Yung Y.L., DeMore W.B. Photochemistry of Planetary Atmospheres
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Appendix
4.1
(continued)
z
51
52
53
54
55
56
57
58
59
60
62
Element
Antimony
Sb
Tellurium
Te
Iodine
I
Xenon
Xe
Cesium
Cs
Barium
Ba
Lanthanum
La
Cerium
Ce
Praseodymium
Pr
Neodymium
Nd
Samarium
Sm
A
119
120
122
124
121
123
120
122
123
124
125
126
128
130
127
124
126
128
129
130
131
132
134
136
133
130
132
134
135
136
137
138
138
139
136
138
140
142
141
142
143
143
144
145
146
148
150
144
147
147
148
149
Atoms
(%)
8.6
32.4
4.56
5.64
57.3
42.7
0.091
2.5
0.89
4.6
7.0
18.7
31.7
34.5
100
0.114
0.111
2.16
27.60
4.34
21.64
26.53
9.69
7.82
100
0.106
0.101
2.42
6.59
7.85
11.2
71.7
0.089
99.911
0.190
0.254
88.5
11.1
100
27.2
12.2
23.8
8.3
17.2
5.7
5.6
3.1
15.1
11.3
13.9
Process
2
S, R
S, R
R
R
S, R
S, R
P
S
S
S
S, R
S, R
R
R
S, R
P
P
S
S, R
S
S, R
S, R
R
R
S, R
P
P
S
S, R
S
S, R
S, R
P
S, R
P
P
S, R
R
S, R
S
S, R
S, R
S, R
S, R
R
R
P
S, R
S
S, R
Abundance
1
"
at./10
6
Si
3.29
x
10-'
1.24
1.74
x
10~
'
2.15
x
10-'
2.02
x
10~'
1.50
x
10-'
4.5
x
10-
3
1.23
x
10-'
4.4 x
10~
2
2.26
x
10-'
3.44
x
10-'
9.18
x
10-'
1.56
1.69
9.0
x
10-'
4.96
x
10~
3
4.83
x
10~
3
9.39
x
10~
2
1.20
1.89
x
10-'
9.4
x
10-'
1.15
4.21
x
10-'
3.4
x
10-'
3.72
x
10-'
4.62
x
10~
3
4.40
x
10~
3
1.06
x
10-'
2.87
x
10-'
3.42
x
10-'
4.88
x
10-'
3.13
4.0
x
10"
4
4.48
x
10-'
2.2
x
10-
3
2.9
x
10-
3
1.026
1.29
x
10-'
1.74
x
10-'
2.27
x
10-'
1.02
x
10-'
1.01
x
10-'
1.99
x
10-'
6.94
x
10~
2
1.44
x
10-'
4.77
x
10~
2
4.68
x
10~
2
8.1
x
10~
3
3.94
x
10~
2
4.06
x
10~
2
2.95
x
10~
2
3.63
x
10~
2
16
Appendix
4.1
(continued)
z
63
64
65
66
67
68
69
70
71
72
73
74
Element
Europium
Gadolinium
Terbium
Dysprosium
Holmium
Erbium
Thulium
Ytterbium
Lutetium
Hafnium
Tantalum
Tungsten
A
150
152
154
Eu 151
153
Gd
152
154
155
156
157
158
160
Tb
159
Dy
156
158
160
161
162
163
164
Ho 165
Er 162
164
166
167
168
170
Tm
169
70
Yb
168
170
171
172
173
174
176
Lu
175
176
776
Hf
174
176
776
177
178
179
180
Ta
180
181
W
180
182
Atoms
(%)
7.4
26.6
22.6
47.9
52.1
0.20
2.1
14.8
20.6
15.7
24.8
21.8
100
0.057
0.100
2.3
19.0
25.5
24.9
28.18
100
0.14
1.56
33.4
22.9
27.1
14.9
100
0.135
3.1
14.4
21.9
16.2
31.6
12.6
97.39
2.61
0.16
5.2
18.6
27.1
13.7
35.2
0.0123
99.9877
0.13
26.3
Process
3
S
R
R
