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Pressure-induced absorption by H2 in the atmosphere of Jupiter and Saturn
ICARUS
27, 391-406
(1976)
Pressure-Induced
Absorption
by Hz in the Atmospheres and Saturn
of Jupiter
TERRY Z. MARTIN,’ DALE P. CRUIKSHANK, CARL B. PILCHER, AND WILLIAM M. SINTON Institute
for Astronomy,
UnCversity of Hawaii, Honolulu, Hawaii, 96822
Received September 23, 1978 The S( 1) line of the pressure-induced fundamental band of H2 was identified and measured in the spectra of Saturn and Jupiter. This broad line at 4750cm-’ lies in a region free from telluric and planetary absorptions. It is about 99% absorbing in the core; the high-frequency wing extends to at least 5100om-‘. We compare the observed line shape to the predictions of both a reflecting-layer model (RLM) and a homogeneous scattering model (HSM). The RLM provides a good fit to the Saturn line profile for temperatures nesr 160K; the derived base-level density is 0.52 (+0.26, -0.17) smagat and the H, abundance is 25 (+10, -Q)km-smagat, assuming a scale height of 48km. The Jupiter line profile is fit by both the RLM and HSM, but for widely differing temperatures, neither of which seems probable. The precise fitting of the observed S(1) line profile to computed models depends critically on the determination of the true continuum level; difficulties encountered in finding the continuum, especially for Jupiter, are discussed. Derived RLM densities and abundances for both planets are substantially lower than those derived from RLM analyses of the H, quadrupole lines, the 3~ band of CH,, and from other sources.
The study of hydrogen absorptions in the atmospheres of Jupiter and Saturn is complicated by the particular nature of the absorptions and the complex atmospheric structure on both planets. For example, in measurements of the collision-narrowed quadrupole lines, high spectral, spatial, and temporal resolutions are required if a meaningful analysis is to be conducted (Hunt and Bergstralh, 1974). Pure rotational pressure-induced absorption by hydrogen is the dominant opacity source in the thermal infrared region of the spectra of Jupiter and Saturn, but its observation is hampered by telluric absorptions. Measurements of the pressure-induced vibrational bands in the near-infrared are similarly hampered by the strong methane absorptions occurring in major planet l Present address: Department of Geophysics and Space Physics, University of California., Los Angeles, California 90024. Copyright 0 1976 by Academic Press, Inc. All rights of reproduction in any form reserved. Printed in Great Britain
spectra in the 4000-12 OOOcm-’ region. Despite this difficulty, some studies of the pressure-induced vibrational bands have been made. The 3-O second overtone band was first noted in the spectra of Uranus and Neptune by Herzberg (1952), thereby establishing the presence of hydrogen on the major planets. A feature at the position of the 4-O band in the Uranus and Neptune spectra is probably not due to hydrogen (Belton and Spinrad, 1973). Neither the 3-O nor 4-O band is evident in the spectra of Jupiter or Saturn, presumably because of the lower abundances existing in their optically accessible atmospheres. The 2-O pressure-induced band in spectra of Jupiter and Saturn has recently been studied by de Bergh et al. (1974). This band is unfortunately coincident with a wide absorption feature of methane that prevents accurate determination of the hydrogen contribution. Upper limits were set to the abundance of H, that would be compatible with the observed absorption 391
392
MARTIN
on both planets, and an estimate was made of pressures in the region of line formation. The 1-O fundamental band was observed in low-resolution (X/AX N 40) spectra of Jupiter obtained by the balloon-borne Stratoscope II experiment (Danielson, 1966). A reflecting-layer model (RLM) analysis of the line shape led to a temperature estimate of about 250K, following the assumption of an H, abundance. However, the low resolution of these data makes quantitative interpretation somewhat uncertain. Cruikshank, Fink, and Larson noticed the probable appearance of the S(1) line of the H, fundamental band at 4800cm-’ in their spectra of Saturn obtained jointly with Kuiper (Fink and Larson, 1971). This absorption was also evident in Jupiter and Saturn spectra by Johnson (1970) but was not identified. Both these Saturn spectra showed contamination by radiation from the rings. Welsh (1969) has presented background material necessary to understand the behavior of the pressure-induced bands. Lines in the bands are broad ; their shapes are predictable, but the relative intensities of lines and the effects of the presence of foreign molecules depend in a complex way upon temperature. The S(1) “line” is comprised primarily of three components: the S,(l), S,(l)+ Q,(O), and X,(l)+Q,(l) lines. Here S denotes a transition in which the molecular rotational quantum number J changes by +2, Q denotes AJ = 0, the subscript is the change in vibrational quantum number, and the term in parentheses is the J value of the initial state. Two of the lines are double transitions-each molecule of the colliding pair contributes one of the noted transitions simultaneously. The absorption coefficient for the S(1) line also contains contributions from the high-frequency tails of other lines in the 1-O band (see Fig. 2 of Welsh, 1969). All component lines are intrinsically broad because of the very brief time during collisions when transitions may occur. We discuss in this paper the S(1) line in spectra of Jupiter and Saturn obtained at Mauna Kea Observatory. This material is
ET AL.
taken from the doctoral dissertation of one of us (TZM). DATA REDUCTION The spectra of Jupiter and Saturn in the interval 4000-6000cm-1 are presented in Fig. 1. Each spectrum has a resolution element of 0.5cm-’ and represents an average of several measurements. Saturn was observed in February and March 1974, and Jupiter in July 1974. Our spectra were obtained with a Digilab Michelson interferometer system mounted at the Cassegrain focus of the Mauna Kea 224~cm telescope. The spectrometer accepted a circular beam 18arcsec in diameter which covered Saturn’s entire disk and a small portion of the rings. For observations of Jupiter the aperture was centered on the planet’s disk. Absorption in the spectra between 4100 and 4700cm-’ is due to both methane and hydrogen, and in the case of Jupiter, ammonia. Methane absorption is also evident beyond 5500 cm-‘. The H, S( 1) line of interest occurs between 4650 and 5050 cm-’ ; telluric water vapor is responsible for most features in the 51005500 cm-’ range. The zero levels of these spectra for analytic purposes are defined by a straight line drawn between the intensity level at and the parts of the 4100-4200 cm-’ 5300cm-* water band where the zero can be well established. The 4100-4200cm-’ interval is totally absorbed by H, and CH,. Because these spectra are unapodized, some absorption features project well below the drawn zero level. We corrected the Saturn spectrum for ring radiation contamination in two ways. At the telescope, a reference aperture of the sky-subtracting foreoptics was positioned partly on the ring to compensate for the ring signal due to light in the main aperture. Guiding errors, however, led to a residual contamination of the disk spectrum. That contribution was subsequently removed as follows. Metha,ne has a strong Q-branch absorption at 4550cm-I, at which frequency the rings are relatively bright. It was assumed that
H,
4000
_LLL
ABSORPTION
ON JUPITER
AND SATURN
I
I
4200
4400
4600
J
Cb m-i
I
I
48co
4600
I
5400
5000'
5200
I
I
5600
5800
6000
FIG. 1. Spectra of Jupiter (top) and Saturn, 4000_6000cm- l. Resolution = 0 6cm-‘. absorption features due to planetary CH,, Jovian NH,, and telluric CO2 are indicated.
Selected
394
MARTIN
any apparent intensity level above zero at 4550cm-’ was in fact due to the rings. Knowing the shape of the ring spectrum measured with the same instrumentation, we subtracted at all frequencies of interest the appropriate fraction of the amount of contamination at 4550 cm-‘. The slope of the solar continuum and the instrumental response function were removed by dividing the data by lunar spectra. These ratios were then multiplied by the lunar spectral reflectivity determined from Apollo 11 soil samples (McCord and Johnson, 1970). In order to plot the planetary absorption as relative intensity (ly/llcontinuum), the continuum level must be determined at some frequency. The continuum is unfortunately not accessible near 5000 cm-’ ; the closest useful window in the methanedominated spectra is at 6300 cm-l. There, the Saturn spectrum resembles that of the moon in detail ; methane lines are weak and isolated. Jupiter’s spectrum at 6300 cm-’ shows considerable absorption from gaseous ammonia. However, that frequency remains the best at which to judge the continuum. If hydrogen itself is an opacity source at 6300cm-‘, then our derived H, line shapes will be in error. However, for reasonable temperatures, the H2 absorption coefficient falls off at such a rate that only about 5% absorption is expected at 6300cm-‘. From the continuum level at. 6300cm-‘, we then calculated the corresponding level at 5100 cm-’ according to 1,(5100) = Ic(6300) r=;]
Moon[=;I,,,,
where c denotes continuum and R is reflectivity. Clearly, if the planetary continuum is not actually flat we produce an erroneous value at 5100 cm-‘. No better method is evident, however. To investigate the shape of the background planetary absorption above 5000 cm-‘, we must account for the effects of telluric H,O, CO,, and planetary CH, and NH,. The comparison with lunar spectra serves to establish the small contribution
ET AL.
