) and embayed by Eratosthenian 

mare units (large black-on-white arrowheads). Mare unit containing sinuous rille has a DL value of 220 ± 45 m (Boyce, 1976), floods southeastern secondary craters of Lambert, and covers older 

mare unit (three black-on-white arrows pointing southeastward). Older unit underlies secondaries of Eratosthenes (E). Therefore, stratigraphic relations of Lambert and Eratosthenes are similar; 

crater densities (fig. 12.4) distinguish Lambert (older) from Eratosthenes (younger). Secondary craters of Copernicus (C) are superposed on Lambert. Apollo 15 frame M-1009.
12. ERATOSTHENIAN SYSTEM

CONTENTS

Page

Introduction ...................................................... 249

Definition ........................................................ 249

Crater materials .................................................. 250 

Recognition criteria ............................................ 250

Frequency .................................................... 257

Mare materials general stratigraphy and distribution .............. 258

Mare Insularum .................................................. 259

Setting of the Apollo 12 landing site ............................. 259

Apollo 12 samples ............................................. 259

Summary and conclusions ...................................... 262

Chronology ...................................................... 262

INTRODUCTION

The two time-stratigraphic units described here and in the next 

chapter are less extensive than the four already described. The contrast in rates of formation is even greater: The Eratosthenian and 

Copernican Systems combined required three times longer to form 

than did the four previous units combined. The Eratosthenian and, to 

a lesser extent, the Copernican Systems continue the geologic style of 

the Upper Imbrian Series. Eratosthenian mare and crater materials 

are both present; the maria are about twice as extensive as the crater 

deposits (pi. 10).

Because the Eratosthenian and Copernican rock units are relatively young, they are little modified. Each crater deposit is visible 

farther from the rim crest than are older crater deposits, because the 

topography has been less eroded and the materials less completely 

mixed with the substrate. Eratosthenian and Copernican units serve 

as models for interpretations of the older units because of their 

freshness and clear stratigraphic relations (fig. 12.1).

In a sense, all the sampling missions obtained samples of Eratosthenian and Copernican material because they returned regolith 

material (Heiken, 1975; Pap ike and others, 1982). Many minor 

impacts have been dated either individually or collectively by determining the duration of exposure of regolith particles (Arvidson and 

others, 1975; Burnett and Woolum, 1977; Crozaz, 1977). However, 

Eratosthenian and Copernican bedrock units (those visible on 

regional photographs) were only sparsely sampled. Mare-basalt flows 

believed to be Eratosthenian were sampled only at the Apollo 12 

landing site. No large Eratosthenian craters were sampled, and large 

Copernican craters were sampled only indirectly (see chap. 13). 

Therefore, the reconstruction of post-Imbrian geologic history 

requires considerable photogeologic extrapolation a process not yet 

complete.

DEFINITION

Shoemaker and Hackman (1962) established the craters Eratosthenes and Copernicus as the type areas of the Moon's two youngest 

time-stratigraphic systems because of stratigraphic relations evident 

on telescopic photographs. Ejecta of Eratosthenes is superposed on 

the nearby mare material, whereas rays of Copernicus are superposed on Eratosthenes (figs. 1.1,1.6,12.2). Eratosthenes is much the 

darker of the two craters on full-Moon photographs (fig. 12.2B) and 

was thought to be rayless. Although one faint ray, in fact, extends

northwestward from the crater (formed by a long linear chain of 

secondary craters visible in fig. 11.3), Eratosthenes nevertheless does 

represent a class of postmare craters that essentially lack bright rays.

Accordingly, postmare craters have been operationally divided in 

most lunar geologic mapping into Eratosthenian and Copernican 

categories according to whether they are nonrayed or bright-rayed, 

respectively. Although ray visibility is neither exact nor correct as a 

criterion of age according to the stratigraphic code,12- 1 it has yielded 

many subsequently confirmed age estimates and is still useful where 

better data are unavailable.

Neither Shoemaker and Hackman (1962) nor later work have 

established the stratigraphic position of Eratosthenes and Copernicus within their respective systems. Wilhelms (1980) attempted to 

define or, at least, to better characterize the two systems on the basis 

of units in southern Mare Imbrium and northern Oceanus Procellarum. A stratigraphic sequence was established (fig. 7.1; table 

7.2), and the existing age assignments of as many units as possible 

were retained. The definitions were based on mare units because, in 

this region, they are more extensive and lithologically uniform than 

impact deposits. Because time-stratigraphic units are contiguous, the 

working criterion for mare units at the top of the Imbrian System, 

DL = 240 ± 10 m, also characterizes the base of the Eratosthenian 

System (table 12.1). Mare units having these or smaller DL values are 

post-Imbrian. Though not a formally acceptable criterion, the DL 

values are at least a secondary (superposed) criterion that is a function of time.

Although fewer crater frequencies than DL values are available 

for individual Eratosthenian units, the frequencies determined by 

Neukum and others (1975a, b) and Neukum and Konig (1976) probably represent the approximate range of Eratosthenian frequencies12 2. A plot of many of their determinations suggests that a crater 

frequency of 2.5 x 10 ~ 3 corresponds approximately to the borderline 

DL value of 240 m (fig. 7.15) and thus should characterize the mare

12 !The stratigraphic code (American Commission on Stratigraphic Nomenclature, 1970) was partly adhered to 

as follows. Rays were considered to be a physical characteristic by which roc^-stratigraphic units may be legitimately defined. Each ray deposit or ejecta deposit from which rays emanate was considered to be a formation and 

could be mapped by objective criteria, as should all formations. Then, the formations were assigned time-stratigraphic significance according to the physical property of the presence or absence of rays; all materials of postmare 

nonrayed craters were lumped on a map as Eratosthenian, and all materials of rayed craters as Copernican. Before 

Lunar Orbiter photographs were available in 1966-67, rayless but topographically fresh craters not in contact with 

mare units were commonly assigned an indefinite Imbrian or Eratosthenian age (map unit EIc). The fact that some 

crater ages were incorrectly estimated by the ray criterion illustrates the validity of the code's insistence on the 

separation of rock and time-rock categories of units.

12 - 2Except as noted, all crater frequencies given in this chapter refer to cumulative numbers of craters 1 km in 

diameter and larger per square kilometer. Neukum and his coworkers obtained these frequencies by normalizing 

their crater size-frequency curves at this 1-km intercept.

249
250 THE GEOLOGIC HISTORY OF THE MOON

TABLE 12.1. Properties of craters superposed on units of the Eratosthenian System

[Frequencies are those of craters at least 1 km in diameter per square kilometer, as determined by normalization to 1 km by means of a calibration curve (see fig. 7.10G; Neukum and others, 1975a);

n.a., not available]

Boundary Substrate Frequency

(m)

Basis Age 

(aeons) Basis

Upper--- Crater material 1(T3 Fiyure 12.4 1BO+20 Table 12.3 - - - 1.1 Chapter 13. 

Mare basalt- ----- n.a. - <165±25 Youngest flows traditionally considered Eratosthenian -- ---- - ----- ---- - ----- -.

Lower Crater material 1-2x10"^ Figure 12.4 >310 Table 12.3 3.2 Between Apollo 12 and 15 ages. 

Mare basalt-- -- 2.5x10" 3 Table 11.1 240+10 Consistent with traditional mapping of crater and mare units (Wilhelms, 1980) - -- -- --- -

materials at the base of the Eratosthenian System (table 12.1). Some 

flows in Oceanus Procellarum and Mare Imbrium with DL values 

smaller than 240 m were found to have frequencies of 1 to 2 x 10 ~ 3 

(figs. 12.3, 12.4A, E; Neukum and Kbnig, 1976; Gerhard Neukum, 

written commun., 1979). A further indication that a crater frequency 

of 2.5 x 10 ~ 3 and a DL value of 240 m are equivalent for the system's 

base is the finding (with a large uncertainty) of a crater frequency of 

2.6 x 10 ~ 3 on the upper Upper Imbrian mare unit(s) at the Apollo 15 

landing site, where DL = 270 ± 15 m (table 11.1; figs. 11.2,12.45, C).

The top of the system is harder to characterize. Wilhelms (1980) 

employed a DL value of 140 m as a working criterion for the Eratosthenian-Copernican boundary, based on values of 165 ± 25 m that were 

measured on the youngest mare units in Mare Imbrium traditionally 

called Eratosthenian (Boyce and others, 1975). However, smaller 

values may eventually be found. Similarly, the small frequencies 

found for some young flows (fig. 12.3) may not be the smallest that 

exist.

In summary, DL values for Eratosthenian mare materials range 

from less than 140 m to 240 ± 10 m, and the frequencies of craters at 

least 1 km in diameter per square kilometer (normalized) range as 

high as to 2.5 x 10 ~ 3 . The values for crater materials of the same age 

differ (table 12.1), as shown in the next section.

CRATER MATERIALS

Recognition criteria

The best criteria for an Eratosthenian age of a crater deposit are 

stratigraphic relations with the mare materials. Eratosthenian mare 

and crater deposits interfinger on much of the central and western 

nearside. Excellent examples of these relations are in Mare Imbrium 

(figs. 7.4, 12.1, 12.5; table 7.2) and Oceanus Procellarum (fig. 12.6). 

Extensive flows withDL values smaller than 230 m overlap Lambert 

and other craters that traditionally have been mapped as Eratosthenian (fig. 12.1; Wilhelms, 1980). However, the fresh bright craters 

Euler and Timocharis, previously considered to be Copernican, are 

also overlapped by such flows (fig. 7.4). These and other stratigraphically inconsistent assignments have been revised according to 

the philosophy of preserving a maximum number of existing assignments. The extensive mare materials remain Eratosthenian, and the 

few embayed bright craters become Eratosthenian (table 12.2).

Crater counts made marginally feasible by the best Lunar 

Orbiter and Apollo photographs (fig. 12.4; Neukum and Kbnig, 1976) 

can date favorably situated craters. Neukum and Kbnig (1976) found 

that the ages assigned to large rayed craters by the criteria of rays and

FIGURE 12.2. Superposition of Copernicus rays and secondary craters on Eratosthenes; 

arrows mark same point in each photograph. 

A. Low-Sun telescopic photograph. 

B. High-Sun telescopic photograph. Eratosthenes is almost invisible. White arrowhead

marks radial ray. CH, Copernicus H (compare fig. 13.6). 

C. Part of Orbiter 4 frame H-l 14.
12. ERATOSTHENIAN SYSTEM 251

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 >.'*' " >$ '. '" '"' *&* *' " &$p»"' '". '. "' ^>-

> 4 -j- .«»? <.-v,"i$r ,\<, ',',.. ,., , .,

 '$ v' 1 .»! $ ';^6Pfe. -V "'' 1 *^'/-'* i*"'

it» :-- ' ; ' *'*" "A. »4 Sh- , :. '' *.«$»* - "'' f^¥%> "V

FIGURE 12.3. Eralosthenian mare materials in Oceanus Procellarum, showing areas included in crater counts by Neukum and others (1975a, b) and 

Neukum and Konig (1976). Courtesy of Gerhard Neukum. 

A. Northern Procellarum. Scene centered at 26.5° N., 62.5° W.

B. Montes Harbinger, at border between northern Procellarum and Mare Imbrium. Crater-count area centered at 27.5° N., 41.5° W. 

C. Southern Procellarum, near crater Letronne (large rugged area at lower left). Scale bar at 10° S., 38.5° W., over island probably belonging to 

Flamsteed-Billy basin.
252 THE GEOLOGIC HISTORY OF THE MOON

morphology were consistently too young, probably because the rays 

and the sharpness and roughness of the coarse-textured crater subunits give the impression of youth. Langrenus (fig. 9.5), Theophilus 

(fig. 11.1), and several other large rayed craters are here accepted as 

Eratosthenian because their crater densities are similar to those on 

Lambert (fig. 12.1) and Bullialdus (fig. 12.7), which are sandwiched 

between Eratosthenian and young Imbrian mare materials.

Crater frequencies as high as 2 x 10 ~ 2 were found to be superposed on craters traditionally mapped as Eratosthenian (fig. 12.4). 

This and other frequencies for Eratosthenian craters are substantially larger than the frequencies of 2.5 x 10 ~ 3 superposed on mare 

materials at the base of the Eratosthenian System. The curve for the 

Apollo 15 units, just below the system's base, lies amidst the frequency plots for many typical Eratosthenian craters and even for

some Copernican craters (fig. 12.45, C). Different crater frequencies, 

therefore, seem to apply to mare and crater materials (Ahrens and 

Watt, 1980).

The 18-km-diameter crater Diophantus is a borderline case near 

the upper boundary of the Eratosthenian System. Diophantus has no 

visible rays (fig. 2.2), is younger than the youngest Eratosthenian 

flows shown in figure 2.2 by an unknown amount, and has not been 

flooded by younger mare units. It has been retained in the Eratosthenian System because that was its longstanding assignment 

(Wilhelms and McCauley, 1971) and because no reason for a change is 

apparent (Wilhelms, 1980). Diophantus is the youngest crater here 

considered to be Eratosthenian for which, to my knowledge, superposed crater frequencies have been accurately determined 

(1.25 ± 0.5 x 10 - 3 ; fig. 12.4D, E; Neukum and Konig, 1976). Thus, this

102

10'

CM

JC

510°

Z

UJ ^

CJ

LUtr.

Lx. 1

a IO 1

LU

^

o:

o

S -2 ^10 Z

 ^

^

^

0

ID'3

10-'

1 1 1 1 llj 1 1 1 1 1 1 1 II 1 III

r

' *

i \

: |4

-

r 'ft

1,1

(TAnl *

V

1 T

- 1* *

I ^ y^

: , |,t ^ US

Tf i,

4 f|| . A T' 

f Mbsting r J H

: Dawes f l <

I Aristarchus L , Stevinus j ,

Philolaus In 

Tycho A 

r Petavius B (total) '

: Petavius B (sub-area)

Copernicus 

; Delisle

Timocharis

Lambert

_ Oceanus Procellarum (total)

A

i i i nil J LJ in nl i i i i

iii| i i i i i ni| i i

-

.

-

-

-

-

~

-

-

-

-

_

t

OT "

 I Ji r

llrfi -

TffllL

1 -:

:

-

i Mil 1 i i i i nil i i

cT 1Cr1

10-2

I I I I I 111( I I

COPERNICAN CRATERS :~ 

e DAWES

« MOSTING

 LANGRENUS 

O THEOPHILUS 

A TIMOCHARIS 

A EULER

 COPERNICUS 

D ARISTARCHUS 

T TYCHO 

V CARPENTER

STEVINUS 

O PHILOLAUS 

« EUDOXUS 

O PETAVIUS B

APOLLO 15 LS

10- 10°

I I I I I I 111 I I I I I I 111 I III! 

\

ERATOSTHENIAN CRATERS

 LAMBERT

O ERATOSTHENES

A DELISLE

A DIOPHANTUS

 ARISTOTELES 

D BULLIALDUS 

T MANILIUS

APOLLO 15 LS

J iiil

10°

10 10' 10 IO1 

CRATER DIAMETER D(km)

Crater Diameter D (km)

102

o

Q

« S

| |

3 1

icr5 

io-4 

io-3 

io-2

m-1

Copernican Craters

~-

- Petavius B (C2)

:

Aristarchus (C2)

- Dawes (C1)

: Hosting (C1)

- Copernicus (CD

Timocharis (CD 

Eudoxus (C1) 

Theophilus (C1)

-

Tycho (C2)

Stevinus (CD

PhilolaUS (not class )

Carpenterinot class)

Euler (CD

Langrenus (CD

Eratosthenian Craters

Diophantus (E)

Eratosthenes (E)

Lambert (E)

Bullialdus (E)

Crater c assification reference : 

Geologic Map of the Near S de of the Moon, 

D.E.Wiihelms and J. F. Me Cauley, 1971

\D

Delisle (E)

Aristoteles (E)

Manilius (E)

younger

FIGURE 12.4. Crater-frequency studies by Neukum and Konig (1976) of large craters 

mapped as Eratosthenian and Copernican by Wilhelms and McCauley (1971). 

A. Plots for some craters relative to curve for mare materials of Oceanus Procellarum (total

of areas in fig. 12.3 and of other areas not shown). 

B. Plots for Copernican craters relative to curve for mare material at Apollo 15 landing site

(LS; compare fig. 11.2). Mare curve is intermediate in group of apparent ages; mare and

crater materials with same superposed-crater frequencies differ in age. 

C. Similar plots for craters considered to be Eratosthenian by Wilhelms and McCauley

(1971). 

D. Age assignments (E, CI, C2) by Wilhelms and McCauley (1971); ranking by Neukum

and Konig (1976), youngest at top. 

E. Same craters plotted by age (youngest at top) and size (horizontal axis), showing

probable bias to too-young age assignments of such large craters as Langrenus and

Theophilus.

10 3

CM

J

V __ , -4

z: 10

o

c

Q)

Jp"

it J _ 3

Jrr ' 10

"CD "

" 2

>

CO Q)

=3 0 -2

CZ =J 10

o ^

1

-

I

;

 ~-~^

i i i i

jPetavius B(C2)

'

^\

11111

Tycho(C2)

__

|C2| :

^~-^ }Aristarchus(C2)

-

Dawes(Cl)

I

- ^.^

^

Diophantus

IE)

--^^

[a] i ' *

"\

v JMosting(CI) Stevmus(CI)

M \.EulerlCH

~ Dehsle(E)' 'f' I / "~ 

'V JTimochanstCII

I

-

I

\Ei

JLambert(E) 

1

1 

|Manilms[E)

|

1 1 I 1

^"^-^^

.

.

^^ Oldest possible age I 

t- _ ,,,,. of young lava flows- Copermcus(CI) , n VtareM i m bnum -

^^ (Em)

-~Eratosthenes(E) ^^^X^

]^

Bullial

Eudoxus(CI)

Anstotelesf 

JuslE) - -'

E (and older)

i I I

1 :

\

T4_Qnqrenus(C1}~ 

r ^~^^ / ~

Average age of

old lava flows in

Mare Imbnum 

(Iml

~

1 l l l l -

1

^_

Q)

C

O

^

0 20 40 60 80 100 120 UO

CRATER DIAMETER D (km)
12. ERATOSTHENIAN SYSTEM 253

TABLE 12.2. Representative Eratosthenian craters

[Cross rules divide diameter ranges mapped differently in plate 10: smaller than 30 km, unmapped; 

30 to 59 km, interiors mapped; 60 km and larger, exterior deposits mapped. UI, possibly Upper 

Imbrian]

P f Di arneter Center 

Lrater (km) (lat) (long)

Taruntius H 

Bessarion B 

Hainzel AEratosthenes 

Bullialdus 

Aristoteles 

Theophilus 

Pythagoras 

3.5 

9 

12 

H 

15 

18 

19 

21 

23 

25 

26 

28

30

30 

32 

34

37 

39 

41 

42 

43 

43 

44 

53 

55 

58

61 

69 

70 

78 

87 

100

130 

132 

167

2° N. 

0.5° N. 

17° N. 

4° N. 

21° N. 

28° N. 

18° N. 

5° S. 

15° N. 

30° N. 

6° N. 

23° N.

26° N.

7° N. 

59° N. 

27° N.

35° N. 

15° N. 

6° S. 

31° S. 

15° N. 

3° N. 

4° N. 

41° S. 

15° S. 

15° N.

21° S. 

47° N. 

28° S. 

43° S. 

50° N. 

11° S.

64° N. 

9° S. 

66° S.

42° 

50° 

42° 

15° 

37° 

34° 

54° 

44° 

55° 

35° 

21° 

29°

21°

55° 

5° 

13°

163°

go

2° 

28° 

24° 

23° 

11° 

34° 

94° 

11°

22° 

39° 

3° 

42° 

17° 

26°

63° 

61° 

88°

W. 

E. 

W. 

E. 

W. 

W. 

E. 

W. 

E. 

U. 

E. 

U.

W.

U. 

E.

U.

E. 

E.

W. 

E. 

E.

W. 

E. 

W. 

W. 

W.

W. 

E. 

E. 

E. 

E. 

E.

U. 

E.

W.

Figure

12.6 

3. 28 

3.14A 

10.38 

12.5 

2.2 

5.23, q.ll 

12. 6A 

9.11 

2.2 

3.2C 

1.6, 3.29, 7.4

1.6, 6.

12, 

10. 

1.6,

12 

1.7, 

7, 

11.5, 

5. 

13, 

10, 

9. 

3.15A, 

1.6, 11,

12 

1.7, 

1.8, 3.26, 

8, 

1.7, 

11

1.6, 

9, 

3.

4, 12.1

.10 

.15 

11.3

.9 

10.16 

.7 

11.7 

17 

.10 

.16 

24 

, 4.4B 

.3, 12.2

.7 

9.3 

7.7, 9.27 

.5 

10.12 

.1

10.2 

.5 

2E

Remarks

Double impact.

Rim flooded. 

See figure 12.4; table 7.2.

Clear stratigraphy. 

See figure 12.4; table 7.2. 

See figure 12.4; table 7.2.

See figure 

12.3.

Asymmetric

12.4; tables 7.2,

See figure 12.4; tables 7.2, 

12.3. 

Farside example. 

See figure 12.4; table 12.3. 

See table 12.3.

See table 12.3. 

Do. 

Do. 

Do. 

Typical. 

Typical (see fig. 12.4; 

tables 7.2, 12.3).

See figure 12.4; table 12.3. 

Hare fill. 

Typical. 

UI? 

See figure 12.4. 

Rayed (see fig. 12.4; table 

12.3).

Rayed (see fig. 12.4). 

Largest young crater.

frequency approximates, but does not specify, the frequency of craters 

superposed on crater materials at the Eratosthenian-Copernican 

boundary. The total range for Eratosthenian crater deposits is, therefore, about 7.5 x 10 ~ 4 to 2 x 10 ~ 2 ; most midpoint values lie between 

10 4 and 10 - 3 craters.

DL values for crater deposits scatter widely (fig. 12.8; table 12.3). 

Including error bars, the range for Eratosthenian units is 130 to 330 

m; midpoints range from 160 to 310 m. Like the frequencies, therefore, the DL values apparently are displaced toward larger values for

TABLE 12.3. Properties of craters superposed on larger crater deposits

[Orbiter frames listed are those used for DL determinations and 

differ from those used for crater-frequency counts.

Ages: C, Copernican; E, Eratosthenian (pis. 10, 11).

Crater frequencies are from Neukum and Kflnig (1976) (see fig. 

12.4); they refer to craters at least 1 km in diameter per 

square kilometer, as determined by normalization to 1 km by 

means of a calibration curve (see fig. 7.10G; Neukum and others, 

1975a).

D|_ values determined by J.M. Boyce]

Tycho 

Aristarchus 

Kepler 

Copernicus 

Gassendi A 

Eudoxus 

Godin 

Aristillus 

Timochari s 

BUrg 

Autolycus 

Herschel 

Manilius 

Horrocks 

Bullialdus 

Plinius 

Cichus 

Reinhold 

Lambert 

Theophilus 

Eratosthenes 

Lexell 

Agrippa 

Hainzel A 

iameter

85

40

32

93

33

67

35

55

34

40

39

41

39

31

61

43

41

43

30

100

58

63

44

53

Orbiter

frame

5 H-128

5 H-201

5 H-138

5 H-146

5 H-180

4 H-98

4 H-97

4 H-103

4 H-114

4 H-86

4 H-103

4 H-108

4 H-97

4 H-96

4 H-125

4 H-85

4 H-124

4 H-125

4 H-126

4 H-77

5 H-134

4 H-107

4 H-97

4 H-136

C

C

C

C

C

C

C

C

E

C

C

E

E

E

E

E

E

E

E

E

E

E

E

E

Crater

frequency

3.7xlO'5

1.5x10"*

 

9.0x10"*

 

3.6xlO' 3

 

 

3.0xlO' 3

 

1.3xlO'2

7.0xlO' 3

 

 

 

4.8xlO" 3

4.6xlO' 3

2.7xlO' 3

 

___

D L

(m)

20+10

50+10

75+10

100+30

100+15

120+30

130+20

140+30

160+30

175+35

180±20

180+20

185±40

200±20

215+25

215+35

215+65

245+35

250+10

260+40

270+60

260+20

265+55

310+20

crater substrates than for mare substrates (compare total range of 140 

to 250 m, table 12.1).

Small craters are identified as Eratosthenian by the criteria of 

Trask (see chap. 7), which were also devised by observing the relations 

of craters to mare materials. Small Eratosthenian craters are buried 

in many places by mare flows on which other Eratosthenian craters 

are superposed (figs. 12.5,12.6).

In most farside and some limb regions, where photographs are 

poor, crater assignments to the Eratosthenian are still based on 

qualitative estimates of a low crater density, coupled with fresh 

morphology and an absence of detectable rays. The fine texture of 

radial ejecta and secondary craters is probably the morphologic feature most sensitive to degradation by small impacts and thus is most

~ '.*#!?.*' && i- t-*$* ^xm:^- C'^i'4fe>- *.'>? ; ^mm^ . ^IsjjfeseTV <iW£ . /v " ' ' i>H "

. '******%& *?&&'-'-' -ym'^i-,^. -^A, 

*»iv ",i *f-~<m& *v :.> :*i.

- - »* -" *v ,v-vy'-f*.-"i-<f» - -- '- ? * [" ' 

V» * ' , !M(. V ^51 : ,-. "J.m W 'r^^|^lftS;''-« '- » * *' *& ' » ... ,/^;*V~f: ;.;^'I"^% - ^~^?*1j': ^

FIGURE 12.5. Extensive Eratosthenian mare materials in southern Mare Imbrium, flooding rim materials of many small Eratosthenian craters, the largest of which is Brayley (B; 15km, 21°N., 

37° W.). Long chain of secondary craters extends northwestward of Brayley radial to Aristarchus, just outside left edge of photograph. Parts of Imbrium-basin rings are exposed as rugged terra 

islands. Mosaic of Apollo 17 frames M-2925, M-2928, and M-2931 (from right to left}
FIGURE 12.6. Part of Oceanus Procellarum near crater Flamsteed.

A. Large ring of rounded hills is Flamsteed P (112 km), a nearly buried pre-Imbrian (Nectarian or pre-Nectarian) crater. Eratosthenian crater Flamsteed (F; 

21 km) is superposed on Imbrian (Im) mare unit in south (white-on-black arrowheads) and embayed by younger Eratosthenian (Em) mare unit in north 

(black-on-white arrowheads). Other embayments of Eratosthenian or Imbrian craters are also shown by black-on-white arrowheads. Dashed lines indicate 

subtle contacts between Em and Im units; mare is closer to rims of small craters in Em than in Im, and crater density is lower in Em. S1, Surveyor 1 landing 

site; A5 (ellipse), potential Apollo landing site 5. Orbiter 4 frame H-143.
mf*. » f.^ ST > » -.'»,» "^T-- fc^" m-.;r"^erT«. - *, -f"^ ST«P^T--^--'7-''F-».T.'-j-;^ TTj'TT* '^iU-' l~i^ I, HW ' 7 ^ 1^ 

i ' "

^' . .v^...->.V--^->-' - . v'<*-'v< ^W^-^f*^*^^

y^^*ri"''*y.*|>iiiiB<jljg^^^^ y'^T^y^il

.<|^!|l^iI.|^!lil!TOpu a.fQ^';4pN^> »»gr^ !>'AllHJIr ^^^g^^'^^^^';1^?^^^^ 

IP*-' .,.; , ; ; »'^ ^ t ,^^?RPilP^^SI|P^^^^^

1 i'^77V-:?T;^g?/^r "'ir?"W' > J'fr^Af > "' y^t ^"yv- '.' ' » vet '^Vr«^r-.^Ji1 ^!^Ji^St:^^^l£&L^

6. Arrow indicates embayment of pair of otherwise fresh-appearing small craters also shown by black-on-white arrowhead above ellipse A5 in A. Largest crater is Maestlin G (3.5 km). 

Extensive level terrain, despite numerous small craters, indicates that no craters larger than 200 m in diameter are degraded to saucer shapes, as would be true on an Imbrian mare surface. 

Heavily cratered streak extending over most of right half of photograph is ray of crater Kepler (Carr and Titley, 1969X 200 km away in direction indicated by herringbone pattern. Orbiter 2 

frame M-207.
256 THE GEOLOGIC HISTORY OF THE MOON

FIGURE 12.7. Eratostheman crater Bullialdus (B; 61 km), superposed on Imbrian mare 

materials (white-on-black arrows) and embayed by Eratosthenian mare materials (black-onwhite arrows). Crater Lubiniezky (L; 44 km) is filled by thin mare unit younger than 

Bullialdus-secondary craters; embayment relations are traceable northwest and northeast of 

Lubiniezky. Orbiter 4 frame H-125.

diagnostic (fig. 12.9). Where neither this fine texture nor crater densities are visible, ray brightness is still used for dating. The inconsistencies in dating uncovered by Neukum and Kbnig (1976) illustrate 

that the morphologies of Late Imbrian, Eratosthenian, and early 

Copernican craters of subequal sizes differ only slightly when seen on 

any but the best photographs. Theophilus and Langrenus, for example, would still be considered Copernican if they were on parts of the 

farside not covered by Apollo photographs and if their rays were 

visible against the terra background.

Among the Eratosthenian craters dated by size-frequency studies is Cavalerius (fig. 12.10; D.B. Snyder, written commun., 1980). The 

peculiar bright feature in Oceanus Procellarum called Reiner gamma 

may be an exceptionally bright part of the otherwise-faint ray pattern 

of Cavalerius (fig. 12.10; Hood and others, 1979). Reiner gamma has 

been thought to be Copernican because of its brightness and superposition on other features (McCauley, 1967a). Interest in the origin of 

this enigmatic feature has been revived by the discovery that it has 

the highest magnetism yet observed from lunar orbit (Hood and 

others, 1979). The Eratosthenian age of Reiner gamma, if confirmed, 

would constrain speculations about the origin of lunar magnetism, 

one of the major unsolved problems raised by Apollo exploration. It 

also would illustrate a failure of the bright-ray criterion for distinguishing Eratosthenian and Copernican ages.

Some of the physical explanations first proposed for the implied 

fading of bright rays over time were solar radiation, cosmic-ray bombardment, and micrometeorite mixing with the darker substrate 

(Shoemaker, 1962b, p. 345; Shoemaker and Hackman, 1962, p. 299). 

Chapter 5 explained that the main darkening agent proved to be the 

accumulation of Fe- and Ti-rich agglutinates generated in the regolith 

by small impacts at the expense of brighter crystalline material. Rays 

that cross mare regoliths rich in these agglutinates are less conspicuous than those superposed on mare and terra materials that are 

poor in Ti and Fe.

The interiors of Eratosthenian craters smaller than about 8 km 

in diameter commonly are more nearly V-shaped than are both older 

and younger craters of the same size (figs. 10.41, 12.3, 12.5-12.7). 

Imbrian craters are shallower and commonly more nearly U-shaped 

(figs. 10.38,10.41). Many Copernican craters still display their original 

floor topography These differences result from time-dependent filling of crater bottoms by debris from the walls. The best Orbiter 

resolutions show that Eratosthenian craters smaller than about 3 km 

in diameter are smoother than their Copernican counterparts. The

300

270

240

210

 180 

E

150

120

90

60

30

\ I r i i i i r

Theophilus

Lambert

Manilius -

Bullialdus

i i i i i i i i i r

Eratosthenes

Timocharis

Eudoxus

Copernicus i

Arista re hus

[Tycho

10 10~2 10 -3

NUMBER OF CRATERS > 1 KM IN DIAMETER/KM 2

10 10-

FlGURE 12.8. Comparison of D^ values and frequencies of small (min I km diam) craters superposed on crater materials, based on figure 12.4 and table 12.3. Cc, Copernican craters; EC,

Eratosthenian craters.
12. ERATOSTHENIAN SYSTEM 257

oldest Eratosthenian craters disappear completely at sizes of about 

300 to 400 m, and the youngest at about 75 m (fig. 7.14).

In summary, crater units of the Eratosthenian System are best 

recognizable where fine-scale features are visible: morphologies and 

frequencies of small superposed craters, fine-scale smoothing by 

invisible craters in the steady state, and superpositional relations 

with mare flows. The brightness of a ray is a function not only of its 

age but also of the material in which the ray-forming secondary

craters were excavated. Better data might require revision of some of 

the age assignments given here (pi. 10), particularly on the farside 

and on the north, south, and east limbs of the nearside.

Frequency

Application of these qualitative criteria apparently has resulted 

in fairly consistent age assignments, judging by the even densities on 

the paleogeologic maps (pis. 10,11) of Eratosthenian and Copernican 

craters larger than 30 km in diameter. A total of 88 Eratosthenian 

craters of this size are recognized.

Wilhelms and McCauley (1971) mapped 256 Eratosthenian craters larger than 10 km in diameter within an area of 12.07 x 106 km2. 

In the same area, they mapped 145 Copernican craters of these sizes. 

Their criteria differed somewhat from those used here because 

brightness depends partly on substrate composition and because 

some sharp craters thought to be primary are actually secondaryimpact craters (Wilhelms, 1976). Nevertheless, the criteria were 

applied consistently, and no more complete compilation of craters of 

this size exists. Smaller craters were mapped at 1:1,000,000 and 

larger scales, but total numbers or frequencies have not been determined; such an exercise would suffer from greatly differing data 

quality, the effects of differing substrates, and subjective differences 

among mappers.

Wilhelms and others (1978) determined the frequencies of Eratosthenian and Copernican craters, combined, larger than 4.5 km in 

diameter (fig. 7.16); craters of these two systems were not distinguished because of the uneven quality of data in the area studied. 

The slope of the cumulative-frequency curve that best fits postmare 

craters 10 to 80 km in diameter is -1.8 (Baldwin, 1964; Shoemaker, 

1965; Neukum and others, 1975a; Basaltic Volcanism Study Project, 

1981, p. 1051-1052). Slopes of the curves for smaller craters depend on 

diameter range, as shown in figures 7.10G and 7.16.

FIGURE 12.9. Crater Stearns (S, 37 km) on lunar farside, identified as Eratosthenian by 

topographic sharpness.

A. Stearns inside box (B). Basin partly filled by mare is Moscoviense (M; 26° N., 147° 

E.). Large craters are Campbell (C; 225 km, probably pre-Nectarian) and D'Alembert (D; 225 km, probably Nectarian) Orbiter 5 frame M-85.

B. Freshness of Stearns' (S) interior is evident, despite oblique viewing angle and obscuration by shadow. Smoothed radial ejecta and secondary craters are also preserved 

(arrows), despite rugged terra substrate; periphery of Stearns is pitted by small craters. 

No exterior textures would be visible around an Imbrian crater of same size and setting; 

fewer pits and sharper texture would characterize Copernican craters of same size and 

setting. Orbiter 5 frame H 85.
258 THE GEOLOGIC HISTORY OF THE MOON

FIGURE 12.10. Reiner gamma (tadpole-shaped bright patch). Arrows indicate secondary craters of crater Cavalerius, centered 220 km away in direction of arrows. Conspicuous crater is Reiner 

(30 km), whose superposition on mare material indicates Eratosthenian or Late Imbrian age; pinch in ejecta pattern northeast and southwest of crater is probably due to oblique impact (see chap. 

3), but possibly is due to embayment by younger (Eratosthenian) mare. North arrow is adjacent to part of Marius Hills (see chap. 5). Orbiter 4 frame H-157.

MARE MATERIALS 

GENERAL STRATIGRAPHY

AND DISTRIBUTION

Eratosthenian mare units occupy two main settings central 

Mare Imbrium and a crescentic zone extending from Mare Frigoris to 

southern OceanusProcellarum (pi. 10; figs. 12.1,12.3,12.5-12.7,12.11). 

The Imbrium and Procellarum tracts join in a zone south of Montes 

Harbinger. Smaller patches of Eratosthenian basalt are scattered in 

the rest of Oceanus Procellarum and in other maria.

The massive, dark, spectrally blue flows in Mare Imbrium 

described in chapter 5 extend northeastward from the connection 

with Oceanus Procellarum near the edge of Mare Imbrium to a 

point north of the Imbrium center (fig. 5.1; Schaber, 1969; Schaber and 

others, 1975). Their stratigraphic relations are among the Moon's 

clearest because they flood small and large Eratosthenian craters and 

cover Imbrian mare units (figs. 12.1,12.5; Wilhelms, 1980). Before the 

Apollo missions, an Eratosthenian age was also established for mare 

units in a large area of equatorial Oceanus Procellarum by superpositional relations with small craters (fig. 12.6) and by evidence that the 

regolith there is thin. Thin regolith was indicated by the blocky 

surface revealed by Surveyor 1 at one point within this area (Shoemaker and Morris, 1970; Offield, 1972) 1970), and by an abundance of 

small blocky craters (Trask, 1969, 1971) and certain nested-crater 

configurations (Oberbeck and Quaide, 1968; Quaide and Oberbeck, 

1968, 1969) that are visible on high-resolution Lunar Orbiter photographs over extensive additional parts of the area. Several sites 

within the area were considered for Apollo landings because of the 

mare's youth (Carr and Titley, 1969; Titley and Trask, 1969; Cummings, 1971; West and Cannon, 1971; Offield, 1972). Many flows in 

northern Oceanus Procellarum are also young; some west of the 

Aristarchus Plateau are only about twice as old as the crater Copernicus (Young, 1977).

The most extensive Eratosthenian mare units of Mare Imbrium 

and Oceanus Procellarum are spectrally blue and, thus, probably rich 

in Fe and Ti (see chap. 5); they also are more radioactive than most 

other mare units (Soderblom and others, 1977). The most common 

spectral type is hDSA, although type HDSA is present in several 

spots in Oceanus Procellarum, including the Flamsteed-Flamsteed P 

region (pi. 4; fig. 12.6; table 5.1; Pieters, 1978; Pieters and others, 

1980). The spectral class (hDSA) and gamma-ray-spectrometer readings of the massive Imbrium flows indicate Ti contents in the range of 

about 2.0 to 4.5 weight percent TiO2 (Pieters, 1978; Davis, 1980).

Small patches of Eratosthenian flows, most spectrally blue but some 

red, are scattered throughout southern Oceanus Procellarum and 

Maria Insularum and Nubium (pi. 10; Whitford-Stark and Head, 

1980). Similar patches, including several with red spectra (Wilhelms, 

1980), are scattered between long 25 W and the central highlands 

"backbone" (fig. 1.8; see chap. 11).

FIGURE 12.11. Sinus Roris, showing flooding of Imbrian mare materials (Im) by dark, thin 

Eratosthenian mare materials (Em). Secondary craters of Eratosthenian crater Pythagoras 

(out of scene to upper left) overlie Imbrian units but are truncated by Eratosthenian unit (white 

arrow). Both mare units flood secondary craters of Iridum crater (black-and-white arrow). 

Large crater flooded by Imbrian units is Markov U (MU; 29 km, 52° N., 60° W.). Orbiter 

4 frame H-170.
12. ERATOSTHENIAN SYSTEM 259

The most abundant spectrally red Eratosthenian flows are on 

opposite sides of the ring of the Imbrium basin that is overlain by the 

Iridum crater and Plato. The flows on the Mare Imbrium side are 

beyond the margins of the spectrally blue Imbrium flows (compare 

pis. 4,10).

I suggest that the concentration of Eratosthenian as well as 

Imbrian eruptions in the Imbrium-Procellarum region is due to the 

thin lithosphere beneath the Procellarum basin (fig. 11.14). Conditions 

elsewhere also favored local late eruptions, however, judging by some 

patches of young basalt in Mare Smythii and other outlying maria 

(Boyce and Johnson, 1978). The factors controlling the distribution of 

spectral types are unclear. The main reason that high-Ti basalt dominates the Eratosthenian System may be its high radioactivity; the 

requisite heat for longlasting melting was greatest in this compositional type (Taylor, 1975,1982).

