of 26
Theory of the Earth
Don L. Anderson
Chapter 1. The Terrestrial Planets
Boston: Blackwell Scientific Publications, c1989
Copyright transferred to the author September 2, 1998.
You are granted permission for indi
vidual, educational, research and
noncommercial reproduction, distribution, display and performance of this work
in any format.
Recommended citation:
Anderson, Don L. Theory of the Earth.
Boston: Blackwell Scientific Publications,
1989.
http://resolver.calt
ech.edu/CaltechBOOK:1989.001
A scanned image of the entire book may
be found at the following persistent
URL:
http://resolver.caltech.edu/CaltechBook:1989.001
Abstract:
Earth is part of the solar sy
stem. Although it is the most
studied planet, it cannot
be completely understood in isolation. The chemistry of meteorites and the Sun
provide constraints on the composition of th
e bulk of the Earth. The properties of
other planets provide ideas for and tests
of theories of planetary formation and
evolution. In trying to understand the or
igin and structure of the Earth, one can
take the geocentric approach or the
ab initio
approach. In the former, one
describes the Earth and attempts to work backward in time. For the latter, one
attempts to track the evolution of the solar nebula through collapse, cooling,
condensation and accretion, hoping that
one ends up with something resembling
the Earth and other planets. In Chapter
1 I develop the external evidence that
might be useful in understanding the Earth. In Chapter 2 I describe the Earth
and Moon.
-
-
-
e Terrestrial
Planets
I
want
to
know how God created
this world.
I
am
not interested in this
or that
phenomenon, in
the
spectrum
of
this or that
element.
I
want
to
know
his
thoughts, the rest are details.
E
arth is part
of
the
solar system. Although it
is the
most
studied planet,
it cannot
be
completely understood
in
isolation. The chemistry
of
meteorites and the Sun provide
constraints
on
the composition
of
the
bulk
of
the Earth. The
properties
of
other planets provide ideas
for
and
tests
of
theories
of
planetary formation
and
evolution. In trying
to
understand
the origin
and
structure
of
the Earth, one can
take the geocentric approach
or
the
ab
initio
approach. In
the former, one describes
the
Earth
and
attempts to
work
backward
in
time. For the latter, one attempts
to
track the
evolution
of
the solar
nebula through
collapse, cooling,
condensation and accretion,
hoping that
one ends
up
with
something resembling the Earth
and
other planets.
In
Chap
-
ter
1
1
develop the external evidence that
might be
useful
in
understanding the Earth.
In
Chapter
2
I
describe the Earth
and
Moon.
THEORIES
OF
PLANETARY
FORMATION
The nature
and evolution
of
the solar
nebula and
the for
-
mation
of
the planets are complex
and difficult
subjects.
The fact that terrestrial planets did in fact form
is
a suffi
-
cient motivation to keep a
few
widely
dispersed scientists
working
on these problems. There are several possible
mechanisms
of
planetary growth. Either the planets
were
assembled from smaller
bodies
(planetesimals), a piece
at a
time, or
diffuse
collections
of
these bodies, clouds, became
gravitationally unstable and collapsed to
form
planetary
sized objects. The planets,
or
protoplanetary nuclei, could
have
formed in
a gas
-
free environment
or
in the presence
of
a large amount
of
gas
that
was
subsequently dissipated.
The planets are
now
generally thought to
have
origi
-
nated
in a slowly rotating disk
-
shaped
"
solar nebula
"
of
gas
and
dust
with
solar composition. The temperature
and
pres
-
sure
in
the hydrogen
-
rich
disk
decrease
both
radially
from
its center
and outward
from
its
plane. The disk cools
by
radiation,
mostly
in
the
direction normal to the plane,
and
part
of
the incandescent gas condenses
to solid
"
dust
"
par
-
ticles.
As
the particles
grow,
they
settle to the
median plane
by
processes
involving
collisions
with
particles
in
other or
-
bits,
by
viscous gas
drag
and
gravitational attraction
by
the
disk. The
total
pressure
in
the vicinity
of
Earth's
orbit
may
have
been
of
the order
of
to
bar. The particles
in
the
plane
probably formed
rings and
gaps. The sedimenta
-
tion
time
is fairly rapid,
but
the processes
and
time scales
involved
in
the collection
of
small objects into planetary
sized objects are
not
clear. The
common thread
of
all
cos-
mogonic
theories
is that
the
planets formed from dispersed
material,
that
is,
from
a protoplanetary nebula. Comets,
some meteorites
and
some
small satellites
may
be
left
over
from these early stages
of
accretion.
