FIRST VERTICAL ION DENSITY PROFILE IN JUPITER S AURORAL ATMOSPHERE: DIRECT OBSERVATIONS USING THE KECK II TELESCOPE

Transcription

1 The Astrophysical Journal, 677:790Y797, 2008 April 10 # The American Astronomical Society. All rights reserved. Printed in U.S.A. FIRST VERTICAL ION DENSITY PROFILE IN JUPITER S AURORAL ATMOSPHERE: DIRECT OBSERVATIONS USING THE KECK II TELESCOPE M. B. Lystrup and S. Miller Department of Physics and Astronomy, University College London, Gower Street, London WC1E 6BT, UK; makenzie@apl.ucl.ac.uk N. Dello Russo and R. J. Vervack, Jr. The Johns Hopkins University Applied Physics Laboratory, Johns Hopkins Road, Laurel, MD and T. Stallard Department of Physics and Astronomy, University of Leicester, University Road, Leicester LE1 7RH, UK Received 2007 August 31; accepted 2008 January 10 ABSTRACT We present the first vertical ion density profiles of Jupiter s upper atmosphere derived directly from ground-based observations. Observations of infrared H þ 3 emissions in Jupiter s auroral/polar regions were collected by the highresolution spectrometer NIRSPEC on the Keck II telescope. We have calculated vertical density profiles for a latitude in the southern auroral region using the measured column densities and a shell model of the Jovian ionospheric H þ 3 emission. We compare our resultant profiles to those generated by a recent one-dimensional model in both local thermodynamic equilibrium (LTE) and non-lte conditions. We find good agreement with the model profiles up to 1800 km. Above that, however, our measurements show that more H þ 3 is produced than is predicted by the model. Our observational method is a new tool for probing Jupiter s upper atmosphere from Earth and can possibly be extended to the study of other gas giant planets. Subject headinggs: infrared: solar system methods: data analysis planets and satellites: individual (Jupiter) techniques: spectroscopic 1. INTRODUCTION Above the Jovian homopause (generally taken around the Jupiter s ionosphere plays an important physical role in both bar pressure level), individual atmospheric components settle out diffusively according to their own scale heights, which are the energy balance of the planet and the coupling to the magnetosphere. Understanding the ionosphere-magnetosphere coupling inversely proportional to their respective masses. Thus H and H requires the determination of the related properties of the ionosphere, such as conductivity and energy inputs. Of particular im- 3 results from 2 (and He) predominate and hydrocarbon densities, for example, become negligible with increasing altitude. H interactions between H portance are the auroral/polar regions of the planet, where various 2 and H þ 2.H 2 is ionized by either energetic electrons (e magnetospheric processes give rise to signature atmospheric ) that are precipitated into the upper atmosphere along magnetic field lines or by solar extreme ultraviolet emissions, especially in the UVand IR wavelength regions of the radiation (h): spectrum. Observations of the vertical electron and ion profiles form a key component of solutions to these questions. There have been many attempts to model electron and ion density H 2 þ e! H þ 2 þ e þ e; ð1þ profiles (e.g., Waite et al. 1983; McConnell & Majeed 1987; Kim et al. 1992; Achilleos et al. 1998) with varying degrees of H 2 þ h! H þ 2 þ e: ð2þ success. Models have consistently shown that the H þ 3 molecular These energetic electrons are abundant in the auroral/polar ion is a major component of the ionosphere at all altitudes, although H + 3 regions of the planet. H þ is subsequently produced via predominates at higher altitudes and on the night side. One recent self-consistent one-dimensional (1D) model of the auroral profiles (Grodent et al. 2001) made use of the in situ measurements H 2 þ H þ 2! Hþ 3 þ H; ð3þ of Galileo (Seiff et al. 1998) to constrain the vertical which is an exothermic reaction releasing about 1.7 ev. In the neutral density profile. The model generated the temperature profile derived from UV and IR measurements of the auroral regions almost instantaneously. The ion is mainly destroyed by Jovian upper atmosphere H þ 3 is formed copiously and rapidly, by fluxes of precipitating electrons with a spectrum of energies from ev to kev. Thus the electron and ion densities are some of the most reliable to have been calculated. H þ 3 þ e! H 2 þ H; ð4þ It has been thought that only in situ instruments could provide measurements of such profiles. Pioneer 10, Pioneer 11, and H þ 3 þ e! H þ H þ H: ð5þ Voyager obtained electron density profiles via radio occultation This shows that the lifetime of H þ 3 depends on the electron observations (Kliore et al. 1974; Fjeldbo et al. 1976; Atreya et al. density. It is on the order of 10Y100 s in the lower auroral ionosphere, 1979; Eshleman et al. 1979). Atmospheric ion profiles have so far not been measured in situ. We present a new method for measuring ion profiles from a ground-based instrument. where peak electron densities reach Y10 12 m 3, rising to >1000 s at lower latitudes and higher altitudes where electron densities are orders of smaller magnitude. 790