S, R
S, R
P
S
S, R
S, R
S, R
S, R
R
S, R
P
P
S
S, R
S, R
S, R
S, R
S, R
P
P,
S
S,R
S, R
S, R
R
S, R
P
S
S, R
S, R
S, R
S, R
R
S, R
S
P
S
S, R
S, R
S,R
S, R
P,
S, R
S, R
P
S, R
Abundance
6
at./10
6
Si
1.93
x
lO"
2
6.94
x
10~
2
5.89
x
10~
2
4.66
x
10~
2
5.06
x
10~
2
6.6 x
10~
4
6.95
x
10~
3
4.90
x
10~
2
6.82
x
10-
2
5.20
x
10~
2
8.21
x
10~
2
7.22
x
10~
2
5.89
x
10-
2
2.27
x
10-
4
3.98
x
10~
4
9.15
x
10-
3
7.56
x
10-
2
1.01
x
10-'
9.91
x
10~
2
1.12
x
10-
1
8.75
x
10-
2
3.54
x
10~
4
3.95
x
10-
3
8.45
x
10-
2
5.79
x
10~
2
6.86
x
10-
2
3.77
x
10"
2
3.86
x
10-
2
3.28
x
10~
4
7.53
x
10~
3
3.50
x
10-
2
5.32
x
10-
2
3.94
x
10~
2
7.68
x
10~
2
3.06
x
10~
2
3.59
x
10-
2
9.64
x
10~
4
1.06
x
10~
3
2.8
x
10-"
9.2
x
10~
3
9.02
x
10~
3
3.27
x
10~
2
4.77
x
1C"
2
2.41
x
10~
2
6.20
x
10-
2
2.78
x
10-
6
2.26
x
10-
2
1.78
x
10~
4
3.60
x
10-
2
(continued)
117
Appendix
4.1
(continued)
Z
Element
75
Rhenium
76
Osmium
77
Iridium
78
Platinum
79
Gold
80
Mercury
81
Thallium
82
Lead
83
Bismuth
90
Thorium
92
Uranium
A
183
184
186
Re
185
187
187
Os 184
186
187
187
188
189
190
192
Ir
191
193
Pt
190
192
194
195
196
198
Au
197
Hg
196
198
199
200
201
202
204
Ti
203
205
Pb
204
206
206
207
207
208
208
Bi
209
Th
232
232
U
235
255
238
238
Atoms
(%)
14.3
30.7
28.6
37.40
62.60
0.018
1.60
1.60
13.3
16.1
26.4
41.0
37.3
62.7
0.013
0.78
32.9
33.8
25.3
7.2
100
0.15
10.0
16.8
23.1
13.2
29.8
6.9
29.5
70.5
1.94
19.12
20.62
58.31
100
100
0.720
99.275
Process
3
S, R
S, R
R
S, R
S, R
P
S
S
S, R
S, R
S, R
R
S, R
S, R
P
S
S, R
S, R
S, R
R
S, R
P
S
S, R
S, R
S, R
S, R
R
S, R
S, R
S
S, R
S, R
S, R
S, R
RA
RA
RA
Abundance
1
"
at./10
6
Si
1.96
x
10~
2
4.21
x
10~
2
3.92
x
1C"
2
1.90
x
10~
2
3.17
x
10~
2
3.43
x
10-
2
1.29
x
10~
4
1.15
x
10~
2
1.15
x
10-
2
8.9 x
10~
3
9.54
x
10~
2
1.15
x
10-'
1.89
x
10-'
2.94
x
10-'
2.46
x
10-'
4.14
x
10-'
1.78
x
10-
4
1.07
x
10~
2
4.51
x
10-'
4.63
x
10-'
3.47
x
10-'
9.86
x
10~
2
1.86
x
10-
1
7.8
x
10-"
5.2
x
10~
2
8.74
x
10~
2
1.20
x
10-'
6.86
x
10~
2
1.55
x
10-'
3.59
x
10-
2
5.42
x
10~
2
1.30
x
10-'
6.12
x
10-
2
6.03
x
10-'
5.94
x
10-'
6.50
x
10-'
6.44
x
10-'
1.838
1.830
1.44
x
10-'
3.35
x
1C"
2
4.20
x
10~
2
6.49
x
10~
5
5.73
x
10-'
8.94
x
10~
3
1.81
x
10~
2
Source:
Anders,
E.,
and
Ebihara,
M.,
1982,
"Solar-System
Abundances
of the
Elements."
