Tectonophysics, 161 (1989) 147-156
Elsevier Science Publishers B.V., Amsterdam - Printed in The Netherlands
147
Growth of planetary crusts
STUART‘ROSS TAYLOR
Research School of Earth Sciences, Australian National University, Canberra (Australia)
(Received October 9,1987; accepted December 2,1987)
Taylor, S.R., 1989. Growth of planetary crusts. In: L.D. Ashwal (Editor), Growth of the Continental Crust.
Tectonophysics, 161: 147-156.
The planets and satellites of the Solar System show much diversity, but most have formed crusts which differ
substantially from their bulk compositions. Three principal types of crust may be broadly distinguished: primary,
formed after accretional heating (e.g., lunar highlands); secondary, formed following partial melting in planetary
mantles (e.g., lunar maria, terrestrial oceanic crust); and tertiary, formed by processing of secondary crusts, of which
the continental crust of the Barth is the only identifiable example. Crustal growth before 3800 m.y. was complicated by
the heavy planetesimal bombardment: active erosion on the Earth is likely to have destroyed the brecziated rubble
resulting from over 200 Mare Orientale-scale events in this period. The Mercurian crust is probably primary, perhaps
analogous to the lunar hi&lands. In contrast, the observable Martian and Venusian crusts are secondary, dominated by
basaltic volcanism. The icy crusts of the satellites of the outer planets are mixtures of both primary and secondary
crusts, with complex histories. Ganymede is of special interest in displaying a secondary (grooved terrain) water ice
crust which has split an older more heavily cratered, possibly primary crust. Slight expansion has resulted probably
from polymorphic transitions from high density ice VIII to low density ice I. The primary lunar crust, 12% of lunar
volume, grew in 10’ years; secondary crusts derived by partial melting of mantles grow at a much slower rate-total
production of terrestrial oceanic crust over 4000 m.y. is only 2% of planetary volume, whereas terrestrial continental
crustal growth is even slower, producing over the same period a volume only about 0.33% that of the Earth. Crustal
growth on the other planets appears to be an irreversible process, without evidence of recycling. No evidence for
massive planetary expansion, as required by expanding earth hypotheses, is apparent on the other planets or satellites
of the Solar System.
The diverse nature of the Solar System
There is a basic philosophical problem in dealing with planetary crusts; all the 8 planets * and the 60 or so satellites differ from one another. Theories which attempt to provide general princi- ples for planetary evolution tend to founder on the rock of stochastic events. Accordingly there
* Pluto (r = 1123&3.5 km) is much smaller than the moon,
and the Pluto-Charon pair is really a double satellite system
rather than a planet and satellite.
~1951/89/$03.50 0 1989 Eisevier Science Publishers B.V.
are difficulties in trying to discover some general patterns of crustal growth in a system in which random events are common. Thus, although there is clear evidence that the moon was melted and formed a primary crust dominated by plagioclase feldspar, this does not necessarily provide us with a model for crustal development in the early Earth, Venus, or Mars, all of which differ, not only from the lunar example, but from one another.
It is assumed here that accretion of the planets was accomplished “brick by brick” following the general lines of the planetesimal hypothesis (Wetherill, 1986). However, the terrestrial planets
148
are not identical in composition (e.g., in density, 0 isotopes, K/U ratios, and rare gases), and it is becoming increasingly clear that they accreted from rather narrow feeding zones which differed in these parameters. Thus, although the inner planets are “chondritic” in composition in a gen- eral sense, it is not possible to match them up with specific meteorite classes and the present popula- tion of meteorites does not appear to be a good analogue for the building blocks of the terrestrial planets (Taylor, 1988). The purpose of this paper is to review briefly our present understanding of the crusts of the terrestrial planets and of some of the planetary satellites, with an emphasis on at- tempting to understand crustal growth. For this reason, many significant petrographical or geo- chemical details peripheral to this theme are treated lightly. There is not space to give an extended discussion on mechanisms of crust-for- ming processes; these require separate treatment. Fuller discussions may be found in the references quoted (see, for example, Basaltic Volcanism on the Terrestrial Planets-L.P.I., 1981).
The solid planets and satellites mostly have relatively thin crusts which differ markedly in composition from their interior and from primor- dial solar nebula compositions. Familiarity with our own crust perhaps has dimmed their remarka- ble nature, such as the tendency to concentrate sizable fractions of the planetary budgets of in- compatible elements. This is seen most notably in the high surficial abundances of the heat-produc- ing elements K, U, and Th in planets and satellites so far sampled.
