The Galileo probe: how it has changed our understanding of Jupiter

New Astronomy Reviews 47 (2003) 1–51 www.elsevier.com / locate / newar

The Galileo probe: how it has changed our understanding of Jupiter Richard E. Young* NASA Ames Research Center, M /S 245 -3, Moffett Field, CA 94035, USA Accepted 3 December 2002 Communicated by J.J. Lissauer

Abstract The Galileo Mission to Jupiter, which arrived in December 1995, provided the first study by an orbiter, and the first in-situ sampling via an entry probe, of an outer planet atmosphere. The rationale for an entry probe is that, even from an orbiter, remote sensing of the Jovian atmosphere could not adequately retrieve the information desired. This paper provides a current summary of the most significant aspects of the data returned from the Galileo entry probe. As a result of the probe measurements, there has been a reassessment of our understanding of outer planet formation and evolution of the solar system. The primary scientific objective of the Galileo probe was to determine the composition of the Jovian atmosphere, which from remote sensing remained either very uncertain, or completely unknown, with respect to several key elements. The probe found that the global He mass fraction is significantly above the value reported from the Voyager Jupiter flybys but is slightly below the protosolar value, implying that there has been some settling of He to the deep Jovian interior. The probe He measurements have also led to a reevaluation of the Voyager He mass fraction for Saturn, which is now determined to be much closer to that of Jupiter. The elements C, N, S, Ar, Kr, Xe were all found to have global abundances approximately three times their respective solar abundances. This result has raised a number of fundamental issues with regard to properties of planetesimals and the solar nebula at the time of giant planet formation. Ne, on the other hand, was found to be highly depleted, probably as a result of it being carried along with helium as helium settles towards the deep interior. The global abundance of O was not obtained by the probe because of the influence of local processes at the probe entry site (PES), processes which depleted condensible species, in this case H 2 O, well below condensation levels. Other condensible species, namely NH 3 and H 2 S, were similarly affected but attained their deep equilibrium mixing ratios before the maximum depth sampled by the probe. Processes that might be capable of producing such effects on the condensibles are still under investigation. Measured isotopic ratios of noble gases and other heavy elements are solar, and (D1 3 He) / H is the same to within measurement uncertainties as in the local interstellar medium. No thick clouds were detected, and in particular no significant water cloud, but the PES location clearly affected the probe measurements of clouds. In fact, the probe data must be understood in the context of the location of the PES, which was within what is termed a 5 micron hot spot, a local clearing in the clouds that is bright near the 5 mm spectral region. The thermal structure at the PES was determined from approximately 1000 km above the 1 bar pressure level (10 29 bars) to 132 km below 1 bar (22 bars). The probe showed the atmosphere to have a generally sub-adiabatic temperature gradient (static stability) of ¯0.1 K km 21 to as deep as the probe made measurements. In the upper atmosphere the probe derived a maximum positive vertical temperature gradient of approximately 5 K km 21 , and maximum temperature of ¯900 K. The energy sources producing the warm upper atmosphere have yet to be completely identified. At first glance, Doppler tracking of the probe indicates that the long *Tel.: 11-650-604-5521; fax: 11-650-604-6779. E-mail address: [email protected]. (R.E. Young). 1387-6473 / 02 / $ – see front matter Published by Elsevier Science B.V. doi:10.1016 / S1387-6473(02)00272-5

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observed cloud level zonal winds extend to levels at least as deep as the probe made measurements. Zonal wind increases from ¯80 m s 21 at pressures less than a bar to about 180 m s 21 near 5 bars, and remains approximately constant with depth thereafter. However, there is a question as to whether the winds measured from probe tracking are representative of the general wind field, or are considerably influenced by localized winds associated with the PES. Published by Elsevier Science B.V. Keywords: Jupiter; Planets: probes

Contents 1. Introduction ............................................................................................................................................................................ 2. Scientific context and motivation for Galileo probe mission........................................................................................................ 2.1. Atmospheric composition ................................................................................................................................................ 2.2. Clouds ........................................................................................................................................................................... 2.3. Thermal structure............................................................................................................................................................ 2.4. Atmospheric circulation and dynamics ............................................................................................................................. 3. Galileo Probe mission background............................................................................................................................................ 3.1. Spacecraft, mission objectives, mission events .................................................................................................................. 3.2. Probe entry site ............................................................................................................................................................... 4. Galileo probe results................................................................................................................................................................ 4.1. Composition ................................................................................................................................................................... 4.1.1. Helium ............................................................................................................................................................... 4.1.2. Carbon and other heavy elements ......................................................................................................................... 4.1.3. Isotopic ratios ..................................................................................................................................................... 4.1.4. Condensible species ............................................................................................................................................ 4.2. Clouds ........................................................................................................................................................................... 4.3. Thermal structure............................................................................................................................................................ 4.4. Winds ............................................................................................................................................................................ 4.5. Energy balance ............................................................................................................................................................... 5. Summary of Galileo probe results and outstanding issues ........................................................................................................... Acknowledgements ...................................................................................................................................................................... References ..................................................................................................................................................................................

1. Introduction On December 7, 1995, the Galileo probe entered the atmosphere of Jupiter, making the first in-situ observations of an outer planet atmosphere. The data returned from the probe changed our understanding of Jupiter and the solar system in important ways. As expected, the probe data also raised new questions, some of which are still unresolved and will likely have to be addressed by future missions. Progress in understanding the origin and evolution of the solar system must involve intensive study of Jupiter. Jupiter is the largest planet in the solar system, containing almost 2.5 times the mass of all the other planets combined, and it has had some influence on almost every object in the solar system.

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A property of Jupiter of particular interest for this review is that the escape time for light elements from Jupiter is longer than the age of the Universe. Therefore, the composition of Jupiter provides important constraints on conditions in the solar nebula at the time of planet formation, on the formation process itself, and on the composition of planetesimals that have contributed to the formation of planets or that have impacted Jupiter over the age of the solar system. Aside from the importance Jupiter assumes in understanding the evolution of the solar system, as discussed below there are many reasons to study Jupiter as being representative of the outer planets as a group. Perhaps the most apparent difference between

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terrestrial and outer planets, aside from the obvious size and composition differences, is that all the outer planets except Uranus are observed to have an internal heat source comparable to the amount of solar energy they absorb (Uranus currently has essentially zero internal heat flux for unknown reasons), in clear contrast to the terrestrial planets, which radiate to space almost the exact amount of energy they absorb from the sun. This outward energy flux is due almost entirely to cooling as gravitational energy released during planet formation or later contraction escapes as heat. For Jupiter and Saturn there is likely a contribution from energy released as helium precipitates from the metallic hydrogen mantle into the deeper planet interior (see discussion in Section 2.1). The internal energy loss rate for each outer planet places a significant constraint on its formation and subsequent evolution, and may play a fundamental role in driving its meteorology. The dynamic meteorology of the outer planets is an especially intriguing but unsolved problem. In some ways the outer planets present simpler systems to study than the Earth. There is no topography or solid planet surface to affect atmospheric motions, no oceans or land–sea contrasts, and in the case of Jupiter small seasonal effects. Atmospheric features tend to be long lived. For example, the banded structure representing the larger than 100 m s 21 cloud level zonal (east–west) jet streams of Jupiter’s atmosphere, and the Great Red Spot, have been observed for literally centuries. Yet it was unclear as to whether the primary energy source of Jupiter’s atmospheric motions was direct solar radiative heating or Jupiter’s internal heat flux. Furthermore, it is known from Voyager spacecraft observations that zonal jets of even larger magnitude than exist on Jupiter exist on Saturn, Uranus, and Neptune. The puzzling aspect of this is that the latter three planets have much less energy available to drive wind systems than does Jupiter. All are considerably farther from the Sun than is Jupiter, and as mentioned above, all except Uranus have measured internal heat fluxes that are comparable to the absorbed solar radiative flux. In fact, the largest range in observed winds in the solar system occurs on Neptune, the farthest of these planets from the Sun. However Jupiter has the best observed

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meteorology of the outer planets, and provides the most accessible opportunity to understand the energetics and dynamical balances associated with outer planet wind systems. Cloud systems on the outer planets differ from those of the terrestrial planets in that multiple distinct global cloud layers that involve different condensible species are thought to exist. As on Earth, however, the effects of clouds on dynamic meteorology, and visa versa, are likely important in producing many of the observed features, such as the belt-zone structure in Jupiter’s atmosphere (the east–west banded structure associated with the zonal winds, bright regions being termed zones, dark regions termed belts). Determining the existence and composition of the deeper Jovian cloud layers, however, has been uncertain based on remote sensing observations, which for the most part have to look through local clearings in the upper cloud layer in limited spectral windows to sample the deeper atmosphere. The thermal structure of outer planet atmospheres at pressures greater than about 1 bar was anticipated to be adiabatic, i.e. temperature as a function of pressure follows an adiabatic relationship. Remote sensing observations, which could only sample thermal structure down to atmospheric depths somewhat below that corresponding to 1 bar pressure, seemed to confirm this expectation. To within the accuracy of these observations, the thermal structure at pressures greater than about 1 bar did appear adiabatic. However, in these atmospheric regions, even small deviations from an adiabatic thermal structure (referred to as static stability of the atmosphere) of magnitude below the precision of remote sensing measurements would be significant dynamically. Atmospheric static stability affects how internal heat is transported outward, the types of atmospheric dynamical modes that could exist, and affects even the rate of internal energy dissipation which is crucial for understanding the evolution and state of the Galilean moons of Jupiter. In this sense, the deep atmosphere thermal structure of outer planet atmospheres, including Jupiter’s, remained unknown from remote sensing. The atmospheric region of the Earth lying above about 80 km altitude, called the thermosphere, is characterized by an increase of temperature with height, with maximum temperatures approaching

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1000 K. Surprisingly, the thermospheres of the outer planets are also characterized by high temperatures of several hundred K to almost 1000 K. Although solar radiation is responsible for heating the Earth’s thermosphere, solar heating is 1–2 orders of magnitude too small to account for the high outer planet thermospheric temperatures observed, and correspondingly, outer planet thermospheric temperatures show no apparent correlation with solar distance. The energy source responsible for upper atmosphere heating on the outer planets is yet to be determined, but the Galileo probe may have provided important clues. Prior to the Galileo Mission four spacecraft had flown by Jupiter: Pioneer 10 and 11 in 1973 and 1974, respectively, and Voyager 1 and 2 in 1979. Each was incredibly successful and contributed tremendously to our knowledge of the Jovian system. There had also been many significant ground-based and airborne observations of Jupiter. However, basic questions remained concerning Jupiter itself which simply could not be adequately addressed by remote sensing. These then provided the motivation for an entry probe mission. For example, prior to the Galileo probe mission there was relatively little composition information below Jupiter’s cloud levels. Important elements for understanding solar system evolution and planet formation were either basically undetected (such as sulfur in the form H 2 S), or the available information concerning the abundance was contradictory (as was the case for oxygen represented by water abundance). Except for helium, Jovian noble gas abundances and their isotopic ratios, considered to be important tracers of planet and atmospheric evolution, were completely unknown. The helium abundance itself, as measured by Voyager, showed a surprising factor of three difference between Jupiter and Saturn. No unambiguous detection of two of the three hypothesized cloud layers had been accomplished, let alone a determination of the abundance profiles of the major condensibles. So basic a property of the atmospheric circulation as the vertical extent of the zonal jets below cloud levels was completely unknown. The static stability of Jupiter’s atmosphere at pressures greater than about 1 bar was undetermined. The upper atmosphere temperature structure was only sparsely sampled, and some of the

temperature determinations appeared contradictory. Thus, there was no shortage of important questions that could only be addressed by in-situ measurements. Making in-situ measurements of Jupiter’s atmosphere, however, was far easier visualized than implemented. The atmospheric entry into Jupiter is the most difficult in the solar system. Jupiter’s escape velocity is nearly 60 km s 21 , an entry speed no probe could survive. So steps had to be taken to minimize the relative entry speed, which essentially meant designing a trajectory in which the arrival vector was near the equator with the probe travelling in the same direction as the planet rotates. This subtracts about 12 km s 21 from the probe entry speed relative to the atmosphere. Even then, the probe must arrive in an extremely precise corridor in order to keep heating loads within survivable limits, which in the case of the Galileo probe meant keeping the probe within an entry corridor having boundaries 61.58 from a nominal flight path angle of 8.68. And for Galileo, this was accomplished with a ballistic trajectory of the probe after release from the orbiter about 80 million kilometers from Jupiter. In the best of circumstances, during entry a Jovian probe has to survive deceleration forces amounting to hundreds of Earth gE ’s (where gE is the gravitational acceleration at the Earth’s surface), while at the same time dissipating enormous heat input during the deceleration, the peak heating for the Galileo probe exceeding 3310 8 W m 22 (for reference, the radiant energy flux immediately above the Sun’s photosphere is 6.5310 7 W m 22 ). It is a testament to all the people involved with the Galileo mission that the Galileo probe survived such obstacles and functioned almost perfectly. The purpose of this paper is to give an up to date assessment of what the Galileo probe data have revealed about Jupiter as a planet and its place in the solar system, illuminated by what was known before the mission. A summary of outstanding issues still to be resolved concerning either interpretation of the probe data or implications of the data for Jupiter will also be presented. It is not the purpose here to present a detailed summary of the probe results. For thorough analyses of the data returned from each of the Galileo probe atmospheric instruments the reader is referred to the papers appearing in Journal of

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Geophysical Research ( Planets), v 103, E10, 1998. Since publication of that special Galileo probe issue, further data analysis papers have appeared regarding particular aspects of the data returned by the probe. The reader is also referred to those papers, which will be referenced at various points in the sections that follow. The structure of this review is the following. Section 2 discusses the scientific context of the Galileo probe mission, summarizing what was known about Jupiter and its atmosphere prior to the probe entry, emphasizing topics that the probe data would address. Section 3 presents particulars of the Galileo probe mission itself, including a discussion of the probe entry site (PES). Section 4 summarizes the most significant aspects of the probe data, and discusses implications of the data for understanding Jupiter. Section 5 identifies outstanding issues as yet unresolved regarding the probe results, and concerning Jupiter itself.

2. Scientific context and motivation for Galileo probe mission

2.1. Atmospheric composition Determining the composition of Jupiter’s atmosphere was one of the primary scientific goals of the Galileo probe. As will be seen from the following discussion, knowledge of Jupiter’s composition was quite limited, even after the Voyager mission. PreGalileo but post-Voyager summaries of the composition of the atmospheres of the outer planets, and in particular Jupiter, are given in Atreya (1986) and Gautier and Owen (1989). Here we summarize the highlights, adding a few updates that have appeared since these summaries appeared and pointing out gaps in what was known that provided motivation for a Jupiter probe mission. Jupiter’s atmosphere is taken to be that region of the planet for which hydrogen is in the molecular form, and for which the hydrogen–helium mixture behaves as a readily compressible gas, i.e., where density is <1 g cm 23 . Density reaches 1 g cm 23 at approximately 0.8 R J , where R J is Jupiter’s equatorial radius at 1 bar pressure (Marley, 1999). Starting near pressures of ¯4310 5 bars, molecular hydrogen

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begins to continuously dissociate, and completely dissociates to monatomic hydrogen by 3 Mbars (cf. Nellis, 2000, and references therein). Nellis also points out that the electrical conductivity of hydrogen reaches that of disordered metals near 1.4 Mbars, corresponding to about 0.9 R J . Thus, considering all the above factors, the Jovian ‘‘atmosphere’’ probably does not extend to pressures much beyond a few times 10 5 bars, a region less than 0.1 R J in extent, and the metallic hydrogen mantle occupies most of the volume of Jupiter. Assuming that elements of interest are not excluded from the mantle for physical chemical reasons (see discussion below), one would not expect significant compositional gradients between the atmosphere and metallic hydrogen mantle. There is evidently no sharp phase boundary (Nellis, 2000), and vertical mixing between the atmosphere and mantle should be efficient because of fluid motions generated by convective heat transport of the interior energy. A crude order of magnitude estimate of mixing times using convective mixing length theory (cf. Cox and Giuli, 1968), and based on numbers taken from the analysis of Goldreich and Nicholson (1977), indicates that mixing would occur over times short compared to the lifetime of Jupiter. Prior to the Galileo mission it had long been established that Jupiter consisted mostly of hydrogen and helium. In 1977 when the instrument complement for the Galileo probe was selected, it was believed that the helium abundance in the solar nebula, and by implication Jupiter, would be the same as the cosmological abundance (e.g. Lewis, 1974; Podolak and Cameron, 1975). Hence, an accurate measurement of the Jovian atmospheric helium abundance would be expected to provide constraints on conditions of the early universe. However, by the time Galileo was launched, it had become apparent that planetary processes might alter the helium abundance in outer planet atmospheres, and the atmospheric helium abundance was seen as placing constraints on outer planet formation and evolution rather than being directly relevant to cosmology. Depending on temperature and solubility of helium in metallic hydrogen, helium is expected to condense and settle towards the center of the planet (cf. Stevenson, 1982). This process, in combination with the vertical mixing expected by convective heat transport driven by Jupiter’s internal heat

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source, could therefore ultimately deplete helium in the Jovian atmosphere relative to the overall helium abundance for the planet as a whole. The question concerning the Galileo probe was what was the Jovian atmospheric helium abundance relative to the protosolar helium mass mixing ratio Y 5 0.27 (Bahcall et al., 2001). Based on Voyager remote sensing observations, it appeared that Jupiter and Saturn were quite different in this respect. The helium mass fraction derived for Jupiter’s atmosphere was significantly less than the protosolar value, implying some helium segregation. Voyager data analyses initially arrived at a Jovian atmospheric helium mass fraction of Y9 5 MHe / (MHe 1 MH2 ) 5 0.1960.05 or 0.2160.06, depending on whether purely spectroscopic data or spectroscopic in combination with radio occultation data were used (Gautier et al., 1981). The Voyager helium mass ratio was subsequently updated to account for the presence of CH 4 to give a value Y 5 MHe / (MHe 1 MH2 1 MCH4 ) 5 0.1860.04 (Conrath et al., 1984), implying significant helium segregation on Jupiter if Jupiter began with close to the protosolar abundance of helium. However, similar analyses for Saturn indicated that there had been considerably more helium segregation there, the helium abundance in Saturn’s atmosphere being smaller than the Voyager value for Jupiter by about a factor of three (Conrath et al., 1984). The implication was that Jupiter and Saturn had evolved differently. But as will be discussed later, new analysis of Voyager data at Saturn, motivated by the Galileo probe results at Jupiter, now indicates a much closer helium abundance for the atmospheres of the two planets. The elements heavier than helium provide additional information on planet formation and evolution. There are currently two classes of models for how Jupiter formed. The most widely accepted is the core accretion model, in which a core of order 10 Earth masses of rock and ices first forms by accretion of smaller bodies, followed by accretion of a hydrogen– helium envelope (Pollack et al., 1996; Wuchterl et al., 2000). The other suggested scenario is that Jupiter formed by gravitational collapse from the gas and dust cloud of the solar nebula via a localized gravitational instability (Boss, 1997, 1998, 2001). In the core accretion model, the heavy element inventory of the atmosphere would consist of ele-

ments outgassed from the core during accretional heating, elements delivered by impacting small bodies over the age of the solar system, and those initially present in the solar nebula gas (Lunine et al., 2000). In the gravitational instability model, atmospheric heavy elemental abundances in excess of solar values would have to be due entirely to impacting bodies, unless some process capable of concentrating elements or their compounds existed in the solar nebula. Within Jupiter, one might consider the possibility that elements heavier than helium could be relatively insoluble in hydrogen beyond some pressure in the interior, and therefore reside mostly in the outer molecular hydrogen envelope. In this case atmospheric abundances would not be representative of abundances for the planet as a whole. However, as will be discussed when the Galileo probe measurements are presented, Saturn seems to have about the same or greater atmospheric abundance of carbon relative to H 2 as does Jupiter (carbon in the atmospheres of Jupiter and Saturn is mostly in the form CH 4 which does not condense on either planet), in conflict with what would be expected on the basis of the interior exclusion hypothesis. This hypothesis would predict that Saturn should have a considerably smaller, rather than similar or larger, atmospheric heavy element enrichment than Jupiter. For example, if carbon enrichment in the atmosphere was due primarily to carbon being excluded from interior regions having pressures greater than that corresponding to the pressure at which hydrogen becomes significantly metallic, about 1.4 Mbars (Nellis, 2000), Saturn would have smaller enrichment of carbon than Jupiter because the 1.4 Mbar surface occurs at a smaller fractional radius from the center of the planet than on Jupiter. A solar abundance of heavy elements in Jupiter corresponds to about 6 M % (Earth masses). Current interior models for Jupiter, using mass, radius, and gravitational moments as constraints, together with the best available composition information including Galileo probe measurements, suggest that the total Jovian heavy element inventory is between 11 and 42 M % (Guillot et al., 1997; Guillot, 1999). One of the principal outputs from the interior models is a computed mass of the Jovian dense core, which as discussed above, is a crucial quantity for understanding how Jupiter formed. Currently, even using the