Geochim.
Cosmochim. Acta
46,
2363-2380.
"U
=
cosmological
nucleosynthesis;
H =
hydrogen burning;
N = hot
hydrogen burning;
He =
helium burning;
Ex =
explosive nucleosynthesis;
E
=
nuclear statistical equilibrium;
S =
j
process;
HeS =
helium-burning
s
process;
R =
r
process;
RA = r
process producing
actinides;
P = p
process;
X =
cosmic
ray
spallation.
b
ltalicized
values
refer
to
abundances
4.55
Gyr
ago.
118
Appendix
4.2
Meteorite classes
and
numbers
Class
Chondrites
CI
CM
CO
CV
H
L
LL
EH
EL
Other
Achondrites
Eucrites
Howardites
Diogenites
Ureilites
Aubrites
Shergottites
Nakhlites
Chassignites
Anorthositic
breccias
Stony irons
Mesosiderites
Pallasites
Irons
IAB
1C
IIAB
IIC
IID
IIE
IFF
IIIAB
IIICD
HIE
HIF
IVA
IVB
Other irons
Falls
5
18
5
7
276
319
66
7
6
3
25
18
9
4
9
2
1
1
0
6
3
6
0
5
0
3
1
1
8
2
0
0
3
0
13
Fall
frequency
(%)
0.60
2.2
0.60
0.84
33.2
38.3
7.9
0.84
0.72
0.36
3.0
2.2
1.1
0.48
1.1
0.24
0.12
0.12
0
0.72
0.36
0.73
0.08
0.45
0.05
0.09
0.10
0.03
1.42
0.14
0.10
0.05
0.39
0.09
1.32
Finds
non-Antarctic
0
5
2
4
347
286
21
3
4
3
8
3
0
6
1
0
2
0
0
22
34
97
11
60
7
12
13
4
189
19
13
6
52
12
175
Antarctic
0
34
6
5
671
224
42
6
1
3
13
4
9
9
17
2
0
0
1
2
1
4
0
6
0
0
0
0
0
0
0
0
1
0
0
Data
from
Sears,
D. W. G., and
Dodd,
R. T.
(1988).
119
5
Jovian
Planets
5.1
Introduction
The
four
giant planets
in the
outer solar system, Jupiter, Saturn, Uranus,
and
Nep-
tune,
are a
distinct group
by
themselves.
The
essential astronomical
and
atmospheric
aspects
of
these planets
are
summarized
in
table
5.1.
The
significance
of
this group
in
the
chemistry
of the
solar system
is
briefly
pointed
out in
chapter
4.
These
planets
are
composed primarily
of the
lightest elements, hydrogen
and
helium, which were
captured
from
the
solar nebula during formation.
The
planets have rocky cores made
of
heavier elements.
In the
case
of
Jupiter
and
Saturn
the
mass
of the gas
greatly
exceeds that
of the
core, whereas
for
Uranus
and
Neptune
the
masses
of gas and
core
are
comparable.
Due to the
enormous gravity
of the
giant planets, little mass
has es-
caped
from
their atmospheres. Hence,
the
bulk composition
of
these planets provides
a
good measure
of the
initial composition
of the
solar
nebula
from
which they were
derived.
Of all
planetary bodies
in the
solar system,
the
constituents
of
giant planets
are the
closest
to the
cosmic abundances
of the
elements.