of telluric components between 5025 and 5075 cm-‘. We used the methane spectra of Cruikshank and Binder (1968) and our own ammonia spectra to assess the importance of these compounds. We thereby derived upper and lower limits to the values of relative intensity that occur in the 50005100cm-i region ; these limiting values have been retained in the analysis below. We determined the lower limit to the relative intensity by drawing the lowest curve that can be smoothly extended from the 4900cm-1 region through the signal level at 5020cm-’ and yet fall above the peaks of intensity that occur in the H,O band. Allowing for no continuum absorption by CH, at 5025 cm-l, we derive a 5100cm-’ lower limit of Iv/I, = 0.45 for Saturn and 0.26 for Jupiter. To obtain an upper limit, we estimated a depression of 5% by CH, at 5025cm-‘, 6% by NH, (for Jupiter), and then arbitrarily doubled these effects. A smooth curve drawn from the 4900cm-’ region through the high points at 5025cm-’ gives a rough value of Iv/I, = 0.63 at 5100cm-’ for Saturn and 0.36 for Jupiter. We may now estimate the percent absorption due to II, at 5100cm-‘. Allowing for 1.0% CH, absorption, the remaining component is 30 to 49% for Saturn. With an additional 11% absorption by NH, on Jupiter, the absorption attributed to H, is between 55 and 67% of the continuum level. Thus, the hydrogen line is apparently more extensive on Jupiter, a fact due either to greater abundance, higher temperature, or a different mode of line formation. The reduced planetary values of IV/I, are given in Table I. We now must establish that methane and ammonia cannot account for the broad absorption we attribute to hydrogen. In Fig. 2 we reproduce Cruikshank and Binder’s (1963) ambient-temperature spectra of various quantities of CH, in the range 4500-5250cm-*. Plotted with their data is our reduced Saturn profile, with the continuum level taken coincident with the top edge of the figure. Methane cannot match the extent and central maximum of the observed absorption. Indeed, at lower temperatures the methane R-branch ex-
H, ABSORPTION TABLE
ON JUPITER
I
RELATIVE INTENSITY OF JUPITER AND SATURN, 4600-5 lOOcm-’ Frequency cm-’ 4600 4625 4660 4675 4700 4725 4750 4775 4800 4825 4850 4875 4900 4925 4950 4975 5000 5025 5050 5075 5100
Jupiter
Saturn
IV/I= 0.009 0.02 0.024 0.02 0.015 0.011 0.011 0.015 0.023 0.0355 0.05 0.0685 0.0885 0.113/0.115 0.135/0.139 0.157/0.169 0.179/0.203 0.20/0.239 0.221/0.28 0.24/0.317 = 0.26/0.36
I”/& 0.0 0.0145 0.0375 0.0365 0.0225 0.008 0.005 0.017 0.0345 0.0555 0.078 0.106 0.145 0.19/0.193 0.23310.243 0.278/0.306 0.319/0.366 0.355/0.438 0.391/0.513 0.42410.579 = 0.4510.63
Lower limit/Upper
limit
tending from 4570 to 4700 cm-’ will change so that maximum absorption occurs at lower frequencies ; the higher frequency 5200
0 1.90
5100
I.95
5030
2.00
395
AND SATURN
lines weaken. Methane bands show fine structure that is not evident in the broad feature of the planetary spectra. Finally, several of the methane R-branch manifolds appear in our spectra between 4620 and 4660cm-L (Fig. 1). These would be t’otally indistinguishable because of saturation if methane were present in sufficient quantities to produce strong absorption above 4700cm-L. Ammonia has several strong bands which may affect the Jovian spectrum in the region of interest. For example, the u, + vq band at 5050cm-’ is clearly visible in Fig. 1. Laboratory spectra indicate that abundances of NH, which produce spectra resembling the Jovian 5050cm-’ band show very little absorption between 4600 and 4900cm-‘. Ammonia frost shows an absorption at 5000 cm-’ corresponding to the vj + v,, band, but the shape and location do not match the broad feature we observe. Ammonium hydrosulfide, a possible cloud component in the lower Jupiter and Saturn atmospheres (Weidenschilling and Lewis, 1973), has a sloping absorption similar in shape to the feature we discuss here (Kieffer and Smythe, 1973), but both NH, and NH,SH have strong bands at 3300 cm-l which are not evident in the Saturn spectrum. We conclude, therefore, that the only consistent explanation of the broad feature CM-’
4900
205
4800
4700
2 IO
46co
2.,5
4500
2 20/A
FIG. 2. Spectra of methane (Cruikshank and Binder, 1968) and the relative intensity of Saturn. The continuum level for Saturn is at the top edge of the figure. 14
396
MARTIN
that we observe is that it is due to the S( 1) pressure-induced line of H,. MODELS AND PREDICTED LINE SHAPES A. Reflecting-Layer
ET
AL.
an adiabatic density distribution, but these varied little from models based on the simpler isothermal density distribution which has the form
Model
In order to interpret the planetary spect.ra, we have computed the pressureinduced spectrum in the S(1) line for a number of model atmospheres using available H, absorption coefficients (Chisholm and Welsh, 1954 ; Hunt and Welsh, 1964; Watanabe and Welsh, 1965, 1967 ; Watanabe, 197 1). We first consider the simple reflecting-layer model (RLM) which has been widely applied to studies of Jupiter and Saturn with varying degrees of success. We characterize the gas by a given density and temperature at the reflecting level, and some distribution of density and temperat,ure with altitude. The intensity relat,ive to the continuum at frequency v for pressure-induced hydrogen absorption in a cell of length 1 is given by Iv/I, = exp (-CC”vp2 I), where p is the density in amagat (agt) and CL\, is the absorption coefficient in agt-*. The strong density dependence implies that the atmospheric layers near the reflecting surface will be heavily weighted. Consequently, the vertical temperature variation will not strongly affect the line shapes. The intrinsic temperat’ure dependence of the aV term is weak over the range of temperatures likely to occur in one scale height. Absorption coefficients are not available in the literature which adequately cover the expected range of important temperatures in the planetary atmospheres (loo15OK). VC’e have therefore calculated absorption coefficients for 130 and 150K, using the approach of Hunt and Welsh (1964). We estimate the absorption coefficient error to be +2.5% at 4800cm-‘, increasing to &100/o at 5000cm-‘. In the absence of the availability of H, data for other temperatures in numerical form, we took data from published curves; the numbers are given together with our calculated data in Table II. We have constructed models based upon
P(L) = poexp (-W),
where p. is the base level density, h is altitude above that level, and H = kT/pg is the scale height, with k being Boltzmann’s constant, T temperature, TVthe mean molecular weight, and g gravitational acceleration. The relative intensity in the vicinity of a pressure-induced line produced by two successive vertical traversals of the model atmosphere is given by Iv/I, = exp [ -va y 7) _f,” p’(h)dh] =
exp (-k
~Hpo2),
(1)
where 7 is an airmass factor having a minimum value of 2. If a more realistic adiabatic bemperature gradient is assumed, the relative intensity has precisely the same form as (1) above, but with the argument of the exponential different by 17% ; here the base-level temperatures are taken to be the same in both cases. Derived values of p. will differ by about So/. In the following discussion we assume the isothermal model for simplicity. The Jovian observations are matched to line profiles computed for the disk center, since the size of the interferometer entrance aperture limited the maximum included value of 7 to 2.22. The Saturn data were compared with a disk-integrated RLM profile, computed from the formula : tlvllc)Disk
=
2
’ exp I--va, po2 H/( 1 - G)“*] r dr s =” (1 - C) ewe - C2 Ei(-C), (2)
where r is a radial integration variable and C = va,po2 H. A computer program was used to calculate RLM profiles from (1) and (2) for va,rious temperatures. Families of line profiles corresponding to different abundances of hydrogen were computed. The results of this analysis are presented in Figs. 3-5 and Fig. 7. The first of these shows disk-integrated RLM profiles for various temperatures and a common value
H, ABSORPTION ONJUPITERAND TABLE
397
SATURN
II
ABSORPTION COEFFICIENTS” FOR NORMAL H,
Frequency cm-l 4600 4625 4650 4675 4700 4726 4760 4775 4800 4826 4850 4875 4900 4926 4950 4975 5000 5025 5050 5075 5100
20K
77K
95K
130R (do.)
150K (talc.)
196K
248K
300K
0.215 0.17 0.135 0.14 0.425 1.00 1.26 0.84 0.485 0.30 0.21 0.16 0.12Ei 0.095 0.08 0.065 0.055 0.05 0.05 0.053 0.06
0.20 0.185 0.215 0.30 0.478 0.61 0.658 0.54 0.37 0.265 0.20 0.15 0.11 0.085 0.069 0.056 0.047 0.039 0.034 0.03 0.026
0.314 0.27 0.262 0.29 0.352 0.415 0.441 0.40 0.342 0.28 0.209 0.15 0.112 0.085 0.06 0.043 0.031 0.023 0.017 0.011 0.006
0.345 0.325 0.348 0.422 0.644 0.634 0.62 0.625 0.422 0.333 0.265 0.214 0.176 0.148 0.126 0.109 0.095 0.084 0.075 0.067 0.061
0.369 0.346 0.372 0.44 0.646 0.618 0.604 0.62 0.426 0.343 0.277 0.227 0.188 0.158 0.136 0.118 0.103 0.091 0.082 0.073 0.067
0.368 0.389 0.43 0.619 0.582 0.608 0.582 0.611 0.426 0.359 0.296 0.25 0.21 0.183 0.158 0.138 0.12 0.104 0.089 0.077 0.066
0.443 0.44 0,477 0.49 0.518 0.625 0.613 0.49 0.44 0.39 0.36 0.31 0.273 0.244 0.219 0.198 0.18 0.162 0.144 0.127 0.112
0.435 0.43 0.45 0.485 0.63 0.666 0.672 0.55 0.495 0.436 0.38 0.336 0.295 0.265 0.235 0.21 0.185 0.17 0.152 0.14 0.13
a Units:amagat-2x 10pq.
RLM
Saturn
-----
FIG. 3. Reflecting-layer model S( 1) line profiles for poZH = 13.21unagt2 and indicated temperetures. The planetary profile is shown as the dashed line.
398
MA.RTIN ET AL.
4600
4700
4900
a00
CM
FIG. 4. Reflecting-layer amagats.
II00
4700
4100
-----
4900 CM
FIG. 5. Reflecting-layer
5000
model line profiles for 16OK, H = 5Okm, and indicated densities p,, in
Jupiter
4600
.’
-’
model line profiles for po2H = 18.2km@.
of po2H. The reduced Saturn profile is included, with upper and lower limits indicated for the far-wing intensity. The value po2 H = 13.2kmagt2 was chosen to give the best fit with the 15OK curve. The general depression of the intensity with increased temperat,ure is demonstrated. Figure 4 shows the RLM disk profiles for constant temperature and varying baselevel dens&y po. H is held constant at 50km. As the den&y changes from 0.4 to 0.7agt, the line profile rapidly becomes totally absorbing in the core. The wings also
so00
5100
and indicated temperatures.
demonstrat,e a strong density dependence. Clearly, if the Saturn line profile is roughly correct, base densities much different from 0.5agt are not permitted. The Jupiter spectrum is presented in Fig. 5 together with RLM center-of-disk profiles at three temperatures. The value of po2 H is chosen to give the best fit at 248K. Calculated profiles at the lower temperatures illustrated do not give satisfactory agreement with the data for any values of po2 H. This problem will be discussed later.