Although the near surface concentration of in- compatible elements was long ago recognized for the Earth, the lunar samples focus& attention on wider aspects of the problem. Pre-Apollo thinking led to the view that the Moon was a primitive undifferentiated object, because of its low density (Urey, 1959). Although this opinion was not uni- versally held (e.g., the mare surfaces were correctly identified as lava flows by Baldwin, 1949), the surprising thing was that the samples both from the maria and the lunar highland crust were very highly fractionated. This is in contrast to estimates of primitive solar nebula values, established from the resemblance between the solar photospheric
and Cl abundances for the non-gaseous elements. Indeed, the highland crustal abundances were so enriched in refractory elements that models ap- peared of crustal formation invoking the late plas- tering on of a refractory-rich layer (e.g., Gast, 1972); however, these were quickly superseded by magma ocean models (Taylor and Jakes, 1974).
It readily became clear that crusts resulted from internal planetary differentiation, and models in- voking late additions of differentiated planetesi- mals to account for them fell into disfavour. Fur- ther exploration revealed that Mercury, Venus, Mars, and many of the larger satellites have surface compositions that differ substantially from any reasonable estimate of their bulk composition; the origin and growth of crusts thus became a general phenomenon in the Solar System.
Crustal origins
Planetary crusts can arise in two basically dif- ferent ways. Firstly, then can form as a result of planetary differentiation consequent upon plane- tary-wide (rather than partial) melting during or shortly following accretion (e.g., lunar highland crust). These may be termed “primary” crusts.
“Secondary” crusts arise later in planetary his- tory as a result of partial melting in planetary interiors. These are t,ypically composed of basalt, the primary melt from silicate mantles (see L.P.I., 1981, for an extensive review). Examples of sec- ondary crusts include the lunar maria, the ter- restrial oceanic crust, the northern hemisphere of Mars, including the great volcanoes, and the Venusian crust. The eucrites, among the meteor- ites, are probable examples of secondary crustal development on an asteroid. Other possibilities include the water-ice crusts of some of the satel- lites composed of rock and ice. In many such examples, the distinction between primary crusts produced by accretional melting and secondary crusts formed by partial melting in the satellite interior, or even solid state resurfacing, must await further study.
“Tertiary” crusts may arise through further melting and differentiation of the extruded material composing the secondary crusts; the con-
tinental crust of the Earth may be the sole exam-
ple of this last type. Growth of primary crusts occurs concomitantly
with or shortly following accretion, and is com- pleted on short time scales (10’ year). Impact induced melting may blur the distinction between primary and secondary crusts. Growth of sec- ondary and tertiary crusts may extend over the lifetime of the planet. The difficulties in producing a tertiary crust are shown by its small mass. Four billion years of growth has resulted in a terrestrial continental crust which is about 0.33% of Earth mass (Taylor and McLennan, 1985). In contrast, the terrestrial secondary oceanic crust is currently formed at a much faster rate, covering most of the planet with a 5 km thick basaltic crust (about 0.1% of Earth mass) within about 200 m.y. Compared with both these examples, the primary lunar bigh- land crust, which amounts to lo-1296 of the moon, was produced in about 100 m-y.
The eariy intense cratering episode
Because of the continuing sweepup of large planetesimals, the growth of primary crusts pro- ceeded in a turbulent environment. Most planets and satellites contain ancient battered surfaces. How long did this embedment last? The age of the Imbrium collision dates, along with the slightly younger Orientale basin formation, the terminal stages of these events on the moon at about 3850 m.y. The cratering record on Mercury and Mars is consistent with the lunar bombardment history, indicating that the flwc is typical of the inner Solar System (Strom, 1987). Thus collisions with plane- tesimais continued for several hundred million years following planetary accretion in the inner solar system. The Iunar highland crust, and pre- sumably other primary crusts on planets and satel- lites, grew in the teeth of this barrage, which accounts for much of the petrographic complexity of the lunar bighhtnd breccias. Since fragments of mare basalt lavas are contained within some of the Apollo 14 breccias (Taylor et al., 1983), some early portions of the secondary lunar crust were caught up in the maelstrom.
The observable surface of the lunar highland crust represents a saturation population of craters
149
and basins (Hartmatm, 1984), and so probably dates from 4100-4200 m.y. Conservative estimates
by W&elms (1985) of the number of objects
which struck the moon between 4440 and 3850 m.y. include about 80 basin-forming events (> 300 km diameter) and over 10,000 craters with diame- ters in the range 30-300 km. Conservative esti- mates for the same interval indicate that over 200 multi-ring basins (> 1000 km diameter) formed on the Earth (Grieve and Parmentier, 1985), prob- ably explaining the absence of identifiable rock units on this planet older than about 3800 m.y.
The bombarding material seems likely to be accretional remnants from rather localized regions of the nebula rather than material from helio- centric orbits; the principal evidence for this is the difference observed in cratering rates. Thus, the rates established for the terrestrial planets differ from those for the Galilean satellites of Jupiter, and for the Uranian satellites, although resem- bling those for the Saturnian satellites (Strom, 1987).