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Galileo probe measurements, the interior models are consistent with a core mass between 0 and 14 M % because of uncertainty in the hydrogen–helium equation of state (Guillot, 1999). It is not possible to identify at present specifically how Jupiter obtained its present inventory of heavy elements. However, the above inferred heavy elemental abundances provide constraints, if not on identification of the delivery process itself, then on particular aspects of it. For example, if all heavy elements in excess of solar abundances were delivered by small impacting bodies, as would almost certainly be the case if Jupiter formed by localized gravitational instability in the solar nebula, then tens of Earth masses must have been delivered. Depending on the efficiency of the delivery process, this would require perhaps hundreds to thousands of Earth masses of small bodies to be present in the outer solar nebula. (For reference, the efficiency for Kuiper Belt objects to impact Jupiter under present conditions is about 1% (Levison and Duncan, 1997).) Another example where elemental abundances can be used as a constraint on the delivery process involves very volatile elements such as N 2 . Such cases will be discussed in detail below, but suffice it to state here that such abundances place constraints on the temperatures, and hence origins, of the bodies delivering heavy elements to Jupiter. Of particular interest regarding heavy elements are the abundances of C, N, O, and S, because after hydrogen and apart from noble gases, these are the most abundant volatile elements in the solar system (cf. Anders and Grevesse, 1989)1 . Their most common compounds in the outer planets are gases at temperatures and pressures accessible to observation. At the same time, these same compounds are principal constituents of outer planet clouds when they condense. In the case of Jupiter, the major com1

There are a number of newer values for solar abundances of some elements, but we have chosen in this paper to use Anders and Grevesse (1989) as the reference because it provides a common baseline for comparison with previous works. As an example of the variation of recent quoted abundances of solar elements, between 1998 and 2001 the best estimate of the solar C / H ratio varied from 3.3310 24 (Grevesse and Sauval, 1998) to 3.9310 24 (Holweger, 2001). However, a very recent determination of the C / H ratio based on three dimensional modeling of the solar photosphere is (2.4560.24)310 24 (Prieto et al., 2002).

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pounds associated with C, N, O, and S are CH 4 , NH 3 , H 2 O, and H 2 S, all of which except CH 4 either condense directly or participate chemically to form components of clouds in the Jovian atmosphere. On Uranus and Neptune CH 4 also condenses. As pointed out by Gautier and Owen (1989), carbon is the most useful standard with regard to heavy element abundances in the atmospheres of Jupiter and Saturn because it does not condense on either planet, and should therefore be well mixed throughout at least the molecular hydrogen region of each planet. As the abundances of the condensible compounds measured by the Galileo probe are discussed later, the significance of this point will become apparent. Post-Voyager determinations of the C / H ratio in Jupiter’s atmosphere ranged from 2.0 to 3.6 times solar (Bjoraker et al., 1986a and references therein). The solar C / H ratio used for these comparisons was taken as 4.7310 24 (Lambert, 1978), but Anders and Grevesse (1989) give a value of 3.6310 24 . As noted by Hunten et al. (1986), NH 3 should exhibit considerable variability with altitude in Jupiter’s atmosphere because of photochemistry in the upper atmosphere, cloud condensation at temperatures below about 145 K, dissolution in any water clouds at depths of several bars, and chemical reaction with H 2 S to form an expected NH 4 SH cloud deck near 2 bars. The deep mixing ratio of ammonia is of special interest, because it presumably represents the mixing ratio of N for Jupiter as a whole. As discussed by Owen et al. (1997), the bulk C / N ratio for Jupiter is an indicator of the source of planetesimals that delivered extra solar abundances of C, S, N, and O. A Jovian C / N ratio greater than solar indicates that such planetesimals formed in a region that was too warm to allow retention of substances as volatile as N 2 , thought to be the primary form of N in interstellar clouds (Womack et al., 1992; van Dishoeck et al., 1993). Owen et al. (1997) point out that most of the comets now in the Oort cloud are believed to have formed in the Uranus–Neptune region of the solar nebula where temperatures appear to have been too warm for forming ices to retain N 2 . Dynamical models also suggest that many Oort cloud comets formed in the Jupiter–Saturn region (cf. Malhotra et al., 2000, and references therein), implying even warmer temperatures. The pre-Galileo retrieved ammonia deep mixing ratios were con-

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sistent with a Jovian C / N about two times solar, and were therefore consistent with heavy element enriching bodies coming from objects similar to those now making up the Oort cloud. The most recent pre-Galileo NH 3 vertical profile was derived by Carlson et al. (1993) using Voyager IRIS spectra in the far infrared in combination with radiative and thermochemical models. They inferred a deep ( p$2 bars) tropospheric mixing ratio (NH 3 / H 2 ) of 4.5310 24 , or very close to twice solar, significantly higher than the mixing ratio of approximately 2.5310 24 deduced at smaller pressures by Marten et al. (1980), Bjoraker et al. (1986a), de Pater and Massie (1985), and de Pater (1986)2 . In order to produce the much smaller (and highly depleted relative to solar) abundances at pressures between about 0.6 and 2 bars determined by Marten et al. (1980), Kunde et al. (1982), and de Pater and Massie (1985), Carlson et al.’s model required super-solar H 2 S abundances. At pressures ,0.5 bars, the NH 3 profile is subsaturated (Kunde et al., 1982; Orton et al., 1982; de Pater and Massie, 1985; de Pater, 1986; Carlson et al., 1993). The ensemble of NH 3 observations at different spectral wavelengths seemed to give NH 3 profiles in fair agreement. Marten et al. (1980), de Pater and Massie (1985), and de Pater (1986) used radio wavelengths, whereas Bjoraker et al. (1986a), Knacke et al. (1982), and Orton et al. (1982) used infrared wavelengths, and Kunde et al. (1982) and Marten et al. (1981) used Voyager IRIS spectra ranging from the near to far infrared. Gierasch et al. (1986) showed that ammonia abundance is correlated with cloud opacity, and de Pater (1986) concluded that latitudinal variations in ammonia abundance associated with belts and zones extend down to the 2 bar level. Prior to Galileo, H 2 S had not been detected, with only highly depleted upper limits near 1 bar existing from high resolution spectra at 2.7 mm (Larson et al., 1984). As mentioned previously, a greater than solar abundance of H 2 S was required in the models of Carlson et al. (1993) near and below 2 bars in order 2

Papers published after 1989 generally used updated solar abundances given in Anders and Grevesse (1989) rather than those given in Cameron (1982) to quote planetary elemental abundances relative to solar.

to produce the highly reduced NH 3 abundances between 0.6 and 2.0 bars reported by Marten et al. (1980), Kunde et al. (1982), and de Pater and Massie (1985). However, a 10 times solar abundance led to difficulties in fitting the IRIS 5 mm spectra. Therefore, there were weak constraints on the upper limit of the H 2 S abundance in the deep atmosphere. However, Carlson et al. pointed out that the amount of H 2 S above solar abundance required to match the NH 3 vertical profile depended on both the chemistry and microphysics of how NH 3 and H 2 S interact to form clouds under Jovian conditions, processes that are poorly understood at present. The water abundance was, and as we shall see when probe results are discussed, to some extent remains, one of the more confusing aspects of Jupiter’s composition. On a global scale water, i.e. O, is expected to be enriched above solar abundance to at least the same extent as the other heavy volatiles. It would be difficult to identify circumstances such that Jupiter is enriched in carbon, sulfur, and nitrogen, but not oxygen. Surprisingly, inversions of Voyager IRIS 5 mm spectra (the spectral region used to sample deeper atmospheric levels through local clearings in the clouds, known as 5 mm hot spots because they are bright near the 5 mm spectral region) produced varying results, yielding mixing ratios near 4 bars varying from less than or about 0.02 times the solar mixing ratio (relative to H 2 ) (Kunde et al., 1982; Drossart and Encrenaz, 1982; Bjoraker et al., 1986b; Lellouch et al., 1989), to at least twice the solar mixing ratio (Carlson et al., 1992, 1993). Airborne observations indicated highly depleted values relative to solar abundance (Bjoraker et al., 1986b). On the other hand, an interpretation of features as gravity waves which propagated outward from the impact sites of the fragments of comet Shoemaker–Levy 9 implied an abundance of water in the deep atmosphere about 10 times solar at the comet impact site near 458 S latitude (Ingersoll and Kanamori, 1995). The above discussion illustrates how meager or uncertain pre-Galileo knowledge was concerning Jovian abundances of C, N, S, and O. But even less was known regarding Jovian noble gas abundances. No noble gases other than helium had been detected, although it was expected they would be present. Noble gases, along with N 2 and carbon in the form

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of CO 2 , are thought to be key diagnostic volatiles for tracing atmospheric evolution for the terrestrial planets (Pepin, 1989). However, terrestrial planet noble gas inventories and their isotopic ratios are turning out to be a rather complex problem, involving perhaps several distinct sources as well as loss by atmospheric escape (see review by Zahnle, 1993). The outer planets, and in particular Jupiter, probably present simpler systems to understand in terms of noble gases and their isotopes. Atmospheric escape, principally hydrodynamic escape which can elementally and isotopically fractionate noble gases on the terrestrial planets, is not thought to have been a significant factor on Jupiter. For example, considering reasonable sets of parameter values, Xe fractionation would be less than 1 part in 10 4 (K. Zahnle, personal communication). Of the several possible sources for providing terrestrial planets with noble gases (cf. Zahnle, 1993), accretion of nebular gases and impacts of cold comets would be expected to be the dominant sources for the outer planets. Outer planet noble gas abundances and their isotopic ratios, particularly those of Jupiter and Saturn, should then provide a baseline for understanding noble gas sources for the terrestrial planets, as well as provide constraints on the source region for planetesimals delivering volatiles to Jupiter and Saturn. Prior to Galileo, no elemental isotopic ratios were known for the Jovian atmosphere other than estimates of the D/ H ratio. The D/ H ratio is a very useful tracer of terrestrial planet atmospheric evolution, especially with regard to the history of water (cf. Donahue et al., 1997; Owen and Bar Nun, 2000; Morbidelli et al., 2000). As discussed by Niemann et al. (1998), D/ H on Jupiter would be expected to be that prevailing in the local interstellar medium at the time of solar system formation, because Jupiter’s hydrogen and helium should have been obtained predominantly from the solar nebula during formation of the Sun. However, if Jupiter acquired on the order of 10–20 times solar abundance of water, as implied by the SL-9 gravity wave model mentioned previously (Ingersoll and Kanamori, 1995), and as suggested by some models of the delivery of heavy elements to Jupiter (see later discussion in Section 4.1.2), the Jovian D/ H ratio could be significantly affected. If the protosolar D/ H is known, and the D/ H in water delivered to Jupiter is known, then the

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extent to which the Jovian D/ H exceeds the protosolar value can yield a measure of the amount of water delivered to Jupiter. A reasonable assumption regarding incoming water might be that the D/ H in the water is the same as the D/ H in cometary water, which has been measured in several comets (cf. Lunine et al., 2000, and references therein). Another measure of the amount of water delivered to Jupiter could be the (D1 3 He) / H ratio. As pointed out in Niemann et al. (1998), if the Jovian D/ H ratio is the protosolar value, a fundamental baseline when studying D/ H in other bodies, then one would also expect that the Jovian (D1 3 He) / H would be the same as that currently in the local interstellar medium (LISM) if D is steadily converted to 3 He in stellar interiors (cf. Geiss et al., 2002, and references therein, for discussion of the production of 3 He). The apparent constancy of (D1 3 He) / H between the protosolar cloud and the LISM (Geiss et al., 2002) supports this view. Thus, the extent to which the Jovian (D1 3 He) / H differs from that in the LISM is a measure of the excess water delivered to Jupiter. Conceptually then, determination of the Jovian (D1 3 He) / H, together with 3 He / H, (or equivalently 3 He / 4 He and 4 He / H), can be used in a comparison with the current value of (D1 3 He) / H in the LISM to place an upper limit on the water abundance of Jupiter. This possibility will be discussed quantitatively when the probe results for isotopic ratios are discussed in Section 4.1.3. A recent summary of the pertinent isotopic ratios in the protosolar cloud and LISM is given in Geiss et al. (2002). Along similar lines of reasoning, the measurement of 3 He itself can be used to derive the protosolar D/ H ratio by subtracting the Jovian 3 He / 4 He ratio from that in the solar wind. The rationale is that the Jovian 3 He / 4 He ratio should be the original value in the solar nebula before nuclear burning in the Sun converts D to He. The method of course depends on determining the correct current solar wind value of 3 He / 4 He. Gautier and Morel (1997) obtained their estimate of the protosolar value of D/ H, (3.060.17)310 25 , using this procedure. The most recent pre-Galileo computation of the Jovian D/ H ratio by Carlson et al. (1993) from CH 3 D/ CH 4 using Voyager IRIS spectra between 1050–1200 cm 21 and 2050–2350 cm 21 yielded a D/ H ratio of (3.660.5)310 25 . Previous values,

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adjusted for differences in CH 4 abundance, had ranged from 1.6310 25 (Bjoraker et al., 1986a) to approximately 3310 25 (Knacke et al., 1982; Kunde et al., 1982). A D/ H of (2.260.5)310 25 was reported from the short-wavelength spectrometer aboard the Infrared Space Observatory (ISO) (Encrenaz et al., 1996). If the Jovian D/ H value is the value corresponding to the protosolar value, then according to the most recent protosolar value given in Geiss et al. (2002), the Jovian D/ H should be (2.060.4)310 25 .

2.2. Clouds Based on the assumptions of solar composition, thermochemical equilibrium, and an adiabatic temperature–pressure relation for the atmosphere below about the 0.2 bar pressure level, three cloud layers in the Jovian atmosphere accessible to the Galileo probe were predicted (Lewis, 1969a,b; Weidenschilling and Lewis, 1973; Atreya and Romani, 1985): an upper NH 3 ice cloud with cloud base at approximately 0.5–0.6 bars, a middle NH 4 SH ice cloud with cloud base at 1.5–2 bars, and a lower H 2 O ice cloud with cloud base at 4–5 bars. For a solar abundance of condensible volatiles, Atreya and Romani (1985) show that an even deeper NH 3 –H 2 O solution cloud is not expected on Jupiter, but is expected on Saturn. However, for water abundances three time solar, such a deeper cloud would be expected on Jupiter (Romani et al., 1995; de Pater et al., 2001). No direct identification of any of these cloud layers has been made, even by the Galileo probe (see later discussion). However, good evidence based on the vertical profile of the NH 3 gas abundance and the existence of aerosols in the region where NH 3 would be expected to condense strongly suggest that the composition of the upper visible cloud is NH 3 ice mixed with a chromophore (cf. West et al., 1986). The major controversy concerning clouds prior to the Galileo mission was whether there was a thick water cloud located in the region near 5 bars. The existence of a water cloud is of course directly related to the abundance of water in the Jovian atmosphere. Much of the information involved in the controversy consisted of spectra taken through 5-mm hot spots, which are local regions of clearing in the

upper cloud (see discussion of 5-mm hot spots in Section 3.2). Summarizing all the evidence, West et al. (1985, 1986) concluded that there was no evidence for a deep H 2 O cloud, and that the atmosphere below about 2 bars is essentially clear down to 6 bars or deeper. This was consistent with the observations discussed previously regarding H 2 O that indicated highly depleted amounts of water relative to solar abundance. However, analysis of Voyager IRIS 5 mm spectra led Carlson et al. (1992, 1993, 1994) to conclude that H 2 O abundance was at least twice solar, and there existed thick water clouds between 5 and 6 bars. This divergent set of conclusions formed the situation at the time of the Galileo probe entry. Since then it has been suggested that there was a calibration problem with the Voyager IRIS data in the spectral region having wavelengths shorter than 5 mm (Roos-Serote et al., 1999), such that in order to fit the IRIS data an additional opacity somewhere between 3 and 8 bars is required. This may explain the discrepancy between the conclusions of Carlson et al. and the previous works. In any event, it will be evident when we discuss results from the Galileo probe that although the controversy as to the existence of a thick water cloud as sensed within 5 mm hot spots has been laid to rest, questions regarding the global abundance of water continue to persist. (Global here means averaged over the Jovian atmosphere as a whole.) We will not go into details here regarding Jovian cloud physics, cloud particle sizes and their distribution, etc. The interested reader is referred to the comprehensive review article by West et al. (1986), and a brief summary of the pre-Galileo probe situation given in Young (1998).