2 JOVIAN ION DENSITY PROFILE 791 Fig. 2. General auroral configuration for exposure 138 at CML ¼ 141:66. The solid curve is the 30 R J oval and the crosses are the Io footprint, both from the VIP4 model of Connerney et al. (1998). Fig. 1. Broadband image of Jupiter s southern hemisphere showing slit position for spectrum 138. The slit on the image has been traced over in white to aid the eye. NIRSPEC s slit was positioned perpendicular to Jupiter s rotational axis during the observation. H þ 3 emission from Jupiter s auroral regions was discovered nearly 20 years ago (Drossart et al. 1989) and has since often been used as a probe of the planet s upper atmosphere. Many studies (e.g., Rego et al. 1999; Stallard et al. 2001, 2004; Melin et al. 2005; Lystrup et al. 2007) have used H þ 3 infrared emissions in the auroral regions and lower latitudes to probe the energetics and dynamics of Jupiter s upper atmosphere. Infrared H þ 3 emission can be considered an indicator of processes linking the ionosphere and magnetosphere. In addition to its usefulness as a probe of these processes, the ion itself plays an important role in Jupiter s upper atmosphere. The ion, which is the most abundant component of the lower Jovian ionosphere, acts as a coolant for Jupiter s thermosphere, reradiating the energy from precipitating particles back into space (Miller et al. 2006). Previous ground-based observational studies have provided column density measurements of H þ 3 emission in the auroral regions and on the disk of the planet (e.g., Lam et al. 1997; Miller et al. 1997; Stallard et al. 2002; Raynaud et al. 2004). These studies assumed that a state of quasi local thermodynamic equilibrium (LTE; Miller et al. 1990) prevails in the Jovian upper atmosphere. At the peak of the H þ 3 density (calculated to be about 500 km above the 1 bar pressure level) the time between collisions is much shorter than the radiative lifetime of the measured rovibrational transitions (on the order of 10 1 to 10 2 s), which is itself much shorter than the overall lifetime of the ion. This allows for collisions to thermalize the H þ 3. However, in the upper reaches of Jupiter s thermosphere where the atmosphere is very rarefied, conditions for LTE do not hold. As a result, the vibrational ground state is overpopulated compared with higher vibrational levels and the H þ 3 densities derived from IR measurements may underestimate the real density. Melin et al. (2005) used a detailed balance calculation (Oka & Epp 2004) to incorporate the effects of non-lte conditions into the Grodent et al. (2001) model. Melin et al. (2005) then used H þ 3 emission observations to investigate the effects of non-lte conditions in Jupiter s upper atmosphere. In this paper we make use of the Grodent et al. (2001) and Melin et al. (2005) work and we present the first H þ 3 local densities as a function of altitude, derived from H þ 3 emission measured off the limb of the planet. 2. OBSERVATIONS, DATA REDUCTION, AND ANALYSIS 2.1. Observations The data presented here were collected on 2006 May 14 using the NIRSPEC echelle spectrometer ( McLean et al. 1998) on the Keck II telescope on Mauna Kea. NIRSPEC was configured to detect a wavelength range of about 3.26Y4.01 m, distributed over 5 orders (NIRSPEC orders 19Y23), within which many emission features from the H þ 3 molecular ion lie. Seeing on the night was NIRSPEC s by slit, with a resolving power of k/k 37; 500, was rotated perpendicular to Jupiter s rotational axis (which had a position angle of on May 14) such that it cut east-west across the planet. The Jovian equatorial diameter on May 14 was At high latitudes (beyond 65 )the24 00 slit was long enough to cut across the entire planet, as seen in Figure 1. In these regions, a limb-to-limb horizontal profile of the entire aurora could be observed, as well as significant emission above the limb of the planet (as defined by the cloud tops at the 300 mbar pressure level). Equatorward of the auroral region, however, the slit (NIRSPEC s longest for an appropriately small slit width) was not long enough to cut across the entire planet. The slit was initially positioned at the northern polar limb of Jupiter. After a spectrum was recorded, a corresponding image was taken showing the position of the slit on the planet. The slit was then moved in steps of about 1 00 toward the equator. The same pattern was repeated in the southern