Yung Y.L., DeMore W.B. Photochemistry of Planetary Atmospheres
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86
Photochemistry
of
Planetary Atmospheres
Table
4.1
Molecules
identified
in the
interstellar
medium
Species Name
Two
ATOMS
A1F
A1CI
C
2
CH
CH+
CN
CO
CO+
CP
cs
CSi
HC1
H
2
KC1
NH
NO
NS
NaCl
OH
PN
SO
SO+
SiN
SiO
SiS
Aluminum
monofluoride
Aluminum
monochloride
diatomic carbon
methylidyne
methyb'dytie
ion
cyanogen
carbon monoxide
carbon
monoxide
ion
carbon phosphide
carbon
monosulfide
silicon carbide
hydrogen chloride
molecular hydrogen
potassium
chloride
nitrogen hydride
nitric oxide
nitrogen
sulfide
sodium
chloride
hydroxyl
phosphorus nitride
sulfur
monoxide
sulfoxide
ion
silicon nitride
silicon monoxide
silicon
sulfide
THREE
ATOMS
C
3
C
2
H
C
2
0
C
2
S
CH
2
HCN
HCO
HCO+
HCS+
HOC+
H
2
O
H
2
S
HNC
HNO
MgCN
MgNC
triatomic
carbon
ethynyl
radical
dicarbon
monoxide
dicarbon
sulfide
methylene
hydrogen cyanide
formyl
formyl
ion
thioformyl
ion
isoformyl
ion
water
hydrogen
sulfide
hydrogen
isocyanide
nitroxyl
magnesium
cyanide
magnesium isocyanide
Species
NH
2
N
2
H+
N
2
O
NaCN
DCS
SO
2
c-SiC
2
FOUR
ATOMS
c-C
3
H
1-C
3
H
C
3
N
C
3
0
C
3
S
C
2
H
2
CH
2
D+
HCCN
HCNH+
HNCO
HNCS
HOCO+
H
2
CO
H
2
CN
H
2
CS
H
3
0
+
NH
3
FIVE
ATOMS
C
5
C
4
H
C4Si
1-C
3
H
2
c-C
3
H
2
CH
2
CN
CH,
HC
3
N
HC
2
NC
HCOOH
H
2
CHN
H
2
C
2
O
H
2
NCN
HNC
3
Name
nitrogen
dihydride
protonated
nitrogen
nitrous oxide
sodium cyanide
carbonyl
sulfide
sulfur
dioxide
silicon
dicarbide
cyclopropenylidyne
propenylidyne
cyanoethynyl
tricarbon
monoxide
tricarbon
monoxide
acetylene
methylium
cyano-methylene
protonated hydrogen
cyanide
isocyanic
acid
isothiocyanic
acid
protonated carbon
dioxide
formaldehyde
methylene-aminylium
thioformaldehyde
hydronium
ion
ammonia
pentatomic
carbon
butadiynyl
silicon
tetracarbide
propenylidene
cyclopropenylidyne
cyanoacetylene
methane
propadienylidene
isocyanoacetylene
formic
acid
methanimine
ketene
cyanamide
cyanoacetylene
Species
SiRt
SIX
ATOMS
C
5
H
C
5
0
C
2
R,
CH
3
CN
CH
3
NC
CH
3
OH
CH
3
SH
HC
3
NH
+
HC
3
HO
HCOCH
2
HCONH
2
1-H
2
C,
SEVEN
ATOMS
C
6
H
CH
2
CHCN
CH
3
C
2
H
HC
5
N
HCOCH
3
NH
2
CH
3
EIGHT
ATOMS
CH
3
C
3
N
HCOOCH
3
NINE ATOMS
CH
3
C
4
H
CH
3
CH
2
CN
(CH
3
)
2
0
CH
3
CH
2
OH
HC
7
N
TEN
ATOMS
CH
3
C
5
N
(CH
3
)
2
CO
ELEVEN
ATOMS
HC
9
N
THIRTEEN ATOMS
HCnN
Name
silane
pentynylidyne
1
,2-cyclobutadiene
ethylene
methyl cyanide
methyl
isocyanide
methanol
methyl
mercaptan
protonated cyanoacetylene
propynal
vinoxy
formamide
butatrienylidene
hexatriynyl
vinyl
cyanide
methylacetylene
cyanodi
acetylene
acetaldehyde
methylamine
methylcyanoacetylene
methyl
formate
methyldiacetylene
ethyl
cyanide
dimethyl
ether
ethanol
cyanohexatriyne
methylcyanotetradiyne
acetone
cyano-octa-tetra-yne
cyano-deca-penta-yne
From Irvine,
W. M., and
Knacke,
R. F.
(1989).