‘Ibe Moon
The lunar highland crust is the best studied example of a primary crust. It is 60-100 km thick on a body whose radius is only 1738 km, thus forming lo-128 of planetary volume. Although it is complex in petrographic detail, mainly because of extensive brecciation during the meteoritic bombardment, it is relatively simple geochem- ically. For example, the rare-earth patterns (chondrite-normallized), for very diverse rock types, are parallel (except for the Eu anomalies) over several orders of magnitude. If the crust was composed of a diversity of igneous rock types, REE patterns would be expected to show a wide range of slopes. The generaBy accepted model is that the lunar highland crust formed during the crystallization of a deep magma ocean, formed concomitantly with accretion of the moon.
Three major components are identified. A principal unit is composed mainly of plagioclase feldspar (ferroan anorthosite) cumulates formed by flotation in a completely dry magma ocean. The Sm-Nd closure age of 4440 f 20 m.y. for the ferroan anorthosite 60025 (Lugmair, 1987) pro-
150
vides a younger limit for plagioclase crystallization
and crustal formation. The final stages of solidifi-
cation of the residual liquid (KREEP) from the
magma ocean may have occurred by about 4350
m.y. (Meyer et al., 1985); in this case about 90
m.y. was apparently required for total crystalliza-
tion. A third component, the Mg-suite appears to
be composed of plutons which intruded the feld-
spathic crust shortly after its formation. The pet-
rogenesis of the Mg suite is still unclear, but they
were intimately mixed with the feldspathic crust
by the meteoritic bombardment. Possibly the Mg
suite represents a secondary crustal component
produced by impact-induced melting. The enig-
matic “KREEP basalts” may also belong in this
category.
The dates for the crystallization of the crust are
close to expected ages for the formation of the
moon. Lunar origin as proposed by the giant
impactor hypothesis (which provides, incidentally,
sufficient energy to produce the massive melting
required for the magma ocean hypothesis), in-
volves collision with the Earth by the next largest
body in the hierarchy (0.1-0.2 earth mass). In
order to account for the low iron content in the
Moon, prior core formation in the impactor is
required. The Moon is derived mostly (80-90s)
from the mantle of the impactor, not from that of
the Earth, accounting for the significant geo-
chemical differences between the two bodies (e.g.,
Taylor, 1987). The energetic nature of the event
provides not only for the lunar magma ocean,
required by the geochemical evidence, but also
probably induces terrestrial mantle melting
(Stevenson, 1987).
However, planetary accretion models based on
the planetesimal hypothesis require up to 100 m.y.
from To (4560) m.y.) to complete planetary assem-
bly. Accordingly, the Moon is likely to form as a
separate entity perhaps only lo-20 m.y. before the
best estimate of the age of 60025. Even allowing
for the large uncertainties in all these ages, it is
apparent that melting of much of the Moon, and
development of the thick lunar highland crust,
proceeded immediately after the formation of the
satellite. Such growth rates are essentially instan-
taneous on a geological time scale.
It is frequently supposed, by analogy with the
Moon, that the Earth formed an early anorthositic
crust. Four reasons make this an unlikely event:
(1) The composition of the moon is probably
richer in Ca and Al than that of the terrestrial
mantle, leading to the early appearance of
plagioclase during crystallization of the lunar
magma ocean. (2) Plagioclase is unstable at shal-
low depths (40 km) in the Earth and will trans-
form to garnet, thus locking up Ca and Al in a
dense phase. In contrast, plagioclase will be stable
in the moon to depths of several hundred kilo-
metres. (3) Plagioclase will not float in a wet
terrestrial basaltic magma. (4) The oldest ter-
restrial anorthosites are not distinct from and
closely resemble younger Archean examples (John
Myers, pers. commun., 1986). They must share the
same petrogenesis; anorthosites derived from a
primordial magma ocean might be expected to
show some petrological and geochemical dif-
ferences.
The observable structure of the lunar crust was
dominated by large basin-forming impacts prior
to 3800 m.y., which produced the circular moun-
tain arcs, so puzzling to early investigators. Later
flooding of the basins by mare basalt derived by
partial melting from the lunar mantle provide a
type example of a secondary crust. Mare basalt
extrusion continued down to about 2500 m.y. The
basalts were derived from cumulates which crys-
tallized from the magma ocean following extrac-
tion of the feldspathic highland crust at about
4400 m.y.
Mercury
Mercury is unique in possessing a metallic
iron-silicate ratio about double that of the other
terrestrial planets. A preferred scenario to account
for this high iron-silicate ratio calls for the loss of
a large portion of the silicate mantle during a
collision with an object about l/6 Mercurian mass
(Cameron and Benz, 1987). Such an event is likely
to have both depleted the planet in volatiles (e.g.