hemisphere. Each spectrum had an integration time of 15 s and two co-adds, resulting in a total exposure time of 30 s for each spectrum. Observations were made in AB and ABA nod-to-sky patterns (where A is an object exposure and B is a sky exposure). For each A exposure an image was taken with 0.1 s integration time and 10 co-adds, resulting in a total exposure time of 1 s for each slit image. These images are useful for checking the slit position but are too broadband to image the aurora itself in detail. In this paper we consider only 1 order of these data (order 19), the wavelength range of 3.944Y4.004 m. It is within this wavelength region that some of the brightest H þ 3 emission features lie. Spectra of an A0 V standard star, HD , were recorded at the end of the Jupiter observations using the wide slit with 2 s integration times and 30 co-adds and using the wide slit with 1 s integration and 5 co-adds. The wider slit was used in order to collect as much light from the star as possible for flux calibration purposes. Although spectra were recorded across the body of the planet, in this study we make use of only exposure 138, taken at 78 south ( latitude of the slit location at the central meridian longitude [CML]) and CML at an airmass of Figure 2 shows that the slit cut across the auroral oval in the midafternoon to dusk sector. The resulting spectral image is seen in Figure 3. Jupiter s infrared continuum is largely, but not entirely, suppressed by stratospheric methane absorption, thus allowing the bright H þ 3 features to be observed from Earth. Strong lines due to transitions in the fundamental 2 Q branch are evident, as

3 792 LYSTRUP ET AL. Vol. 677 Fig. 3. Spectral image for exposure 138. Clear H þ 3 emission is seen extending many pixels beyond the body of the planet. This emission originates in the ionosphere. Q branch fundamental H þ 3 emission lines are labeled; full Hþ 3 emission feature assignments are given in Table 1. The continuum emission is an indicator of the planetary limbs. well as 2 2! 2 hotband lines, including several that have not previously been observed outside the laboratory (see Table 1). These lines are all brightest in the midafternoon to dusk sector, as expected, which corresponds to the main auroral oval. Lowintensity continuum emission from the body of the planet can also be seen in the spectral image; this continuum emission is useful, as we shall discuss later Reduction and Spectral Fitting Wavelength calibration, sky subtraction, and order straightening were completed using the REDSPEC reduction package. 1 The flux calibration was performed separately using the recorded spectra of the standard star. A Gaussian function with FWHM equal to the slit width was used to model the star and calculate the fraction of the star flux entering the wide slit. The H þ 3 line fitting program of Melin (2005) was used to calculate the temperature, column density, and total emission of the H þ 3. In brief, the program uses Hþ 3 transitions from Neale et al. (1996) and the H þ 3 partition function of Neale & Tennyson (1995) and iteratively creates a theoretically predicted H þ 3 spectrum, matching this to the observed spectrum exploiting a minimizing technique according to Cramer s rule. (Further details can be found in Melin et al ) The result is a fitted spectrum and calculated temperatures, column densities, total H þ 3 emission, and errors for these parameters. The procedure was performed for each spatial position on the spectral image for which H þ 3 emission was strong enough for the program to produce a reliable fit. To improve signal to noise we used a three-pixel 1 Details regarding the REDSPEC reduction package may be found at www2.keck.hawaii.edu /inst /nirspec/redspec/index.html. (equivalent to ) rolling boxcar average. This results in an effective seeing of according to qffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffi seeing ea ¼ seeing 2 þ rolling average 2 : ð6þ 2.3. Determination of Planetary Latitude The Jovian surface is generally defined as the 300 mbar pressure level, which corresponds to the tops of the visible cloud decks. The planetary limbs (coincident with this surface) can be seen in the images from NIRSPEC s slit imager SCAM and in the extent of the continuum infrared emission from the body of the planet on the spectral images (Fig. 3). When Jupiter is not near opposition the geometry is such that one limb is illuminated and the other is not. This defect of illumination causes images of the planet to show a terminator that does not coincide with the actual planetary limb. Jupiter was at opposition on 2006 May 4, just 10 days prior to these observations, so this defect of illumination is negligible ( ) and we can ascertain the limb position from the terminator on the spectral