conditions
normal
neutral
chemistry
would
be too
slow.
The
primary
driving
forces
of
disequilibrium
chemistry
are
ionization
by
cosmic
rays
and
photoionization
by
starlight.
A
detailed
discussion
of the
chemistry
of the
interstellar
medium
is
beyond
the
scope
of
this
chapter.
We
show
the
essential
aspects
of
this
chemistry
by a few
examples.
A
schematic
diagram
showing
the
formation
of
HaO,
CO,
and
CC>2
is
given
in figure
4.6.
The
bulk
of the gas in the
interstellar
medium
is
H
2
.
Cosmic
ray
or
photon
ionization
leads
to the
production
of
H2
+
,
which
is
removed
primarily
by
Origins
87
Figure
4.6 The
most important chemical pathways involving oxygen-bearing species
in
the
interstellar medium.
After
van
Dishoeck,
E.
R,
Black,
G.
A.,
and
Lunine,
J.
I.,
1993, "The Chemical Evolution
of
Protostellar
and
Protoplanetary
Matter,"
in
Levy
and
Lunine
(1993;
cited
in
section 4.1),
pp.
163-241.
As
shown
in figure
4.6,
H3
+
can now
react
with
O:
This
is
followed
by
The final
product
is
H
2
O,
synthesized from
HZ
and O.
Similar reaction pathways,
as
illustrated
in figure
4.6, produce
CO and
CO
2
.
The
synthesis
of
hydrocarbons
may be
carried
out
using similar
ion
chemistry.
The
precursor
species
C and
C
+
can
react with hydrogen
by
88
Photochemistry
of
Planetary Atmospheres
followed
by
The net
result
is the
formation
of
CH4.
Higher hydrocarbons containing
two
carbon
atoms
may be
formed
in the
following reactions:
These reactions ultimately produce acetylene
and
ethylene.
Simple nitrogen compounds
may be
produced
from
N
+
as
follows:
followed
by
where
n = 1, 2, or 3.
Ammonia
and HCN can
then
be
formed
by
We
should point
out the
great
efficiency
of
organic synthesis
in the
interstellar
medium. First,
the
chemistry
is
dispersed over
a
large volume
of
space.
The
size
of a
giant
molecular cloud
is of the
order
of 100 pc (1 pc = 3 x
10
18
cm = 2 x
10
s
AU).
This increases
the
collection
efficiency
of
cosmic rays
and
ultraviolet photons. Second,
the
chemical synthesis
is
primarily
via
ion-molecule reactions that have little
or no
activation
energy,
and
these
are
extremely
fast
even
at low
temperatures. Third,
the
ambient
radiation
field is not
intense enough
to
destroy
the
synthesized products. These
conditions
are
difficult
to
duplicate
in
planetary atmospheres.
The
only
atmosphere
in
which
organic synthesis approaches
the
richness
and
complexity
of the
interstellar
medium
is
that
of
Titan
and
presumably
the
ancient atmospheres
of
Earth
and
Mars.
Origins
89
Table
4.2
Deuterium
fractionation
in
interstellar
clouds
Molecule Observed ratio
DCN/HCN
DCO+/HCO+
N
2
D
+
/N
2
H
+
DNC/HNC
NH
2
D/NH
3
HDCO/H
2
CO
C
2
D/C
2
H
DC
3
N/HC
3
N
DC
5
N/HC
5
N
C
3
HD/C
3
H
2
0.002-0.02
0.004-0.02
~0.01
0.01-0.04
0.003-0.14
-0.01
0.01-0.05
~0.02
0.02
0.03-0.15
From Irvine,
W.