H2
ABSORPTION
ON JUPITER
B. Homogeneous Scattering Model The case of line formation in a hazy atmosphere has been treated by Chamberlain (1970), following the work of Chandrasekhar (1950), van de Hulst (1952), and others. We shall contrast the homogeneous scattering model (HSM) of Chamberlain with the RLM, although both of these are clearly physical oversimplifications. The HSM assumes a semi-infinite atmosphere ; isotropic scattering occurs with no redistribution of photon frequency. All parameters are independent of depth. The relative intensity is given by I,/& = & H(%
CL)HP,,
PO)/
3, H(G 14 H(%, PO), where the single-scattering albedo is & in the continuum and 6, at frequency v in the line, p. and p are, respectively, the direction cosines for rays entering and leaving the atmosphere, and H(&, p) are the Chandrasekhar (1950) H-functions for isotropic scattering. Jupiter and Saturn were near opposition during our observations, so we assumed p = pO. Employing the van de Hulst (1952) approximation to the H-functions H(G, P) = HP, cl)/{1 + /43(1 - 41”2), expression for relative intensity :
we get the
I,,Z, = 2
l/2
2
1 +- I-43(1 - a1 WC ( 1 $ /.43(1 - &,)]“2 I *
(3)
AND
399
SATURN
This formula was used to prepare model line profiles for differing values of p and integrated over the disk to provide profiles for comparison with the Saturn spectrum. The single-scattering albedo 8, is given by Gl, = (c$/u + l/f&)-‘, where u is the volume scattering coefficient. A computer program was used to calculate line profiles from (3), with various values of ay, CJ,and G,. For Saturn, since the whole disk was measured, the disk-integrated equation analogous to (3) was computed from the following expression : Vvl~chMc
=
J-L (4 0
2 55 where
P = [3( 1 - GJ]“2,
1
1 + P(l r2)“2 1 + &(l - G)“2
1rdr, 2
Q = [3(1 - t~J]l’~.
The result is (JvlZ&isk = 2(%&&)
{P/Q2)[In (1 + Q) - Q/(1 + &)I + W’/Q3)[-21n(l +&I +&-k&/(1 +&)I + (P2/Q4) [3ln(l + Q)+&'/2- 2Q - Q/U + Q,l>. The results of the HSM analysis are presented in Figs. 6-9. The upper part of
FIG. 6. Behavior of homogeneous-scattering model line profiles for 77 K, p = 1 agt. Upper scattering coefficient 0 = 5 x lo-‘cm-‘, various single-scattering elbedos d, indicated. 8, = 1, various crindicated.
: volume Lower:
400
MARTIN
ET
AL.
150 K
0.1.
0.6 -
0.4. 0.2 : ; : go., ’ 0.6-
HSM
0.4 *
0.1* 4600
4100
4.00
4900 CM
.’
FIG. 7. Center-to-limb variation in the RLM and HSM at 150K. ~_tindicated. Lower : HSM, o = 6 x lo-‘cm-‘, d, = 1, p = 1 apt.
0.1 .
0.6 .
5000
SlOO
Upper: RLM, direction cosines
HSM
Saturn
-----
: ; E 0.4 -
I’
,’ ,’
/ C
,,=-
z
FIG. 8. Homogeneous-scettering model line profiles (disk-integrated) u = 2 x lo-‘cm-‘. 130K: 0 = 2 x lO”cm-‘.
Fig. 6 displays the effect of varying the continuum singlescattering albedo at constant a for a temperature of 77K and a density of 1 agt. Fixing 9, and changing the volume scattering coefficient produces similar changes in line shape, shown in the lower part of Fig. 6. The center-to-limb variation of the RLM and HSM line profiles is shown in Fig. 7. In the former, absorption increases near the limb, exactly as it would with a change in base-level density. The HSM behavior is the opposite ; the relative intensity increases near the limb. Figure 7 emphasizes
for p = lagt,
3, = 1. 77K:
the diagnostic importance of observations of the center-to-limb variation on Jupiter and Saturn. In Figs. 8 and 9 the HSM line profiles which best match the Jupiter and Saturn spectra are plotted. In Fig. 8, curves for 77 and 130K and density 1 agt are shown. These temperatures bracket the expected range of Saturn temperatures near the cloud layers. Scattering coefficients were chosen to give a reasonably good overall fit to the planetary profile, although the choice is obviously somewhat arbitrary. The value of 3, is 1 .O ; curves with smaller
H, ABSORPTION
ON JUPITgR
401
AND SATURN
HSM
Jupiter
4600
,700
4800
-----
,900 CM
-’
5000
5100
FIG. 9. Homogeneous scattering model line profiles (center of disk) for p = 1. 77K: cm-‘, 6, = 0.94. 130K: o = 5 x 10-Scm-l, c3, = 0.5.
values will fall above the ones that are plotted. Figure 9 shows the Jupiter profile and the 77K and 13OK HSM curves. The 77K line gives a good match for p = 1, u = 2 x 10-7cm-‘, and G, = 0.94; the 130K curve does not give a good match for any value of these parameters. DISCUSSION
A. Xatum The RLM appears to give the best fit to the Saturn line profile. The HSM profiles are not sufficiently narrow even at 77K. The core of the line is more strongly absorbing than the HSM profiles (Fig. 8), so that allowance for errors in determination of the continuum level will not be able to produce agreement over more than a fraction of the observed line width. The RLM profiles produce a best fit for a temperature in the range 150--200K. That is higher than expected cloud-top temperatures for Saturn (Weidenschilling and Lewis, 1973), and higher than measured rotational temperatures that presumably refer to similar levels (Encrenaz and Owen, 1973; Trafton, 1973; the RLM was employed for both these investigations). However, we do not expect accurate temperatures from the H, line study, since the temperature-dependence of ay is small. Most likely, our result implies
CT= 2 x lo-’
that the continuum level used at 5100cm-’ was too high. A lower continuum gives a narrower line shape, which is then matched For lower-temperature models. by example, if the continuum is just 15% lower, an excellent fit is obtained using the 130K profile. Any continuous absorption that is stronger near 5100cm-’ than at 6300 cm-’ could thus help explain the RLM temperature discrepancy. Ammonia frost and ammonium hydrosulfide both have absorptiona near 5000cm-’ that could be partly responsible. However, the absence of the 3300cm-’ absorptions of these compounds in the Saturn spectrum has already been mentioned. The possibility exists that solar radiation scattered back from hotter (and hence lower) regions of the atmosphere contributes to the 4600-5100cm-L intensity. In such a case, the line profile might have a broader shape than has a line formed purely above the clouds. However, this possibility seems fairly unlikely in light of current multiple-layer cloud models for the planets. The abundances of H, existing below the upper cloud are sufficient to absorb totally any incoming radiation in the 4600-5100cm-’ range before it can again escape from the clear space between the upper and lower cloud decks. The argument is supported by Fig. 4 ; lower cloud level densities are probably at least 3agt for both planets, and the intercloud distance is well over a scale height. The
402
MARTIN
product p21 for the between-cloud region is then on the order of 200 kmagt2. If all the radiation observed is reflected from the upper cloud alone, and scattering effects are unimportant, we can derive a cloud-level density and above-cloud H, abundance from the RLM analysis. The parameter X = po2H is determined from the fitting of the 150K curves to the planetary spectrum. We may then solve for p,, and the H, abundance W = poH: po
=
GVV’2,
W = (XH)“2.
The scale height H is calculated, using g = 1000cmsec-2, p = 2.25amu (following Weidenschilling and Lewis for a solarcomposition atmosphere), and T = 13OK. Then H = 48km, and with X= 13.2km agt2, we have p. = 0.52agt,
W = 252kmagt.
The effects of errors in X and Hare reduced by the square-root dependence of p. and W. The best way to assess the total error in these quantities is to let all contributing errors fall so that the derived values are maximized. These are listed in Table III. We cannot, of course, assess errors arising from the assumption of a flat continuum between 5 100 and 6300 cm-‘. By contriving in all respects to maximize the derived po, we arrive at a value po(max) = 0.78agt using X = 22.5kmagt2 and H = 37km. In the other direction, we have X = 8.1 km agt2, H = 55km, and po(min) = 0.38agt. Applying the same principles to the derivation of W, we get finally base-level density : p. = 0.52 ?t:$amagat, H, abundance above reflecting layer : W = 25.2+i” km amagat. It should be emphasized that the values of p. and W derived depend little on the temperature used to select absorption Although one temperature coefficients. gives a best fit, the curves for other temperatures which match the planetary profile to some degree also yield densities and abundances near those quoted. These values are discrepant with existing atmospheric models and with conclusions about pressure and abundance derived
ET AL.