Many of the Eratosthenian flows occupy, or flowed from, the 

margins of several maria in addition to Imbrium (pi. 10; Boyce and 

Johnson, 1978). Small patches are perched high on the peripheries of 

Imbrium, Serenitatis, Crisium, and, probably, other basins (figs. 

5.10D, 5.16,5.25). Solomon and Head (1980) suggested that the changing lunar stress pattern resulting from global cooling and consequent 

lithospheric thickening favored closing of conduits everywhere except 

at mare margins. However, the thick existing sections of Imbrian 

basalt in the basin centers may have deflected the rising Eratosthenian magmas to the mare margins.

MARE INSULARUM

Setting of the Apollo 12 landing site

Apollo 12 landed in a topographically complex region, rich in 

islands of terra material, which at that time (November 1969) was 

considered part of Oceanus Procellarum but is now known as Mare 

Insularum (lat 3.2° S., long 23.4° W.; pi. 1; fig. 12.12; U.S. National 

Aeronautics and Space Administration, 1970; Marvin and others, 

1971). As mentioned above, the existence of the blue Eratosthenian 

mare material in the west-central equatorial belt of the nearside was

FIGURE 12.12. Area around Apollo 12 landing site (LS) whose size-frequency distribution 

was determined by Neukum and others (1975b). Two units (at least) were included in counts 

(fig. 11.2B); less densely cratered unit around landing site gives lower frequency plotted in 

figure 11 .2B. Fra Mauro B is 7 km across. Courtesy of Gerhard Neukum. Orbiter 4 frame 

H-125.

known when the site was selected. These maria were called "western," as opposed to the "eastern" maria east of long 25° W., which are 

mostly Imbrian. After the succesful Apollo 11 mission to an "eastern" 

site, geologists wanted to explore a "western" mare. However, the 

Apollo 12 landing site ended up being selected entirely for operational 

reasons. The desire was to demonstrate the ability of the Apollo 

system to land a LM at a predesignated point, and the Surveyor 3 

spacecraft, which had landed 2l/2 years earlier, provided an appropriate target. Surveyor 1 farther to the west offered a similar target, and 

its site was called Apollo landing site 6; but complex requirements of 

launch opportunity and other factors favored the Surveyor 3 site. As 

predictable from its position near the 25° W. dividing longitude, the 

geologic context of the site is less well known than that of the Eratosthenian Surveyor 1 mare (Offield, 1972).

The vicinity of the landing site and almost all the sampled 

terrain are occupied by the interiors or ejecta of several large craters 

(figs. 12.13-12.15). Most of this area is theoretically within the ejecta 

blanket of the largest of these, Middle Crescent Crater, 400 m across 

and originally about 80 m deep (Sutton and Schaber, 1971). Surveyor 

Crater (the 200-m crater in which Surveyor 3 landed) and Head 

Crater (110 m, the head of a crater configuration that resembles a 

snowman) also have strongly influenced the site. Samples were also 

taken from several spots on the ejecta of Bench Crater and near Halo 

Crater (fig. 12.14).

The overlapping ejecta blankets contribute to considerable 

uncertainty about the source strata of the basaltic samples and the 

significance of remotely sensed spectra. A spectrum centered over the 

landing site is of type mIG-, that is, intermediate in Ti, intermediate 

in albedo, and average or lacking in the two long-wavelength absorption bands (Pieters and McCord, 1976; Pieters, 1978). The TiO2 content of 2 to 3 percent suggested by this spectrum approximately 

matches that of the average regolith sample and falls between those of 

the two main types of sampled basalt. This averaging is expectable 

wherever surface deposits are composed of the ejecta blankets of 

craters that penetrate more than one unit. The crater frequencies and 

DL values that characterize the sampled material are also unclear 

because the few named craters dominate the local crater population. 

Soderblom and Lebofsky (1972) suggested that two units are present 

in the general vicinity an older unit at the site proper and a younger 

one about 1 km both west and east of the site. The unit at the site 

proper has a DL value of 210 to 215 m (table 11.1; Soderblom and 

Lebofsky, 1972; Boyce, 1976). TheZ^ value of the other unit is smaller 

by an undetermined amount (Soderblom and Lebofsky, 1972). These 

values fall in the Eratosthenian System as defined here. The midpoint of size-frequency counts in the unit at the site is 2.4xlO~ 3 

craters larger than 1 km in diameter per square kilometer (table 11.1, 

figs. 11.2,12.12). The relative and absolute ages are the youngest from 

any sampled mare site, including the Apollo 15 units, which lie near 

the top of the Imbrian System. Therefore, I consider the Apollo 12 

mare materials to be Eratosthenian.

Apollo 12 samples

A total of 36 rock-size basalt samples were returned by Apollo 12 

(Warner, 1971; James and Wright, 1972). Four compositional types of 

basalt have been recognized, starting with the early studies. They are 

now thought to compose two major and one minor unit at the site 

(Rhodes and others, 1977).

Olivine and pigeonite basalt. Two of the compositional 

groups olivine basalt characterized by modal olivine, and pigeonite 

basalt characterized by especially abundant modal pigeonite are 

probably parts of the same flow, which fractionated at the surface or in 

a shallow magma chamber (James and Wright, 1972; Rhodes and 

others, 1977). The pigeonite basalt has the finer textures, which 

indicate more rapid cooling, and so is presumably the upper fraction 

of the flow. The distribution of samples near the landing site is 

consistent with this stratigraphic relation (Rhodes and others, 1977). 

The olivine basalt apparently was excavated only by the largest crater, 

Middle Crescent (original depth, approx 80 m; fig. 12.15), whereas the 

smaller crater, Surveyor (original depth, approx 40 m), excavated the 

pigeonite basalt but not the olivine basalt.

The radiometric ages of samples of these two compositional 

groups have about the same broad range (table 12.4). Rb-Sr ages from 

a given laboratory favor an age 0.1 aeon older for the olivine than for 

the pigeonite basalt 3.25 versus 3.15 aeons (Papanastassiou and 

Wasserburg, 1970, 1971a), or 3.22 versus 3.12 aeons (Nyquist and
260 THE GEOLOGIC HISTORY OF THE MOON

FIGURE 12.13. Apollo 12 landing site and vicinity. Arrow marks location of Lunar Module in Surveyor Crater. Mare unit around landing site contains more large, subdued, old craters than unit

outside line and thus is the older unit. Orbiter 3 frame H 154.
12. ERATOSTHENIAN SYSTEM 261

others, 1977, 1979a). The fewer Ar-Ar ages cluster more tightly, 

however, and those determined by Turner (1971,1977) in each group 

are the same 3.21 aeons. In view of the overlapping analytical 

errors and the inconsistent discrepancies between the two methods 

(for example, sample 12065 is older by Ar-Ar, and sample 12002 older 

by Rb-Sr), the absolute ages may be considered consistent with the 

chemical, textural, and petrologic evidence for a common origin of the 

two compositional groups (James and Wright, 1972; Rhodes and others, 1977). Rhodes and others (1977) estimated a flow thickness of 30 

to 50 m. The average age for the combined compositional groups, from 

determinations whose stated error ranges are less than 0.1 aeon, is 

3.16 aeons, which is adopted here as the age of the flow.

Ilmenite basalt. The ilmenite basalt contains 2.7 to 5.5 weight 

percent TiO2, more than the other Apollo 12 groups, though less than 

many mare-basalt samples from other sites. Accordingly, it has been 

called either "low-Ti basalt" or "intermediate-Ti basalt" (table 5.2). 

Ilmenite basalt has been considered a minor component near the 

landing site, possibly introduced in crater ejecta from a nearby Tirich spectral unit approximately corresponding to the younger unit 

recognized by Soderblom and Lebofsky (1972). In this interpretation 

(Pieters and McCord, 1976), the olivine-pigeonite flow would be the 

local bedrock at the landing site. However, the restudy by Rhodes and 

others (1977) of the surficial distribution of recovered samples showed 

that ilmenite basalt is, in fact, quite common among the rock-size 

samples and is probably local to the site. The ilmenite basalt is 

apparently the only inplace basalt excavated by craters smaller than 

Surveyor, such as Head, whose original depth was about 20 m, 

whereas only the larger Surveyor and Middle Crescent Craters excavated inplace olivine and pigeonite basalt (fig. 12.15). Although the 

smaller craters also reejected some fragments of the olivine-pigeonite 

flow, these fragments resided in the ejecta of the larger craters.

FIGURE 12.14. Astronaut traverses at Apollo 12 landing site. Small circles, sampling 

localities; pennant, landing site; ALSEP, geophysical instruments. Airbrush map by U.S. 

Geological Survey.

Therefore, the ilmenite basalt flow overlies the olivine-pigeonite flow. 

On the basis of the likely original penetration depths of the craters 

(fig. 3.3), Rhodes and others (1977) estimated that the ilmenite basalt 

flow is about 40 m thick. Compositional data indicate derivation of 

this flow from a different magma from the parent magma of the 

olivine and pigeonite groups, which originated in a different mantle 

source (James and Wright, 1972; Rhodes and others, 1977). However, 

the ages for the ilmenite basalt listed in table 12.4, which average 3.15 

aeons from the best determinations, indicate no large gap between 

times of eruption of the two magmas.

Feldspathic basalt. Only one or, possibly, two rock-size samples 

of Al-rich, feldspathic basalt were recovered (12031?, 12038). The RbSr ages overlap those of the other groups. The average of the best RbSr age determinations (stated error ranges, less than 0.1 aeon) and 

the one young Ar-Ar age (3.08 aeons) is, again, 3.16 aeons. The 

petrologic relation of the feldspathic basalt to the other compositional 

groups is uncertain (Rhodes and others, 1977). The paucity of samples 

suggests a provenance in a distant unit or in a deeply buried unit at 

the landing site proper. The latter view (Nyquist and others, 1981b) is 

supported by collection of sample 12038 from the largest crater, Middle Crescent (by way of the superposed Bench Crater). Nyquist and 

others (1981b) favored derivation of this and other aluminous mare 

basalt from sources containing plagioclase in addition to the olivine 

and pyroxene generally thought to compose mare-source zones.

TABLE 12.4. Radiometric ages of large Eratosthenian basalt samples

[All ayes in aeons, recalculated using the radioactive-decay constants recommended by the International Union of 

Geological Sciences (IUGS) (Steiger and J3ger, 1977).

References: AD74, Alexander and Davis (1974; incorporates revisions of earlier determinations by the same laboratory [University of California, Berkeley], which are not given here); B81, Basaltic Volcanism Study Project 

(1981, p. 95D-952, table 7.3.1); C71, Cliff and others (1971); Co7l, Compston and others (1971), as modified 

by de Laeter and others (1973); H75, Horn and others (1975); Murthy and others (1971); N77, Nyquist (1977); 

N79a, Nyquist and others (1979a); N81b, Nyquist and others (1981b); PW70, Papanastassiou and Wasserburg 

(1970); PW71a, Papanastassiou and Wasserburg (1971a); PW71b, Papanastassiou and Uasserburg (1971b); S73, 

Stettler and others (1973); T71, Turner (1971); T77, Turner (1977)]

Group

Apollo 12 

feldspathic 

basalt.

Sample

'12031 

12038

Station

N. rim of Head ------ 

SW. of Bench 

Rb-Sr aye 

(reference)

3.13+0.03 (N81b), 

3.16+0.07 (N79a).

3.23+0.04 (B81),

Ar-Ar age 

(reference)

2 3. 08+0.05 (T77)

Sm-Nd age 

(reference)

3.28*0.23 (N81b)

3.28+0.09 (N81b).

Apollo 12 

ilmenite 

basalt.

Apollo 12 

pigeonite 

basalt.

Apollo \i 

olivine 

basalt.

12051

12063 

12017

12064 

12065

12002

1202C

S. rim of Surveyor -

Block (NE. rim of 

Surveyor).

NE. of Head 

SE. rim of Surveyor 

SE. rim of SurveyorSE. rim of Middle 

Cre cent.

SE. r m of Middle 

Cre cent.

3.09+.0.04 (N77), 

3.19±0.10 (PW70, PW71a).

3.23+0.13 (PW71a), 

3.27+0.10 (M71).

3.26+0.06 (PW71a).

3.13+0.03 (N81b).

3.13+0.03 (N81b). 

3.1HC.09 (PW71a) 

3.09+D.09 (PW71a)

3.29±0.10 (PW70, PW71a)

3.23+0.04 (PW71a).

3.13+0.06 (S73), 

3.21+0.05 (AD74), 

3.23±0.05 (T71).

3.15+0.05 (T77) 

3.16+0.07 (H75)

3.15+0.04 (H75) 

3.13+0.04 (AD74), 

3.21±0.05 (T71)

3.18±0.04 (AD74), 

3.21+0.05 (T71)

3.09+.0.04 (AD74)

'Nyquist and others 

Nyquist and others

(1981b) considered this sample to he piyeonite basalt.

(1981b) believed that this aye is too young because of disturbance of the argon system.

NW

MIDDLE CRESCENT CRATER HEAD CRATER

tr

LLJ 

UJ

S 10CH

LM SURVEYOR CRATER

Ilmenite basalt Pigeonite basalt 

Olivine basalt

^Ilmenite basalt

100 METERS

FRAGMENTAL UNITS 

i'.'-^ :.-'] Regolith

i^B Head Crater ejecta (ilmenite and pigeonite basalt fragments) 

IHIIHIIIIIII Surveyor Crater ejecta (ilmenite and pigeonite basalt fragments) 

INVSVNI Middle Crescent Crater ejecta (pigeonite and olivine basalt fragments)

FIGURE 12.15. Geologic cross section northwest-southeastward across area sampled by Apollo 12 (compare fig. 12.13). Interpretations after Sutton and Schaber (1971) and Rhodes and others

(1977). Substrate of Fe-Mg-rich basalt may be another mare unit or Fra Mauro Formation.
262 THE GEOLOGIC HISTORY OF THE MOON

Summary and conclusions

The best determined ages on Apollo 12 samples cluster between 

3.08 and 3.26 aeons and average 3.16 aeons. The stratigraphic 

sequence, from lowest to highest, is olivine basalt, pigeonite basalt, 

and ilmenite basalt. The bottom two compositional groups are parts 

of the same flow, and the ilmenite basalt is from a different flow. 

Absolute ages of these compositional groups or of the rare feldspathic 

basalt, whose provenance is uncertain, are indistinguishable.

Each mare-sampling site has increased our knowledge of the 

petrogenetic processes and source materials of lunar mare basalt 

after intensive and diligent study lasting more than a decade. After 

Apollo 11, there was talk of a single episode of mare extrusion; after 

Apollo 12, the number of episodes rose to two; and after each subsequent mission, the history of lunar volcanism appeared increasingly 

complex. We should, therefore, be chastened by the fact that at least 

two-thirds of the major mare spectral types remain unsampled 

(Pieters, 1978).

In particular, Pieters (1978) cautioned against drawing conclusions about the rather poorly denned spectral class (mIG-), to 

which the Apollo 12 basalt samples apparently belong. This class of 

material is widespread on the Moon (pi. 4) and contains many age and 

spectral subunits. Nevertheless, its general distribution may be of 

interest in attempts to discern first-order features of the lunar crust. 

It lies in areas of known thick crust outside the Procellarum basin 

(Crisium, Fecunditatis), as well as inside Procellarum in zones nearest the central terra "backbone." Eratosthenian basalt is also less

abundant there than elsewhere in Procellarum. Therefore, the 

lithosphere may have been thicker, and the crust or upper mantle 

different in composition, here than elsewhere in southern, western, 

and northern Procellarum, where Ti-rich basalt flows are abundant. 

The unusual north-south trend of the Eratosthenian patches and of 

such terra features as the "Fra Mauro peninsula" and the "central 

backbone" itself (fig. 1.8) may be additional evidence for a major 

structural trend along the central meridian of the nearside.

CHRONOLOGY

The Apollo 12 basalt flows, which lie near the base of the Eratosthenian System, average about 3.16 aeons in age. No younger or 

older Eratosthenian materials were sampled, and so absolute ages of 

other Eratosthenian units must be extrapolated from the dated samples by means of superposed-crater frequencies and DL values. The 

Eratosthenian Period began some time between 3.16 aeons ago and 

the 3.26-aeon formational time of the youngest Apollo 15 basalt samples. An approximate intermediate date of 3.2 aeons ago is adopted 

here.

The end of the Eratosthenian Period is harder to estimate. As 

discussed in detail in the next chapter, a date of 1.1 aeons is adopted on 

the assumption that the cratering rate has been constant since 3.2 

aeons ago. Thus, in this hypothesis, the Eratosthenian Period lasted 

2.1 aeons the longest of the six lunar time divisions adopted in this 

volume and nearly half the age of the Moon.
»i.*

TI N\/OINil3dOO liy]lSAS 
FIGURE 13.1 (OVERLEAF). Fresh, sharp-textured Copernican crater Aristarchus (40 km, 24° N., 47° W.) Orbiter 5 frame M-198.
13. COPERNICAN SYSTEM

CONTENTS

Introduction ....

Crater materials 

Mare materials . 

Structure .......

Chronology ....

Page

265

265

269

269

269

INTRODUCTION

Most early astronomic and geologic observers who carefully studied the Moon realized that rays are the youngest lunar features 

because they are superposed on all other terrains. The stratigraphically minded geologists Shoemaker and Hackman (1962) furthermore knew that the rays' youth indicates that the topographically 

expressed ejecta blankets of the source craters are also among the 

Moon's youngest. They accordingly established the Copernican System as the Moon's youngest assemblage of rock units.

No formal definition of the base of the Copernican System exists 

because no extensive stratigraphic-datum horizons exist near the 

lower system boundary. The rays of Copernicus are good early Copernican markers, but Copernicus does not mark the base of the system; 

many craters traditionally mapped as Copernican are older. The 

young, Eratosthenian mare lavas in Mare Imbrium and Oceanus 

Procellarum are near, but evidently not exactly at, the base of the 

system. Chapter 12 and table 12.1 give tentative crater-frequency and 

DL criteria for dividing Eratosthenian and Copernican units, but 

these criteria, also, are not definitive. This chapter further specifies 

the position of the system's base. The end of the Copernican Period is 

defined as the present; a crater formed today would be Copernican.

Copernican craters are sparsely scattered over the entire Moon; 

Copernican mare units are concentrated in northwestern Oceanus 

Procellarum (pi. 11). There are half as many craters larger than 30 km 

in diameter as in the Eratosthenian System (44 versus 88; pis. 10,11), 

and still fewer than in older time-stratigraphic units. The rays, 

however, make the Copernican System more conspicuous than the 

number of its units would indicate. Some Copernican deposits are 

beautifully fresh (figs. 13.1,13.2; table 13.1). They serve as analogs of 

more degraded features and hint at the hidden complexity of the older 

rock record. In other words, as in terrestrial geology, we can say that 

the present is the key to the past if the lunar "present" means 

hundreds of millions of years.

CRATER MATERIALS

Rays and extreme topographic freshness remain the principal 

criteria for the assignment of a Copernican age to craters (table 13.1). 

In favorable circumstances, ages are assigned more rigorously. 

Superposition of a given unit on the rays of Copernicus establishes its 

Copernican age (fig. 7.4; table 7.2). Crater counts can correlate a few 

units with Copernicus or with other bright-rayed craters that also are 

certainly Copernican (figs. 12.4,13.3; Neukum and Konig, 1976; Guinness and Arvidson, 1977; Young, 1977). Blocky ejecta and other features detectable at high resolutions that are used as age criteria by 

Trask (1969,1971) distinguish small Copernican craters (see chap. 7; 

figs. 7.12-7.14,13.2).

Problems in distinguishing Copernican from Eratosthenian craters arise when the craters are older than Copernicus, do not contact 

Copernicus, are faintly rayed, or are poorly photographed. DL values 

and crater frequencies (figs. 12.4,12.8; tables 12.1,12.3) can be determined only on superior photographs (fig. 13.4); even the best nearside

TABLE 13.1. Representative Copernican craters

[Cross rules divide diameter ranges mapped differently in plate 11: smaller than 30 km, unmapped; 

30 to 59 km, interiors mapped; 60 km and larger, exterior deposits mapped]

Crater Diameter 

(km)

Center 

(lat) (long) Figure

South Ray - 

North Ray 

Linne -- -- 

Copernicus H- -

Lansberg B 

MessierGambart A -

Goddard A --

Sulpicius Gallus -

Dawes- - 

Dionysius 

Li chtenberg - 

Pytheas 

Thebit A 

Conon 

Giordano BrunoMBsting 

Triesnecker 

Prod us0.8

1.0

2.5

4.3

9

11

12

12

12

18

18

20

20

20

22

22

25

26

28

9.2° S.

8.8° S.

28° N.

7° N.

3° S.

2° S.

1° N.

17° N.

20° N.

17° N.

3" N.

32° N.

21° N.

22° S.

22° N.

36° N.

1° S.

4° N.

16° N.

15.3° E.

15.5° E.

12° E.

18° W.

28° W.

48° E.

19° W.

90° E.

12° E.

26° E.

17° E.

68° W.

21° W.

5" W.

2° E.

103° E.

6° W.

4° E.

47° E.

9.9

9.9

3.2A

12. 2B, 13.6

10.41

3.11A

13.5

4.7

5.16

11.13

10.38

10.2, 13.7

3.4D

7.7

10.7, 10.14

9.28

7.7

6.20

3.33, 9.11, 9.15

See table 13.2.

See table 13.2.

Young.

Dark-haloed.

Elongate.

Very young.

See figure 12.4.

Young.

Rim flooded.

See table 7.2.

___

Very young.

See figure 12.4.

Young; asymmetric rays.

Kepler 32

Petavius B-- 

Godin-- 

Autolycus 

AristarchusOlbers A 

Crookes- - 

AnaxagorasAristillusTaruntius 

20° S.

2° N.

31° N.

24° N.

8° N. 

10° S. 

73° N. 

34° N.

6° N.

38° W.

57° E.

10° E.

2° E.

47° W.

78° W.

165° W.

10° W.

1° E.

47° E.

13.4 

10.16

1.6, 1.7, 2.5B 

5.12, 7.4, 13.1

10.3 

9.21

3.34, 10.15

1.6, 1.7, 2.5B

6.13B, 6.18, 9.3

Young (see tables 7.2,

12.3).

Young (see fig. 12.4). 

See table 12.3.

Do. 

Young (see figs. 12.4,

13.3; tables 7.2,

12.3). 

Young.

See table 12.3. 

Floor upl if ted.

Eudoxus- 67

King-- 

Tycho5° N. 121° E. 

43° S. 11° W.

Copernicus1.7, 10.12

1.2, 3.23, 3.32, 3.36 

1.1, 1.8, 2.3, 3.2D

1.1, 1.6, 3.4, 3.30, 

3.35, 7.4, 12.2, 13.10

See figure 12.4; table

12.3.

Farside example. 

Young (see figs. 12.4,

13.3; tables 12.3,

13.2. 

Typical (see figs. 12.4,

13.3; tables 7.2,

12.3, 13.2.

Lunar Orbiter 4 H-frames are barely adequate to date Copernican 

crater units. Because the DL method is valid only on level surfaces, the 

only crater materials it can date reliably are impact-melt pools, which 

are relatively small. The traditional criteria of ray brightness and 

topographic freshness are also hard to apply in poorly photographed 

areas, particularly on most of the farside poleward of lat 40° N. and lat 

40° S. (pi. 2). Some observational bias may account for the excess of 

nearside over farside craters mapped in plate 11 (26 versus 18) and the 

excess of northern- over southern-hemisphere craters on the nearside 

(17 versus 9). Nevertheless, the identification, during two separate 

iterations, of half as many Copernican as Eratosthenian craters in 

two size ranges (min 10 km diam [see chap. 12] and min 30 km diam 

[pis. 10,11]) suggests that the criteria have been consistently applied 

and correctly discriminate ages to a good approximation.

Copernican craters can commonly be identified by remote sensing. Temperatures measured in the infrared during a total lunar 

eclipse or from orbit are valuable for detecting young blocky craters 

and other fresh surfaces covered by little insulating fragmental material (see chap. 5, subsection entitled "Other Properties"). The eclipse 

infrared values were extensively used during the lunar geologicmapping program before the advent of spacecraft photography to 

distinguish between Copernican ("hot") and Eratosthenian ("cool")

265
266 THE GEOLOGIC HISTORY OF THE MOON

craters (fig. 13.5). Some of the crater-age assignments in this volume 

still rest on this distinction (for example, Agrippa and Godin, fig. 

10.16). Radar has also been used to detect blockiness in craters and 

thus to estimate cohesiveness and age-related state of degradation 

(Thompson and others, 1980, 1981). Color spectra may also reveal 

youthful surfaces in craters; enhanced reflectivity near 1 and 2 |xm 

indicates high proportions of fresh crystalline material to agglutinitic 

glass, whose formation is a function of exposure age (see chap. 5). 

Color contrasts seen on color-difference images also are strongest 

around the youngest craters (fig. 5.20).

Age interpretations of low albedo have changed for crater materials as they have for mare materials. The dark color of the volcanic 

maria and the existence of the irregular, dark-haloed endogenic craters in Alphonsus (fig. 5.10F) led to volcanic interpretations for the 

dark halos of certain circular craters as well (Shoemaker and 

Hackman, 1962, p. 297; Shoemaker, 1964; Salisbury and others, 1968). 

Impact origin was also entertained for these dark-haloed circular 

craters (Carr, 1965b) and was later substantiated by the impactlike 

morphology (deep floor, rough ejecta) of a typical example, Copernicus 

H (figs. 12.25,13.6). Most circular craters with impactlike morphology 

and dark ejecta are superposed on bright rays or other thin bright 

materials that, in turn, overlie maria or dark terra plains (see chap. 9; 

figs. 13.5, 13.6). Thus, the halos are only dark by contrast with their 

surroundings and probably contain basaltic materials brought to the 

surface by impacts (Scott and others, 1971, p. 276-277; Lucchitta, 

1972; Hodges, 1973a; Lucchitta and Schmitt, 1974; Wolfe and others, 

1975; Schultz and Spudis, 1979). Whereas the interiors of most endogenic craters are dark, the interiors of dark-haloed impact craters are 

as bright as those of any other impact crater. The dark-haloed craters 

superposed on Copernican rays are, of course, Copernican.

Other strong contrasts in albedo may also indicate youth because 

albedos become neutral with advancing age (see chap. 5). The rim of 

the young crater Tycho is surrounded by both bright and dark zones 

(fig. 1.1), which are unobserved around most craters. Dark and bright 

rays were artificially generated by impacts of spacecraft on the Moon 

(Whitaker, 1972a) and presumably also surrounded many natural

10C

o 

5 ic

(J

o

ic

10

(J

10-

T i ir i

EXPLANATION

nOceanus Procellarum 

o Aristarchus

 Copernicus

 Tycho

j__i

10 -2 10-1 10° 

CRATER DEPTH, IN KILOMETERS

10 1

FIGURE 13.3. Frequencies of craters superposed on three Copernican craters, in comparison 

with mare units in Oceanus Procellarum shown in figure 12.3 (compare fig. 12.4A). 

Courtesy of Gerhard Neukum.

FIGURE 13.2. Fresh Copernican crater, about 0.5 km in diameter, in central Mare Seremtatis. Blocks ejected from crater in direction of 

bright rays are preserved as far as three crater diameters from rim. In contrast, blocks are preserved only in interior and on rim crest of 

older, larger crater in upper left. Mare surface is brightened by Copernican rays except in sector of nondeposition left of crater 

(probably owing to oblique impact). Apollo 15 frame P-9337.
13. COPERNICAN SYSTEM 267

impacts, but most have been homogenized. Original albedos of target 

strata uplifted in the crater may be preserved in some young craters 

(fig. 3.29).

In summary, the brightest rays, most highly contrasting albedos 

of other crater materials, highest thermal anomalies, freshest morphologies, most coherent ejecta blocks, deepest floors, and fewest

superposed craters indicate a Copernican crater age. The relative 

degree of development of these attributes also may be used to subdivide the Copernican System (Trask, 1969, 1971; Wilhelms and 

McCauley, 1971). Mapping and dating of older craters is aided by 

reference to the appearance of Copernican primary and secondary 

craters.

B

FIGURE 13.4. Fresh Copernican crater Petavius B (33 km, 20° S., 57° E.), illustrating 

crater counts made on small areas from superior photographs. Counts made at each scale are 

combined and included in figure 12.4. Courtesy of Gerhard Neukum. 

A. Regional view. Ejecta is indistinct in north sector because it is superposed on mare

material and probably because of oblique primary impact. Orbiter 4 frame M 37. 

B. Near northern "excluded zone." Orbiter 5 frame H 37. 

C. Southeast sector. Orbiter 5 frame H 37.
268 THE GEOLOGIC HISTORY OF THE MOON

FIGURE 13.5. Small dark-haloed craters (arrows) superposed on rays of Copernican crater Gambart A (12 km), a thermal-infrared "hotspot." Mare material underlies Gambart A. Apollo 12

frame H-7737.

FIGURE 13.6. Dark-haloed crater Copernicus H (4.3 km; compare fig. I2.2B), showing subconcentric dunes characteristic of small

impact craters. Orbiter 5 frame H 147.
13. COPERNICAN SYSTEM 269

MARE MATERIALS

Northern Oceanus Procellarum contains some very young mare 

materials (pi. 11), the largest patch of which covers the southeastern 

ejecta of the bright-rayed crater Lichtenberg (fig. 13.7; Moore, 1967; 

Wilhelms, 1970b). Thus, the mare basalt is Copernican, and the 

Copernican Period had started by the time it was extruded. Three 

additional patches that I estimate to have the same or lower frequencies as this post-Lichtenberg unit are also mapped in plate 11. This 

part of northern Oceanus Procellarum was apparently the most 

active, or only, site of Copernican volcanism.

Dating of these patches would greatly help in determining the 

duration of lunar volcanism. The density of craters superposed on the 

patch near Lichtenberg has been determined to be about equal to that 

on Copernicus (P.H. Schultz and PD. Spudis, written commun., 1982). 

If, as discussed previously, a given impact superposes a larger crater 

on breccia than on basalt, the cumulative size-frequency curve for a 

crater deposit will be displaced upward from that for a mare deposit of 

the same age (see chap. 12; fig. 12.4; tables 7.1, 12.1; Schultz and 

Spencer, 1979; Ahrens and Watt, 1980). The flow that embays Lichtenberg would thus be younger than Copernicus.

The smallest DL value determined by Boyce and others (1975), 

150 ± 20 m, probably pertains to one of the Copernican patches. This 

value is somewhat smaller than 165 ±25 m, the DL value for the 

extensive neighboring Eratosthenian flows.

The small Copernican mare units are among the spectrally 

bluest on the color-difference photograph reproduced in figure 5.20. 

They are slightly bluer than the adjacent Eratosthenian units, whose 

spectral class is hDSA (pi. 4); the Copernican units are too small to 

show in plate 4 or on the spectral map of Pieters (1978). These units 

overlie the extensive Eratosthenian flows in the middle trough of the 

Procellarum basin (pi. 10). Other Copernican mare units may eventually be found in maria superposed on the central Procellarum 

basin, for example, Mare Imbrium; or young magmas melted there 

may not have been able to rise through the thick section of Imbrian 

and Eratosthenian basalt already present.

Otherwise, no Copernican volcanic materials are known. Darkmantling materials that were thought to be young on the basis of 

smoothness have been proved to be Imbrian in age (see chap. 11). 

Volcanism or emissions of gas were among the early-suggested explanations for Reiner gamma (fig. 12.10; McCauley, 1967b) and similar 

but more extensive bright "swirls" on the farside (fig. 4.7; El-Baz, 

1972). Some connection with lunar magnetism and an origin by 

Copernican (Schultz and Srnka, 1980) or pre-Copernican (Hood and 

others, 1979) impacts have been suggested more recently (see chap. 10, 

section entitled "Imbrium-Secondary Craters," and chap. 12, section 

entitled "Crater Materials").

STRUCTURE

Copernican tectonism was minor. Narrow gashes in the mare 

west of the Apollo 17 landing site (fig. 13.8) are young but probably 

resulted from drainage of regolith into older voids (B.K. Lucchitta, in 

Masursky and others, 1978, p. 209). The existence of mascons indicates that isostatic compensation of mare-filled basins has ceased or is 

very sluggish. Continuing subsidence and the resulting mare-ridge 

compression (see chap. 6) are unlikely because most endogenic moonquakes do not correlate with maria (Nakamura and others, 1979, fig. 

3; Solomon and Head, 1979). The only definite Copernican tectonic 

features are the uplifted floors of such craters as Taruntius (fig. 

6.13S). Binder (1982), however, suggested that young-appearing lobate scarps on the farside terra (Masursky and others, 1978, p. 96) are 

thrust faults caused by currently increasing global contraction.

Lunar seismic energy of endogenic origin is minuscule, only 

10 - 12 to 10 - 7 that of the Earth (French, 1977, p. 228; Lammlein, 1977, 

p. 266). The most energetic, but rarest, endogenic moonquakes probably originate in the upper mantle from small thermal stresses 

(Nakamura and others, 1979). The most common, but very small, 

moonquakes originate at depths of 1,000 ± 100 km from tidal stresses 

that do not deform the surface (Lammlein, 1977). The third type of 

moonquake, intermediate in frequency but unlimited in magnitude, 

is induced by impacts.

CHRONOLOGY

Establishing the chronology of the Copernican Period is hindered 

by the small number of well-dated stratigraphic units of regional 

extent. At 1.29 aeons old (Bernatowicz and others, 1978), sample 15405 

from the Apollo 15 landing site is the youngest rock-size lunar sample 

and the oldest that could be Copernican (fig. 13.9; table 13.2). This 

sample contains KREEP-rich "granitic" and "monzodioritic" material 

(Ryder and others, 1975a; Ryder, 1976) that is exotic to the landing 

site. A source in the Copernican craters Aristillus or Autolycus, 

centered 250 and 150 km north of the landing site, respectively, is 

considered possible by most investigators. Autolycus is marginally 

favored for several reasons: (1) It formed at least partly on the 

required target, KREEP-rich plains (Metzger and others, 1979), 

whereas mare basalt probably constitutes the uppermost target 

material of Aristillus; (2) it is closer than Aristillus to the collection 

site, easing somewhat the objection that the 1-m boulder from which 

sample 15405 was taken should have disintegrated in a long flight; 

and (3) it is older than Aristillus, making it the more likely source of 

1.29-aeon-old material if Aristillus is younger than Copernicus (as 

tentatively suggested by Guinness and Arvidson, 1977) and if Copernicus is about 0.8 aeon old.

Most estimates of the post-Imbrian cratering rate are based on 

an 0.8- to 0.85-aeon age for Copernicus that was determined on lightcolored KREEP-rich material dug from a shallow trench in the Apollo 

12 mare regolith (Hubbard and Gast, 1971; Hubbard and others, 1971; 

Marvin and others, 1971; Meyer and others, 1971). This material was 

identified with Copernicus by the situation of the landing site along a 

Copernicus ray (fig. 13.10; Pohn, 1971) and by the fact that part of the 

Copernicus target material was the KREEP-rich (see chap. 10) Fra 

Mauro Formation (Schmitt and others, 1967). U-Pb-Th systematics 

yielded a date of 0.85 ± 0.10 aeon for a thermal event affecting regolith 

fragments (Silver, 1971). Although U-Pb-Th techniques have proved to 

be ambiguous in many lunar stratigraphic applications, the age of 

this event seemed to be substantiated by Ar-Ar determinations of 

0.81 ± 0.04 aeon on sample 12033 (Eberhardt and others, 1973b; Alexander and others, 1976). Alexander and others (1977) obtained similar 

results for another KREEP-rich sample, 12032.

The identification of this approximately 0.8-aeon age with Copernicus has been doubted from several standpoints. Wasson and 

Baedecker (1972) doubted that primary ejecta would be preserved so 

far (340 km) from Copernicus. Quaide and others (1971, p. 715)

TABLE 13.2. Absolute ages of Copernican rock units

[Methods: Ar-Ar, the 40 Ar- 39 Ar method, applicable to old rocks (Turner, 1977). KrKr, dating of duration of exposure to the space environment (exposure age) by measuring the ratio of the radioactive isotope 81 Kr generated by cosmic-ray spallation 

to stable krypton isotopes, generally 83 Kr; applicable to relatively recent exposures (Marti, 1967; Arvidson and others, 1975). U-Th-Pb, applicable to major reequilibrations in old rocks; the entire system of U, Th, and Pb isotopes was examined by Silver (1971).

References: A76/77, Alexander and others (1976, 1977); B73, Behrmann and others 

(1973); B78, Bernatowicz and others (1978); C72, Crozaz and others (1972); D74, 

Drozd and others (1974); D77, Drozd and others (1977); E73b, Eberhardt and others 

(1973b); E77, Eugster and others (1977); LM72, lugmair and Marti (1972); M73, 

Marti and others (1973); S71, Silver (1971). Summarized by Arvidson and others 

(1975)]

Unit

Tycho (bright landslide) 

Autolycus or Aristillus

(ray).

Apol lo 

sampl ing 

site

16

17

14

16

17

12

15

Sample

do.

14306 (dark)

14306 (light)

14321, FM 1+2

14321, FMS

do.

Others

15405,90

Age

(m.y.)

2.D4+0.08

19

23.4+1.4

25.4±2.9

23.8+0.6

27.2+0.5 

48.9+1.7

50.3+0.8 

109+4

810+40

01 A+'iOO 810-50

850+100

1,290+40

Method

Kr-Kr

Kr-Kr

Kr-Kr

Kr-Kr

ICr \!r

Kr-Kr 

Kr-Kr

Ar-Ar

U-Th-Pb

Ar-Ar

Reference

074 

E77

C72

LM72

LM72

D74 

D77

A76/77

S71

B78
270 THE GEOLOGIC HISTORY OF THE MOON

observed that the Copernicus ray begins at a cluster of secondary 

craters only 45 km north of the landing site; the samples would thus 

include material excavated at that point (fig. 13.10). Alexander and 

others (1977) pointed out that non-KREEP particles yield the same 

ages as the KREEP-rich samples and questioned the identity of the 

high-temperature event. Finally, Alexander and others (1976,1977), 

though favoring the 0.8-aeon age as the more significant, showed that 

a plateau around 1.2 to 1.5 aeons is also quite well defined both in their 

data and in those of Eberhardt and others (1973b).

If the samples are not from Copernicus or if one of these older 

ages is that of the crater, most current estimates for the cratering rate 

would have to be revised. Earlier discussions in this volume suggest, 

however, that abandonment of the 0.8-aeon age would be premature. 

The argument that primary Copernicus ejecta is not present at the 

Apollo 12 landing site is based largely on the upper equation in figure 

10.26, which is highly questionable (see chap. 10, section entitled "Fra 

Mauro Formation"). Spectral studies show that primary ejecta is 

present in rays (see chap. 3; Pieters and others, 1982). Mixtures of 

diverse shock grades and compositions (KREEP-rich and KREEPpoor) are the rule rather than the exception in impact deposits (see 

chaps. 3, 8 10). Furthermore, the secondary impacts 45 km north of 

the site could have reexcavated the Fra Mauro Formation (fig. 13.10) 

and reset its age (although resetting by secondary impacts is doubted 

by most investigators). In summary, the presence of 0.8-aeon-old 

primary Copernicus ejecta or some other influence of Copernicus on 

the age of the analysed material is possible, despite the objections.