M., and
Knacke,
R. F.
(1989).
4.4.2
Isotopic
Fractionation
A
distinctive signature
of
interstellar chemistry
is a
high
degree
of
isotopic
fractiona-
tion.
Table
4.2
summarizes
the
deuterium fractionation
for
hydrogen-bearing species
observed
in
molecular clouds.
These
values
are
orders
of
magnitude higher than
the
cosmic ratio
D/H = 1.6 x
10~
5
.
It is a
triumph
of
ion-molecule reaction studies
to
provide
a
satisfactory account
of
these large fractionations.
We
consider
two
examples.
The
isotope exchange reactions
are
exothermic
by E\
=
0.46
and
£
2
=
0.74
kcal/mole,
respectively.
The
amount
of
fractionation
may be
estimated using
the
formulas developed
in
section 3.2:
where
/ =
[D]/[H]
is the
cosmic abundance. Equations (4.23)
and
(4.24)
are
good
ap-
proximations because
H
2
and HD are the
major
reservoirs
for
hydrogen
and
deuterium,
respectively,
in the
interstellar medium,
and the
exchange reactions (4.21)
and
(4.22)
are
sufficiently
fast
that isotopic equilibration
is
rapidly established.
The
fractionation
may
be
expressed
as
where
the
fractionation
factors
are
given
by
90
Photochemistry
of
Planetary Atmospheres
Note
that
R\
and
/?
2
would
be
unity
if the
composition were given
by
cosmic abun-
dances.
At low
temperatures
(10-50
K), the
fractionation factors
can be
very large.
For
example,
at 30 K, the
values
of
R\
and
/?
2
are 1.4 x
10
3
and 1.5 x
10
5
,
respectively.
Reactions
of the
type
where
X is a
neutral molecule such
as OH,
CHj,
or
NH
2
favor
the
right side,
and E
is
positive.
XD
+
is
usually more tightly bound
due to a
reduced zero point energy.
We
can now
appreciate
the
beauty
of a
fractionation
process
such
as
(4.24).
The
fractionation
factor
is
very large
due to the
large energy difference
E
compared with
RT at low
temperatures.
At the
same time reactions such
as
(4.22)
are
ion-exchange
re-
actions that
can
readily
proceed
without activation energy,
as
discussed
in
section
3.4.4.
Thus,
the
observed isotopic fractionations
in
table
4.2 are
simply explained.
4.4.3 Grain
Chemistry
and
Radiative
Association
In
addition
to ion
chemistry, catalysis
on the
surface
of
grains
is
believed
to
play
a
major
role
in
interstellar chemistry.
The
most important reaction that
is
believed
to be
catalyzed
on the
surface
is the
recombination
of
hydrogen atoms:
H
+
grain
-»
H(ads) (4.30)
H
+
H(ads)
-»
H
2
(4.31)
Note
that
the first
atom
is
adsorbed
on the
grain
and
waits
for the
second atom
for the
reaction
to be
completed.
The
efficiency
of the
recombination depends
on the
rate
of
collision between
H
atoms
and
grains
and on the
accommodation
coefficient
of H on
the
surface.
The
latter
is
highly uncertain
and may
depend
on the
number
of
available
sites
for
chemi-adsorption.
Organic synthesis
may
also
be
carried
out on the
surface
of
grains.
Organic synthesis
is
initiated
by
reactions such
as
(4.1)
involving
H
2
.
The
sequence
of
reactions
(4.1)-(4.5)
does
not
result
in a net
increase
of the
number
of
chemical
bonds,
but an
exchange
of
chemical bonds. Therefore,
it is
extremely important
to first
form
the
H—
H
bond.
One may ask how
H
2
was
formed
in the
primordial universe
without
heavy elements
and
dust grains.
The
following
are
alternative
processes:
These
processes
involving radiative recombination
are
very
inefficient
compared
with
grain
catalysis,
but in the
absence
of
grains they
may
have been important during
the
early
development
of the
universe.
Origins
91
Figure
4.7
Schematic
illustration
of the
physical
and
chemical processes
in the
solar nebula.