from RLM analyses of the methane 3v, band and the H, quadrupole lines. The models (Weidenschilling and Lewis, 1973 ; Palluconi, 1972) indicate densities near 4amagat at the level of NH, cloud formation. Measurements of the widths of collision-broadened lines have implied that pressures are around 2-3 atm (de Bergh et al., 1973) where the 3vJ lines are formed. Encrenaz and Owen (1973) derived an H, abundance of 76 f 20kmagt from the quadrupole lines. However, these have been found to exhibit considerable time variation for Jupiter (Hunt and Bergstralh, 1974). A recent study of the 2-O pressureinduced band (de Bergh et al., 1974) yielded upper limits for pressure and abundance of 0.6atm and 63kmatm, respectively, which are not inconsistent with our derived values. Clearly, we are presented with a problem of interpretation. If the models and most other RLMderived values are meaningful, then there should be no intensity evident in the Saturn spectrum between 4600 and 5100 cm-‘. The pressure-induced opacity of the expected abundances of H, would cause total absorption in that range. Since we do detect radiation at these frequencies and can measure the H, line shape, either the existing models and interpretation of measurements are in error, or the RLM is not adequate to explain line formation near 5000cm-‘. We have already seen that the HSM does not do well. The next logical step is to investigate the effects on the line shape of anisotropic scattering functions or vertical inhomogeneities. An inhomogeneous scattering model has been applied with success to Jupiter (Bergstralh, 1973 ; Hunt, 1973). At the same time, however, it is imperative to observe the center-to-limb variation of the 4750cm-’ line, since that information bears strongly on the validity of any model. It is of interest to note that Binder and McCarthy (1973) found the near-infrared limb-darkening of Saturn and Jupiter to be incompatible with isotropic-scattering atmospheres. The nature of pressure-induced line formation should permit certain simplifications in inhomogeneous scattering models. Because of the strong density
H2
ABSORPTION
ON JUPITER
TABLE
AND
403
SATURN
III
SOURCES OF ERROR
Size or nature of error
A. Absorption coefficients 1. Errors in laboratory measurements 2. Errors in determining coefficients from curves 3. Errors in calculating CC,for intermediate temperatures
Unknown +3% of a, at 4800cm-’ +2.5% of CC,. at 4800; *lo%
at 5000
B. Models (RLM) 1. Effect of including enhancement by 20% He 2. Effect of reducing ortho/pera ratio from 3.0 to 2.47 (equilibrium at 145K) 3. Effect of dropping isothermal assumption 4. Effect of using adiabatic T(h) 5. Errors in scale height (primarily due to T)
+0.55% of G(,for 195K -5% of ix, Roughly equivalent to using Q, for temperature lower by -1OK -8% of po, W +20% of pO, W if T E (100, 150K) for Saturn
C. Data reduction 1. 2. 3. 4.
Errors in estimating line profile due to noise Errors in estimating zero level Error in estimating Saturn ring contribution Estimation of 5100cm-’ continuum level a. Error in measured lunar continua b. Error in assumed lunar reflectivity c. Error in estimating planetary 6300cm-’ continuum d. Maximum possible instrumental effects 5. Error in estimating planetary intensity at 5100cm-’
dependence, one need not consider the low-lying cloud deck ; only the upper cloud should be important. Where holes appear in the clouds, no radiation will emerge from the dense atmosphere below. We expect radiation in the 4600-5100cm-1 region to arise only from within or just above clouds near the 0.5agt level. B. Jupiter The Jovian pressure-induced S(1) line spectrum is more problematic than that of Saturn. Both the RLM and HSM provide good fits to the data, but for widely differing temperatures, neither of which is expected. Extensive study of the forma-
+2% of 5025cm-’ intensity +2% of 5025om-’ intensity _t20% of maximum corm&ion ~20% of ratio 1(6300)/1(5100) <5% of ratio +5% of I(6300) ” ” ~28% of ratio I(6300)/1(6100) Upper/lower limits sre carried through to plots
tion of the H, quadrupole lines has shown the need for complex models incorporating vertical structure of the cloud layers (Hunt, 1973). In Fig. 5 we see that the best fit with the RLM occurs for a temperature near 250K. Temperatures below 200K produce a line that is too narrow. Expected cloudtop temperatures for Jupiter are in the range lOO-150K ; it is unlikely that radiation from lower, hotter regions could survive transit through the considerable abundance of H, below the NH, cloud deck. There is, in fact, direct evidence that the measured radiation near 5000 cm-’ comes entirely from the upper cloud level or above. Figure 1 shows the 5050cm-’ NH, band clearly. The appearance of that band
404
MARTIN ET AL.
is closely matched by laboratory spectra having an abundance of about lOcm-atm. This value is considerably lower than abundances derived from the 64508 band, which is probably formed in part below the upper cloud deck. Formation of both bands may be complicated by the condensation of ammonia. The small infrared abundance is however evidence that the 5000cm-’ radiation has not traversed the regions below the upper cloud. One means of explaining the fact that an RLM curve corresponding to too high a temperature matches the Jupiter data is to assume an error in the continuum level at 5100cm-‘. If the continuum were lower, the values of I,/I, would be higher, the line would have a more narrow profile and would be fit by lower-temperature models. Such a lowering of the continuum would result from a broad absorption near 5000 cm-l that does not occur at 6300 cm-’ where the continuum is derived. As we candidates for such an noted above, absorption are NH, frost and NH,SH in suspended solid form. We ruled out the possibility of either of these two substances in the case of Saturn because both compounds have spectral features near 3300 where the Saturn spectrum is cm-l undisturbed. In the case of Jupiter, however, the 3300cm-’ region is strongly absorbed by NH, gas. Consequently, we cannot eliminate the possibility that either of the solid constituents is absorbing near 5000cm-‘. A change in the 5000 cm-’ continuum of approximately 50% would be necessary to make the Jupiter profile match curves for 150K. We will not, however, assume that the continuum is lowered by such a large amount in the absence of specific information. We shall not derive a base-level density and H, abundance from the RLM analysis for Jupiter because of the unrealistically high temperature that gives the best fit. It is unfortunate that the RLM gives this result; the pressure-induced line is intrinsically a much better indicator of pressure abundance properties than are and collisionally broadened lines. The HSM profile that best fits the Jovian line (Fig. 9) has a temperature of 77K.
Existing models of the atmosphere do not show temperatures that low at any level, and certainly not near the levels where pressure-induced lines probably form. Because of the shape of the 130K curve, it is also unlikely that a change in the continuum level could make it agree with the data in detail. We conclude that it is apparently necessary to develop more realistic scattering models for the Jovian line analysis, and to observe the center-to-limb variation.
C. C’onclusions The S(I) line of the 1-O pressure-induced band of H, is readily observed in the spectra of Jupiter and Saturn. Its great width makes it possible to perform useful line-shape analysis even with much lower spectral resolution than was employed here (0.5 cm-‘). The line shape is very sensitive to the (density)2 x (path) product. As a result, it is likely that all detected radiation in the 4600-5100 cm-’ interval is coming from atmospheric levels in or above the upper cloud layers thought to exist on Jupiter and Saturn. The line shapes measured do not give consistency between the RLM or HSM and existing models and measurements of pressure and H, abundance. The reason for the disagreement is probably a lack of sophistication in the RLM and HSM. Although presumption of certain errors in the analysis can bring the line shapes into agreement regarding temperature with the results of previous investigations using the RLM, the densities and abundances then derived are inconsistent with pressure and abundance information from those and other studies. Therefore, the following efforts are necessary or desirable : 1. Observation of the center-to-limb variation of the pressure-induced H, line shape on Jupiter and Saturn. Spectral resolution for this task can be as low as lOcm-’ ; the measurements are possible with existing instrumentation. 2. Application of existing anisotropic and/or inhomogeneous scattering models to the problem of the H, pressure-induced
Hz ABSORPTION ON JUPITER AND SATURN absorption. Derivation of center-to-limb variations of line shape in the models. 3. Recalculation or measurement of accurate absorption coefficients for H,. 4. Observation of the spatial distribution of 5000cm-L radiation on the disk of Jupiter. Does the radiation indeed arise only from regions above the upper cloud layer? This experiment could best be performed with an imaging system similar to that used by Westphal et al. (1974) for 5 pm studies. 5. Study of the formation of the 5050 cm-’ band of ammonia. This absorption probably forms at the same levels as the H, line. High-resolution (0.1 cm-‘) spectra would be valuable. ACKNOWLEDGMENTS We thank Messrs. James Engel and Harry Chase and Dr. G. Wijntjes of Block Engineering/ Digilab, Inc. for their assistance with the interferometer system. We also thank Dr. Nancy D. Morrison for her unflagging assistance at the telescope when most of the observations were made. This work was supported by NASA Grant NGL 12-001-057.
REFERENCES BELTON, M. J. S., ANIJ SPINRAD, H. (1973). H, pressure-induced lines in the spectra of the major planets. Astro@y8. J. 185, 363-372. DE BEROH, C., VION, M., COMBES, M., LECACHEUX, J., AND M AILLARD, J. P. (1973). New infrared spectra of the Jovian planets from 12 000 to 4000cm-’ by Fourier transform spectroscopy. II. Study of Saturn in the 3vs CH, band. As&on. Astrop/&ys. 28, 457-466. DE BERQH, C., LECACHEUX, J., COMBES, M., AND MAILLARD, J. P. (1974). New infrared spectra of the Jovian planets from 12 000 to 4000cm-’ by Fourier transform spectroscopy. III. Firstovertone pressure-induced H, absorption in the atmospheres of Jupiter and Saturn.
Astron. Astrophys.
35, 333-337.
BERGSTRALH, J. (1973). Methane absorption in the Jovian atmosphere. II. Absorption line formation. Icam. 19, 390-418. BINDER, A. B., AND MCCARTHY, D. W. (1973). IR spectrophotometry of Jupiter and Saturn. A&on. J. 78, 939-950. CHAMBERLAIN, J. W. (1970). Behavior of absorption lines in a hazy planetary atmosphere. Aatrophy8. J. 159, 137-158.
405
CHANDRASEKHAR,S. (1950). Radiative Transfer. Clarendon Press, Oxford. CHISHOLM, D. A., AND WELSH, H. L. (1954). Induced infrared absorption in hydrogen and hydrogen-foreign gss mixtures at pressures up to 1500 atmospheres. Can. J. Phya. 32, 291-
311. CRUIKSHANK, D. P., AND BINDER, A. B. (1968). The infrared spectrum of Jupiter, 0.95-1.60 microns, with laboratory calibrations. Conana.
Lunar Planet. Lab. 6, 275-288. DANIELSON, R. E. (1966). The infrared spectrum of Jupiter. Astrophys. J. 143, 949-960. ENCRENAZ, TH., AND OWEN, T. (1973). New observations of the hydrogen quadrupole lines on Saturn and Uranus. Aetron. Astrophys. 28,
119-124. FINK, U., AND LARSON, H. (1971). Fourier spectroscopy at the Lunar and Planetary Laboratory at the University of Arizona. In Aspen International Conference on Fourier Spectroscopy (G. Vanasse et al., Eds.), pp. 453461. AFCRL-71-0019. HERZBERU, G. (1962). Spectroscopic evidence of molecular hydrogen in the atmospheres of Uranus and Neptune. Aetrophy8. J. 115,
337-340. VAN DE HULST, H. C. (1952). Scattering in the atmospheres
of the Earth
and planets.
