Yung Y.L., DeMore W.B. Photochemistry of Planetary Atmospheres
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Figure
7.21
Loss
rates over
the
history
of the
Martian
atmosphere,
(a)
Total loss
of
H2O.
(b)
Loss
rate
of O by
electronic recombination
(R96).
(c)
Loss
rate
of O by
solar wind-induced
sputtering,
(d)
Loss
rate
of
CC>2.
(e) and (f)
correspond
to
loss rates
(c) and (d)
as
computed
by the
model
of
Luhmann,
J.
G.,
Johnson,
R.
E.,
and
Zhang,
M
H.
G.,
1992, "Evolutionary Impact
of
Sputtering
of the
Martian
Atmosphere
by
O
+
Pickup
Ions."
Geophys.
Res.,
19,
2151.
From
Kass,
D.
M.,
and
Yung,
Y. L.,
1995,
"Loss
of
Atmosphere
from
Mars
Due to
Solar Wind Induced Sputtering."
Science,
268,
697.
276
Mars
277
Table
7.8 Net
escape
fluxes
from
Mars
due to
solar wind-induced
sputtering
for
three epochs
6
EUV
(1
Gyr)
3 EUV (2
Gyr)
1
EUV
(4.5 Gyr)
Sputtered
O
Exospheric
O
Pickup
O
+
Sputtered
CO
2
Escaped
H2O
3.1
x
10
28
1
x
10
27
3 x
10
27
7.8
x 10"
1.9
x
10
28
1.8
x
10
27
5 x
10
26
4 x
I0
26
6.4 x
10
26
1.4
x
10
27
4.7
x
I0
24
8 x
I0
25
6 x
10
24
2.4
x
10
24
8.6 x
10
25
From
Kass,
D.
M.,
and
Yung,
Y. L.
(1995).
species,
the
sputtering
by
solar wind-accelerated ionized exospheric
O
+
is
more
im-
portant.
The
loss rates over
the
history
of
Mars
are
summarized
in
figure
7.21
and
table
7.8.
As
much
as 3 bar of
CC>2
and 80
m
of
water could have been lost
from
Mars,
with
most
of the
loss having occurred
in the
remote past.
The
escape
rates
are
higher
in the
past
due to the
greater solar
wind
flux. The
uncertainties
in
these calcu-
lations
are
illustrated
by the
difference
in the
loss rates predicted
by
different
models
(e.g., between
figure
7.21,
c and e).
7.4.2 Isotopic
Signature
Table
7.3
shows that
the
D/H
ratio
of
water vapor
on
Mars
is
enriched
by a
factor
of 5
relative
to
terrestrial water.
The
I5
N/
14
N
ratio
of N2 on
Mars
is
1.62 times
the
terrestrial
value. (The isotopic composition
of
Earth's atmosphere
is
given
in
table 9.3.)
Note that
in
both
cases
the
heavier isotope
is
enriched relative
to the
lighter isotope.
The
other isotopic ratios,
17
O/
16
O,
18
O/
16
O,
and
13
C/
I2
C,
are not
fractionated relative
to
their terrestrial counterparts.
We
adopt
the
simplifying
assumptions that
the
material
outgassed from
the
interior
of
Mars
has the
same
isotopic composition
as
Earth
and
that
the
observed isotopic
fractionation
is
caused
by
difference
in the
escape
efficiencies
of the
isotopes.
As an
illustration
of the
proposed mechanism
of
fractionation,
let us
consider
the
escape
of
hydrogen
and
deuterium
from
the
planet.
Let H and D be the
total
atmospheric
abundances
of
hydrogen
and
deuterium
in the
atmosphere.
The
equations
that
govern
the
time evolution
of H and D are
where
we
have assumed that
the
rates
of
escape
of H and D are
proportional
to
their
abundances,
and / is the
fractionation
factor describing
the
efficiency
of
escape
of D
compared
with
H. For / = 1 the
escape
process
is not
fractionating
and
D/H
ratio
would
remain constant
with
time.
For / = 0
only hydrogen
can
escape
and the
D/H
ratio
would increase
at the
maximum possible rate.