CMC
=
giant
molecular
nebula.
After
van
Dishoeck,
E. F.,
Black,
G.
A.,
and
Lunine,
J.
I.,
1993,
"The Chemical Evolution
of
Protostellar
and
Protoplanetary
Matter,"
in
Levy
and
Lunine
(1993; cited
in
section 4.1),
pp.
163-241.
4.5 The
Solar Nebula
The
nebula hypothesis
for a
common origin
of the sun and the
planets
can be
traced
back
to
Kant
and
Laplace
in the
eighteenth
century.
The
formation
of the
solar
system
from
a
molecular cloud
was
most probably triggered
by a
neighboring supernova.
The
shock from this explosion initiated
an
instability that
led to the
collapse
of the
molecular cloud
and the
formation
of a
slowly rotating
disklike
object known
as the
solar nebula.
A
schematic diagram illustrating
the
physical
and
chemical
processes
in
the
solar nebula
is
given
in
figure
4.7.
The
physical extent
of the
nebula
is
about
100 AU,
roughly
the
size
of the
present
solar
system.
The
total amount
of
material
is
uncertain. Estimates range from barely
1
solar
mass
to
2—4
solar
masses.
The
bulk
of
the
mass
is
concentrated
at the
center,
the
protosun.
The
"flared"
geometry
of the
nebula
is due to the
weakening
of the
component
of
gravity that
is
perpendicular
to
the
ecliptic plane
in the
outer part
of the
nebula. Both
gas and
solid
are
present
in
the
nebula.
The
inner nebula (within
5 AU
from
the
center)
is
relatively warm,
and
the
solid material
is
mostly refractory.
At
larger distances from
the
center
the
nebula
is
cooler,
and
ices
can
form.
As
shown
in figure
4.7,
the
nebula
is
hardly
a
"static"
system.
Without
going
into
the
details,
we
mention
three
dynamical
processes
that
might
have been important
for
chemistry
in the
nebula:
(a)
large-scale convection
that
mixes material between
the
inner
and
outer nebula,
(b)
viscous shear motion
that transfers most
of the
mass
of the
nebula
to the
protosun
but
leaves
most
of the
92
Photochemistry
of
Planetary Atmospheres
Figure
4.8
Equilibrium
condensation
sequence
for a
system
of
solar
elemental
composition.
After
Lewis,
J. S., and
Prinn,
R. G.,
1984, Planets
and
Their
Atmospheres (New York:
Academic
Press).
angular
momentum
to
Jupiter,
and (c)
supersonic shock
as the
free-falling
material
enters
the
"boundary"
of the
nebula.
The
effect
of (a) and (c) on
chemistry
is
obvious.
Process
(b) may
require hydromagnetic coupling, which
in
turn requires
a
high state
of
ionization
of the
nebula material. Under these unproven assumptions,
ion
chemistry
may
be
important.
The
physical
and
chemical conditions
of the
solar nebula
are no
longer observable
today.
Models
of the
solar nebula
can be
constructed using information based
on the
composition
of the
planets, meteorites,
and
comets (see sections
4.6-4.8).
An
example
of
this class
of
model
is
given
in
figure
4.8.
The
pressure
and
temperature conditions
in
the
solar nebula
are
approximately given
by the
adiabat.
The
positions that correspond
to the
locations
for the
formation
of the
planets
are
indicated
in
figure
4.8.
For
example,
at
1 AU
where Earth
is
formed,
the
temperature
is
about
600 K and the
pressure
is
10~
3
bar.
In the
region
close
to the
sun,
the
pressure
and
temperature
are
higher
and
only
the
most refractory oxides
of
elements such
as
Al,
Ca, Mg, and the
rare earths
condense, followed
by Fe and Ni. The
minerals
of the
alkali metals condense
next
in
the
sequence.
FeS
condenses
at
about
700 K, at a
location between
Venus
and the
Earth.
Beyond
1 AU
hydrated minerals such
as
tremolite
and
serpentine
will
condense.
Finally,
as we
move
out to the
outer solar system, volatile material such
as
HaO,
NHs,
and
CH
4
condense
as ice and
hydrates,
NHs
• H2O and
CH4
•
IHzO.