The following observations are the
main
constraints
on
theories
of
planetary origin:
1.
Planetary orbits are
nearly
circular, lie virtually in a
single plane,
and
orbit
in
the same sense
as
the
Sun's
rotation. The
Sun's
equatorial plane
is close
to the orbital
plane. The planets exhibit a preferred sense
of
rotation.
2.
The distribution
of
planetary distances is regular
(Bode's
Law).
3.
The planets group into compositional classes
related
to
distance from the Sun. The
inner,
or terrestrial
planets
(Mercury,
Venus,
Earth
and
Mars), are small,
have high
density,
slow
rotation rates
and few
satellites. The
Moon
is often
classified
as
a terrestrial planet. The giant
plan
-
ets (Jupiter, Saturn,
Uranus and
Neptune)
are large,
have
low
density, rotate rapidly
and
have numerous
satellites.
Although the
Sun
contains
more than
99
percent
of
the
2
THE
TERRESTRIAL
PLANETS
mass
of
the solar system, the planets contain more
than
98 percent
of
the angular momentum.
Apart from
the
mechanisms
of
accretion
and
separa
-
tion
of
planetary from solar material,
there
are
several
im
-
portant
unresolved
questions.
How
dense
was
the protoplanetary nebula? A
lower
limit
is found
by
taking
the present mass
of the
planets
and
adding
the
amount
of
light elements necessary
to
achieve
solar composition.
This gives about
lo-,
solar mass.
Young
stars
in
the
initial stages
of
gravitational contraction expel
large quantities
of
matter, possibly accounting for several
tens
of
percent
of
the star's
mass. Some theories therefore
assume
a massive early nebula that
may
equal twice
the
mass
of
the
Sun
including the
Sun's
mass.
T
-
Tauri
stars, for
example, expel
about
solar
mass per year
for
lo5
to
lo6
years.
What
are
the time scales
of
cooling, separation
of
dust
from gas,
growth
of
asteroidal size bodies,
and
growth
of
planets from meter
-
to kilometer
-
size objects?
If cooling is
slow
compared to the
other processes,
then
planets
may
grow during
cooling
and
will
form
inhomogeneously.
If
cooling
is fast,
then the planets
may
form
from cold mate
-
rial
and
grow
from
more
homogeneous material.
The accretion
-
during
-
condensation,
or
inhomogeneous
accretion, hypothesis
would lead to
radially
zoned
planets
with
refractory
and
iron
-
rich cores,
and
a compositional
zoning
away
from the Sun;
the
outer planets
would
be
more
volatile
-
rich. Superimposed
on
this effect
is a size ef
-
fect; the larger planets,
having
a larger gravitational cross
section, collect
more
of
the
later condensing (volatile)
material.
The
Safronov (1972)
cosmogonical theory
is currently
the
most
popular. It is assumed that the Sun initially pos
-
sessed
a uniform gas
-
dust nebula. The nebula
evolves
into
a torus
and
then
into a disk. Particles
with
different eccen
-
tricities
and
inclinations collide
and
settle
to
the
median
plane within
a few
orbits.
As
the disk gets denser, it goes
unstable
and
breaks
up
into
many
dense accumulations
where the
self
-
gravitation exceeds the disrupting tidal force
of
the Sun.
As
dust is
removed
from the
bulk
of
the nebula,
the
transparency
of
the
nebula
increases,
and
a large tem
-
perature gradient is established
in
the nebula.
The
mechanism
for
bringing
particles together
and
keeping
them together
to
form large planets
is obscure. A
large
body, with
an
appreciable gravity field,
can
attract
and
retain planetesimals. Small particles colliding at
high
speed
disintegrate
and
have such small
gravitational cross section
that
they can
attract
only nearby
particles. Large collections
of
co
-
rotating particles,
with
minimum
relative velocities,
seems to be
a prerequisite condition. Self
-
gravitation
of
the
aggregate can
then bring
the particles together. Small
bod
-
ies might
also act
as
condensation nuclei
and
therefore
add
material
directly from
the gaseous phase.