The best established date of a large Copernican crater is probably 

the one that appears to rest on the flimsiest evidence. The landslide at 

the Apollo 17 landing site was probably triggered by the impact of 

projectiles from crater Tycho, 2,250 km away (fig. 13.11; Wolfe and 

others, 1975; Arvidson and others, 1976; Lucchitta, 1977a). The duration of the landslide's exposure to cosmic rays and thus its time of 

formation has been accurately bracketed at 0.1 aeon (Arvidson and 

others, 1976; Drozd and others, 1977). Similar exposure ages of 

regolith materials in the central area of the landing site may date 

excavation by the "Central Cluster" craters (for example, Camelot and 

Sherlock, figs. 9.17,11.15), which are thought to be additional Tycho 

secondaries (Wolfe and others, 1975,1981; Lucchitta, 1977a). The age 

of Tycho is, therefore, widely accepted as about 100 million years and 

probably can be pinpointed to 109 million years (Drozd and others, 

1977).

The three best established Copernican ages (Arvidson and others, 1975) were obtained from the exposure ages of small craters (table 

13.1). Two young craters at the Apollo 16 landing site, North Ray and 

South Ray, are 50 million and 2 million years old, respectively. Cone 

Crater at the Apollo 14 landing site lies between the two, at about 25 

million years (fig. 13.12). These absolute age differences are consistent 

with the scheme of Trask (1969, 1971; Moore and others, 1980a). 

Several other craters have been dated by exposure ages, including 

Shorty at the Apollo 17 landing site (19 million years; Eugster and 

others, 1977) and, less directly, West at the Apollo 11 landing site (100 

million years; see chap. 11; Beaty and Albee, 1980).

FIGURE 13.7. Crater Lichtenberg (L; 20 km), northwestern Oceanus Procellarum. Blackand-white north arrow shows same point in A and B. 

A. Telescopic view suggesting sharp truncation of Lichtenberg materials by discrete dark

unit.

B. Lichtenberg ejecta flooded by mare basalt (white arrow) and thinly mantled by darkmantling material (d). Or biter 4 frame H-170.
13. COPERNICAN SYSTEM 271

Because no stratigraphic units older than 0.8 or 1.29 aeons and 

younger than the 3.16-aeon-old Apollo 12 mare units have been absolutely dated even tentatively, the ages of the many stratigraphic units 

that formed during this long gap in the record must be derived by 

interpolation based on the cratering-rate curve (fig. 13.13). However, 

the shape of this curve is poorly constrained because the known ages 

are so widely separated and the best Copernican ages are too young to 

affect it significantly.

Independent knowledge of the cratering rate would provide the 

necessary basis for interpolation. The terrestrial cratering rate is 

used to calibrate the impact rate in the Earth-Moon system, with 

corrections for the sampling problem caused by a geologically active 

Earth and for Earth-Moon differences in gravity and target properties (Baldwin, 1949, 1963, 1964, 1971; Opik, 1960; Shoemaker and 

others, 1962a, 1979; Hartmann, 1965b; Grieve and Dence, 1979; 

Grieve and Robertson, 1979; Basaltic Volcanism Study Project, 1981, 

chap. 8; Shoemaker, 1981). Some authors believe that the cratering 

rate has been constant since the Late Imbrian (Hartmann, 1972c; 

Soderblom and Lebofsky, 1972; Neukum and Konig, 1976, p. 2881; 

Guinness and Arvidson, 1977; Young, 1977). The shapes of the lunar 

and terrestrial curves suggest to others that the rate has declined 

during the past aeon (Trask, 1972; Soderblom and Boyce, 1972; Neukum and Konig, 1976, p. 2880; Shoemaker and others, 1979).

A curve corresponding to a constant rate since 3.2 aeons ago and 

fitted to a point between the Apollo 12 and 15 crater-frequency midpoints (curve a, fig. 13.13A) passes near, but not through, the bar 

representing the 0.8-aeon age and the frequencies of small craters 

superposed on Copernicus. The fit would be better if (1) the cratering 

rate has varied (curve b, fig. 13.13A), (2) the cratering rate was constant but has declined more steeply than shown because the crater

C. Detail. Apollo 15 frame P-0370.

FIGURE 13.8. Narrow gashes in surface of eastern Mare Serenitatis just west of Apollo 17 

landing site. Width on the order of tens of meters and apparent transection of rim of bright 

blocky crater (arrow) suggest a Copernican age. Gashes probably formed by drainage of 

regolith into larger bedrock grabens (B.K. Lucchitta, in Masursky and others, 1978, p. 

209). Apollo 17 frame P-2313.

FIGURE 13.9. Sample 15405, the youngest dated lunar rock, from Apennine front (sta. 

6A), Apollo 15 landing site.
272 THE GEOLOGIC HISTORY OF THE MOON

FIGURE 13.10. North-south strip from crater Copernicus to Apollo 12 landing site. 

A. Arrow indicates Apollo 12 landing site and points along ray toward Copernicus (C). 

Other craters include Reinhold (R; 43 km, 3° N., 23° W., probably Eralosthenian), 

and Lansberg (L; 39 km, 0°, 27° W., Upper Imbrian). Monies Carpatus (MC), part of 

Imbrium basin rim, divide Maria Imbrium (above) and Insularum (below). Obiter 4 

frame H-126.

frequencies determined for the Apollo 12 and 15 basalt units are too 

small (curve c, fig. 13.13A), (3) the crater-frequency values determined 

for Copernicus are too large, or (4) Copernicus is about 1.2 aeons old 

(point a', fig. 13.13A), one of the possibilities mentioned by Alexander 

and others (1976,1977). The Copernicus and Apollo 12 and 15 frequencies could be reconciled if substrate properties affect the size-frequency curves, as has been suggested. For example, if the frequency of 

small craters (normalized to min 1 km diam) superposed on Copernicus were about 6.5 x 10^ 4 per square kilometer, Copernicus would 

fall on the constant-rate curve, which applies to mare substrates, 

without any adjustment of the absolute age (point a", fig. 13.13A). The 

DL values could be reconciled in similar ways (fig. 13.135). Because of 

the many uncertainties in both the absolute and relative ages of 

Copernicus, the currently available data appear to be equally consistent with either a constant or a varying cratering rate since 3.2 aeons 

ago. I tentatively accept the simpler model, the constant rate.

Granted a constant rate, wide latitude still remains in assessing 

the age of the Eratosthenian-Copernican boundary because of the gap 

in the record and the uncertain significance of the crater frequencies 

andDL values. Because most assessments of the stratigraphic significance of the crater frequencies andDL values given in this volume are 

based on previous calibrations with the large craters traditionally 

mapped as Eratosthenian or Copernican, I return to this basis for 

estimating the duration of the two periods. If the craters larger than 

30 km in diameter mapped here (pis. 10,11) formed at a constant rate 

since 3.2 aeons ago, formation of the 88 Eratosthenian and 44 Copernican craters would require 2.13 and 1.07 aeons, respectively. Similarly, the 256 Eratosthenian and 145 Copernican craters larger than 

10 km in diameter mapped on the nearside by Wilhelms and 

McCauley (1971) (see chap. 12) would require 2.04 and 1.16 aeons, 

respectively. Averages of 2.1 and 1.1 aeons are tentatively adopted here 

for the Eratosthenian and Copernican Periods.

In summary, a 1.1-aeon duration for the Copernican Period is 

consistent with: (1) a constant impact rate for primary craters of all 

sizes since the beginning of the Eratosthenian Period 3.2 aeons ago, 

between the emplacement of the Apollo 12 and 15 basalt units; (2) 

existing age assignments of craters based on the criteria of stratigraphic superpositions, superposed craters, rays, morphology, and 

remote sensing; (3) a 0.8-aeon age for Copernicus; and (4) a position of 

Copernicus somewhat above the base of the Copernican System, in 

accord with stratigraphic estimates (fig. 7.1; table 7.2). Future data 

may demonstrate that the Copernican Period was longer or shorter 

than 1.1 aeons.

Apollo 12

MARE INSULARUM 

Am,Em

COPERNICUS 

Cc

MARE IMBRIUM 

N

200 KILOMETERS

B. Geologic cross section drawn along length of A from base of arrow to upper right corner. 

Units, from oldest to youngest: If, Fra Mauro Formation; Im, Imbrian mare basalt; Em, 

Eratosthenian mare basalt; Cc, Copernicus ejecta. Fra Mauro Formation composes 

submare section from Monies Carpatus lo Apollo 12 landing sile and ihus would be 

excavated by poslmare craters, including Copernicus and ils secondaries, anywhere 

along ihis profile.
13. COPERNICAN SYSTEM 273

FIGURE 13.11. Landslide (light-colored area at bottom) triggered from slope of South Massif, bordering the Taurus-Littrow Valley, by impact of projectiles from 

young Copernican crater Tycho, 2,200 km to southwest; arrow indicates probable Tycho secondaries (Lucchitta, 1 977a). Massif is about 2,000 m high. South 

at top. Apollo 15 frame P-9297.

10 3

10"

10&

10 6

1 ' 1 . i i i ,_

: NORTH RAY -

"

.

\

:\

- ^\

: 5

: \

\|

; 1

, ,

10

\

CRATER ;

T

\

\

!

:

-

-

i

:

-

--

\

\

-i i i i i

103

10"

10 6

in 6

' '\' ' ' '" 1

: 5 CONE CRATER:

^ \

I \

\

: \

\

\i

: \

; \

:

-

\ -

\T

\

i i i i i i i

100 3 10

\

\:

L \:

\

"

-

,

10 4

10 5

10 6

m 7

-'

I

-

V

_

-

-

I

:

-

"

-

30

I 1 1 1 1 1 1 1 1_

SOUTH RAY CRATER :

-

:

.

"

\

-1-

-

\ :

\ ;

-

10 100

DIAMETER (meters)

A

FIGURE 13.12.

B c

 Size-frequency distributions of very small craters superposed on the three

sampled and absolutely dated Copernican craters, North and South Ray Craters at Apollo

16 landing site and Cone Crater at Apollo 14 landing site (Moore and others, 1980a).

FIGURE 13.13. Alternative cratering rates since 3.2 aeons ago. Eratosthenian-Copernican 

boundary is estimated from numbers of large Eratosthenian and Copernican craters, assuming a constant cratering rate.

A. Based on frequencies of craters at least 1 km in diameter. Copernicus and Tycho 

frequencies from fig. 12.4, absolute ages from table 13.2. Mare-basalt frequencies from 

table 11.1, absolute ages from tables 11.3 and 12.4. Curve a, constant cratering rate 

since 3.2 aeons ago, based on crater frequencies on Apollo 12 and 15 mare units. Point 

a', intersection with 1.2-aeon age of Copernicus; point a", intersection with 0.81-aeon 

age of Copernicus (Eberhardt and others, 1973b; Alexander and others, 1976). Curve 

b, varying cratering rate since 3.2 aeons ago, passing through midpoint of Copernicus 

age. Curve c, constant cratering rate since 3.2 aeons ago, passing through midpoint of 

Copernicus age.

!5x10

3x10~3 -

£1x10

500

.55 4.0 3.0 2.0 

ABSOLUTE AGE, IN AEONS

1.0

B. Based on DL values. Copernicus and Tycho values from table 12.3 and Moore and 

others (1980b); mare-basalt values from table 11.1. Absolute ages as in A. Curve a, 

constant cratering rate since 3.2 aeons ago, based on average of Apollo 12 and 15 DL 

values. Curve b, varying cratering rate since 3.2 aeons ago, drawn through midpoint of 

Copernicus age. Curve c, one of many possible curves representing DL values on crater 

substrates (compare fig. 12.8).
FIGURE 14.1. Representative area of lunar surface (composite of several regions) and inferred subsurface structure. Mare basalt, gray in plan view and black in cross section, overlies impact melt in 

basins. Basin and crater ejecta blankets (white in cross section) overlap in several places, in accord with superposed-crater densities and degradational morphologies of surface exposures. 

Deformation of crustal material beneath basin rings is shown in accord with model in chapter 4; crust is thin, and mantle correspondingly uplifted, beneath basins. Stippling denotes possible 

unstratified lower crust that may have been reached only by the very largest impacts. True curvature; no vertical exaggeration. Painting by Donald E. Davis, courtesy of the artist.
14. SUMMARY

CONTENTS

	Page

Geologic style of the Moon ....................................... 276

Pre-Nectarian period ............................................. 277

Nectarian Period ................................................. 278

Early Imbrian Epoch .............................................. 279

Late Imbrian Epoch ............................................... 279

Eratosthenian Period .............................................. 279

Copernican Period ............................................... 280

Remaining problems .............................................. 280

275
276 THE GEOLOGIC HISTORY OF THE MOON

GEOLOGIC STYLE OF 

THE MOON

Two decades of study have shown that two major processes, 

impact and basaltic volcanism, have shaped the major physical features of the present lunar crust. Impact deposits and basalt presumably also constitute much of the subsurface. Tectonism has modified 

the original depositional geometry far less than on the Earth.

A typical section of the upper lunar crust consists of laterally 

continuous, interfingering beds of basin and crater materials that are 

pierced, deformed, and redistributed by later impacts (fig. 14.1). 

Deposits of ringed impact basins are the main constituents of this 

section. Crustal material originally created by extensive plutonism 

was redistributed far and wide over the surface, partly preserved and 

partly newly modified, by the great impacts. The upper crust was 

removed as far as each topographic basin rim. Ejected material, 

emplaced largely by turbulent surface flow, thickly mantled the surface outside the topographic basin rims for average distances of one 

basin radius. Beyond this distance, secondary-impact craters, as 

large as 2 to 4 percent of the basin size, were formed in abundance. 

Their ejecta incorporated some of the primary basin ejecta that 

formed them and greater amounts of the circumbasin terrain that 

they excavated. The primary-ejecta deposits contain much melted 

and incompletely melted target material, mixed with less highly 

shocked debris, whereas the secondary deposits contain little or no 

newly melted material except that from the primary ejecta which 

formed them. Plains deposits containing the same breccia types as 

other primary and secondary ejecta are concentrated in the primary/ 

secondary transition zone and among the secondary craters.

Basin interiors contain rings uplifted by deformation of the basin 

floor, knobby deposits ejected late in the excavation sequence, and 

impact-melt rocks. The rings increase in spacing, number, and complexity with increasing basin size. The outermost zones of large 

basins' interiors are especially complex. The knobby deposits also 

extend locally beyond the rim.

Smaller primary impacts have disturbed smaller amounts of the 

crust and redistributed it as lesser strata and fragmental regoliths. 

These ejecta deposits are proportional to those of basins in average 

extent; ejecta both of basins and of these smaller craters commonly 

have asymmetric, lobate or bow-tie map patterns. The secondary 

craters of basins and craters have similar size-frequency distributions, and the secondary-to-primary size ratio is only slightly smaller 

for the secondaries of basins. Primary-impact craters larger than 16 to 

21 km in diameter have central peaks consisting of severely deformed 

material derived during crater excavation from beneath the crater 

center. Unlike the chaotic ejecta of larger craters and basins, ejecta of 

small simple craters may preserve the bedding of the target, inverted 

from the original sequence.

Basaltic volcanism probably has altered this basic impact architecture throughout lunar history. Extrusions of Fe-rich, alkali-poor, 

dry, and reduced basaltic magmas generated in the mantle not only 

formed the visible maria but also were probably extruded earlier. 

Relatively silicic basalt, rich in trace elements characteristic of 

KREEP and originating in the crust, also may have contributed 

surface extrusions. Basaltic magmas that did not reach the surface 

must have solidified as extensive dikes and sills in the crust (fig. 14.1). 

Most lunar basalt has formed planar lava flows, some with lobate 

frontal scarps; minor eruptions have constructed positive landforms, 

such as domes and cones.

The position and volume of the extrusions were determined by 

the crustal and mantle structure established by the basins. The 

mantle rose beneath each basin and, probably, beneath each large 

crater to compensate for loss of the crustal overburden. Thus, the 

crust-mantle interface probably resembles a series of domelike 

swells, proportional in size to the basins and superposed where 

basins are superposed. Far more basalt was extruded into basin 

depressions through the thin crusts that cap the mantle uplifts than 

through the thick crusts outside the large basins. Except for the 

Apollo 11 suite from southern Mare Tranquillitatis (fig. 14.2), each 

suite of basalt samples returned by the Apollo and Luna missions 

apparently was formed by magmas extruded within a short time from 

a small, compositionally heterogeneous zone of the mantle.

Pyroclastic glass fountained as liquid droplets from fissures and 

irregular craters at the mare margins. These pyroclastic deposits are 

now observed as dark mantles on the nearby terra and mare lavas,

to £

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14. SUMMARY 277

and probably also inter-finger with the lavas. The total volume of 

extruded mare basalt and pyroclastic material amounts to only 0.05 

to 0.2 percent of the upper 75 km of the Moon.

Most lunar faulting was caused by the subsidence of mare basalt, 

which was greatest where basin excavations had thinned and weakened the elastic lithosphere. The lithosphere was probably equivalent 

to the petrologic terra crust during most of the volcanic and tectonic 

activity. Arcuate grabens were opened by stretching of the margins of 

the subsiding basalt sections. Mare surfaces were folded into complex 

archlike and spinelike ridges where the surface area was shortened; 

the shortening apparently was greatest where basalt sections of 

unequal thickness settled differentially. The degree of subsidence is 

also recorded in the gravity structure of maria; maria and basins that 

have been fully compensated isostatically are gravitationally neutral, 

whereas others appear as mascons that represent enduring superisostatic loads of basalt on the lithosphere. Crater floors were uplifted 

and fractured most extensively where the lithosphere was weakest. 

Otherwise, internal activity contributed very little to crater development except by generating the mare fillings of many craters and 

creating small vents for the pyroclastic eruptions.

In summary, geologic history at any given point on the Moon has 

advanced in a series of catastrophic impacts separated by much 

longer periods of gradual degradation and intermittent volcanism. 

Regoliths accumulated on each of the larger deposits in proportion to 

the duration of exposure and the impact flux at the time. Innumerable 

small impacts blurred depositional textures unremittingly, but subsurface geometry persisted until disturbed by large impacts. The 

ejecta from large impacts incorporated and mixed materials from the 

crater and basin units, the regoliths, the volcanic lavas, and subsurface intrusions. Lunar geology is simple in general style, though 

commonly very complex in detail.

This account pictures a uniformitarian Moon in which old features once resembled the more recent. Since crustal solidification, the 

face of the Moon has always displayed craters and ringed basins, and 

probably has always been spotted by dark maria. Major changes have 

occurred, however, in the proportions of units. The details of mare 

distribution constantly shifted as new basins appeared, were flooded, 

and were covered by later ejecta. Because volcanism continued after 

the impact rate had decreased sharply, more basalt remained exposed

late than early in lunar history (fig. 14.3). The rest of this chapter 

traces the changing scene from pre-Nectarian to Copernican time 

and concludes with a summary of unsolved problems.

PRE-NECTARIAN PERIOD

Lunar and pre-Nectarian history began with the formation of the 

Moon about 4.55 aeons ago. Similar oxygen-isotope ratios of the Earth 

and the Moon favor accretion in the same general part of the Solar 

System (Clayton and Mayeda, 1975). The process of accretion has not 

been agreed upon (Wood, 1979; Cadogan, 1981; Glass, 1982; Taylor, 

1982; Hartmann, 1983), despite an intense search for chemical clues 

in the lunar samples (Dreibus and others, 1977; Kaula, 1977; Anders, 

1978; Delano and Ringwood, 1978; Wanke and others, 1978). The 

accreted material is commonly assumed to have been chemically 

similar to that of chondritic meteorites, depleted in volatile elements.

The newly accreted material thoroughly differentiated. A feldspathic crust, now averaging about 75 km in thickness, and the upper 

several hundred kilometers of a denser, ultramafic, much more voluminous mantle are believed to have separated from an extensive early 

pre-Nectarian magma system. This magma is commonly considered 

to have formed a global ocean. However, recent modeling of crustal 

petrogenesis envisions a succession of bodies, which may or may not 

have been emplaced into the crystallization products of a primordial 

magma ocean. A chemically distinct core may or may not exist (Levin 

and Mayeva, 1977; Wiskerchen and Sonett, 1977; Taylor, 1982). A 

barrage of impacts must have deposited impact-melt pools, heated 

and mixed the environment of the crustal magmas to considerable 

depths, and generally influenced the igneous petrogenesis and the 

sites of magma migration. The differentiation ended at a time commonly estimated at from 4.4 to 4.2 aeons ago; it may have ended at 

different times in different places.

The cumulative effect of a heavy impact rate during pre-Nectarian time is evident in the highly brecciated, impact-melted, siderophile-rich material that constitutes the Nectarian and Imbrian 

deposits from which the returned samples were collected. The sample 

record of individual pre-Nectarian events that affected the solid terra 

crust, however, is less complete than the photogeologic record.

Mare basalt 

l Crater deposits 

Basin deposits

I

Obliterated crater 

and basin deposits 

(inferred)

100

4.55 4.0 3.0 2.0 

AGE, IN AEONS

FIGURE 14.3. Evolution of the lunar surface. Vertical axis, surface area inferred to have been originally covered by each group of deposits. Basin and crater deposits (approximately, the continuous 

ejecta) are assumed to extend one diameter from the crater or basin center; mare deposits of each age are assumed to extend beneath entire area of younger maria. Procellarum and South PoleAitken basins (pre-Nectarian basin group 1) are shown diagrammatically by narrow vertical bars; ages are estimated. Observed and inferred pre-Nectarian crater deposits cover more than 100 

percent of the Moon. Early peak in lunar geologic activity is evident.
278 THE GEOLOGIC HISTORY OF THE MOON

Mountainous rings of the giant Procellarum and South PoleAitken basins apparently constitute the earliest preserved record. 

Together, these two basins excavated crustal material from 40 percent 

of the Moon's area, and this material must have covered or impacted 

most of the Moon. This redistribution of material apparently resulted 

in terra surfaces more magnesian than average inside the basins and 

more aluminous outside; furthermore, KREEP-rich material of the 

deep crust may have been exposed in the deepest parts of Procellarum 

or of both basins. These newly exposed compositional provinces later 

became the targets of other basin impacts. The mantle does not seem 

to have been reached by impacts except, possibly, by the combined 

excavations of Procellarum and Imbrium. The crustal structure 

resulting from the giant impacts also set the stage for later mare 

volcanism and tectonism by determining local thicknesses of the 

lithosphere and the depths to the mare source regions in the underlying mantle.

Thirty pre-Nectarian basins, including Procellarum and South 

Pole-Aitken, have been identified or tentatively identified (table 8.2). 

They are ranked according to crater-density and superposition data 

into nine age groups. The oldest basins are obscure, whereas the 

youngest display textures and secondary craters similar to those of 

Nectarian and Imbrian basins. The spatial distribution of the preNectarian basins seems to be random (pi. 6).

The ages of these basins and of the oldest preserved surface 

features are only approximately known because few ancient units 

have been dated absolutely and because extrapolations of frequencies 

of craters that are saturated to large crater diameters are uncertain. 

If the impact rate was constant during the pre-Nectarian and Nectarian Periods and if the Nectarian Period lasted 0.07 aeon, about 28 

basins (groups 2-9) and 3,400 craters larger than 30 km in diameter 

formed from about 4.1 aeons to 3.92 aeons ago. This estimate is too 

young if the Nectarian Period was longer than 0.07 aeon, and too old if 

the early impact rate declined steeply as is, in fact, commonly 

assumed. The Procellarum and South Pole-Aitken basins are older 

than the group-2 basins by an unknown amount; to construct figure 

14.3,1 assume ages of 4.15 and 4.10 aeons, respectively, for these two 

giant basins. Many additional craters and basins that are now 

obscured also formed between the time of crustal solidification and 

the first group-2 basins.

Although volcanic materials are probably interbedded with the 

impact deposits, the pre-Nectarian volcanic record is much less evident than the rich impact record. The absence of observed tectonic 

deformation of pre-Nectarian units suggests that the crust had cooled 

before the oldest observed basins formed. The lithosphere presumably thickened throughout the pre-Nectarian.

NECTARIAN PERIOD

The Nectarian System continues the geologic style of the visually 

detectable part of the pre-Nectarian, but with fewer units (pi. 7). By 

definition, the Nectarian Period began with the impact that created 

the Nectaris basin and ended with the Imbrium-basin impact 3.85 

aeons ago. Geochronology and photogeology have suggested, but not 

definitely established, that feldspathic, melt-poor Nectaris-basin 

material was sampled and dated at the Apollo 16 landing site (fig. 

14.2). The age of Nectaris is tentatively estimated at 3.92 aeons on the 

basis of one interpretation of the ages determined on small samples of 

this material (table 14.1). If correct, this age implies that the Nectarian Period lasted 0.07 aeon and the pre-Nectarian 0.63 aeon.

From 10 to 12 basins formed during the Nectarian Period; the 

doubtful ones are Mendel-Rydberg, which may be pre-Nectarian, and 

Sikorsky-Rittenhouse, which may not actually be a basin. The Nectarian basins are ranked individually by age, but only two distinct 

age groups are recognized (table 9.4). If 11 basins are Nectarian and 

the Nectarian Period lasted 0.07 aeon, these basins formed at the rate 

of 157 per aeon. An extrapolation from the observed number of craters 

to the whole Moon indicates that 1,330 craters from 30 to 300 km in 

diameter formed during the Nectarian Period. Their formation rate 

was 19,000 per aeon, and 121 craters formed for each Nectarian basin. 

These Nectarian impact rates and crater-to-basin ratio are the basis 

for my estimates of pre-Nectarian ages and areal coverage (fig. 14.3).

At least one, possibly two, Nectarian basins in addition to Nectaris were sampled (fig. 14.2). Abundant fragment-laden impact melt 

was collected by Apollo 17 from massifs near the rim of the Serenitatis

TABLE 14.1. Adopted absolute ages of major sampled lunar geologic units

[Sampling sites: L, Luna; numbers only, Apollo]

System or 

series Unit or rock type Samplinc site

Age 

(aeons)

Copernican 

Erato s them' anUpper Imbrian-----

Lower Imbrian-- -

Nectarian - -

Nectarian and

pre-Nectarian.

Pre-Nectarian - 

Small craters- --- 

Tycho 

Copernicus - 

Basalt flows - -

Olivine basalt----- -

Pigeonite basalt- ----

Pyroclastic glass-- 

VLT basalt 

Feldspathic basalt 

High-K, high-Ti basalt 

Pyroclastic glass ---- 

Low-K, high-Ti basalt 

Thick basalt section-----

Low-K, high-Ti basalt- -

Low-K, high-Ti basalt -

Apennine Bench Formation (KREEP-rich

basalt).

Imbrium-basin ejecta 

Crisium basin--- - 

Serenitatis basin- - -

Nectaris basin ---

Volcanic-rock clasts-----

Plutonic-rock clasts-- 

14, 16, 17 

17

12

12

15

15

15

L24

L16

11

17

11

17

11

11

15

14, 15

L20

17

16

14, 17

15, 16, 17

<0.05 

.109

.81

3.16

3.26

3.30

3.30

3.30

3.40

3.57

3.64

3.67

3.70

3.72

3.79

3.84

3.85

3.85

3.84+0.04

 3 07

3 O /

3.92

3.85-3.98

4.17-4.54

Exposure ages-- 

do 

Average of best

Rb-Sr and Ar-Ar

ages for all

flows.

do 

do 

Ar-Ar age- - - 

do 

Rb-Sr and Sm-Nd

ages.

Averages of best

Ar-Ar ages for

two f 1 ows .

Average of best

Rb-Sr and Ar-Ar

ages.

Average of best

Ar-Ar ages.

Average of best

Ar-Ar ages.

Average of best

Rb-Sr, Sm-Nd,

and U-Th-Pb

ages.

Young cluster

(Apollo 14) and

black-and-white

breccia (Apollo

15).

Ar-Ar age 

Youngest Ar-Ar 

ages.

feldspathic

fragment-laden

impact-melt

rock from North

Ray Crater.

Individual ages -

Individual ages--- 

Table 13. 

Do.

Do.

Table 12.

Table 11.

Do.

Do.

Do.

Do.

Do.

Do. 

Do.

Do.

Do.

.__

Table 11.

Table 10.

Do.

Chapter 9 

Table 9.2

Table 9.1

Table 9.5

Table 8.4

2.

4.

3.

3.

1.

basin. Isotopically requilibrated clasts have been dated at 3.87 aeons. 

Although some doubts persist about the provenance of the best dated 

material, 3.87 aeons is probably the age of Serenitatis (table 14.1). 

Luna 20 returned small feldspathic fragments of probable Crisiumbasin material. Their determined age is a little younger than that of 

Serenitatis but not so precise.

These young isotopic ages, a concentration of other isotopic ages 

of terra samples between 3.85 and 3.95 aeons, the relative paucity of 

older ages, and the apparently old morphologic age of Serenitatis led 

to the suggestion that a "terminal impact cataclysm" formed most 

lunar basins. However, the large number of identified and inferred 

pre-Nectarian and Nectarian basins would create the same effects 

and would intensely rework the lunar crustal material, as is 

observed. Although the record does not require a terminal cataclysm, 

basins may have formed in groups because clusters of projectiles 

derived by breakup of a single large object temporarily maintained 

similar orbits after breakup (Wetherill, 1981); these projectiles may 

have impacted approximately along lunar paleoequators (Runcorn, 

1982).

An undetermined amount of volcanism occurred during the Nectarian and pre-Nectarian Periods. If the Nectarian surface was covered by mare basalt at the same average rate as was the Late Imbrian 

surface (approx 6.4 x 106 km2 in 0.6 aeon), 0.75 x 106 km2 of the lunar 

surface was covered from 3.92 to 3.85 aeons ago. This is enough basalt 

to cover a circular area about 295 km in diameter slightly smaller 

than the smallest basin in each of the 11 Nectarian basins. Large 

basins and all basins inside the giant basins would receive more than 

this average, other basins less. Extrapolation of this rate to the preNectarian depends on the source of heat for the volcanism. If most of 

it was leftover primordial heat (Hubbard and Minear, 1975), volcanism declined steadily since the crust solidified, and the pre-Nectarian basins would have been filled somewhat more extensively than 

the Nectarian basins. Failure to detect abundant pre-Nectarian volcanic geologic units and samples would then be explainable by the 

intensity of the early impact bombardment. If, alternatively, most of 

the heat came from radioactivity, premare volcanism may have been 

relatively uncommon (Taylor, 1982, p. 317), and the pre-Nectarian 

basins would contain relatively little pre-Nectarian basalt. The data 

are insufficient to choose definitively between these possible early 

rates of volcanism. Sparse data from mare-basalt clasts found in 

breccia deposits suggest that more Nectarian than pre-Nectarian
14. SUMMARY 279

basalt was formed (table 9.5). Furthermore, more craters with dark, 

Mg-rich ejecta, indicating burial of basalt flows by basin debris, are 

superposed on light-colored planar deposits of Nectarian than of preNectarian age. The buried basalt, however, may be pre-Nectarian.

EARLY IMBRIAN EPOCH

During the Early Imbrian, two large impacts left effects on the 

terra that still dominate its appearance (pi. 8). The creation of 

Imbrium, the Moon's third largest known basin (after Procellarum 

and South Pole-Aitken), uplifted rings that are the largest lunar 

mountains after those of South Pole-Aitken, and launched ejecta over 

most of the nearside and part of the farside. Thick primary ejecta was 

deposited within several hundred kilometers of the basin, and some 

probably reached more than 1,000 km from the rim. Extensive lightcolored plains consisting of primary and secondary ejecta were deposited beyond the primary ejecta. Secondary craters formed over vast 

outlying areas, possibly even at the antipode of the impact point. The 

KREEP-rich breccias recovered from the Fra Mauro Formation at the 

Apollo 14 landing site and lesser amounts of breccia from Montes 

Apenninus at the Apollo 15 landing site yield an average age of about 

3.85 aeons. The same age was obtained for fragments of KREEP-rich 

volcanic basalt or impact melt probably derived from a planar deposit 

near the Apollo 15 landing site (Apennine Bench Formation), which is 

contemporaneous with or slightly younger than the basin (tables 10.1, 

14.1). The Imbrian Period, therefore, began 3.85 aeons ago. The 

Imbrium basin was the site for the most voluminous preserved extrusions of lunar mare basalt.

The other large Early Imbrian basin, Orientale, formed about 50 

million years later on the west limb of the Moon and ended the era of 

enormous impacts. In accord with the younger relative age of the 

Orientale basin and its distance from later maria, its effects on the 

terra are even more striking than those of Imbrium. The Orientale 

ridged ejecta, impact melt, ejecta plains, secondary craters, and ring 

structure serve as models for the interpretation of older basin materials. A similar role at the small end of the size series of impact basins 

is played by the third Lower Imbrian basin, Schrodinger.

The Early Imbrian Epoch was too brief, and the effects of Orientale and of later mare extrusions too severe, for smaller impacts or 

volcanism to have left a conspicuous record. The oldest high-K, highTi flow excavated from the subsurface at the Apollo 11 landing site may 

be Early Imbrian (fig. 14.2; table 14.1). Orientale formed on a thick 

terra crust not previously thinned by a giant basin, and so, unlike 

Imbrium, it was not the site of abundant later basalt extrusions. 

Volcanism of terra composition may have formed some terralike 

plains and domical landforms during the Early Imbrian Epoch.

LATE IMBRIAN EPOCH

The record of the rest of lunar geologic history differs greatly 

from that of the 0.75 aeon already summarized. The cessation of giant 

impacts enabled the mare-basalt flows, which were continuously 

extruded, finally to remain preserved on the lunar surface. Previously, they had been catastrophically excavated and redistributed 

along with terra materials of the target or, if not disrupted, at least 

were blanketed by impact ejecta so as to loose their distinctive surface 

albedo. More visible mare units are Late Imbrian than any other age. 

Basalt of this age constitutes parts of all the larger maria (pi. 9).

Late Imbrian craters are also numerous and are surrounded by 

moderately well preserved ejecta blankets (pi. 9). Late Imbrian basalt 

is the Moon's most abundantly sampled and most reliably dated 

material (fig. 14.2;tables 11.3,14.1). Three age groups are mapped here 

(pi. 9).

Units of the oldest group, which formed within about 0.05 aeon, 

occupy elevated terrain in the outer Procellarum-basin shelf and 

terrain outside Procellarum, localities where they were protected 

from later volcanism. Their subsurface extent is probably much 

greater. Samples 3.79 aeons old were obtained from a buried unit at 

the Apollo 11 landing site. Units of the intermediate-age group, which 

span the interval from about 3.75 to 3.50 aeons ago, lie at the surface 

in all maria and are especially abundant outside Procellarum. They 

probably formed in all basins but remained uncovered where a thick 

lithosphere blocked ascent of magmas and prevented sufficient basin

subsidence to allow subsequent basalt extrusion. This and the older 

group of mare units commonly are faulted inside Procellarum but 

rarely outside, a further indication of relative subsidence. Apollos 11 

and 17 collected Ti-rich basalt samples of this group 3.57 to 3.70 aeons 

old and 3.72 aeons old, respectively. The sampled flows were extruded 

intermittently and sparsely in the shallow pre-Nectarian Tranquillitatis basin but much more rapidly in the Taurus-Littrow Valley 

marginal to the Nectarian Serenitatis basin. Pyroclastic dark-mantling materials fountained from the marginal zones of Serenitatis and 

other basins while the intermediate-age and, probably, the older 

groups of basalt units were forming; samples of such pyroclastic 

material collected by Apollo 17 are 3.64 aeons old.

Addition of the third and youngest group of mare basalt units 

further depressed the older units and the underlying basin, further 

fracturing the older mare and the adjacent basin material. These 

young basalt units appear in all maria, most abundantly in those 

whose total basalt section is thick because of a thin lithosphere. Units 

of this young group cover as much area as those of the other two 

groups combined, but are probably not proportionally voluminous 

because the older units underlie many or all of their exposures. If the 

rate of volcanism declined continuously during the Late Imbrian 

Epoch, the young group of basalt units is only one-seventh as voluminous as the intermediate-age group. Al-rich mare-basalt extrusions outside Procellarum are represented by fragments collected by 

Luna 16 from Mare Fecunditatis (3.40 aeons old) and very-low-Ti 

fragments collected by Luna 24 from Mare Crisium (3.30 aeons old). 

Abundant Ti-poor basalt samples from an older, relatively silicic flow 

and a younger, thinner, relatively mafic flow were collected by Apollo 

15 from a marginal mare of Imbrium (Palus Putredinis) and were 

dated at 3.26 to 3.30 aeons. Pyroclastic material also formed at that 

time. In the convention adopted in this volume, the Apollo 15 and 

Luna 24 units formed late in the Late Imbrian Epoch.

The Late Imbrian Epoch lasted about 0.6 aeon, from about 3.8 

aeons (Orientale impact) to 3.2 aeons ago (about midway between the 

older Apollo 15 and younger Apollo 12 absolute ages). Imbrian marebasalt units probably covered the entire area now covered by maria 

(including the mare area covered by younger craters), that is, about 17 

percent of the lunar surface, or 6.5 x 106 km2 .

During the same epoch, the continuous deposits of craters larger 

than 30 km in diameter covered about 1.7 x 106 km2 of the whole lunar 

surface, or about 25 percent of the mare area. Late Imbrian craters of 

this size formed at the rate of 280 per aeon, only about 1.5 percent of 

the Nectarian rate. If crater age assignments to the Lower Imbrian 

Series are correct (pi. 8), the Early Imbrian cratering rate (approx 

3,000 per aeon) is transitional between the Nectarian and Late 

Imbrian rates, closer to the latter. Thus, the crater-size impacts 

declined simultaneously with the basin-size impacts.

ERATOSTHENIAN PERIOD

Volcanism continued during the Eratosthenian Period, generating extensive flows in Oceanus Procellarum and Mare Imbrium and 

less extensive flows in other maria. These lavas are predominantly Tirich except in northern Mare Imbrium and Mare Frigoris. At least the 

Ti-rich types were derived from sources rich in radioactive elements 

and were extruded through thin lithospheres. Graben formation had 

ceased by the Eratosthenian Period except in uplifted crater floors, 

but thick sections of basalt and the underlying basin floors continued 

to sink sufficiently to crumple the mare surfaces into ridges. The only 

absolutely dated Eratosthenian basalt samples are those of four 

compositional types collected from at least three flows at the Apollo 12 

landing site. These rocks, which are 3.16 aeons old, are atypical of the 

mare units surrounding this site, most of which are Imbrian in age.

Interfingering of crater and mare materials is more obvious in 

the Eratosthenian System than in any other time-stratigraphic unit. 

Most Eratosthenian craters are rayless; some superposed on low-Fe, 

low-Ti mare materials and terrae are rayed. Exploration has confirmed an early hypothesis that a time-dependent exogenic process 

fades rays probably the manufacture by small impacts of Fe- and 

Ti-rich agglutinitic glasses at the expense of crystalline materials.

Ages younger than those of the Apollo 12 basalt samples and 

older than Copernicus must be interpolated by data on small superposed craters calibrated with models of the cratering rate. Available 

terrestrial and lunar data are consistent with a constant impact rate
280 THE GEOLOGIC HISTORY OF THE MOON

since the start of the Eratosthenian Period until the present, if mare 

and crater substrates of the same age display different crater frequencies. This constant rate would create 41 of the 132 observed 

Eratosthenian and Copernican craters larger than 30 km in diameter 

per aeon, a rate that is substantially lower than even the Late 

Imbrian rate and that signals a final leveling out of the early bombardment which shaped so much of the Moon's face. The Eratosthenian Period thus lasted about 2.1 aeons. The only impacting bodies 

available to create craters since the beginning of the Eratosthenian 

Period were similar to those still present in Earth-Moon space today 

(Shoemaker and others, 1979; Shoemaker, 1981).