If we
assume
that
the
coefficients
A
and / in
equations
(7.57a)
and
(7.57b)
are
constant,
the
solution
for the
evolution
of
the
D/H
ratio
is
278
Photochemistry
of
Planetary
Atmospheres
Figure
7.22
Effusion
velocities
from
the
exosphere
of
Mars
as
a
function
of
exospheric
temperature.
The
crosses
indicate
the
approximate
effusion
velocities
in the
present
epoch. From
Nair,
H.,
Allen,
M.,
Anbar,
A. D.,
Yung,
Y. L., and
Clancy,
R. T.,
1994,
"A
Photochemical
Model
of the
Martian
Atmosphere."
Icarus
111,
124.
where
t = 0 is the
initial
time,
and
H(0)
and
H(?)
denote
the
total hydrogen reservoir
at
t = 0 and r,
respectively. Since hydrogen
is
lost from
the
planet
we
have H(0)
>
H(0-
Equation
(7.58)
states that there
is a
corresponding enrichment
in the
D/H
ratio.
This result
is
known
as
Rayleigh distillation.
The
fractionation factor
/ in
(7.57b)
is
affected
by a
number
of
physical
and
chemical
processes
in the
atmosphere.
We
briefly
discuss
the
most important
of
these
processes.
First,
the
effusion
velocities given
by the
Jeans
escape
formula
are ex-
tremely
sensitive
to
mass difference,
as
shown
in figure
7.22. Second,
all the
atomic
and
molecular
species
are
diffusively
separated above
the
homopause according
to
their
own
scale heights. This means
the
lighter
species
will
be
enriched
at the
exobase
relative
to the
bulk atmosphere below.
The
fractionation factor
/ for
deuterium
has
been estimated from
a
detailed model
of
hydrogen
escape.
For
Mars
in the
current
epoch
/ is
0.32 i.e.; that
is,
deuterium
is
escaping
at 32% of the
efficiency
of hy-
drogen. Assuming
a
uniform rate
of
escape,
the
observed
D/H
enrichment implies
that
the
initial
reservoir
of
hydrogen
is
equal
to 3.6
m
of
water, most
of
which
has
escaped, leaving
a
residue
of 0.2 m
water
on
Mars today. This represents
a
very con-
servative
estimate
of
water (both
for the
primordial
and for the
remaining reservoir).
The
reservoir sizes would obviously
be
bigger
if
escape
rates have been faster
in the
past. Also,
this
analysis completely ignores
loss
of
volatiles
by
atmospheric cratering
and
sequestering
in
underground reservoirs, because
the
latter
processes
are not
likely
to
leave
an
isotopic signature.
Mars
279
Figure 7.23 Fsotopic fractionation
for
I5
N
as a
function
of
time.
The
dashed
line
is for an
initial
atmosphere
of 2 bar of
CC>2
that
is
lost exponentially
with
a
time constant
of 700
Myr.
From
Fox,
J. L.,
1993,
"Production
and
Escape
of
Nitrogen from
Mars."
J.
Geophys.
Res.
98,
3297.
A
similar analysis
may be
applied
to the
oxygen isotopes.
All the
above results
apply
if we
replace
D by the
heavy isotope
18
O
and H by the
light isotope
16
O.
(There
is
no
additional information
from
analysis
of the
pair
I7
O
and
16
O.)
The
fractionation
factor
/
equals 0.75
in
this
case.
The
lack
of an
observed isotopic fractionation implies
that
there
is a
large oxygen reservoir
on
Mars,
of the
order
of 1
bar, that
is
capable
of
diluting
the
isotopic enrichment caused
by
oxygen
escape.
We
cannot distinguish
whether
this
reservoir
is CO2 or
H2O,
but
from
the
previous analysis,
the
oxygen
reservoir
is
probably
not
H
2
O.
This leaves
CO
2
as the
most
likely
candidate.
A
similar analysis
may be
applied
to
nitrogen:
15
N
and
14
N.
The
fractionation
factor
in
this
case
is
0.82.
The
results
for the
evolution
of
isotopic fractionation
of
nitrogen
are
given
in figure
7.23.
The
total enrichment over
the age of the
solar system
(1.4
x
10
17
s) is
about 2.6,
a
value
significantly
higher
than
the
Viking observation
of
1.62.
The
reason
is
that
as we go
back
in
time
N2
became
the
dominant
gas in the
atmosphere. This implies large
escape
rates, hence
a
greater enrichment factor.