H,O) and triggered mantle-wide melting, so that a
primary lunar-like crust is conceivable. The alter-
native hypothesis that the low silicate/iron ratio
resulted from preferential evaporation (Fegley and
Cameron, 1987) requires such extensive removal
of silicate that the alkali elements are completely
lost. This is at variance with the possible presence of
sodium ions, which are presumably derived by sputtering from the crust, in the tenuous Mercurian atmosphere (Hunten et al., 1988). The presence of this sodium cloud around Mercury is consistent with the surficial presence of sodium-containing minerals, such as plagioclase, or perhaps pyrox- ene. The limited spectral reflectance information from this planet indicates a crust possibly similar in such properties to the lunar highlands (McCord and Clark, 1979). The identification of an absorp- tion feature due to Fe2 +, indicating an Fe0 con- tent of about 5% resembling that of the lunar highlands, however has not yet been confirmed (McCord and Vilas, 1986).
The observable crust was heavily cratered and must be older than 4000 m.y. by analogy with the dated lunar highland crust. The presence of lobate fault scarps of about the same age indicates a slight contraction of about 2 km in radius. Since these scarps both cut through some craters, and are in turn cut by other craters, they must predate the final stages of the massive bombardment. From the lunar analogy, this indicates a date for the contraction prior to 4000 m.y. Accordingly, the surface of Mercury provides evidence of stability since the contraction ceased. In summary, the crust is very likely to be primary (a result of initial planetary melting) although the origin of the smooth plains (whether ejecta blankets or lava flows) is unresolved.
Mars
North of a boundary inclined at about 28’ to the equator, the Martian surface consists of volcanic plains and large volcanoes, all probably basaltic (Carr, 1981). In contrast, the southern hemisphere of Mars is broadly composed of an Ancient Cratered Terrain, older than about 4000 m.y. based both on crater counting and the lunar analogy. Its composition is unknown, but it seems unlikely to be acidic or very different in composi- tion from the basaltic plains. The rationale for this is that the Viking Lander XRF data at the two sites 4000 km apart were both similar, and basaltic
151
in composition (Table 1). The fine material analysed presumably represents a planetary-wide
dust average (analogous to loess) and could be expected to provide some kind of average sample of the surface. It contains no component suggestive that the average crust is more silicic than basalt.
The SNC meteorites suggest a similar scenario (McSween, 1985). If, as appears likely, they are derived from Mars by meteoritic impact, then they must come from within a few kilometres of the surface and so also represent crustal or uppermost mantle material. Their composition (Table 1) is ultrabasic and does not encourage the speculation that Mars possesses a low density granitic crust analogous to the terrestrial continental crust. Accordingly, the Martian crust appears to be dominated by rocks with low silica contents. The ages of the volcanic plains, based on crater count- ing, extends over much of geologic time and so they form another example of a secondary crust, derived by partial melting from the Martian man- tle.
The origin of the Ancient Cratered Terrain, and whether it represents primary, or as judged here, secondary crust, must await further data. How- ever, one might expect that it has contributed significantly to the composition of the fine material
TABLE 1
Major element composition of Mars, Venus and lunar surface
material (all data in wt.% normalized to XXX)
Mars Venus Moon
(1) (2) (3) (4) (5) (6)
SiO, 55.1 50.5 46.8 49.8 51.5 45
TiO, 0.78 0.89 1.66 1.28 0.23 0.56
AlaO, 9.1 7.74 16.4 18.3 18 24.6
Fe0 19.7 20.2 9.7 9.0 8.7 6.6
MS0 7.5 9.13 11.7 8.3 13 6.8
CaO 7.1 9.82 7.4 10.5 8.5 15.9
Na,O - * 1.50 (2.1) b (2.5) b - a 0.45
GO < 0.6 0.19 4.2 0.2 0.1 0.075
x loo 100 loo loo loo 100
Analyses, (1) Viking Lander, Chryse Planitia (Clarke et al.,
1982); (2) Shergotty meteorite bulk analysis (Lam et al., 1986,
table 1, Maim. data); (3) Venera 13 (Surkov, 1984); (4) Venera
14 (Surkov, 1984); (5) Vega 2 (Surkov et al., 1986); (6) Lunar
highland crust (Taylor, 1982). a_ no data. b (f calculated value.
152
analysed by the Viking Landers. The terrain is heavily cratered, and so can be expected to possess a high proportion of dust and finely comminuted debris. This material is likely to be a significant component in the planetary-wide dust storms. Accordingly, the composition of the martian fine material might even be biased toward that of the Ancient Cratered Terrain, unless dust derived from that terrain was locked up early in polar layered deposits. If this interpretation is correct, then the Ancient Cratered Terrain in basaltic in composi- tion
In summary, the observable crust on Mars ap- pears to be basaltic grading into ultramafic, and is here interpreted as being derived from the mantle by partial melting, hence secondary.