In
the
Safronov
theory,
accumulation
of
97
-
98 percent
of
the Earth oc
-
curred
in
about
los
years. In
other
theories the accretion
time
is much
shorter,
lo5-
lo6
years.
If
the relative
velocity between
planetesimals
is
too
high, fragmentation rather
than
accumulation
will
dominate
and
planets
will not grow.
If relative velocities are too
low,
the
planetesimals
will
be
in
nearly
concentric orbits
and
the
collisions
required
for
growth will not
take place.
Safronov
(1972)
showed that
for plausible assumptions regarding dis
-
sipation
of
energy
in
collisions
and
size distribution
of
the
bodies,
mutual
gravitation causes the
mean
relative
veloci
-
ties
to
be
only
somewhat
less
than
the escape velocities
of
the larger bodies. Thus, throughout the entire course
of
planetary growth, the
system
regenerates
itself such
that the
larger bodies
would always grow.
The initial stage
in
the
formation
of
a planet
is the
con
-
densation
in
the
cooling
nebula. The
first
solids appear
in
the
range
1750-1600
K
aaac:
are
oxi;
a,
silicates
and
titan
-
ates
of
calcium
and
aluminum
(such
as
Al,O,,
CaTiO,,
Ca,Al,Si,O,)
and
refractory
metals such
as the platinum
group. These minerals (such
as
corundum, perovskite, mel
-
ilite)
and
elements are found
in
white inclusions (chon
-
drules)
of
certain meteorites,
most notably
in
Type
I11
car
-
bonaceous chondrites. Metallic iron condenses
at relatively
high
temperature
followed
shortly
by
the bulk
of
the silicate
material as forsterite
and
enstatite.
FeS
and
hydrous miner
-
als appear
only at very
low
temperature, less
than
700
K.
Volatile
-
rich
carbonaceous chondrites
have
formation
tem
-
peratures
in
the range 300
-
400
K,
and
at least
part
of
the
Earth
must have accreted
from material
that
condensed
at
these
low
temperatures. The presence
of
CO,
and
H,O
on
the Earth
has led some to
propose that the Earth
was
made
up
entirely
of
cold carbonaceous chondritic
material
-
the
cold
accretion hypothesis. Turekian
and
Clark
(1969) assume
that
volatile
-
rich
material
came
in as
a late
veneer
-
the
in
-
homogeneous
accretion
hypothesis.
Even
if
the Earth ac
-
creted slowly, compared
to cooling
and
condensation times,
the later stages
of
accretion
could
involve
material
that
con
-
densed
further
out
in
the nebula and
was
later perturbed into
the inner solar system. The Earth
and
the
Moon
are
deficient
in
not
only the
very volatile
elements
that make
up
the
bulk
of
the
Sun
and
the outer planets, but also
the
moderately
volatile elements such as sodium, potassium, rubidium
and
lead.
A large amount
of
gravitational energy is released as
the particles fall onto a growing Earth, enough to raise the
temperature
by
tens
of
thousands
of
degrees
and
to
evapo
-
rate the Earth
back
into space
as
fast as it forms. There are
mechanisms for buffering the temperature rise
and
to retain
material even
if it vaporizes,
but melting and
vaporization
are likely once the proto
-
Earth
has
achieved
a given
size,
say
lunar size. The mechanism
of
accretion
and
its
time
scale determine the fraction
of
the
heat that
is retained,
and
therefore the temperature
and heat
content
of
the growing
Earth. The
"
initial
"
temperature
of
the
Earth
is
likely to
have been high even
if it formed from cold planetesimals.
A rapidly
growing
Earth retains more
of
the gravitational
energy
of
accretion, particularly
if there are large
impacts
that can
bury
a large fraction
of
their gravitational energy.
Evidence
for
early
and
widespread melting on
such small
THEORIES
OF
PLANETARY
FORMATION
3
objects
as
the
Moon
and
various meteorite parent bodies
attests
to
the importance
of
high
initial temperatures,
and
the energy
of
accretion
of
the Earth
is
more than
15 times
greater
than
that
for
the
Moon. The intensely
cratered
sur
-
faces
of
the solid planets
provide
abundant
testimony
of
the
importance
of
high
-
energy
impacts in the later stages
of
accretion.