In
The Atmospheres of the Earth and Planets, p. 102. Univ. of Chicago Press, Chicago, Ill. HUNT, G. E. (1973). Interpretation of hydrogen quadrupole and methane observations of Jupiter and the radiative properties of the visible clouds. ilfon. Not. Roy. Astron. Sot.
161,347-363. HUNT, G. E., AND BEROSTRALH, J. T. (1974). Structure of the visible Jovian clouds. Nature 249, 636-637. HUNT, J. L., AND WELSH, H. L. (1964). Analysis of the profile of the fundamental infrared band of hydrogen in pressure-induced absorption. Can. J. Phya. 42, 873-885. JOHNSON, H. L. (1970). The infrared spectra of Jupiter and Saturn at 1.2-4.2microns. Astrophya. J. 159, Ll-5. KIEFFER, H. H., AND SMYTHE, W. D. (1973). Unpublished spectra. MCCORD, T. B., AND JOHNSON, T. V. (1970). Lunar spectral reflectivity (0.30 to 2.60 microns) and implications for remote mineralogical analysis. Science 169, 855-857. PALLUCONI, F. D. (1972). The planet Saturn (1970). NASA SP-8091. SPINRAD, H. (1963). Pressure-induced dipole lines of molecular hydrogen in the spectra of Uranus and Neptune. Astrophys. J. 138,
1242-1245.
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TRAFTON, L. M. (1973). Saturn: A study of the 3v3 methane band. Astrop@9. J. 182,615~636. WATANABE, A. (1971). Pressure-induced infrared absorption of gaseous hydrogen and deuterium at low temperatures. III. Further analysis of the fundamental and first overtone bands of hydrogen. Can. J. Phys. 49, 1320-1326. WATANABE, A., AND WELSH, H. L. (1965). Pressure-induced absorption of gaseous hydrogen and deuterium at low temperatures. I. The integrated absorption coefficients. Can. J. Phys. 43, 818-828. WATANABE, A., AND WELSH, H. L. (1967). Pressure-induced infrared absorption of gas-
eous hydrogen and deuterium at low temperatures. II. Analysis of the band profiles for hydrogen. Can. J. Phys. 45,28iifb2871. WEIDENSCHILLINU, S. J., (1973).
Atmospheric
AND LEWIS,
5.
S.
and cloud structures of
the Jovian planets. Icarus 20, 465-476. WELEUI, H. infrared
L.
(1969).
spectrum
The of
pressure-induced
hydrogen
and
its
application to the study of planetary atmospheres. J. &rnoa. Sci. 26, 836-840. WESTPHAL, J. A., MATTHEWS, K., AND TERRILE, R. J. (1974). Five-micron Astrqhya.
pictures of Jupiter.
J. 188, Lll l-L1 12.
27, 391-406
(1976)
Pressure-Induced
Absorption
by Hz in the Atmospheres and Saturn
of Jupiter
TERRY Z. MARTIN,’ DALE P. CRUIKSHANK, CARL B. PILCHER, AND WILLIAM M. SINTON Institute
for Astronomy,
UnCversity of Hawaii, Honolulu, Hawaii, 96822
Received September 23, 1978 The S( 1) line of the pressure-induced fundamental band of H2 was identified and measured in the spectra of Saturn and Jupiter. This broad line at 4750cm-’ lies in a region free from telluric and planetary absorptions. It is about 99% absorbing in the core; the high-frequency wing extends to at least 5100om-‘. We compare the observed line shape to the predictions of both a reflecting-layer model (RLM) and a homogeneous scattering model (HSM). The RLM provides a good fit to the Saturn line profile for temperatures nesr 160K; the derived base-level density is 0.52 (+0.26, -0.17) smagat and the H, abundance is 25 (+10, -Q)km-smagat, assuming a scale height of 48km. The Jupiter line profile is fit by both the RLM and HSM, but for widely differing temperatures, neither of which seems probable. The precise fitting of the observed S(1) line profile to computed models depends critically on the determination of the true continuum level; difficulties encountered in finding the continuum, especially for Jupiter, are discussed. Derived RLM densities and abundances for both planets are substantially lower than those derived from RLM analyses of the H, quadrupole lines, the 3~ band of CH,, and from other sources.
The study of hydrogen absorptions in the atmospheres of Jupiter and Saturn is complicated by the particular nature of the absorptions and the complex atmospheric structure on both planets. For example, in measurements of the collision-narrowed quadrupole lines, high spectral, spatial, and temporal resolutions are required if a meaningful analysis is to be conducted (Hunt and Bergstralh, 1974). Pure rotational pressure-induced absorption by hydrogen is the dominant opacity source in the thermal infrared region of the spectra of Jupiter and Saturn, but its observation is hampered by telluric absorptions. Measurements of the pressure-induced vibrational bands in the near-infrared are similarly hampered by the strong methane absorptions occurring in major planet l Present address: Department of Geophysics and Space Physics, University of California., Los Angeles, California 90024. Copyright 0 1976 by Academic Press, Inc. All rights of reproduction in any form reserved. Printed in Great Britain
spectra in the 4000-12 OOOcm-’ region. Despite this difficulty, some studies of the pressure-induced vibrational bands have been made. The 3-O second overtone band was first noted in the spectra of Uranus and Neptune by Herzberg (1952), thereby establishing the presence of hydrogen on the major planets. A feature at the position of the 4-O band in the Uranus and Neptune spectra is probably not due to hydrogen (Belton and Spinrad, 1973). Neither the 3-O nor 4-O band is evident in the spectra of Jupiter or Saturn, presumably because of the lower abundances existing in their optically accessible atmospheres. The 2-O pressure-induced band in spectra of Jupiter and Saturn has recently been studied by de Bergh et al. (1974). This band is unfortunately coincident with a wide absorption feature of methane that prevents accurate determination of the hydrogen contribution. Upper limits were set to the abundance of H, that would be compatible with the observed absorption 391
392
MARTIN
on both planets, and an estimate was made of pressures in the region of line formation. The 1-O fundamental band was observed in low-resolution (X/AX N 40) spectra of Jupiter obtained by the balloon-borne Stratoscope II experiment (Danielson, 1966). A reflecting-layer model (RLM) analysis of the line shape led to a temperature estimate of about 250K, following the assumption of an H, abundance. However, the low resolution of these data makes quantitative interpretation somewhat uncertain. Cruikshank, Fink, and Larson noticed the probable appearance of the S(1) line of the H, fundamental band at 4800cm-’ in their spectra of Saturn obtained jointly with Kuiper (Fink and Larson, 1971). This absorption was also evident in Jupiter and Saturn spectra by Johnson (1970) but was not identified. Both these Saturn spectra showed contamination by radiation from the rings. Welsh (1969) has presented background material necessary to understand the behavior of the pressure-induced bands. Lines in the bands are broad ; their shapes are predictable, but the relative intensities of lines and the effects of the presence of foreign molecules depend in a complex way upon temperature. The S(1) “line” is comprised primarily of three components: the S,(l), S,(l)+ Q,(O), and X,(l)+Q,(l) lines. Here S denotes a transition in which the molecular rotational quantum number J changes by +2, Q denotes AJ = 0, the subscript is the change in vibrational quantum number, and the term in parentheses is the J value of the initial state. Two of the lines are double transitions-each molecule of the colliding pair contributes one of the noted transitions simultaneously. The absorption coefficient for the S(1) line also contains contributions from the high-frequency tails of other lines in the 1-O band (see Fig. 2 of Welsh, 1969). All component lines are intrinsically broad because of the very brief time during collisions when transitions may occur. We discuss in this paper the S(1) line in spectra of Jupiter and Saturn obtained at Mauna Kea Observatory. This material is
ET AL.
taken from the doctoral dissertation of one of us (TZM). DATA REDUCTION The spectra of Jupiter and Saturn in the interval 4000-6000cm-1 are presented in Fig. 1. Each spectrum has a resolution element of 0.5cm-’ and represents an average of several measurements. Saturn was observed in February and March 1974, and Jupiter in July 1974. Our spectra were obtained with a Digilab Michelson interferometer system mounted at the Cassegrain focus of the Mauna Kea 224~cm telescope. The spectrometer accepted a circular beam 18arcsec in diameter which covered Saturn’s entire disk and a small portion of the rings. For observations of Jupiter the aperture was centered on the planet’s disk. Absorption in the spectra between 4100 and 4700cm-’ is due to both methane and hydrogen, and in the case of Jupiter, ammonia. Methane absorption is also evident beyond 5500 cm-‘. The H, S( 1) line of interest occurs between 4650 and 5050 cm-’ ; telluric water vapor is responsible for most features in the 51005500 cm-’ range. The zero levels of these spectra for analytic purposes are defined by a straight line drawn between the intensity level at and the parts of the 4100-4200 cm-’ 5300cm-* water band where the zero can be well established. The 4100-4200cm-’ interval is totally absorbed by H, and CH,. Because these spectra are unapodized, some absorption features project well below the drawn zero level. We corrected the Saturn spectrum for ring radiation contamination in two ways. At the telescope, a reference aperture of the sky-subtracting foreoptics was positioned partly on the ring to compensate for the ring signal due to light in the main aperture. Guiding errors, however, led to a residual contamination of the disk spectrum. That contribution was subsequently removed as follows. Metha,ne has a strong Q-branch absorption at 4550cm-I, at which frequency the rings are relatively bright. It was assumed that
H,
4000
_LLL
ABSORPTION
ON JUPITER
AND SATURN
I
I
4200
4400
4600
J
Cb m-i
I
I
48co
4600
I
5400
5000'
5200
I
I
5600
5800
6000
FIG. 1. Spectra of Jupiter (top) and Saturn, 4000_6000cm- l. Resolution = 0 6cm-‘. absorption features due to planetary CH,, Jovian NH,, and telluric CO2 are indicated.