image within about 1 pixel. The positions of the planetary limbs, as given by the terminators on the spectral image, are then used with the known pixel resolution of NIRSPEC to provide a measure of the width of the planet at the instrument s slit position. Assuming Jupiter is an oblate spheroid, Jupiter s disk appears as an ellipse with known major and minor axes (the equatorial and polar diameters, respectively). Thus the planetary latitude (at the CML) at which the spectrum was taken can be calculated, accounting for the sub- Earth angle of The nonzero sub-earth angle results in the slit sampling a range of latitudes; for the purposes of our calculations in the following section we use the latitude at the limb, which differs from the latitude at CML by 4. H þ 3 TABLE 1 Emission Features Feature: Band! E 0 (cm 1 ) J 0, G 0, U 0 J 00, G 00, U 00 (cm 1 ) E 00 (cm 1 ) A ij (s 1 ) R(5,2) a :2 2 (2) 2 (1) , 2, +2 5, 2, R(3,4): 2 2 (0) 2 (1) , 4 3, 4, Q(1,0): 2 (1) , 0, 1 1, Q(2,1): 2 (1) , 1, 1 2, R(1,1) a :2 2 (2) 2 (1) , 1, 2 1, 1, R(2,1) a :2 2 (2) 2 (1) , 1, 2 2, 1, R(3,1) a :2 2 (2) 2 (1) , 1, +2 3, 1, Q(3,0): 2 (1) , 0, 1 3, R(4,1) a :2 2 (2) 2 (1) , 1, +2 4, 1, Q(3,1): 2 (1) , 1, 1 3, Q(3,2): 2 (1) , 2, 1 3, Note. After Kao et al. (1991). a Hotband emission features not previously observed outside the laboratory.

4 No. 1, 2008 JOVIAN ION DENSITY PROFILE 793 Fig. 5. Extracted spectra for exposure 138. At the top is the spectrum for the outermost position used in the density profile calculation and at the bottom is the spectrum for the position just beyond the limb, the innermost position used in the calculation. Ticks on the y-axis of each plot indicate the average background and the peak of the brightest emission line. Fig. 4. Temperature, column density, and total H þ 3 emission for exposure 138. These parameters are derived from the Melin et al. (2005) H þ 3 fitting program. 3. RESULTS AND DISCUSSION 3.1. Temperature, Column Density, and Total Emission Profiles across the Planet Figure 4 shows the temperatures T(H þ 3 ), column densities N(H þ 3 ), and total emission E(Hþ 3 ), along the spatial direction of the slit, as fitted by our routines. In each of the three panels, the dawn limb (planetary west) is on the left and the dusk limb (planetary east) on the right. The top panel shows T(H þ 3 )compared with N(H þ 3 ). The plot suggests a strong anticorrelation between the two parameters that is well outside of the error bars of our fitting. This may be explained as follows. The largest column densities are seen in the midafternoon to dusk sector corresponding to the auroral region (Fig. 2). In this region high energy (10Y100 kev) precipitating electrons penetrate deep into the Jovian thermosphere, into the pressure region between 1 and 0.1 bar where the temperature is considerably lower than the exospheric temperature, around 900Y1000 K (Lam et al. 1997; Grodent et al. 2001). There they generate high concentrations of H þ 3. The temperature that we derive in this region is 1050 K. The corresponding value of N(H þ 3 ) is around ð3:7y4:1þ ; m 2. (Note that this value is not corrected for the line of sight.) It should also be noted that the very small error bars for the line-of-sight temperatures and column densities are due the high signal-to-noise ratio (S/N) in our data. With such high S/N the constraints on these two parameters are tight, at least until regions high in the atmosphere (very far from the limb of the planet) are reached and the S/N is smaller (Fig. 5). Westward (on the planet) of the main auroral oval in the dawn to midafternoon sectors the temperature climbs to around 1350 K and remains fairly stable at that value all the way to the dawn limb. (Note that the wavelike structure in the temperature

5 794 LYSTRUP ET AL. Vol. 677 Fig. 6. Diagrams of experimental setup. Left: Face-on schematic (not to scale) diagram of Jupiter (in dark gray with dotted line indicating its rotational axis) with the slit cutting across the southern auroral/polar region of the planet. We have defined a shell of H þ 3 emission for each pixel (n ¼ 1toN ) that extends beyond the planet s dusk limb. In this schematic only a few shells (light gray shading) are shown. Each pixel corresponds to 462 km at Jupiter; thus 462 km is the thickness of each shell. Right: Schematic diagram of a slice of Jupiter (in dark gray with black dot indicating the planet s rotational axis coming out of the page) and the orientation of the slit. The perspective is looking down on the slice from above the northern rotational pole. Each pixel (n ¼ 1toN ) beyond the body of the planet samples emission from various altitudes along the line of sight (example lines of sight are shown as dotted lines). Only pixels extending beyond Jupiter s dusk limb are marked; the