2.3. Thermal structure The thermal structure of Jupiter is one of the fundamental properties of the planet that affects almost every other Jovian phenomenon, be it a phenomenon in the atmosphere or in the interior. Here, discussion of thermal structure will be limited to T( p), where T is temperature and p is pressure. Horizontal temperature variations are important, especially for understanding Jupiter’s atmospheric dynamics, but since the Galileo probe could only directly measure T( p) at one location, that will be

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emphasized. It is well established that near and somewhat above cloud levels, horizontal temperature contrasts are small, being only a few degrees (Ingersoll et al., 1976; Gierasch et al., 1986). As will be discussed later, although not directly measured by the probe, the probe can provide some meaningful constraints on processes that generate horizontal temperature structure. In the following discussion we shall occasionally refer to various regions of the atmosphere using terms such as troposphere, tropopause, stratosphere, and thermosphere. By and large, an atmosphere’s troposphere is a region in which T( p) is an increasing function of p (decreasing function of altitude), and is a region where atmospheric motions considerably affect the globally averaged (here meaning averaged over horizontal position) thermal structure. Therefore, the troposphere is also a region that is usually well mixed in terms of chemical species. The stratosphere lies immediately above the troposphere, and can be considered to be that region where the globally averaged T( p) is approximately constant with height. The vertical temperature structure there is dominated by radiative processes. The tropopause is the boundary, not always precisely defined, between troposphere and stratosphere. The thermosphere is a region lying above the stratosphere in which T( p) is a decreasing function of p (increasing function of altitude). Generally, thermospheric temperatures can greatly exceed temperatures in the upper troposphere. As discussed in Section 1, the energy source producing the high thermospheric temperatures has been identified for the Earth, but has not yet been well established for the outer planets. In the case of the Earth, the thermosphere is separated from the stratosphere by a region called the mesosphere, a region for which T( p) is an increasing function of p (decreasing function of altitude). However, there is not a region analogous to the mesosphere in the outer planet atmospheres; each of their thermospheres lies directly above the stratosphere. The generally prevailing view prior to Galileo was that at pressures exceeding approximately 0.5 bar, T( p) would correspond to an adiabatic relationship. This expectation was based on radiative–convective models of the atmosphere (e.g., Appleby and Hogan, 1984), and the fact that Jupiter had a sizable globally averaged internal heat flux of about 5 W m 22 (Hanel

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et al., 1981). Computations by Stevenson and Salpeter (1977) indicated that neither conduction nor radiative transport would likely be capable of transporting this amount of internal energy in the molecular hydrogen region of Jupiter. Assuming that the heat flux is transported by convection, convective mixing length theory (cf. Cox and Giuli, 1968) applied to Jupiter indicates that deviations from an adiabatic thermal structure would be quite small, u0.01%. Observationally, temperature versus depth had been deduced from stellar and spacecraft occultations for the upper atmosphere (see discussion below), and in the stratosphere and troposphere between approximately 0.001–1 bar from radio occultation of the Voyager spacecraft (Lindal et al., 1981). Inversion of Voyager IRIS spectra also produced vertical temperature profiles between 0.003 and 0.5 bar at various horizontal locations, however the vertical resolution was relatively coarse, on the order of a scale height (Flasar et al., 1981). The key features of the Voyager radio occultation temperature profiles (Lindal et al., 1981) taken near the equator (ingress at 128 S, egress at 08) and between 0.001 bar and 1 bar are: temperature at 1 bar of 165 K with an uncertainty of 3%; an adiabatic temperature gradient at 1 bar to within an uncertainty of 60.2 K km 21 ; a tropopause located at approximately 140 mbar; generally increasing temperature above the tropopause but with significant oscillations with height; and significant temperature contrasts above the tropopause between the ingress and egress profiles. Thus, both energy transport considerations and the best observations available of the thermal structure indicated that at pressures greater than a few tenths of a bar, T( p) was adiabatic. To the extent T( p) is adiabatic, measurement of T( p) in the atmosphere fixes the thermal profile for the interior. However, utilizing updated determinations of the radiative opacities of hydrogen and helium, Guillot et al. (1994) reexamined how energy could be transported in the outer planets, and concluded that the updated hydrogen and helium opacities, combined with the opacities of CH 4 , NH 3 , and H 2 O, were insufficient to ensure convection everywhere in the molecular hydrogen region. Their results for Jupiter indicated a radiative transport zone between about 1200 K and 3000 K (pressure |3000 bar). Including additional

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opacities of alkali metals may, however, be sufficient to ensure convection throughout the region (Guillot, personal communication). Thus at present, there seems to be no conclusive indication from radiative energy transport modeling that the thermal structure in the molecular hydrogen region of Jupiter at pressures greater than about 0.5 bar is anything but adiabatic. On the other hand, atmospheric static stability, measured here as the difference between the actual lapse rate of temperature and the adiabatic lapse rate in K km 21 , is such a critical parameter for determining the dynamical behavior of both the atmosphere and interior of Jupiter that even small positive values would be significant. The general dynamical behavior of the atmosphere in terms of wave modes and wind fields depends on whether the temperature structure is statically stable or neutrally stable (adiabatic T( p)). Furthermore, Jupiter’s energy dissipation, which is directly connected with such phenomena as the intense volcanism of Jupiter’s satellite Io and the atmosphere’s response to gravitational tides raised by the Jovian moons, could be intimately connected with the static stability (Ioannou and Lindzen, 1993, 1994). Positive static stability, i.e., a sub-adiabatic temperature lapse rate, represents stability of the atmosphere to vertical overturning or mixing. Molecular weight gradients, which can be established by condensation of trace constituents, can also contribute to the stability of the atmosphere against vertical overturning. Generally these contributions are small, or are zero if considering well mixed regions at depths below the clouds. Although energy transport considerations and spacecraft remote sensing of Jupiter’s atmosphere seem to indicate an adiabatic thermal structure much below the tropopause, there are implications from observed dynamical features in the Jovian atmosphere that at least to pressures of several bars, there must be a small, but dynamically significant, positive static stability of the order of tenths K km 21 , a value beneath the ability of the Voyager radio occultation measurements to detect. Features interpreted in terms of atmospheric waves have been observed in the Jovian atmosphere. Oscillations in Voyager radio occultation temperature profiles (Lindal et al., 1981) and wave like features seen at cloud levels have been used by Allison (1990) to deduce properties of

equatorially trapped planetary scale waves. Flasar and Gierasch (1986) identified mesoscale wave like features at the equator in Voyager images which they interpreted in terms of trapped inertia-gravity waves. These studies imply a static stability below the visible clouds which would be consistent with such waves, indicating that below the ammonia cloud deck the atmosphere should have a positive static stability to a depth of at least 20 km (about a scale height). A statically stable atmosphere in this region is consistent with the conclusion of Allison and Del Genio (1995), who deduced that the atmosphere had to have a Richardson number Ri ¯ 1.5 based on latitude variation of the mean zonal winds, where Ri is given by

S

g dT Ri ; ] ] 1 G T dz

DYS]dudz D

2

(1)

g is the Jovian gravitational acceleration taken at the 1 bar level, u is zonal wind velocity, z is height, and G is the adiabatic temperature lapse rate, which for a perfect gas is just g /c p , where c p is the specific heat 21 at constant pressure. G is approximately 2 K km in the Jovian atmosphere. In light of the apparent difference between what radiative equilibrium and internal energy transport models and atmospheric dynamical models would indicate regarding static stability near and below cloud levels, there was a clear need for very accurate in-situ measurement of the Jovian atmospheric thermal structure. Several pressure scale heights above the troposphere the upper atmosphere of Jupiter, and in fact each of those of the outer planets, exhibits a considerable rise in temperature above temperatures in the troposphere or stratosphere, such that all reach temperatures exceeding 400 K, with the Jovian upper atmosphere reaching about 1000 K (cf. Yelle et al., 1996, and references therein). As pointed out in Yelle et al., large upper atmospheric temperatures are an outstanding problem in our understanding of outer planet atmospheres. Thermal models based on absorption of solar EUV radiation and thermal conduction, which work well for explaining Earth’s high thermospheric temperatures, predict temperature increases in the upper atmospheres of the outer planets of only about 10 K or less (Strobel and Smith, 1973), an increase at least an order of magnitude too

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small. At present there is no adequate explanation of the hot outer planet thermospheres. In some ways the outer planet thermospheric heating problem is analogous to the classical problem of heating the solar corona. In both situations there is a steep rise in temperature in a tenuous gas above the bulk of the parent body and its atmosphere. In all likelihood, both situations involve transport of energy upwards from the lower atmosphere. When the Galileo probe thermal structure data are discussed, some suggested explanations will be described, but none has yet been shown to be sufficient. Fig. 1 summarizes observations of Jupiter’s near equatorial upper atmosphere thermal structure prior to the Galileo probe entry. It can be seen that there is a considerable difference, several hundred K, between the temperature profiles due to Festou et al. (1981) and Yelle et al. (1996). Each analysis is based on analyses of stellar occultation data taken by the ultraviolet spectrometer (UVS) aboard the Voyager 1 spacecraft, but Festou et al. and Yelle et al. used different boundary condition assumptions about the altitude at which the maximum temperature is reached. Hubbard et al. (1995) determined a temperature of 176612 K at 1.8 mbar near 88 N latitude from the occultation of the star SAO 78505 by

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Jupiter. The semi-empirical thermal profile of Yelle et al. passes through the temperature point derived by Liu and Dalgarno (1996) from analysis of H 2 emissions in the Jovian UV dayglow, and is essentially consistent with all the observations except the uppermost temperature–pressure point of Festou et al. The temperature derived by Marten et al. (1994) is based on analysis of emissions which yield temperature but rely on ionospheric model predictions to determine pressure level, hence the large pressure uncertainty associated with that point. The solar occultation experiment by the UVS on Voyager established that at pressures of |10 211 bar, T5 11006200 K (Atreya et al., 1979; Festou et al., 1981; McConnell et al., 1982). Thus, there was little doubt that temperatures become large in the upper atmosphere near the base of the exosphere (the region where the mean free path for molecular collisions is approximately equal to a density scale height). The question was how rapidly did T increase between the relatively cool atmosphere near 10 26 bar and the hot upper region. The profile of temperature rise is directly related to the source of heating of the upper atmosphere, and hence is a critical parameter when trying to identify the source of upper atmosphere heating. The dispa-

Fig. 1. Summary of observational data and semi-empirical profiles of temperature in Jupiter’s upper atmosphere, adapted from Yelle et al. (1996). The observations and profiles are discussed in the text.

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rate profiles of Festou et al. and Yelle et al. implied quite different energy deposition in the upper atmosphere. Here again was a situation that required direct in-situ measurements of the atmospheric thermal structure.

2.4. Atmospheric circulation and dynamics The banded appearance of Jupiter’s disc was first described in 1630 (Rogers, 1995), and amazingly has apparently changed very little since then. We now know that the banded structure, in the form of belts (dark regions) and zones (bright regions), reflects a cloud level pattern of zonal winds alternating in direction and speed as a function of latitude, and that this entire system has been remarkably steady over the entire time the planet has been observed. The meridional (latitudinal) profile of zonal winds has been determined by tracking visible cloud features in what are presumably the NH 3 clouds. The best information has come from the Hubble Space Telescope (HST) (Beebe et al., 1996) and Voyager spacecraft (e.g. Ingersoll et al., 1981; Limaye et al., 1982; Limaye, 1986). Fig. 2 illustrates the zonal winds inferred from tracking of cloud features observed by Voyager (Limaye, 1986), superimposed on part of an image taken by the Cassini spacecraft on October 31, 2000 as it approached Jupiter on its way to Saturn. The zonally (longitudinally) averaged zonal winds are the dominant component of the circulation that has been observed. Except for certain features such as the Great Red Spot and some bright ovals, deviations from the zonal average, termed eddies, are generally relatively small, the order of 10 m s 21 (Ingersoll et al., 1981). Reviews of outer planet atmospheric dynamics can be found in Ingersoll (1990), Gierasch and Conrath (1993), and references therein. While the dependence of the cloud level mean zonal winds on latitude had been well established prior to the Galileo mission, there was no direct observational data as to how the winds behaved vertically. In particular, so basic a property of the winds as their vertical extent below the clouds was unknown. Voyager cloud images taken in different spectral filters implied vertical variations in the zonal ˜ et al., 1990; Limaye, 1986), but winds (Magalhaes these were difficult to make quantitative. A coarse

Fig. 2. Jovian zonal wind field inferred from tracking of cloud features observed by Voyager (Limaye, 1986) superimposed on part of a Jupiter cloud top image taken by the Cassini spacecraft on October 31, 2000 as it approached Jupiter. Historically, zonal jets tend to be located at the boundaries between belts and zones, with westward jets (negative zonal wind) at the poleward edges of belts and eastward jets (positive zonal wind) at the poleward edges of zones. Even though the time span between the time the Cassini image was taken and the time that Voyager measured the winds is over 20 years, the above correlation is still apparent.

estimate of the vertical gradient of the zonal wind (termed vertical wind shear) above cloud levels indicated that the mean zonal wind decreased with height at pressure levels of 0.15 and 0.27 bars (Gierasch et al., 1986). The rate of decrease would predict disappearance of the zonal winds 3–4 scale heights above the tropopause. These sorts of studies used the thermal wind equation, which relates horizontal variations of temperature on a constant pressure surface to the vertical gradient (vertical shear) of the wind. In the following discussion horizontal variations in molecular weight will be neglected because below cloud levels it is expected that atmospheric constituents would be well mixed. (As will be discussed when probe results are presented, there are local regions where this may not be true, but on a global scale one would expect the atmosphere to be well mixed below condensation levels.) On a rotating planet the atmospheric equations of motion can be simplified if the Coriolis force domi-

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nates inertial forces. A measure of their relative importance is given by the Rossby number, Ro, where Ro ; u /(2V L) V is the planet angular rotation rate, and L is a typical horizontal length scale of the atmospheric circulation. If Ro < 1, then inertial forces are small compared to the Coriolis force. If Ro < 1, and atmospheric dissipation is small and can be neglected, a good approximation to the atmospheric meridional (latitudinal) equation of motion for a rotating planet is ≠u R ≠T ]] 5 ]]] ] ≠ ln( p) 2aV sin u ≠u

(2)

where R is the gas constant, u is latitude, and a is the planet radius. The meridional derivative of T is taken at constant pressure. Except within a few degrees of the equator, Eq. (2), termed the thermal wind equation, is a good approximation for Jupiter, and probably all the outer planets. By Eq. (2), the smallness of meridional temperature contrasts on the outer planets (cf. Fig. 5 of Ingersoll, 1990) implies small vertical variation in u. On this basis, Ingersoll (1990), and works cited therein) concluded that mean zonal winds extend deep into the atmospheres of the outer planets. The weakness in the argument, as Ingersoll pointed out, is that horizontal temperature variations were known at only a few pressure levels, all at pressures ,1 bar. Horizontal temperature contrasts at pressures exceeding 1 bar were, and still are, unknown. The expectation, however, is that thermal differences in the deeper regions would be even smaller than near cloud levels. All in all, based on the thermal wind relation, expectations were that the zonal winds extended well below the cloud features used to track the winds, not only in Jupiter’s atmosphere but on the other outer planets as well. One can turn the problem around. If the vertical shear ≠u / ≠ ln( p) can be measured as a function of pressure, information is then obtained from Eq. (2) about horizontal temperature contrasts as a function of depth. Assuming that at a particular site the dominant winds are the zonally averaged winds, which should be a good approximation outside obvious features such as the Great Red Spot, various bright ovals, hot spots, and convective plumes, then one vertical profile at the site provides insights into processes producing meridional thermal differences.

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(The probe was intended to avoid such features as hot spots, but as we will discuss, it did not.) For example, if vertical shear is confined to upper cloud levels where solar heating is significant, the implication is that solar heating predominantly maintains large scale temperature or density differences on constant pressure surfaces. According to the radiative model of Hunten et al. (1980), solar heating should not penetrate much below 2 bars. If vertical shear exists below approximately 2 bars but is found only above condensation levels where clouds form, latent heat release would be implicated as a contributing driving mechanism of the shear. For solar or larger abundance of water in the Jovian atmosphere, water would dominate latent heat release in terms of the fractional change in temperature produced (cf. tables in Gierasch and Conrath, 1985), the temperature perturbation occurring where water condenses at approximately 5 bars or so. Anywhere vertical shear exists, the atmosphere must be stably stratified, assuming that what is being measured is representative of average conditions. It is well known that if the Richardson number Ri is ,0.25 (cf. Eq. (1)), Kelvin–Helmholtz, or shear, instabilities quickly develop and eliminate the shear. Therefore, if vertical shear persists to well below cloud levels, the implication would be that the atmosphere there is stably stratified, i.e., has positive static stability. This would imply that, at least at the levels sampled, internal heat is transported outward by radiation rather than by convection, because convective heat transport would drive the interior to an adiabatic state which has neutral stability. Furthermore, as discussed by Ingersoll (1990), an adiabatic state has constant temperature on constant pressure surfaces. Hence the existence of a measureable zonal wind vertical shear deep in the atmosphere implies convection does not dominate atmospheric motions there. Non-convective heat transport of Jupiter’s internal energy much below cloud levels would be a surprise, but if the water abundance is highly depleted, near temperatures of 400 K the radiative models of Guillot et al. (1994) show there would be stable stratification and radiative energy transport. Thus, a single vertical profile of zonal wind has the potential of not only giving the depth dependence of the winds at the entry site, it might also yield information on important physical

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and energy conversion and transport processes. One of the primary Galileo probe objectives was to measure the zonal wind as the probe descended. Gierasch and Conrath (1993) point out that the exact nature of the forcing of large scale flow on Jupiter may not be important in producing the structure of the winds observed, for example the latitudinal spacing of the belts and zones. It may be that Jupiter’s size and rotation rate, coupled with perhaps just a few other parameters, are the controlling influences no matter how the overall flow is forced. Nevertheless, from a general point of view it is of interest to know whether the ultimate driving of the Jovian winds, and for that matter winds of the other outer planets which have even more intense wind systems, is solar heating or internal heat flux. One indication is the depth to which the winds extend. Deep flows are consistent with, but do not necessarily prove that, internal heat flux is the driver.

3. Galileo Probe mission background

3.1. Spacecraft, mission objectives, mission events The Galileo spacecraft consisted of two main components: an atmospheric entry probe and a planetary orbiter. Together they provided the first in-situ sampling of an outer planet atmosphere via the entry probe, and the first extended study of an outer planet, its satellites, and magnetosphere from an orbiter. Comprehensive descriptions of the Galileo Mission, including science objectives, instruments, and trajectory, have been given by papers

appearing in Space Science Reviews, 60, Nos. 1–4, 1992. The Galileo probe was designed and built by Hughes Space and Communications Group under contract to NASA Ames Research Center. The Jet Propulsion Laboratory designed and built the Galileo orbiter and had overall responsibility for mission design and operations. The German government was a partner in the mission through its provision of the orbiter propulsion subsystem and two science experiments. Scientists from six nations participated in the mission. The principal scientific objectives of the mission were: (1) to study the composition, structure, and dynamics of the Jovian atmosphere, to a depth having a pressure of at least 10 bars; (2) to study the composition, geology, surface morphology, and interior structure of the planet-sized Galilean moons Io, Europa, Ganymede, and Callisto; and (3) to determine the structure of, and characterize the plasma physics in, the vast Jovian magnetosphere. Objective (1) was primarily associated with the probe, while the last two were primary objectives of the orbiter. However, the orbiter together with ground based platforms provided key sets of observations that helped to establish the global context of the single atmospheric vertical profile measured by the probe. In this review we will concentrate on probe aspects of the mission and those other observations directly relevant to interpreting probe data. The instrument suite aboard the probe, together with the principal science objective of each instrument, are given in Table 1. Detailed descriptions of each instrument and its scientific objectives can be

Table 1 Galileo probe scientific payload Investigation

Purpose

Atmospheric Structure Instrument (ASI) Energetic Particle Instrument (EPI) Lightning / Radio Emission Detector (LRD) Helium Abundance Detector (HAD) Nephelometer (NEP) Net Flux Radiometer (NFR) Neutral Mass Spectrometer (NMS) Doppler Wind Experiment (DWE) Radio Scintillation Experiment (RSE) Ground Based Doppler Wind* (GDWE)

Jovian atmosphere thermal structure Jovian inner radiation belt parameters Jovian atmosphere optical / radio lightning properties Jovian atmosphere relative helium abundance Jovian atmosphere cloud parameters, location Solar / planetary radiative flux in Jovian atmosphere Jovian atmosphere composition, 2–150 amu Zonal wind profile in Jovian atmosphere Turbulence, signal absorption in Jovian atmosphere Zonal wind profile in Jovian atmosphere

* Not an originally selected Galileo investigation, but endorsed by the Galileo Project Steering Group.