Selected
394
MARTIN
any apparent intensity level above zero at 4550cm-’ was in fact due to the rings. Knowing the shape of the ring spectrum measured with the same instrumentation, we subtracted at all frequencies of interest the appropriate fraction of the amount of contamination at 4550 cm-‘. The slope of the solar continuum and the instrumental response function were removed by dividing the data by lunar spectra. These ratios were then multiplied by the lunar spectral reflectivity determined from Apollo 11 soil samples (McCord and Johnson, 1970). In order to plot the planetary absorption as relative intensity (ly/llcontinuum), the continuum level must be determined at some frequency. The continuum is unfortunately not accessible near 5000 cm-’ ; the closest useful window in the methanedominated spectra is at 6300 cm-l. There, the Saturn spectrum resembles that of the moon in detail ; methane lines are weak and isolated. Jupiter’s spectrum at 6300 cm-’ shows considerable absorption from gaseous ammonia. However, that frequency remains the best at which to judge the continuum. If hydrogen itself is an opacity source at 6300cm-‘, then our derived H, line shapes will be in error. However, for reasonable temperatures, the H2 absorption coefficient falls off at such a rate that only about 5% absorption is expected at 6300cm-‘. From the continuum level at. 6300cm-‘, we then calculated the corresponding level at 5100 cm-’ according to 1,(5100) = Ic(6300) r=;]
Moon[=;I,,,,
where c denotes continuum and R is reflectivity. Clearly, if the planetary continuum is not actually flat we produce an erroneous value at 5100 cm-‘. No better method is evident, however. To investigate the shape of the background planetary absorption above 5000 cm-‘, we must account for the effects of telluric H,O, CO,, and planetary CH, and NH,. The comparison with lunar spectra serves to establish the small contribution
ET AL.
of telluric components between 5025 and 5075 cm-‘. We used the methane spectra of Cruikshank and Binder (1968) and our own ammonia spectra to assess the importance of these compounds. We thereby derived upper and lower limits to the values of relative intensity that occur in the 50005100cm-i region ; these limiting values have been retained in the analysis below. We determined the lower limit to the relative intensity by drawing the lowest curve that can be smoothly extended from the 4900cm-1 region through the signal level at 5020cm-’ and yet fall above the peaks of intensity that occur in the H,O band. Allowing for no continuum absorption by CH, at 5025 cm-l, we derive a 5100cm-’ lower limit of Iv/I, = 0.45 for Saturn and 0.26 for Jupiter. To obtain an upper limit, we estimated a depression of 5% by CH, at 5025cm-‘, 6% by NH, (for Jupiter), and then arbitrarily doubled these effects. A smooth curve drawn from the 4900cm-’ region through the high points at 5025cm-’ gives a rough value of Iv/I, = 0.63 at 5100cm-’ for Saturn and 0.36 for Jupiter. We may now estimate the percent absorption due to II, at 5100cm-‘. Allowing for 1.0% CH, absorption, the remaining component is 30 to 49% for Saturn. With an additional 11% absorption by NH, on Jupiter, the absorption attributed to H, is between 55 and 67% of the continuum level. Thus, the hydrogen line is apparently more extensive on Jupiter, a fact due either to greater abundance, higher temperature, or a different mode of line formation. The reduced planetary values of IV/I, are given in Table I. We now must establish that methane and ammonia cannot account for the broad absorption we attribute to hydrogen. In Fig. 2 we reproduce Cruikshank and Binder’s (1963) ambient-temperature spectra of various quantities of CH, in the range 4500-5250cm-*. Plotted with their data is our reduced Saturn profile, with the continuum level taken coincident with the top edge of the figure. Methane cannot match the extent and central maximum of the observed absorption. Indeed, at lower temperatures the methane R-branch ex-
H, ABSORPTION TABLE
ON JUPITER
I
RELATIVE INTENSITY OF JUPITER AND SATURN, 4600-5 lOOcm-’ Frequency cm-’ 4600 4625 4660 4675 4700 4725 4750 4775 4800 4825 4850 4875 4900 4925 4950 4975 5000 5025 5050 5075 5100
Jupiter
Saturn
IV/I= 0.009 0.02 0.024 0.02 0.015 0.011 0.011 0.015 0.023 0.0355 0.05 0.0685 0.0885 0.113/0.115 0.135/0.139 0.157/0.169 0.179/0.203 0.20/0.239 0.221/0.28 0.24/0.317 = 0.26/0.36
I”/& 0.0 0.0145 0.0375 0.0365 0.0225 0.008 0.005 0.017 0.0345 0.0555 0.078 0.106 0.145 0.19/0.193 0.23310.243 0.278/0.306 0.319/0.366 0.355/0.438 0.391/0.513 0.42410.579 = 0.4510.63
Lower limit/Upper
limit
tending from 4570 to 4700 cm-’ will change so that maximum absorption occurs at lower frequencies ; the higher frequency 5200
0 1.90
5100
I.95
5030
2.00
395
AND SATURN
lines weaken. Methane bands show fine structure that is not evident in the broad feature of the planetary spectra. Finally, several of the methane R-branch manifolds appear in our spectra between 4620 and 4660cm-L (Fig. 1). These would be t’otally indistinguishable because of saturation if methane were present in sufficient quantities to produce strong absorption above 4700cm-L. Ammonia has several strong bands which may affect the Jovian spectrum in the region of interest. For example, the u, + vq band at 5050cm-’ is clearly visible in Fig. 1. Laboratory spectra indicate that abundances of NH, which produce spectra resembling the Jovian 5050cm-’ band show very little absorption between 4600 and 4900cm-‘. Ammonia frost shows an absorption at 5000 cm-’ corresponding to the vj + v,, band, but the shape and location do not match the broad feature we observe. Ammonium hydrosulfide, a possible cloud component in the lower Jupiter and Saturn atmospheres (Weidenschilling and Lewis, 1973), has a sloping absorption similar in shape to the feature we discuss here (Kieffer and Smythe, 1973), but both NH, and NH,SH have strong bands at 3300 cm-l which are not evident in the Saturn spectrum. We conclude, therefore, that the only consistent explanation of the broad feature CM-’
4900
205
4800
4700
2 IO
46co
2.,5
4500
2 20/A
FIG. 2. Spectra of methane (Cruikshank and Binder, 1968) and the relative intensity of Saturn. The continuum level for Saturn is at the top edge of the figure. 14
396
MARTIN
that we observe is that it is due to the S( 1) pressure-induced line of H,. MODELS AND PREDICTED LINE SHAPES A. Reflecting-Layer
ET
AL.
an adiabatic density distribution, but these varied little from models based on the simpler isothermal density distribution which has the form
Model
In order to interpret the planetary spect.ra, we have computed the pressureinduced spectrum in the S(1) line for a number of model atmospheres using available H, absorption coefficients (Chisholm and Welsh, 1954 ; Hunt and Welsh, 1964; Watanabe and Welsh, 1965, 1967 ; Watanabe, 197 1). We first consider the simple reflecting-layer model (RLM) which has been widely applied to studies of Jupiter and Saturn with varying degrees of success. We characterize the gas by a given density and temperature at the reflecting level, and some distribution of density and temperat,ure with altitude. The intensity relat,ive to the continuum at frequency v for pressure-induced hydrogen absorption in a cell of length 1 is given by Iv/I, = exp (-CC”vp2 I), where p is the density in amagat (agt) and CL\, is the absorption coefficient in agt-*. The strong density dependence implies that the atmospheric layers near the reflecting surface will be heavily weighted. Consequently, the vertical temperature variation will not strongly affect the line shapes. The intrinsic temperat’ure dependence of the aV term is weak over the range of temperatures likely to occur in one scale height. Absorption coefficients are not available in the literature which adequately cover the expected range of important temperatures in the planetary atmospheres (loo15OK). VC’e have therefore calculated absorption coefficients for 130 and 150K, using the approach of Hunt and Welsh (1964). We estimate the absorption coefficient error to be +2.5% at 4800cm-‘, increasing to &100/o at 5000cm-‘. In the absence of the availability of H, data for other temperatures in numerical form, we took data from published curves; the numbers are given together with our calculated data in Table II. We have constructed models based upon
P(L) = poexp (-W),
where p. is the base level density, h is altitude above that level, and H = kT/pg is the scale height, with k being Boltzmann’s constant, T temperature, TVthe mean molecular weight, and g gravitational acceleration. The relative intensity in the vicinity of a pressure-induced line produced by two successive vertical traversals of the model atmosphere is given by Iv/I, = exp [ -va y 7) _f,” p’(h)dh] =
exp (-k
~Hpo2),
(1)
where 7 is an airmass factor having a minimum value of 2. If a more realistic adiabatic bemperature gradient is assumed, the relative intensity has precisely the same form as (1) above, but with the argument of the exponential different by 17% ; here the base-level temperatures are taken to be the same in both cases. Derived values of p. will differ by about So/. In the following discussion we assume the isothermal model for simplicity. The Jovian observations are matched to line profiles computed for the disk center, since the size of the interferometer entrance aperture limited the maximum included value of 7 to 2.22. The Saturn data were compared with a disk-integrated RLM profile, computed from the formula : tlvllc)Disk
=
2
’ exp I--va, po2 H/( 1 - G)“*] r dr s =” (1 - C) ewe - C2 Ei(-C), (2)
where r is a radial integration variable and C = va,po2 H. A computer program was used to calculate RLM profiles from (1) and (2) for va,rious temperatures. Families of line profiles corresponding to different abundances of hydrogen were computed. The results of this analysis are presented in Figs. 3-5 and Fig. 7. The first of these shows disk-integrated RLM profiles for various temperatures and a common value
H, ABSORPTION ONJUPITERAND TABLE
397
SATURN
II
ABSORPTION COEFFICIENTS” FOR NORMAL H,
Frequency cm-l 4600 4625 4650 4675 4700 4726 4760 4775 4800 4826 4850 4875 4900 4926 4950 4975 5000 5025 5050 5075 5100
20K
77K
95K
130R (do.)
150K (talc.)