A
major success
of
this
simple model
is its
ability
to
account
for the
bulk
density
and
composition
of
the
planetary bodies
in the
solar system.
Origins
93
What
type
of
chemistry prevails
in the
solar nebula depends
on the
energy source
available. Prinn
and
Fegley (1989) define
a
ratio that measures
the
usable energy
flux
available
for
driving
chemical
reactions
relative
to the
thermal
flux:
where
(j>
(thermal)
=
aT
4
is the
thermal
flux at
temperature
T.
At 1 AU the
estimated
ratios
for
various chemical processes
are as
follows:
These
estimates
are
made under
the
assumptions that
the
activation energies
for
chem-
ical reactions
are
about
10
kcal/mole,
the
ambient temperature
is
about
600
K,
and
the
nebula
is
largely opaque
to
visible
and
ultraviolet radiation.
The
source
of
energy
due to
lightning
is
highly uncertain. Hence,
the
only proven energy source
is
ther-
mal
energy, derived
from
the
gravitational accretion energy
of the
nebula itself. Other
forms
of
energy
may
have existed
but
need
to be
verified
by
observational evidence.
We
should note that
in
(4.38)
/?(sunlight)
=
0 not
because
of the
lack
of
ultraviolet
radiation
from
the sun but
because
of the
enormous opacity
due to
absorption
by
dust.
As
discussed
in
section 4.7, there
may be an
indirect source
of
ultraviolet radiation
arising
from
a hot
corona above
the
nebula.
An
example
of a
thermochemical
calculation
for the
partitioning
of the
biogeo-
chemical elements
C,
H,
O, and N is
shown
in figure
4.9.
The
calculations
are
based
on
equilibrium chemistry along
the P, T
adiabat
in the
solar nebula.
H2, the
dominant
species,
has a
mixing ratio close
to
unity
and is not
plotted
on
this
figure
4.9.
The
most important part
of the
equilibrium chemistry concerns
the
following partitioning:
The
equilibrium chemistry tends
to
favor
CH
4
and NH3 in the
outer
solar
nebula
and
CO and
CO
2
in the
inner solar nebula. Independent information
is
required
to
estimate
the
temperature
at
which these reactions
are
"quenched."
As
shown
in figure
4.9,
the
quenching
temperatures
for
N
2
->•
NH
3
,
CO
-»•
CH
4
,
and CO
->
CO
2
are
1600, 1470,
and 830 K,
respectively.
Thus,
for
example,
the
predicted conversion
of CO to
CHLt
at
the
orbit
of
Jupiter
or
beyond
(T < 300 K) can
never
be
realized using homogeneous
equilibrium
chemistry
alone,
because
the
time required would
be too
long compared
with
the age of the
nebula.
94
Photochemistry
of
Planetary
Atmospheres
Figure
4.9
Calculated equilibrium abundances
of
important
H, C, N,
and
O
gases
along
the
solar nebula adiabat.
After
Prinn,
R. G., and
Fegley,
B.,
Jr., 1989,
"Solar
Nebula Chemistry: Origin
of
Planetary,
Satellite
and
Cometary
Volatiles,"
in
Atreya
et
al.
(1989;
cited
in
sec-
tion
4.1),
pp.
78-136.
There
are
obviously other
processes
that would
modify
the
equilibrium chemistry
discussed
in the
preceding section. Catalysis
by
transition metals
on the
surface
of
grains
may
accelerate
the
rate
of
these reactions.
For
example,
Fe-catalyzed
conversion
of
N2 to
NHs
and of CO to
CUt
may be
important
at
temperatures
as low as 500 K.
Additional
processes
that would
modify
the
simple chemistry outlined here include
mixing
in the
nebula,
and
similar
but
distinctive chemistry that
can
take place
in the
planetary subnebulae.
4.6
Meteorites
Meteorites
are
rocks that
are of
extraterrestrial origin. They
are
compositionally
similar
to the sun (in the
heavy elements). Radiometric dating places their
age at
4.55 Gyr,
about
the
same
age as the
Moon
and
Earth. There
is
strong
evidence
that meteorites
are
samples
of
solid material
left
over
from
the
solar nebula.