COPERNICAN PERIOD

By the time the Coperican Period began, volcanism had yielded 

the stage almost entirely back to impact. Copernican mare units are 

known in only a few spots in the intermediate trough of the Procellarum basin that were also the favored sites of Eratosthenian 

volcanism. Later, further global cooling, depletion of radioactive-heat 

sources by extrusion of magma, and consequent lithospheric thickening and strengthening halted surface volcanism. The only possible 

melted zones today are below the zone of deep moonquakes, about 

1,000 ±100 km deep, in the zone where seismic shear waves are 

attenuated (Goins and others, 1979). Whatever core dynamo may have 

generated an early magnetic field (Runcorn, 1980) has been extinct for 

aeons.

The beginning of the Copernican Period has been accurately 

dated on neither the relative nor the absolute time scale. Absolute 

ages thought to date Copernican units include a questionable age of 

0.8 aeon for Copernicus, a probably accurate age of 0.1 aeon for Tycho, 

and accurate ages of 2 to 50 million years for small craters dated by 

Apollo 14 and 16 samples (tables 13.2, 14.1). An estimate for the 

duration of the Copernican Period, assuming a constant impact rate 

and a nonequivalence of the crater frequencies on basalt and crater 

substrates, is 1.1 aeon. If this calculation is correct, limited volcanism 

lasted until at least 1 aeon ago a late date by lunar, though not by 

terrestrial, standards.

Today, the Moon's surface is changing very slowly. Regoliths have 

formed during the past aeon at the sluggish rate of less than 1 mm per 

million years (Shoemaker, 1970; Quaide and Oberbeck, 1975). Downslope movement of debris is continuing to degrade slopes, but at a rate 

too slow to obliterate even such fine-scale features as the secondaryimpact craters of Tycho (fig. 3.6), which formed during the Earth's 

Mesozoic Era. Freshly exposed crystalline material is being very 

gradually converted to darker glasses. No Copernican units except a 

few small craters are being added to the Moon. The meteorite impacts 

recorded by the Apollo seismometers prove that impacts still do occur; 

one large enough to probe the deep interior of the Moon struck the 

farside in July 1972 (Nakamura and others, 1973). One to three 

craters at least 10 km in diameter may form every 10 million years 

(Shoemaker, 1981). Volcanism and almost all tectonism have ceased.

REMAINING PROBLEMS

This account has shown that much has been learned about the 

Moon during the past two decades. The origin of the maria and of 

most craters is settled. Key deposits exposed at the surface have been 

dated on both the relative and absolute time scales, and the antiquity 

of the Moon's face has been established. The composition of the Moon 

and many of its geologic units has been learned to a good first 

approximation.

It should also be clear that some key questions remain 

unanswered. Lunar remotely based studies and direct exploration, 

though in a lull, have not completed their job. The Moon's mode and 

place of origin are still unknown. Its subsurface structure is very

poorly known even in relatively well explored areas. Because 

Imbrium is the only basin that is well dated on both the relative and 

absolute time scales, the premare impact rate is too poorly calibrated 

before 3.85 aeons ago to establish such important points as the time of 

crustal solidification and the origin and lifetime of large Solar-System 

projectiles. The dates of volcanism before 3.8 and after 3.2 aeons ago 

are uncertain. The ages of most of the rayed craters are uncertain 

within broad limits. The relation among composition, source depth, 

and extrusion site of the mare basalt flows is hypothetical. The 

compositions of the farside maria are unknown. Terra compositions 

are known only very crudely from a few spot samples extrapolated 

from orbital measurements made at low resolutions of a small percentage of the surface. The origin of central peaks and shallow floors 

of complex craters is uncertain. The origin of basin rings and even of 

the position of the boundary of basin excavation central questions 

in studies of impact mechanics, lunar petrology, and stratigraphy is 

frustratingly elusive.

Some of these questions can probably be answered by continued 

experimental and field study on the Earth. The origin of complex 

craters and basin rings might yield to further study of terrestrial 

craters, laboratory and large-scale explosive experiments, and physical theory. Other questions, such as the distribution of mare-basalt 

compositions, can be partly answered by continued geologic mapping, 

crater-frequency counts, telescopic spectral studies, and petrologic 

theory. Still other questions may be answerable by continued examination of the 80 percent of lunar samples that have not yet been 

thoroughly analyzed.

Answers to most of the remaining geologic questions, however, 

require resumption of lunar spaceflights. A global orbiter could 

gather important data concerning: (1) mare compositions; (2) terra 

compositions (an even more serious gap from the point of view of the 

petrologist, the geochemist, and the cosmogonist attempting to learn 

the origin of the Moon and the Solar System); (3) the Moon's gravitational field; (4) the topography of basins, especially Orientale; (5) the 

puzzling problem of lunar magnetism; and (6) the stratigraphy of 

poorly photographed regions (particularly the polar regions above lat 

40°, a zone along long 100°-120° W., and the east limb on both the 

nearside and farside hemispheres). Knowledge of the Moon's third 

dimension could be greatly improved by this remote exploration. 

Mars is better photographed than the Moon.

Other problems will require additional samples from the Moon 

itself. Table 14.2 lists my recommendations for the landing sites of 

unmanned sample-returning spacecraft that could answer some of 

the key geologic questions. Objectives fall into five main categories: (1) 

absolute ages needed to calibrate the stratigraphic column, (2) compositions and textures to decide such genetic problems as the hypothetical terra volcanism, (3) crustal compositions at points of known 

stratigraphic context that can be extrapolated to larger areas, (4) 

mantle compositions inferrable from samples of currently unsampled 

color and age units of mare basalt, and (5) compositions and ages of 

premare volcanic basalt. Data from most of the selected targets can be 

extrapolated by means of currently available or future orbital sensing. Petrologists and geochemists might have a different list. The 

highest priority is given here to dating the Nectaris basin, whose 

relative age is well known and which, therefore, would provide the 

needed calibration for the premare cratering rate. Manned missions 

could achieve multiple goals at single landing sites but are not likely 

to be resumed within the foreseeable future.

Whether these programs should be undertaken is, of course, a 

matter of priorities. Very few spaceflights are undertaken by any 

nation or consortium of nations for purely scientific purposes. Moreover, some scientists give the Moon a low priority despite the easy 

accessibility of this primitive, cool, internally inactive, silicate planet 

unmodified by air or water that has preserved an unsurpassed record 

of the early Solar System. These recommendations are offered here in 

the belief that the Moon contains additional clues to our understanding of the Solar System, including our own home planet.
14. SUMMAKY

TABLE 14.2. Potential landing sites for future unmanned lunar sampling missions

[Each probe is considered capable of returning a single sample of regolith randomly selected from within the designated area. Contributions from Paul D. Spudis]

281

Priority Stratigraphic unit Landing area Figure Objective

Nectaris basin (a) Ejecta, near lat 35° S., long 42° E 9.1

(b) Plains (impact melt?), near lat 2?° S., long 41° E Q.2A

Copernican mare- Southeast of Lichtenberg, near lat 31° N., long 67° W 13.7

Terra plains (see also priority 11)- (a) Albategnius 1.8, 10.36

(b) Ptolemaeus l.B, 7.7

Terra domesFarside mare-

(b) Mare Ingenii 

Maunder Formation---- --- - -- South of Mare Orientale

(a) Gruithuisen gamma or delta- 10.45A

(b) Hansteen alpha 10.45B

(a) Floor of Tsiolkovskiy 1.3, 3.27A 

 . 1.4

 4.4B

 3.35

 3.32

 8.1

 9.21A

 7.6

(b) North of Balmer 1.5, 9.5

Early Late Imbrian mare ----- Mare Marginis, in Ibn Yunus-- ------ . 4.7

Eratosthenian mare--- - (a) Southwestern Mare Imbrium 5.1

(a) Surveyor 1 region, near lat 2.5° S., long 43.5° W - 12.6

Central Mare Serenitatis - Between Bessel and Dawes 11.13

Orientale lobate ejecta Near lat 53° S., long 79° W 4.4D

Alpes Formation Southeast of Vallis Alpes, near lat 45° N., long 5° E 10.12

Copernicus------------------- ---- Impact melt on floor ------- . .._--.-_..

King-- ----- - -- .. impact melt on floor- ----_-_-------.-

Ancient crust Near lat 30° N., long 160° E----------------

South Pole-Aitken massifs South of Korolev, lat 21.5° S., long 160° WPre-Late Imbrian mare(?) basalt- (a) Center of SchickardApennine Bench Formation- Near lat 27° N., long 8° WFissured crater-floor deposits- Floor of Murchison, lat 1° W., long 5° N6.10

10.28

(a) Absolute age.

(b) Composition.

(a) Absolute age.

(b) Composition of source.

Nonmare volcanism or buried mare basalt flows(?), 

Nonmare volcanism(?). 

Composition of source.

(a) Age of Orientale basin.

(b) Composition of thick crust.

(a) Absolute age.

(b) Deep crustal composition(?).

(a) Farside crustal composition.

(b) Absolute age.

(a) Composition.

(b) Absolute age.

(a) Farside crustal composition.

(b) Absolute age.

(a) Absolute age.

(b) Composition.

(a) Absol ute age.

(b) Composition (KREEP-rich?).

(a) Absol ute age.

(b) Calibration of color spectra.

(a) Calibration of color spectra (standard spectrum).

(b) Absolute age (near Imbrian-Eratosthenian boundary).

Impact melt or other ejecta.

(a) Impact-melt/debris content.

(b) Composition of deep ejecta.

(a) Impact melt or KREEP-rich volcanic materials(?).

(b) Absolute age.

(c) Calibration of gamma-ray data.

Ejected impact melt (Imbrium?).

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USGS Map 1-599 (RLC-14), scale 1:250,000.

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 1970, Geologic map of the Maskelyne DA region of the 

Moon, Lunar Orbiter site II P-2, southwestern Mare Tranquillitatis, including Apollo landing site 1: USGS Map 1-616 

[ORB 11-2(100)], scale 1:100,000.

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rilles and martian channels: Icarus, v. 22, no 1, p. 1-23.

Carr, M.H., and Meyer, C.E., 1974, The regolith at the Apollo 15 

site and its stratigraphic implications: Geochimica et Cosmochimica Acta, v. 38, no. 7, p. 1183-1197.

Carr, M.H., and Titley, S.R., 1969, Geologic map of the Maestlin G 

region of the Moon, Lunar Orbiter site II P-13, Oceanus 

Procellarum, including Apollo landing site 5: USGS Map 

1-622 [ORB 11-13 (100)], scale 1:100,000.

Carr, M.H., Howard, K.A., and El-Baz, Farouk, 1971, Geologic 

maps of the Apennine-Hadley region of the Moon; Apollo 15 

pre-mission maps: USGS Map 1-723, scales 1:250,000, 

1:50,000, 2 sheets.

Casella, C.J., 1976, Evolution of the lunar fracture network: 

Geological Society of America Bulletin, v. 87, no. 2, p. 

226-234.

Chabai, A.J., 1977, Influence of gravitational fields and 

atmospheric pressures on scaling of explosion craters, in 

Roddy, D.J., Pepin, R.O., and Merrill, R.B., eds., Impact and 

explosion cratering: Planetary and terrestrial implications: 

New York, Pergamon, p. 1191-1214.

Chao, E.C.-T., 1967, Shock effects in certain rock-forming minerals: Science, v. 156, no. 3772, p. 192-202.

 1973, Geologic implications of the Apollo 14 Fra Mauro 

breccias and comparison with ejecta from the Ries crater, 

Germany: USGS Journal of Research, v. 1, no. 1, p. 1-18.

 1974, Impact cratering models and their application to 

lunar studies a geologist's view: LSCP 5, v. 1, p. 35 52. 

-1977, The Ries crater of southern Germany, a model for

large basins on planetary surfaces: Geologisches Jahrbuch, 

ser. A, no. 43, 85 p.

Chao, E.C.-T, and Minkin, J.A., 1977, Impact cratering phenomenon for the Ries multiring structure based on constraints of geological, geophysical, and petrological studies 

and the nature of the impacting body, in Roddy, D. J., Pepin, 

R.O., and Merrill, R.B., eds., Impact and explosion cratering: 

Planetary and terrestrial implications: New York, Pergamon, 

p. 405-424.

Chao, E.C.-T., Hodges, C.A., Boyce, J.M., and Soderblom, L.A., 

1975, Origin of lunar light plains: USGS Journal of Research, 

v. 3, no. 4, p. 379-392.

Chao, E.C.-T, Minkin, J.A., and Best, J.B., 1972, Apollo 14 breccias: General characteristics and classification: LSCP 3, v. 1, 

p. 645-659.

Chao, E.C.-T, Minkin, J.A., and Thompson, C.L., 1976, The 

petrology of 77215, a noritic impact ejecta breccia: LSCP 7, v. 

2, p. 2287-2308.

Chao, E.C.-T, Shoemaker, E.M., and Madsen, B.M., 1960, First 

natural occurrence of coesite: Science, v. 132, no. 3421, p. 

220-222.

Chapman, C.R., and Haeffner, R.R., 1967, A critique of methods 

for analysis of the diameter-frequency relation for craters 

with special application to the Moon: Journal of Geophysical 

Research, v. 72, no. 2, p. 549-577.

Chapman, C.R., Aubele, J.C., Roberts, W.J., and Cutts, J.A., 1979, 

Sub-kilometer lunar craters: Origins, ages, processes of degradation, and implications for mare basalt petrogenesis 

tabs.]: LPSX, pt. 1, p. 190-191.

Charette, M.P, McCord, T.B., Pieters, C.M., and Adams, J.B., 

1974, Application of remote spectral reflectance measurements to lunar geology classification and determination of 

titanium content of lunar soils: Journal of Geophysical 

Research, v. 79, no. 11, p. 1605-1613.

Cintala, M.J., Head, J.W., III, and Veverka, Joseph, 1978, Characteristics of the cratering process on small satellites and 

asteroids: LPSCP 9, v. 3, p. 3803-3830.

Cintala, M.J., Wood, C.A., and Head, J.W., 1977, The effects of 

target characteristics on fresh crater morphology: Preliminary results for the moon and Mercury: LSCP 8, v. 3, p. 

3409-3425.

Clark, RE., and Hawke, B.R., 1981, Compositional variation in 

the Hadley Apennine region: LPSCP 12, pt. B, sec. 1, p. 

727-749.

Clayton, R.N., andMayeda, T.K., 1975, Genetic relations between 

the moon and meteorites: LSCP 6, v. 2, p. 1761-1769.

Cliff, R.A., Lee-Hu, C., and Wetherill, G.W., 1971, Rb-Sr and U, 

Th-Pb measurements on Apollo 12 materials: LSCP 2, v. 2, p. 

1493-1502.

Colton, G.W., Howard, K.A., and Moore, H.J., 1972, Mare ridges 

and arches in southern Oceanus Procellarum, pt. S of Photogeology, [chap.] 29 of Apollo 16 preliminary science report: 

NASA Report SP-315, p. 29-90 to 29-93.

Compston, W., Berry, H., Vernon, M.J., Chappell, B.W., and Kaye, 

M., 1971, Rubidium-strontium chronology and chemistry of 

lunar material from the Ocean of Storms: LSCP 2, v. 2, p. 

1471-1485.

Compston, W., Foster, J.J., and Gray, C.M., 1975, Rb-Sr ages of 

clasts within boulder 1, station 2, Apollo 17: The Moon, v. 14, 

no. 3-4, p. 445-462.

Compston, W., Vernon, M.J, Berry, H., Rudowski, R., Gray, C.M., 

Ware, N.G, Chappell, B.W, and Kaye, M.J, 1972, Apollo 14 

mineral ages and the thermal history of the Fra Mauro Formation: LSCP 3, v. 2, p. 1487-1501.

Conca, James, and Hubbard, N.J., 1979, Evidence for early volcanism in Mare Smythii: LPSCP 10, v. 2, p. 1727-1737.

Conel, J.E., and Nash, D.B., 1970, Spectral reflectance and 

albedo of Apollo 11 samples: Effects of irradiation and

vitrification and comparison with telescopic observations: 

LSCP 1, v. 3, p. 2013-2023.

Cooper, H.F., Jr., 1977, A summary of explosion cratering phenomena relevant to meteor impact events, in Roddy, D.J, 

Pepin, R.O., and Merrill, R.B., eds., Impact and explosion 

cratering: Planetary and terrestrial implications: New York, 

Pergamon, p. 11-44.

Cooper, H.F., and Sauer, F.M., 1977, Crater-related ground 

motions and implications for crater scaling, in Roddy, D.J, 

Pepin, R.O., and Merrill, R.B., eds., Impact and explosion 

cratering: Planetary and terrestrial implications: New York, 

Pergamon, p. 1133-1163.

Cooper, M.R., Kovach, R.L., and Watkins, J.S., 1974, Lunar nearsurface structure: Reviews of Geophysics and Space Physics, 

v. 12, no. 3, p. 291-308.

Crater Analysis Techniques Working Group, 1979, Standard 

techniques for presentation and analysis of crater size-frequency data: Icarus, v. 37, no. 2, p. 467 474.

Croft, S.K., 1978, Lunar crater volumes: Interpretation by models of impact cratering and upper crustal structure: LPSCP 9, 

v. 3, p. 3711-3733.

 1980, Crateringflow fields: Implications for the excavation 

and transient expansion stages of crater formation: LPSCP 

11, v. 3, p. 2347-2378.

-1981, The modification stage of basin formation: Conditions of ring formation, in Multi-ring basins: LPSCP 12, pt. A,

p. 227-257. 

Crozaz, Ghislaine, 1977, The irradiation history of the lunar soil:

Physics and Chemistry of the Earth, v. 10, no. 3, p. 197-214. 

Crozaz, Ghislaine, Drozd, R.J, Hohenberg, C.M., Hoyt, H.P,

Ragan, D., Walker, R.M., and Yuhas, D., 1972, Solar flare and

galactic cosmic ray studies of Apollo 14 and 15 samples: LSCP

3, v. 3, p. 2917-2931. 

Cruikshank, D.P., and Wood, C.A., 1972, Lunar rilles and

Hawaiian volcanic features: Possible analogues: The Moon, v.

3, no. 4, p. 412-446. 

Cummings, David, 1971, Geologic map of the Wichmann CA

region of the Moon, Lunar Orbiter site III P-ll, Oceanus

Procellarum, including Apollo landing sites 4 and 4R: USGS

Map 1-624 [ORB III-P11 (110)], scale 1:100,000.

 1972, Geologic map of the Clavius quadrangle of the Moon:

USGS Map 1-706 (LAC-126), scale 1:1,000,000. 

Danes, Z.F., 1965, Rebound processes in large craters: USGS

ASAPR, July 1,1964 to July 1,1965, pt. A, p. 81-100. 

Davis, PA., Jr., 1980, Iron and titanium distribution on the moon

from orbital gamma ray spectrometry with implications for

crustal evolutionary models: Journal of Geophysical

Research, v. 85, no. B6, p. 3209-3224. 

De Hon, R.A., 1971, Cauldron subsidence in lunar craters Ritter

and Sabine: Journal of Geophysical Research, v. 76, no. 23, p.

5712-5718.

 1974, Thickness of mare material in the Tranquillitatis 

and Nectaris basins: LSCP 5, v. 1, p. 53-59.

 1979, Thickness of the western mare basalts: LPSCP 10, v.

3, p. 2935-2955. 

De Hon, R.A., and Waskom, J.D., 1976, Geologic structure of the

eastern mare basins: LSCP 7, v. 3, p. 2729-2746. 

De Laeter, JR., Vernon, M.J, and Compston, W, 1973, Revision of

lunar Rb-Sr ages: Geochimica et Cosmochimica Acta, v. 37,

no. 3, p. 700-702. 

Delano, J.W., 1975, Petrology of the Apollo 16 mare component:

Mare Nectaris: LSCP 6, v. 1, p. 15-47. 

 1979, Apollo 15 green glass: Chemistry and possible

origin: LPSCP 10, v. 1, p. 275-300.

-1980, Chemistry and liquidus phase relations of Apollo 15

red glass: Implications for deep lunar interior: LPSCP 11, v. 1,

p. 251-288. 

Delano, JW, and Ringwood, A.E., 1978, Siderophile elements in

the lunar highlands: Nature of the indigenous component

and implication for the origin of the Moon: LPSCP 9, v. 1, p.

111-159. 

Dence, M.R., 1964, A comparative structural and petrographic

study of probable Canadian meteorite craters: Meteoritics, v.

2, p. 249-270.

 1965, The extraterrestrial origin of Canadian craters, in 

Whipple, H.E., ed., Geological problems in lunar research: 

New York Academy of Science Annals, v. 123, art. 2, p. 

941-969.

 1968, Shock zoning at Canadian craters: Petrography and 

structural implications, in French, B.M., and Short, N.M., 

eds., Shock metamorphism of natural materials: Baltimore, 

Mono, p. 169-184.

-1971, Impact melts: Journal of Geophysical Research, v.

76, no. 23, p. 5552-5565.

Dence, M.R., and Grieve, R.A.F., 1979, The formation of complex 

impact structures [abs.]: LPS X, p. 292-294.

Dence, M.R., and Plant, A.G., 1972, Analysis of Fra Mauro samples and the origin of the Imbrium Basin: LSCP 3, v. 1, p. 

379-399.

Dence, M.R., Grieve, R.A.F., and Robertson, P.B., 1977, Terrestrial impact structures: Principal characteristics and 

energy considerations, in Roddy, D.J, Pepin, R.O., and Merrill, R.B., eds., Impact and explosion cratering: Planetary and 

terrestrial implications: New York, Pergamon, p. 247 275.

Dennis, J.G., 1971, Ries structure, southern Germany, a review: 

Journal of Geophysical Research, v. 76, no. 23, p. 5394-5406.

De Paolo, D.J., McCulloch, M.T., Papanastassiou, D.A., Huneke, 

J.C., and Wasserburg, G. J, 1978, Ages, evolution and neutron 

effects of Luna 24 samples [abs.]: LPS IX, p. 244-246.

Dodd, R.T., Jr., Salisbury, JW, and Smalley, V.G., 1963, Cratering 

frequency and interpretation of lunar history: Icarus, v. 2, no.

5-6, p. 466-480. 

Dowty, Eric, Keil, Klaus, Prinz, Martin, Gros, J, and Takahashi,

H., 1976, Meteorite-free Apollo 15 crystalline KREEP: LSCP

7, v. 2, p. 1833-1844. 

Dowty, Eric, Prinz, Martin, and Keil, Klaus, 1974a, Ferroan

anorthosite: A widespread and distinctive lunar rock type:

Earth and Planetary Science Letters, v. 24, no. 1, p. 15-25.

 1974b, "Very high alumina basalt": A mixture and not a 

magma type: Science, v. 183, no. 4130, p. 1214-1215.

Dreibus, Gerlind, Kruse, H., Spettel, B., and Wanke, H., 1977, 

The bulk composition of the moon and the eucrite parent 

body: LSCP 8, v. 1, p. 211-227.

Drozd, R.J, Hohenberg, C.M., Morgan, C.J, Podosek, FA., and 

Wroge, M.L., 1977, Cosmic-ray exposure history at TaurusLittrow: LSCP 8, v. 3, p. 3027-3043.

Drozd, R.J, Hohenberg, C.M., Morgan, C.J, and Ralston, C.E., 

1974, Cosmic-ray exposure history at the Apollo 16 and other 

lunar sites: Lunar surface dynamics: Geochimica et Cosmochimica Acta, v. 38, no. 10, p. 1625-1642.

Duncan, A.R., Grieve, R.A.F., and Weill, D.F, 1975, The life and 

times of Big Bertha: Lunar breccia 14321: Geochimica et 

Cosmochimica Acta, v. 39, no. 3, p. 265-273.

Dvorak, John, and Phillips, R.J, 1978, Lunar Bouguer gravity 

anomalies: Imbrian age craters: LPSCP 9, v. 3, p. 3651-3668.

 1979, Gravity anomaly and structure associated with the 

Lament region of the moon: LPSCP 10, v. 3, p. 2265-2275.

Dymek, R.F, Albee, A.L., and Chodos, A.A., 1975, Comparative 

petrology of lunar cumulate rocks of possible primary origin: 

Dunite 72415, troctolite 76535, norite 78235, and anorthosite 

62237: LSCP 6, v. 1, p. 301-341.

 1976, Petrology and origin of Boulders #2 and #3, Apollo 

17 Station 2: LSCP 7, v. 2, p. 2335-2378.

Eberhardt, P., Geiss, J, Graf, H., Grogler, N., Krahenbuhl, U., 

Schwaller, H., Schwarzmuller, J, and Stettler, A., 1970, Correlation between rock type and irradiation history of Apollo 11 

igneous rocks: Earth and Planetary Science Letters, v. 10, no. 

1, p. 67-72.

Eberhardt, P., Geiss, J, Grogler, N., Maurer, P., and Stettler, A., 

1973a, 39Ar-40Ar ages of lunar material [abs.]: Meteoritics, v. 

8, no. 4, p. 360-361.

Eberhardt, P., Geiss, J, Grogler, N., and Stettler, A., 1973b, How 

old is the crater Copernicus?: The Moon, v. 8, no. 1-2, p. 

104-114.

Eggleton, R. E., 1964, Preliminary geology of the Riphaeus quadrangle of the Moon and definition of the Fra Mauro Formation: USGS ASAPR, August 25,1962 to July 1,1963, pt. A, p. 

46-63.

 1965, Geologic map of the Riphaeus Mountains region of 

the Moon: USGS Map 1-458 (LAC-76), scale 1:1,000,000.

Eggleton, R.E., and Marshall, C.H., 1962, Notes on the Apenninian Series and pre-Imbrian stratigraphy in the vicinity of 

Mare Humorum and Mare Nubium: USGS ASSPR, February 

26,1961 to August 24,1961, p. 132-137.

Eggleton, R.E., and Offield, T.W., 1970, Geologic maps of the Fra 

Mauro region of the Moon: Apollo 14 premission maps: USGS 

Map 1-708, scales 1:250,000,1:25,000, 2 sheets.

Eggleton, R.E., and Schaber, G.G., 1972, Cayley Formation interpreted as basin ejecta, pt. B of Photogeology, [chap.] 29 of 

Apollo 16 preliminary science report: NASA Report SP-315, 

p. 29-7 to 29-16.

Eggleton, R.E., Schaber, G.G., and Pike, R.J, 1974, Photogeologic 

detection of surfaces buried by mare basalts [abs.]: LSV, pt. 1, 

p. 200-202.

Eichhorn, G., McGee, J.J., James, O. B., and Schaeffer, O.A., 1979, 

Consortium breccia 73255: Laser 39Ar- 40Ar dating of 

aphanite samples: LPSCP 10, v. 1, p. 763-788.

El-Baz, Farouk, 1972, The Alhazen to Abul Wafa swirl belt: An 

extensive field of light-colored, sinuous markings, pt. T of 

Photogeology, [chap.] 29 of Apollo 16 preliminary science 

report: NASA Report SP-315, p. 29-93 to 29-97.

El-Baz, Farouk, and Roosa, S.A., 1972, Significant results from 

Apollo 14 lunar orbital photography: LSCP 3, v. 1, p 63-83.

El-Baz, Farouk, Worden, A.M., and Brand, V.D., 1972, Astronaut 

observations from lunar orbit and their geologic significance: 

LSCP 3, v. 1, p. 85-104.

Eliason, E.M., and Soderblom, L.A., 1977, An array processing 

system for lunar geochemical and geophysical data: LSCP 8, 

v.l, p. 1163-1170.

Elston, D.P., 1972, Geologic map of the Colombo quadrangle of 

the Moon: USGS Map 1-714 (LAC-79), scale 1:1,000,000.

Engelhardt, Wolf von, 1967, Neue Beobachtungen im Nordlinger 

Ries [New observations in the Nordlingen Ries]: Geologische 

Rundschau, v. 57, no. 1, p. 165-188.

 1971, Detrital impact formations: Journal of Geophysical 

Research, v. 76, no. 23, p. 5566-5574.

Engelhardt, Wolf von, and Stengelin, Rudolf, 1979, Normative 

composition and classification of lunar igneous rocks and 

glasses, I. Lunar igneous rocks: Earth and Planetary Science 

Letters, v. 42, no. 2, p. 213-222.

Engelhardt, Wolf von, Arndt, J, Stoffler, D., and Schneider, H., 

1972, Apollo 14 regolith and fragmental rocks, their compositions and origin by impacts: LSCP 3, v. 1, p. 753-770.

Eugster, O., Eberhardt, P., Geiss, J, Grogler, N., Jungck, M., and 

Morgeli, M., 1977, The cosmic ray exposure history of Shorty 

Crater samples: The age of Shorty Crater: LSCP 8, v. 3, p. 

3059-3082.

Evans, R.E., and El-Baz, Farouk, 1973, Geological observations 

from lunar orbit, [chap.] 28 of Apollo 17 preliminary science 

report: NASA Report SP-330, p. 28-1 to 28-32.

Evensen, N.M., Murthy, V.R., and Coscio, M.R., 1973, Rb-Sr ages 

of some mare basalts and their isotopic and trace element
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Gooley, R.C., Brett, Robin, Warner, Jeff, and Smyth, JR., 1974, A 

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Greeley, Ronald, 1971, Observations of actively forming lava tubes 

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 1973, Comparative geology of crater Aratus CA (Mare 

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Greeley, Ronald, and Carr, M.H., eds., 1976, A geological basis for

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Greeley, Ronald, and Gault, D.E., 1970, Precision size-frequency

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 1973, Crater frequency age determinations for the proposed Apollo 17 site at Taurus-Littrow: Earth and Planetary

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Greeley, Ronald, and Spudis, P.D., 1978, Mare volcanism in the

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1975, Experimental petrology of Apollo 17 basalts: LSCP 6, v.

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1972, Experimental petrology and petrogenesis of Apollo 14

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Green, Jack, 1971, Copernicus as a lunar caldera: Journal of

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 1976, Review of "Planetary Geology," by N.M. Short: Sky 

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 1978, The melt rocks at Brent Crater, Ontario, Canada: 

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 1970b, Geologic map of the Sabine D region of the Moon, 

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 1973, Stratigraphy of ejecta from the lunar crater 

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 1976, Volcanic features of the nearside equatorial lunar 

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 1964, Stratigraphy and structure of the.Montes Apenninus region of the Moon: USGS ASAPR, August 25,1962 to

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Hackman, R.J, and Mason, A.C., 1961, Engineer special study of 

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1978, Thorium concentrations in the lunar surface. II: Deconvolution modeling and its application to the regions of 

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Effects of vapor-phase deposition processes on the optical, 

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 1964a, On the distribution of lunar crater diameters: Tucson, University of Arizona, Lunar and Planetary Laboratory 

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 1964b, Radial structures surrounding lunar basins, II: 

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 1965a, Secular changes in meteoritic flux through the 

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 1965b, Terrestrial and lunar flux of large meteorites in the 

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 1966, Early lunar cratering: Icarus, v. 5, no. 4, p. 406-418.

 1967, Lunar crater counts. II: Three lunar surface type 

areas: Tucson, University of Arizona, Lunar and Planetary 

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 1968, Lunar crater counts. VI: The young craters Tycho, 

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 1972a, Interplanet variations in scale of crater morphology Earth, Mars, Moon: Icarus, v. 17, no. 3, p. 707-713.

 1972b, Moons and planets: Belmont, Calif., \Vadsworth, 

404 p.

 1972c, Paleocratering of the Moon: Review of post-Apollo 

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 1973, Ancient lunar mega-regolith and subsurface structure: Icarus, v. 18, no. 4, p. 634-636.

 1975, Lunar "cataclysm": A misconception?: Icarus, v. 24, 

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 1980, Dropping stones in magma oceans: Effects of early 

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 1981, Discovery of multi-ring basins: Gestalt perception in 

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 1983, Moons and planets (2d ed): Belmont, Calif., 

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 1977b, Pre-Imbrian history of the Fra Mauro region and 

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 1978a, Criteria for the identification of the Imbrium 

ejecta and local components in the Apollo 14 samples [abs.]: 

LPS IX, pt. 1, p. 477-479.

 1978b, Lunar KREEP volcanism: Geologic evidence for 

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 1974a, Lunar dark-mantle deposits: Possible clues to the 

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 1974b, Morphology and structure of the Taurus-Littrow 

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 1974c, Orientate multi-ringed basin interior and implications for the petrogenesis of lunar highland samples: The 

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 1974d, Stratigraphy of the Descartes region (Apollo 16): 

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 1976a, Lunar volcanism in space and time: Reviews of 

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 1976b, The significance of substrate characteristics in 

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 1977, Origin of outer rings in lunar multi-ringed basins: 

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 1979a, Lava flooding of early planetary crusts: Geometry, 

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X, pt. 2, p. 516-518.

 1979b, Lava flooding of early planetary crusts: Geometry, 

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Head, J.W., and Gifford, Ann, 1980, Lunar mare domes: Classification and modes of origin: The Moon and the Planets, v. 22, 

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Head, J.W., and Goetz, A.F.H., 1972, Descartes region: Evidence 

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Head, J.W., and Hawke, B.R., 1975, Geology of the Apollo 14 

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Head, J.W., and McCord, T.B., 1978, Imbrian-age highland volcanism on the Moon: The Gruithuisen and Mairan domes: 

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Head, J.W., and Wilson, Lionel, 1979, Alphonsus-type dark-halo 

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Head, J.W., Adams, J.B., McCord, T.B., Pieters, C.M., and Zisk, 

S.H., 1978a, Regional stratigraphy and geologic history of 

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Head, J.W.,, Pieters, C.M., McCord, T.B., Adams, J.B., and Zisk, 

S.H., 1978b, Definition and detailed characterization of lunar 

surface units using remote observations: Icarus, v. 33, no. 1, p. 

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Head, J.W., Robinson, Edmund, and Phillips, Roger, 1981, 

Topography of the Orientale basin [abs.]: LPS XII, pt. 2, p. 

421-423.

Head, J.W., Settle, Mark, and Stein, R.S., 1975, Volume of material ejected from major lunar basins and implications for the 

depth of excavation of lunar samples: LSCP 6, v. 3, p. 

2805-2829.

Heiken, G.H., 1975, Petrology of lunar soils: Reviews of 

Geophysics and Space Physics, v. 13, no. 4, p. 567-587.

Heiken, G.H., McKay, D.S., and Brown, R.W., 1974, Lunar deposits of possible pyroclastic origin: Geochimica et Cosmochimica Acta, v. 38, no. 11, p. 1703-1718.

Herbert, Floyd, Drake, M.J., Sonett, C.P, and Wiskerchen, M.J., 

1977, Some constraints on the thermal history of the lunar 

magma ocean: LSCP 8, v. 1, p. 573-582.

Hertogen, Jan, Janssens, M.-J, Takahashi, H., Palme, Herbert, 

and Anders, Edward, 1977, Lunar basins and craters: Evidence for systematic compositional changes of bombarding 

population: LSCP 8, v. 1, p. 17-45.

Herzberg, C.T., 1978, The bearing of spinel cataclasites on the 

crust-mantle structure of the moon: LPSCP 9, v. 1, p. 

319-336.

Herzberg, C.T., and Baker, M.B., 1980, The cordierite- to spinelcataclasite transition: Structure of the lunar crust, in Lunar 

and Planetary Institute, compiler, Conference on the Lunar 

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Hess, PC., Rutherford, M.J, and Campbell, H.W., 1977, Origin 

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 1973, Apollo 17 site selection, [chap.] 1 of Apollo 17 preliminary science report: NASA Report SP-330, p. 1-1 to 1-5.

Hodges, C.A., 1972, Descartes highlands: Possible analogs 

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 1973a, Geologic map of the Langrenus quadrangle of the 

Moon: USGS Map 1-739 (LAC-80), scale 1:1,000,000. 

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Hodges, C.A., and Muehlberger, W.R., 1981, A summary and 

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Apollo 16 area, central lunar highlands: USGS Professional 

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Hodges, C.A., and Wilhelms, D.E., 1978, Formation of lunar basin 

rings: Icarus, v. 34, no. 2, p. 294-323.

Hodges, C.A., Muehlberger, W.R., and Ulrich, G.E., 1973, Geologic setting of Apollo 16: LSCP 4, v. 1, p. 1-25.

Holcomb, Robin, 1971, Terraced depressions in lunar maria: Journal of Geophysical Research, v. 76, no. 23, p. 5703-5711.

Holt, H.E., 1974, Geologic map of the Purbach quadrangle of the 

Moon: USGS Map 1-822 (LAC-95), scale 1:1,000,000.

Hood, L.L., Coleman, P.J., Jr., and Wilhelms, D.E., 1979, Lunar 

nearside magnetic anomalies: LPSCP 10, v. 3, p. 2235-2257.

Horn, P., Kirsten, T, and Jessberger, E.K., 1975, Are there A 12 

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Horz, Friedrich, 1978, How thick are lunar mare basalts?: 

LPSCP 9, v. 3, p. 3311-3331.

 1981, The "Bunte Breccia" of the Ries: Implications for the 

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Workshop on Apollo 16: Houston, Tex., Lunar and Planetary 

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Horz, Friedrich, and Banholzer, G.S., Jr., 1980, Deep seated 

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Hbrz, Friedrich, and Ostertag, Rolf, 1979, The transient crater of 

the Ries crater, Germany [abs.]: LPS X, pt. 2, p. 570-572.

Howard, K.A., 1970, Mascons, mare rock and isostasy: Nature, v. 

226, no. 5249, p. 924-925.

 1974, Fresh lunar impact craters: Review of variations 

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1-840, scale 1:250,000.

Howard, K.A., and Masursky, Harold, 1968, Geologic map of the 

Ptolemaeus quadrangle of the Moon: USGS Map 1-566 

(LAC-77; RLC-13), scale 1:1,000,000.

Howard, K.A., and Wilshire, H.G., 1975, Flows of impact melt at 

lunar craters: USGS Journal of Research, v. 3, no. 2, p. 

237-257.

Howard, K.A., Carr, M.H., and Muehlberger, W.R., 1973, Basalt 

stratigraphy of southern Mare Serenitatis, pt. A of Stratigraphic studies, [chap.] 29 of Apollo 17 preliminary science 

report: NASA Report SP-330, p. 29-1 to 29-12.

Howard, K.A., Head, J.W, and Swann, G.A., 1972, Geology of 

Hadley rille: LSCP 3, v. 1, p. 1-14.

Howard, K.A., Wilhelms, D.E., and Scott, D.H., 1974, Lunar 

basin formation and highland stratigraphy: Reviews of 

Geophysics and Space Physics, v. 12, no. 3, p. 309-327.

Hubbard, N.J., 1979, Regional chemical variations in lunar 

basaltic lavas: LPSCP 10, v. 2, p. 1753-1774.

Hubbard, N.J., and Gast, P.W., 1971, Chemical composition and 

origin of nonmare lunar basalts: LSCP 2, v. 2, p. 999-1020.

Hubbard, N.J., and Minear, J.W, 1975, A physical and chemical 

model of early lunar history: LSCP 6, v. 1, p. 1057-1085.

Hubbard, N.J., Meyer, Charles, Jr., Gast, P.W., and Wiesmann, 

Henry, 1971, The composition and derivation of Apollo 12 soils: 

Earth and Planetary Science Letters, v. 10, no. 3, p. 341-350.

Hubbard, N.J., Rhodes, J.M., and Gast, P.W., 1973, Chemistry of 

lunar basalts with very high alumina contents: Science, v. 181, 

no. 4097, p. 339-342.

Hubbard, N.J., Rhodes, J.M., Wiesmann, H., Shih, C.-Y, and 

Bansal, B.M., 1974, The chemical definition and interpretation of rock types returned from the non-mare regions of the 

Moon: LSCP 5, v. 2, p. 1227-1246.