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found in the Space Science Reviews articles referenced above, and the pre-launch review article by Hunten et al. (1986). All instruments functioned well throughout the mission. The Galileo spacecraft arrived at Jupiter on December 7, 1995, after having been launched on October 18, 1989 by the space shuttle Atlantis from Kennedy Space Center. The original launch date for the mission was to have been January, 1982, but various delays, many not associated with the spacecraft, led to a target launch date of May, 1986. However, the Challenger shuttle accident occurred in January, 1986, delaying Galileo’s launch for yet another 3 years. The launch delays meant that the trajectory to Jupiter had to be completely redesigned because of Shuttle payload, safety, and launch limitations. The ingenious trajectory redesign resulted in a trajectory that included a close flyby of Venus (February, 1990) and two close flybys of Earth (December, 1990 and 1992) for gravity assists to gain sufficient energy to reach Jupiter. By serendipity this trajectory allowed the first ever flyby reconnaissance of two asteroids (Belton et al., 1994, 1996) and the direct viewing of the impacts with Jupiter of the kilometer-sized fragments of comet Shoemaker–Levy 9 (Orton et al., 1995). The probe was separated from the orbiter on July 13, 1995, and continued on a ballistic trajectory to Jupiter. At the time of separation, both the orbiter and probe were about 80 million kilometers from Jupiter, and 5 months away from arrival. The probe plunged into the atmosphere of Jupiter on December 7, 1995, 14:04 PST, at a latitude of 6.58 N, surviving the most difficult atmospheric entry ever attempted in terms of relative atmospheric entry speed. Table 2

17

lists probe entry parameters. The values are based on the final trajectory reconstruction up to the point of entering the Jovian atmosphere (Bell, 1996). All probe entry parameters were close to nominal and within specifications. The probe returned data to the overflying orbiter for just over an hour, until it reached a pressure level in the atmosphere near 22 bars. On the same date, after having received and recorded the probe data, the orbiter was placed in orbit about Jupiter, and is still returning data about the Jovian system as of November, 2002. Table 3 gives an event time line for the probe entry and descent. Time zero is taken as the start of descent mode operations on the probe, descent mode being basically that portion of the probe mission during which the probe made direct atmospheric measurements. Prior to descent mode, the thermal structure of the upper Jovian atmosphere was indirectly measured by accelerometers which monitored the deceleration of the probe. From trajectory reconstruction using the measured probe deceleration, atmospheric density could be determined. Then assuming a hydrostatic atmosphere, pressure could be deduced, followed by temperature using the perfect gas equation of state and a model of upper atmosphere composition. Enroute to Jupiter the spacecraft suffered a major anomaly on April 11, 1991, when the high gain antenna (HGA) failed to deploy properly. Until this point, the HGA had been folded against a central mast and was protected by sun shields so that heat loads on the HGA would be within tolerances. Despite repeated attempts to solve the problem, the HGA was not functional for the rest of the mission.

Table 2 Probe entry parameters. Longitude is referenced to 1985 IAU convention (System III, see Section 3.2) Parameter Entry time (450 km above 1 bar pressure level), UTC Entry speed relative to atmosphere, km s 21 Initial relative flight path angle, deg Planetocentric latitude, deg Longitude, deg

Target

Requirement

Achieved

1s in achieved

22:04:26

63 min

22:04:44

3s

47.4 28.60 6.57 N 5.02 W

,47.8 (27.2)–(210.0) 66.6 (but not within 61) –

47.4

–

28.41 6.53 N

0.04 0.01

4.88 W

0.07

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18

Table 3 Event time table. Descent time is measured from beginning of probe descent operation. (See also text.) Descent time,

Pressure,

min 2187* 2145* 2106* 268*

bars *** *** *** ***

22.8 21.82 0.0 0.058 0.079

7310 28 0.007 0.4 0.4 0.41

0.23 1.0 3.0

Event LRD/ EPI data acquired at 5 R J LRD/ EPI data acquired at 4 R J LRD/ EPI data acquired at 3 R J LRD/ EPI data acquired at 2 R J EPI continues until just after entry Entry (450 km above 1 bar pressure level) Peak deceleration (228 g) Start of descent mode operation Pilot chute deployed Aft heat cover separated Main parachute deployed Start of direct atmospheric measurements Orbiter locks on to probe telemetry signal Begin first NMS sampling of ambient atmosphere composition Passage of 1 bar pressure level End of first NMS sampling of ambient atmosphere composition Begin second NMS sampling of ambient atmosphere composition Passage of 10 bar pressure level End of second NMS sampling of ambient atmosphere composition Begin third NMS sampling of ambient atmosphere composition End of probe signal

0.43 0.54 0.88

3.6 15.4

1.0 3.8

29.8

8.6

33.4 38.5

10.0 12.1

46.2

15.6

58.6

21.7

* To the nearest minute.

Mission communications were carried out using the low gain broad beam antenna, with a subsequent reduction in communication rate by a factor of 1000. It is to the credit of all the mission teams that in spite of this severe problem, Galileo met almost all its fundamental science objectives. The only effect on return of probe data to Earth, which data was initially stored onboard the orbiter after probe transmission, was to extend the time interval of data transmission to Earth to the order of a few months rather than almost real time. But all probe data were successfully received on the ground. There was, however, one significant impact on the overall probe mission of the failure of the HGA combined with other spacecraft complications. Planned high resolution approach images to be taken by the orbiter of the probe entry site just before probe entry were canceled in the mission sequence. This had the potential of leaving unknown the particular cloud

and atmospheric features through which the probe descended, thereby leaving the global context of the probe measurements uncertain. As it turned out and as described below, ground based measurements were able to identify the atmospheric feature into which the probe entered, and this identification has been crucial for trying to understand and interpret various aspects of the probe data.

3.2. Probe entry site It was impossible to choose a-priori the cloud or atmospheric features through which the probe would descend. A precise entry latitude and longitude (System III)3 were of course targeted, and as Table 2 3

System III is a coordinate system for the definition of Jovian longitude. In System III, the rotation of Jupiter is 9 h 55 m 29.7 s, and is the rotation of the Jovian interior as deduced by radio observations (cf. Rogers, 1995).

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illustrates, achieved almost exactly. However, even as long lived as many of Jupiter’s cloud features are, there was virtually no possibility of targeting a particular atmospheric feature at the time of probe entry and descent. The rationale for choosing the particular probe entry coordinates involved many factors (cf. D’Amario et al., 1992), three of the most important constraints being: (a) because of the Jovian ring discovered by Voyager, the probe had to be at least 3.8 R J from the center of Jupiter when it crossed the equatorial plane as it approached the planet; (b) the entry speed relative to the atmosphere could not be greater than 47.8 km s 21 ; and (c) the probe must enter the atmosphere within an entry angle from the horizontal between 2 7.2 to 2 10.0 degrees. These and other constraints dictated that the probe entered the atmosphere near 6.58 N latitude. At this latitude it would be likely that the probe would meet another desired mission objective, namely to enter in a relatively ‘‘typical’’ and calm region of the Jovian atmosphere. However as luck would have it, the probe entered the southern region of what is known as a 5-micron hot spot. Such regions are thought to be localized thinings in the clouds, and are relatively transparent to planetary radiation near 5 mm emitted from deeper atmospheric levels. The hot spots are therefore bright at 5 mm, being defined as having brightness temperatures near 5 mm in excess of 240 K (Ortiz et al., 1998; Orton et al., 1998). The hotspots cover only about 1% of the total surface area of Jupiter, but they are concentrated near and north of the equator. However, even then they cover only about 15% of the equatorial surface area. The vertical extent of hot spots is unknown. Significantly, 5 mm hot spot spectra were used in almost all pre-Galileo remote sensing attempts to make deductions about the Jovian atmosphere below the upper clouds. Having entered one of the hotspots, the probe has provided ground truth for those observations and their interpretation, and in several instances probe data have been at variance with the conclusions based on remote sensing. On the other hand, there is no doubt that conditions in hot spots are not typical of Jupiter as a whole. Just the fact that hot spots are regions of local clearing in the clouds is a clear indication that they are not typical. As will be discussed later, certain aspects of the probe data are

19

likely to be understood only in terms of hot spot dynamics, whereas other aspects of the data can with confidence be related to Jupiter as a whole. Ground based observations conducted with the NASA Infrared Telescope Facility (IRTF) between December, 1993 and July, 1997 (Orton et al., 1998; Ortiz et al., 1998), a time encompassing the December 7, 1995 entry date of the Galileo probe, show that during this approximately 3.5 year interval there were 8–10 quasi-evenly spaced hot spots at any given time between latitudes 68–88 N. All the hot spots moved eastward as a group with a velocity near 103 m s 21 in System III. Fig. 3 illustrates the appearance of hot spots in the visible in an image of Jupiter taken by the Cassini spacecraft. Individual hot spot morphology is temporally variable, but the overall pattern is relatively more steady. Ortiz et al. (1998) discuss longitudinal locations, morphology, and evolution of 5 mm hotspots at the probe entry latitude of 6.58 N. Ortiz et al. suggest that the hot spots are the result of equatorial Rossby waves having zonal wavenumbers 8–11. Rossby waves are planetary scale atmospheric waves which depend for their existence on the variation with latitude of the vertical component of the angular planetary rotation vector, ] V. Twice the vertical component, 2V sin u, where u is latitude, is termed the Coriolis parameter. The atmospheric equations of motion on a rotating planet involve the Coriolis force, 2 V ] 3 u, where u is the atmospheric wind vector. For the situation where the atmosphere can be considered to be shallow, i.e., the vertical extent of the atmosphere is small compared to the planet radius, the Coriolis force for planetary scale motions depends on the vertical component of V ] more than the horizontal component, and is therefore proportional to the Coriolis parameter. It is the variation of the Coriolis parameter through sin u that is critical for the existence of Rossby waves (a treatment of Rossby waves can be found in any textbook on dynamic meteorology, e.g. Holton, 1992). The zonal wavenumber characterizing a wave refers to a wave having a dependence on longitude proportional to e im f , where f is longitude, and m is the zonal wavenumber. Rossby waves are ubiquitous features in the Earth’s atmosphere. Under certain circumstances these waves can be trapped near the equator, mean-

20

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Fig. 3. Cylindrical projection of Jupiter from 608 N to 608 S taken in the visible by the Cassini spacecraft on October 31, 2000 as it approached Jupiter. Just north of the equator (denoted by the solid line), at the northern boundary of the bright equatorial zone and southern boundary of the north equatorial belt (dark band), can be seen 8–10 dark features which are probably 5 mm hot spots. These features are located very near the entry latitude of the probe, 6.58 N. Hot spots appear dark in visible images, but not all dark spots in the visible are necessarily 5 mm hot spots. Hot spots appear dark in the visible because of local thinning of the clouds, and hence there is less reflected visible light. Bright plume features are readily discernable adjacent to some of the spots. See Orton et al. (1998) for detailed comparison of Hubble Space Telescope visible images of particular hot spots and NASA Infrared Telescope Facility (IRTF) infrared images near 5 mm.

ing that the wave amplitude dies off rapidly with latitude on both sides of the equator. This, and the fact that hot spots are only observed near the Jovian equator, formed part of the basis for the suggestion by Ortiz et al. (1998) that the hot spots are equatorially trapped Rossby waves. Ortiz et al. also mention the possibility that the bright plumes associated with, and located between, hot spots may simply be the result of a different phase of the same Rossby waves. Allison (1990) had previously considered a model of equatorially trapped waves on Jupiter, and suggested that the plumes could be due to such waves. As will be discussed later, the wave interpretation of the hot spots is attractive. However, recent modeling of hot spots (Showman and Dowling, 2000) suggests that the waves cannot have small amplitude, and hence are not self-consistently described by linear wave theory, which was used by Ortiz et al. (1998). Nonlinear effects associated with large wave amplitudes must be important in maintaining the coher-

ence of any particular hot spot for an appreciable time. Hot spot brightness temperatures are typically near 255 K (Orton et al., 1996), corresponding to an emission pressure level of about 4 bars. For reference, the bottom boundary of the upper NH 3 ice cloud is near 0.5 bar. The hot spot into which the probe descended had a 4.78 mm brightness temperature of 252 K (Orton et al., 1998). Fig. 4, taken from Fig. 3 in Orton et al. (1998), illustrates the probe entry and descent trajectory, projected on NASA IRTF 4.78 mm images of the probe entry site hot spot. Several points are apparent from the figure. First, the probe apparently descended in the southern region of the hot spot. The location in the hot spot will become significant when probe wind measurements are discussed. Second, the probe was within the hot spot (at least as far as horizontal position) throughout the entire descent portion of the mission, descent portion meaning that

R.E. Young / New Astronomy Reviews 47 (2003) 1–51

21

Fig. 4. Location of the Galileo probe entry site within the entry hot spot, taken from Orton et al. (1998). Images were taken by the facility near-infrared camera at the NASA IRTF at the summit of Mauna Kea in Hawaii. Colors in each image are proportional to brightness at 4.78 mm, white being the brightest. Each image is scaled to the brightest pixel in that image. A single 0.5830.58 pixel is used to denote the 6.58N planetocentric latitude and System III longitude of the probe entry, and a 1.58 longitude (3 pixel) extent depicts the longitudinal extent of the probe entry path starting from 450 km above the 1 bar pressure level. A 1 2 s circle shows the effects of pointing uncertainty of the telescope on the location of the final portion of the probe entry. The panels from different dates were aligned together using a drift rate of 103 ms 21 relative to System III.

part of the mission where the probe was making direct atmospheric measurements. Immersion in the hot spot will be an important factor when the vertical profiles of condensible species are discussed. Third, the hot spot maintained its integrity for the two month period illustrated in the figure, and actually did so for much longer (see Orton et al., 1998). This is a crucial property of hot spots that must be matched by theoretical models, and as will be discussed, the primary reason that nonlinear rather than linear wave theory seems to be required to understand hot spots. In the sections of the paper that follow, in which Galileo probe observations are presented and dis-

cussed, a theme running through the entire discussion will concern what aspects of the probe results were significantly influenced by the fact that the probe descended into a 5 mm hot spot, and which aspects can be considered to be representative of Jupiter’s atmosphere as a whole.

4. Galileo probe results

4.1. Composition In general, abundance measurements involving non-condensible trace species would be expected to

22

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be representative of global abundances. Again, with regard to species abundances or mixing ratios, global means averaged over the Jovian atmosphere as a whole. In the absence of sources or sinks of a particular species acting over time scales smaller than atmospheric mixing times, atmospheric motions cause non-condensible species to be well mixed throughout the atmosphere (except at very high altitudes, above the homopause), resulting in a mixing ratio invariant with position. In the case of the Earth, for example, because of condensation and evaporation, water vapor abundance in the atmosphere varies considerably, both temporally and spatially. On the other hand, on an annual mean basis the non-condensible CO 2 is generally well mixed, i.e., has constant mixing ratio, in the terrestrial troposphere, in spite of there being significant sources and sinks (U.S. Standard Atmosphere, 1976). In the case of Jupiter, we can have confidence that the abundances of particular key non-condensible trace species measured by the Galileo probe represent global atmospheric values. These would include CH 4 , He and other noble gases, together with isotopic ratios of carbon, noble gases, and hydrogen. As will be shown later however, the probe measurements of the condensibles NH 3 , H 2 O, and H 2 S clearly are influenced by local effects, but even in these cases there is good reason to believe that the global abundances of N and S were determined.

4.1.1. Helium Fig. 5 and Table 4 present the He mass fraction Y ; MHe /(MHe 1 MH2 1 MZ ) and show how it compares to that of other bodies in the solar system. Z represents elements heavier than helium. Two independent Galileo probe measurements establish that to two significant figures, in Jupiter’s atmosphere Y 5 0.23: Y 5 0.23460.005 was obtained by the helium abundance detector (HAD) (von Zahn et al., 1998), and Y 5 0.23460.04 was obtained by the NMS (Niemann et al., 1998). A Y ¯ Y9 value of 0.23 is below the protosolar value of 0.27, but significantly above the Voyager value of 0.18. The published errors for the Voyager and Galileo measurements preclude overlap of the Voyager and Galileo values, but not by much. The motivation for revising the Saturn helium mass fraction (Conrath and

Fig. 5. Helium mass fraction in the Sun and outer planets, adapted from von Zahn et al. (1998). Sun (P) indicates the protosolar value, Sun (S) indicates the value in the convective zone of the Sun from helioseismology observations. For numerical values and references see Table 4.

Gautier, 2000) came from being unable to reconcile the Jovian Voyager and Galileo He measurements. As Fig. 5 and Table 4 illustrate, Jupiter is somewhat depleted in helium in its outer envelope relative to the protosolar helium abundance, probably because of some amount of helium separation near the Jovian molecular-metallic hydrogen boundary and associated gravitational settling of helium (see discussion in Section 2.1). The Galileo value for Y would then imply that there is of the order 10 M % more helium in the molecular hydrogen region of Jupiter than indicated by the Voyager value. For a fixed bulk Jovian helium mass mixing ratio, presumably the protosolar value, there would then be 10 M % less helium in the deeper interior, a mass figure which is the same order as estimates of the total mass of heavy elements for Jupiter from Jovian interior models (see previous discussion of total heavy element abundances). The Galileo helium mixing ratio Y thus represents a significant change in helium abundance in the molecular hydrogen en-

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23

Table 4 Helium mass fraction of the Sun and outer planets, adapted from von Zahn et al. (1998). Except where noted, the mass fraction is the mass ratio of helium to the total mass, assuming a heavy element (heavier than helium) mass fraction Z50.019 Source

Y

Reference

Galileo probe

0.23460.005 0.23460.04

von Zahn et al. (1998) Niemann et al. (1998)

Uranus Neptune Sun: Helioseismology

0.1860.04 0.0660.05 (original) 0.21*60.04 (revised) 0.26*60.05 0.3260.05 0.2460.01

Sun: Protosolar Primordial

0.27 0.23

Conrath et al. (1984) Conrath et al. (1984) Conrath and Gautier (2000) Conrath et al. (1987) Conrath et al. (1991) Kosovichev et al. (1992) Perez-Hernandez and Christensen-Dalsgaard (1994) Bahcall et al. (2001) Olive and Steigman (1995)

Voyager Jupiter Saturn

*Computed with Z50.

velope of Jupiter from that based on Voyager. At best, it is only marginally possible to get Jovian interior and evolution models to be consistent with known constraints using the Voyager Y (cf. Guillot, 1999), whereas the Y determined by Galileo is in the range for which current Jovian interior and evolution models agree with the known constraints of planetary mass, radius, gravitational moments, and internal heat flux (Guillot et al., 1997; Guillot, 1999; Hubbard et al., 1999; Gudkova and Zharkov, 1999). These same models also suggest that the original Voyager determination of Y for Saturn is considerably too low. The revised Y9 for Saturn shown in Fig. 5 and Table 4, 0.18 , Y9 , 0.25, (Conrath and Gautier, 2000) represents a three to fourfold increase over the original Voyager value. It is based on Voyager IRIS spectra alone, rather than the combination of IRIS spectra and radio occultation temperature measurements used originally. The reason for the discrepancy between the early Saturn He determinations and the revised values may be that there are systematic errors in the Voyager radio occultation measurements. As mentioned previously, it has been suggested that there was a calibration problem with the Voyager IRIS data in the spectral region having wavelengths shorter than 5 mm (Roos-Serote et al., 1999). However, the spectral region used for the He derivation is beyond 5 mm, having wavelengths between 15 and 50 mm, which, based on Voyager IRIS methane retrievals using only IRIS spectra and

their agreement with the Galileo probe methane measurement, seems to be free of systematic error (Conrath and Gautier, 2000). Hence, Conrath and Gautier conclude that it is probable there is a systematic error in the Voyager Saturn radio occultation temperature retrievals, which of course do not involve the IRIS instrument. Their technique of using Voyager IRIS spectra alone to derive the helium abundance cannot be applied to Jupiter because of strong NH 3 gas and cloud opacity, but can be applied to Saturn because such NH 3 contributions are smaller. On the basis of the revised helium mass fraction, it now appears that there is no need to invoke greatly different evolutionary paths for Jupiter and Saturn. Previously, when it appeared that the Saturn helium mass fraction was so much smaller than that of Jupiter, it was necessary to consider scenarios by which helium separation and subsequent gravitational settling occurred much more readily within Saturn than within Jupiter. Such separation, if it does occur, should probably occur more readily on Saturn because Saturn is colder than Jupiter. However, the revised He determination makes it possible to account for helium separation and at the same time produce Saturn interior and evolutionary models which are more reasonable, and are consistent with observed abundances of heavy elements such as methane (Hubbard et al., 1999; Guillot, 1999; Gudkova and Zharkov, 1999). Thus, the determination of the Jovian helium mass mixing ratio by the Galileo

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probe led not only to better understanding of Jupiter and its evolution, but indirectly led to what is almost certainly a much more accurate view of Saturn.