The
classification
of
meteorites
and
their
fall
frequencies
are
summarized
in the
appendix 4.2. Most
of the
meteorites
are
chondrites, primitive clastic rocks that formed
in the
early history
of
the
solar nebula. Although most
meteorites
consist
of
unprocessed material, that
is
Origins
95
not
true
of
iron meteorites.
These
must have been formed
in
planetesimals that
had
melted
and
were subsequently broke
up. The
Shergottites,
Nakhlites,
and
Chassignites
are
most
likely
from
Mars.
4.6.1
Meteorite
Chronology
Meteorites contain radioactive nuclei that
are
extremely
useful
for
dating important
events
that took place
in the
solar nebula.
The
principle
of the
classical method
of
dating
using radionuclides
may be
illustrated
as
follows:
The
time-dependent concen-
tration
of a
radioactive element
N(f)
is
where N(0)
is the
initial
concentration
and
AT
1
is the
mean life.
In
deriving this result
we
have assumed that there
is no new
source
of the
radioactive nuclei
after
t = 0.
Consider
the
example
of the two
most abundant uranium
isotopes.
The
mean lifetimes
of
235
U
and
238
U
are 1 and 6.5
Gyr, respectively.
The
initial ratio
N(
235
U)/N(
238
U)
is
1/3;
the
current ratio
is
1/137.8. From this
and
(4.44),
we may
infer
the
sample
is 4.6
Gyr
old.
The
ages
of the
oldest terrestrial
and
lunar rocks
are
determined
by
this
or
variants
of
this method.
There
is a
class
of
short-lived nuclei that
are
extremely important
for
understanding
the
chronology
of the
early solar system.
These
nuclei have mean lives between
0.1
and
100
Myr, long-lived enough
for
them
to
survive
from
nucleosynthesis
(in
stars)
to
incorporation
in the
meteorites
but not
long-lived enough
to
survive
to the
present.
For
this reason they
are
known
as
extinct nuclei. Since these nuclei
are by
definition
extinct,
their
abundances
can be
inferred only
from
their decay products.
An
example
is
26
A1,
which
has a
mean
life
of
1.1
Myr, decaying
to
26
Mg.
Figure 4.10 shows
the
variation
of the
ratio
26
Mg/
24
Mg
in an
inclusion
in the
Allende meteorite.
The
range
of
variation
is of the
order
of
10%,
and
there
is
good correlation
with
the
ratio
27
Al/
24
Mg.
Since
26
A1
decays
to
26
Mg,
it is a
good hypothesis that
the
excess
26
Mg
is
all
derived from
the now
extinct
26
A1.
The
amount needed
is
26
A1/
27
A1
= 5 x
10~
5
.
Given
the
short
life
of
26
A1,
we
conclude that
the
meteorite must have formed within
a
few
lifetimes
of
26
A1,
or a few
million years. Since this pioneering work,
the
existence
of
six
extinct nuclei
in
meteorites
has
been confirmed,
as
summarized
in
table 4.3.
The
results
from
the
other nuclei generally
confirm
the
conclusion
based
on
26
A1.
In
addition,
the
large number
of
nuclei allows
us to
infer
their stellar
sources.
As
shown
in
table 4.3,
for
example,
the
nuclei
26
A1,
^Fe,
107
Pd,
135
Cs,
and
182
Hf
are
derived
from
a
low-mass
AGB
star
2 Myr
before
the
solar system (BSS).
At 25 Myr
BSS,
an
OB
association
is
responsible
for
53
Mn.
4.6.2 Environment
of
Formation
The
meteorites
are
good probes
of the
denser
and
warmer parts
of the
solar
neb-
ula.
Based
on the
fractionation
of
refractory elements,
the
maximum temperature
of
the
location where chondrites were formed must
be
about
2000
K.
Calcium-
and
aluminium-rich
inclusions
(CAIs)
were most
likely
formed
at
temperatures between
1800
and
2000
K. In
fact,
the
minerals
in
CAIs were
found
to
have been formed