Hubbard, N.J., Vilas, Faith, and Keith, J.E., 1978, From Serenity 

to Langemak: A regional chemical setting for Mare Crisium, 

in Lunar and Planetary Institute, compiler, Mare Crisium: 

The view from Luna 24: Conference on Luna 24, Houston, 

Tex., 1977, Proceedings: New York, Pergamon (Geochimica et 

Cosmochimica Acta, supp. 9), p. 13-32.

Huneke, J.C., 1978, 40Ar-39Ar microanalysis of single glass balls 

and 72435 breccia clasts: LPSCP 9, v. 2, p. 2345-2362.

Huneke, J.C., and Wasserburg, G.J., 1975, Trapped 40Ar in troctolite 76535 and evidence for enhanced 40Ar- 39Ar age 

plateaus [abs.]: LS VI, pt. 1, p. 417-419.

Huneke, J.C., Jessberger, E.K., Podosek, FA., and Wasserburg, 

G.J, 1973, 40Ar/39Ar measurements in Apollo 16 and 17 samples and the chronology of metamorphic and volcanic activity 

in the Taurus-Littrow region: LSCP 4, v. 2, p. 1725-1756.

Huneke, J.C., Jessberger, E.K., and Wasserburg, G.J, 1974, The 

age of metamorphism of a highland breccia (65015) and a 

glimpse at the age of its protolith [abs.]: LS V, pt. 1, p. 

375-377.

Huneke, J.C., Podosek, FA., and Wasserburg, G.J, 1972, Gas 

retention and cosmic-ray exposure ages of a basalt fragment

from Mare Fecunditatis: Earth and Planetary Science Letters, v. 13, no. 2, p. 375-383.

Husain, Liaquat, 1974, 40Ar-39Ar chronology and cosmic ray 

exposure ages of the Apollo 15 samples: Journal of 

Geophysical Research, v. 79, no. 17, p. 2588-2606.

Husain, Liaquat, and Schaefer, O.A., 1973, Lunar volcanism: Age 

of the glass in the Apollo 17 orange soil: Science, v. 180, no. 

4093, p. 1358-1360.

 1975, Lunar evolution: The first 600 million years: 

Geophysical Research Letters, v. 2, no. 1, p. 29-32.

Husain, Liaquat, Schaeffer, O.A., Funkhouser, J, and Sutter, JF, 

1972, The ages of lunar material from Fra Mauro, Hadley 

Rille, and Spur Crater: LSCP 3, v. 2, p. 1557-1567.

Irving, A.J, 1975, Chemical, mineralogical and textural systematics of non-mare melt rocks: Implications for lunar 

impact and volcanic processes: LSCP 6, v. 1, p. 363-394.

 1977, Chemical variation and fractionation of KREEP

basalt magma: LSCP 8, v. 2, p. 2433-2448. 

Jackson, E.D., 1971, The origin of ultramafic rocks by cumulus

processes: Fortschritte der Mineralogie, v. 48, no. 1, p.

128-174. 

Jackson, E.D., Sutton, R.L., and Wilshire, H.G., 1975, Structure

and petrology of a cumulus norite boulder sampled by Apollo

17 in Taurus-Littrow valley, the Moon: Geological Society of

America Bulletin, v. 86, no. 4, p. 433-442. 

James, O.B., 1972, Lunar anorthosite 15415: Texture, minerology,

and metamorphic history: Science, v. 175, no. 4020, p.

432-436.

 1973, Crystallization history of lunar feldspathic basalt 

14310: USGS Professional Paper 841, 29 p.

 1976, Petrology of aphanitic lithologies in consortium 

breccia 73215: LSCP 7, v. 2, p. 2145-2178.

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 1974, Apollo laser altimetry and inferences as to lunar 

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regolith at the Apollo 12 site: LSCP 2, v. 1, p. 701-718.

Raedeke, L.D., and McCallum, I.S., 1980, A comparison of fractionation trends in the lunar crust and the Stillwater Complex, in Lunar and Planetary Institute, compiler, Conference 

on the Lunar Highlands Crust, Houston, Tex., 1979, Proceedings: New York, Pergamon (Geochimica et Cosmochimica 

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Reed, V.S., 1981, Geology of the areas near South Ray and Baby 

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Muehlberger, W.R., eds., Geology of the Apollo 16 area, central lunar highlands: USGS Professional Paper 1048, p. 

82-105.

Reed, V.S., and Wolfe, E.W, 1975, Origins of the Taurus-Littrow 

massifs: LSCP 6, v. 3, p. 2443-2461.

Reedy, R.C., 1978, Planetary gamma-ray spectroscopy: LPSCP 9, 

v. 3, p. 2961-2984.

Rehfuss, D.E., 1974, Glass production differences for equaldiameter impact craters: The Moon, v. 11, no. 1-2, p. 19-28.

Reid, A.M., Duncan, A.R., and Richardson, S.H., 1977, In search 

of LKFM: LSCP 8, v. 2, p. 2321-2338.

Reid, A.M., Ridley, W.I, Harmon, R.S., Warner, J.L., Brett, Robin, 

Jakes, P., and Brown, R.W, 1972a, Highly aluminous glasses 

in lunar soils and the nature of the lunar highlands: Geochimica et Cosmochimica Acta, v. 36, no. 8, p. 903-912.

Reid, A.M., Warner, J.L., Ridley, W.I., and Brown, R.W, 1972b, 

Major element composition of glasses in three Apollo 15 soils: 

Meteoritics, v. 7, no. 3, p. 395-415.

Rhodes, J.M., and Hubbard, N.J., 1973, Chemistry, classification, 

and petrogenesis of Apollo 15 mare basalts: LSCP 4, v. 2, p. 

1127-1148.

Rhodes, J.M., Blanchard, D.P., Dungan, M.A., Brannon, J.C., and 

Rodgers, K.V., 1977, Chemistry of Apollo 12 mare basalts: 

Magma types and fractionation processes: LSCP 8, v. 2, p. 

1305-1338.

Rhodes, J.M., Hubbard, N.J., Wiesmann, H., Rodgers, K.V., 

Brannon, J.C., and Bansal, B.M., 1976, Chemistry, classification, and petrogenesis of Apollo 17 mare basalts: LSCP 7, v. 2, 

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Ridley, W.I., 1975, On high-alumina mare basalts: LSCP 6, v. 1, p. 

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Rinehart, J.S., 1975, Stress transients in solids: Santa Fe, N. 

Mex., HyperDynamics, 230 p.

Ringwood, A.E., and Essene, E., 1970, Petrogenesis of Apollo 11 

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LSCP1, v.l, p. 769-799.

Roberts, W.A., 1966, Shock a process in extraterrestrial sedimentology: Icarus, v. 5, no. 5, p. 459-477.

 1968, Shock crater characteristics, in French, B.M., and 

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Baltimore, Mono, p. 101-114.

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Roddy, D.J., 1968, The Flynn Creek crater, Tennessee, in French, 

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 1976, High-explosion cratering analogs for bowl-shaped, 

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 1977, Pre-impact conditions and cratering processes at the 

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277-308.

-1979, Structural deformation at the Flynn Creek impact

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Ronca, L.B., 1965, Selenology vs geology of the Moon etc: Geotimes, v. 9, no. 9, p. 13. 

Rowan, L.C., 1971a, Geologic map of the Oppolzer A region of the

Moon, Lunar Orbiter site II P-8, Sinus Medii, including

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(100)], scale 1:100,000.

 1971b, Geologic map of the Rupes Altai quadrangle of the 

Moon: USGS Map 1-690 (LAC-96), scale 1:1,000,000.

Runcorn, S.K., 1980, An iron core in the moon generating an 

early magnetic field?: LPSCP 10, v. 3, p. 2325-2333.

 1982, Primeval displacements of the lunar pole: Physics of 

the Earth and Planetary Interiors, v. 29, p. 135-147.

Ryder, Graham, 1976, Lunar sample 15405: Remnant of a 

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 1979, The chemical components of highlands breccias: 

LPSCP 10, v. 1, p. 561-581.

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Ryder, Graham, and Bower, J.F, 1976, Poikilitic KREEP impact 

melts in the Apollo 14 white rocks: LSCP 7, v. 2, p. 1925-1948.

 1977, Petrology of Apollo 15 black-and-white rocks 15445 

and 15455 fragments of the Imbrium impact melt sheet?: 

LSCP 8, v. 2, p. 1895-1923.

Ryder, Graham, and Norman, M.D., 1980, Catalog of Apollo 16 

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Ryder, Graham, and Spudis, P.O., 1980, Volcanic rocks in the 

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Ryder, Graham, and Taylor, G.J., 1976, Did mare-type volcanism 

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Ryder, Graham, and Wood, J.A., 1977, Serenitatis and Imbrium 

impact melts: Implications for large-scale layering in the 

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Ryder, Graham, Stoeser, D.B., Marvin, U.B., and Bower, J.F, 

1975a, Lunar granites with unique ternary feldspars: LSCP 

6, v. 1, p. 435-449.

Ryder, Graham, Stoeser, D.B., Marvin, U.B., Bower, J.F, and 

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Ryder, Graham, Stoeser, D.B., and Wood, J.A., 1977, Apollo 17 

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Saari, J.M., Shorthill, R.W, and Deaton, T.K., 1966, Infrared and 

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Sabaneyev, P.P., 1962, Some results deduced from simulation of 

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The Moon: London, Academic Press, p. 419-431.

Saito, Kazuo, and Alexander, B.C., 1979, 40Ar-39Ar studies of 

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Salisbury, J.W, and Glaser, RE., eds., 1964, The lunar surface 

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Salisbury, J.W, Adler, J.E.M., and Smalley, V.G., 1968, Darkhaloed craters on the Moon: Royal Astronomical Society 

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Sanchez, A.G., 1981, Geology of Stone mountain, [chap.] D4 of 

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Geology of the Apollo 16 area, central lunar highlands: USGS 

Professional Paper 1048, p. 106-126.

Sato, Motokai, 1979, The driving mechanism of lunar pyroclastic 

eruptions inferred from the oxygen fugacity behavior of 

Apollo 17 orange glass: LPSCP 10, v. 1, p. 311-325.

Saunders, R.S., and Wilhelms, D.E., 1974, Geologic map of the 

Wilhelm quadrangle of the Moon: USGS Map 1-824 

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Schaber, G.G., 1969, Geologic map of the Sinus Iridum quadrangle of the Moon: USGS Map 1-602 (LAC-24), scale 

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290 THE GEOLOGIC HISTORY OF THE MOON

 1973, Lava flows in Mare Imbrium: Geologic evaluation 

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Geology of the Apollo 16 area, central lunar highlands: USGS 

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Schaber, G.G., Boyce, J.M., and Moore, H.J., 1976, The scarcity of 

mappable flow lobes on the lunar maria: Unique morphology 

of the Imbrium flows: LSCP 7, v. 3, p. 2783-2800.

Schaber, G.G., Thompson, T.W., and Zisk, S.H., 1975, Lava flows 

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Schaeffer, G.A, and Schaeffer, O.A., 1977, 39Ar-40Ar ages of lunar 

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Schaeffer, O.A., and Husain, Liaquat, 1973, Early lunar history: 

Ages of 2 to 4 mm soil fragments from the lunar highlands: 

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 1974, Chronology of lunar basin formation: LSCP 5, v. 2, p. 

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Schaeffer, O.A., Bence, A.E., Eichhorn, G., Papike, J.J., and Vaniman, D.T., 1978, 39Ar-40Ar and petrologic study of Luna 24 

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Schaeffer, O.A., Husain, Liaquat, and Schaeffer, G.A, 1976, Ages 

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Schmidt, R.M., 1977, A centrifuge cratering experiment: 

Development of a gravity-scaled yield parameter, in Roddy, 

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 1981, Scaling crater time-of-formation [abs.]: Eos (American Geophysical Union Transactions), v. 62, no. 45, p. 944.

Schmitt, H.H., Lofgren, G., Swann, G.A, and Simmons, G., 1970, 

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Schonfeld, Ernest, 1976, Chronology of the early lunar crust: 

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 1977, Comparison of orbital chemistry with crustal thickness and lunar sample chemistry: LSCP 8, v. 1, p. 1149-1162.

Schonfeld, Ernest, and Bielefeld, M.J., 1978, Correlation of dark 

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Schonfeld, Ernest, andMeyer, Charles, Jr., 1972, The abundance 

of components of the lunar soils by a least-squares mixing 

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Schultz, P.H., 1976a, Floor-fractured lunar craters: The Moon, v. 

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 1976b, Moon morphology: Austin, University of Texas

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Schultz, PH., and Gault, D.E., 1975a, Seismic effects from major

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 1975b, Seismically induced modification of the lunar surface features: LSCP 6, v. 3, p. 2845-2862.

Schultz, PH., and Mendell, Wendell, 1978, Orbital infrared 

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impact ejecta emplacement: LPSCP 9, v. 3, p. 2857-2883.

Schultz, P.H., and Mendenhall, M.H., 1979, On the formation of 

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Schultz, PH., and Spencer, J., 1979, Effects of substrate strength 

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Schultz, PH., and Spudis, P.H., 1979, Evidence for ancient mare 

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Schultz, P.H., and Srnka, L.J., 1980, Cometary collisions on the 

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Schultz, PH., Greeley, Ronald, and Gault, D.E., 1977, Interpreting statistics of small lunar craters: LSCP 8, v. 3, p. 

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 1972b, Geologic map of the Maurolycus quadrangle of the 

Moon: USGS Map 1-695 (LAC-113), scale 1:1,000,000.

 1973, Mare Serenitatis cinder cones and terrestrial analogs, pt. B o/Volcanic studies, [chap.] 30 of Apollo 17 preliminary science report: NASA Report SP-330, p. 30-7 to 30-8.

 1974, The geologic significance of some lunar gravity 

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Scott, D.H., and Eggleton, R.E., 1973, Geologic map of the 

Rumker quadrangle of the Moon: USGS Map 1-805 

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Scott, D.H., and Pohn, H.A., 1972, Geologic map of the Macrobius 

quadrangle of the Moon: USGS Map 1-799 (LAC-43), scale 

1:1,000,000.

Scott, D.H., Diaz, J.M., and Watkins, J.A., 1977, Lunar farside 

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Scott, D.H., Lucchitta, B.K., and Carr, M.H., 1972, Geologic maps 

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Scott, D.H., McCauley, J.F., and West, M.N., 1977, Geologic map of 

the west side of the Moon: USGS Map 1-1034, scale 

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Scott, D.H., West, M.N., Lucchitta, B.K., and McCauley, J.F, 1971, 

Preliminary geologic results from orbital photography, pt. B 

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Settle, Mark, and Head, J.W, 1979, The role of rim slumping in 

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Shih, C.-Y., 1977, Origins of KREEP basalts: LSCP 8, v. 2, p. 

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Shoemaker, E.M., 1960, Penetration mechanics of high velocity 

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 1962a, Exploration of the Moon's surface: American Scientist, v. 50, no. 1, p. 99-130.

 1962b, Interpretation of lunar craters, in Kopal, Zdenek, 

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 1963, Impact mechanics at Meteor Crater, Arizona, in 

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 1964, The geology of the Moon: Scientific American, v. 211, 

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 1965, Preliminary analysis of the fine structure of the 

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 1966, Progress in the analysis of fine structure and 

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 1971, Origin of fragmental debris on the lunar surface and 

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 1972, Cratering history and early evolution of the Moon 

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Shoemaker, E.M., and Hackman, R.J., 1962, Stratigraphic basis 

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Shoemaker, E.M., Bailey, N.G., Batson, R.M., Dahlem, D.H., 

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H.E., Larson, KB., Rennilson, J.J., Schaber, G.G., Schleicher, 

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 1967b, Television observations from Surveyor V, [chap.] 3 

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 1968, Television observations trom Surveyor VII, [chap.] 3 

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1969b, Observations of the lunar regolith and the earth

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Shoemaker, E.M., Hackman, R.J., and Eggleton, R.E., 1962a, 

Interplanetary correlation of geologic time: Advances in the 

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Shoemaker, E.M., Hackman, R.J., Eggleton, R.E., and Marshall, 

C.H., 1962b, Lunar stratigraphic nomenclature: USGS 

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Shoemaker, E.M., Hait, M.H., Swann, G.A., Schleicher, D.L., 

Schaber, G.G., Sutton, R.L., Dahlem, D.H., Goddard, E.N., 

and Waters, AC., 1970, Origin of the regolith at Tranquillity 

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Shoemaker, E.M., Williams, J.G., Helin, E.F, and Wolfe, R.F., 

1979, Earth-crossing asteroids: Orbital classes, collision 

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Short, N.M., 1975, Planetary Geology: Englewood Cliffs, N.J., 

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Short, N.M., and Forman, M.L., 1972, Thickness of impact crater 

ejecta on the lunar surface: Modern Geology, v. 3, no. 2, p. 

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Shorthill, R.W, 1973, Infrared atlas of the eclipsed Moon: The 

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Shorthill, R.W, and Saari, J.M., 1969, Infrared observation on 

the eclipsed Moon: Seattle, Boeing Scientific Research Laboratories Document Dl-82-0778, 73 p.

Silver, L.T., 1971, U-Th-Pb isotope systems in Apollo 11 and 12 

regolithic materials and a possible age for the Copernican 

impact [abs.]: Eos (American Geophysical Union Transactions), v. 52, no. 7, p. 534.

Simonds, C.H., 1975, Thermal regimes in impact melts and the

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Simonds, C.H., Floran, R.J., McGee, P.E., Phinney, W.C., and 

Warner, J.L., 1978a, Petrogenesis of melt rocks, Manicouagan 

impact structure, Quebec [Canada]: Journal of Geophysical 

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Simonds, C.H., Phinney, W.C., McGee, P.E., and Cochran, Ann, 

1978b, West Clearwater, Quebec impact structure, part I: 

Field geology, structure and bulk chemistry: LPSCP 9, v. 2, p. 

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Simonds, C.H., Phinney, W.C., and Warner, J.L., 1974, Petrography and classification of Apollo 17 non-mare rocks with 

emphasis on samples from the Station 6 boulder: LSCP 5, v. 1, 

p. 337-353.

Simonds, C.H., Phinney, W.C., Warner, J.L., McGee, P.E., 

Geeslin, Jill, Brown, R.W, and Rhodes, J.M., 1977, Apollo 14 

revisited, or breccias aren't so bad after all: LSCP 8, v. 2, p. 

1869-1893.

Simonds, C.H., Warner, J.L., and Phinney, W.C., 1973, Petrology 

of Apollo 16 poiklilitic rocks: LSCP 4, v. 1, p. 613-632.

 1976a, Thermal regimes in cratered terrain with emphasis on the role of impact melt: American Mineralogist, v. 61, 

no. 7-8, p. 569-577.

Simonds, C.H., Warner, J.L., Phinney, W.C., and McGee, P.E., 

1976b, Thermal model for impact breccia lithification: Manicouagan and the Moon: LSCP 7, v. 2, p. 2509-2528.

Sjogren, W.L., and Smith, J.C., 1976, Quantitative mass distribution models for Mare Orientale: LSCP 7, v. 3, p. 2639-2648.

Sjogren, W.L., and Wollenhaupt, W.R., 1976, Lunar global figure 

from mare surface elevations: The Moon, v. 15, no. 1-2, p. 

143-154.

Sjogren, W.L., Wimberly, R.N., and Wollenhaupt, W.R., 1974, 

Lunar gravity via the Apollo 15 and 16 subsatellites: The 

Moon, v. 9, no. 1-2, p. 115-128.

Smith, E.I., and Hartnell, J.A., 1978, Crater size-shape profiles 

for the Moon and Mercury: Terrain effects and interplanetary 

comparisons: The Moon and the Planets, v. 19, no. 4, p. 

479-511.

Smith, E.I., and Sanchez, A.G., 1973, Fresh lunar craters: Morphology as a function of diameter, a possible criterion for 

crater origin: Modern Geology, v. 4, no. 1, p. 51-59.

Smith, J.V., 1974, Lunar mineralogy: A heavenly detective story. 

Presidential address, part I: American Mineralogist, v. 59, 

no. 3-4, p. 231-243.

 1979, Mineralogy of the planets: A voyage in space and 

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especially the Earth and Moon: Journal of Geology, v. 90, p.

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Smith, J.V., and Steele, I.M., 1975, Lunar mineralogy: A heavenly 

detective story. Part II: American Mineralogist, v. 61, no.

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Smith, J.V., Anderson, A.T., Newton, R.C., Olsen, E.J., Wyllie, 

P.J., Crewe, A.V., Isaacson, M.S., and Johnson, D., 1970, 

Petrologic history of the moon inferred from petrography, 

mineralogy, andpetrogenesis of Apollo 11 rocks: LSCP 1, v. 1, p. 

897-925.

Soderblom, L.A., 1970, A model for small-impact erosion applied 

to the lunar surface: Journal of Geophysical Research, v. 75, 

no.14, p. 2655-2661.

Soderblom, L.A., and Boyce, J.M., 1972, Relative ages of some 

near-side and far-side terra plains based on Apollo 16 metric 

photography, pt. A of Photogeology, [chap.] 29 of Apollo 16 

preliminary science report: NASA Report SP-315, p. 29-2 to 

29-6.

 1976, Distribution and evolution of global color provinces 

on the moon [abs.]: LS VII, pt. 2, p. 822-824.

Soderblom, L.A, and Lebofsky, L.A, 1972, Technique for rapid 

determination of relative ages of lunar areas from orbital 

photography: Journal of Geophysical Research, v. 77, no. 2, p. 

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Soderblom, L.A, Arnold, J.A, Boyce, J.M., and Lin, R.P., 1977, 

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 1978, The nature of isostasy on the moon: How big a Prattfall for Airy models?: LPSCP 9, v. 3, p. 3499-3511.

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6579-6609.

Wilbur, C.L., 1978, Volcano-tectonic history of Tsiolkovskij [abs.]: 

LPS IX, pt. 2, p. 1253-1255.

Wilhelms, D.E., 1964, Major structural features of the Mare 

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1964, pt. A, p. 1-16.

 1968, Geologic map of the Mare Vaporum quadrangle of 

the Moon: USGS Map 1-548 (LAC-59), scale 1:1,000,000.

 1970a, Geologic map of Apollo landing site 1: USGS Map 

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 1970b, Summary of lunar stratigraphy telescopic observations: USGS Professional Paper 599-F, p. F1-F47.

 1972a, Geologic map of the Taruntius quadrangle of the 

Moon: USGS Map 1-722 (LAC-61), scale 1:100,000.

 1972b, Geologic mapping of the second planet: Flagstaff, 

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 1972c, Reinterpretations of the northern Nectaris basin, 

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science report: NASA Report SP-315, p. 29-27 to 29-30.

 1976, Secondary impact craters of lunar basins: PLSC 7, v. 

3, p. 2883-2901.

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east side of the Moon: USGS Map 1-948, scale 1:5,000,000.

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near side of the Moon: USGS Map 1-703, scale 1:5,000,000.

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539-562.

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1:5,000,000.

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Professional Paper 785, 26 p.

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David, 1972a, Geology of the Sierra Madera cryptoexplosion 

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599-H, p. H1-H42.

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Jackson, E.D., 1972b, Geologic setting and petrology of Apollo 

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PLSC 8, v. 1, p. 515-535.

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 1972a, Fragments of terra rock in the Apollo 12 soil samples and a structural model of the Moon: Icarus, v. 16, no. 3, p. 

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 1972b, Thermal history and early magmatism in the 

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 1973, Bombardment as a cause of lunar asymmetry: The 

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 1975a, Glass compositions as a clue to unsampled mare 

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 1975b, Lunar petrogenesis in a well-stirred magma ocean: 

PLSC 6, v. 1, p. 1087-1102.

 1975c, The moon: Scientific American, v. 233, no. 3, p. 

92-102.

 1975d, The nature and origin of boulder 1, station 2. Apollo 

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Lunar anorthosites and a geophysical model of the moon: 

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 1977, The lunar impact flux, radiometric age correlation, 

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3457-3473.

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the interpretation of lunar mare stratigraphy and relative 

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The Moon, v. 10, no. 1, p. 17-50.
293

SUPPLEMENTARY TABLE. Craters superposed on basins

[Diameters of superposed craters larger than 100 km are listed individually. A/K, Al-Khwarizmi/King; Ap, Apollo; B, Birkhoff; C, Crisium; F-S, Freundlich-Sharonov; 

Hb, Hunboldtianum; Hu, Humorum; Hz, Hertzsprung; I, Imbrium; K, Korolev; K-H, Keel er-Heaviside; L, Lorentz; M-R, Mendel -Rydberg; Md, Mendeleev; Mo, Moscoviense; 

N, Nectaris; 0, Orientale; Sc, Schrodinger; Sm, Smythii. Photographic coverage: G, good; GP, half of area good and half poor (Orientale, two-thirds good); P, 

poor; M, numbers of craters probably diminished by mutual obliteration; R, most reliable data because of large sample, good photographs, and absence of niutual obliteration or blanketing]

Pre-Nectarian

Crater count

(/10 6 km*).

Coverage - 

Diameter

(km)

>100

99

98

97

95

95

94

93

92

91

90

89

88

87

85

85

84

83

82

81

80

79

78

77

76

75

74

73

72

71

70

69

68

67

66

65

64

63

62

61

60

59

58

57

56

55

54

53

52

51

50

49

48

47

46

45

44

13

42

41

40

39

38

37

36

35

34

33

32

31

30

29

28

27

26

25

?4

23

22

21

20

- A-K

0.320

- G, M

118,

115,

106

13

_

-

1

-

_

-

-

-

-

-

_

-

-

-

-

-

1

-

1

-

_

-

1

-

1

-

1

1

-

-

1

1

1

1

-

-

-

2

1

1

_

-

1

-

1

1

2

1

-

2

1

2

1

-

-

-

1

2

2

1

1

1

1

-

2

1

-

2

2

1

1

1

2

-

4

1

3

-

2

2

K-H

0.371

G, M

177,

172,

156,

7,

135,

126,

122,

112

_

-

_

_

_

-

_

-

-

-

1

-

-

-

-

-

1

1

1

-

_

-

-

-

-

-

2

-

1

-

1

-

_

-

_

-

2

-

-

-

_

1

1

-

-

1

-

]

-

1

_

1

_

3

-

-

_

1

-

]

2

4

1

-

1

1

-

1

-

3

2

3

2

4

1

3

3

S

2

2

Sn

0.445

P, M

152,

125,

110,

106

_

-

_

.

1

-

1

-

-

-

_

-

-

-

-

-

-

-

_

-

1

2

_

-

_

_

1

-

_

-

_

-

1

-

1

1

_

-

1

2

1

1

_

-

-

1

3

2

1

-

_

1

1

1

1

-

2

2

_

1

2

-

-

1

5

1

1

-

1

1

3

5

2

3

?

i

2

3

4

3

L

0.208

P, M

123,

120

_

.

.

.

_

1

_

-

1

-

.

1

-

-

.

-

-

-

-

-

_

-

_

.

.

1

-

_

-

-

_

2

_

-

_

-

1

-

_

-

_

.

1

-

.

_

-

-

-

-

1

-

_

-

-

1

2

_

_

-

_

1

-

-

-

2

-

-

1

3

1

1

-

-

2

-

1

1

2

4

B

0.401

G, M?

164,

148,

135,

117

_

.

_

_

_

1

1

-

_

-

.

-

-

-

-

-

_

-

-

-

1

-

1

.

_

.

.

-

.

-

_

1

1

-

_

-

_

-

-

-

_

2

.

-

.

1

-

-

-

-

1

1

.

2

_

1

1

1

.

?

1

1

3

1

4

-

-

2

-

3

_

1

3

3

3

1

1

-

1

-

F-S

0.629

P

110

1

-

.

.

_

-

1

1

-

-

_

-

-

-

-

-

-

-

-

2

_

-

-

1

-

i

-

-

.

1

1

1

.

.

_

-

_

1

1

-

1

2

.

2

-

.

2

-

-

2

_

-

.

2

1

_

2

1

1

-

2

-

2

3

2

-

2

3

1

7

_

3

5

3

H

2

4

3

1

5

Ap

0.480

GP

209

_

-

.

.

.

-

-

-

-

-

_

-

1

-

-

-

1

_

-

-

_

-

_

_

.

.

1

_

.

-

_

1

1

_

1

_

-

_

1

-

1

_

_

_

1

_

2

.

1

-

1

-

.

-

-

1

.

.

1

2

1

3

1

-

4

-

5

-

4

2

_

4

1

-

1

3

3

2

3

2

N

1.286

G R

126,

117,

100

2

.

.

_

.

-

-

-

-

1

_

-

2

-

1

-

-

_

1

-

_

-

_

_

.

_

1

_

2

-

1

_

1

-

1

_

1

_

-

-

_

_

.

_

.

1

-

.

2

1

2

1

3

-

2

5

1

2

2

4

1

2

_

2

1

3

2

8

4

2

4

6

4

3

2

5

2

2

2

4

M-R

0.247

P

_

_

_

.

_

-

-

-

-

-

1

-

-

-

-

-

_

-

_

-

_

_

_

_

_

.

_

_

-

1

1

_

_

_

_

.

_

.

.

-

_

.

.

.

_

1

-

_

_

1

.

-

_

-

_

]

_

.

_

1

_

1

1

1

1

-

-

-

_

2

_

1

1

-

1

-

.

_

1

1

K

1.113

'

.

_

.

_

_

1

.

2

-

-

.

_

_

_

_

-

-

1

.

-

_

.

_

.

_

.

_

1

1

-

1

1

_

.

.

.

.

1

1

-

_

2

3

1

1

1

_

3

1

-

1

2

_

1

1

2

_

.

.

2

.

2

2

2

3

1

1

2

i

4

5

]

5

_

4

-

7

7

8

2

Nectarian

Mo

0.609

P

.

_

_

_

_

-

.

_

_

-

.

-

_

_

-

-

-

.

_

-

_

1

1

.

.

.

.

.

_

-

.

1

_

1

.

.

1

_

.

-

_

.

_

1

_

.

-

.

_

1

1

.

_

-

_

1

_

_

_

1

1

1

3

2

2

_

1

1

1

2

4

3

_

1

2

3

4

3

5

2

Hb

0.515

GP

127

_

.

_

.

.

-

_

-

-

-

_

-

-

-

-

-

-

_

1

-

_

_

-

_

1

.

-

.

-

1

_

_

_

_

1

.

1

1

_

-

2

_

1

.

.

_

1

_

_

-

1

_

_

-

_

_

.

.

_

-

1

-

-

-

-

_

1

2

1

-

_

_

1

1

2

2

_

1

2

2

Md

0.569

122,

103

_

.

.

_

.

-

-

-

-

-

_

-

-

-

-

-

1

-

_

-

_

.

_

_

1

_

-

.

-

-

1

_

_

2

_

.

_

_

_

1

_

.

_

.

_

.

-

.

_

-

.

-

2

-

1

_

_

.

2

3

1

_

-

_

_

-

1

7

.

I

1

4

2

3

.

1

1

2

1

-

C

0.843

GP

138

_

.

_

_

_

-

.

-

-

-

_

-

-

-

-

-

1

-

_

-

_

-

-

2

.

_

-

_

_

-

_

.

1

_

_

.

_

_

.

1

_

_

_

2

.

.

1

1

-

-

.

_

2

-

.

_

1

3

.

1

.

2

.

1

-

1

1

3

-

1

2

_

.

2

2

1

4

4

1

3

Hu

0.428

110

_

_

_

.

_

-

_

-

_

-

_

-

-

-

-

-

_

1

_

-

_

-

-

-

_

_

_

.

.

-

_

.

_

.

_

.

-

1

.

-

_

-

.

1

.

.

-

_

-

-

_

_

1

1

.

.

_

_

1

1

1

_

1

_

1

_

-

2

1

-

_

1

.

_

1

?

2

1

1

2

Hz

0.883

G j R

105

_

-

_

.

_

1

_

-

1

-

_

-

-

-

1

-

_

-

_

-

_

-

_

-

.

.

-

1

.

-

_

-

.

_

_

.

_

.

_

-

_

-

1

-

1

.

-

.

-

-

1

-

.

1

2

1

_

.

1

1

1

1

.

1

3

1

2

3

1

1

1

1

5

-

1

1

4

4

4

2

I

2.491

G D , K

260,

105,

102,

100

1

-

.

_

_

-

-

1

-

-

_

-

-

1

-

-

-

1

-

-

_

-

-

-

_

.

-

-

-

-

1

-

_

1

1

-

_

.

_

-

_

-

.

2

1

_

-

_

-

-

.

1

.

1

-

_

1

1

.

3

3

5

_

1

2

1

2

4

2

1

2

2

2

_

4

3

4

/I

3

3

Imbrian

Sc 0

0.594 3.001

G fiP R *

.

_

.

_

.

_

-

.

.

-

i

_

-

.

i

i

i

i

-

.

-

.

-

.

-

_

.

.

.

-

2

_

-

-

-

-

-

2

-

-

-

_

1

_

-

-

3

1

2

2

1 3

_

-

-

-

-

_

.

2

1

3

2

1 1

-

3

4

1 1

-

1

-

1

_

1 1

1 4

1

1 2

3

1 ?

4

5

1 9

INDEX

[Italic page numbers refer to principal discussions]

Abbe, pis. 4, 9 

Abulfeda, 41, 180, pi. 7 

Abulfeda chain, 41, 115 

Adams, 167, 180, pi. 7 

Aeon, defined, viii, 19

Aeronautical Chart and Information Center, 124 

Age, meaning, 19-22, 49, 129, 139, 163, 277 

Ages, relative, 121-136. See also Crater dating, Crater frequencies, 

DL method, DL values, Superpositions, and individual 

series and systems

Agglutinates, 95, 96, 256, 266, 279. See also Regolith 

Agrippa, 205, 253, 266, pi. 10 

Airy, 155, 244, pi. 5 

Airy model of isostasy, 143 

Aitken, 120, pis. 4, 9 

Al Biruni, pis. 4, 9 

Al-Khwarizmi, 154, 180, pi. 7 

Al-Khwarizmi/King basin, 65, 148, 154, pis. 3, 6

crater frequency, 136, 146, 148, 157, 160 

Albategnius, 10, 213, 214, 219, pis. 7, 8 

Albedo, 2, 3, 49, 85, 94, 256, 265-267

color, 96, 97

magnetism, 215, 256

rays, 94-96, 249, 250, 256, 265-267, 279

slopes, 48, 49, 96 

Aliacensis, 48, 127, 180, 189, pi. 7 

Alkalis, mare basalt, 101, 276

Alpes Formation, 15, 73, 81, 82,119,168,171,174, 200, 203, pi. 8 

Alphonsus, 10, 89, 90, 94, 96, 128, 213, pis. 5, 7 

Alpine Valley, 113-116, 119, 204 

ALSEP, defined, 23, 165, 240 

Altai scarp, 164, 233 

Altimetry, Apollo and Zond, 77, 120, 145 

Alumina (aluminum)

Apollo 16 units, 165, 168

mare basalt, 101-103

orbital sensing, 98-103, 139-144, 165, 168

terra crust, 139-144

Amundsen-Ganswindt basin, 7, 65, 148, pis. 3, 6 

Anaxagoras, 51, 204, pi. 11 

AndSl M, 218 

Anorthite, 140

Anorthosite and anorthositic norite, 140-143, 156, 165, 168 

ANT suite, 140-143, 171, 201 

Antipodes, basins, 67, 76, 82, 181, 215-217, 279 

Antoniadi, 7, 58, 65, 81, 231, pis. 4, 9 

Apennine Bench Formation, 8, 73, 113, 193, 197, 231

age, 198, 212, 224, 243, 278

geologic maps and sections, 15, 19, 200, 243, 276, pi. 8

samples, 197, 224, 278, 279 

Apenninian Series, 123-125 

Aphanitic texture, 23, 174, 177, 178 

Apianus, 189, pi. 7 

Apollo basin, 55, 60, 64, 79-82,103,145, 147,160, 178-181, pis. 3, 6

mare filling, 103, pis. 4, 9

Apollo missions, 12, 19-23, 82, 101, 140, 163, 195, 229, 235, 249, 

277-280, pis. 2-4, 7-11. See also individual missions

unvisited sites, 125, 131, 254, 255, 258, 259 

Apollo photographs, types, 4, pi. 2 

Apollo 8, 12, 145, 152, 160 

Apollo 10, 12 

Apollo 11

crater frequency and DL values, 230, 235, 238

crater morphology, 135

landing site, 12, 13, 99, 101, 169, 227, 235, pis. 2, 4, 9

mare samples, 19, 85, 101, 224, 235, 245, 276-279

regolith, 235

spectral reflectance, 96, 99

terra samples, 140, 142 

Apollo 12

crater frequency and DL values, 230, 259, 271-273

crater morphology, 135

landing site, 12, 99, 101, 132, 207, 259, 272, pis. 2, 4, 10

mare samples, 19, 23, 101, 102, 249, 259, 276-279

objectives, 259

spectral reflectance, 96, 99, 261

terra and Copernicus-ray samples, 140, 269, 270, 278 

Apollo 13, 12, 195 

Apollo 14

crater frequency, 136, 230

Fra Mauro Formation samples, 82, 195, 200, 204, 276-279

landing site, 10, 12, 22, 82, 144, 204, pis. 2, 3, 8

mare-basalt clasts, 101, 190, 200, 278

objectives, 195, 204

sampling stations, 190, 200, 204, 205, 209

terra samples, general, 19, 23, 140-144

Apollo 15

Apennine Bench Formation samples, 197, 224, 278, 279

crater frequency and DL values, 136, 197, 230, 231, 252, 262, 

271-273

landing site, ii, 8, 9, 12-15, 22, 88, 91, 99, 101, 193, 200, pis. 

2-4, 8, 9

mare samples, 19-23, 101, 102, 237, 243, 276-279

massif samples, 19, 23, 49, 82,140-144,195, 198, 224, 276-279

objectives, 88, 198

sampling stations, 201, 237, 243, 244

spectral reflectance, 96, 99 

Apollo 16

crater frequency, 197, 219, 222, 230

crustal thickness, 12, 143

landing site, 12, 21, 22, 101, 164, 218, pis. 2, 3, 7, 8

mare samples, 101, 237, 238, 245, 276, 278

objectives, 21, 195

sampling stations, 165, 168-170, 237

significance, 21, 57, 72, 127, 146, 164, 216

spectral reflectance, 97

terra samples, 21, 57, 82, 140-144, 163, 216-220, 276, 278 

Apollo 17

crater frequency and DL values, 230, 273

landing site, 9,12, 22, 93, 99,101,105,109, 171, 239, pis. 2-4, 

7-9

mare and dark-mantle samples, 101, 102, 237, 239, 245, 

276-279

massif samples, 21, 46, 49, 82, 140-144, 156, 157, 163, 173, 

190, 195, 201, 270, 276, 278

objectives, 173

sampling stations, 174, 176, 178, 190, 237, 239, 240

spectral reflectance, 96-99, 240 

Apollonian metamorphism, 212 

Apparent crater, 43 

Arago, 28, 114 

Aratus CA, 90 

Archimedes, 8, 15, 19, 21, 113, 124, 125, 156, 193, 216, 221, 231,

pi. 9

Archimedian Series, 123, 124 

Archytas, 204, pi. 10 

Ariadaeus B, 220, 221

Aristarchus, 50, 91, 121, 125, 126, 156, 252, 253, 263, pi. 11 

Aristarchus Plateau, 86, 89, 91, 125, 233, 240, 244, pi. 4 

Aristillus, 8, 9, 22, 48, 253, 269, pi. 11 

Aristoteles, 9, 202, 252, pi. 10 

Artamonov, pis. 4, 9 

Arzachel, 10, 128, 221, pis. 5, 8 

Asclepi, 41, pi. 6 

Ash-flow tuff, 21, 85 

Asthenosphere, 115 

Astroblemes, 77 

Astrogeology, vii 

Atlas, 9, 32, 166, 231, pis. 5, 9 

Augite, 140, 143, 243

Australe basin, 7, 62, 64, 103, 145, 148, 179, 245, pis. 3, 6 

Autolycus, 8, 9, 22, 200, 253, 269, pi. 11

B

Baco, 75, pi. 6

Bailly basin, 11, 59, 64, 80, 179, 180, pis. 3, 7

Balmer-Kapteyn basin, 7, 65, 148, 165, 167, 190, pis. 3, 6

Barbier, 149, 151

Barocius G, 37

Barringer, 181, pi. 7

Barrow, 149, 199, 204, pi. 6

Basalt. See Fra Mauro basalt, KREEP basalt, Mare-basalt samples

terrestrial, 101, 102 

Base surge, 43

Basin-mare distinction, 3, 19, 85, 86 

Basin materials, 55, 276-279, pis. 3, 6-8,12. See also basins listed

on p. 64, 65, 148, 179 

ages, relative, 64, 65, 148, 157, 179, 190, 224, 276-279, pis.