4.1.2. Carbon and other heavy elements An up to date overview of the composition measurements from the Galileo probe NMS is given in Atreya et al. (2002), and the abundances of the major and principal minor species are summarized in Table 5, which is adapted from that paper. The Galileo determination of CH 4 / H 2 5(2.160.4)3 10 23 ¯2.9 solar can be considered to be the global volume mixing ratio. This result is in the range

anticipated from pre-Galileo remote sensing observations (see Section 2.1). Prior to Galileo, a key question was whether the mixing ratios of other elements heavier than helium were similar to that of carbon. Elements of obvious interest are sulfur, nitrogen, and oxygen. Thermochemical models of the Jovian deep atmosphere were based on the assumption that heavy elements would have the same mixing ratio as CH 4 (cf. Fegley and Lodders, 1994), but as discussed in Section 2.1, there were observational reasons to believe that each of S, N, and O could have smaller mixing ratios. Conversely, there was some evidence (see Section 2.1) that O could be

Table 5 Composition of Jupiter’s atmosphere as measured by the Galileo probe (adapted from Atreya et al. (2002) Major species

Mixing ratio relative to H 2

H2 He

1.0 0.15760.003 a

Principal minor species (Element represented) H2O (O)

Elemental abundance Jupiter / Sun † 0.80760.02

#10 26 (#4 bar, hotspot) (5.662.5)310 25 (12 bar, hotspot) 6.0(13.9, 22.8)310 24 (19 bar, hotspot)b

0.35 (10.23, 20.16) (19 bar, hotspot)*

CH 4 (C)

(2.160.4)310 23

2.960.5

NH 3 (N)

,10 26 (0.4 bar, hotspot)d (2.861.2)310 24 (4 bar, hotspot)e (7.161.2)310 24 (8 bar, hotspot)e (7.163.2)310 24 (9–12 bar, hotspot)b,f

3.660.5 (8 bar, hotspot)* 3.361.5 (9–12 bar, hotspot)*

H2S (S)

,7310 27 (#4 bar, hotspot)c 7310 26 (8.7 bar, hotspot)c (7.760.5)310 25 (16 bar, hotspot)c

2.560.15 (16 bar, hotspot)*

(1563)310 26 g (9.261.7)310 29 g (8.8461.7)310 210

2.560.5 2.760.5 2.660.5

Noble gases 36 Ar Kr Xe

c

g

† Uncertainties do not include uncertainties in solar abundances. *The mixing ratio of the deepest measurements for these species (elements) is taken as the global mixing ratio, except for the probe radio link determination of NH 3 abundance which is taken at 8 bars. a von Zahn et al. (1998). b Atreya et al. (2002). c Niemann et al. (1998). d Sromovsky et al. (1998). e Folkner et al. (1998). f Mahaffy et al. (1999). g Mahaffy et al. (2000).

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even more enriched than C. Even for in-situ measurements, the difficulty with S, N, and O is that these elements are involved in condensation processes in Jupiter’s atmosphere at pressures u5 bars, and as will be discussed later, the fact that the Galileo probe descended within a 5 mm hot spot clearly affected the probe measurements of these elements. Nevertheless, there is good reason presented later to think that the mixing ratios of NH 3 and H 2 S obtained by the probe at pressures greater than about 8–12 and 15 bars, respectively, represent global mixing ratios, while that of H 2 O does not at any pressure sampled by the probe. The deep volume mixing ratio of H 2 S as measured by the Galileo probe is (7.760.5)310 25 ¯2.5 solar (Niemann et al., 1998), and that of NH 3 in the region 8–12 bars is (7.161.2)310 24 ¯3.6 solar (Folkner et al., 1998; Mahaffy et al., 1999; Atreya et al., 2002). The results for H 2 O will be considered later when the vertical distributions of condensible species observed by the probe are discussed. The most surprising result from the above measurements in terms of global abundance is that for NH 3 , see Table 5. As discussed in Section 2.1, prior to Galileo it was thought that the Jovian C / N ratio was about 2 times solar. However the Galileo probe value of C / N derived from the abundances of CH 4 and NH 3 is approximately solar. A C / N. solar was consistent with the idea that very volatile elements such as N 2 would not be delivered efficiently to Jupiter because thermal conditions were too warm, both in the solar nebula region where Jupiter formed and in regions where impacting plantesimals were formed. This idea was supported by C / N. solar in Halley’s Comet (Owen and Bar-Nun, 1995). Apparently at least part of this picture has to be revised. In order that the probe global abundance of NH 3 measured in the deep atmosphere be consistent with Jupiter’s microwave spectrum, de Pater et al. (2001) have shown that the NH 3 abundance must decrease at pressures u4 bars, and reach subsolar abundance at pressures u2 bars, a result consistent with the probe net flux radiometer (NFR) measurements at pressures ,2 bars (Sromovsky et al., 1998). However this result immediately raises a problem. How is it that the NH 3 abundance could decrease generally with decreasing pressure at pressures well above the pressure where the NH 3 ice clouds form? Carlson et

25

al. (1993) had modeled a scenario whereby a super solar abundance of H 2 S chemically combined with NH 3 to form NH 4 SH ice particles, which are thought to make up the second Jovian cloud layer near 2 bars. This same scenario is discussed in de Pater et al. (2001), however both Carlson et al. and de Pater et al. point out that there is little actual knowledge of how this process would occur, and that a good deal more work needs to be done before any real confidence can be placed in the cloud chemistry and physics associated with producing NH 4 SH ice clouds from NH 3 and H 2 S. Furthermore, the probe measurements of H 2 S show that at pressures ,4 bars there is virtually no detectable H 2 S, although the fact that the probe was in a hot spot must be kept in mind. Thus, at the present time it is not clear what mechanism would cause the abundance of NH 3 to decrease with decreasing pressure at pressures u4 bars. The variation of NH 3 abundance with depth apparent from Table 5 will be discussed at more length when the abundance profiles of all the condensibles are considered in Section 4.1.4. The Galileo probe NH 3 result also raised interest in whether other very volatile elements, such as noble gases heavier than helium, were also globally enriched to the same extent as N. Owen et al. (1999), Mahaffy et al. (2000), and Atreya et al. (2002) report noble gas mixing ratios from the probe NMS, and give abundances relative to H 2 for argon, krypton, and xenon of 2.560.5 solar, 2.760.5 solar, and 2.660.5 solar, respectively. In contrast to the above three noble gases, neon is highly depleted, being measured to have a volume mixing ratio near 0.1 solar (Niemann et al., 1998; Mahaffy et al., 2000). The current thinking is that helium is slightly depleted in the molecular envelope for reasons discussed previously, and as helium ‘‘rain drops’’ gravitationally settle toward the deeper interior, they carry neon with them because neon is soluble in helium (Roulston and Stevenson, 1995). A slight bit of helium separation can remove most of the neon because of their relative mixing ratios. Based on Table 5, Fig. 6 summarizes the heavy volatile element global abundances derived from the Galileo probe, excluding neon. It is clear that with the exception of neon, and O in the form of H 2 O which is affected by local processes within the PES, all the heavy element abundances measured by the

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Fig. 6. Summary of principal heavy element and noble gas global abundances relative to solar as measured by the Galileo probe (adapted from Owen et al., 1999); see also Table 5. Only upper limits have been placed on the oxygen abundance, as denoted by the crooked arrow. The dashed error bar for N represents the lower portion of the error bar in the NMS determination of N, which error was 61.5 solar units (Mahaffy et al., 1999).

probe are ¯3 times solar. Based on the probe measurements, a reasonable generalization would seem to be that all elements heavier than helium have Jovian global mixing ratios about three times solar, although there may be reasons to think O in the form of H 2 O could be enriched beyond three times solar (see discussion below). There are at least two aspects of these results which raise significant issues. First, very volatile elements are enriched rather than being depleted (see previous discussion of pre-Galileo NH 3 measurements and the expected C / N ratio). Second, how is it that the enrichments are all approximately three times solar? At the present time there are two concepts for trying to understand the observed volatile enrichments. One is based on the laboratory work of Bar-Nun et al. (1988) which examined the trapping of gas mixtures containing CH 4 , Ar, and either CO or N 2 in amorphous water ice (and that group’s previous work treating a single gas at a time, see Bar-Nun et al., 1988, for discussion and references). The second concept is the delivery of volatiles in clathrate hydrates, and is based on calculations of the thermodynamics of clathrate hydrates at low pressures applicable to the solar nebula (Lunine and Stevenson, 1985). As stated by Lunine and Steven-

son, clathrate hydrates are water ice compounds in which a distinctive open lattice structure of the ice forms ‘‘cages’’ stabilized by the inclusion of molecules of other chemical species, e.g. noble gases. The implications of each idea are discussed below. We consider first the idea that volatiles delivered to the outer planets were trapped in amorphous ice. Prior to Galileo it had been generally assumed that planetesimals involved in the formation of the outer planets were similar, perhaps identical, to comets observed in the Oort cloud at present. Because these Oort cloud comets are generally thought to have formed in the region where the giant planets were forming (see discussions in Mumma et al., 1993; Malhotra et al., 2000), and then been scattered to where they are today, Oort cloud comets are expected to have originated with temperatures too warm to trap N 2 or the noble gases effectively (Owen and Bar-Nun, 1995; Bar-Nun et al., 1988). This expectation is consistent with measurements of the composition and elemental ratios of Comet Halley and the lack of N 2 in comets (Wyckoff et al., 1991; Womack et al., 1992; Krankowsky, 1991): in Comet Halley, N 2 / NH 3 ¯0.1, and C / N¯6 times the solar ratio. On this basis, the giant planets, and in particular Jupiter, would not be substantially enriched with respect to N or the heavier noble gases. Following this same line of reasoning, if planetesimals which formed near the H 2 O ice line in the solar nebula (near 3–5 AU according to models, see later discussion) dominated the delivery of heavy elements to Jupiter, there would be considerable depletions of the more volatile elements. As Owen et al. (1999) conclude based on the concept of volatile trapping in amorphous ice, the only apparent explanation of the situation is that super solar abundances of elements were delivered to Jupiter in very cold (,30 K) icy planetesimals. The question would then become what scenarios might be plausible for very cold planetesimals to deliver heavy elements? Three potential scenarios were identified by Owen et al. (1999), but as they point out without elaboration, each has associated difficulties. Each is described below, together with a crude assessment of the likelihood that each could occur. One possibility is that the responsible planetesimals could have formed in the interstellar medium during the early phases of molecular cloud fragmentation and col-

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lapse. There are at least two reservations regarding such a possibility. First, it is not clear icy grains could grow in the interstellar medium during the early stages of cloud collapse (Weidenschilling and Ruzmaikina, 1994). Second, it seems likely that cometary material has been processed in the solar nebula (Irvine et al., 2000; Lunine et al., 2000), so the question arises as to how much of the original interstellar material was retained. There are no known solid objects in the solar system that have the required composition, although it has been suggested that Kuiper Belt objects might (Owen and Bar-Nun, 1995). However as mentioned earlier, Kuiper Belt objects currently have about a 1% efficiency for impacting Jupiter (Levison and Duncan, 1997). If a similar impact efficiency applied in the solar nebula, hundreds to thousands of Earth masses of objects would be required in order to deliver Jupiter’s super solar inventory of volatiles. On the other hand, it is conceivable that gas drag from the solar nebula could have caused small objects in the outer solar system to migrate inwards, perhaps resulting in a higher impact efficiency (cf. Kary et al., 1993; Kary and Lissauer, 1995). Another conceivable scenario is that perhaps Jupiter formed beyond 30 AU where the solar nebula had temperatures ,30 K and migrated to its present position. But this too has difficulties. The currently favored method of forming giant planets is accretion of a rock and ice core, the order of 10 M % , followed by accretion of the gaseous envelope (cf. Wuchterl et al., 2000). One of the major hurdles of this model has been forming Jupiter within the lifetime of the solar nebula, O(10 7 ) years (Pollack et al., 1996). The models have to be pushed somewhat in terms of increased density of the solar nebula in order to achieve formation times u10 7 years (Lissauer et al., 1995). A formation region for Jupiter much farther out in the solar nebula would exacerbate this problem considerably because of the increased time between planetesimal encounters, due to both longer orbital periods and larger distances between objects (Thommes et al., 2002). A second method of forming giant planets, localized gravitational instability of the nebula (Boss, 1997, 1998, 2001), might have short giant planet formation times, but a computation following the stages from initial instability to planet formation has yet to be done. Therefore, it is not

27

clear this mechanism is viable. Whichever formation process is correct, migration of a giant protoplanet from tens of AU to 5 AU may not be all that difficult. It is certainly clear that giant planets forming in the Jupiter–Saturn region can readily migrate inwards in a protoplanetary disk, in some cases all the way to the parent star (cf. Trilling et al., 1998). The problem would seem to be how to understand a sudden cessation of inward migration of Jupiter near 5 AU with a circular orbit, coupled with migration of one or more of the other outer planets. And of course, a Jupiter forming at tens of AU rather than 5 AU should also have implications for the formation and evolution of the terrestrial planets. The third possibility for explaining Jupiter’s heavy element inventory mentioned by Owen et al. (1999) is that solar nebula thermal structure might be colder than expected. The canonical argument is that Jupiter formed near the water ice line, corresponding to a temperature of about 150 K, because that should be the region where the density of solid material in the solar nebula was largest (interior to the ice line water is in vapor form and the density of solids is therefore less; on the outer side of the ice line it is expected that nebula density continues its general decrease with distance from the Sun). Observations that objects within the asteroid belt seem to be principally rocky material (cf. Cruikshank et al., 2001; Millis and Dunham, 1989; Miller et al., 2002; Belton et al., 1996) are consistent with this view. There is also observational evidence that temperatures near Uranus and Neptune were not cold enough to produce planetesimals that retained the most volatile elements, i.e., as mentioned previously, observations of Comet Halley indicate that it is depleted in N 2 , perhaps but not necessarily due to inefficient trapping of N 2 at formation temperatures .30 K (Krankowsky, 1991; Wyckoff et al., 1991). On the other hand certain classes of meteorites, in particular CM carbonaceous chondrites, show evidence of aqueous alteration, although the initial water / ice to rock ratio is unknown (cf. Rosenberg et al., 2001 and references therein). This aqueous alteration could be considered evidence that the water ice line in the solar nebula was inside 2.5–3 AU, assuming that the meteorites originated in the asteroid belt. Thus, at the present time observational evidence as to the temperature profile in the early solar nebula, and the

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location of the water ice line, does not appear conclusive. Current theoretical solar nebula thermal models are uncertain and have unresolved issues. A discussion of model nebula thermal structure at very early times, less than about 1 million years, when accretional heating dominates stellar luminosity, can be found in Bell et al. (2000) and references therein. At what should pertain to later times, Chiang and Goldreich (1997) have modeled passive disks around T Tauri stars. At present, theoretical models do not seem to either preclude or firmly support the possibility that cold enough temperatures to trap volatiles such as N 2 in amorphous water ice existed near Uranus and Neptune after about 1 million years. We now consider the idea put forth by Lunine and Stevenson (1985) that volatiles were delivered to Jupiter in the form of clathrate hydrates. Gautier et al. (2001) recently expanded upon this idea to include evolutionary models of the solar nebula, and attempted to explain the Galileo probe measurements by modeling the delivery of volatile material to Jupiter in clathrate hydrates from the ‘‘feeding zone’’ of Jupiter during Jupiter’s formation. Gautier et al. point out that at the 150 K temperature at which water vapor condenses to form ice in the solar nebula near Jupiter’s formation region, the ice formed is mostly the crystalline form of ice, not amorphous ice. Under these conditions, surface adsorption trapping of volatiles in amorphous ice would not be significant. Based on their thermodynamic clathrate hydrate model, Lunine and Stevenson had predicted the Jovian ratios of Ar, Kr, and Xe, e.g. Xe /Ar, as a function of the enrichment of CH 4 over solar abundance. For the enrichment of CH 4 of 2.9 times solar observed by the probe, their model predicts a Xe /Ar ratio of approximately 9 times solar, when in fact the probe measurements show Xe /Ar to be solar. Mahaffy et al. (2000) argue that this discrepancy is evidence that at least the noble gases were not delivered to Jupiter in clathrate hydrates. However to within error limits of the measurements, the correct enrichments of C, N, Ar, Kr, and Xe can be obtained with clathrate hydrates if the evolution of the solar nebula near Jupiter’s formation region is taken into account according to the model of Gautier et al. (2001). Gautier et al. use the thermodynamic stability

curves of clathrate hydrates computed by Lunine and Stevenson (1985), and show that the final Jovian abundances of noble gases, together with the abundances of C, N, S, and O, depend on the stability curves, but just as importantly, depend on the timing of the final hydrodynamic collapse of material in the feeding zone of Jupiter to form the planet. In their model, once clathrated ices agglomerate and form planetesimals of centimeter to meter sizes, the planetesimals decouple from the gas of the solar nebula, and the clathrate hydrate mass no longer increases with time. However the mass of H 2 gas is continuously decreasing as the nebula evolves. Therefore the final Jovian abundance of volatiles relative to H 2 is determined by the timing of the final hydrodynamic collapse of material in the feeding zone to form Jupiter. Their model predicts about the correct abundances of the elements measured by the Galileo probe, except that it predicts too high an abundance of S. Gautier et al. hypothesize that H 2 S will corrode Fe alloy grains, and by assuming a smaller than solar abundance of H 2 S in the nebula for this reason, they can also match the S abundance measured by the probe. However, Atreya et al. (2002) argue that Gautier et al.’s explanation for the calculated overabundance of S, namely that a substantial amount of sulfur is consumed in the inner solar nebula through the reaction of H 2 S with Fe alloy grains to form troilite (FeS), is quantitatively unsupportable. A major uncertainty associated with the clathrate hydrate hypothesis is that the needed laboratory experiments have not yet been done to clearly establish that such a process does indeed occur effectively under conditions appropriate to the solar nebula. Lunine and Stevenson (1985) suggest a number of focused laboratory studies. At the present time there is also no observational evidence of clathrate hydrates in solar system bodies to support this theory of delivery of volatiles to Jupiter. The question of how fresh ice would be continuously exposed in the solar nebula so that clathrate hydrate formation could continue long enough to be effective is addressed by Lunine and Stevenson and Gautier et al. by arguing that collisions among planetesimals would continuously expose fresh ice. There is a potential discriminator between the amorphous ice delivery concept and the clathrate hydrate hypothesis, namely the abundance of O in

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the form H 2 O on Jupiter. Because the clathrate hydrate process would require considerably more water ice to deliver the observed elemental abundances to Jupiter than would the amorphous ice delivery process, the model of Gautier et al. (2001) predicts that O should be enriched on Jupiter by about 8 times solar. However, Gautier et al. assume a 100% efficiency in clathrate formation, i.e., every cage of crystalline ice gets filled with a molecule of another species. The actual efficiency is unknown, and to the extent that the efficiency is really less than 100%, more water ice, and hence more O, would be required. Therefore, one could reasonably expect that delivery of heavy elements by clathrate hydrates would imply .8 times solar amount of O (H 2 O), whereas considerably less O would be expected if the elements were delivered bound in amorphous ice (cf. Gautier et al.). As has been discussed previously, however, the Galileo probe did not determine the global abundance of O, and such a determination likely awaits a future Jupiter mission. To summarize, there seems little reason to doubt that the Galileo probe results concerning the heavy element abundances apply to the global composition of Jupiter’s atmosphere, with the possible exception of oxygen (H 2 O) for which there are identified qualifications. The finding that, with the exception of O and Ne (see Section 4.1.2), all measured global heavy element abundances exceed their respective solar abundance by a factor ¯3 has forced rethinking about conditions associated with the formation and evolution of Jupiter, the solar nebula, and the outer planets, but no clear scenario has emerged. As Atreya et al. (2002) state, whatever the mechanism of heavy element enrichment on Jupiter, no solid objects have been observed in the solar system that have solar ratios of Ar, Kr, Xe, S, and N relative to C, as does Jupiter. Yet as Owen and Encrenaz (2002) point out, such material accounts for about 18 M % on Jupiter if the total Z50.02 and is well mixed, implying that such material must have been the most abundant solid material in the early solar system. So where did it go? There appears to be no ready answer, indicating something is amiss in our understanding of Jupiter and the solar system.

