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

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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

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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.

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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

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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.

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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.

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