One
simple
way to
resolve
this discrepancy
is to
postulate that
the
Martian atmosphere
was
denser
in the
past such that
N
2
was
"protected"
from escaping.
The
dashed
line
in
figure
7.23 represents
the
result
of a
calculation assuming
that
there
was an
initial
COa
atmosphere
of 2
bars
of CO2
that
is
lost exponentially
with
a
time constant
of
700
Myr.
The
predicted isotopic enrichment
agrees
with
the
Viking value.
7.5
Terraforming
Mars
Our
interest
in
Mars
is
more than
scientific.
Of all the
planets
in the
solar system,
Mars
has a
physical
and
chemical environment
that
most closely resembles that
of
the
prebiotic Earth
and
presents
the
best opportunity
for
future human habitation.
280
Photochemistry
of
Planetary
Atmospheres
Although
the
present average surface temperature
of
Mars,
—
60°C,
is
much colder
than
that
of
Earth
(15°C),
the
fluvial
features
on the
surface suggest that Mars
was at
one
time much warmer. Perhaps
if we can
understand
the
causes
for the
ancient warm
climate,
we can
engineer
the
return
of the
Martian atmosphere
to
that state.
The
major problem
in
terraforming
Mars
is to
create
an
atmosphere that
is
dense
enough
to
raise
the
mean surface temperature
to
around
0°C.
To
achieve
this
requires
about
1 bar of
COa
in the
atmosphere. From
the
isotopic composition analysis given
in
section 7.4,
we
know that
this
amount
of CO2
probably exists
in the
surface reservoirs
on
Mars today.
One
efficient
way to
raise
the
atmospheric temperature
is to
inject
into
the
atmosphere molecules that have strong infrared absorption
bands,
such
as the
chlorofluorocarbons
(CFCs).
According
to
recent estimates,
the
amount that
is
needed
in
the
atmosphere
is of the
order
of
10~
6
.
Assuming
a CFC
lifetime
of
about
100
years,
the
rate
of
industrial production
to
maintain this amount
of
CFCs
in
steady
state
is
10
8
ton
yr~',
which should
be
compared
with
the
current global production
rate
of
10
6
ton
yr~'.
The CFC
production must
be
carried
out in
situ, since planetary
transportation
of
materials
at
this rate
is
impractical.
The
assumption
of a 100 yr
lifetime
for CFC on
Mars
assumes that
Mars
is
protected
by an
ozone layer. Since
the
CFCs
are
highly
detrimental
to
Os,
an
alternative ultraviolet screen must
be
provided.
The
simplest
one may be a
layer
of
polysulfur,
like that
in the
atmosphere
of
Venus.
The
source material
that
needs
to be
injected into
the
atmosphere
is
carbonyl
sulfide
(COS).
We
estimate that
an
annual production
of
10
7
tons would
be
sufficient
to
create
an
effective ultraviolet screen.
The
above estimates
are
somewhat pessimistic
because
we are
thinking
in
terms
of
current technology
and
currently available materials.
Ultimately,
the
long-term global environment
of
Mars
must
be
sustained with min-
imal
human input. Once
the
planet
is
modified
to the
point capable
of
sustaining
life
(we
assume there
is no
life
on
Mars
now),
we
could
use
Mars
to
test
the
Gaia
hy-
pothesis, which states
that
the
biosphere
is a
living
entity
and is
capable
of
modifying
the
environment
for its own
welfare.
Of
course,
we do not
know what
the
threshold
is
for a
Gaia
runoff
of
biological modification
of a
planet. Only
an
experiment will
tell.
If
successful, Gaia
may
then
be the
path
for the
future
colonization
of the
solar
system,
and
beyond.
But
the
prospects
of
this happening soon
are not
good.
Indeed,
the
Viking land-
ing
on
Mars
may be
compared
to the
historical discovery
of
North America
by the
Vikings
some
five
centuries before Columbus.
The
political, social,
and
technological
developments
of
Europe
at
that time were
not
ready
for the
colonization
of
North
America. Indeed,
the
word
"colonialism"
had not yet
been invented.
There
was a
long
wait—so
long that Columbus
had to
rediscover what
the
Vikings
did
many years
earlier.
We
hope that
the
contents
of
this section
can
serve
as an
inspiration
to a
future
Columbus
of the
solar system.