4.1.3. Isotopic ratios For reasons discussed previously, measurement of isotopic ratios, especially of the noble gases, was

29

high priority for the Galileo probe. As with elemental abundances derived from non-condensible species, the isotopic ratios measured by the probe should be representative of global values. The particular isotopic ratios obtained by the probe NMS were, D/ H and 3 He / 4 He (Mahaffy et al., 1998; Niemann et al., 1998), 13 C / 12 C (Niemann et al., 1998), 15 N / 14 N (Owen et al., 2001; Atreya et al., 2002), 20 Ne / 22 Ne, 36 Ar / 38 Ar, and the ratios of Xe 128–132, 134, and 136 to total Xe (Mahaffy et al., 2000); the numerical values are summarized in Table 6, taken from Atreya et al. (2002). For the elements heavier than He, one word describes the isotopic ratios: solar. Although there are contradictory results for the solar value of 15 N / 14 N, ranging from less Table 6 Isotopic ratios measured by Galileo probe NMS (taken from Atreya et al., 2002) Elements

Sun

13

0.011 #2.8310 23

15

12

C/ C N / 14 N

36

Ar / 38 Ar Xe / Xe 134 Xe / Xe 132 Xe / Xe 131 Xe / Xe 130 Xe / Xe 129 Xe / Xe 128 Xe / Xe 20 Ne / 22 Ne 3 He / 4 He 136

D/ H

a

Jupiter b

5.7760.08 d 0.0795 0.0977 0.265 0.217 0.0435 0.274 0.022 13.8160.08 d 1.560.3310 24 (meteoritic) (2.060.4)310 25 e 3.060.17310 25 f protosolar values

0.010860.0005 (2.360.3)310 23 (0.8–2.8 bar)c 5.660.25 a 0.07660.009 a 0.09160.007 a 0.29060.020 a 0.20360.018 a 0.03860.005 a 0.28560.021 a 0.01860.002 a 1362 a 1.6660.05310 24 2.660.7310 25 g 2.2560.35310 25

h

Mahaffy et al. (2000). Normalized to 1.0 for xenon isotopes measured. Only 126 Xe and 124 Xe, which together make up 0.2% of the total xenon (solar) could not be measured by the GPMS. The xenon error bars in the Mahaffy et al. (2000) paper are with respect to the ratio of each isotope to its non-radiogenic terrestrial value. b Hashizume et al. (2000); a review of earlier extensive solar wind 15 N / 14 N measurements from the lunar record is given by Kerridge (1993); this result does not agree with a recent SOHO result of (5.561.38)310 23 from Kallenbach et al. (1998). c Owen et al. (2001), GPMS measurement samples largely from below the NH 3 condensation level. d Pepin et al. (1999). e Geiss et al. (2002). f Gautier and Morel (1997). g Mahaffy et al. (1998). h Lellouch et al., 2001).

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than 2.8310 23 (Hashizume et al., 2000) to about 5310 23 (Kallenbach et al., 1998), there are reasons to expect that the Jovian 15 N / 14 N measured by the Galileo probe, (2.360.3)310 23 , is the protosolar value (Owen et al., 2001). Thus, so far the Jovian elemental isotopic ratios have been pretty much as one would expect. The Jovian value of (D1 3 He) / H5(4.361.6)3 10 25 (Niemann et al., 1998) is the same to within the error bars as that currently in the LISM, (3.660.38)310 25 , derived from the LISM D/ H5 (1.560.18)310 25 and the LISM 3 He / H5 (2.160.33)310 25 (cf. Geiss et al., 2002, and references therein). On the other hand, as discussed near the end of Section 2.1, conceptually the Jovian (D1 3 He) / H could place a constraint on the enrichment of H 2 O over solar abundance in Jupiter’s atmosphere. Similarly, a comparison of the protosolar D/ H value5(2.060.4)310 25 listed in Geiss et al. (2002) with the Jovian D/ H5(2.660.7)310 25 obtained by the probe NMS (Niemann et al., 1998) places a limit on the enrichment of H 2 O. Using the nominal values of all the measurements, and assuming that the D/ H in H 2 O is the same as that measured in comets, approximately 3310 24 (cf. Lunine et al., 2000, and references therein), yields a Jovian enrichment of H 2 O between 20 and 30 times solar. However, the Jovian value of (D1 3 He) / H has significant uncertainty, about 637%, while the Jovian D/ H ratio has an uncertainty of 627%. Zero Jovian H 2 O enrichment is allowed by the uncertainties. Aside from measurement uncertainties, there has to be some uncertainty as to how much the atmospheric 3 He / H has been affected by helium settling to the deep Jovian interior, i.e., does helium settling towards the deep Jovian interior fractionate He? In any case, a correction for helium settling must be made when using the atmospheric (D1 3 He) / H ratio. Thus, until all the uncertainties can be reduced, including uncertainty in the appropriate value of D/ H in the H 2 O delivered to Jupiter, the Jovian H 2 O abundance is not well constrained as derived from the Jovian D/ H or (D1 3 He) / H ratios.

4.1.4. Condensible species Unlike non-condensible species discussed above, each of the three condensible species observed by the

Galileo probe, namely NH 3 , H 2 S, and H 2 O, exhibited a vertical abundance profile which was clearly affected by dynamical processes in the PES. In fact, the probe observations of these three species are still not clearly understood. Fig. 7 illustrates the vertical distribution of the abundance for NH 3 derived from the attenuation of the amplitude of the probe to orbiter radio signal (Folkner et al., 1998), and H 2 O and H 2 S abundance profiles measured by the probe NMS (Niemann et al., 1998), together with the cloud structure predicted from the thermochemical equilibrium cloud model of Romani et al. (1995) (see also de Pater et al., 2001). It should be noted that measurement of NH 3 by the NMS is very difficult, due to interactions of metal surfaces of the walls of the ionization and detector chambers with the gas. Thus, the NMS was unable to obtain a vertical profile of NH 3 , but did obtain one reading deep in the atmosphere which is in basic agreement with the values derived from the probe-orbiter radio link at pressures above about 8 bars (Mahaffy et al., 1999). There are several features of the NH 3 , H 2 S, and H 2 O abundance variations with depth that are quite puzzling: first, the lack of correlation of the abundance profiles with computed cloud (condensation) levels; second, the continued abundance increase with depth to pressures of 8–10 bars and greater; and third, the respective profiles for NH 3 , H 2 S, and H 2 O approach (presumably constant) deep mixing ratios at different fractional rates. Clearly the local composition of Jupiter’s atmosphere below cloud levels, in so far as condensible species are concerned, is affected by atmospheric dynamics occurring over regional spatial scales. Below the condensation level of each of NH 3 , H 2 O, and the level at which H 2 S presumably combines with NH 3 to form NH 4 SH ice clouds near 2 bars, it had been anticipated that each species, respectively, would have a constant mixing ratio, i.e. be well mixed due to convective motions. For example, water was expected to have a constant mixing ratio below the bottom of the water clouds, but decrease with height above the water cloud more or less along the water vapor saturation curve until water vapor was depleted. The fact that this simple picture did not apply to any of the condensibles at the PES was a considerable surprise. Atreya et al. (1997) and Owen et al. (1997) have

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Fig. 7. Condensible species abundance profiles measured by the Galileo probe, together with thermochemical equilibrium cloud model of Romani et al. (1995). Figure is adapted from de Pater et al. (2001). Cloud profiles correspond to cloud model of Romani et al. that best fits the radio observations of Jupiter between 0.05 and 100 cm wavelength. The water data points represent upper limits to the water abundance.

discussed how the depletions of NH 3 , H 2 S, and H 2 O below cloud levels might be understood in terms of a substantial downdraft in the hot spot through which the probe descended, in analogy to desert regions on Earth where dry air descends after having been dried out during ascent. However, the above papers did not quantify how such a downdraft could be maintained to depths reaching 10–20 bars or more, as required by the probe data. In contrast to the Earth, dry atmosphere on Jupiter is less dense than moist atmosphere. The buoyancy force, 2 g dr /r, where r is atmospheric density and dr is the density difference between the downdraft and surrounding atmosphere, is therefore working against the downdraft if the downdraft is drier than the surrounding atmosphere. More generally, Showman and Ingersoll (1998) address four questions regarding a downdraft and how it might explain the probe abundance observations: could the downdraft remain dry and avoid mixing with a moist deep atmosphere; under what conditions does the downdraft have a dry adiabatic (neutrally stable) thermal profile with depth (as essentially measured by the probe atmospheric struc-

ture instrument, see discussion under thermal structure, and as might be expected because of long atmospheric radiative relaxation time scales); could a downdraft penetrate to the 10–20 bar region; and finally, how could the downdraft produce different NH 3 , H 2 S, and H 2 O fractional increases with depth. Showman and Ingersoll suggest and discuss possible scenarios for each of these questions, to which the reader is referred. However, they were unable to construct a completely consistent picture that simultaneously satisfied the atmospheric equations of motion and produced the observed different fractional rates of increase of NH 3 , H 2 S, and H 2 O with depth. Baker and Schubert (1998) suggest that less dense dry fluid parcels can be carried several scale heights downwards into the deep Jovian atmosphere by convective entrainment. In other words, a dry fluid parcel is entrained in a downdraft of moist fluid parcels. They base their suggestion on 2-dimensional model simulations of high Rayleigh number convection, where the Rayleigh number is a measure of convective intensity. They also argue that convection could explain the regular spacing of hot spots if the

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hot spots are associated with regions of convective downdrafts. Not addressed in their paper is the question of how different species could have different mixing ratio profiles as a function of depth. Perhaps the most promising concept for understanding the probe measurements is that hot spots are the manifestation of large amplitude atmospheric waves, in particular equatorially trapped Rossby waves (see Section 3.2). In their 3-D numerical simulations of 5 mm hot spots, Showman and Dowling (2000) show that many properties of hot spots can apparently be understood in terms of nonlinear Rossby waves. In particular, they show that linear (small amplitude) Rossby waves are too dispersive and do not retain the integrity of a hot spot over any substantial length of time. On the other hand, highly nonlinear (large amplitude) simulations do produce relatively long-lived features that have many, although so far not all, of the observed characteristics of hot spots. So how might large amplitude Rossby waves explain the probe observations? The basic idea can be understood in terms of potential temperature, Q, which is defined as

S D,

p0 Q ;T ] p

k

R k 5 ], cp

(3)

where T is temperature, p is pressure, p0 is an arbitrary reference pressure, R is the atmospheric gas constant, and c p is specific heat at constant pressure per unit mass. For an ideal gas, c p d(ln Q ) 5 dQ,

(4)

where Q is entropy per unit mass. Thus, isentropic surfaces are surfaces of constant potential temperature. To a good approximation atmospheric motions can be considered adiabatic in the Jovian atmosphere, i.e. exchange of energy of fluid parcels with their surroundings is negligible over time scales associated with atmospheric motions. Thus, following an atmospheric fluid parcel, dQ / dt 5 0, and therefore each fluid parcel conserves its potential temperature as it moves. The important point is that in this situation there is no mixing of any conserved property of the atmosphere across isentropic, i.e. constant potential temperature, surfaces. In the absence of chemical or physical sources or sinks, the

mixing ratio of any trace species is such a conserved quantity. Therefore, an isentropic surface initially coincident with a surface of constant mixing ratio of a particular trace species, such as would occur in an undisturbed region of a plane parallel stratified atmosphere, will distort under atmospheric motions the same way as does the surface of constant mixing ratio. Fig. 8 illustrates qualitatively what may be happening in a hot spot. A large amplitude disturbance, made up of large amplitude Rossby waves propagating as a wave packet through a background plane parallel stratified atmosphere, pushes isentropic surfaces downwards locally. This distorted region defines the extent of the hot spot. The disturbance amplitude is a function of depth such that, for example, atmospheric parcels initially near the NH 3 , H 2 S, and H 2 O condensation levels near 0.5, 2, and 5 bars respectively are deflected to about 8, 16, and .20 bars respectively. Above the condensation level of a given condensible, its abundance decreases rapidly because condensation depletes the vapor phase from the atmosphere. So as the isentropic surface corresponding to the condensation level gets pushed downward in the hot spot, the atmosphere in the hot spot above that particular isentrope is depleted in the corresponding condensible species. Thus the simplified picture depicted in Fig. 8 might be the basis for understanding why NH 3 abundance increased from near 1 bar to 8 bars, H 2 S increased from about 8 to 16 bars, and H 2 O was still increasing when the probe ceased transmitting data near 22 bars. This is the qualitative picture that emerges from the simulations of Showman and Dowling (2000), however in their actual calculations they could only produce deflections of isentropic surfaces by a factor of 2 in pressure, considerably less than that required to match the probe observations. Nevertheless, at the present time this seems the most promising picture for understanding what is going on at the PES. In spite of the obvious effects of localized dynamics, there is confidence that the Galileo probe descended sufficiently deep that measurements of the global mixing ratios of NH 3 and H 2 S were obtained, since their abundance profiles with depth appeared constant to within the error bars at pressures greater than about 8 and 16 bars respectively. The same cannot be said for H 2 O, since at the end of the probe

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Fig. 8. Qualitative depiction of what may be occurring in 5 mm hotspots. Distance from hotspot center is normalized to the size of the hot spot. The solid curves represent the isentropic surfaces associated with the condensation levels of the labeled condensibles. Isentropic surfaces are equivalent to constant potential temperature surfaces for an ideal gas with constant specific heats (see text). The dashed line represents the pressure, pQ , marking the boundary between the deep atmosphere which is presumed neutrally stable (temperature gradient5adiabatic lapse rate5g /c p ), and the atmosphere above which has positive static stability. In the hotspot the isentropic surfaces are deflected downwards, resulting in depletion of the condensibles well below their condensation level in the surrounding environment.

mission there was no indication that the H 2 O abundance was leveling off or approaching a constant mixing ratio. Thus we are left with the rather unfortunate situation of still not knowing what the Jovian global abundance of H 2 O, or equivalently, O, really is. It is tempting to assume that like the other measured heavy elements, O / H 2 will be approximately 3 times solar. However as discussed in Section 4.1.2, there are plausible arguments that O / H 2 could be atypically large. If in fact Jupiter formed essentially at the water ice line in the solar nebula, water would be condensing in that region but more volatile elements would not. In that case it seems possible that Jupiter could have received an inventory of H 2 O significantly enhanced over that of other volatiles. In any case the determination of the Jovian O / H 2 ratio will have to await future missions to Jupiter. There are however other observations,

both ground based and from the Galileo orbiter, that imply significant amounts of water. Observations at 5.18 mm from the NASA Infrared Telescope Facility (IRTF) using the CSHELL spectrometer imply that H 2 O abundance between 3 and 6 bars is variable in the north equatorial belt (NEB), located between about 7–188 N, and at the latitude of the Great Red Spot (Hewagama et al., 2001). However these same observations imply that water clouds are not rare, and exist in the NEB, zones, and in particular the equatorial zone (EZ) (Bjoraker et al., 2002). There are also Galileo orbiter images of what appear to be cumulus water clouds in certain regions, particularly the Great Red Spot (Gierasch et al., 2000), and orbiter observations of lightning presumably generated in water clouds associated with cyclonic features apparent in the uppermost NH 3 cloud layers (Little et al., 1999). These ob-

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servations support the possibility that had the probe been able to provide measurements to even greater depths, it is likely that H 2 O would have reached an abundance i solar.

4.2. Clouds It is known that Jovian 5 mm hot spots are associated with local clearing or thinning of clouds to depths corresponding to several bars pressure. The fact that the Galileo probe descended in a hot spot immediately implies that the probe cloud observations are affected by processes associated with the PES, and are not representative of Jupiter as a whole. Expectations with regard to Jovian cloud structure prior to Galileo were summarized in Section 2.2. The probe mission sequence was designed to start descent mode data collection near the 0.1 bar pressure level, well above the expected uppermost NH 3 cloud layer located near 0.5 bars. However the probe began descent mode measurements 53 s (¯0.3 bars) later than planned (Young et al., 1996). Thus, probe measurements began either in the lower part or completely below the NH 3 cloud. Given the location of the PES, the NH 3 cloud would have been tenuous regardless. But tenuous is the word that describes the entire set of cloud properties derived from the probe cloud observations, which were made principally by the onboard nephelometer (NEP). Between 0.4 and 0.6 bars the NEP results are consistent with a particle population having particle sizes exceeding about 0.75 mm radius up to about 5 mm radius (Ragent et al., 1998), but uncertainties in inversion of the data probably still allow a population of smaller particles. The NEP detected what are probably the remnants of the NH 3 cloud between 0.46 and 0.55 bars, a more distinct but still tenuous cloud layer between 0.76 and 1.34 bars which appears consistent with an NH 4 SH ice cloud given the probe composition measurements (Atreya et al., 1997; Ragent et al., 1998), and a small vertically thin cloud layer near 1.65 bars which could be detached from the cloud above it. This lower level detached cloud might be a tenuous water ice cloud consistent with the highly depleted water abundance observed by the probe (Atreya et al., 1997; Ragent et al., 1998). The NEP did not detect any dense water cloud, but did detect a weak but distinctive structure

extending from 1.9 to 4.5 bars. This deeper structure remains to be identified, but Ragent et al. (1998) suggest that it could be a very tenuous water cloud. If the model explaining the vertical distribution of condensibles discussed in Section 4.1.4 is correct, it is easy to understand why only tenuous or no clouds were encountered by the Galileo probe. According to that model, within the hot spot cloud particles, if they formed at all, would be carried downwards below their respective condensation levels, and hence would evaporate. Under these circumstances it is pretty much impossible, based on the probe observations, to make general statements about the global characteristics of Jovian clouds. Characterizing Jovian cloud layers was one science objective that clearly suffered because of the PES location.