3, 6-9, 12 

ages, absolute, 157,160,168-171,177,178,186,190,191, 200,

201, 212, 224, 276

asymmetry, 66, 81, 82, 127, 171, 173, 179, 276, pis. 3, 6-8 

crater frequencies, 148, 149, 160, 179, 186, 293 

degradation, 65, 81, 82, 145, 146, 163-165 

dunelike deposits ("deceleration dunes"), 67, 71,130, 204, 218,

219

ejection angles, 67, 73, 80, 81 

emplacement, 66-82, 174-177, 195, 202-204, 210-212,

218-220, 276

excluded zone, 66, 164, 165

extent, 64-82, 123, 127, 129, 145-147, 171, 179, 190, 195, 

213-215, 218, 276, pis. 3, 6-8, 12

Basin materials Continued

Imbrian, 178, 179, 195, 279, pis. 3, 8, 12. See also Imbrium

basin, Orientale basin, Schrodinger basin 

impact melt, 66-70, 73, 76, 77, 82, 203, 211, 212, 276 

knobby, 73, 81, 82, 168, 171, 172, 276 

lateral transitions, 66, 67, 72, 73, 81, 82, 163, 164, 190, 204,

205, 214, 215, 218, 219, 224 

lineate, 66, 67, 80-82, 127, 147, 150, 163, 164, 179, 180, 188,

189, 204, 206, 209 

mapping conventions, 127-129, pi. 6 

massifs, 6-11, 49, 66-73, 76-82,105,127-129,145,146,152,

155, 171-178, 181, 193, 198-203 

Nectarian, 161, 276-278, pis. 3, 7, 12 

pre-Nectarian, 139, 143, 178, 180, 277, 278, pis. 3, 6, 12 

sampling points, 82, 144, 163, 171, 195, 276, pis. 3, 7, 8. See 

also Apollo 14, Apollo 15, Apollo 16, Apollo 17, Luna 

20

secondary-crater relations. See Basin-secondary craters 

stratigraphic mapping conventions, 121-129, pi. 6 

superpositions. See Superpositions 

thickness, 66, 82 

Basin rings, 57-73, 77, 110, 276, 280, pi. 3. See also Basin materials,

massifs

external, 78-81

Basin-secondary craters, 32, 67, 74, 81,137,151,159,160,181-186,

196, 213-215, 276. See also Imbrium basin, secondary

craters; Nectaris basin, secondary craters; Orientale

basin, secondary craters

distribution and extent, 32, 64,127, 163,171, 213-215, pis. 3,

6-8

ejecta relation, 42, 66, 67, 72, 81, 82, 211-213, 217-220, 276 

linearity and radiality, 3,32, 67, 80, 81,127,147,161-170,179,

186-188, 196, 213-216, 225 

mapping conventions, 127-129, pi. 6

plains relation, 67, 71, 72, 128, 146, 164, 190, 215-220, 224 

provenance of material, 42, 43, 81, 82, 129, 211, 212, 217-220,

276

"rays," 74, 75, 82, 127 

stratigraphic relations, 127, 145-151, 158, 161-167,175, 179,

196, pis. 6-8

Basins (ringed basins, multi-ringed basins), 3, 19-23, 55, 276, pis. 

3, 6-8. See also other basin headings and basins listed 

on p. 64, 65, 148, 179. 

antipodes, 67, 76, 82, 181, 215-217, 279 

craters, comparisons, 57, 58, 65, 66 

crater-uplift control, 145, 149 

crustal composition, 143-145 

defined, 3, 57, 65 

deposits. See Basin materials 

depth, 77, 80, 143-145 

ejecta. See Basin materials 

excavation cavity, 66, 77-81 

gravity, 77, 79, 117

inventory, 57, 64, 65, 81, 148, 179, pi. 3 

mare-extrusion control, 102, 103, 145, 229, 238-245 

Moho control, 77, 80, 102, 103, 115, 120, 143-145, 229,

240-245, 276 

rings. See Basin rings

secondary-impact craters. See Basin-secondary craters 

size series, 59-65 

topographic rim, defined, 66 

topography, 77, 100, 280 

Bedrock, 12, 13, 21, 45 

Beer, 8, 113 

Bench Crater, 259, 261 

Berosus, 166, pi. 7 

Berzelius, 175 

Bessarion B, 39 

Bianchini, 34, pi. 9

Big Backside Basin. See South Pole-Aitken basin 

Birkhoff basin, 59, 64, 147, 148, 158, 179, 188, 293, pis. 3, 6 

Birkhoff X, 159

"Black-and-white rocks," 23, 165, 203, 224, 278 

Blancanus, pi. 7 

Blanchinus, 149, 189

Blocks, ejected, 29, 46, 47, 99, 168, 258, 265 

Bonpland, 22, 189, 208 

Bonpland D, 112, 221 

Boscovich Formation, 205 

Boulder tracks, 174, 178 

Bowen's reaction series, 140, 142 

Boyce-Dial crater-dating technique, 133. See also DL method, DL

values

Brayley, 253 

Breccia, 12, 45-47

complexities, 21-23, 45-47, 177 

dating, 22, 47, 156, 169, 177, 178 

dikes, 44, 46

295
296 THE GEOLOGIC HISTORY OF THE MOON

Breccia Continued

dilithologic (dimict), 46, 165, 168

friable fragmental feldspathic, 165-171

granulitic, 165, 168-170, 174

lens, crater-floor, 43

melt-poor, 46, 47, 165, 168

melt-rich, 47, 82, 157, 165, 168

mixing and recycling, 21-23, 139-143, 157, 163

provenance, 22, 23, 47, 205, 211, 212, 217-220, 276

regolith, 45

samples, 45-47

shock grades, 44-47

terrestrial, 45

textures, 23, 45-47, 171, 174 

"Bright swirls," 76, 215, 216, 256, 258, 269 

Bullialdus, 252, 253, 256, pi. 10 

Bunte Breccia, 45, 46, 73 

Burg, 253, pi. 11 

Buys-Ballot, pi. 4

c

Calcium, 101, 102, 139, 140

Calderas, 17, 32, 40, 88, 113

Camelot Crater, 239, 270

Campanus, 112, pi. 8

Campbell, 58, 257, pis. 4, 6, 9

Cardanus, 197, 231, 234, pi. 9

Carnot, 159, 180, pi. 7

Caroline Herschel, 8, 32

Carpenter, 252, pi. 11

Cassini, 8, 202, 216, 221, pi. 8

Cataclastic anorthosite, 141, 165, 168

Cataclysm, 190, 191, 278

Catena (pi., catenae), defined, 3

Catena Artamonov, 156

Catena Dziewulski, 156

Catena Mendeleev, 188

Catharina D, 37

Cauchy, 119

Cauchy structures, 115, 119, pi. 5

Cavalerius, 256, 258, pi. 10

Cayley, 220

Cayley Formation (Cayley plains), 21, 165, 195, 216, 224, pi. 8

crater frequency, 197, 219, 221, 230

samples, 165-170

thickness, 220

type area, 218-221 

Central Cluster, 240, 270 

Central peaks. See Crater materials, peak; Crater processes, peak

formation

Chao, E.C.-T., vii, 42, 43 

Chaplygin, 6, 120, 180, pi. 7 

Chemistry

mare basalt, 101, 102, 276

orbital, 97, 139, 142, 143, 156, 190, 198, 241, 245, 258, pi. 2

terra crust, 139, 140 

Chretien, pis. 4, 9 

Cichus, 253, pi. 10 

Clavius, 58, 74, 180, pis. 7, 8 

Clearwater craters, Quebec, 32, 45 

Cleomedes, 166, 180, pis. 4, 5, 7 

Clinopyroxene, 101, 102, 140 

Collapse craters, 87-90 

Colluvium, 49, 174, 201, 202 

Color, 86, 96-99. See also Spectral reflectance 

Compass directions, lunar conventions, 2 

Compton, 58, 65, 81, 166, 186, 199, pis. 4, 5, 8 

Cone Crater, 190, 205, 209-212, 269, 270, 273 

Cones, 86-89

Congreve and Congreve U, 180, 184, pi. 7 

Conon, 200, 203

Continuous deposits, defined, 43 

Copernican Period

defined, 221, 265

duration, 249, 269-273, 280

impact rate, 246, 271-273, 280 

Copernican System, 263, 280, pi. 11

chronology, 269, 280

crater frequency and DL values, 130, 136, 249-253, 256, 265

crater materials, 130, 236, 249, 252, 253, 263, 279, 280, pi. 11

defined, 121-125, 249, 265

distribution, 265, 280, pi. 11

mare materials, 196, 265, 269, pis. 11, 12

remote-sensing properties, 265-267

structures, 269, 271 

Copernicus, 1, 8, 29, 49, 52, 92, 250, 265, 269, pi. 11

age, 269-273, 278, 280

composition, 156, 269, 270

crater frequency and DL value, 136, 252, 253, 266, 271-273

depth-diameter ratio, 80

impact energy, 29

impact melt, 49, 50, 52, 76

rays, 1, 29, 30, 48, 250, 269-272

secondary craters, 8, 29, 125, 126, 247, 250, 270

stratigraphic relations, 121, 125, 126, 247, 250, 265, 272 

Copernicus H, 31, 250, 265-268 

Cordillera ring. See Montes Cordillera 

Core, 12, 277, 280 

Coriolis L, 149, 151, pi. 6

Coulomb, 158, 159, pi. 7

Coulomb-Sarton basin, 60, 64, 72, 79, 145, 148, 158, 179, pis. 3, 6

Crater dating, 129-136

"counts." See Crater frequencies

DL method, 133, 136, 216, 250, 253, 265, 272, 273

morphologic basis, 129, 145, 149, 180, 222, 232, 249, 253-257

Pohn-Offield method, 129-131, 143, 145

superpositional basis, 17, 27, 47, 125. See also Superpositions

Trask's method, 131-135, 236, 243, 253

Crater frequencies, 18, 19, 129-136, 145-147, 180, 221, 230-232. 

See also individual Apollo and Luna missions, basins, 

formations, maria, series, and systems

basins, 136, 148, 160, 179, 186, 293

diameter dependence, 129, 132, 136, 257

DL correlations, 130, 136, 216, 230, 256, 273

endogenic craters, 29

secondary-impact craters, 29, 32, 132, 276

slopes of size-frequency curves, 29, 32, 129-134, 145, 146, 216, 

257

small-on-large superpositions, 129, 133-135, 145, 250-253, 

266, 267, 273

substrate influence, 231, 234, 239, 241, 250, 252, 269, 272

time-stratigraphic correlations, 130, 136

Crater materials, 3, 17-19, 25-53. See also individual craters, 

series, and systems, and craters listed on p. 149, 180, 

221, 231, 253, 265

asymmetric, 32, 38, 39, 258, 266, 267

dunelike, 27, 28, 268

endogenic interpretations, 17-19

erosional degradation, 19, 20, 27, 47, 48, 129, 145, 149, 180, 

185, 189, 222, 232, 249, 253, 256, 265-267

extent, 27, 127

floor, 27, 28, 43, 45, 50-52, 113, 116, 117

geologic mapping, 15, 27, 47-50, 67, 125, 127, pis. 6-12

herringbone, 28-35, 47, 48, 67, 125, 211

mixing, 44-48

morphology, 3, 27, 149, 180, 222, 223, 232, 253-256

origin, 17-19

peak, 27, 43-51, 58, 65, 77-80, 280

remote-sensing properties, 96-99, 265-267

rim, 27-29, 42, 45-48

shock zoning, 41, 42, 45-48

subunits, 27, 28, 45-53

wall, 27-29, 45-51 

Crater processes. See also Ejection process

floor uplift, 19, 58, 65, 93, 99, 113, 129, 145, 149, 196, 245, pi. 5

formation mechanics, 33, 40-45

formation times, 43, 45

impact melting, 44. See also Impact melt

interference, 29-32, 35, 37, 48, 220

origin, general, 17-19, 27, 32, 33, 41, 47

peak formation, 43, 44, 49, 50, 77

"push" and "pull" mechanisms, 43

terrace formation, 43, 44, 48, 77

thrusting, 42, 47 

Crater properties, 3, 17-19, 27-33

circularity, 29

depth-diameter ratios, 27, 29, 42, 43, 66, 80

excavation cavity, 41-43

morphology, 3, 27-29

simple-to-complex transition, 3, 27-29, 43, 65

size classes, 27-29, 57, 58

size-frequency distributions. See Crater frequencies

spatial distributions, 17-19, 29, 32, 139, 163, 265, 276, pis. 

6-12

target infuence, 43, 44, 47, 133, 239, 241 

Crater types

atypical, 19, 32, 41

collapse, 87-92

complex, 25-29, 43, 44, 47-53, 77, 78, 280

dark-haloed. See Dark-haloed craters

delta-rim, 32, 33, 92

dimple, 236

double-ring, 113, 116

endogenic. See Endogenic craters

explosion, 32, 77-81

gravity, 43, 81

hybrid, 33

irregular, 32

laboratory, 31-33, 38, 39

missile, 29, 39

nested, 78, 113, 116, 258

noncircular, 32

paired, 8, 32, 39

primary, defined, 17

ringed, 116

satellitic (secondary). See Secondary-impact craters

simple, 25-29, 33

smooth-rimmed, 32, 33, 40

strength, 43, 81 

Cratering flow, 41, 80 

Cratering rate. See Impact rate

Craters, 3. See also crater and cratering headings, individual 

craters, and craters listed on p. 149,180, 221, 231, 253, 

265 

Crisium basin, 7, 63, 65, 100, 103, 165, 170, 242, pis. 3, 5, 7

age, 171, 179, 186, 278

antipode, 215-217

composition, 144

crater frequency, 179, 186, 293

cross sections, 82, 192, 235, 243, 276

Crisium basin Continued

deposits, 65, 164-167, 171-175, 190, 238, 243, pis. 3, 7

impact, 82, 171

mare fill, 103. See also Mare Crisium

mascon, 117

rim and rings, 65, 78-81, 100, 103, 110, 117, 171, 214, 235, 

241, 242, pi. 3

samples, 82, 163, 171, 276, 278

secondary craters, 167, 171, 175

stratigraphic relations, 148, 165, 171, 179, 192, 235, 243, pis. 

7, 12

structures, tectonic, 112, 115

topography, 77, 78, 100 

Crocco, pis. 4, 9

Crookes and Crookes D, 185, pi. 11 

Cross sections

diagrammatic (schematic), 44, 45, 82, 100,102, 112, 115, 118, 

120, 144, 200, 210, 211, 222, 240, 274, 275

geophysical, 13, 79

local, 19, 200, 210, 238, 240, 243, 261

regional, 21, 144, 192, 235, 240, 243, 272, 274-276 

Crozier, 40

Cruger, 67, 71, 127, 130, pis. 4, 8 

Crust (terra)

age (time of solidification), 156, 157, 277, 280

basin effects, 77, 80, 102, 103, 115, 120, 143-145, 229, 

240-245, 276

composition, 139, 165, 201, 280

density, 12, 140, 143

differentiation, 142, 143, 156, 157

igneous and impact processes, importance, 140-143

layering, 12, 13, 99, 143, 274-276

mare-extrusion influence, 102, 103, 115, 145, 229, 238-245

mineralogy, 140

mixing, 139-143

petrology and petrogenesis, 140-143

thickness, 12, 77, 102, 103, 115, 143, 145, 278, 279

zonation, 143 

Cumulates, 141-143, 174 

Curie, 226, pi. 6 

Cyrano, 149, 151, pi. 6

D

"D caldera," 90, 203

Daedalus R, S, U, 149, 151, 231, pi. 6

D'Alembert, 257, pi. 7

Damoiseau, pis. 5, 9

Daniell, 40, pi. 5

Dark-haloed craters

endogenic, 89, 90, 266

exogenic (impact), 19, 176, 190, 191, 223, 266, 268, 279 

Dark-mantling materials, 3, 8, 9, 40, 89, 109, 113, 116, 117, 174, 

203, 205, 233, 239, 270, pi. 4. See also Glass

age, 174, 234-237, 241, 245, 269, 270, 278, 279

Apollo 17, 234, 237-241, 278

color (spectral reflectance), 96, 97, 240

composition, 101

distribution, 15, 89, 94, 95, 113, 174, 234, 240-245, 276, 279, 

pi. 4

mantle source, 102

origin and eruptive mode, 89, 102, 240, 241, 245, 276

radar properties, 99

source vents, 89-93, 102, 113, 234, 240, 276

stratigraphic relations, 93-95, 109, 233-235, 239, 240, 245,

276-279

Darwin, 130, pi. 6 

Davy, pis. 5, 9 

Dawes, 239, 252

Debris surge (flow). See Ground surge 

Deceleration dunes, 67, 130, 204, 218, 219 

Delaunay, 37

Delisle, 18, 19, 121, 125, 252, pi. 10 

Densities

magma, 102, 103, 241

mantle, 12

mare basalt, 12, 102, 103

Moon, 12

terra crust, 12, 143 

Depth-diameter ratio

basins, 77, 80

craters, 27, 29, 42, 43, 66, 80 

Descartes and Descartes A, 22, 218, pi. 7 

Descartes Formation (Mountains), 21, 165-171, 195, 216 

Deslandres, 10, 149, pi. 6 

Diameters, data source, 4

Differentiation, magmatic, 102, 142, 143, 156, 157 

Dikes and veins

breccia, 44, 46

igneous, 101, 102, 157, 276

impact melt, 44-47, 165 

Dionysius, 220, 221, 265 

Diophantus, 18, 19, 121, 125, 252 

DL method, 133, 136, 216, 250, 253, 265, 272, 273 

DL values

Copernican-Eratosthenian boundary, 250, 265

Copernican maria, 130, 136, 250, 269

craters, 250, 253, 256, 273

crater-frequency correlations, 130, 136, 216, 230, 256, 273

Eratosthenian-Imbrian boundary, 229-231, 243, 250
INDEX 297

DL values Continued

Eratosthenian maria, 125, 130, 136, 230, 231, 247-250, 253, 

259, 269, 273

farside maria, 245

Imbrian maria, 125, 130, 136, 229-233, 243, 245, 250, 273

Orientale basin, 216, 230

plains, 136, 216 

Dobrovol'skiy, 47 

Dollond B and C, 218 

Domes

intercrater, 29, 35, 127, 214, 215

mare, 86-89

terra, 110, 111, 127, 146, 214, 215, 222-225 

Doppelmayer, pis. 5, 7 

Doppelmayer Formation, 244 

Doppler, 184, pi. 8

Dorsum (pi., dorsa), defined, 3. See also Mare ridges 

Downslope movement, gradual, 49, 89, 96, 110, 111, 133, 280. See

also Colluvium, Slumping 

Dune Crater, 201, 237, 243 

Dunite, 143, 157, 174

E

Early Imbrian Epoch. See also Lower Imbrian Series

defined, 121, 196

duration, 224, 279

igneous activity, 190, 198, 224, 232, 238, 279

impact rate, 160, 246, 279 

"Eclipse temperatures," 99, 113, 265-268 

Einstein, 35, pi. 6

Ejecta, defined, 3. See also Basin materials, Crater materials 

Ejection process

angle, 41-43, 72, 80, 211

ballistic, 41-43, 47, 48, 211, 212

basins, 66, 67, 72, 73. 77-81

complex craters, 47

deposition, 42

nonballistic, 43, 46, 47, 202, 211, 212. See also Ground surge

overturned flap, 42, 73

range, 42, 48, 211, 212

sequence, 41, 42

target effects, 41, 47

velocity, 42, 43, 72, 203 

Elastic waves, 81 

Elbow Crater, 201, 237, 243 

Elements. See also individual elements and oxides

large-ion-lithophile (LIL), 140

major, 96-103, 139-144, 165, 280

radioactive, 98, 99, 102, 140, 190, 198

rare-earth. See Europium anomaly, KREEP, Rare-earth 

elements

siderophile, 101, 141, 143, 157, 158, 202

trace, 102, 140-142, 156, 276

volatile, 101 

Elevation

average, 12

Crisium (basin and mare), 77, 78, 99, 100

farside, 12, 120

maria, 93, 115, 120

nearside, 12, 120

Orientale (basin and mare), 77 

Endogenic craters, 87-93, 239

impact-crater distinction, 19, 27, 29, 32, 33, 47, 88, 89, 113, 266

size-frequency distribution, 29 

Endymion, 199, 166, 180, pis. 4, 7, 9 

Energy coupling, 41, 43, 81, 212 

Energy partitioning, 43, 44, 157, 212 

Energy scaling. See Scaling laws 

Engelhardt, 184, pi. 8 

Epigenes, 204, pi. 7 

Epoch, defined, 123 

Equilibrium. See Steady state 

Eratosthenes, 8, 121, 125, 193, 231, 247, pi. 10 

Eratosthenian Period

defined, 121, 125

duration, 249, 262, 280

impact rate, 246, 271-273, 280 

Eratosthenian System, 247, 279, pi. 10

chronology, 262, 279, 280

crater frequency, 130, 136, 230, 249-253, 256, 257, 265, 272, 

280

crater materials, 130, 236, 247, 265, 280, pi. 10

defined, 121-125, 249, 250

DL values, 125, 130, 136, 229-231, 247-250, 253, 256, 259, 

269, 273

mare materials, 15, 121-125, 131-135, 207, 230, 231, 234, 243, 

244, 247-256, 258, 265, 269, 272, 278, 279, pis. 10, 12

samples, 249, 259-262, 276-279

structures, 279 

Escape velocity, 29 

Esnault-Pelterie, 149, 159 

Eudoxus, 9, 202, 252, 253, pi. 11 

Euler, 8, 49, 121, 125, 126, 250-253, pi. 10 

Europium anomaly, 101, 142 

Excavation cavity

basins, 23, 66, 77, 80, 280

defined, 43

depth and volume, 42, 43, 66, 77, 80

excavation and growth, 41-43

Explorer 35, 12

Explosions, 41

Exposure ages, 235, 238, 269, 270

Fabricius, 147, 253, pi. 10 

Fabry, 4, pi. 6 

Farside

crustal and lithospheric thickness, 12-14, 115, 145

defined, 2, 3

maria, 3, 103, 115, 233, 245, 280

southern, 3

structures, 107, 112, 115, 245, pi. 5 

Faults, 107. See also Fractures, Rilles

craters, 42, 47, 113-118

mare ridges, 110-112

solitary, 113, 115

thrust, 115, 269 

Fecunditatis basin, 40, 65, 117, 145, 146, 148, 179, 235, 241, 243,

pis. 3, 6

Feldspar, crustal varieties, 140 

Felsite, 177, 178 

Fermi, 5, pi. 6

Ferroan anorthosite, 142, 143, 156 

Fersman, 158, pi. 7 

Feuillee, 8, 113 

Fizeau, 179, pi. 9 

Flammarion, 128, 213, pi. 6

Flamsteed and Flamsteed P, 111, 253, 254, 258, pi. 6 

Flamsteed-Billy basin, 65, 146, 148, 244, 251, 254, 258, pis. 3, 6 

Fleming, 4, 156, 180, 191, pi. 7 

Flynn Creek Crater, Tenn., 43 

Fontenelle, pis. 5, 8 

Formation, defined, 123 

Fowler, 158, 159, pi. 6 

Fra Mauro, 22, 89, 149, 208, pi. 6 

"Fra Mauro basalt," 140 

Fra Mauro breccia, 200, 211, 212 

Fra Mauro Formation, 204, 218, pi. 8

age, 200, 212, 224, 279

composition, 143, 144, 211, 269, 279

crater frequency, 136, 230

crater morphology, 131, 135

definition and type area, 73, 124, 125, 196, 208

distribution, 15, 204, 207, 213, pi. 8

emplacement process, 195, 204, 211, 212

Hevelius Formation analogy, 73, 82, 204, 212, 213, 216, 218

lateral gradations, 190, 204, 208, 213

samples, 195, 200, 204, 278, 279

stratigraphic relations, 127, 196, 207, 235, 272

stratigraphic significance, 19, 21, 121, 125, 127, 195, 196

thickness, 205, 210

Fra Mauro peninsula, 10, 22, 115, 262 

Fracastorius, pis. 5, 7 

Fractures (fissures, gashes). See also Faults, Rilles

crater-floor, 113-118. See also Crater processes, floor uplift

impact-melt, 113, 117

regolith, 269, 271 

Fraunhofer E and J, 147, 149 

Freundlich, pi. 7

Freundlich-Sharonov basin, 60, 64, 82, U7, 179, pis. 3, 6 

Furnerius, 149, 167, pis. 4-6

G

Gabbro, 140

Gabbronorite, 142

Gagarin, 6, 31, 58, 120, 149, 191, pi. 6

Galois and Galois Q, 183, pi. 6

Gambart, 40, 92, 208

Gambart A, 268

Gamma-ray spectrometer, 97-99, 156

Gargantuan basin. See Procellarum basin

Gassendi, 113, 116, pis. 5, 7

Gassendi A, 253, pi. 11

Gaudibert, 113, 117, pi. 5

Gauss, 58, 166, pis. 5, 7

Geminus, 166, pi. 10

Geo-, vii

Geochronology

breccia dating, 22, 47, 156, 169, 177, 178

decay constants, 19, 168

exposure ages, 235, 238, 269, 270

mare-basalt dating, 101, 235-244, 261, 269, 270

methods, 156, 177, 235-238, 269

model ages, 156

pyroclastic dating, 241, 244 

Geologic cross sections. See Cross sections 

Geologic mapping, vii, 17-21, 123-129

basins, 15, 66, 67, 73, 76, 127, 129, 169, 174, 207, pis. 3, 6-8,12

craters, 27, 47, 50, 176, pis. 6-12

maria, 86, 94, 95, 174, 236, pis. 4, 9-12

systematic program, vii, 123-125 

Geologic maps

local, 176, 236

regional, 15, 76, 94, 95, 169, 174, 207, 217, pis. 3-12 

Geologic units, vii, 17-23, 121-129. See also Rock-stratigraphic 

units, Time-stratigraphic units, Time units

Geologic units Continued

correlation of types, 121-125

correlation with samples, 21-23, 101, 139, 143, 276, 280

interfingering, 27, 121, 124

lateral continuity, 17, 21, 27

nomenclature, 121-125

superpositions, 17, 21, 125-127. See also Superpositions

three-dimensionality, 17 

Giordano Bruno, 191, 265 

Glass. See also Dark-mantling material

Apollo 15, green, 101, 237, 244, 245, 278

Apollo 15, red, 101, 244

Apollo 17, orange and black, 101, 237, 240, 241, 245, 278

coatings, 46

darkening, 95, 96, 280

devitrification, 240

regolith, 95, 96, 101 

Goclenius, 119, pi. 7 

Goddard A, 76, 265 

Godin, 205, 253, 266, pi. 11 

"Gold dust," 85 

Goldschmidt, 49, 204, pi. 6 

Gosses Bluff, Australia, 43 

Grabens. See Rilles 

Granite, 140, 223, 269 

Granodiorite, 140

Granulitic and granoblastic texture, 46, 157 

Granulitic breccia, 165, 168-170, 174 

Granulitic impactite, 141, 212 

Gravity, cratering effects, 43, 44, 81

Gravity anomalies, 77-80, 117. See also Mascons and individual 

basins and maria

craters, 77, 113 

Grimaldi (basin and mare), 11, 59, 64, 71, 77, 117, 120, 147, 148,

pis. 3-6, 9, 10

Grissom-White basin, 65, 148, 160, 179, pis. 3, 6 

Ground surge (debris flow, debris surge, ground flow, surface flow)

basins, 66, 164, 204, 211-213, 217-220, 222, 276

craters, 42, 43, 47, 48

secondary-crater relation, 42, 67, 211-213, 217-220, 222, 276 

Group, defined, 123 

Gruithuisen domes, 222, 225 

Gruithuisen K, 92

H

Hackman, R.J. See Shoemaker-Hackman stratigraphic scheme 

Hadley rille, 13, 88, 91, 200, 201, 237, 243 

Hahn, 166, pi. 8 

Hainzel A, 187, 253, pi. 10 

Halo Crater, 259, 261

Hansteen and Hansteen alpha, 225, pis. 5, 9 

Hausen, 28, 253, pi. 10 

Hayn, 199, pi. 11 

Head Crater, 259, 261 

Heat flow, 198

Heat sources, 102, 157, 190, 198, 212, 277-280 

Heaviside, 6, 120, 150, pi. 6 

Heis, 8, 32 

Helicon, 8, 32

Henbury craters, Australia, 29, 31, 32 

Hercules, 9, 32, 166, pi. 10 

Herodotus, 91, pi. 9 

Herschel, 128, 253, pi. 10

Hertzsprung basin, 60, 64, 79-82, 148, 149, 158, 159, 178-183, 188, 

pis. 3, 7, 8

crater frequency, 179, 180, 186, 293

elevation, 120

plains, 190, 245, pi. 8 

Hevelius, 67, 71, 120, 180, pis. 5, 7

Hevelius Formation, 66-73, 79-82, pis. 3, 8. See also Orientale 

basin

crater analogy, 57, 66, 79

crater frequency, 216

definition and type area, 66, 67, 71, 196, 197

extent, 66, 67, 76, 79-82, pi. 8

Fra Mauro Formation analogy, 73, 82, 204, 212, 213, 216, 218

inner facies, 66, 76, 81, 197

Montes Cordillera relation, 67-71, 79, 232

nonlineated facies (member), 70, 72

outer facies, 67, 76, 82

plains relation, 67, 164, 190, 215, 218, 219, 224

stratigraphic significance, 66, 121-125, 196

transverse facies, 67, 71, 73, 76, 204 

"Highland basalt," 140, 141, 156 

Hilbert, 4, 5, 154, 191, pi. 7 

Hipparchus, 19, 37, 213, pis. 6, 8 

Horrocks, 253

Hortensius-Milichius domes, 86 

Bubble, pis. 4, 7, 9

Humboldt, 7, 58, 102, 113, 167, 223, 226, 231, pis. 4, 5, 9 

Humboldtianum basin, 61, 64, 79-81, 117, 148, 166,179, 186, 199, 

pis. 3-5, 7

crater frequency, 179, 186, 293

impact, 79, 82, 171

stratigraphic relations, 148,171,175,179,186,199, 204, pi. 7 

Humorum basin, 57, 61, 64, 82, 148, 153, 179,187, 192, 214, pis. 

3-5, 7

crater frequency, 179, 186, 293

mascon, 117, 244
298 THE GEOLOGIC HISTORY OF THE MOON

Humorum basin Continued

structures, tectonic, 112, 117, 244, pi. 5 

Hypervelocity, 17, 40

I

Ilmenite, 95, 101, 140, 235

Ilmenite basalt, 101, 235, 238, 261

Imbrian-Imbrium distinction, 124

Imbrian Period. See also Early Imbrian Epoch, Late Imbrian

Epoch

defined, 124, 125 

duration, 245, 279 

impact rate, 160, 191, 245, 246, 279 

Imbrian System. See also Lower Imbrian Series, Upper Imbrian

Series

crater frequency, 136, 186, 197 

crater materials, 15, 131-136, 186, 221, 222, 232 

defined, 124, 125, 193, 196 

extent, 195, 229, 279, pis. 8, 9 

plains materials. See Plains 

type area, 196, 208, 231

Imbrium basin, 8, 19, 21, 36, 57, 63, 65, 79-82, 193, pis. 3, 8, 12 

age, 186, 191, 201, 212, 224, 277, 279 

antipode, 215-217 

center, 103, 198, 202, 204 

composition, 143, 144, 168, 201, 211

crater frequency, 135, 136, 160, 191, 197, 219-222, 230, 293 

deposits, 19, 57, 65, 127, 129, 164, 165, 168, 174, 175, 190,

193, 279, pis. 3, 8. See also Alpes Formation, Apennine Bench Formation, Fra Mauro Formation,

"Material of Montes Apenninus" 

depth, 143, 240

excavation cavity, 81, 197, 198, 204, 211 

fissured plains, 203 

impact angle, 82 

impact melt, 198, 203, 211, 212 

mapping conventions, 127, 129 

mare fill, 103, 240, 243-245, 258, 259, 269, 279. See also Mare

Imbrium 

mascon, 117

Orientale basin, comparisons, 196, 197, 202, 203, 213, 218 

plains, 3, 21,164, 213, 215, 279, pi. 8. See also Apennine Bench

Formation, Cayley Formation 

plains/secondary-crater relation, 215-219 

radials. See Imbrium sculpture 

rim and rings, 57, 65, 79, 81, 110, 117, 193, 196, 207, 211, 218,

240, 253, pi. 3. See also Montes Alpes, Montes Apenninus, Montes Archimedes, Montes Carpatus, Montes

Caucasus 

samples, 19, 21, 82, 192, 200, 276, 278. See also Apennine

Bench Formation, samples; Fra Mauro Formation,

samples; Montes Apenninus, samples 

secondary craters, 32, 36, 40, 41, 65, 74, 116, 127-131, 155,

156, 164-168, 172, 175, 187, 189, 196, 205, 206, 213,

232, 233, 243, pis. 3, 8, 12 

secondary craters, sampling, 164, 165, 173,174, 204, 205, 211,

218 

stratigraphic relations, 15,19, 21,121, 127,129,148, 173, 174,

179, 196, 199, 204-209, 213-219, 233, 243, 276-279,

pis. 3, 8, 12 

stratigraphic significance, 19, 21,121-125,195,196, 204, 221,

279

structures, external, 81 

structures, tectonic, 112, 113, 117, pi. 5 

Imbrium sculpture, 10, 19, 37, 41, 57, 89, 113, 115, 118, 127, 187,

196, 199, 204, 208, 213 

Impact breccia. See Breccia 

Impact melt

basins, 66-70, 73, 76, 82, 164, 198, 203, 211, 274-276 

cooling time, 45

ejected, 41, 47, 48, 113, 117, 174, 177, 212 

fissured (fractured), 50-52, 113, 117, 203, 212 

flows, 44, 47-50, 53 

homogenization, 50, 177, 178 

injected, 44, 46 

origin, 40, 44, 50, 212 

pools, 28, 44, 46-48, 50, 53, 212, 265 

projectile effects, 44, 47 

stratigraphic relations, 44, 45, 50 

target effects, 47, 212 

terrestrial, 45

volcanic interpretation, 50, 76, 140 

Impact-melt-rock samples, 22, 23, 45, 82, 140, 157, 210 

aphanitic, 201 

Apollo 14, 211, 212, 278 

Apollo 15, 198, 201, 203, 278, 279 

Apollo 16, 165, 168-170, 219, 220, 278 

Apollo 17, 174, 177, 278 

bomb, 177, 178

composition, 101, 140-142, 165, 174, 177, 201, 211, 212 

dating, 22, 168-170, 177, 178, 212, 278 

fragment-laden, 141, 165-170, 174, 177, 278 

mare-basalt samples, comparison, 101 

poikilitic, 46, 168, 174, 177

texture, 23, 46, 140, 165, 174, 177, 178, 201, 203, 210 

uncertain sources, 22, 23, 168-170, 174, 177, 211. 219, 220 

Impact rates

Copernican and Eratosthenian, 246, 271-273, 280 

Imbrian, 160, 191, 245, 246, 279

Impact rates Continued

Nectarian and pre-Nectarian, 151,157,160,180,190,191, 246, 

278-280

present, 280

terrestrial, 271

Impacts, 3, 17-19, 27, 33, 38-42. See also Projectiles and basin 

and crater headings

clustered, 180

crustal disruption, 142, 143, 157, 190, 279

energy, 17, 29, 43

experimental, 31, 33, 38, 39, 280

layering effects, 40, 43, 44, 77, 78, 81

missile, 29, 39

oblique, 19, 32, 38, 39, 81, 82, 171, 173, 258, 266, 267

petrogenetic effects, 143, 157, 277

primary, defined, 17

recent, 280

scaling laws, 43, 81, 211, 212

secondary, defined, 18. See also Secondary-impact craters

simultaneous, 8, 19, 30-33, 39, 81, 171, 197

spacecraft, 266

velocities, 17, 18, 29

Infrared, 96, 97. See also "Eclipse temperatures" 