Venus
Although the low lying areas of Venus are thought to be basaltic, the existence of small con- tinental-sized areas such as Aphrodite Terra and Ishtar Terra raises the possibility that these high standing areas might result from the presence of low density rocks, analogous to the terrestrial con- tinents. The early gamma-ray data from Venera 8 (Table 2) indicated high K (4%), U (2.2 ppm) and Th (6.5 ppm) values. These are typical of abun- dance levels observed in terrestrial granites, and immediately led to the speculation that granite was present on the surface of Venus. This would imply an evolution for Venus resembling that of the Earth.
The Venera 9 and 10 gamma-ray data, in con- trast, gave low values for K, U, and Th, consistent with the levels seen in terrestrial basaltic rocks. The Venera 13 and 14 XRF experiment resolved this problem by providing data on the major
TABLE 2
Potassium, uranium and thorium concentrations in Venus
surface rocks data from Surkov et al. (1987)
Venera Vega _
8 9 10 1 2
K (%) 4.0 0.47 0.30 0.45 0.40
U (ppm) 2.2 0.60 0.46 0.64 0.68
Th (ppm) 6.5 3.65 0.70 1.5 2.0
element composition of the surface material (Surkov, 1984). The Venera 13 analysis resembles terrestrial alkali-basalt while the Venera 14 data is
close in composition to MORB. (The high SO, and Cl values must represent surficial volcanic additions, in the absence of an oceanic sink for these elements, as is also the case for Mars which shows similar high values.) The high K,O value of 4% in the alkali-basaltic composition revealed by Venera 13 most probably indicates that the high K, U and Th values in the Venera 8 gamma-ray data (Surkov, 1977) are also due to alkali basalt and do not indicate the presence of granite.
All these analyses are from near Beta Regio, generally interpreted as a volcanic construct, or from low-lying areas, thought to be basaltic on topographic grounds. However, the Vega 2 XRF experiment provided major-element data from the eastern flank of the high-standing Aphrodite Terra. This analysis also turned out to be basaltic (Table l), with a very low value for K,O confirmed by the gamma ray data (Surkov et al., 1987) (Table 2). Gamma ray data from Vega 1 (Table 2) which also landed in a mountainous area (Mermaid Val- ley) showed similar low values for K, U and Th. These data provide no support for low density silica-rich highland areas on Venus.
Thus the Venusian crust appears to be dominated by basaltic lavas and the presence of extensive areas of more fractionated rocks is con- jectural, Based upon this interpretation, the high standing regions of Aphrodite Terra and Ishtar Terra are due to tectonic rather than composi- tional controls (Basilevsky et al., 1986). This is generally supported by the topographical informa- tion, which appears to indicate extensional defor- mation in equatorial regions (e.g., Aphrodite Terra) and compressional deformation in northern high latitudes (Atalanta Planitia and Ishtar Terra) (e.g.,
Schaber, 1982; Crumpler et al., 1986; Head, 1986), indicating lateral movement of the basaltic crust over a weak substratum.
The surface of Venus appears to be relatively young, based on crater counting studies, but estimates vary widely from about one billion years (Ivanov et al., 1986) to as young as about 100 m.y. (Schaber et al., 1987). If the surface is typical of the crust at depth, then Venus possesses a sec-
ondary crust, produced by partial melting in the
planetary interior. It is interesting that the heat-
producing elements, K, U and Th, are being con-
centrated near the planetary surface in a manner
different from that observed terrestrially, where
they are concentrated in the siliceous continents.
As is the case for Mars, the interpretation of
crustal composition is based on surface character-
istics. Over the considerable range in elevations
observed, the composition appears to be basaltic.
This, coupled with the relative youth of the surface
(< 1 b.y.), indicates that the observable crust is
secondary.
Icy crusts
Most of the satellites of the outer planets have
low densities, consistent with rock-ice mixtures
(Bums and Matthews, 1986). The physical proper-
ties of ice have led to some interesting and unique
features. The Galilean satellites of Jupiter are of
special interest (10 is a separate category, with its
extensive volcanic activity driven by tidal heating
due to its proximity (6 Jupiter radii) to the planet).
Although Europa is mainly rock (r = 1536 km;
r = 2.99 g/cm3), it appears to have a water ice
surface. Melting of ice (ice-ammonia mixtures
seem less likely in the Jovian sub-nebula than in
the cooler Saturnian system) in the interior, prob-
ably due to accretional heating, has led to forma-
tion of an ice surface, overlying a rock-ice inter-
ior. Whether it is a primary crust, due to accre-
tional heating, or a secondary crust derived by
melting of ice within the satellite is not clear. It is
essentially free of craters and so must be continu-
ally resurfaced, possibly by tidal heating from
Jupiter. The third Galilean satellite, Ganymede is
discussed below. Callisto, in contrast to Europa,
has the most heavily cratered surface in the solar
system, indicating that the crust dates from very
early in solar system history, and so may represent
a primary crust formed during accretion. As the
outermost Galilean satellite, it apparently escaped
heating by tidal interaction with Jupiter.