The
initial
temperature distribution in a planet can
be
estimated
by
using
the equation
of
conservation
of
energy
acquired during accretion:
where
t is the time,
p
is the density
of
the accreting parti
-
cles,
G
is the gravitational constant,
M(r)
is the
mass
of
a
growing
planet
of
radius
r,
a
is the
Stefan
-
Boltzmann con
-
stant,
E
is
the emissivity,
C,
is the specific
heat,
T
is the
temperature
at
radius
r,
and
T,
is
the
blackbody radiation
temperature. The equation
gives
the balance
between
the
gravitational energy
of
accretion, the energy radiated into
space
and
the thermal
energy produced
by
heating
of
the
body.
Latent heats associated
with
melting
and
vaporization
are also involved
when
the surface temperature gets
high
enough. The ability
of
the
growing body
to radiate
away
part
of
the heat
of
accretion depends on
how much
of
the
incoming material remains near the surface
and
how
rapidly
it is covered
or
buried.
An
impacting
body
cannot
bury
all
of
its
heat since heat is transferred to the planetary material,
and both
the projectile
and
target fragments are
thrown
large distances
through
the atmosphere, cooling during
transit
and
after spreading
over
the surface. Parts
of
the
pro
-
jectile
and
ejecta
are
buried and
must,
of
course, conduct
their heat to the surface before the heat can
be
radiated
back
to space. Gardening
by
later
impacts
can also
bring
buried
hot material to
the
surface. Devolatization
and
heating as
-
sociated
with
impact can
be
expected to generate a hot,
dense atmosphere that serves to keep the surface tempera
-
ture hot
and
to
trap
solar radiation.
Intuitively,
one would
expect the early stages
of
accre
-
tion to
be slow,
because
of
the
small
gravitational cross sec
-
tion
and
absence
of
atmosphere,
and
the terminal stages to
be slow,
because the particles are
being
used
up. A conve
-
nient expression for the rate
of
accretion
that has
these char
-
acteristics is (Hanks
and
Anderson, 1969)
dr
-
-
-
k,t
2
sin
k,t
dt
where
k,
and
k2
can
be
picked to give a
specified final
radius
and
accretion time. The temperature
profile
resulting
from
this
growth
law
gives a planet
with
a cold interior, a tem
-
perature
peak
at
intermediate depth,
and
a cold outer layer.
Superimposed on this
is the temperature increase
with
depth
due
to
self
-
compression
and
possibly higher temperatures
of
the early accreting particles.
However,
large late im
-
pacts,
even
though infrequent, can heat
up
and melt
the
up
-
per mantle.
One limiting case
is that at every stage
of
accretion the
surface temperature
of
the Earth is
such
that
it radiates en
-
ergy
back
into the
dust cloud
at precisely the rate
at
which
gravitational energy
is released
by
dust
particles free
-
falling
onto its surface.
By
assuming
homogeneous accretion
spread out
over
lo6
years, the
maximum
temperature is
1000
K.
For this type
of
model
a short accretion time is
required to generate
high
temperatures. The Earth,
how
-
ever, is
unlikely to
grow
in radiative equilibrium. Higher
internal temperatures can
be achieved
if
the Earth
accu
-
mulated
partly
by
the continuing capture
of
planetesimal
swarms of meteoritic bodies. These
bodies hit the
Earth at
velocities considerably higher
than
free fall
and,
by
shock
waves, generate heat
at depth
in
the impacted
body.
Modern accretional calculations,
taking
into account
the energy partitioning during impact,
have
upper
-
mantle
temperatures in excess
of
the
melting
temperature
during
most
of
the accretion time (Figure
1-1).
If melting
gets too
extensive,
the melt moves toward the
surface,
and some
fraction reaches the surface
and
radiates
away
its
heat.
A
hot atmosphere, a thermal boundary
layer
and
the presence
of
chemically
buoyant
material at
the Earth's
surface,
how
-
ever, insulates
most
of
the interior,
and
cooling is
slow.
Ex
-
tensive cooling
of
the
upper mantle can
only
occur
if cold
surface material is subducted into the mantle. This requires
an
unstable surface layer, that
is,
a very cold,
thick
thermal
boundary layer
that
is
denser
than
the underlying mantle.
An
extensive accumulation
of
basalt
or olivine near the
Earth's
surface during accretion forms a
buoyant layer
that
resists subduction.