rest of the slit has been blacked out. profile on the body of the planet is an artifact of the fitting routine failing to account completely for the Jovian continuum background. This is not important for this paper but future publications will address it.) Here the column densities are smaller than in the auroral region: they are 10%Y20% of the auroral densities. Although it is at the same geographical latitude as the duskside aurora, this region corresponds to lower magnetic latitudes than the auroral oval (Connerney et al. 1998). H þ 3 is therefore produced by a combination of nonauroral lower-energy particle precipitation and solar extreme-uv ( EUV) ionization. As a result the ions are produced higher in the thermosphere, closer to the exosphere, where the temperature is also higher. As the slit crosses the dusk limb and extends off the planet itself, T(H þ 3 ) increases to around 1450 K and N(Hþ 3 ) falls rapidly to just a few percent of the maximum. This is the region of chief interest for this paper since we are interested in deriving an altitude profile of local ion densities (H þ 3 ), and the corresponding temperatures. For the purposes of our analysis, we have set the upper limit of our altitude profiles to the point at which T(H þ 3 ) begins to fall monotonically. This temperature drop may be a real feature, indicating a region in the atmosphere which is being cooled adiabatically by expansion ( D. Strobel 2007, private communication). Although Jeans escape is thought not to be significant at Jupiter because the exospheric temperatures are not high enough to overcome the planet s deep gravitational potential well, hydrodynamic escape from Jupiter may be possible ( D. Strobel 2007, private communication) as has been found with Titan (Strobel 2008a) and Pluto (Strobel 2008b). Nonthermal escape processes may also be important. Hamilton et al. (1980) did find H þ 3 in the Jovian magnetosphere and very recently the New Horizons mass spectrometer detected H þ 3 in the Jovian magnetotail (McComas et al. 2007). The Jovian ionosphere is almost certainly the source for this magnetospheric H þ 3. The region of falling temperatures corresponds to altitudes higher than those normally considered by ionospheric models (5000 km); we will not consider it further in this study. Consistent with previous findings by Raynaud et al. (2004), the plots also indicate that the total emission E(H þ 3 )ismost strongly correlated with N(H þ 3 ), but not on a simple one-to-one basis. This is because the higher temperatures associated with lower values of N(H þ 3 ) result in a greater emission per molecule, partially compensating for the diminishing numbers of ions. Thus the peak emission persists across the auroral region and up into the higher altitude/ higher temperature regions on either side Calculation of Vertical Ion Density Profiles The fitted column densities for pixels extending from the limb of the planet outwards reported in the previous section provide the starting point for deriving our vertical local density profiles. The NIRSPEC pixel angular resolution corresponds to 462 km at Jupiter. Pixels that are off the limb of Jupiter correspond to an altitude of n ; 462 km ; cos, where n is the pixel number (with n ¼ 1 at the limb), and is the observer angle, which depends on the latitude and the sub-earth angle. Because the sub-earth angle was not zero for these observations, the planetary latitude that the slit samples is not constant across the length of the slit. For the purposes of this study we have taken the latitude a which the slit crosses the limb as our latitude in our definition of the observer angle. We model the ionosphere of the planet as a series of concentric shells enveloping the planet, with the shell thickness corresponding to one pixel. This type of onion peel approach has successfully been used previously for Earth observations (e.g., Sofieva et al. 2004). Figure 6 shows a simplified view of these shells (only six are shown) as seen from the observer s perspective. Since the shells are oblate spherical, each individual spectrum along the slit represents emission that originates from a range of altitudes along the line of sight to the observer. The total number of pixels used in calculating the density profile presented herein, from the limb of the planet outward to where the H þ 3 emission can no longer be detected, is 22. At each pixel in the spatial direction on the array we measure a column density N(H þ 3 ) that can be thought of as the sum of