The
chemistry
of the
atmospheres
of the
giant planets
is
interesting
for the
following
reasons:
1.
chemistry
in a
dominantly reducing atmosphere
2.
interplay between photochemistry
and
equilibrium chemistry
3. ion
chemistry
in
polar auroral regions
4.
heterogeneous chemistry
of
aerosols
5.
chemistry
of
meteoritic
debris
6.
lack
of a
planetary
"surface"
120
Jovian
Planets
121
Table
5.1
Astronomical
and
atmospheric
data
Characteristic
Jupiter
Saturn
Uranus
Neptune
Radius
(equatorial) (km)
Mass
(kg)
Mean
density
(g
cm~
3
)
Gravity
(surface, equator)
(m
s~
2
)
Semimajor
axis (AU)
Obliquity
(relative
to
orbital plane) (deg)
Eccentricity
of
orbit
Period
of
revolution (Earth days)
Orbital
velocity
(km
s~')
Period
of
rotation
(h)
Mass
of
atmospheric column
(1
bar)
(kg
cm"
2
)
Total
atmospheric
mass
2
(1
bar) (kg)
Equilibrium
temperature
(K)
Temperature
at 1
bar
11
(K)
Annual
variation
in
solar
insolation
0
Atmospheric
scale height (km)
Atmospheric
lapse rate
( K
km"
1
)
Escape velocity
(km
s~')
71492
1.899xl0
27
1.33
23.12
5.2028
3.08
0.0483
4332.71
13.06
9.841
0.447
2.87
x
10
20
128
165
1.21
27.0
1.8
59.5
60268
5.688
x
10
26
0.69
8.96
9.5388
26.73
0.0556
10759.5
9.64
10.233
1.15
5.26
x
10
20
98
134
1.25
59.5
6.9
35.6
25559
8.68
x
10
25
1.29
8.69
19.1914
97.92
0.0461
30684
6.81
17.9
1.19
9.76
x
10
19
56
76
1.21
27.7
6.7
21.22
24764
1.027x
10
26
1.64
11.00
30.0611
28.80
0.0086
60190
5.43
19.21
0.94
7.24
x
10
19
57
71.5
1.04
19.1
8.5
23.3
"Total
mass
of the
atmosphere
above
1 bar
pressure
level.
b
Data
from
Linda!
et
al.
(1992).
c
Based
on the
ratio
of
aphelion
to
perihelion
distance
squared.
We
briefly
comment
on
these reasons
in
this section. Each topic
will
receive
a
more
detailed treatment
in
later sections. First
of
all,
the
atmospheres
of the
Jovian planets
are
more than
90%
hydrogen
and
helium. Since helium
is
inert,
the
atmospheric
chemistry
is
dominated
by
hydrogen. Therefore,
we
would expect
the
most stable
compounds
of
carbon, oxygen, nitrogen,
and
phosphorus
to be
CHt,
H
2
O,
NHa,
and
PHs.
This
is in
fact
confirmed
by the
available observed composition
of the
bulk
atmospheres
of
these planets. However,
in the
upper atmospheres
of
these planets,
the
composition
is
controlled
by
photochemistry.
For
example,
one
consequence
of the
dissociation
of CH4
following
the
absorption
of
ultraviolet radiation
is the
production
of
C
2
He,
as
summarized schematically
by
(A
discussion
of the
detailed chemical pathways
is
given
in
section
5.3.1.)
The
fate
of
the H
atoms produced
in
(5.1)
is to
recombine
to
form
H2-
Thus, photochemistry
is
capable
of
driving
the net
endothermic
reaction
Here
we
point
out a
fundamental
difference between
the
giant planets
and
other
planetary
bodies
in the
solar system.
If the H and
H
2
produced
in
(5.1)
and
(5.2)
escape
from
planetary atmospheres, this results
in
irreversible change
in the
chemistry
of
these atmospheres.
As
discussed
in
chapter
1,
H and
H
2
do not
escape
from
giant
planets.