7.6
Unsolved
Problems
We
list
a
number
of
outstanding unresolved problems related
to the
atmosphere
of
Mars
and its
comparison
with
Earth's atmosphere.
1.
Is
heterogeneous chemistry important
in the
"dustiest"
atmosphere
in the
solar
system?
Mars
281
2. Are the
properties
of the
high-altitude
water
ice
clouds
similar
to
those
of the
polar
stratospheric
clouds
on
Earth?
3. Do the
high-altitude
water
ice
clouds play
a
role
in the
NO
X
chemistry?
4. Can the
HO,
chemistry
on
Mars
and in the
terrestrial mesosphere
be
"fixed"
by
adjusting
of the
values
of a
similar
set of
reactions such
as
R42:
CO+OH
->
CC>2+
H
and
R40:
OH +
HO
2
->
H
2
O
+
O
2
?
The odd
nitrogen model suggests that
the
principal sink
for
odd
nitrogen, R61:
N
+
NO
->
N
2
+ O, may be
slower
at
lower
temperatures
(150
K).
What
are
these
implications
for
laboratory chemical kinetics?
5. How
much
volatile
material
has
escaped
from
Mars?
By
atmospheric cratering?
By
thermal
and
nonthermal
escape?
By
solar wind-induced sputtering?
6. How
much
volatile
material
is
still
sequestered
on the
surface
and the
polar caps
of
Mars?
7. Is
Mars
still
degassing reducing
gases
from
the
interior?
8.
Mars
has
undergone dramatic
climatic
variations,
as can be
seen from seasonal
patterns
of
dust storms
and
billion-year-old river beds. Were
the
changes driven
by
variations
of the
solar constant
or
variations
of the
orbital elements?
9. Is the
current arid
and
dusty atmosphere
of
Mars
a
good
analog
of the
terrestrial
atmosphere
during
the
last glacial
maximum,
about
18,000
yr
ago?
10. Is
there
life
on
Mars today?
Was
there ever
life
on
Mars?
11.
Is it
possible (and ethical)
to
carry
out
terraforming
of
Mars
and
engineer
the
planet
for
human
habitation?
12.
Are the
escape rates
of
oxygen
and
hydrogen given
by the
stoichiometric
ratio
of
H
2
O
today?
Is
this
relation true
only
when averaged over
a
Milankovitch
climatic
cycle
of
Mars?
8
Venus
8.1
Introduction
Venus
has
been visited
by a
number
of
spacecraft.
Those
most important
in
advancing
our
understanding
of the
atmosphere include Mariner
10
(1974), Pioneer Venus
(1978-
1992),
and the
series
of
Venera probes
by the
former USSR
(1982-1986).
The
Pioneer
Venus
orbiter conducted more
than
a
decade
of
monitoring
of the
upper atmosphere
of
Venus.
The
spacecraft data
are
supplemented
by
telescopic
and
satellite
(in
Earth
orbit)
observations.
The
astronomical data
for
Venus
are
summarized
in
table 8.1. Venus
is a
close
but
slightly
smaller sibling
of
Earth.
The
radius
is
6051
km, as
compared
to
Earth's
6371
km. The
mass
is
4.87
x
10
24
kg, a
little
less than Earth's 5.98
x
10
24
kg. The
gravity
of
Venus
is
8.87
m
s~
2
,
as
compared
to
Earth's 9.82
m
s""
2
.
The
dynamical
and
orbital
parameters
of
Venus
are
very
different
from
those
of
Earth.
The
rotation
of
Venus
is
retrograde,
with
a
period
of 225
days.
The
planet
has
little
obliquity,
and
its
orbit
is
close
to
being circular. Thus there
is
little seasonal variation
in
insolation
over
a
Cytherian year. Perhaps
the
greatest surprise about Venus
is its
dense, dry,
and
hot
atmosphere
of 92 bar
(see
figure
8.1).
This
is all the
more surprising because
the
planet
is
completely covered
by
thick
clouds. Figure
8.2
shows
the
altitude profiles
of
three
modes
of
cloud particles
in the
middle atmosphere.