6 No. 1, 2008 JOVIAN ION DENSITY PROFILE 795 local density contributions from each shell multiplied by the lineof-sight path length through that shell. Inverting this relation, we can derive local densities in each shell. At the top of the ionosphere we are sampling only through one shell along the line of sight. For practical purposes, we assume that the top of the ionosphere coincides with the last pixel before the derived temperature begins to fall. For this outer pixel the measured column density N(H þ 3 ) is (with respect to Fig. 6 in which the outermost pixel is pixel six) N 6 H þ 3 ¼ 6 H þ 3 l6;6 ; ð7þ and subsequent densities are given by N 5 H þ 3 ¼ 5 H þ 3 l5;5 þ 6 H þ 3 l5;6 ; ð8þ and so on for pixels down to n ¼ 1. We calculate the path lengths through each shell and invert these relations to derive the local densities (H þ 3 ). In practice, this process results in (H þ 3 ) for the outermost pixel being slightly too large, as it contains a small contribution duetoweakh þ 3 emission from shells that we ignore. With this caveat, we can now calculate (H þ 3 ) recursively inward toward the limb of the planet using the process described above. The result is a vertical profile of H þ 3 local density centered on a specific planetary latitude at the given CML. A vertical temperature profile was created in a similar way (Fig. 9). The initial measured temperatures T 0 (H þ 3 ) are density weighted along the line of sight. The temperature in the outermost shell is taken to be the average in that shell. T 6 H þ 6 H þ 3 ¼ 3 l6;6 T6 0 Hþ 3 6 H þ ; ð9þ 3 l6;6 which is simply the measure temperature for the outermost shell. Subsequent temperatures are given by T 5 H þ 5 H þ 3 ¼ 3 l5;5 T5 0 Hþ 3 þ 6 H þ 3 l5;6 T6 0 Hþ 3 5 H þ 3 l5;5 þ 6 H þ ; ð10þ 3 l5;6 and so on for pixels down to n ¼ 1 where the T 0 (H þ 3 )aretheshell temperatures and the T(H þ 3 ) are the line-of-sight average temperatures. This profile is shown in the bottom plot of Figure Comparison of Derived Local Density Profile with Model The vertical profile of the local density derived from our data is shown in Figure 7 as a solid line. For comparison, we show the model of Grodent et al. (2001) as a dashed line. It is difficult to fix the planetary limb better than 1 pixel,sothereissome uncertainty in our altitude scale. We have therefore shifted our measured profiles in both Figure 7 and Figure 8 such that the altitude of the peak density matches that of the Grodent model peak density. Thus our profile provides an independent measure of the topside ionosphere. Figure 7 shows that our peak H þ 3 density is approximately 2.5 times smaller than the Grodent et al. (2001) model. Figure 7 also indicates that at the top of the ionosphere, at an altitude around 3500 km, our profile appears to show local ion densities around 50 times smaller than that of Grodent et al. (2001). However, as discussed in Melin et al. (2005), the excited vibrational levels are underpopulated at high altitudes due to the breakdown of LTE. As a result, our derived local densities underestimate the actual (H þ 3 ) values. Fig. 7. Comparison of measured and model profiles. The measured vertical density profile (solid line) is plotted with the predicted profile from the Grodent et al. (2001) model (dashed line) and the predicted profile resulting from the non- LTE correction to the Grodent model (dotted line). In the middle ( log scale) and bottom ( linear scale) plots the measured profile has been scaled to the maximum of the Grodent model and the peak density has been shifted to match the Grodent model peak density. Thus a more meaningful comparison is to scale the model profile to take account of non-lte effects; this is also shown in Figure 7 (dotted line). The effective (H þ 3 ) densities produced by non-lte scaling of the model profile are much smaller at higher altitudes than our profile measures. This effect is clearly seen in the comparison of normalized profiles shown in the middle and lower panels of Figure 7. (It is also worth noting that our measured profile falls off smoothly with altitude and does not show the rather constant ion densities of the model profile between 1100 and 1400 km. This feature in the model profiles may be an artifact due to the double Maxwellian electron flux that Grodent et al. [2001] employed.) There is, however, another factor that must be taken into account in comparing the model and measured profiles: the effective seeing that tends to increase the apparent value of (H þ 3 ) at higher altitudes. The combination of actual seeing conditions on the night and the rolling average we have taken to improve signal to noise gives an effective seeing of Deconvolving the seeing from our profile is not straightforward, however, given the shell model we have used to derived (H þ 3 )values. Figure 8 therefore compares our measured profiles with those derived from Grodent et al. (2001) which have been convolved with the effective seeing. We have done this both for the original Grodent profile and for the profile scaled for non-lte effects.