All the
products
of
photochemistry
are
preserved
in the
atmosphere. Since
1
22
Photochemistry
of
Planetary Atmospheres
is
stable
in the
atmosphere once
it is
removed
from
the
photochemically
active
region (mesosphere),
its
ultimate
sink
is in the
deep atmosphere, where
thermodynamic
equilibrium
chemistry favors
the
reverse
of
(5.2):
Reactions such
as
(5.3) that tend
to
restore
the
atmospheric composition
to
thermo-
dynamic
equilibrium occur
at
pressures greater than
1
kbar
and at
temperatures above
1000
K.
From this discussion
we may
infer
that
the
chemistry
of the
giant planets
consists
of two
distinct regimes: photochemistry
in the
upper atmosphere driven
by
solar ultraviolet light,
and
equilibrium chemistry
in the
interior
of the
planet driven
by
thermal energy. Mixing
processes
deliver
the
photochemical products
from
the up-
per
atmosphere
to the
interior,
and the
equilibrium products
are
recycled back
to the
upper atmosphere. This
is a
closed loop because there
is no
loss
of
even
the
lightest
molecules.
All
giant planets have extensive magnetic fields, giving rise
to
complex interactions
with
the
solar wind.
The
deposition
of
energetic particles
in the
polar regions gives
rise
to
aurorae that have been observed
from
both spacecraft
and
Earth.
The
dominant
chemistry
that
can
occur
in the
auroral zones
is ion
chemistry, reminiscent
of
what
is
known
to be
important
in the
molecular clouds (section 4.4). Organic synthesis
is
believed
to be
important.
The net
result
of the
photochemical destruction
of
CFLt
is its
conversion
to
higher
hydrocarbons. Ultimately
the
product species become heavy enough
to
condense,
giving
rise
to a
ubiquitous aerosol layer
in the
atmospheres
of the
giant planets known
as
Axel-Danielson
dust.
The
presence
of
this aerosol layer
has the
most profound
consequence
for the
thermal budget
of the
upper atmosphere. Indeed,
aerosol
heating
is a
primary cause
of the
thermal inversion
in the
atmospheres
of the
giant planets
above
the
tropopause (see
figure
1.2).
The
aerosols
may
also present
a
surface
for
possible heterogeneous chemical reactions.
The
giant planets, with
their
strong gravitational
fields,
tend
to
focus
meteoritic
materials into their
atmospheres.
In
addition, extensive ring
systems
are
also
sources
of
extraplanetary
material
as the
orbits
of the
ring particles decay
due to
internal
collisions,
tidal interactions,
or
atmospheric drag.
The
nature
of
this
influx
of
material
is
unknown,
but we can
assume that
it is
chemically very
different
from
that
of the
ambient atmosphere.
This
material
is
certainly rich
in the
heavy
elements
such
as O, C,
N,
S, Fe, Na, K, Mg,
Al,
and Si. The
fate
of
these elements
in a
reducing atmosphere
is
poorly understood
and may
give rise
to
species
that
we
would
not
expect
from
an
atmosphere
in
thermodynamic equilibrium.
The
Jovian planets
are
characterized
by the
lack
of
planetary surfaces.
The
atmo-
sphere merges continuously
with
the
bulk
of the
planet.
"Geochemistry"
as
defined
on
the
surface
of
smaller planetary bodies
is
absent.
The
visible atmosphere inter-
acts directly
with
the
deep interior
with
time constants that
are
orders
of
magnitude
smaller than
"geological"
time.
As a
consequence,
the
interior
has a
direct impact
on
the
composition
of the
upper atmosphere.
In
this chapter
the
composition
of the
atmospheres
of the
Jovian planets
is
critically
examined
and
discussed
in
terms
of the
physical
and
chemical
processes
that
are
known
or
remain
to be
identified.
In the
general discussions
the
four
planets will
be
taken
as
Jovian
Planets
123
a
group.
In
presenting quantitative results
of
model calculations
we first
report them
for
Jupiter, followed
by
similar results
for the
other giant planets.
5.1.1 Voyager Observations
Our
knowledge
of the
atmospheres
of the
giant planets
has
advanced
by a
quantum
jump
from
the two
spectacularly successful missions Voyager
1 and 2 to the
outer
solar system.
We
show examples
of the
Voyager observations
to
convey
the
sense
of
excitement
of
these discoveries.