Ingenii basin, 6, 60, 64, 146, 148, pis. 3, 6 

Inghirami and Inghirami A, 71, 180, 221, 224, pis. 7, 8 

Insularum basin, 65, 146, 148, 153, 190, 207, 244, pis. 3, 6 

Intrusions, igneous, 101, 102, 113, 140-143, 157, 276, 277 

Iridum crater, 8, 32, 34, 125, 197, 221-224, 231, 232, 258, 259, pi. 9 

Iron, 95-101, 140, 141, 276 

Isaev, pis. 4, 9 

Isostasy, 77, 112-115, 143, 145, 245, 277. See also Mascons

J, K

J. Herschel, pis. 5, 6

Janssen, 37, 149, 162, 164, pis. 5, 6

Janssen Formation, 73, 121, 143, 163-167, pi. 7

Jenner, 5, 7, 223, 226, 231, pi. 9

Jetting, 41-46

Joliot, 76, 166, 191, pis. 4, 6, 9

Jules Verne, 6, pis. 4-6, 9

Julius Caesar, 9, 95, 218, pi. 6

Kant Plateau, 164, 165, 218

Karpinskiy, pis. 5, 8

Keeler, 6, 25, 120, 150, pi. 8

Keeler-Heaviside basin, 6, 25, 61, 64, 146, 148, 150, 179, pis. 3, 6

Kepler, 121, 125, 253, 255, 265, pi. 11

Kibal'chich, 180, 183, pi. 7

King, 4, 46, 50, 53, 154, 265, pi. 11

Kipukas, 86, 110

Kirchhoff, 173, 175

Kohlschiitter, 85, 245, pis. 4, 7

Kopff, 33, 40, pis. 5, 9

Korolev basin, 59, 64, 79-82, 148, 178, 179, 181, 190, 245, pis. 3,

7, 8

Korolev M, 182 

Kovalevskaya, 158, pi. 9 

Krafft, 197, 231, 234, pi. 9 

KREEP

basalt, 140, 156, 190, 198, 202, 211, 224, 276-279

composition, 140

crustal component, 140, 143, 144, 156, 211

defined, 140

high-K (HKFM), 140

low-K (LKFM), 140, 141, 168, 171, 174, 201

medium-K (MKFM), 140, 198

orbital chemistry, 143, 156, 190, 198, 245

origin, 142, 198

samples, 165, 168, 171, 174, 190, 198-203, 211, 269, 270

volcanic, 144, 156, 198, 211 

Krieger, 121, 125, 231-233 

Krusenstern, 149, 189

La Caille, 128

La Condamine, 34, 221, 224, pis. 5, 8 

Lacus, defined, 3 

Lacus Autumni, 245, pis. 4, 9 

Lacus Felicitatis, 90 

Lacus Mortis, 96, pis. 4, 9 

Lacus Odii, 93

Lacus Solitudinis, 86, 154, 245, pis. 4, 9 

Lacus Somniorum, 40, 96, 99, 244, pis. 4, 9 

Lacus Veris, 245, pis. 4, 9 

Lade, 205, 206, 213, pi. 6 

Lalande A, 49

Lambert, 8, 110, 121, 125, 247, 250-253, pi. 10 

Lambert R, 110, pi. 8 

Lament, 115, 227, 235 

Landau, 58, 158, pi. 6

Landslide, Apollo 17 landing site, 105, 174, 176, 240, 269, 270, 273 

Langemak, 5, pis. 4, 9 

Langrenus, 167, 252-256, pi. 10 

Lansberg, 221, 223, 272, pi. 9 

Lansberg B and C, 223, 231, 265 

Lassell, 40, pi. 5

Late Imbrian Epoch. See also Upper Imbrian Series 

defined, 121, 196

Late Imbrian Epoch Continued

duration, 245, 246, 279

impact rate, 246, 279

Latitude, equivalence in kilometers, 4, 98, pi. 1 

Lava. See also Magma, Mare-basalt samples, Mare units

channels and tubes, 89, 92

flow lobes, 83, 86 

Layered intrusions, 141, 143 

Layering, effect on impacts, 40, 43, 44, 77, 78, 81 

Lebedinskiy, 180, 184, pi. 8 

Lee-Lincoln scarp, 105-107, 176 

Leeuwenhoek, pis. 4, 7, 9 

Leibnitz, 103, 160, pis. 4, 6, 9 

Le Monnier, 109, 173, pi. 7 

Letronne, 110, 221, 251, pi. 8 

Le Verrier, 8, 32 

Lexell, 253, pi. 10 

Licetus, 216, pi. 6

Lichtenberg, 196, 265, 269-271, pi. 11 

Light plains. See Plains 

Limb, defined, 2, 4 

Lindenau, 232, pi. 9 

Lineated terrain, pre-Imbrian, 127, 205 

Linne, 28, 265 

Liouville DA, 117

Lithosphere, 12, 77, 102, 115, 145, 243-245, 259, 277-279 

Littrow, 171-175

Lomonosov, 4, 76, 156, 191, 221, 222, pis. 4, 8, 9 

Lomonosov-Fleming basin, 65, 146, 148,156, 179, 190, 191, pis. 3, 6 

Lonar Crater, India, 43 

Longitude, lunar conventions, 2 

Longomontanus, 74, pis. 7, 8

Lorentz basin, 59, 64, 72, 118, 147, 148, 158, 179, 196, pis. 3, 6 

Lower Imbrian Series, 193, 279. pi. 8. See also Early Imbrian 

Epoch, Schrodinger, and Imbrium- and Orientalebasin headings

chronology, 224, 279

crater frequency, 130, 135, 160, 186, 197, 221, 222, 230, 246

crater materials, 130-136, 221, 246, 279, pi. 8

definition and type area, 71, 121-123, 196, 208

domes, 222-225

extent, 195, 279, pi. 8

mare materials, 223, 224, 232, 238, 279

plains materials, 156, 197, 198, 204, 208, 215, 222, pi. 8 

Lubiniezky, 256, pi. 7 

Luna missions (U.S.S.R.), 12 

Luna 3, 3, 12, 94 

Luna 16

crater frequency and DL values, 230

landing site, 12, 99, 101, 242, pis. 2, 4, 9

samples, 101, 140, 237, 241, 245, 276-279

spectral reflectance, 99 

Luna 20

landing site, 12, 101, 167, 214, 242, pis. 2, 3, 7

samples, 82, 101, 140, 143, 144, 171, 276-278 

Luna 24

landing site, 12, 99, 101, pis. 2, 4, 9

samples, 101, 237, 241, 245, 276-279

spectral reflectance, 99 

Lunar Astronautical Charts (LAC), 124 

Lunar grid, 107 

Lunar Module (LM), 235, 259 

Lunar orbiter, future, 280 

Lunar Orbiter missions (U.S.A.), 5, 12, 57, 196 

Lunar Orbiter photographs, types, 5, pi. 2 

Lunar Receiving Laboratory (LRL), 23 

Liitke, 47 

Lyot, 7, 103, 226, pis. 4, 6, 9

M

Macrobius, 172, 175, pi. 8 

Maestlin G, 255 

Maginus, 74, 149, pi. 6 

Magma

conduits and extrusion sites, 102, 103, 115, 240, 259

density, 102, 103, 241

differentiation, 102, 142, 143, 156, 157

fluidity, 86

fractionation, 102, 238, 243

high-Al, 102, 103, 241

high-Ti and low-Ti, 102, 103

mixing, 102

primary, mare, 102

primary, terra, 141 

Magma ocean, 142, 143, 244, 277 

Magma pods (ponds), 143

Magnesium, 98-101, 140-143. See also Mg suite 

Magnetism, 168, 215, 256, 269, 280 

Main sequence, craters, 27, 32 

Mairan 34, pi. 9 

Mairan domes, 222 

Maksutov, pis. 4, 9 

Manicouagan Crater, Quebec, 45 

Manilius, 8, 9, 205, 252, 253, pi. 10 

Manned Spacecraft Center (MSC), Johnson Space Center, Houston,

Tex, 23 

Mantle

composition, 102, 277, 280

density, 12
INDEX 299

Mantle Continued

mare-source zones, 101, 238, 240, 259, 261, 276, 280 

origin, 102

samples, 73, 80, US, 278 

structure, 12, 13, 102, 103 

uplifts, 77, 102, 115, 120, 145, 240, 274-276 

Maraldi, 111, 172-175, pi. 7 

Maraldi B, 87 

Mare(pL, maria), 3, 19, 21, 83,117, 229, 232-245, 249, 258-262,

269, pis. 4, 9-12. See also Mare-basalt samples, Mare

units, and lacus, mare, and sinus headings 

area, 3, 85

basin relations, 3, 85, 100-103, 115, 117, 145 

defined, 3, 85, 86

elevations, 77, 78, 100, 115, 117, 120, 241 

morphology and landforms, 85-93 

origin and emplacement, 19, 21, 85, 86, 102, 276 

thickness and volume, 12, 77, 99, 117, 190, 277 

Mare Australe, 7, 62, 103, 167, 223, 226, 233, 245, pis. 4, 9 

Mare-basalt samples, 19, 23, 101. See also Mare Crisium, Mare

Fecunditatis, Mare Imbrium, Mare Insularum, Mare

Nectaris, Mare Serenitatis, Mare Tranquillitatis, Palus

Putredinis

ages, absolute, 235-245, 259, 261, 273, 278 

Al-rich (high-Al), 101-103, 190, 200, 210, 238, 241, 242, 245,

279

Apollo 11, 19, 85, 96, 101, 224, 235, 245, 276-279 

Apollo 12, 19, 23, 96, 101, 102, 249, 259, 276-279 

Apollo 15, 19-23, 96, 101, 102, 237, 243, 276-279 

Apollo 16, 96, 101, 237, 238, 245, 278 

Apollo 17, 96, 101, 102, 239-241, 245, 276-279 

classification, 101, 235, 239-242, 262 

clasts in breccia, 101, 156, 190, 200, 210, 212, 278 

compositions, 96-103 

defined, 85, 101 

density, 12

Eratosthenian. See Apollo 12 (above) 

europium anomaly, 142 

feldspathic, 101, 237, 241, 242, 245, 261, 278 

high-K, high-Ti, 101, 237, 238, 245, 278 

ilmenite, 235, 238, 261, 262 

KREEP-rich, 156, 190, 198, 245 

low-K, high-Ti, 101, 235-241, 245, 278 

low-K, low-Ti, 101, 235-238, 278 

Lower (Early) Imbrian, 224, 279 

Luna 16, 96, 101, 237, 241, 245, 276-279 

Luna 24, 96, 101, 241, 245, 276-279 

mantle sources, 102, 103, 238, 240, 259, 261, 276, 280 

Mg-rich, 242 

mineralogy, 101

olivine (olivine-normative), 237, 239, 243, 245, 259-262, 278 

pigeonite, 237, 243-245, 259-262 

pre-Imbrian (Nectarian and pre-Nectarian), 190, 200, 212, 238,

278, 280

quartz-normative, 243 

reduction, 101

terra-melt-rock samples, comparison, 101 

textures, 23, 101, 102, 235-244, 259, 261 

Ti-poor (low-Ti), 96, 101-103, 235-238, 245, 278 

Ti-rich (high-Ti), 96, 101-103, 235-241, 245, 259, 262, 278 

Upper (Late) Imbrian. See Apollos 11, 15-17 and Lunas 16,

24 (above)

very low titanium (VLT), 101, 102, 237, 242, 278 

Mare Cognitum, 86, 112, 120, 244, pis. 4, 9 

Mare Crisium, 4, 63, 78, 100, 103, 165-167, 172, 241, pis. 4, 5, 9 

age, 230, 241, 245, 279 

color (spectral reflectance) and composition, 96, 99, 241, 242,

pi. 4

crater frequency and DL values, 230, 231 

cross section, 100 

dark-mantling material, 89 

diameter, 103, 117 

elevation, 77, 78, 100, 117, 120 

Mare Imbrium, comparison, 243 

mascon, 99, 117, 241

peripheral "lakes" and eruption sites, 78, 103, 241, 259 

samples, 241, 245, 276-279 

stratigraphy, 100, 241, 243, pi. 9 

structures, 117, 241, pi. 5 

thickness, 99, 100, 103, 117, 241, 242 

Mare Fecunditatis, 119, 167, 241, pis. 4, 9 

age, 241, 245, 279

color (spectral reflectance), 96, 99, 241, pi. 4 

composition, 241 

crater frequency, 230 

diameter, 117, 241 

elevation, 78, 117, 120 

mascon, absence of, 99, 117 

samples, 241, 245, 276-279 

stratigraphy, 233, 241, pi. 9 

structures, 115-119, 241, pi. 5 

thickness, 99, 100, 103, 117, 242 

Mare Frigoris, 8, 96, 99, 125, 199, 204, 224, 243, 258, 279, pis.

4, 9, 10

basin-related trough, 81, 198, 243 

Mare Humboldtianum, 9, 61, 103, 117, 166, pis. 4, 9 

Mare Humorum, 61, 112, 117, 187, 244, pis. 4, 5, 9, 10 

color (spectral reflectance), 96, 99, 244, pi. 4 

dark-mantling material, 89, 116 

mascon, 99, 117, 244 

structures, 107, 112, 116, 117, 244, pi. 5

Mare Humorum Continued

thickness, 99, 100, 117

Mare Imbrium, 8, 63, 83, 92, 99,125-127,193, 224, 231, 243, 247, 

253, 258, 259, 272, pis. 4, 5, 9, 10. See also Palus 

Putredinis

age, 243-245

color and composition, 96, 99, 243, 244, 258, 259, 279, pi. 4

composition, 99, 244, 258, 259, 279

diameter and elevation, 117, 120

DL values, 231, 243

flow lobes, 83, 86, 258

mascon, 99, 117, 243

stratigraphy, 94, 121, 125-127, 243, 247-253, 258, 259, 265, 

269, pis. 9, 10

structures, 113, 117, 243, pi. 5

thickness, 100, 117, 243, 244 

Mare Ingenii, 6, 60, 216, pis. 4, 9 

Mare Insularum, 153, 259, 272, pis. 4, 9, 10

crater frequency and DL values, 230, 259, 271-273

domes, 86

samples, 259, 276

Mare Marginis, 4, 67, 76, 166, 191, 224, 245, pis. 4, 9, 10 

Mare Moscoviense, 59, 245, 257, pis. 4, 9 

Mare Nectaris, 1, 62, 164, 238, pis. 4, 5, 9

age, 245

color (spectral reflectance), 96, 99, 238, pi. 4

dark-mantling material, 89

diameter and elevation, 115, 117, 238

DL value, 230

mascon, 99, 115, 117, 238

samples, 238, 245, 276, 278

stratigraphy, 233, 238, pi. 9

structures, 112-117, 238, pi. 5

thickness, 100, 103, 117 

Mare Nubium, 10, 96, 103,112, 116, 117, 120, 153, 232, 244, 256,

258, pis. 4, 5, 9, 10

Mare Orientale, 11, 62, 69, 77, 99, 103, 115, 117, 233, 245, pis. 4, 

5, 9

name, 2 

Mare ridges (dorsa), 3, 8, 103, 107, 153, 277, pi. 5

age, 115

basin-ring markers, 110, 196 

Mare Serenitatis, 9, 61, 92, 108, 171, 174, 238, pis. 4, 5, 9, 10

age, 230, 241, 245

color (spectral reflectance), 96, 99, 239, 244, 245, pi. 4

crater frequency and DL values, 230, 231

dark-mantling material, 89, 92-95, 239-241, 245

diameter, 117

elevation, 93, 117, 120

geologic maps, 94, 95, 174, 176, pis. 9, 10, 12

gravity structure (mascons), 79, 99, 117, 171, 239

landforms, 86, 90, 239

samples, 237, 239, 245, 276-279

sources, 259

stratigraphy, 19, 94, 95, 99, 238, 244, 245, pis. 9, 10

structures, 107, 117, 239, pi. 5

thickness, 99, 100, 117, 173, 239, 241 

Mare Smythii, 4, 7, 62, 116, 154, 245, pis. 4, 5, 9, 10

composition, 99, 103, 245

dark-mantling material, 89

diameter, 117

elevation, 115, 117, 120

Mare Orientale, comparison, 245

mascon, 99, 115, 117, 245

stratigraphy, 245, 259, pis. 9, 10

structures, 115-117, 245, pi. 5

thickness, 99, 103, 117, 245 

Mare Spumans, 78, 100, 241 

Mare Tranquillitatis, 9, 39, 114, 119, 172, 227, 235, pis. 4, 5, 9

color (spectral reflectance), 96, 99, 235, pi. 4

crater frequency and DL values, 230, 235, 238

diameter, 117, 235

elevation, 93, 115, 117, 120

landforms, 86, 87

mascon, absence of, 99, 117

samples, 224, 235, 245, 276

stratigraphy, 233, 235, 238, pi. 9

structures, 39, 114-119, pi. 5

thickness, 99, 100, 117, 235, 238, 241 

Mare Undarum, 78, 100, 241

Mare units. See also individual locus, mare, sinus, series, and 

system, headings

ages, 96, 102, 124, 125, 190, 223, 224, 229-262, 265, 269,

278-280, pis. 9-12 

albedo, 94-96

aluminous, 98, 99, 241, 242, 245 

basin relations, 223, 232-235, 238-245,258, 259, 269,276-280,

pi. 4

buried, 156, 190, 243, 245, 278, 279

chemistry, orbital, 97, 139, 143, 156, 190, 198, 241, 245, 258 

color (spectral reflectance), 86, 96, 235, 238-245, 258, 259, 262,

279-281, pi. 4 

crater frequencies, 130, 136, 160, 229-235, 238, 252, 259,

271-273

DL values, 130, 136, 230-233, 243, 259, 271-273 

"eastern" and "western," 259 

eruption rates, 86, 238-241 

extent, 86, 127, 232, 233, 277 

farside, 3, 85, 103, 233, 245, 280, pis. 4, 9-12 

flows, 83, 86, 127 

geologic mapping, 86, 94, 95, 125, 174, 232, 236, pis. 4, 9-12

Mare units Continued

KREEP-rich, 156, 190, 245

mapping properties, 86-99

Mg-rich, 100, 241

oldest, 76, 197, 223

radioactivity, 96-99, 245, 259

remote-sensing properties, 94-99

sources, 102, 103, 240, 259

stratigraphic relations. See Superpositions and individual 

maria

subdivision, 85, 86, 94-96, 232-235

thickness, 86, 99, 153, 238-241, 261

Ti-poor (low-Ti) and Ti-rich (high-Ti), 96, 99, 101-103, 235, 

238-245, 258-262, pi. 4

youngest, 197

Mare Vaporum, 86, 92, 99, 127, 214, 244, pis. 9, 10 

Marginis basin, 65, 148, 179, pis. 3, 6 

Marius Hills, 86, 88, 89, 244, 258, pi. 4

Mascons, 77, 99, 100, 112, 115, 117, 171, 235, 238-245, 269, 277 

Mass wasting. See Downslope movement, Slumping, Talus,

Terraces

Massifs. See Basin materials

"Material of Montes Apenninus," 15, 73, 200-203, pi. 8 

Maunder, 30, 33, 40, 69. 76, 253, pi. 10 

Maunder Formation, 69, 76, 82, 164, 174, 198, 216, pi. 8 

Maupertuis, 34, 221, 224, pis. 5, 8 

Maurolycus, 155, 214, pi. 7 

Maxwell, pis. 4, 7 

McCauley, J.F., vii, 66-77 

McClure, 40

Mechnikov, 180, 183, pi. 7 

Megaregolith, 45 

Megaterracing, 78, 79, 81 

Member, defined, 123 

Mendel-Rydberg basin, 11, 61, 64, 148, 178, 179, 186, 278, 293,

pis. 3, 4, 7 

Mendeleev basin, 6, 59, 64, 85, 117, 120, 148, 179, 186, 188, 219,

230, 293, 25 pis. 3, 7 

Mercator, 112, pi. 7 

Mersenius, 180, 187, pis. 5, 7 

Messala, 166, pi. 6

Messier and Messier A, 32, 38, 171, 173, 265 

Metamorphism

Apollonian, 212

shock, 45-48, 211

textures, 46, 141

thermal, 46, 157, 165 

Meteor Crater, Ariz., vii

Meteorites, 101, 141, 142, 277. See also Projectiles 

Metius, 147, 180, pi. 7 

Mg suite, 141-144, 156, 157, 174, 201 

Middle Crescent Crater, 259, 261 

Milichius domes, 86, 87 

Milne, 5, 58, 65, 145-149, 154, pi. 6 

Minerals. See also Ilmenite, Olivine, Plagioclase, Pyroxene, Spinel

crustal, 140

mare-basalt, 101

minor, 101, 140

silica, 140 

Mixing models, 142

Modal mineral (mode), defined, 102, 140 

Modification stage, craters, 43 

Moho (crust-mantle boundary), 12, 77-81, 102, 115, 120,143, 229,

240, 245, 262, 274-278 

Mohorovicic A, 152 

"Moldings," debris, 110, 111 

Mons (pi., montes), defined, 3 

Mons Rumker, 86, 244 

Montes Alpes, 8, 9, 81, 125, 198, 199

Montes Apenninus, 8, 9, 13, 15, 22, 36, 81, 90, 108, 193, 197, pi. 

8. See also "Material of Montes Apenninus"

crater frequency, 136, 197, 222, 230

samples, 19, 23, 49, 82,140, 143, 144, 195, 198, 224, 278, 279

slumps, 81, 193, 198, 200 

Montes Archimedes, 113, 193, 198 

Montes Carpatus, 197, 198, 204, 207, 211, 272 

Montes Caucasus, 8, 9, 81, 198, 199, 202 

Montes Cordillera, 66-73, 76-81, 179, 197 

Montes Haemus, 9, 171, 174, 205, 218 

Montes Harbinger, 89, 91, 233, 251, 258 

Montes Recti, 224 

Montes Riphaeus, 90, 112, pi. 6 

Montes Rook, 66, 68-70, 73, 76, 77, 80, 171 

Montes Rook Formation, 69, 164, 168, 171-174, 203, pi. 8 

Montes Taurus, 171-176 

Monzodiorite, 140, 223, 269 

Moon

age, 156, 157, 277

astronomic properties, vii

comparative planetology, vii, 280

cooling and contraction (global), 115, 269, 280

geologic style, viii, 276, 277

global properties, 12, 13

origin, vii, 156, 157, 280

subsurface, 12-14, 274, 275

surface, 1-13

surface area, 180 

Moonquakes, 13, 115, 269, 280 

Moscoviense basin, 59, 64, 81, 82,148,171,178, 179,186, 245, 257,

293, pis. 3, 7 

Mosting, 128, 252
300 THE GEOLOGIC HISTORY OF THE MOON

"Mug shot," defined, 23

Miiller chain, 41, 115

Multi-ringed basin, defined, 3. See also basin headings and basins

listed on p. 64, 65, 148, 179 

Multispectral images, 99 

Mutus, 155, pi. 6 

Mutus-Vlacq basin, 65, 145-148, 155, 164, 179, 190, 244, pis. 3, 6

N

Nansen, 199, pi. 7 

Nearch, 147, 149, pi. 6 

Nearside

defined, 2, 3

mass offset, 103 

Nectarian Period

defined, 121, 163

duration, 157, 190, 278

geologic history, 190, 191, 277-279

igneous activity, 190, 277-279

impact rate, 151, 157, 160, 180, 190, 191, 246, 277-280 

Nectarian System, 161, 278, 279, pi. 7

basin materials, 161, 276-278, pis. 3, 7

chronology, 190, 277-279

crater frequency, 130, 136, 160, 178-180, 186, 191, 197, 230

crater materials, 130, 136, 149, 180, 246, 278, pi. 7

defined, 121-125, 163

distribution, 163, pi. 7

plains materials, 154-156, 163-167, 184, 188-191, 217, 219, 

230, pi. 7

structures, 278

type area, 146, 147

volcanic rocks, 190, 243, 278 

Nectaris basin, 1, 3, 62, 64, 146, 161, 179, 238, pis. 3, 5, 7, 12

age, 168, 178, 179, 186, 190, 278

composition, 143, 144

crater frequency, 136, 160, 179, 197, 293

deposits, 64, 161, 189, 219, 220, 232, pis. 3, 7, 12

impact angle, 82

impact melt, 164

knobby terrain, 168

mascon, 117, 238

Orientale basin, comparison, 163, 164, 168

radials, 7, 131, 143, 189

rim and rings, 22, 36, 64, 81, 155, 163, 235, pi. 3

samples, 82, 163, 216, 220, 276

secondary craters, 36, 37, 64, 74, 131, 143, 161

stratigraphic relations, 121-125, 131, 143,161, 179,189, 192, 

235, 243, 276-280, pis. 3, 7, 12

stratigraphic significance, 3, 121-125, 143, 145, 163, 280

structures, tectonic, 112, 115, 117, 238, pi. 5 

Neper, 166, 180, pis. 4, 7, 9 

Nernst, pis. 5, 7 

Nicolai, 37, 164 

Nishina, pis. 4, 9 

Nomenclature

petrologic, 12, 101, 102, 139-142, 156

selenographic, 3

stratigraphic, 121-125 

Nonproportional growth, of craters, 43, 44 

Norite, 140-142, 157, 174, 201, 203 

Normal albedo, defined, 94 

Normative mineral (norm), defined, 102 

North Complex, 91 

North Massif, 174, 175

North Ray Crater, 165, 168-171, 269, 270, 273 

Nubium basin, 65, 103, 146, 148, 153, 179, 244, pis. 3, 6

0

Oceanus Procellarum, 11, 103, 115, 145, 153, 244, pis. 4, 5, 9-12

color (spectral reflectance), 96, 99, 232, 244, 258, 259, 269, pi. 4

crater frequency, 250-252, 266

dark-mantling material, 89, 234

elevation, 120

landforms, 86, 89, 110

localization factors, 103, 115, 145

mascons, 99, 115, 244

stratigraphy, 121-127, 135, 196-198, 232-234, 244, 250, 265, 

269-271, pis. 9-12

structures, 107, 110, 115, 118, 196, pi. 5

thickness, 99, 100, 153 

Olbers and Olbers A, 197, 265, pis. 5, 7, 11 

Olivine, 101, 102, 140-142, 259

Olivine (olivine-normative) basalt, 101, 237, 239, 242-244, 261, 278 

Ophitic texture, 23, 46, 174, 235 

Oppenheimer, 160, pis. 5, 7 

Orbital geochemistry. See Chemistry, orbital 

Orientale basin, 3, 11, 55, 196, 279, pis. 3, 8, 12. See also Hevelius 

Formation, Maunder Formation, Montes Rook 

Formation

age, 178, 186, 196, 224, 245, 277, 279

antipode, 67, 76

central uplift, 69, 81

concentric facies, 66, 73

corrugated facies, 73, 76

crater frequency, 135, 136, 197, 216, 219-222, 230, 293

DL value, 216, 230

domical facies, 73

Orientale basin Continued

dunelike deposits, 67, 71, 130, 218

exterior, 66, 81, 82

fissured facies, 66, 73

formation, 78-82

geologic maps, 76, pis. 3, 8, 12

global effects, 67, 72, 190, 216, 220, 223

gravity structure, 77, 80, 117

grooved facies, 67, 73

impact angle, 82

impact melt, 66-72, 76, 77, 82, 212

interior, 68-73, 76-82

mare fill, 190, 245, 279, pis. 4, 9. See also Mare Orientale

massifs, 66, 69, 73, 79, 80

model for basins, 65, 66, 73, 77-82, 279

pitted terrain, 67, 71

plains, 3, 71, 128, 164, 179, 182, 184, 190, 197, 204, 213, 224, 

279, pi. 8

radial facies, 73

"rays," 66, 74, 75, 127

rings, 64-73, 77, pi. 3. See also Montes Cordillera, Montes Rook

secondary craters, 32, 35, 66, 86, 127, 128, 151, 158, 159, 

181-183, 188, 196, 205, 224, 225, 233, pis. 3, 8

stratigraphic relations, 127, 128, 148,158, 159, 179-183, 188, 

196, 224, 225, 233, 277, pis. 3, 8, 12

stratigraphic significance, 121-123, 195, 196, 221, 229, 279

structures, tectonic, 117, 243, 245, pi. 5

topography, 77, 280 

Orientale Group, defined, 66, 72 

Origin, meanings, 19, 21, 139, 163 

Orontius, 10, 20, 149, pi. 6 

Orthopyroxene, 140-142 

Oscillatory uplift, 78-81 

"Overturned flap," ejecta, 42, 73

Palus, defined, 3

Palus Putredinis, 8, 22, 113, 200, 243, pis. 4, 9

age, 243, 244

color and composition, 243

crater frequency and DL values, 136, 197, 230, 243

elevation, 120

samples, 243, 276-279 

Parry, 22, 208, pi. 7 

Partial melting, 101, 102, 140 

Paschen, 183, pi. 6 

Pasteur, 4, 120, 149, 154, 191, pi. 6 

Pauli, 5, pis. 4, 8, 9 

Pavlov, 5, pi. 7

Peak rings, 28, 50, 58, 65, 77-81 

Peaks. See Basin materials, massifs; Crater materials, peak; Crater

processes, peak formation 

Peirce, 100, 172 

Peirescius, 147, 149, pi. 6 

Period, defined, 123

Petavius, 32, 58, 102, 113, 167, 232, pis. 4, 5, 8 

Petavius B, 252, 265, 267, pi. 11 

Petrogenesis, early impacts, 157, 277 

Petrologic nomenclature, 139, 140, 156 

Philolaus, 252, pi. 11 

Phocylides, pi. 7, 8

Photogeology, vii, 17, 72, 139, 143, 280 

Photographs

identification scheme, 4, 5, pis. 1, 2

lighting effects, 2, 86, 94, 110, 133-135, 173, 219-221, 250

quality, 145, 265, 280, pi. 2

resolution effects, 12, 13 

Picard, 100, 172, 241 

Piccolomini, 231-233, pi. 9 

Pictet, 20

Pigeonite, 101, 102, 140, 243, 259 

Pigeonite basalt, 101, 102, 237, 243, 259, 278 

Pingre-Hausen basin, 65, 148, 179, pis. 3, 6 

Pitatus, 10, 116, 180, 215, pis. 5, 7 

Pitiscus, 164, 180, pi. 7 

Plagioclase. See also Anorthosite

composition, 140, 142

mare basalt, 101

terra crust, 12, 140, 142

textures, 23, 46, 141, 171

Plains (light plains, terra plains), 3, 4, 10, 11, 21, 71, 113, 128, 131, 

154-156, 164, 172, 184, 188, 197, 200, 204, 213, 215, 

276, 279, pis. 7, 8, 12

crater frequencies, 191, 197, 216, 219, 222, 230

distribution, 3, 15, 21, 169, 190, 215, 217, 276, pis. 7, 8

DL values, 136, 216

Imbrian, 15, 156, 191, 198, 208, 213-224, 279, pi. 8. See also 

Apennine Bench Formation; Cayley Formation; 

Imbrium basin, plains; Lower Imbrian Series, plains; 

Orientale basin, plains

Nectarian, 154-156, 163, 164, 167, 184, 188-191, 217, 219, 

230, pi. 7

pitted, 146, 164

pre-Nectarian, 146, 154, 156

secondary-crater relations, 67, 71, 72, 128, 146, 164, 190, 

213-220, 224

volcanic, buried, 139, 146, 190, 191, 243, 245, 279 

Planck basin, 5-7, 59, 64, 103, 147, 148, pis. 3, 4, 6, 9 

Planetology, defined, vii

Plante, 8, pi. 10

Plato, 8, 125, 221, 224, 231, 259, pis. 4, 9, 10

Playfair and Playfair G, 10, 189, pis. 6, 7

Plinius, 93, 253, pi. 10

Plutonic rocks, 140-143, 156, 157, 278

Pohn-Offield crater classification, 129-131, 143

Poikilitic texture, 23, 46, 165, 168, 174, 177

Poincare basin, 6, 7, 59, 64, 103, 146, 148, pis. 3, 6

Pontecoulant, 167, 180, pi. 7

Posidonius, 9, 22, 93, 113, 231, pis. 5, 9

Potassium, 98, 99, 102, 140, 190. See also KREEP; Mare-basalt

samples, high-K, low-K 

Pratt model of isostasy, 143 

Pre-Imbrian (materials and time), 143-145, 196 

Pre-Nectarian period

defined, 121-125, 139, 145

duration, 157, 278

geologic history, 156, 157, 277

igneous activity, 139-143, 156, 157, 190, 277, 278

impact rate, 151, 157, 160, 180, 191, 246, 277-280 

Pre-Nectarian system, 137, 277, 278, pi. 6

basin materials, 139, 143, 178, 180, 277, 278, pis. 3, 6

chronology, 156, 277, 278

crater frequency, superposed, 130, 136, 145-149, 186, 197

crater materials, 130, 136, 139, 145, 180, 186-189, 278, pi. 6

defined, 121-123, 125, 139, 143, 145

distribution, 139, pi. 6

plutonic rocks, 141-143, 156, 157, 278

recognition criteria, 145-147

samples, 139-143, 156, 157, 163, 195

structures, 139, 278

type area, 146, 147

typical terrain, 137

volcanic rocks, 156, 190, 238, 278, 279 

Primary-impact crater, defined, 17 

Prinz, 91, 125, pi. 8 

Pristine rocks, 141-144, 156, 157 

Procellarian System, 124, 125 

Procellarum basin, 65, 103, 115, 120, 145, 153, pis. 3-6

age, 157, 277, 278

crust and lithosphere, 77,103, 115,120,143, 235, 238-245, 259, 

262, 278, 279

depth, 80, 120, 143-145, 192

hemispheric asymmetry, 103, 143, 145

maria, 99, 103, 115, 145, 232-245, 259, 262, 278, 279

rings and troughs, 65, 81, 103, 112,115,120, 145,153, 171-173, 

207, 227, 240, 243, 244, 278, pis. 3, 5

structures, tectonic, 107, 112, 115, 145, 238, pi. 5

volcanic filling, early, 156, 190, 211 

Procellarum Group, 125 

Proclus, 51, 172, 175, 265 

Projectiles

basin association, 143

density, 43, 47

ejection, 41

mass-frequency distribution, 18, 129

melting, 40

multiple, 19, 30-33

penetration, 40-44, 81

present-day, 280

secondary, 29, 211

size effects, 43, 77, 81, 177

velocities, 17, 18, 29, 177 

Promontorium, defined, 3 

Proportional growth, craters, 43 

Protagoras, 204

Ptolemaeus, 10, 32, 128, 155, 213, 219, pis. 6, 8 

Purbach, 10, 128, 149, pi. 6

Pyroclastic material. See Dark-mantling material, Glass 

Pyroxene

augite, 140, 143

clinopyroxene, 101, 102, 140

high-Ca, 97, 140

low-Ca, 97, 102, 140

orthopyroxene, 140-142

pigeonite, 101, 102, 140, 243, 259

textures, 23, 46, 141, 171, 177 

Pythagoras, 8, 196, 258, pi. 10 

Pytheas, 30, 121, 125

Q,

Quartz, 140

Quartz monzodiorite, 140, 223, 269

Quartz-normative basalt, 102

Rabbi Levi, 232, pi. 6

Radar, 99, 100

Radioactivity, 98, 99, 143, 190, 198, 212, 259, 280

Ranger missions (U.S.A.), 12, 90, 125, pi. 2

Rare-earth elements, 101, 140, 142. See also Europium anomaly,

KREEP 

Raspletin, 31 

Rays, 249-257, 265-267

albedo, 94-96, 249, 250, 256, 265-267, 279

asymmetric, 32, 38, 39, 266

basins, 74, 75, 82, 127

fading, 95, 96, 249, 256, 279

length, 3, 29

offcenter relation to primary crater, 29, 30

origin, 29-32, 42, 48
INDEX 301

Rays Continued

primary-ejecta content, 48, 269, 270

samples, 269, 270

stratigraphic relations, 125-127, 249-252, 265-267 

Reduction, chemical, 101 

Regiomontanus, 10, 128, pi. 6 

Regolith (soil, surficial material), 12, 13, 95, 96

age, 249

agglutinates, 95, 96, 256, 266, 279

albedo, 95, 96

Apollo 11 landing site, 13, 235, 238

Apollo 12 landing site, 259, 261, 269

Apollo 14 landing site, 205

Apollo 15 landing site, 13, 198, 243

Apollo 16 landing site, 165, 168, 219, 220

Apollo 17 landing site, 240, 241

bedrock substrate, 12, 13, 21, 45

blockiness, 99, 258

breccia, 45

buried, 238

craters, 78, 132, 133, 258

drainage, 269, 271

formational rate, 280

glasses, 95, 96, 101, 280

Luna 16 landing site, 101, 241

Luna 20 landing site, 171

Luna 24 landing site, 101, 241, 242

maturity, 95, 96

origin, 17, 277

sampling, 21, 249

Surveyor 1 landing site, 258

thickness, 12, 165, 205, 235, 239, 240, 258 

Reimarus R, 147, 149 

Reiner, 258, pi. 10 

Reiner gamma, 256, 258, 269 

Reinhold, 253, 272, pi. 10

Remote sensing, 94-99. See also Albedo; Chemistry, orbital; Color; 

Radar; Spectral reflectance

defined, 86

integration of types, 99 

Repsold and Repsold C, pis. 5, 6 

Residual liquids, 140 

Rheita, 147, 161, 164, 180, pi. 7 

Rhysling Crater, 243 

Riccioli, 67, 71, 197, pis. 4, 6 

Riccius, 37, 149, 214, 232, pi. 6 

Richardson, 156, pis. 4, 6, 9 

Ries Crater, Germany, vii, 43, 45, 73, 78, 165 

Rilles (rimae), 3, 86-93, 107, 167, 168, 239. See also individual 

rimae

ages, 115, 127

arcuate, 3, 107, 145

distribution, 3, 107, 112, pi. 5

sinuous, 3, 18, 77, 86, 200, 233, 247

straight, 3, 107, 113, 127, 145, 167, 168 

Rima (pi., rimae), defined, 3 

Rima Bode II, 90, 115, pi. 5 

Rima Hadley, 13, 88, 91, 200, 201, 237, 243 

Rima Hyginus, 9, 50, 88, 90, 115, 118, 120, 127, 205, pi. 5 

Rima Sharp I, 117

Rimae Ariadaeus, 115, 118, 220, pi. 5 

Rimae Cauchy, 115, 119, pi. 5 

Rimae Hypatia, 39, pi. 5 

Rimae Stadius, 30, 31 

Rimae Sulpicius Callus, 93, 108, pi. 5 

Rimae Triesnecker, 115, 120, pi. 5 

Ringed basins. See basin headings and basins listed on p. 64, 65,

148, 179

Ritter, 9, 32, 39, 113, pi. 8 

Rocca, 71, 180, pi. 7 

Roche, 5, pis. 4, 7

Rock-stratigraphic units, 121-129. See also Basin materials, Crater 

materials, Mare units, Orientale Group, and individual 

formations

defined and named, 123-125

interfingering, 27, 121, 124

lateral continuity, 17, 21, 27

sample correlations, 21-23, 101, 139

three-dimensionality, 17

time-stratigraphic relations, 121-125 

Romer, 172, pi. 11 

Romer A, 173, 175 

Rontgen, 118, pi. 6

Rosenberger, Rosenberger B, Rosenberger C, 147, 149, pi. 6 

Rothmann, 232, 233, pi. 10 

Rumford, pis. 4, 9 

Rumker hills, 86, 89, 244, pi. 4 

Rupes, defined, 3 

Rupes Altai, 131, 155, 164, 233 

Rupes Recta, 10, 115, pi. 5 

Rydberg, pi. 10

S

Sabine, 32, 39, 113, pi. 8

Sacrobosco, 131, 149, pi. 6

St. George Crater, 91

Samples (material from Moon), vii, 17-23, 280. See also samples

listed on p. 168, 177, 200, 237, 253, 261 

correlation with stratigraphic units, 21, 101,139, 143, 276, 280

Samples (material from Moon) Continued 

mare basalt, 101, 102, 276, 278 

numbering system, 23, 168, 177, 190, 200 

terra breccia, 21-23, 45-47, 139-143, 276-279 

10003, 236, 237 

10029, 236, 237 

10050, 235-238 

10062, 236-238, 245 

12002, 261

12031. 261

12032. 269

12033. 269

12038, 261

12051, 23

12065, 261

14053, 190

14063, 101, 190

14066, 212

14072. 190

14073. 211, 212 

14161, 212 

14167, 212 

14276, 211, 212 

14306, 269 

14310, 23, 211, 212 

14321, 190, 210, 269 

15016, 244 

15382, 198, 200 

15386, 198, 200, 202 

15405, 269, 271 

15445, 23, 200, 201 

15455, 156, 200-203 

15538, 23 

15555, 244 

60025, 141 

60315, 170 

62195, 168

67015. 171

67016. 169

67435, 141, 156

67603, 169

67667, 156

68415. 220

68416. 220

72255, 156

72275, 190

72415-72418, 156, 157, 174

72435, 178

73215, 177

73255, 156, 177, 178

74001, 241

74220, 241

76215, 177

76535, 141, 143, 156, 157, 174

77215, 156, 174

78236, 156

Sarton, 158, 159, pi. 7

Satellitic craters, 30. See also Secondary-impact craters 

Saussure, 20 

"Scablands," 86

Scaling laws, impacts, 43, 81, 211, 212 

Schickard, 11, 128, pis. 4, 6, 9 

Schiller, 11, 32, 38, pi. 7 

Schiller C, 128

Schiller-Zucchius basin, 11, 32, 38, 59, 64, 147, 148, 190, pis. 3, 6 

Schlesinger, 159, pi. 6 

Schluter, 231, 232 , pis. 4, 5, 9 

Schrodinger basin, 7, 59, 64, 67, 79-82, 148, 155, 179, 186, 222,

225, 279, 293, pis. 3, 5, 8 

Schwarzschild, 58, 65, 180, 199, pi. 7 

Sculptured hills, 174, 176 

Sechenov, 180, 183, pi. 7

Secondary-impact craters, of basins. See Basin-secondary craters;

Imbrium basin, secondary-craters; Nectaris basin,

secondary craters; Orientale basin, secondary craters

Secondary-impact craters, 3, 18, 19, 28, 125, 224, 226, 231, 247,

256-258, 276

circularity, 29, 42, 130

cratering process, 29, 31, 32

endogenic hypotheses, 17, 29, 32

formational sequence, 42, 48

ground-surge relation, 42, 47

impact velocities, 18, 29

laboratory simulation, 31, 33

primary-crater size, 29, 32, 211, 276

primary-ejecta content, 42, 48, 276

projectiles, 29, 42

shock grades, 48

size-frequency distribution, 29, 32, 129, 132, 276

source direction, 125, 255

spatial distribution, 18, 19, 29-35, 125, pis. 3, 6-11

superpositions, 125-127. See also Superpositions

target-structure influence, 29 

Seismic velocities, 12, 13, 205, 210 

Seismicity, 12, 13, 269 

Seleno-, vii

Selenographic coordinates, 2 

Septa, 29, 32, 33, 37, 39, 214, 220 

Serenitatis basin, 3, 8, 61, 64, 93, 108, 109, 171, pis. 3, 5, 7

age, absolute, 177, 190, 191, 278

age, relative, 19, 173, 178-180, 190

Serenitatis basin Continued

antipode, 181, 215

cross section, 82, 192, 240, 243, 276

deposits, 174, pi. 7

Imbrium-basin interaction, 81, 192, 198, 201

mare extrusion, 103, 279

mascons, 117, 171

modification, 173

Orientale basin, comparison, 171, 173

rim and rings, 64, 81, 110, 117, 171, 218, pi. 3. See also North 

Massif, South Massif

samples, 82, 143, 144, 157, 163, 173, 201, 278

secondary craters, 172, 175

stratigraphic relations, 148, 164, 173,175,179, 192, 235, 238, 

243, 276, pis. 7, 12

structures, tectonic, 108, 109, 112, 117, pi. 5 

Series, defined, 123 

Sharonov, pi. 11 

Sharp, 34, pi. 9 

Sherlock Crater, 270 

Shield volcanos, 86 

Shirakatsi, 47, pi. 10

Shock compression and decompression, 40-42 

Shock grade, in ejecta, 41, 42, 45-48, 211, 212, 219 

Shock melting. See Impact melt 

Shock pressures, 40, 211 

Shock waves, 40-42, 45, 80, 81 

Shoemaker, E.M., vii, viii 

Shoemaker-Hackman stratigraphic scheme, vii, 19, 123-125, 143,

249, 265

Shorty Crater, 237-240, 269, 270, 273 

Siderophile elements, 101, 141, 143, 157, 198, 202 

Sikorsky-Rittenhouse basin, 7, 65, 148, 179, 180, 278, pis. 3, 7 

Silberschlag, 118 

Silica

mare basalt, 101, 102

orbital detection, 98

terra rocks, 140 

Sills, 101, 102, 276

Sinuous rilles, 3, 18, 77, 86, 200, 233, 247 

Sinus, defined, 3

Sinus Aestuum, 92, 244, pis. 4, 9 

Sinus Amoris, 172, 235, pis. 4, 9 

Sinus Asperitatis, 120, 235, 238, 244, pis. 4, 9 

Sinus Iridum, 8, 34, pis. 4, 9. See also Iridum crater 

Sinus Medii, 81, 86, 95, 115, 213, 244, pis. 4, 5, 9 

Sinus Roris, 99, 117, 222, 243, 258, pis. 4, 9, 10 

Size-frequency distributions. See Crater frequencies 

Slope material, 49, 96, 110, 111 

Slumping, 48, 51, 81, 193, 198, 200. See also Downslope movement,

Megaterracing, Terraces 

Smoky Mountain, 165, 170, 219 

Smythii basin, 7, 62, 64, 77, 112-117, 146, 148, 154, 179, 245, pis.