An
extensively
molten
upper mantle is
FIGURE
1
-
1
Schematic temperatures as a function
of
radius
at three stages
in
the accretion
of
a planet (heavy lines). Temperatures in the inte
-
rior
are
initially low because of the low energy
of
accretion.
The
solidi
and
liquidi
and
the melting zone in the upper mantle are
also shown. Upper
-
mantle melting
and
melt
-
solid separation
is
likely
during
most
of
the accretion process. Silicate melts, en
-
riched
in
incompatible elements, will be concentrated toward the
surface throughout accretion. Temperature estimates provided
by
D.
Stevenson (personal communication).
4
THE
TERRESTRIAL
PLANETS
therefore likely
during
accretion.
As
a thick
basalt
crust
cools, the
lower
portions
eventually
convert to dense
eclo-
gite,
and
at this point portions
of
the upper mantle can
be
rapidly cooled. A dense primitive atmosphere is also
an
ef
-
fective insulating agent
and
serves
to
keep the crust
and
upper mantle from cooling
and
crystallizing
rapidly.
All things
considered, it
is
likely that impact
melt
-
ing
and
gravitational separation combined
with
internal ra
-
dioactive heating
resulted
in
terrestrial planets that
were
being
differentiated while
they were
accreting. Extensive
upper
-
mantle
melting
gives
an
upward
concentration
of
melts
and
the incompatible
and
volatile elements
and
burial
of
dense refractory crystals
and
melts. Dense melts include
Fe,
FeS
and
FeO-MgO-rich
melts
that form
at
high pres
-
sure. The enrichment
of
volatiles
and
incompatible ele
-
ments
in
the
Earth's
atmosphere
and
crust and
the depletion
of
siderophile elements in the upper mantle point
toward
a
very effective
chemical separation
of
this type, as does the
presence
of
a ligh crust
and
a dense core. There is evidence
from both
the Earth
and
the
Moon
that these bodies
were
covered
by
deep
magma
oceans early
in
their history. The
lifetime
of
such an
ocean depends on the temperature
of
the
atmosphere, the thickness
of an
insulating crust
and
the rate
of
energy
delivery, from outside,
by
impacts and, from
in
-
side,
through
the thermal boundary layer at the
base
of
the
ocean.
Removal
of
crystals from a crystallizing
magma
ocean and
drainage
of
melt from a cooling crystal
mush
(also, technically a
magma)
are very
much
faster processes
than
cooling
and
crystallization times. Therefore,
an
ex
-
pected
result
of
early planetary differentiation is a stratified
composition.
The
following
is a plausible variant
of
the
inhomoge-
neous
accretion hypothesis. Planets accrete as the nebula
cools,
and
the accreting material
has
the composition
of the
solids that are in equilibrium
with
the
nebula
at that tem
-
perature
plus
the
more
refractory material
that
has con
-
densed
earlier
and
escaped accretion. The
mean
composi
-
tion
of
a planet therefore becomes less refractory
with
time
and with
radius. After dissipation
of
the nebula, the terres
-
trial planets continue
to slowly
accrete material that
has
condensed
in
their
vicinity
and
the
more
volatile material
that condensed farther
out
in
the solar system. It is
not
nec
-
essary that all
of
the refractories
and iron
be
accreted before
the silicates. There is
always
unaccreted material available
for interaction
with
the
gas. Iron, for example, is accreted
as
metal at
the early stages
but
reacts
with
the silicates to
form
the
ferromagnesian silicates that are accreted later.
Likewise, calcium, aluminum, uranium
and
thorium are
available
for incorporation into the later condensates,
but
they
are enriched in the early condensates.
FeS
will con
-
dense
and
accrete
at low
temperatures unless all the
Fe
metal has been removed
by
earlier processes.
The viability
of
the inhomogeneous accretion hypothe
-
sis depends
on
the relative time scales
of
nebula
cooling
and
accretion
in
the
early
stages
of
condensation.
If cooling is
slow relative
to
accretion rates,
then
the
iron
and
refracto
-
ries
will
form the initial nuclei
of
the planets. Alternatively,
cooling can
be rapid
if temperatures
in
the vicinity
of
the
Earth
do
not
drop far into the olivine
-
pyroxene stability
field
before dissipation
of
the nebula.