The
Voyager spacecraft each
carried
11
instruments,
of
which
the
most important
for
atmospheric studies
are the
ultraviolet spectrome-
ter
(UVS),
the
radio science (RSS),
and the
infrared radiometer, interferometer,
and
spectrometer (IRIS).
The
ultraviolet Spectrometer probes
the
uppermost regions
of the
atmosphere,
the
thermosphere
and the
mesosphere where
the
sun's extreme ultraviolet
(EUV)
energy
is
absorbed
and
scattered.
The
thermosphere
is
also
the
region where
a new
source
of
ultraviolet radiation,
known
as the
electroglow,
may be
produced
by an
unknown
mechanism.
The
Voyager
UVS is a
grating spectrometer
with
128
channels covering
a
range
from
512
to
1690
A. The
spectral resolution
is 33 A. The
instrument
can
operate
in
the
airglow mode,
by
looking
at the
radiation emanating
from
the
disk
of the
planet,
or in the
occultation mode,
by
looking
at the sun or a
bright star
as it is
occulted
by
the
planet. Figure
5.1
shows
the
airglow spectrum
of
Jupiter (and Saturn) obtained
by
Voyager
1. The
single brightest feature
in the
spectrum
is
Lyman
a at
1216
A,
with
disk-center
intensity
of 15
kR
(1
kilo-rayleigh
= 1 x
10
9
photons
cm~
2
s~
]
),
arising
from
the
resonance scattering
of the
solar Lyman
a
line
by H
atoms
in the
upper
atmosphere
of
Jupiter.
To first
order
we may
regard
the H
atoms
as a
layer
of
white
scatterers
lying
above
a
deep
layer
of
perfect
absorbers—CH4.
The
column density
of
H
atoms required
to
reflect
this much Lyman
a is of the
order
of
10
17
cm~
2
.
The
strong emission features between Lyman
a and 800 A are
mostly
due to
Ha
Lyman
and
Werner bands.
The
integrated intensity
is
about
3 kR and
cannot
be
readily accounted
for
by the
scattering
of
solar
EUV
radiation.
In
view
of
this
difficulty,
this intense
EUV
radiation
from
Jupiter (and other Jovian planets)
has
been named
"electroglow."
An
obvious energy
source
is the
precipitating
electrons from
the
magnetosphere,
but
then
why is
this emission
not
observed
on the
dark side
of the
planets?
The
upper
limit
of H2
bands
on the
night
side
is
about
70
times below
the
observed brightness
on
the
day
side. Thus,
the
discovery
of the
intense
EUV
emissions raised
a
fundamental
question
as
deep
as the
discovery itself.
In
addition
to the
electroglow, Jupiter also
has
intense auroral emissions caused
by
energetic
particles
from
the
magnetosphere.
The
latter phenomenon
has a
terrestrial analog
and is
better understood.
Figure
5.2
gives
the
spectrum
of the sun
during occultation
as the
tangent
ray
descends
deeper
and
deeper
into
the
atmosphere
of
Jupiter.
At the
altitude
of 920 km
above
an
arbitrary reference level, there
is
virtually
no
attenuation
of the
solar
beam.
Attenuation
increases
as the
altitude
decreases.
Most
of the
absorption
is by
H
2
under
1000
A, but CH4 and
other hydrocarbons
become
important
at
longer wavelengths.
The
exact wavelengths where absorption
by
various molecules
become
important have
been discussed
in
chapter
2
(see
figures 2.8 and
2.10)
At
—80
km,
nearly
all the
solar
flux has
been extinguished
by the
atmosphere.
The
nature
of the
absorption spectrum
yields
information
on the
atmospheric composition.
By
modeling
the
rate
of
change
124
Photochemistry
of
Planetary
Atmospheres
Figure
5.1
Extreme ultraviolet
airglow
observations
of
Jupiter
and
Saturn
from
Voyager
1
encounters
in
July 1979
and
November
1980,
respectively.