4.3. Thermal structure The Galileo probe determined the thermal structure of the Jovian atmosphere at the PES from approximately 1000 km above the 1 bar pressure level, corresponding to p ¯ 10 29 bar, to 132 km below the 1 bar level, corresponding to p 5 22 bars. During the atmospheric entry phase of the probe mission, accelerometers measured the deceleration of the probe, from which, knowing the aerodynamic characteristics of the probe, atmospheric density could be derived. Using the hydrostatic equation, pressure could then be obtained. From a composition model and the equation of state for an ideal gas, the temperature could then be derived consistent with both the composition model and the equation of state. The entry phase extended down to 23 km (0.35 bars), while descent mode data for which pressure and temperature were directly recorded and saved began at 17.3 km (0.47 bars). Thus there is a gap of 5–6 km over which temperature had to be interpolated, but the descent mode temperatures and entry mode temperatures match very well across the gap (Seiff et al., 1998). There is no reason to believe that upper atmosphere temperatures derived from the probe were affected significantly by local processes occurring at or within the PES, i.e. there is no evidence to suggest that 5 mm hot spots extend well into the upper atmosphere. On the other hand, one can expect planetary scale regional variations in upper atmos-

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phere thermal structure. Thus the probe upper atmosphere temperature profile should be representative of the equatorial region of Jupiter, but not perhaps profiles at mid and high latitudes. There could also be significant variations with time if upward propagating large amplitude atmospheric waves are present, or temperatures are affected by charged particlemagnetic field interactions. In the troposphere the general temperature profile at pressures near and greater than 1 bar is very likely to be a good representation of the tropospheric temperature profile elsewhere on Jupiter, in the sense that the thermal structure is close to being adiabatic. However as will be discussed, there is good reason to believe that the static stability derived from the probe data is affected significantly by local processes within the PES. The reason that the overall temperature profile should still be generally adiabatic and representative of other regions is that the static stability measured by the probe is tenths of K km 21 , about 10% of the adiabatic, or total measured, lapse rate of ¯ 2 K km 21 .

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Fig. 9 presents the entire temperature profile derived from the probe Atmospheric Structure Instrument (ASI) data (Seiff et al., 1997a, 1998). There are three aspects of the figure that will be emphasized in the following discussion: (1) deviations from an adiabatic T( p) in the troposphere, which represent the static stability; (2) the hot upper atmosphere and the rapid rise in temperature between about 3 and 0.025 mbar (¯300 km to ¯ 550 km altitude); and (3) the apparent presence of atmospheric waves in the thermal structure. The importance of atmospheric static stability was discussed in Section 2.3. The Jovian atmosphere above the tropopause, located at ¯0.26 bars (¯30 km altitude), is obviously highly stable, since the temperature is either mostly constant with height or increasing with height. The question is what is the static stability in the troposphere? Seiff et al. (1998) show that T( p) is very close to being adiabatic. They also derive the difference between the actual lapse rate of temperature and the adiabatic lapse rate, i.e. ˜ et al. (2002) rederive the static stability. Magalhaes

Fig. 9. Entire profile of temperature as a function of pressure as determined by the Galileo probe atmospheric structure instrument (see Seiff et al., 1998). At pressures less than 0.35 bar, temperature and pressure were derived from accelerometer data during the entry phase of the probe mission. At pressures greater than 0.47 bar, pressure and temperature were measured directly. The dashed portion of the curve represents that portion of the profile that is uncertain because of the effects of having to choose a boundary value of temperature at the topmost level sampled.

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the static stability from the Galileo probe ASI using only temperature data. The justification is that the accuracy of the pressure sensors suffered to some extent due to unanticipated thermal conditions in the probe interior (Young et al., 1996; Seiff et al., 1998). In order to derive the static stability using just temperature data, certain well justified assumptions had to be made, such as assuming that the probe was in equilibrium descent (for which atmospheric aerodynamic drag is balanced by gravity). Although the static stabilities derived by Seiff et al. using p ˜ et al. using and T and those derived by Magalhaes only T are qualitatively similar, there are significant ˜ et al. argue that quantitative differences. Magalhaes the static stability profile derived from T only is significantly more accurate than that based on both p and T since pressure sensor errors are eliminated. Static stability derived from T only is used to ˜ et al., 2002). Illustrated produce Fig. 10 (Magalhaes ¨ frequency, N 2 , where is the square of Brunt-Vaisalla 2 N 5 gS /T and S 5 static stability5(dT / dz 1 G ), where G is the adiabatic temperature gradient. The

quantity fp is the number fraction of H 2 molecules in the para-H 2 state (anti-parallel nuclear spins). The value of fp affects the value of the adiabatic lapse rate. Conrath and Gierasch (1984) retrieved fp on Jupiter near the 0.3 bar pressure level as a function of latitude from Voyager IRIS data. Near the equator fp was determined to be 0.29, and so that is the value of fp taken in the computation of N 2 . Other symbols are as defined previously. Fig. 10 illustrates that the ASI temperature data indicate that the Jovian atmosphere at the PES is by and large statically stable, at least to pressures near ˜ et al. (2002) there 20 bars. As derived by Magalhaes are, however, variations in both S and N 2 with depth, the physical cause of which remains to be determined. Near the region where N 2 appears slightly negative, it is possible that atmospheric molecular weight gradients produced by the variation of condensible species abundance with depth are sufficient ˜ et al., 2002). to stabilize the atmosphere (Magalhaes One consequence of the variations in N 2 with depth is that regions in which N is a local maximum can

¨ frequency, N 2 5 g /T S, S 5 (dT / dz 1 G ), as derived from the Galileo probe atmospheric structure instrument using Fig. 10. Brunt-Vaisalla ˜ et al. (2002). fp is the number fraction of H 2 molecules in the only temperature sensor data (see discussion in text). Based on Magalhaes para-H 2 state (see text).

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act as waveguides for atmospheric waves such as gravity waves. Gravity waves are atmospheric waves which depend on buoyancy to act as the restoring force when a fluid parcel is displaced in a statically stable atmosphere (cf. Holton, 1992). Atmospheric gravity waves propagating from regions of higher N to a region of lower N can be partially or even totally reflected when they encounter the smaller N region. Thus the waves can effectively be confined to the region of relatively high N. Allison and Atkinson (2001) argue that residuals in the Doppler wind measurements from the Galileo probe (Atkinson et al., 1998) are most plausibly interpreted as the manifestation of vertical winds associated with gravity waves, which imply a weak but downward increasing static stability between 5 and 20 bars. They deduce a downward increasing static stability S which scales as p 1.7 below 1 bar within the 1–20 bar region of Jupiter’s atmosphere. ˜ et al. This contrasts with the results from Magalhaes (2002) in which S does not show a consistent increase with depth, which can be seen from Fig. 10 and the expression for N 2 in terms of S. The quantitative discrepancy between the two analyses is still to be resolved. From an analysis of Kelvin–Helmholtz instability (see Section 2.4) using the zonal wind profile derived from the probe Doppler wind measurements (Atkinson et al., 1998), Bosak and Ingersoll (2002) conclude that if N 2 is ,10 26 s 22 , the small scale linear wave features observed near the equator (Flasar and Gierasch, 1986), mentioned in Section 2.3, could be the result of instability associated with the vertical shear of the zonal wind. However, as Fig. 10 shows, N 2 is generally greater than 10 26 s 22 , at least at the PES. On the other hand, Fig. 10 also shows that there are particular regions where N 2 is very small, and perhaps instability could occur in these regions. Based on the discussion of condensible vertical abundance profiles in Section 4.1.4, 5 mm hot spots would be expected to have smaller average static stability than surrounding regions external to the hot spot. This follows from the fact that g dQ ] ] 5 N2 Q dz

(5)

and as the surfaces of constant Q are deflected downward and farther apart in the hot spot, N 2

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would decrease from what it was originally outside the hot spot. However as Fig. 8 illustrates, in the background atmosphere there could be a particular value of Q, corresponding to a pressure pQ , such that at pressures . pQ the background atmosphere is neutrally stable. At pressures . pQ the static stability in the hot spot could then actually exceed that in the surrounding atmosphere as descending fluid parcels in the hot spot deflect isentropic surfaces downward. The fact that the probe measurements indicate generally positive static stability in a hot spot, together with the fact that they also indicate that water abundance was increasing with depth but had not yet leveled off by the end of the probe mission, implies that if the scenario depicted in Fig. 8 is correct, the Jovian atmosphere surrounding the PES is statically stable to depths corresponding to at least the condensation level of water. For a solar abundance of water, condensation in the form of water ice occurs near 5 bars. For greater water abundances the condensation is in the form of liquid water and occurs at higher pressures, for example near 7 bars for three times solar water, and near 9 bars for 10 times solar water. (It should be noted that in reality a deep solution cloud of NH 3 in liquid water would form instead of pure liquid water for water abundances a few times solar (cf. Romani et al., 1995; de Pater et al., 2001), but the above condensation level estimates for pure liquid water should be a reasonable approximation). As mentioned previously, certain dynamical features in the Jovian atmosphere imply positive static stability in the atmosphere well below the clouds, and not just in the atmosphere near the PES (cf. Allison, 2000). Hence, based on the probe measurements and arguments associated with the dynamics of observed atmospheric features, it seems reasonable to believe that the equatorial regions of the Jovian atmosphere are statically stable to depths below cloud levels, and this may be true at other latitudes as well. Radiative–convective thermal models of Jupiter’s atmosphere published prior to Galileo have generally predicted that the radiative equilibrium temperature profile becomes super-adiabatic at pressures greater than 0.5–0.7 bars (cf. Appleby and Hogan, 1984, and references therein; Guillot et al., 1994). In other words, near and below 1 bar, convective heat

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transport is predicted to dominate the vertical transport of energy. In contrast, the probe measurements indicate positive, although relatively small, static stability to nearly 20 bars at the PES, and by implication to i5 bars in the surrounding atmosphere. If this is representative of the situation in general, and dynamic meteorological arguments suggest that it is (see discussion in Section 2.3; also cf. Allison, 2000), the implication is that the radiative– convective models are missing an important contributor to the energy balance. In all likelihood aerosol opacity due to cloud particles has probably not been adequately treated. Heating within the clouds would stabilize the cloud level atmosphere relative to the atmosphere below. One would expect that at deeper levels, where pressures are significantly greater than the pressure corresponding to the condensation level of water or an H 2 O–NH 3 solution cloud, the radiative equilibrium temperature profile would be superadiabatic because of gaseous opacity. Convection would then dominate energy transport much below the deepest condensation level, driving the atmosphere there to be neutrally stable. As discussed above, such a situation is not contradicted by the

probe measurements. On the other hand, we cannot know the actual situation external to hot spots until in-situ measurements are obtained there. Turning to the Jovian upper atmosphere, Fig. 11 shows the thermal structure of the upper atmosphere derived from the ASI accelerometer readings during the probe high speed entry portion of the mission. This region extends from about 1000 km altitude to approximately 23 km altitude, beyond which point the descent portion of the probe mission began and T, p were measured directly (Seiff et al., 1996, 1997a, 1998). Recall that the hydrostatic equation is used to derive pressure from density in the entry phase of the probe mission. Thus a boundary value of pressure or temperature has to be assumed at the upper boundary of entry measurements. The effect of this boundary condition becomes negligible at altitudes below about 600 km. Above 700 km, temperature is clearly affected by choice of upper boundary temperature at 1000 km. However, above some altitude energy deposition should become small and the atmosphere should approach being isothermal, so that there is no net energy transport either upwards or downwards. Apart from the appar-

Fig. 11. Upper atmosphere temperature derived from the Galileo probe ASI accelerometer readings during the entry portion of the probe mission. Altitude z is measured from the 1 bar pressure level. The various curves, which all converge below z5600 km, correspond to different choices of the boundary temperature near z51000 km.

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ent existence of wave-like features in the temperature profile, discussed below, an isothermal state above 800 km is suggested when the upper boundary temperature is chosen to be between 800 and 1000 K. The upper atmosphere ASI results verify the existence of large vertical temperature gradients, reaching a maximum slightly under 5 K / km at about 435 km altitude (Seiff et al., 1996, 1997a, 1998). A derived temperature of 870630 K at 800 km is consistent with choice of the upper boundary temperature between 800 and 1000 K, with the uncertainty rapidly decreasing with decreasing altitude. The stratosphere is located between approximately 300 km (T ¯ 160 K, p ¯ 0.003 mbars) and the tropopause near 30 km (T ¯ 111 K, p ¯ 0.26 bars). Temperature between 100 and 300 km is nearly constant at approximately 160 K. Superposed on both the thermospheric and stratospheric mean temperature profiles are temperature oscillations reaching several tens of degrees K. Young et al. (1997) have shown the thermospheric temperature oscillations to be consistent with a discreet number of gravity waves propagating from below, being damped by molecular viscosity and thermal conductivity above about 350 km. The upper atmosphere temperature profile determined by the probe is consistent with that predicted by Yelle et al. (1996), in particular the large temperature gradient between approximately 350–550 km and temperatures reaching 900 K at the highest altitudes sampled. Having determined the rapid temperature increase with height above the stratosphere and the high upper atmospheric temperatures, the probe results immediately raise the question as to the heating mechanism of the upper Jovian atmosphere. As discussed previously, all the outer planets have hot upper atmospheres, and absorption of solar energy is completely insufficient to account for the heating. At first it seemed that the gravity waves identified by Young et al. (1997) in the probe upper atmosphere data were sufficient to provide the heating. However, further analysis by Matcheva and Strobel (1999) and Hickey et al. (2000) show that these atmospheric waves would not produce the heating required. Waite et al. (1997) show that if observed X-ray emissions near the Jovian equator are due to heavy ion precipitation from Jupiter’s inner

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radiation belt, such ions might contribute about half of Jupiter’s upper atmosphere heating, given assumptions about where the heat is deposited, etc. Within the range of model parameters they consider, it appears that it might be possible to provide all the heating. Although heavy ion precipitation may contribute substantially to heating Jupiter’s thermosphere, it would not be expected to be a significant contributor to the energy budget of the hot upper atmospheres of the other outer planets. Each of their surface magnetic fields is over an order of magnitude smaller than that of Jupiter, and would therefore trap charged particles much less energetic than the trapped precipitating particles that contribute to heating Jupiter’s upper atmosphere. Schubert et al. (2001) have suggested that acoustic waves, perhaps generated by thunderstorm activity, could heat the Jovian thermosphere. However, at the present time it is not known if the amplitudes and energy fluxes are sufficient to provide the required heating, nor is it known if acoustic waves might be possible heat sources for the upper atmospheres of the other outer planets. Thus, if all outer planet hot upper atmospheres including that of Jupiter are the result of the same heating mechanism, that mechanism is yet to be identified.

4.4. Winds As the Galileo probe descended making direct atmospheric measurements, it was also tracked via line-of-sight Doppler tracking from the overflying Galileo orbiter. Because the line of sight between probe and orbiter was substantially vertical, the vertical descent velocity of the probe had to be known in order to derive horizontal winds, which are assumed with well founded justification to be essentially zonal. As it turned out, the date and time of probe entry also allowed the probe signal to be tracked directly from the ground using the Very Large Array (VLA) set of radio telescopes in New Mexico (Folkner et al., 1997). The line of sight from the VLA to the probe was almost entirely along the zonal direction with respect to Jupiter, so that the VLA was mainly sensing zonal winds. Even with the VLA, however, the probe signal at Earth was very difficult to detect, and use had to be made of the probe telemetry symbol stream as relayed by the

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orbiter to the Earth to derive the Doppler shift in the direct probe-Earth signal. Doppler tracking of the probe from the Galileo orbiter clearly established that at the PES the zonal winds observed at cloud levels extend well below the visible clouds, to at least as deep as the approximately 22 bars reached by the probe (Atkinson et al., 1996, 1997, 1998). This is well below both the lowest cloud condensation level and the depth to which solar energy penetrates to any significant degree (cf. Sromovsky et al., 1996, 1998). Fig. 12 shows the zonal wind profiles derived from the orbiter and VLA Doppler data sets. The offset of 30–40 m s 21 between the VLA and orbiter-tracked winds would probably be significantly reduced if the VLA analysis was redone to incorporate the most recent determinations of probe descent velocity, which the VLA analysis uses at the beginning to derive the zonal wind profile. On the other hand, the winds near 1 bar are subject to significant error in the orbiter Doppler tracking method because of the almost vertical orientation of the orbiter-probe geometry, and derived values range from about 80 to 120 m s 21 . At pressures higher than 4 bars the winds level off to a more or less constant speed of about

170 m s 21 with relatively small uncertainties near 13 m s 21 (Atkinson et al., 1998). The VLA measurements (Folkner et al., 1997), which extend to pressures just above 4 bars, reach somewhat higher wind speeds, about 200 m s 21 . There is the possibility, of course, that meridional winds are present and are being interpreted to some extent as zonal winds in the orbiter tracked results. In any event, both Doppler tracking from the orbiter and from the ground using the VLA (Folkner et al., 1997) show that there is substantial vertical shear in the winds between the start of probe descent measurements near 0.4 bars to about 4 bars. The sense of the shear is the same as that deduced from Voyager temperatures and thermal wind balance at 0.15 and 0.27 bars (Gierasch et al., 1986). Analysis of accelerometer data from the ASI has yielded wind profiles that are basically consistent, to within the ranges given above, with the Doppler tracking (Seiff et al., 1997b). Dynamical stability of the vertical shear between upper cloud levels and 4 bars implies that the atmosphere should be statically stable there, consistent with the discussion in the previous section. If the atmosphere were not statically stable in this

Fig. 12. Jovian zonal wind profiles at the PES as derived from Doppler tracking of the probe. Data taken from Atkinson et al. (1998) and Folkner et al. (1997). See text in Section 4.4 for discussion of differences between the curves and uncertainties in the measurements.