196K
248K
300K
0.215 0.17 0.135 0.14 0.425 1.00 1.26 0.84 0.485 0.30 0.21 0.16 0.12Ei 0.095 0.08 0.065 0.055 0.05 0.05 0.053 0.06
0.20 0.185 0.215 0.30 0.478 0.61 0.658 0.54 0.37 0.265 0.20 0.15 0.11 0.085 0.069 0.056 0.047 0.039 0.034 0.03 0.026
0.314 0.27 0.262 0.29 0.352 0.415 0.441 0.40 0.342 0.28 0.209 0.15 0.112 0.085 0.06 0.043 0.031 0.023 0.017 0.011 0.006
0.345 0.325 0.348 0.422 0.644 0.634 0.62 0.625 0.422 0.333 0.265 0.214 0.176 0.148 0.126 0.109 0.095 0.084 0.075 0.067 0.061
0.369 0.346 0.372 0.44 0.646 0.618 0.604 0.62 0.426 0.343 0.277 0.227 0.188 0.158 0.136 0.118 0.103 0.091 0.082 0.073 0.067
0.368 0.389 0.43 0.619 0.582 0.608 0.582 0.611 0.426 0.359 0.296 0.25 0.21 0.183 0.158 0.138 0.12 0.104 0.089 0.077 0.066
0.443 0.44 0,477 0.49 0.518 0.625 0.613 0.49 0.44 0.39 0.36 0.31 0.273 0.244 0.219 0.198 0.18 0.162 0.144 0.127 0.112
0.435 0.43 0.45 0.485 0.63 0.666 0.672 0.55 0.495 0.436 0.38 0.336 0.295 0.265 0.235 0.21 0.185 0.17 0.152 0.14 0.13
a Units:amagat-2x 10pq.
RLM
Saturn
-----
FIG. 3. Reflecting-layer model S( 1) line profiles for poZH = 13.21unagt2 and indicated temperetures. The planetary profile is shown as the dashed line.
398
MA.RTIN ET AL.
4600
4700
4900
a00
CM
FIG. 4. Reflecting-layer amagats.
II00
4700
4100
-----
4900 CM
FIG. 5. Reflecting-layer
5000
model line profiles for 16OK, H = 5Okm, and indicated densities p,, in
Jupiter
4600
.’
-’
model line profiles for po2H = 18.2km@.
of po2H. The reduced Saturn profile is included, with upper and lower limits indicated for the far-wing intensity. The value po2 H = 13.2kmagt2 was chosen to give the best fit with the 15OK curve. The general depression of the intensity with increased temperat,ure is demonstrated. Figure 4 shows the RLM disk profiles for constant temperature and varying baselevel dens&y po. H is held constant at 50km. As the den&y changes from 0.4 to 0.7agt, the line profile rapidly becomes totally absorbing in the core. The wings also
so00
5100
and indicated temperatures.
demonstrat,e a strong density dependence. Clearly, if the Saturn line profile is roughly correct, base densities much different from 0.5agt are not permitted. The Jupiter spectrum is presented in Fig. 5 together with RLM center-of-disk profiles at three temperatures. The value of po2 H is chosen to give the best fit at 248K. Calculated profiles at the lower temperatures illustrated do not give satisfactory agreement with the data for any values of po2 H. This problem will be discussed later.
H2
ABSORPTION
ON JUPITER
B. Homogeneous Scattering Model The case of line formation in a hazy atmosphere has been treated by Chamberlain (1970), following the work of Chandrasekhar (1950), van de Hulst (1952), and others. We shall contrast the homogeneous scattering model (HSM) of Chamberlain with the RLM, although both of these are clearly physical oversimplifications. The HSM assumes a semi-infinite atmosphere ; isotropic scattering occurs with no redistribution of photon frequency. All parameters are independent of depth. The relative intensity is given by I,/& = & H(%
CL)HP,,
PO)/
3, H(G 14 H(%, PO), where the single-scattering albedo is & in the continuum and 6, at frequency v in the line, p. and p are, respectively, the direction cosines for rays entering and leaving the atmosphere, and H(&, p) are the Chandrasekhar (1950) H-functions for isotropic scattering. Jupiter and Saturn were near opposition during our observations, so we assumed p = pO. Employing the van de Hulst (1952) approximation to the H-functions H(G, P) = HP, cl)/{1 + /43(1 - 41”2), expression for relative intensity :
we get the
I,,Z, = 2
l/2
2
1 +- I-43(1 - a1 WC ( 1 $ /.43(1 - &,)]“2 I *
(3)
AND
399
SATURN
This formula was used to prepare model line profiles for differing values of p and integrated over the disk to provide profiles for comparison with the Saturn spectrum. The single-scattering albedo 8, is given by Gl, = (c$/u + l/f&)-‘, where u is the volume scattering coefficient. A computer program was used to calculate line profiles from (3), with various values of ay, CJ,and G,. For Saturn, since the whole disk was measured, the disk-integrated equation analogous to (3) was computed from the following expression : Vvl~chMc
=
J-L (4 0
2 55 where
P = [3( 1 - GJ]“2,
1
1 + P(l r2)“2 1 + &(l - G)“2
1rdr, 2
Q = [3(1 - t~J]l’~.
The result is (JvlZ&isk = 2(%&&)
{P/Q2)[In (1 + Q) - Q/(1 + &)I + W’/Q3)[-21n(l +&I +&-k&/(1 +&)I + (P2/Q4) [3ln(l + Q)+&'/2- 2Q - Q/U + Q,l>. The results of the HSM analysis are presented in Figs. 6-9. The upper part of
FIG. 6. Behavior of homogeneous-scattering model line profiles for 77 K, p = 1 agt. Upper scattering coefficient 0 = 5 x lo-‘cm-‘, various single-scattering elbedos d, indicated. 8, = 1, various crindicated.
: volume Lower:
400
MARTIN
ET
AL.
150 K
0.1.
0.6 -
0.4. 0.2 : ; : go., ’ 0.6-
HSM
0.4 *
0.1* 4600
4100
4.00
4900 CM
.’
FIG. 7. Center-to-limb variation in the RLM and HSM at 150K. ~_tindicated. Lower : HSM, o = 6 x lo-‘cm-‘, d, = 1, p = 1 apt.
0.1 .
0.6 .
5000
SlOO
Upper: RLM, direction cosines
HSM
Saturn
-----
: ; E 0.4 -
I’
,’ ,’
/ C
,,=-
z
FIG. 8. Homogeneous-scettering model line profiles (disk-integrated) u = 2 x lo-‘cm-‘. 130K: 0 = 2 x lO”cm-‘.
Fig. 6 displays the effect of varying the continuum singlescattering albedo at constant a for a temperature of 77K and a density of 1 agt. Fixing 9, and changing the volume scattering coefficient produces similar changes in line shape, shown in the lower part of Fig. 6. The center-to-limb variation of the RLM and HSM line profiles is shown in Fig. 7. In the former, absorption increases near the limb, exactly as it would with a change in base-level density. The HSM behavior is the opposite ; the relative intensity increases near the limb. Figure 7 emphasizes
for p = lagt,
3, = 1. 77K:
the diagnostic importance of observations of the center-to-limb variation on Jupiter and Saturn. In Figs. 8 and 9 the HSM line profiles which best match the Jupiter and Saturn spectra are plotted. In Fig. 8, curves for 77 and 130K and density 1 agt are shown. These temperatures bracket the expected range of Saturn temperatures near the cloud layers. Scattering coefficients were chosen to give a reasonably good overall fit to the planetary profile, although the choice is obviously somewhat arbitrary. The value of 3, is 1 .O ; curves with smaller
H, ABSORPTION
ON JUPITgR
401
AND SATURN
HSM
Jupiter
4600
,700
4800
-----
,900 CM
-’
5000
5100
FIG. 9. Homogeneous scattering model line profiles (center of disk) for p = 1. 77K: cm-‘, 6, = 0.94. 130K: o = 5 x 10-Scm-l, c3, = 0.5.
values will fall above the ones that are plotted. Figure 9 shows the Jupiter profile and the 77K and 13OK HSM curves. The 77K line gives a good match for p = 1, u = 2 x 10-7cm-‘, and G, = 0.94; the 130K curve does not give a good match for any value of these parameters. DISCUSSION
A. Xatum The RLM appears to give the best fit to the Saturn line profile. The HSM profiles are not sufficiently narrow even at 77K. The core of the line is more strongly absorbing than the HSM profiles (Fig. 8), so that allowance for errors in determination of the continuum level will not be able to produce agreement over more than a fraction of the observed line width. The RLM profiles produce a best fit for a temperature in the range 150--200K. That is higher than expected cloud-top temperatures for Saturn (Weidenschilling and Lewis, 1973), and higher than measured rotational temperatures that presumably refer to similar levels (Encrenaz and Owen, 1973; Trafton, 1973; the RLM was employed for both these investigations). However, we do not expect accurate temperatures from the H, line study, since the temperature-dependence of ay is small. Most likely, our result implies
CT= 2 x lo-’
that the continuum level used at 5100cm-’ was too high. A lower continuum gives a narrower line shape, which is then matched For lower-temperature models. by example, if the continuum is just 15% lower, an excellent fit is obtained using the 130K profile. Any continuous absorption that is stronger near 5100cm-’ than at 6300 cm-’ could thus help explain the RLM temperature discrepancy. Ammonia frost and ammonium hydrosulfide both have absorptiona near 5000cm-’ that could be partly responsible. However, the absence of the 3300cm-’ absorptions of these compounds in the Saturn spectrum has already been mentioned. The possibility exists that solar radiation scattered back from hotter (and hence lower) regions of the atmosphere contributes to the 4600-5100cm-L intensity. In such a case, the line profile might have a broader shape than has a line formed purely above the clouds. However, this possibility seems fairly unlikely in light of current multiple-layer cloud models for the planets. The abundances of H, existing below the upper cloud are sufficient to absorb totally any incoming radiation in the 4600-5100cm-’ range before it can again escape from the clear space between the upper and lower cloud decks. The argument is supported by Fig. 4 ; lower cloud level densities are probably at least 3agt for both planets, and the intercloud distance is well over a scale height. The
402
MARTIN
product p21 for the between-cloud region is then on the order of 200 kmagt2. If all the radiation observed is reflected from the upper cloud alone, and scattering effects are unimportant, we can derive a cloud-level density and above-cloud H, abundance from the RLM analysis. The parameter X = po2H is determined from the fitting of the 150K curves to the planetary spectrum. We may then solve for p,, and the H, abundance W = poH: po
=
GVV’2,
W = (XH)“2.