7 796 LYSTRUP ET AL. Vol. 677 Fig. 9. Temperature profile. The predicted vertical temperature profile from the Grodent et al. (2001) model is the dashed curve. The solid curve is our temperature fit through the line of sight for the off-limb pixels. Fig. 8. Comparison of measured and convolved models. Similar to Fig. 7, but the model profiles have been convolved with a Gaussian kernel with FWHM equal to the effective seeing. As in Fig. 7, the measured profile has been shifted such that the altitude of the peak density matches that of the Grodent model peak density. Note that the sharp peak in the measured profile is an artifact of the calculation. This convolution of the model profiles brings them into much better agreement with the measured values of (H þ 3 ), particularly for the model non-lte profile. The agreement between this profile and the measured profile is excellent up to 1800 km. At higher altitudes, however, the measured profile indicates that there is considerably more H þ 3 than the model predicts. This may be due to a flux of lower energy electrons precipitating at altitudes higher than either of the two Maxwellian distributions used by Grodent et al. (2001). Figure 9 compares temperature as a function of altitude for our profile and that of Grodent et al. (2001). The measured temperature profile is a density-weighted average along the line of sight. The measured temperatures are consistently higher for the same altitude. In particular, the (average) exospheric temperature that we measure is 1450 K, approximately 150 K higher than the Grodent et al. (2001) temperature. One consequence of this is that the model heights above 1500 km might have to be increased by about 12% for a better comparison with the measured density profile in order to allow for the increase in scale height. This would not, however, be enough to account for the discrepancy between the modeled non-lte convoluted (H þ 3 ) and that derived from our data above 1800 km (Fig. 8). 4. CONCLUSIONS We present the first measured H þ 3 ion density profile ever produced for the Jovian auroral regions. Comparison with work by Grodent et al. (2001) shows that their model reproduces the actual ion densities very well up to 1800 km. After that, our data indicate that further fluxes of lower energy electrons may be required to account for the measured densities. In the future we plan to use this observational method to probe ion densities at a full range of planetary latitudes. One of the key issues we will attempt to address is the distribution of H þ 3 at lower latitudes. Previous studies (Lam et al. 1997; Miller et al. 1997; Rego et al. 2000) found that the measured ion column densities were greater than could be explained by solar EUV ionization alone and that there seemed to be some correlation with the magnetic field (Connerney et al. 1998). These workers concluded that additional low latitude particle precipitation was needed to account for the measured values of N(H þ 3 ). Altitude profiles of (Hþ 3 ) will place important constraints on the nature and intensity of the particle fluxes required. The UK authors have previously used infrared studies of H þ 3 emission to investigate the other giant planets of our solar system, in particular long-term monitoring of Saturn. It may be possible to investigate the Saturnian ionosphere with the method presented herein. The H þ 3 emission originating in Saturn s ionosphere is far weaker than in the case at Jupiter, and thus long exposures would be necessary. We thank Grant Hill and Gary Puniwai for their expert assistance during these observations. Many thanks to Roger Yelle and Darrell Strobel for very helpful discussions. The data presented herein were obtained at the W. M. Keck Observatory, which is operated as a scientific partnership among the California Institute of Technology, the University of California, and the National Aeronautics and Space Administration. The Observatory was made possible by the generous financial support of the W. M. Keck Foundation. The authors wish to recognize and acknowledge the very significant cultural role and reverence that the