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region, Kelvin–Helmholtz (shear) instabilities would rapidly grow and destroy the wind shear. The amount of stability does not have to be large, and is well below the stability derived from the ASI presented in the previous section. Only about 0.01 K km 21 subadiabatic lapse rate is sufficient to produce a Richardson number .0.25, the critical value commonly used to judge the stability of shear flows. If all the stability were due to a molecular weight gradient, it appears that the molecular weight gradient implied by the probe condensible measurements would be marginally sufficient to produce a ˜ et al., 2002). Richardson number .0.25 (Magalhaes Thus, there appears more than enough static stability present to stabilize the wind shear against shear instabilities, except perhaps in the regions where N 2 is small (cf. Fig. 10). However a surprising aspect of the wind shear and static stability is that they seem to be anti-correlated ˜ et al., 2002), compare Figs. 10 and 12. (Magalhaes In other words, in the 0.5–3 bar region where the wind shear is largest, N 2 decreases. Conversely, in the region near 3–5 bars where wind shear de˜ et creases, N 2 appears to increase again. Magalhaes al. (2002) present a more detailed comparison between S and u that illustrates this inverse correlation to an even greater extent. This relationship between the horizontal wind and the static stability is quite different than that which one would expect in general. For example, in the atmosphere of Venus (cf. Gierasch et al., 1997), static stability and vertical wind shear associated with the high speed Venusian planetary scale zonal wind field are directly correlated (i.e., regions of higher wind shear in the mean zonal wind correlate with regions of higher static stability). In the case of Venus, the decreased mean zonal wind shear in regions of low static stability is attributed to increased vertical mixing of momentum due to Kelvin–Helmholtz instability in regions where Ri u 0.25. The origin of the inverse correlation at the Galileo PES is unclear, but it may be related to local dynamical effects associated with the fact that the probe descended into a 5 mm hot spot. An indication that hot spot dynamics could be significantly affecting the winds derived from Doppler tracking comes from the numerical hot spot model of Showman and Dowling (2000). Comparison of the vertical wind profiles computed from their

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model with the probe Doppler wind profile, shows that near the southern boundary of the model hot spot, computed winds are almost entirely zonal, and the vertical variation of the zonal wind resembles that seen in the Doppler wind profile. Since the probe did in fact descend in the southern region of the hot spot it entered, the possibility has to be considered that the probe Doppler winds and wind shear between about 0.4 and 4–5 bars represent mostly dynamical processes in the hot spot rather than the general wind field. The water abundance profile indicates that the probe hot spot extends to at least 20–30 bars. Is the probe wind profile at these depths also more representative of the hot spot than the general wind field? Aside from a probe mission into a non-hot spot region of Jupiter, more detailed modeling of hot spots will be required to resolve the issue. If the deeper winds from the Doppler probe tracking do represent the planetary winds near the equator, the probable implication is that the winds are being driven from below by the internal heat flux. Solar heating becomes negligible at pressures greater than several bars (Hunten et al., 1980; Sromovsky et al., 1998). Latent heat release, if it occurs in clouds outside the hot spot, would likely occur at pressures less than 10 bars, but the actual pressure level would depend on the water abundance. It is not impossible that such forcing could drive deeper winds from above. But if the winds are driven in the cloud layers, dynamical mechanisms must transport momentum from the upper levels downwards over at least 2–3 scale heights, and maybe many more scale heights depending on how deep the winds extend beyond 22 bars. No models have yet been constructed to identify dynamical modes that might accomplish such downward transport, nor are we aware of existing analogous situations. It seems more likely that the energy source for the winds comes principally from the interior heat flux, and that the cloud level winds are the upper extension of deeper convective wind fields, influenced by solar heating and perhaps latent heat release at cloud levels.

4.5. Energy balance The probe entered a region that is bright near

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spectral wavelengths of 5 mm, having a brightness temperature near 5 mm of 252 K (Orton et al., 1998). Most of the 5 mm emission to space from Jupiter comes from 5 mm hot spots located near the probe entry latitude, but the 5 mm emission to space contributes very little to the overall energy balance of Jupiter (cf. Hanel et al., 1981). However as will be discussed below, this is not to say that the 5 mm spectral region could not be important for understanding outward transport of heat flux and differential heating which drives winds. Based on Voyager 1, the effective temperature of Jupiter is 124.460.3 K, and the thermal radiation emitted to space peaks near 35 mm (Hanel et al., 1981). The globally averaged thermal emission is 13.6 W m 22 (Hanel et al., 1981), yielding a ratio of total emitted thermal radiation to absorbed solar radiation of 1.7. Thus Jupiter has an internal heat source comparable in magnitude to the solar heating. The globally averaged internal heat flux is 5.4 W m 22 (Hanel et al., 1981). The question to be addressed by the Galileo probe was how radiative fluxes behaved as a function of atmospheric pressure, i.e., what was the vertical profile of radiative energy transport. The profile directly affects heating of the atmosphere which drives atmospheric motions, and is a factor in understanding the overall energy budget of Jupiter. During the descent portion of the mission the Net Flux Radiometer (NFR) on the Galileo probe measured solar and infrared radiative fluxes at the PES (Sromovsky et al., 1998). Interpretation of measured radiative fluxes must take account of the fact that the probe entered a 5 mm hot spot. Hot spots are bright at 5 mm because relatively little cloud material is present there and gaseous opacity in the Jovian atmosphere is relatively low near 5 mm. However in order to be as bright as measured near 5 mm, certain radiatively active trace species such as H 2 O had to be depleted at the PES over much of the probe descent. For these reasons care must be taken in applying the probe radiative measurements to other regions on Jupiter. The NFR carried two spectral channels to measure solar flux: channel B (0.3–3.5 mm) and channel E (0.6–3.5 mm). The purpose of channel E was to address methane absorption. The net downward solar flux measured by the NFR varies from about 3 W

m 22 at pressure 0.5 bars to less than 1 W m 22 at 5 bars, and is essentially zero by 10 bars (Sromovsky et al., 1996, 1998). Not all of this decrease is due to atmospheric opacity, since at the start of probe descent measurements near 0.4 bars the solar zenith angle was 678, and by the time the probe reached a pressure of 15 bars the zenith angle was 908 due to planetary rotation. Interestingly, by a pressure of 5 bars absorption at wavelengths .0.6 mm was complete, exceeding the expected absorption based on the abundance and opacity of CH 4 . The reason for this has not yet been identified, but may be due to NH 3 (Sromovsky et al., 1998). The NFR also carried three spectral channels to measure longwave thermal flux: channel A (3–200 mm), channel C (3.5–5.8 mm), and channel D (14– 150 mm) (Sromovsky et al., 1998). Channel C sampled the 5 mm window in which gaseous opacity is relatively low, while channel A was a broadband channel to measure thermal radiation as a whole. Channel D sampled the hydrogen dominated part of the thermal spectrum. Beyond 8–10 bars pressure each channel seemed to be affected by the probe internal temperatures discussed previously, so the observations are probably only valid at pressures smaller than about 8 bars. The results from the NFR clearly show that at pressures between 4 and 8 bars, the net upward longwave flux at the PES is confined almost entirely to the 5 mm window. Sromovsky et al. (1998) compute the total net upward longwave radiative flux profile down to 20 bars assuming an atmospheric composition and cloud characteristics consistent with the probe measurements. Between 4 and 20 bars the computed net upward thermal flux is between 3 and 4 W m 22 . This can be compared to the globally averaged internal heat flux of 5.4 W m 22 (Hanel et al., 1981). It should be noted that T( p) at the PES is basically an adiabat. Therefore the NFR observations and derived total thermal flux shows that significant radiative fluxes are possible in particular regions even if convection produces an adiabatic temperature structure. Whether or not the net upward longwave radiative heat flux below upper cloud levels is significant in Jupiter’s energy budget, or in producing differential heating which drives winds, depends on the areal extent over which such radiative fluxes exist, and that depends on the distribution of key trace species.

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At the PES two atmospheric condensibles were depleted between 4 and 8 bars, H 2 S and H 2 O, and the third, NH 3 , had not yet attained its deep mixing ratio. Between 4 and 8 bars NH 3 abundance increased from about 1.5 to 3 times solar (Folkner et al., 1998). The deep mixing ratio of NH 3 is about 3.5 solar, and at that mixing ratio produces the same or greater opacity as CH 4 in the 5 mm window. The measured deep global mixing ratio of H 2 S would provide relatively small opacity in the 5 mm window, and hence H 2 S is not a significant factor in understanding the radiative fluxes. The deep mixing ratio of H 2 O is undetermined, and at the PES H 2 O was highly depleted to pressures of at least 12 bars. Even between 15 and 20 bars the H 2 O abundance is less than solar (Niemann et al., 1998; Atreya et al., 2002; see also Section 4.1.4, Fig. 7). Given Jovian abundances of other trace species, in the 5 mm window H 2 O is the constituent limiting the radiative flux. If the abundance of H 2 O is i solar, the gaseous opacity in the 5 mm window would be dominated by H 2 O, sufficient to effectively close the window. On the other hand, if there exist regions below the upper clouds where H 2 O is depleted, the NFR results show that net upward longwave radiative fluxes associated with an adiabatic temperature profile can reach a significant fraction of the flux corresponding to the globally averaged internal heat flux. The 5 mm hot spots cover only about 1% of Jupiter’s area, so if H 2 O is depleted only in hot spots radiative transport below cloud levels would not be significant for global transport of internal heat flux, and radiative thermal forcing from below probably would not be significant for driving winds. Guillot et al. (1994) investigated the radiative transport of energy from Jupiter’s deep interior to near cloud levels, and concluded that unless H 2 O is depleted, the radiative equilibrium temperature profile near and below cloud levels is super-adiabatic, implying that convection occurs. The NFR results show that even if convection does occur and drives the temperature profile to being adiabatic, depending on the abundance of H 2 O, radiative fluxes can still be significant. If the H 2 O abundance is sufficient to form thick water clouds, there might be regions above and between the clouds where water vapor could be depleted and radiative transport be important. A-priori it would be expected that water vapor

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would be well-mixed below cloud levels, but everyone was surprised by the probe measurements, and until we have a good understanding of the depletion of H 2 O at the PES and composition measurements elsewhere, the question remains open regarding the distribution of H 2 O in other regions of Jupiter. Although there are indirect indications of water clouds near the Great Red Spot, in zones including the EZ, and the NEB (see discussion of H 2 O in Section 4.1.4), no direct detection of thick water clouds exists yet. Guillot (1995) has discussed the possibility that diffusive convection can act to dry out the atmosphere below the saturation level of a condensible. Diffusive convection in the outer planets would occur as a result of radiative heat transport down a vertical superadiabatic temperature gradient coupled with diffusion of condensible species (in the Jovian case, water) in an environment in which mean molecular weight increases with depth (recall that unlike the Earth, on the outer planets the presence of a condensible vapor increases the mean molecular weight of the atmosphere). If diffusive convection were wide spread, it could be the case that Jupiter’s atmosphere would be dry well below the expected condensation level for water over regional horizontal spatial scales, in which case radiative transport in the sub-cloud troposphere could contribute significantly to Jupiter’s outward heat flux at pressures ¯5–10 bars.

5. Summary of Galileo probe results and outstanding issues The Galileo probe data lived up to expectations that direct sampling of Jupiter’s atmosphere could fundamentally affect our understanding of Jupiter, the outer planets, and the formation and evolution of the solar system. The Galileo probe established the global abundances of key atmospheric elements in Jupiter’s atmosphere, namely He, C, N, S, Ne, Ar, Kr, and Xe, together with many of their isotopic ratios. With the exception of He and Ne, all these elements have been found by the probe to be enriched in the Jovian atmosphere with respect to solar abundance by a factor of ¯3. As discussed in Section 4.1.2, these results are perplexing in many ways, and require substantial revisions to current

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views of how the outer planets formed and how the solar system evolved during and after planet formation. Isotopic ratios were measured to be basically the same as those for the Sun, indicating that as Jupiter formed from the solar nebula, there were not processes which isotopically fractionated the material forming the planet to any significant degree. The Jovian helium mass fraction was measured to be 0.23, about 30% greater than the value derived from Voyager data, and therefore closer to the protosolar value of 0.27 than previously thought. Thus, there has been less helium settling towards the deep interior of Jupiter than implied by the Voyager results. On this basis, theoretical Jovian interior models can now produce results more consistent with the known constraints of planet mass, radius, gravitational moments, and internal heat flux. However, one of the principal outputs from the interior models is a computed mass of the Jovian core, which is a crucial quantity for understanding how Jupiter formed (see Section 2.1). Currently, even using the updated He Jovian mass fraction, the interior models are consistent with a core mass between 0 and 14 M % because of uncertainty in the hydrogen–helium equation of state. Ne was measured to be depleted in the Jovian atmosphere, having an abundance relative to solar of 0.1. The reason for the Ne depletion is probably that Ne is soluble in liquid helium, and is removed as helium in the liquid form settles towards the deep interior (see Section 4.1.2). The Galileo helium result for Jupiter has also led to a re-analysis of the Voyager helium mass fraction for Saturn, resulting in an increase of the derived Saturnian value from about 0.06 to between 0.18 and 0.25 (see Section 4.1.1). The helium mass fractions for Jupiter and Saturn now appear much less different than they did previously. One result has been that interior models of Saturn, which are based on the same physics as the Jovian models, now yield enrichments in heavy elements that are reasonable, and which are consistent with spectoscopic data of Saturn’s atmosphere, whereas before the models did not. The probe determined the thermal structure at the PES from approximately 1000 km above the 1 bar pressure level (10 29 bars) to 132 km below 1 bar (22 bars). In the troposphere the thermal profile is very close to an adiabat that passes through 260 K at 4.18

bars. Nevertheless, the probe measurements showed that the troposphere is generally stably stratified at the PES as deep as the probe made measurements, with a typical static stability of ¯0.1 K km 21 . In the upper atmosphere the probe data indicate a maximum positive vertical temperature gradient of about 5 K km 21 near 435 km altitude (height above the 1 bar pressure level), and a maximum temperature of ¯900 K at 800 km altitude. Thus, Jupiter has a hot upper atmosphere with regions of large positive vertical temperature gradients. No thick clouds were detected by the probe, and in particular, no thick water cloud. However, the entry of the probe into a 5 mm hot spot clearly affected the probe measurements regarding clouds, and the cloud structure at the PES cannot be considered to be representative of Jupiter’s cloud layers in general. Doppler tracking of the probe, both from the orbiter and from the ground using the VLA set of radio telescopes, showed that as the probe descended in the troposphere it encountered zonal winds the order of 100 m s 21 , rising to and leveling off at ¯170 m s 21 at pressures greater than 4–5 bars. Thus, at the PES, the probe determined the vertical zonal wind profile, and established that large winds extend well below the visible clouds. As in almost all planetary exploration missions, new questions arise from the data, and some puzzling aspects of the returned data cannot be resolved without either further missions or modeling studies, or perhaps both. The major such issues from the Galileo probe mission are the following. Although the probe determined the abundances of almost all the key elements of interest, the major missing element is oxygen. Because of the probe entry location, the deep mixing ratio of H 2 O, which presumably represents the global abundance of O in at least the molecular hydrogen outer envelope of Jupiter, was not determined. Unlike the other condensible species, NH 3 and H 2 S, the water abundance did not reach a steady value with depth even at the end of the probe mission near 22 bars. It is tempting to assume that O is also enriched to about 3 times solar, but as discussed in Sections 4.1.2 and 4.1.4, there are plausible arguments as to why O could be even more enriched in Jupiter than are the other heavy elements. On the other hand, aside from some indirect indications that substantial water clouds exist

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at least regionally, the possibility that O is depleted cannot be ruled out. This of course would be very difficult to understand in light of the enrichment of the other heavy elements. In all likelihood, resolution of this question will require another mission to Jupiter, perhaps with multiple probes to ensure that a representative sample is obtained. Because of the implications for the formation and evolution of Jupiter, determination of the global O abundance is probably the most important of the unresolved issues left from the Galileo probe mission. Another unresolved issue is why the elements so far measured by the probe that are heavier than helium (excepting O and Ne as discussed) should each be enriched over its solar abundance by a factor of ¯3 in Jupiter’s atmosphere. The possible scenarios suggested to date (discussed in Section 4.1.2) by which such enrichments might occur all have substantial difficulties. As with the oxygen abundance, resolution of this issue will almost certainly have considerable impact on our ideas of solar system evolution and planet formation. Related to the question of the abundance of H 2 O is the question of the strange behavior of all the condensibles, NH 3 , H 2 S, and H 2 O at the PES. All these condensibles exhibited the same qualitative behavior: depleted abundance at depths well below their respective condensation levels, increasing abundance with depth at different fractional rates, and asymptotic and approximately constant mixing ratio with depth at large pressures. In the case of H 2 O the probe did not reach depths sufficient to measure the asymptotic mixing ratio. For these reasons we are left without a clear picture of the cloud structure on Jupiter, but it may be that no single probe would provide a globally accurate picture. Although there are promising models and ideas of how to understand the vertical profiles of condensibles at the PES, at the present time there is not yet an adequate theoretical description. No model can yet reproduce the probe observations. Directly connected to the above discussion is the relationship of 5 mm hot spots to global processes. It would appear that without understanding the general processes that produce hot spots there is little prospect of being able to fully understand some important aspects of the probe measurements. That the hot spots are part of planetary scale processes is

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evident from the fact that, although confined near the equator (presently existing only in the northern hemisphere), they are generally distributed more or less evenly in longitude around the planet, move together as a group in the zonal direction, and have lifetimes the order of a terrestrial year. They do not appear to be isolated features. Even if large-amplitude atmospheric waves can be shown to generate the hot spots and be capable of inducing the observed condensible species vertical profiles, one must ask what acts to generate such large amplitude waves. Perhaps moist convection could be the driving mechanism if there is sufficient H 2 O, and if so this would also have implications for Jupiter’s dynamic meteorology as a whole. In considering hot spots, the question arises as to whether the zonal wind, and in particular the vertical shear in the zonal wind, measured by the probe represents the general wind field or a more localized manifestation of hot spot dynamics. Recent hot spot model simulations would seem to indicate that at least the vertical shear could be the result of dynamical processes associated with the probe entry hot spot. Thus, we are left with some ambiguity as to the degree to which the probe provided insight into the general Jovian wind field. Because of the above considerations one might argue that the chance descent of the probe through one of the 5 mm hot spots was unfortunate. On the other hand, it is also true that much of past and future remote sensing of Jupiter’s deeper atmosphere has or will be done through hot spots. Thus, as understanding and modeling of hot spots improves, the ground truth provided by the probe measurements will improve analysis of remote sensing measurements of Jupiter’s atmosphere at depths below cloud levels. One aspect of the probe measurements still not theoretically explained and not expected to be related to hot spots is the high temperature of the Jovian upper atmosphere. No current models have clearly established a mechanism that is sufficient for heating Jupiter’s upper atmosphere, let alone identifying mechanisms sufficient for heating the upper atmospheres of the other outer planets. The charged particle precipitation models that currently indicate significant heating in the Jovian upper atmosphere do not appear to be viable on the other outer planets.

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Ideally, one would like to construct a model that could explain not only Jupiter’s hot upper atmosphere, but also each of those of the other outer planets. So we are still left with an interesting puzzle. In spite of such unresolved issues as mentioned above, from a large perspective the Galileo probe mission was an immense success. The probe observations have forced us to think anew about how Jupiter formed and evolved, with all the implications that may have for understanding the entire solar system. As seems to be typical of planetary missions, there were several major surprises, surprises which will require considerable future work to explain and understand. No doubt the Galileo probe results will alter the framework by which we view the outer planets and the solar system, which is what the probe mission was designed to do.

Acknowledgements We wish to thank John Chambers, Jeffrey Cuzzi, Richard Freedman, Jack Lissauer, Mark Marley, and Kevin Zahnle for many very helpful discussions on a number of topics. We also wish to thank the referee, Tobias Owen, for making several valuable suggestions and points which have significantly improved the paper.

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