The
high albedo
of
Venus
implies
that
the
planet receives less energy from
the sun
than
Earth despite
its
closer
proximity
to the
sun. However,
the
surface
of
Venus
is
hot,
with
a
temperature
of 733
K,
attributed
to the
greenhouse
effect.
282
Table
8.1
Astronomical
and
atmospheric
data
for
Venus
Radius
(km)
a
Mass
(kg)
b
Mean density
(g
cm~
3
)
b
Gravity
(surface, equator)
(m
s~-)
c
Escape
velocity
(km
s~')
c
Semimajor
axis
(AU)
Eccentricity
of
orbit
c
Obliquity
(deg)
b
Sidereal
period
of
revolution (Earth
days)
c
Mean orbital velocity
(km
s~')
c
Sidereal period
of
rotation
(Earth
days,
retrograde)*
1
Annual
variation
in
solar
insolation
d
Almospheric pressure
at 70 km
altitude
(above cloud tops)
(mbar)
e
Temperature
at 70 km
altitude
(K)
c
Atmospheric
scale
height
at 70 km
altitude
(km)
Adiabatic
lapse
rate
at 70 km
altitude
(K
knr
1
)
Atmospheric lapse rate
at at 70 km
altitude
(K
km~')
e
Atmospheric
pressure
at
surface
(bars)
e
Mass
of
atmospheric column
(kg
m"
2
)
Total atmospheric
mass
(kg)
Surface temperature
(K)
c
Atmospheric
scale
height
at
surface (km)
Adiabatic
lapse rate
at
surface
(K
km"
1
)
Atmospheric lapse rate
at
surface
(K
km~')
e
605
1.
8
± 1.0
4.869
x
10
24
5.24
8.88
10.4
0.72333
0.0068
177.36
224.701
35.03
-243.0187
1.028
30
238
5.2
10.2
1.5
93
1.0
x
10
6
4.8 x
10
20
735
15.8
7.8
8
"Davies
et
al.
(1992)
b
The
astronomical
almanac
for the
year
1994
(1993)
c
Beatty
and
Chaikin
(1990)
d
Based
on
the
ratio
of
aphelion
to
perihelion
distance
squared
e
Seiff
(1983)
Figure
8.1
Altitude
profile
of
temperature
and
pressure
at
30
C
'
latitude
for the
atmosphere
of
Venus,
based
on the
standard
atmosphere.
After
Seiff,
A.,
1983,
"Models
of
Venus'
Atmospheric
Structure,"
in
Hunten
et al.
(1983;
cited
in
section
8.1),
p.
1045,
as
quoted
by
Prinn,
R.
G.,
1985, "The
Photochemistry
of
the
Atmosphere
of
Venus,"
in
Levine
(1985;
cited
in
section
8.1).
p.
281.
283
Figure
8.2
Altitude profile
of the
mass density
of
three
modes
of
cloud particles
in the
middle
atmosphere
of
Venus.
After
Knollenberg,
R.,
and
Hunten,
D.
M.,
1980, "The
Microphysics
of the
Clouds
of
Venus: Results
of the
Pioneer Venus
Particle
Size
Spectrometer
Experiment."
J.
Geophys.
Res.
85,
8039,
as
quoted
by
Prinn,
R. G.,
1985, "The Photochemistry
of the
Atmosphere
of
Venus,"
in
Levine
(1985;
cited
in
section 8.1).
p.
281.
Table
8.2
Composition
of the
atmosphere
of
Venus
Species
Mole
fraction
References
and
remarks
CO
2
N
2
H
2
O
Ar
CO
0
2
He
Ne
HC1
COS
S0
2
Kr
SO
HF
Xe
96.5 ±0.8%
3.5
±
0.8%
2 x
10~
5
6 x
10^
5
1.5
x
lO"
4
7
±2.5
x
10~
5
1.7
± 1 x
10~
6
3.0
± 1.8 x
10-
5
4.5 ± 1 x
10~
5
1
x
10~
3
< 3 x
10~
7
12+f
x
10~
6
7±3 x
10~
6
6 x
10~
7
3
x
10~
7
5-lOOx
10~
9
2.5
x
10-
8
2± 1 x
10~
8
5+2%
x
10~
9
1.9
x
10~
9
surface
22km
42 km
surface
42 km
cloud
top
100 km
(D
(2);
cloud
top
(3);
cloud
top
(4)
(5)
(6)
(3)
Unless
otherwise stated,
the
composition data
is
taken
from
von
Zahn
el
al.
in
Hunten
et
al.,
eds.