8 No. 1, 2008 JOVIAN ION DENSITY PROFILE 797 summit of Mauna Kea has always had within the indigenous Hawaiian community. We are most fortunate to have the opportunity to conduct observations from this mountain. The UK authors are part of the Europlanet European planetary science network, which is supported by the European Union s Framework 6 program. This material is based on work supported by AURA through the National Science Foundation under AURA Cooperative Agreement AST as amended. The authors wish to thank the referee, Emmanuel Lellouch, whose timely and thoughtful comments improved this paper. Facilities: Keck: II. Achilleos, N., Miller, S., Tennyson, J., Aylward, A. D., Mueller-Wodarg, I., & Rees, D. 1998, J. Geophys. Res., 103, Atreya, S. K., Donahue, T. M., & Waite, J. H., Jr. 1979, Nature, 280, 795 Connerney, J. E. P., Acuña, M. H., Ness, N. F., & Satoh, T. 1998, J. Geophys. Res., 103, Drossart, P., et al. 1989, Nature, 340, 539 Eshleman, V. R., Tyler, G. L., Wood, G. E., Lindal, G. F., Anderson, J. D., Levy, G. S., & Croft, T. A. 1979, Science, 204, 976 Fjeldbo, G., Kliore, A., Seidel, B., Sweetnam, D., & Woiceshyn, P. 1976, in IAU Colloq. 30, Jupiter: Studies of the Interior, Atmosphere, Magnetosphere and Satellites, ed. T. Gherels ( Tucson: Univ. Arizona Press), 238 Grodent, D., Hunter Waite, J., & Gérard, J.-C. 2001, J. Geophys. Res., 106, Hamilton, D. C., et al. 1980, Geophys. Res. Lett., 7, 813 Kao, L., Oka, T., Miller, S., & Tennyson, J. 1991, ApJS, 77, 317 Kim, Y. H., Fox, J. L., & Porter, H. S. 1992, J. Geophys. Res., 97, 6093 Kliore, A., Cain, D. L., Fjeldbo, G., Seidel, B. L., & Rasool, S. I. 1974, Science, 183, 323 Lam, H. A., Achilleos, N., Miller, S., Tennyson, J., Trafton, L. M., Geballe, T. R., & Ballester, G. E. 1997, Icarus, 127, 379 Lystrup, M. B., Miller, S., Stallard, T., Smith, C. G. A., & Aylward, A. 2007, Ann. Geophys., 25, 847 McComas, D. J., Allegrini, F., Bagenal, F., Crary, F., Ebert, R. W., Elliott, H., Stern, A., & Valek, P. 2007, Science, 318, 217 McConnell, J. C., & Majeed, T. 1987, J. Geophys. Res., 92, 8570 McLean, I. S., et al. 1998, in Proc. SPIE, 3354, 566 Melin, H. 2005, Ph.D. thesis, Univ. College London REFERENCES Melin, H., Miller, S., Stallard, T., & Grodent, D. 2005, Icarus, 178, 97 Miller, S., Achilleos, N., Ballester, G. E., Lam, H. A., Tennyson, J., Geballe, T. R., & Trafton, L. M. 1997, Icarus, 130, 57 Miller, S., Joseph, R. D., & Tennyson, J. 1990, ApJ, 360, L55 Miller, S., et al. 2006, Philos. Trans. R. Soc. London A, 364, 3121 Neale, L., Miller, S., & Tennyson, J. 1996, ApJ, 464, 516 Neale, L., & Tennyson, J. 1995, ApJ, 454, L169 Oka, T., & Epp, E. 2004, ApJ, 613, 349 Raynaud, E., Lellouch, E., Maillard, J.-P., Gladstone, G. R., Waite, J. H., Bézard, B., Drossart, P., & Fouchet, T. 2004, Icarus, 171, 133 Rego, D., Achilleos, N., Stallard, T., Miller, S., Prangé, R., Dougherty, M., & Joseph, R. D. 1999, Nature, 399, 121 Rego, D., Miller, S., Achilleos, N., Prangé, R., & Joseph, R. D. 2000, Icarus, 147, 366 Seiff, A., et al. 1998, J. Geophys. Res., 103, Sofieva, V., Tamminen, J., Haario, H., Kyrölä, E., & Lehtinen, M. 2004, Ann Geophys., 22, 3411 Stallard, T., Miller, S., Millward, G., & Joseph, R. D. 2001, Icarus, 154, , Icarus, 156, 498 Stallard, T. S., Miller, S., Trafton, L. M., Geballe, T. R., & Joseph, R. D. 2004, Icarus, 167, 204 Strobel, D. F. 2008a, Icarus 193, b, Icarus, 193, 612 Waite, J. H., Cravens, T. E., Kozyra, J., Nagy, A. F., Atreya, S. K., & Chen, R. H. 1983, J. Geophys. Res., 88, 6143