3-6

Soda (sodium), mare basalt, 101 

Soderblom-Lebofsky crater-dating technique, 133. See also DL

method

Soil, defined, 12. See also Regolith 

South Massif, 105, 174, 175, 240, 273

South Pole-Aitken basin, 11, 65, 115, 143-148, 152, 181, 245, 278, 

pis. 3, 6

age, 157, 277, 278

crust and lithosphere, 115, 143, 245, 278

deposits, 146, 278

elevation, 120, 145

mare extrusion, 103, 115, 245, 278

rim and rings, 65, 145, 181, pi. 3

stratigraphic relations, 145, 148 

South Ray Crater, 165, 168, 170, 269, 270, 273 

Spacecraft impacts, 266 

Spaceflights, 12 

Spalling, 81 

Spectral classes, 96, pi. 4

hDG-, 96, pi. 4

HDSA, 96, 99, 258, pi. 4

hDSA, 96, 99, 258, pi. 4

hDSP, 96, 99, pi. 4

HDWA, 96, 99, pi. 4

hDWA, 96, 99, 239, 244, pi. 4

LBG-, 96, 99, 244, pi. 4

LBSP, 96, pi. 4

LIG-, 96, 99, pi. 4

LISP, 96, 99, pi. 4

mBG-, 96, 99, 238, pi. 4

mIG-, 96, 99, 244, 259, 262, pi. 4

mISP, 96, 99, 239, 244, pi. 4

sampling, 101, 262

symbols, defined, 96 

Spectral reflectance, 96-99, pi. 4

craters, 266

dark-mantling material, 97-99, 240

maria, 96-99, 232, 235, 238-245, 258, 259, 269. See also 

individual maria

terra soil, 97 

Spinel, 140-143 

Spur Crater, 203 

Stadius chains, 30, 31 

Steady state, cratering, 129-135, 157, 230 

Stearns, 257, pi. 10 

Stebbins, 149, 159, pi. 6 

Steinheil, 147, 180, pi. 7
302 THE GEOLOGIC HISTORY OF THE MOON

Steno Crater, 239 

Steptoes, 86, 110 

Stevinus, 252, pi. 11 

Stiborius, 232, 233, pi. 9 

Stillwater Complex, Mont., 143 

Stofler, 10, 216, pi. 6 

Stone Mountain, 165, 170, 218, 219, 237 

Straight Range, 224 

Straight Wall, 10, 115 

Stratigraphic code, 123-125 

Stratigraphic column, 121-125, 280

Stratigraphic units, 17. See also Rock-Stratigraphic units, Timestratigraphic units, Time units, and individual formations, series, and systems 

Stratigraphy, vii, 15, 121. See also Superpositions

basin-mare distinction, 19

crater interpretation, 17-19

defined, vii

nomenclature, 121-125

regional context, 21

Shoemaker-Hackman scheme, vii, 19, 123-125, 143, 249, 265

terra-material interpretation, 19-21 

Stratovolcanoes, 86 

Stratton, 25, pi. 7 

Strength crater, 43, 81

Structures, tectonic, 105, 277, pi. 5. See also Faults, Fractures, 

Mare ridges, Rilles

Copernican, 115, 269, 271

crater-floor, 113-118, pi. 5

distribution, 107, 112-115, pi. 5

Eratosthenian, 279

far side, 245

Imbrian, 238, 241-245

lithospheric thickness, 115, 277

maria and basins, 107-115, pi. 5

pre-Imbrian, 139, 278 

Struve L, 35

Subophitic texture, 23, 46, 165, 168, 174, 177, 202, 235 

Suevite, 45, 46, 73, 165 

Sulpicius Callus, 93 

Superpositions (superpositional relations), 17, 125

basin-basin, 127, 145, 175, 178, 192, 196, 199, 222, 225, 235, 

240, 243, 276, pis. 3, 6-8, 12

crater-basin, 127, 135-138, 145-151, 154-168, 172-175, 

179-191, 193-199, 202-208, 213-234, 274, pis. 6-8, 12

crater-crater, 18, 25-27, 31, 34, 47, 48, 125, 149,180,185,189, 

191, 196, 224, 226, 248-253, 256, 266-268, 273

mare-basin, 19, 21, 77, 85, 94, 95, 108, 109, 115, 121, 153, 

165-169, 174, 193-200, 209, 223-227, 231-245, 259, 

274-279, pis. 4, 12

mare-crater, 18, 21, 76, 85, 99, 121, 125, 131-135, 191, 197, 

200, 204, 205, 223, 226, 231-234, 238, 239, 247-261, 

265-272, pis. 9-12

mare-mare, 94, 95, 109,110, 226, 231-245, 247, 254-262, 269 

Surficial layer. See Regolith 

Surveyor Crater, 259-261 

Surveyor missions (U.S.A.), 12

chemical analyses, 12, 85, 139

landing sites, 12, 254, 258, 259, pi. 2 

System, defined, 123 

Szilard, 156, pi. 6

Talus, 89, 93 

Tamm, 117

Taruntius, 113, 116, 119, 166, 190, 265, 269, pis. 5, 11 

Taruntius H, 28

Taurus-Littrow Valley, 22, 171, 175, 239, 240, 273, 279 

Tectonism. See Structures 

Tektites, 45 

Terminator, defined, 2

Terra (pi., terrae) (highlands, uplands), 3,12,19-23. See also Crust 

"backbone," 10, 19, 244, 258, 262

Terra (pi., terrae) (highlands, uplands) Continued

defined, 3

regional differences, 3

volcanism, 21, 146, 190, 215, 222, 223, 280 

Terra plains. See Plains 

Terraces

basins, 64, 78-81

craters, 27-29, 43, 44, 48 

Textures

aphanitic, 23, 174, 177, 178

breccia, 23, 45-47, 171, 174

cooling-rate relation, 46, 102, 174, 259

cumulus, 141-143, 174

fragment-content relation, 46, 165, 174, 177, 212

granoblastic and granulitic, 46, 157

igneous, terra rocks, 23, 46, 47, 140, 141, 174, 198

impact melt, 23, 46, 165, 174, 177

intergranular, 165, 168

intersertal, 23, 202, 238

mare basalt, 23, 101, 102, 235-244, 259, 261

metamorphic, 46, 141

ophitic, 23, 46, 174, 235

poikilitic, 23, 46, 165, 168, 174, 177

seriate, 141, 171

shock intensity, 45-47, 165, 211

subophitic, 23, 46, 165, 168, 174, 177, 202, 235

vitrophyric, 243 

Thalassoid, defined, 3 

Thebit, 128, pis. 5, 9

Theophilus, 45, 227, 235-238, 252-256, pi. 10 

Thorium, 98, 99, 102, 140, 190, 198 

Tidal bulge and deformation, 115 

Timaeus, 204, pi. 9

Time-stratigraphic units, 121-125, 130, 133, 136. See also 

individual series and systems

areal extent, 163, 195, 229, 249, 265, 276-280, pis. 3, 4, 6-12

defined, 121-123 

Time units, defined, 121-123. See also individual epochs and

Timiryazev, 180, 183, pi. 7

Timocharis, 8, 121, 125, 231, 250-253, pi. 10

Tisserand, 172

Titanium, 95-103, 243. See also Mare-basalt samples, Mare units

Tranquillitatis basin, 227, 238, pis. 3, 6

mare fill, 103, 235, 238, 279, pi. 9

mascon, 115, 117

rings, 174, 235, pi. 3

samples, 157

Stratigraphic relations, 103, 145-148, 179, pi. 6

structures, tectonic, 112-117, 235, pi. 5 

Transient crater, 43

Trask, N.J., 131-135. See also Crater dating, Trask's method 

Triesnecker, crater and rille system, 115, 120, pi. 5 

Troctolite, 140-143, 156, 157, 174, 201 

True crater, 43 

Tsander, 183, pi. 6

Tsiolkovskiy, 4-6, 47, 48, 120, 188, 231-233, 245, pis. 4, 9 

Tsiolkovskiy-Stark basin, 65, 148, 245, pis. 3, 6 

Tsunami mechanism, 78 

Tycho, 1, 10, 20, 28, 265, 270, pi. 11

age, 269, 270, 273, 278, 280

albedo, 266

crater frequency and DL values, 252, 253, 266, 273

impact melt and knobs, 28, 73, 76

rays and secondary craters, 1, 20, 29, 32, 72, 125, 176, 240, 

270, 273, 280

Stratigraphic relations, 20, 121, 125, 270, 273

U

Ukert, 193

Uniformitarianism, 19, 20, 265

Unnamed basin A. See Sikorsky-Rittenhouse basin

Unnamed basin B. See Coulomb-Sarton basin

Unnamed craters A and B, 219, 222

Upper Imbrian Series, 227, 279, pi. 9. See also Late Imbrian Epoch

chronology, 245, 279

crater frequency, 130, 135, 136, 160, 221, 222, 229

crater materials, 130-136, 221, 222, 2S1, 246, 254-256, 279, 

pi. 9

dark-mantling materials, 109,174, 233-241, 244,245,269,278, 

279

definition and type area, 121-123, 229-231

DL values, 125, 130, 136, 229-233, 243, 245, 249, 250, 273

extent, 229, 279

mare materials, 121, 125, 135, 160, 227, 252, 254, 256, 278, 

279, pis. 9,12. See also locus, mare, and sinus headings

structures, 245 

Uranium, 98, 99, 140, 190

V

Vallis (pi., valles), defined, 3

Vallis Alpes, 113, 115, 116, 119, 204, pi. 5

Vallis Bohr, 35

Vallis Bouvard, 70, 73, 163, 187, 213

Vallis Palitzsch, 167, pi. 7

Vallis Rheita, 147, 161, 163, 164, 167

Vallis Schroteri, 91

Van de Graaff, 6, 39, 120, 180, 216, pis. 4, 7, 9

Van den Bos, 117

Vega, 147, 149, pi. 6

Vendelinus, pis. 4, 6

"Very high alumina (VHA) basalt," 141, 165, 168

Very low titanium (VLT) basalt, 101, 237, 242, 278

Viscous relaxation, 113, 115

Vitello, 113, 116, 187, pis. 5, 8

Vitello Formation, 214

Vitruvius, 86, 175, 221, pi. 9

Vitruvius front, 171, 175

Vlacq, 147, 155, pi. 6

Volatiles, 101, 102, 243, 245, 277

Volcanic craters. See Endogenic craters

Volcanic plains, buried, 139,156, 164, 190, 215, 232, 233, 238, 243,

245, 266, 276-281 

Volcanic rocks. See also KREEP basalt, Mare-basalt samples

Nectarian, 190, 243, 278

pre-Nectarian, 156, 190, 243, 278 

Volcanism, 19. See also Magma and mare headings

eruptive styles, 86-89

rates, 190, 238, 240, 241, 245, 278-280

terra, 21, 146, 190, 215, 219, 222, 223, 280 

Volcanotectonic craters, 32 

Volta, pis. 5, 6 

Von Karman and Von Karman M, 103, 160, pis. 4, 6, 9

W, X, Y, Z

W. Bond, 9, 199, 204, pi. 6

Wallace, 231

Walter, 10, pi. 7

Wargentin, 11, 72, pis. 7, 8

Water, 101

Waterdrop experiments, 80, 81

Werner, 10, 48, 127, 128, 155, 189, 253, pi. 10

Werner-Airy basin, 65, 145, 148, 155, 179, 244, pis. 3, 6

West Crater, 235-238, 270

Wrinkle ridges. See Mare ridges

X-ray-fluorescence spectrometer, 97-100

Young D, 147, 180, pi. 7

Zagut, 232, pi. 6

Zeno, 166, 171, pi. 7

Zero phase, 94

Zhukovskiy, 180, 184, pi. 7

Zond missions (U.S.S.R.), 12, 145

Zucchius, 38, pi. 11
PLATES
-(A

PLATE 1. INDEX TO PHOTOGRAPHIC ILLUSTRATIONS

30°, 30°

20" 20°

 */**-&* x .-I-:,{,

Pythagoras^ I 7

-^*" ,, , \ / "

 Pabbage ab'Ty'u ' ' i J>

^\ \Herechel A f

S^ \ f~ * .-<. - / V

NEARSIDE

70"

SOUTH POLE

WEST EAST 

LONGITUDE LONGITUDE

1 2 3

lw=.-x-*.<o. .

4

Chapter

: 7 ,i 8 9 10 11 12 13

Area of photograph and figure number

Chapter number precedes decimal point, figure number within chapter follows. 

Photographs of small areas, individual craters, and basin centers (fig. 4.3) are not 

shown. Boundaries dashed where approximate.

Landing site
NORTH POLE

80°

20°

8 10°

Birkhoff M 

^Rowland \\

Carnot \V Sarton

- y \ \v (

' , \ / \ Oppenheirifier

/ \ £eibn>fe \

' >*> \

70°

80° FARSIDE

SOUTH POLE

EAST 

LONGITUDE

WEST 

LONGITUDE

Each 10° of latitude equals about 300 km

Shaded-relief base by J. L. Inge, U.S. Geological Survey; reproduced 

by courtesy of the National Geographic Society.
2A PLATE 2. PHOTOGRAPHIC COVERAGE OF THE MOON, SHOWING

70°^

NORTH ^=^~ one ^...-,, _ - -i - ~ "' 

^^ ,

wv &

x 7 /;

POLE

V

X,

-*H 

<V 92

^

«D 

, lOS * -C

- * %

7 ^ X

h-APOLLO Surveyor 3 

16 Apollo 12

NEARSIDE

Areas multiply photographed by Lunar Orbiter; resolution, 10-20 m

1 M 44, 50-67; 3 M 25, 31, 33; 5 M 38, 41. 42, 44-51

1 M 48, 49, 68-83; 2 M 5, 7, 9, 11-20, 25-32, 35-42; 3 M 5-19 (odd

nos.); 5 M 52, 55-63 

1 M 85-100; 2 M 67-74, 76-91; 3 M 58, 60, 63, 66, 68, 70; 5 M

64,71-78

1 M 105-112; 2 M 94; 3 M 80-83 

1 M 118-133; 2 M 93, 95-111, 113-136; 3 M 84, 86-101; 5 M 108-

115

1 M 141-148, 3 M 116-119

1 M 137, 139, 140; 2 M 163-178; 3 M 120, 124-131 

1 M 157-172; 3 M 136-160

1 M 176-183; 3 M 171, 173-180; 5 M 169-176

2 M 195, 197-212; 3 M 161, 163-170 

1 M 185-215; 3 M 172, 181-212

 Actual boundary of photographs

SOUTH POLE

WEST 

LONGITUDE

EAST 

LONGITUDE

An H-frame is centered within each M-frame: 

each H-frame has an eightfold better resolution 

than the corresponding M-frame. X-ray 

spectrometer data acquired along Apollo 15 track 

between long 155° E. and 40° W. and along 

Apollo 16 track between long 140° E. and 30° W.; 

gamma-ray spectrometer data acquired along 

entire Apollo 15 and 16 tracks including 

unphotographed regions (see chap. 5). Orbiter 

resolutions after Hansen (1970); complete plots of 

Orbiter coverage given by Hansen (1970), 

Kosofsky and EI-Baz (1970), and Bowker and 

Hughes (1971). Apollo data after Masursky and 

others (1978).

Other Orbiter 1, 2, 3, 5 lowaltitude frames; resolution, 

10-40 m. Numbers of 

Orbiter 5 M-frames are 

inclusive

Orbiter 4 low-altitude Hframes, resolution, 60-100 

m. Denoted by midpoint; 

adjacent rows overlap in 

zone centered on heavy 

dashed line

Other Orbiter; resolution, 150-

Approximate boundary of best coverage Other Orbiter; resolution, 

>300m
BEST PHOTOGRAPH OR SET OF PHOTOGRAPHS OF EACH AREA 2B

NORTH POLE

/.7/^ // ^V^^s

/ X /~5 M 5-12 \ V

70°

21,23/^7 / I

X500

K40°

/ 1 

/ 5 M 181^ /(

/ / 5 M 158 / 5 M 124 / 5 M 103 /5 M 851 5 M 79

' S

7 \7*\ 

/ )/>

FARSIDE SOUTH POLE

EAST WEST 

LONGITUDE I LONGITUDE

APOLLO 15 APOLLO 16 APOLLO 17 Overlap Apollo 16 frames taken after No coverage

... , .A . , ^ , . . . trans-Earth injection; Apollo orbital photographs (panoramic and map- ... ^ cr. o^

. ,r , , \ V « o on resoluton, 150-300 m ing; near-vertical only); resolution 2-30 m

Landing site

Shaded-relief base by J. L. Inge, U.S. Geological Survey; reproduced 

by courtesy of the National Geographic Society.
3A PLATE 3. GEOLOGIC MAP OF RINGED BASINS

NORTH POLE

10e

70° 70°

50°

40° 40"

30° 30°

20°

10°

30° 30°

40°

\A "'^ N^O

V^S^N ^\\HPH\ Vi >y 60°

> ,-> f^"t

.- ;/ *,i

NEARSIDE

70°

SOUTH POLE

WEST EAST 

LONGITUDE LONGITUDE

Q

Undivided Inner Outer

Undivided Inner Outer

Undivided

Undivided

Imbrian

Nectarian

Pre-Nectarian

Pre-Nectarian

Basin ring 

oor ring arc

Dotted where inferred. Names of definite basins with mappable deposits 

are in capital letters, and their deposits are shown in color; mapped 

units include all but minor deposits of each named basin and include 

inferred buried extensions. Names of definite basins without mappable 

deposits are in lowercase letters. Indefinite basins are indentified by 

initials: AG, Amundsen-Ganswindt; AK, Al-Khwarizmi /King; BK, 

Balmer-Kapetyn; FB, Flamsteed-Billy; Fe, Fecunditatis; GW, GrissomWhite; In, Insularum;LF,Lomonosov-Fleming; Ma, Marginis; MV, MutusVlacq; Nu, Nubium; PI, P2, P3, Procellarum; PH, Pingre-Hausen; 

SR, Sikorsky-Rittenhouse; Tr, Tranquillitatis; TS, Tsiolkovskiy-Stark; 

WA, Werner-Airy
3B

NORTH POLE

70°

60" 60°

50°

40°

,10°

/C"/J ^ >.v'"'.'/ X^T",

f> 

wh^\^^\^ I I I.V ', .' , ^' '._ -)( ''- * «v ( < \

* Af / /., Ar<r " TS ~f^ ' :

. I'. ' '.*. (» , ^ !, r

JUC X AitUer/

>^^ft; :mA 

r*\. Bn°

ti^

40°

70°

80°AG0

SOUTH POLE

EAST WEST 

LONGITUDE LONGITUDE

FARSIDE

Contact 

Queried where indefinite or buried

Approximate contact between inner and 

outer deposits. Queried where doubtful

Landing site

Shaded-relief base by J. L. Inge, U.S. Geological Survey; reproduced 

by courtesy of the National Geographic Society.
4A PLATE 4. MARIA

NORTH POLE

NEARSIDE

SOUTH POLE

WEST I EAST 

LONGITUDE I LONGITUDE

SPECTRAL CLASS"

HDWA HDSA hDSA hDSP hDG- hDWA mISP mIG- mBG- LISP LIG- LEG- LBSP

Dark Cone Undivided

mantling and dome (complex

material material or unknown)

Terra

Projection of mare-terra contact 

beneath younger crater deposits

Basin rim. Dashed where uncertain Inner ring of Procellarum basin
4B

NORTH POLE

80°

10°

60°

20°

Richardson 

Joliot SJ V 

/ A Maxwell

n// //Lomonosov

OArtamor ov

X?Langemak f \ 

i t *

. Leibnitz

O

Nishina

rman (~)

SCHRODINGER)

FARSIDE

SOUTH POLE

EAST 

LONGITUDE

WEST 

LONGITUDE

^NOTE:

Spectral classes (see table 5.1) from Pieters (1978; see 

Basaltic Volcanism Study Project, 1981, pi. 2.8), 

redrawn here on the basis of color-difference image 

(fig. 5.20). Mare Humorum units from Pieters and 

others (1975); Mare Crisium units from Head and 

others (1978a). Lowercase names and names in parentheses refer to basins or craters containing otherwise 

unnamed maria

Landing site

Shaded-relief base by J. L. Inge, U.S. Geological Survey; reproduced 

by courtesy of the National Geographic Society.
5/4 PLATE 5. STRUCTURAL FEATURES

NORTH POLE

70° 70°

60°

50° 50°

40° 40°

30°

20° 20°

40°

60°

NEARSIDE

70°

WEST EAST 

LONGITUDE LONGITUDE

Mare ridge

Fault (mostly graben)

Fractured crater floor

Basin rim Dashed where uncertain

Inner ring of Procellarum basin
NORTH POLE

70° /

M

y^-""""

50° _ -/Compton

Karpinskiy

40°

'V f . ', <? V J - - ' '-' .< J >.-v. '>' 40°

20°

10°

20°

. ^ ^°U

 '&w

?-.^4 ^ -'.

< r ^ / -. t ' ? 1 « r^< \ .- 

4" (, ^-^w (

r 'V<:

I*

SOUTH POLE

J, '^^ 10° ^.'

MARE 

ORIE1)^ALE

30°

FARSIDE

EAST 

LONGITUDE

WEST 

LONGITUDE

Shaded-relief base by J. L Inge, U.S. Geological Survey; reproduced 

by courtesy of the National Geographic Society
6A PLATE 6. PRE-NECTARIAN SYSTEM

NORTH POLE

PROCELLARUM

Julius CaesarQ TRANQUILLITATIS 

(~) Murchison

Ramsteed P \~_

/'Rothmann G 

Zagut

70°

NEARSIDE

SOUTH POLE

WEST 

LONGITUDE

EAST 

LONGITUDE

Nectarian and younger units

Interior Exterior 

Basin materials, younger groups (6-9)

Basin interior materials, older groups (2-5)

6? 0

Rim crest of buried crater 

Dashed where possibly Nectarian

O

Basin rim

Crater interior 

materials

Intercrater terrane

Probably consists mostly

of crater deposits and

undetected basin deposits

Interior materials of South Pole-Aitken basin
6B

NORTH POLE

OoO\O u\0 n ,r>aoHalnc R \

Zwicky \^^M -77 / ^

- ^-Wa9!< ill*

FARSIDE

EAST I WEST 

LONGITUDE I LONGITUDE

Plates 6-11 are paleogeologic maps that show units of the indicated age in color, older units only by 

the shaded-relief base, and younger units by blank areas. Exterior basin and crater deposits that are 

sufficiently well expressed to display stratigraphic relations with other units are mapped; divided 

where possible into near-rim deposits consisting of thick primary ejecta and outer discontinuous 

deposits including secondary craters and their ejecta, small patches of plains, terra-mantling material, 

and lineate ejecta. Craters >30 km in diameter are mapped; exterior deposits mapped around craters 

>120 km (pi. 7) or >60 km (pis. 8-11) in rim-crest diameter. Basin-interior materials (pis. 6-8) 

include massifs,knobby ejecta, and impact-melt plains. Number of superposed craters varies with age 

of underlying basin; number of visible buried craters varies with thickness of overlying deposit and 

with age of the substrate.

Shaded-relief base by J. L. Inge, U.S. Geological Survey; reproduced 

by courtesy of the National Geographic Society.
7A PLATE 7. NECTARIAN SYSTEM

NORTH POLE

SERENITATIS C1 Le Monnier

TaylorQ 

O Apollo 16

Albategnius /" X

K ) Descartes 

Alphonsus

. 

Geber OCathaT

WaurolycLis. < , Metius ^^ Rheita

10°

NEARSIDE

SOUTH POLE

WEST EAST 

LONGITUDE I LONGITUDE

Plains

May overlie Nectarian 

or pre-Nectarian 

volcanic deposits

Post-Nectarian 

units

i

Interior Exterior 

Basin materials, 

younger group

Interior Exterior

Basin materials,

older group

Pre-Nectarian 

units

Interior Exterior 

Crater materials
IB

NORTH POLE

ts DnuyiiMii ^^^ , ,". ,'. ^

^.-. .'^}^?°a-^ °"*V*'

^L- o ;--/( \) ^Sv_.0^«^y.'r] ;;^ 

x^^*7 . x- ^ ^. o ' ^»*,A'

SlKORSKY-RITTENHOUSE

vC^

Amundsen 

SOUTH POLE

EAST WEST 

LONGITUDE LONGITUDE

FARSIDE

o o

Rim crest of buried crater

Dashed where possibly

pre-Nectarian or

Lower Imbrian

O

Basin rim

Shaded-relief base by J. L. Inge. U S. Geological Survey; reproduced 

by courtesy of the National Geographic Society.
SA PLATE 8. LOWER IMBRIAN SERIES

NORTH POLE

70°

20»

NEARSIDE 80°

SOUTH POLE 

WEST LONGITUDE EAST LONGITUDE

Post-Lower 

Imbrian units

Plains

May overlie Nectarian 

or Lower Imbrian 

volcanic deposits

Montes Rook and Hevelius Outer deposits

Maunder Formation and secondary

Formations (inner) craters

Interior Exterior

Alpes Formation Fra Mauro Outer deposits

and material of Formation and secondary

Montes Apenninus craters

Pre-Imbrian 

units

I- Orientale basin materials

Materials of craters 

and Schrodinger basin

Imbrium basin materials
8B

NORTH POLE

80°

30°

20°

10° 10°

Lebedinskiy 

LebedinskiyK Engelhardt (HERTZSPRUNG)

SOUTH POLE

EAST WEST 

LONGITUDE I LONGITUDE

FARSIDE

o o

Rim crest of buried crater 

Dashed where possibly Nectarian or Upper Imbrian

O.

Basin rim

Buried contact 

location inferred

Note:

Pre-Imbrian basins (capitals) and craters 

(lowercase) containing otherwise unnamed plains patches are named in 

parentheses

Shaded-relief base by J. L. Inge, U.S. Geological Survey; reproduced 

by courtesy of the National Geographic Society.
9,4 PLATE 9. UPPER IMBRIAN SERIES

NORTH POLE

^60°

C^%§&

T-maeus ' £ ^ J*"^^^.^

* Vr Endyrnion A

MARE FRIGORIS X-^-p ^^^j^^^^^,-^ (jN(Endyrnion) 

atoVw^^W^ ^.HercuIeXrV

i n>i 1 

r H

Hansen

*^

Luna 24+/

(Ncper)l

^ J'L-B'jndtinU u^ 

)Sharp»/"^>

7a^. \/^

r I . » *"

/ rt. ; 5-t 1'7*". . >< v:,^;/'-',.

MARE

NEARSIDE

lLa Per

Schorr^

r20°

SOUTH POLE

WEST EAST 

LONGITUDE I LONGITUDE

Post-Imbrian 

Units

Mare basalt _

Volcanic Interior Exterior 

province Crater materials

Pre-Upper 

Imbrian units
NORTH POLE

/ i'"""" ' ""^^Bintpr, -'- / ., yr/ \ -^ ' > "i ^^f i *$? Zr -iffwrfdfon> '^Pf^; > > " y\l

* , / V i" " t f , .'.' - . *? 

'Aitken

dmann ., Kearpns

 ^ ftiJn

LACUSf 

VERIS /

- 

(Lceuwenhqek) * .-'

», 1 Golytsin

.

-

P\(JuleTVerne) Oresmc V

' ' 

MaksUtov 

O /30°

(Chretien)

VonKfirmanJ ^ ^^^ ^White $,- < f/' \t , '/' <r f* ^ ' ^^tektoV't' 

w , n ,umsnen| /^__ /? ^O /> 'f/'./^'^' ' .-"(^ ; ^/*'' -' ,4" f- V\^%^ ^.T^ v ^Sco, ^ (voK^M) <^ .v ^fcl^Vf /- ;; Txyaht ^ -''/ 

 ;vv?^>«*?'^^.'- a^-^OHess, n-' ( l&j-;rf£K^-^&{£.$ % ;%Kt^v^ x i r^L >^^ t;*^-^ ^v v% . ^ v ^C^" VX-;^ 1^

^Sjlv-3?^^^

O i 

'(Abbe)/ kC *$T,;TI&* (£^^'*r;*9. x,. ^

'f,nn^rf /^A^J-yrl^J ^ ,/.'- WpeRoy^X ^> '^/ '^<4'A

>w - ',* /<->.. * / --'Jr /¥ * * ; * " ''f, * L^^'- /ff *'& ' '^>V- ' '/ff- / 

- M^ M^L/-.S~' "^*?/" < - ' ^ '/

'?:,

,niadi y-rc2/*' -""^^ "' ^? 

>''/ *"?-£-'*.. ''* f "

'-«- .*- -^Si > -^->' '<& ' 

t

SOUTH POLE

EAST I WEST 

LONGITUDE I LONGITUDE

*X C?7 DoerfelS

. , '^' ^-- » est 'x>*' ^

-fc-'Doerfel

FARSIDE

NOTE:

Mare mapping preliminary (pis. 9-12); contacts 

generalized from plate 4; small patches omitted; ages 

on central farside from Walker and El-Baz (1982). 

Volcanic provinces include numerous domes, cones, 

and dark-mantling deposits; other dark-mantling 

deposits are not shown (see pi. 4). Pre-Late Imbrian 

basins (capitals) and craters (lowercase) containing 

otherwise unnamed mare patches are named in 

parentheses

Shaded-relief base by J. L. Inge, U.S. Geological Survey; reproduced 

by courtesy of the National Geographic Society.
10,4 PLATE 10. ERATOSTHENIAN SYSTEM

70°

40°

20° 20°

/ - $" - 

Cepheus ' Geminus*

^Timochans , 4^K MARE T

J i< , , J|^^ SERENITATIS

, v " . *^F

A ' . ,»<'->.

 i / \ ^

ib v^f*, ^\ t' fc

. Belmptte \ ^ /

, ' » -^. \X

"' >. ' i.

F

^^-. - ^'-~j**.y-^ I / ~i N #

^5^ L- ^ / ^ ^Timocharis 

EulerxLambeiit * ..;y

f X

Apoiio 12^?^^

--Vi -4

Langrenus .';<(

 

t, ( * 

% : * *^« 

^

*

:E ^ ' , *' *fA-- ¥ - ^M ^ , , ; v /-,C{- ^

. if Wemer (f ' . S

l, ' X ^

10°

30°

v \N

"V'

^ i 'V »Cichus:, .»;

' "T,,V * f<

. J : '

» /i' 1^^11261 A|(

^. V ^ ' A"

 f'*t

 Reichenaach^/y

' 20°

\ 

40° <;

 * 

50°

ichenaachj/y / '.J'',A% , ,''/.

? :| - A^' ^^I/^^.:^!^^'l^' ^^. { / / *^-*: '"^^V-'K»V/i^r *

-^ '; r^v^^^>^^f4^^K V ^ < v's-,-,/-. $ £?& '$ '' */'--*!,W*<'

*?,',.V * -\ 'y-'-'<< -*-L"^' f f' ' r (' '''-4 ' '!'-* *'C s'- -4- '-^

\

-^ ^_ < ; . c.t , 'jr< < --"- -I.*:, '" '" K * **:/. &f? 

c f A _, -, v. s -'t 'S f-, '*" t- -" .- ' -Jt^ '

NEARSIDE

Post-Eratosthenian 

units

0 

Mare 

basalt

Interior Exterior 

Crater materials

Pre-Eratosthenian 

units
80°

20°

?%mwqmwis-z-, W4<- -}> ; %-H- W< '^,*mf^ :fe>':* if;. -J ft: ^y t(, -.V\. ~ .. #\i!Ss: -^^?P"« :3|¥^' t? :^' ' '"' ^;; 0*

?^p- f ^ : vW* i f VC"^ '^H-/^-^-^-*-^! ' ; ;.1%^' ! '' '^f ' » orse' ) v / -v--/^-V Foster 'p> A Vc 'm® ..'v " \

', /./ -;-,rC''V ^ . J'^.f4' ** * v ~r-^-- f-- --. v^_./ ^-^y ,c_ v, ^ '-t'v-V ^ '»

' (% / Xf-V'^"'^ 1 '^- ^ ^.;^: Van'GentX». -* , . " ". , l ' \f ^ , ' V /. ' i. ^- \ jf> ,*' M C ', ( '''\ *

j a* jf|| *C $ *-Vt " f*' "' ' 1 ) * /* ^^ ^ * ^* ** V j ' £'* '^'^^ * *'f f* '' ' **t", J 4 '$'~ * f~' 'it* ^ *r l j^V* ' ^ *«^* l?^u|- -/M ( J; 4 -^c/A' ; ^H^^I/^^F <v-> !' *,/ h^/qrv /*i v;;; ,/: vt

feMYTHIJ%^J^ (^r , 5 , % X" '# ^''^ V^'' <|] Y" ^iV^'^ ' '«^'V' ^' H '' ->^\ li i' <il '^'

^^flfll^^ ^^^f ^- ^*'^ ^4; ;«.« v-v«^ f > <^ -: ,^^ '^$^^^^

^^^^

- : , ^>-^*^^iV5* -*& ^-^> ^

iff -* ^ ^ ^^^^:^|^-.w^c^r^ %^r; 70°

FARSIDE

SOUTH POLE

EAST 

LONGITUDE

WEST 

LONGITUDE

Note:

Pre-Eratosthenian basins (capitals) and 

craters (lowercase) containing otherwise 

unnamed mare patches are named in 

parentheses

Shaded-relief base by J. L. Inge, U.S. Geological Survey; reproduced 

by courtesy of the National Geographic Society.
PLATE 11. COPERNICAN SYSTEM

NORTH POLE

70°

Anaxagoras

Philolaus

50°

70°

>v Haym

": ^'"'~>-

 Thales

60°

50°

40° 'Harpalus

!^ "* x : ^ ' ( ^

" ",Vs Sx^ 40°

30°

»Burg

^ .v

30°

Eudoxus

^Aristillus 

>Autolycus

f. ^Aristarchus Romer

Kepler*

MK '

UN V

^ ^Taruntius , */; ,-.? '' '*-.'»

 Godin

10° <v t, Gassendi A 

t

,10°

20°

NEARSIDE

SOUTH POLE

WEST 

LONGITUDE

EAST 

LONGITUDE

Mare 

basalt

Interior Exterior 

Crater materials

*- ,. " ( "Wnl' 1

-r-;.U'L g-ter

Pre-Copernican 

units
70°

30°

 :>'*., *.<$( *, - : '-'^X/ v. w*, ft Kfate:::'£&;; <?'

70°

FARSIDE

SOUTH POLE

EAST WEST 

LONGITUDE I LONGITUDE

Shaded-relief base by J. L. Inge, U.S. Geological Survey; reproduced 

by courtesy of the National Geographic Society.
1Z4 PLATE 12. GEOLOGIC MAP OF THE PRESENT MOON

NORTH POLE

70°

40°

30"

20°

10° 10°

60°

70°

NEARSIDE

ERATOSTHENIAN

Mare basalt Crater material

Mare basalt Crater material J SYSTEM

SOUTH POLE 

WEST LONGITUDE I EAST LONGITUDE

COPERNICAN 

SYSTEM

Mare Volcanic Crater material 

basalt province

Plains 

material

Crater material 

(including Schrodinger basin)

Imbrium -basin 

material

Q)

o _J

Plains 

material Older basin 

material

Nc h NECTARIAN SYSTEM

Crater 

material

Basin material 

(including Milne)

- IMBRIAN SYSTEM Undivided 

material

- PRE-NECTARIAN SYSTEM
NORTH POLE

70°

60°-

40° 40°

30°

20°

10° 10°

10°

; 20°

50°

70°

FARSIDE

SOUTH POLE

EAST 

LONGITUDE

WEST 

LONGITUDE

Shaded-relief base by J. L. Inge, U.S. Geological Survey; reproduced 

by courtesy of the National Geographic Society.