In
this case, the
ma
-
jority
of
the
mantle would be added
by
material perturbed
into Earth orbit from cooler parts
of
the nebula. The earliest
condensates also
have more
time for accretion
and
possibly
experience
more
viscous drag
in
the
dense, early nebula.
The presence
of
water
at the
Earth's
surface
and
the sidero
-
phile content
of
the
upper
mantle
indicate that the later
stages
of
accretion
did
not involve substantial amounts of
metallic iron (Lange
and
Ahrens,
1984).
Thus,
there are
several arguments supporting the
view
that the proto
-
Earth
was
refractory
and
became
more
volatile
-
rich with time.
The atmospheres
of
the terrestrial planets are apparently
secondary, formed
by
outgassing
of
the interior
and
devola-
tilization
of
late impacts. Primitive atmospheres are gener
-
ally thought
to
have been blown
away
either
by
a strong
solar
wind
or
by
giant impacts. Giant impacts in early Earth
history
may
have blasted
material into orbit
to
form the
Moon
(see Chapter
2)
and,
in later Earth history,
been re
-
sponsible for the various extinctions
that
punctuate the
paleontological record. Some meteorites found on Earth are
thought
to have been
derived from the surface
of
the
Moon
and
Mars
by
large
impacts
on
these bodies.
There is no particular reason to
believe
that there were
originally only four or
five
terrestrial planets
of
Mars
or
Moon
size or greater. The
sweeping
up of
multiple small
planets
by
the
remaining
objects is,
in
effect, a mechanism
of
rapid
accretion. The early history
of
the surviving terres
-
trial planets is therefore
violent
and
characterized
by
melt
-
ing and
remelting events.
METEORITES
Using
terrestrial samples,
we
cannot see
very
far
back
in
time or very deep into a
planet's
interior. Meteorites
offer
us
the opportunity to extend
both
of
these dimensions.
Some meteorites, the chondrites, are chemically
primitive,
having compositions
-
volatile
elements
excluded
-
very
similar to that
of
the sun. The volatile
-
rich carbonaceous
chondrites are samples
of
slightly altered, ancient planetesi
-
mal
material that condensed at moderate to
low
tempera
-
tures in the solar nebula. The nonchondritic meteorites are
differentiated
materials
of
nonsolar composition
that have
undergone chemical processing like that
which has affected
all known terrestrial
and
lunar rocks.
Meteorites are assigned to three
main
categories. Irons
(or siderites) consist primarily
of
metal; stones
(or
aerolites)
consist
of
silicates
with
little metal; stony irons (or sidero
-
lites)
contain abundant
metal and
silicates. These are further
subdivided in various classification schemes,
as
listed in
Table
1
-
1.
METEORITES
5
TABLE
1
-
1
Classification and Characteristics
of
Stony Meteorites and
Iron-rich
Meteorites
(Elements,
Weight Percent; Ratios Based
on
Atomic Percent)
Fe
Meteorite
class
Minerals
S
0
Fe
T
FeO
Fe+Mg
AIISi
CaISi
MgISi
Remarks
Stones (95
percent
of
meteorite
population)
Chondrites (86 percent)
Carbonaceous (5 percent)
C1
=
CI
e.g.
Ivuna; Orgueil
C2
=
CM
=
CII
e.g.
Mighei; Murchison
C3
=
CIII
e.g.
CO
=
Ornans
CV
=
Vigarano
Allende
Ordinary
(8
1 percent)
E
=
Enstatite
e.g.
Abee;
Khairpur
H
=
High iron
=
Bronzite
L
=
Low iron
=
"
Hypersthene
"
LL
=
Very
low
iron
=
Amphoterite
=
Soko-Banja
Achondrites
(9
percent)
Calcium
-
poor
Ureilites
=
Olivine-
pigeonite
Chassignites
=
Basaltic
Aubrites
=
Enstatite
*
Diogenites
=
"
Hypersthene,
"
"
Cumulates
"
Calcium
-
rich
*
Howardites and
*
Eucrites
=
Basaltic
e.g.
Juvinas,
Pasarnonte,
Stannern,
Moore
Co.
Angrites
=
Augite
'Nakhlites
=
Basaltic
'Shergottites
=
Basaltic
Stony Irons
(1
percent)
*
Mesosiderites
Pallasites
Irons
(4
percent)
Octahedrites
Hexahedrites
Nickel
-
rich Ataxites
Layer
-
silicates
(
"
clays
"
)
18
-
22
pct.