After
Broad-
foot,
L. et
al.,
1981,
"Extreme
Ultraviolet Observations
from
Voyager
1
Encounter
with
Saturn." Science 212, 206.
of
atmospheric
opacity
with altitude, valuable information
concerning
the
temperature
(scale height)
and
distribution
of the
absorbing
species
may be
extracted.
One of the
important results
of
this analysis
is the
confirmation
of a hot
thermosphere
of
Jupiter.
The RSS
observations
on
Voyager were carried
out in the
occultation mode using
the
S-band
(A.
=
13 cm, v
=
2.3
GHz)
and the
X-band
(X
= 3.5 cm, v
=
8.6
GHz)
channels
of the
spacecraft's telemetry
as the
signal, which were received
on
Earth.
As
the
spacecraft went behind
the
planet, bending
of the
tangent
ray
occurred
due to the
difference
between
the
index
of
refraction
of the
atmosphere
and the
vacuum. Three
regions
of the
atmosphere that
can be
most sensitively probed
by
this technique
are
the
ionosphere,
the
stratosphere,
and the
troposphere. Unfortunately,
the
mesosphere
does
not
have enough electrons
or
neutral molecules
to
produce
a
detectable
signal
in
this
measurement.
An
illustration
of the
ionospheric measurements
by RSS is
given
in
figure 5.3 for
Saturn. Vertical profiles
of the
density
of
electrons
are
derived
from
the
radio occultation experiments
on
Voyager
1 and 2
(labeled
Vi
and V2 in the
figure). For
comparison,
the
results
from
the
earlier Pioneer mission (labeled
P/S in
the
figure) are
included. This
figure
shows that Saturn
has an
extensive ionosphere
that
exhibits considerable spatial
and
temporal variations (the spacecrafts' encounters
were separated
by a
number
of
years).
The
topside scale height
for the
ionosphere
is
in the
range
of
2000-3000
km.
This determines
the
value
r
p
/»j
(
,
where
T
p
is the
Jovian
Planets
125
Figure
5.2
Extreme
ultraviolet
solar
occultation
observations
of
Jupiter.
The
altitude
scale
is
arbitrary.
Large
Z
values
refer
to
high
atmosphere.
Lower
Z
values
refer
to
deeper
regions
of
atmosphere
where
solar
UV
radiation
is
progressively
absorbed.
After
Festou,
M.
C. et
al.,
1981, "Composition
and
Thermal
Structure
Profiles
of
the
Jovian
Upper
Atmosphere
Determined
by the
Voyager
Ultravi-
olet
Stellar
Occultation
Experiment."
J.
Geophys.
Res.
86,
5715.
sum of the
electron
(T
e
)
and the ion
(7})
temperature
and
mi
is the
mean
mass
of the
ion.
The
most likely
ion at
this high altitude
is a
proton. Therefore,
we may
conclude
that
the
temperature
of the
ionosphere (assuming
T
e
=
T
f
)
of
Saturn
is in the
range
of
300-600
K.
This
is
much greater
than
that expected
on the
basis
of
solar
EUV
heating.
The new
energy
source
has not
been
identified
to
date.
A
similar problem
exists
with
the
thermospheric
temperatures
of all
other Jovian planets.
The
lower part
of
the
electron density profile reveals
the
existence
of
sharp
layers.
These
features
have
been observed
in all
radio occultations
of the
giant planets.
One
explanation
is
that
they
are
layers
of
metal ions derived
from
meteoritic debris.
Figure
5.4
shows
an
example
of
what
can be
deduced
from
RSS
measurements
in
the
stratosphere
and the
troposphere.
The RSS
signal yields
T/m as a
function
of
altitude, where
T =
mean temperature
and m =
mean molecular mass. From
an
independent knowledge
of
atmospheric composition (e.g.,
from
UVS and
IRIS),
the
thermal structure
of the
atmosphere
may be
determined. Figure
5.4
shows that
the
RSS-derived
profiles
exhibit
considerably more vertical structure
than
those
from
IRIS measurements
and
model
computations.
The
reason
is
that
the
latter give mean