Several of the satellites of Saturn (e.g., Encel-
adus) show evidence of some resurfacing and
tectonic activity, perhaps connected with melting
of ice or ice-ammonia mixtures. The Uranian
153
satellites, especially Miranda and Ariel, display
similar features. These represent a fertile field for
research into crustal evolution on icy bodies.
Ganymede
Ganymede has two distinct crustal types. An
ancient dark cratered terrain, which is probably a
mixture of ice and silicate, is fractured and in-
truded by younger grooved terrain, apparently
composed of water ice, which appears to have
filled wide rift zones in the older crust. The most
reasonable interpretation is that the younger
grooved terrain is derived from the interior by
later internal melting. The older more heavily
cratered and darker crust appears to have been
split and intruded by the younger event, accompa-
nied by some minor planetary expansion (e.g.,
Murchie et al., 1986).
A probable explanation was presented by
Squires (1980). This assumes that Ganymede was
formed from a mixture of rock and water ice. The
internal pressure is sufficient to convert ice to high
pressure polymorphs (e.g., Ice Vl or Vlll, with
densities up to 1.67 g/cm3, Fig. 1). Heating in the
interior, due to the radioactive elements K, U and
Th in the silicate components, will melt the water
ice, which will then rise to the surface and freeze
150 200 250 300 350 4
Temperature (I<)
Fig. 1. The temperature-pressure phase diagram for water ice.
The numerals refer to the various polymorphs in each stability
field. Ice VIII has a density of 1.67 g/cm3, so that a change
from ice VIII to ice I will result in a volume expansion.
154
as Ice 1 leading to a significant density decrease and to minor planetary expansion of a few percent in radius (Squires, 1980).
The grooved terrain is thus an example of a secondary crust, while the ancient more heavily cratered dark terrain may represent the undif- ferentiated bulk satellite composition, or less likely, a primary crust. Ganymede, the largest satellite in the solar system, is thus of extraordinary interest. It displays evidence of slight expansion, which is explicable from the known physical properties of ice.
As noted above, the outermost Galilean satel- lite, Callisto, possesses, in sharp contrast to Ganymede, a heavily cratered apparently primor- dial crust, unaltered since the termination of the early massive bombardment. Callisto (r = 2410 km; r = 1.83 g/cm3) is both smaller and less dense than Ganymede (r = 2638 km; r = 1.93 g/cm3). According to Sotin et al. (1987) this difference is likely to lead to a warmer mantle in Ganymede, leading to melting of the ice Vl layer a few hundred million years following accretion, while the smaller size and density of Callisto have inhibited this evolutionary development.
Tertiary crusts
The continental crust of the Earth appears to be the only example known of a tertiary crust (Taylor and McLennan, 1985). Its composition is far removed from the bulk composition of the planet or of primary melts from the mantle. Al- though it comprises only 0.57% of the mass of the mantle, it contains over 20% of the abundances of several of the incompatible elements (e.g., Cs, Rb, Tl, K, Ba, U and Th; Taylor and McLennan, 1985, fig. 11.6).
It probably owes its existence to the presence of extensive liquid water at the surface (Campbell and Taylor, 1983) and to subduction of the sec- ondary basaltic crust and accordingly is likely to be unique. According to Taylor and McLennan (1985), its growth rate is slow, having taken about 4000 m.y. to reach its present mass, although much of this growth appears to have taken place in an episodic manner, particularly rapid in the late Archean.
It also seems to have undergone intra-crustal differentiation, leading to geochemically distinct upper and lower crusts, principally through the extraction of a low-melting granitic fraction. In this survey of planetary crusts, no candidates from other planets or satellites appear to resemble the terrestrial continental crust, and the familiar granite is most likely a very rare rock type elsewhere in the solar system. This serves to re- mind us of the unique nature of this planet, and of the hazards of extrapolating the laboriously gathered knowledge of terrestrial geology and geo- chemistry to elsewhere in the solar system.
In conclusion, a few general features about the growth rate of crusts may be gleaned through the complex series of crust-forming events on planets and satellites in the solar system.
There are very different growth rates for primary, secondary and tertiary crusts. The lunar highland crust, forming lo-12% of the moon was produced within about 100 m.y. Secondary crusts are produced at much slower rates. Thus the lunar maria, produced by partial melting from the lunar interior over a period exceeding lo9 years are only 0.1% of lunar mass. The oceanic crust of the Earth, which is 0.1% of the mass of the planet, is presently forming from the mantle in about 200 m.y. If the present rate is extrapolated back, the total volume of oceanic crust produced is about 2% of the mass of the planet. Over 4000 m.y. of growth of the continental crust of the Earth has resulted in the production of a crust of 0.33% of planetary mass. The latter process is clearly not efficient.
Recycling of crust appears to be rare on a planetary scale, so that crustal growth in the rest of the solar system is essentially an irreversible process. This survey provides no evidence in sup- port of the presence of primary siliceous crusts analogous to those of the terrestrial continental crust, and thus does not favour no-growth re- cycling models.