The scale height H is calculated, using g = 1000cmsec-2, p = 2.25amu (following Weidenschilling and Lewis for a solarcomposition atmosphere), and T = 13OK. Then H = 48km, and with X= 13.2km agt2, we have p. = 0.52agt,
W = 252kmagt.
The effects of errors in X and Hare reduced by the square-root dependence of p. and W. The best way to assess the total error in these quantities is to let all contributing errors fall so that the derived values are maximized. These are listed in Table III. We cannot, of course, assess errors arising from the assumption of a flat continuum between 5 100 and 6300 cm-‘. By contriving in all respects to maximize the derived po, we arrive at a value po(max) = 0.78agt using X = 22.5kmagt2 and H = 37km. In the other direction, we have X = 8.1 km agt2, H = 55km, and po(min) = 0.38agt. Applying the same principles to the derivation of W, we get finally base-level density : p. = 0.52 ?t:$amagat, H, abundance above reflecting layer : W = 25.2+i” km amagat. It should be emphasized that the values of p. and W derived depend little on the temperature used to select absorption Although one temperature coefficients. gives a best fit, the curves for other temperatures which match the planetary profile to some degree also yield densities and abundances near those quoted. These values are discrepant with existing atmospheric models and with conclusions about pressure and abundance derived
ET AL.
from RLM analyses of the methane 3v, band and the H, quadrupole lines. The models (Weidenschilling and Lewis, 1973 ; Palluconi, 1972) indicate densities near 4amagat at the level of NH, cloud formation. Measurements of the widths of collision-broadened lines have implied that pressures are around 2-3 atm (de Bergh et al., 1973) where the 3vJ lines are formed. Encrenaz and Owen (1973) derived an H, abundance of 76 f 20kmagt from the quadrupole lines. However, these have been found to exhibit considerable time variation for Jupiter (Hunt and Bergstralh, 1974). A recent study of the 2-O pressureinduced band (de Bergh et al., 1974) yielded upper limits for pressure and abundance of 0.6atm and 63kmatm, respectively, which are not inconsistent with our derived values. Clearly, we are presented with a problem of interpretation. If the models and most other RLMderived values are meaningful, then there should be no intensity evident in the Saturn spectrum between 4600 and 5100 cm-‘. The pressure-induced opacity of the expected abundances of H, would cause total absorption in that range. Since we do detect radiation at these frequencies and can measure the H, line shape, either the existing models and interpretation of measurements are in error, or the RLM is not adequate to explain line formation near 5000cm-‘. We have already seen that the HSM does not do well. The next logical step is to investigate the effects on the line shape of anisotropic scattering functions or vertical inhomogeneities. An inhomogeneous scattering model has been applied with success to Jupiter (Bergstralh, 1973 ; Hunt, 1973). At the same time, however, it is imperative to observe the center-to-limb variation of the 4750cm-’ line, since that information bears strongly on the validity of any model. It is of interest to note that Binder and McCarthy (1973) found the near-infrared limb-darkening of Saturn and Jupiter to be incompatible with isotropic-scattering atmospheres. The nature of pressure-induced line formation should permit certain simplifications in inhomogeneous scattering models. Because of the strong density
H2
ABSORPTION
ON JUPITER
TABLE
AND
403
SATURN
III
SOURCES OF ERROR
Size or nature of error
A. Absorption coefficients 1. Errors in laboratory measurements 2. Errors in determining coefficients from curves 3. Errors in calculating CC,for intermediate temperatures
Unknown +3% of a, at 4800cm-’ +2.5% of CC,. at 4800; *lo%
at 5000
B. Models (RLM) 1. Effect of including enhancement by 20% He 2. Effect of reducing ortho/pera ratio from 3.0 to 2.47 (equilibrium at 145K) 3. Effect of dropping isothermal assumption 4. Effect of using adiabatic T(h) 5. Errors in scale height (primarily due to T)
+0.55% of G(,for 195K -5% of ix, Roughly equivalent to using Q, for temperature lower by -1OK -8% of po, W +20% of pO, W if T E (100, 150K) for Saturn
C. Data reduction 1. 2. 3. 4.
Errors in estimating line profile due to noise Errors in estimating zero level Error in estimating Saturn ring contribution Estimation of 5100cm-’ continuum level a. Error in measured lunar continua b. Error in assumed lunar reflectivity c. Error in estimating planetary 6300cm-’ continuum d. Maximum possible instrumental effects 5. Error in estimating planetary intensity at 5100cm-’
dependence, one need not consider the low-lying cloud deck ; only the upper cloud should be important. Where holes appear in the clouds, no radiation will emerge from the dense atmosphere below. We expect radiation in the 4600-5100cm-1 region to arise only from within or just above clouds near the 0.5agt level. B. Jupiter The Jovian pressure-induced S(1) line spectrum is more problematic than that of Saturn. Both the RLM and HSM provide good fits to the data, but for widely differing temperatures, neither of which is expected. Extensive study of the forma-
+2% of 5025cm-’ intensity +2% of 5025om-’ intensity _t20% of maximum corm&ion ~20% of ratio 1(6300)/1(5100) <5% of ratio +5% of I(6300) ” ” ~28% of ratio I(6300)/1(6100) Upper/lower limits sre carried through to plots
tion of the H, quadrupole lines has shown the need for complex models incorporating vertical structure of the cloud layers (Hunt, 1973). In Fig. 5 we see that the best fit with the RLM occurs for a temperature near 250K. Temperatures below 200K produce a line that is too narrow. Expected cloudtop temperatures for Jupiter are in the range lOO-150K ; it is unlikely that radiation from lower, hotter regions could survive transit through the considerable abundance of H, below the NH, cloud deck. There is, in fact, direct evidence that the measured radiation near 5000 cm-’ comes entirely from the upper cloud level or above. Figure 1 shows the 5050cm-’ NH, band clearly. The appearance of that band
404
MARTIN ET AL.
is closely matched by laboratory spectra having an abundance of about lOcm-atm. This value is considerably lower than abundances derived from the 64508 band, which is probably formed in part below the upper cloud deck. Formation of both bands may be complicated by the condensation of ammonia. The small infrared abundance is however evidence that the 5000cm-’ radiation has not traversed the regions below the upper cloud. One means of explaining the fact that an RLM curve corresponding to too high a temperature matches the Jupiter data is to assume an error in the continuum level at 5100cm-‘. If the continuum were lower, the values of I,/I, would be higher, the line would have a more narrow profile and would be fit by lower-temperature models. Such a lowering of the continuum would result from a broad absorption near 5000 cm-l that does not occur at 6300 cm-’ where the continuum is derived. As we candidates for such an noted above, absorption are NH, frost and NH,SH in suspended solid form. We ruled out the possibility of either of these two substances in the case of Saturn because both compounds have spectral features near 3300 where the Saturn spectrum is cm-l undisturbed. In the case of Jupiter, however, the 3300cm-’ region is strongly absorbed by NH, gas. Consequently, we cannot eliminate the possibility that either of the solid constituents is absorbing near 5000cm-‘. A change in the 5000 cm-’ continuum of approximately 50% would be necessary to make the Jupiter profile match curves for 150K. We will not, however, assume that the continuum is lowered by such a large amount in the absence of specific information. We shall not derive a base-level density and H, abundance from the RLM analysis for Jupiter because of the unrealistically high temperature that gives the best fit. It is unfortunate that the RLM gives this result; the pressure-induced line is intrinsically a much better indicator of pressure abundance properties than are and collisionally broadened lines. The HSM profile that best fits the Jovian line (Fig. 9) has a temperature of 77K.
Existing models of the atmosphere do not show temperatures that low at any level, and certainly not near the levels where pressure-induced lines probably form. Because of the shape of the 130K curve, it is also unlikely that a change in the continuum level could make it agree with the data in detail. We conclude that it is apparently necessary to develop more realistic scattering models for the Jovian line analysis, and to observe the center-to-limb variation.
C. C’onclusions The S(I) line of the 1-O pressure-induced band of H, is readily observed in the spectra of Jupiter and Saturn. Its great width makes it possible to perform useful line-shape analysis even with much lower spectral resolution than was employed here (0.5 cm-‘). The line shape is very sensitive to the (density)2 x (path) product. As a result, it is likely that all detected radiation in the 4600-5100 cm-’ interval is coming from atmospheric levels in or above the upper cloud layers thought to exist on Jupiter and Saturn. The line shapes measured do not give consistency between the RLM or HSM and existing models and measurements of pressure and H, abundance. The reason for the disagreement is probably a lack of sophistication in the RLM and HSM. Although presumption of certain errors in the analysis can bring the line shapes into agreement regarding temperature with the results of previous investigations using the RLM, the densities and abundances then derived are inconsistent with pressure and abundance information from those and other studies. Therefore, the following efforts are necessary or desirable : 1. Observation of the center-to-limb variation of the pressure-induced H, line shape on Jupiter and Saturn. Spectral resolution for this task can be as low as lOcm-’ ; the measurements are possible with existing instrumentation. 2. Application of existing anisotropic and/or inhomogeneous scattering models to the problem of the H, pressure-induced
Hz ABSORPTION ON JUPITER AND SATURN absorption. Derivation of center-to-limb variations of line shape in the models. 3. Recalculation or measurement of accurate absorption coefficients for H,. 4. Observation of the spatial distribution of 5000cm-L radiation on the disk of Jupiter. Does the radiation indeed arise only from regions above the upper cloud layer? This experiment could best be performed with an imaging system similar to that used by Westphal et al. (1974) for 5 pm studies. 5. Study of the formation of the 5050 cm-’ band of ammonia. This absorption probably forms at the same levels as the H, line. High-resolution (0.1 cm-‘) spectra would be valuable. ACKNOWLEDGMENTS We thank Messrs. James Engel and Harry Chase and Dr. G. Wijntjes of Block Engineering/ Digilab, Inc. for their assistance with the interferometer system. We also thank Dr. Nancy D. Morrison for her unflagging assistance at the telescope when most of the observations were made. This work was supported by NASA Grant NGL 12-001-057.
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