(1983).
(1)
See
discussion
in
Yung
and
DeMore
(1982);
(2)
Trauger,
J. T. and
Lunine,
J.
I.,
1983, Icarus
55,
272;
(3)
Connes,
P. et
al.,
1967,
Astrophys.
J.
147, 1230;
(4)
Crisp,
D.
et
al.,
1990. Nature 345, 508;
(5)
Zasova,
L. V. et
al.,
1993,
Icarus
105,
92; (6) Na. C. Y. et
al.,
1990,
J.
Geophys. Res.
95,
7485.
284
Venus
285
Table
8.3
Isotopic composition
of the
atmosphere
of
Venus
Species
D/H
I2
C/'
3
C
'
4
N/
15
N
I6
0/
I8
0
20
Ne/
22
Ne
35
C1/
37
CI
36
Ar/
38
Ar
40
Ar/
36
Ar
Ratio
1.6
± 2 x
10'
3
1.9
±
6x
I0~
3
86
± 12
88.3
±
1.6
273
± 56
500
± 25
500
± 80
11.8
± 0.7
2.9
± 0.3
5.56
±
0.62
5.08
±
0.05
1.03
±
0.04
1.19
±
0.07
Remarks
3
PVMS
IR
PVMS
From
von
Zahn
el
al.
in
Hunten,
D.
M.
et
al,
eds.
(1983).
a
PVMS=Pioneer
Venus
Mass
Speccromeler.
IR=Infrared
ground-
based observations.
The
bulk
of our
knowledge
of the
atmosphere
of
Venus
is
derived from observations
in
the
middle atmosphere
(60-100
km
altitude).
At
cloud-top levels
(65
km),
ground-
based
ultraviolet (UV) observations revealed
a 4-5 day
period, east-to-west circulation
that
is 60
times faster than
the
solid surface.
The
mechanism
for
generating this
superrotation
is not
well
understood.
The net
result
of
this
rotation
is
that
it
gives
the
upper atmosphere
an
effective
diurnal
cycle
of 4-5
days.
The
composition
of the
atmosphere
of
Venus
is
summarized
in
table
8.2.
Isotopic
composition
is
given
in
table
8.3.
The
major constituent
of the
atmosphere
is
CC>2
(96.5%).
The
next
most abundant
species
is N2
(3.5%).
Since Venus
is so
similar
to
Earth
in
size,
it is
useful
to
compare
the
atmospheres
of the two
planets. Table
8.4
lists
the
essential features that
are
compared.
The
greatest difference between Venus
and
Earth concerns
the
amount
of
water
on
these
two
planets.
The
mixing
ratio
of
water vapor
in the
lower atmosphere
of
Venus
is
variable,
with
a
maximum
value
of
1.5
x
10~
4
.
The
column-integrated water abundance
is
equivalent
to a
layer
of
2-10
cm of
water,
uniformly
spread over
the
surface
of the
planet.
For
comparison, Earth
contains
an
average layer
of 2.7 km of
water, residing mostly
in the
oceans. Whether
Venus
was
formed
dry or
whether most
of its
water
has
escaped
is
still
the
subject
of
debate.
The
lack
of an
ocean
on
Venus
has at
least three dramatic consequences
for
the
atmosphere. First, most
of the
planet's
CO2
remains
in the
atmosphere,
in
contrast
to
Earth, where most
of the 50
bars
of CO2 is
sequestered
as
carbonate rock
in
the
sediments. Second,
the
atmosphere
of
Venus contains large quantities
of
SC>2.
On
Earth, most
of the
volatile
sulfur
resides
in the
ocean
as
sulfate
ions.
The
presence
of
this
large amount
of
SC>2
in the
atmosphere
is
largely responsible
for the
production
of the
dense
H
2
SO4
clouds
on
Venus. Third,
the
atmosphere
of
Venus contains large
amounts
of
HC1.
On
Earth,
the
bulk
of
chlorine
is in the
form
of
salt
(NaCl)
in the
oceans.
The
photochemistry
of the
atmosphere
of
Venus
is
driven
by the
catalytic
chemistry
of
chlorine
in the
atmosphere.