Models which involve rifting and separation of an initial world encircling crust by massive plane- tary expansion receive no support from this study
155
(see also Taylor, 1983). The lunar crust preserves a
frozen fossil surface for over 4000 m.y. The
Mercurian crust, of similar age, provides evidence
of slight contraction, while the unequivocal expan-
sion of a few percent of Ganymede is explicable
by the polymorphic transitions of ice.
Finally, the early intense cratering history has
affected all planetary crusts prior to about 3800
m.y. The principal effect, as well shown in the
lunar highland samples, is extensive brecciation.
The record of such events will be retained on most
planets and satellites because of the paucity of
effective eroding agents. On the Earth however,
the record is unlikely to be preserved because of
the efficiency of terrestrial weathering, eroding,
and transporting agents working on the debris
resulting from several hundred Mare Orientale-
scale impacts. The breccias would have been easy
meat.
Acknowledgements
I am grateful to Graham Ryder and Larry
Taylor for composing perceptive reviews.
References
Baldwin, R.B., 1949. The Face of the Moon. Chicago Univ.
Press, Chicago, Ill.
Basilevsky, A.T., Pronin, A.A., Ronca, L.B., Kryuchkov, V.P.,
Sukhanov, A.L. and Markov, M.S., 1986. Styles of tectonic
deformation on Venus: analysis of Venera 15 and 16 data.
J. Geophys. Res., 91: D3999D411.
Bums, J.A. and Matthews, M.S., 1986. Satellites. Univ. Arizona
Press, Tucson, Ariz., 1021 pp.
Cameron, A.G.W. and Benz, W., 1987. Planetary collision
calculations: origin of Mercury. Lunar Planet. Sci., XVIII:
151-152.
Campbell, I.H. and Taylor, S.R., 1983. No water, no granites -
no oceans, no continents. Geophys. Res. Lett., 10:
1061-1064.
Carr, M.H. 1981. The Surface of Mars. Yale Univ. Press, New
Haven, Conn., 232 pp.
Clark, B.C., Baird, A.K., Weldon, R.J., Tsuaki, D.M., Schna-
bel, L. and Candelaria, M.P., 1982. Chemical composition
of martian fines. J. Geophys. Res., 87: 10059-10067.
Crumpler, L.S., Head, J.W. and Campbell, D.B., 1986. Oro-
genie belts on Venus. Geology, 14: 1031-1034.
Fegley, B. and Cameron, A.G.W., 1987. A vaporization model
for iron/silicate fractionation in the Mercury protoplanet.
Earth Planet. Sci. Lett., 82: 207-222.
Gast, P.W., 1972. The chemical composition and structure of
the moon. Moon, 5: 121-148.
Grieve, R.A.F. and Parmentier, E.M., 1985. Considerations of
large scale impact and the early Earth. L.P.I. Tech. Rep.,
85-01: 23-24.
Hartmann, W.K., 1984. Does crater “saturation equilibrium”
occur in the solar system? Icarus, 60: 56-74.
Head, J.W., 1986. Ishtar Terra, Venus: a simple model of
large-scale tectonic convergence, crustal thickening and
possible delamination. Proc. Lunar Planet. Sci. Conf., 17:
323-324.
Hunten, D.M., Shemansky, D.E. and Morgan, T.M., 1989. The
Mercury atmosphere. In: C.R. Chapman, F. Vilas and M.S.
Mathews (Editors), Mercury. Univ. Ariz. Press, Tucson,
Ariz., in press.
Ivanov, B.A., Basilevsky, A.T., Kryuchkov, V.P. and Chemaya,
I.M., 1986. Impact craters on Venus: analysis of Venera 15
and 16 data. J. Geophys. Res., 91: D413-D430.
Johnson, R.E., Nelson, M.L., Hawke, B.R., Bell, J.F. and
Cintala, M.J., 1986. Mineralogical implications of Mercury’s
sodium cloud. Abstr. Conf. on Mercury, Tucson, Ariz.
Laul, J.C., Smith, M.R., Wanke, H., Jagoutz, E., Dreibus, G.,
Palme, H., Spettel, B., Burghele, A., Lipschutz, M.E. and
Verkouteren, R.M., 1986. Chemical systematics of the
Shergotty meteorite and the composition of its parent body
(Mars). Geochim. Cosmochim. Acta, 50: 909-926.
L.P.1, 1981. Basaltic Volcanism on the Terrestrial Planets.
Lunar and Planetary Institute, Houston, Tex., 1286 pp.
Lugmair, G.W., 1987. The age of the lunar crust:
60025-Methuselah’s legacy. Lunar Planet. Sci., XVIII:
584-585.
McCord, T.B. and Clark, R.N., 1979. The Mercury soil: pres-
ence of Fe’+. J. Geophys. Res., 84: 7664-7668.