H20
6
-
16
pct.
H,O
Olivine,
refractory
minerals
<4
pct.
H,O
Enstatite, Fe
-
Ni
Olivine,
bronzite, Fe
-
Ni
Olivine,
bronzite, Fe
-
Ni
Olivine
Olivine,
0.5
pigeonite,
Fe
-
Ni
Enstatite
Orthopyroxene
0.4
Orthopyroxene,
0.27
pigeonite,
0.20
plagioclase
Augite
0.45
cpx,
01
0.06
cpx,
~1%
Orthopyroxene,
1.1
plagioclase,
Fe
-
Ni
Olivine,
Fe
-
Ni
0.19
Fe-Ni
0.02-
0.09
Fe-Ni
0.06
-
Fe-Ni
0.08
1.05-
N
O
1.06
chondrules
*The
eucritic association.
'The
SNC
association.
TABLE
1
-
2
Major
Minerals
of Calcium
-
Aluminum
-
Rich Inclusions
(CAI)
Volume
Condensation
Mineral
Percent
Temperature
*
Spinel
(MgAl,O,)
15
-
30
1513
-
1362
K
Melilite
(Ca,Al,SiO,)
0
-
85
1625
-
1450
Perovskite
(CaTiO,)
0
-
2
1647
-
1393
Anorthite
(CaA1,Si
,O,)
0
-
50
1362
Pyroxenes
0
-
60
1450
Grossman
(1972).
*Lower
temperature
is
temperature
at
which phase reacts with
nebular
gas
to
form
new
phase.
Carbonaceous Chondrites
Carbonaceous chondrites contain
unusually high
abun-
dances
of
volatile components
such
as
water
and
organic
compounds,
have
low
densities,
and
contain the heavier ele
-
ments
in
nearly
solar proportions.
They
also contain carbon
and
magnetite. These characteristics
show
that
they have
not been
strongly heated, compressed
or
altered since their
formation;
that
is,
they
have not been
buried deep inside
planetary objects.
The
C1
or
CI
meteorites are the
most
extreme
in
their
primordial characteristics
and
are
used
to supplement solar
values
in
the
estimation
of
cosmic composition. The other
categories
of
carbonaceous chondrites, CII (CM) and CIII
(CO
and
CV), are less volatile
-
rich.
Some carbonaceous chondrites contain
calcium-alu-
minum
-
rich inclusions (CAI),
which
appear to
be
high-
temperature condensates from the solar nebula. The miner
-
als
(Table
1
-
2) include anorthite
(CaAlSi,O,),
spinel,
diop-
side,
melilite, perovskite
(CaTiO,),
hibonite
(CaAl1201,)
and
the
Al
-
Ti
pyroxene, fassaite. These inclusions are
found
in
CV
and
CO chondrites,
most notably
(because
of
the total
volume
of
recovered material) the Allende
mete
-
orite. Theoretical calculations
show
that compounds rich
in
Ca,
A1
and
Ti,
including the
above minerals, are
among
the
first
to condense
in a cooling solar nebula. Highly refractory
elements
are
strongly
enriched in the CAI compared to
C1
meteorites, but they
occur in
C1,
or
cosmic, ratios.
Chondrites are
named
after the rounded fragments, or
chondrules, that
they
contain. Some
of
these chondrules ap
-
pear to
be
frozen drops
of
silicate liquid
and
others resemble
hailstones
in
their internal structure. Whatever their origin,
the presence
of
chondrules indicates the composite nature
of
meteorites
and
the
melting
or
remelting episodes that
characterized the history
of
at least some
of
their compo
-
nents.
C1
"
chondrites
"
are
fine
grained
and
do
not
contain
chondrules.
They
are chemically similar,
however,
to the
true chondrites (see
Table
1
-
3).
Ordinary Chondrites
As
the
name
suggests, ordinary chondrites are
more
abun
-
dant,
at
least in Earth
-
crossing orbits,
than
all other types
of
meteorites.
They
are chemically similar but differ
in their
contents
of
iron
and
other siderophiles,
and
in
the ratio
of
TABLE
1
-
3
Compositions
of Chondrites
(Weight
Percent)
Ordinary
Carbonaceous
Enstatite
H
IL
CI
CM
CO
CV
Mason (1962).