McCord, T.B. and Vilas, F., 1986. Reflectance spectropho-
tometry of Mercury: implications for the surface composi-
tion. Abstr. Conf. on Mercury, Tucson, Ariz.
McSween, H.Y., 1985. SNC meteorites: clues to Martian petro-
logic evolution. Rev. Geophys., 23: 391-416.
Meyer, C., Compston, W. and Williams, I.S., 1985. Lunar
zircon and the closure age of the lunar crust. Lunar Planet.
Sci., XVI: 557-558.
Murchie, S.L., Head, J.W., Helfenstein, P. and Plescia, J.B.,
1986. Terrain types and local-scale stratigraphy of grooved
terrain of Ganymede. J. Geophys. Res., 91: E222-E238.
Schaber, G.G., 1982. Venus: limited extension and volcanism
along zones of lithospheric weakness. Geophys. Res. Lett.,
9: 499-502.
Schaber, G.G., Shoemaker, E.M. and Kopzak, R.C., 1987. The
surface age of Venus: use of the terrestrial cratering record.
Solar System Res., in press.
Sotin, C., Parmentier, E.M., Murchie, S.L. and Head, J.W.,
1987. Convective instability of phase transitions: applica-
tion to the icy satellites Ganymede and Callisto. Abstr.
IUGG, 19th Meet., Vancouver, UlO-41: 100.
Squires, S.W., 1980. Volume changes in Ganymede and Calli-
sto and the origin of grooved terrain. Geophys. Res. Lett.,
7: 593-596.
156
Squires, S.W., 1987. Resurfacing and tectonism on icy satel-
lites. Abstr. IUGG, 19th Meet., Vancouver, UlO-28: 96.
Stevenson, D.J., 1987. Origin of the moon-the collision hy-
pothesis. Amm. Rev. Earth Planet. Sci., 15: 271-315.
Strom, R.G., 1987. The solar system cratering record: Voyager
2 results at Uranus and implications for the origin or
impacting objects. Icarus, 70: 517-535.
Surkov, Yu.A., 1977. Geochemical studies of Venus by Venera
9 and 10 automatic interplanetary stations. Proc. Lunar Sci.
Conf., 8: 2665-2689.
Surkov, Yu.A., 1984. New data on the composition, structure
and properties of Venus rock obtained by Venera 13 and
Venera 14. J. Geophys. Res., 89: B393-B402.
Surkov, Yu.A., Moskalyova, L.P., Kharyukova, V.P., Dudin,
A.P., Smimov, G.G. and Zaitseva, S. Ye., 1986. Venus rock
composition at the Vega 2 landing site. Lunar Planet. Sci.,
XVII: 849-850.
Surkov, Yu.A., Kimozov, F.F., Glazov, V.N., Dunchenko,
A.G., Tatsy, L.P. and Sobomov, O.P., 1987. Uranium,
thorium and potassium in the Venusian rocks at the land-
ing sites of Vega 1 and 2. J. Geophys. Res., 92: E537-E540.
Taylor, L.A., Shervais, J.W., Hunter, R.H., Shih, C.-Y., Bansal,
B.M., Wooden, J., Nyquist, L.E. and Laul, J.C., 1983.
Pre4.2 AE mare-basalt volcanism in the lunar highlands.
Earth Planet. Sci. Lett., 66: 33-47.
Taylor, S.R., 1982. Lunar and terrestrial crusts: a contrast in
origin and evolution. Phys. Earth Planet. Inter., 29:
223-241.
Taylor, S.R., 1983. Limits to Earth expansion from the surface
features of the Moon, Mercury, Mars and Ganymede. In:
Expanding Earth Symp. Univ. Tasmania, Hobart, pp.
343-347.
Taylor, S.R., 1987. The unique lunar composition and its
bearing on the origin of the Moon. Geochim. Cosmochim.
Acta, 51: 1297-1309.
Taylor, S.R., 1988. Planetary compositions. In: J.F. Kerridge
and M.S. Matthews (Editors), Meteorites and the Early
Solar System. Section 7.8. Univ. Ariz. Press, Tucson, Ariz.,
512-534.
Taylor, S.R. and Jakes, P., 1974. The geochemical evolution of
the moon. Proc. Lunar Sci. Conf., 5: 1287-1305.
Taylor, S.R. and McLennan, S.M., 1985. The Continental
Crust: Its Composition and Evolution. Blackwell, Oxford,
312 pp.
Urey, H.C., 1959. Primary and secondary objects. J. Geophys.
Res., 64: 1721-1737.
Wetherill, G.W., 1986. Accumulation of the terrestrial planets
and implications concerning lunar origin. In: W.K. Hart-
mann et al. (Editor), Origin of the Moon. Lunar and
Planetary Institute, Houston, Tex., pp. 519-550.
Willhelms, D.E., 1985. Lunar impact rates reconsidered. Lunar
Planet. Sci., XVI: 904-905.