The three-dimensional structure of Saturn s equatorial jet at cloud level

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1 Icarus 187 (2007) The three-dimensional structure of Saturn s equatorial jet at cloud level A. Sánchez-Lavega, R. Hueso, S. Pérez-Hoyos Departamento de Física Aplicada I, Escuela Técnica Superior de Ingeniería, Universidad del País Vasco, Alda. Urquijo s/n, Bilbao, Spain Received 1 August 2006; revised 28 September 2006 Available online 20 December 2006 Abstract The three-dimensional structure of Saturn s intense equatorial jet from latitudes 8 Nto20 S is revealed from detailed measurements of the motions and spectral reflectivity of clouds at visible wavelengths on high resolution images obtained by the Cassini Imaging Science Subsystem (ISS) in 2004 and early Cloud speeds at two altitude levels are measured in the near infrared filters CB2 and CB3 matching the continuum (effective wavelengths 750 and 939 nm) and in the MT2 and MT3 filters matching two methane absorption bands (effective wavelengths 727 and 889 nm). Radiative transfer models in selective filters covering an ample spectral range ( nm) require the existence of two detached aerosol layers in the equator: an uppermost thin stratospheric haze extending between the pressure levels 20 and 40 mbar (tropopause level) and below it, a dense tropospheric haze-cloud layer extending between 50 mbar and the base of the ammonia cloud (between 1 and 1.4 bar). Individual cloud elements are detected and tracked in the tropospheric dense haze at 50 and 700 mbar (altitude levels separated by 142 km). Between latitudes 5 N and 12 S the winds increase their velocity with depth from 265 m s 1 at the 50 mbar pressure level to 365 m s 1 at 700 mbar. These values are below the high wind speeds of 475 m s 1 measured at these latitudes during the Voyager era in , indicating that the equatorial jet has suffered a significant intensity change between that period and or that the tracers of the flow used in the Voyager images were rooted at deeper levels than those in Cassini images Elsevier Inc. All rights reserved. Keywords: Atmospheres, dynamics; Atmospheres, structure; Saturn, atmosphere; Radiative transfer 1. Introduction The strong equatorial jets of Jupiter and Saturn represent a major challenge to the understanding of the dynamics and of the general circulation of their planetary atmospheres (Ingersoll et al., 2004; Vasavada and Showman, 2005). Saturn has a broad (extending between latitudes ±40 ) and intense zonal jet blowing eastward with maximum speed at cloud level of 475 m s 1 as measured during the Voyagers encounters in (Ingersoll et al., 1984; Sánchez-Lavega et al., 2000). The jet showed much lower speeds during (Sánchez- Lavega et al., 2003, 2004) and vertical wind shears were proposed to explain the anomalous velocities (Porco et al., 2005). However, a reanalysis of Saturn s cloud vertical structure during the Voyager encounters suggested that real changes could have occurred in the jet at cloud level (Pérez-Hoyos and Sánchez- * Corresponding author. Fax: address: agustin.sanchez@ehu.es (A. Sánchez-Lavega). Lavega, 2006). Distinguishing between dynamical variability (temporal changes) and permanent vertical wind shears or the existence of both, is a fundamental step to understand the mechanisms and energy sources, external (solar) or internal (deep convection), that intervene in creating this strong jet. The Cassini-ISS instrument is taking images of the planet regularly since 2004, employing a variety of filters that cover the whole spectral visual range (Porco et al., 2004). These images can be used to track the atmospheric features (Porco et al., 2005; Vasavada et al., 2006; Sánchez-Lavega et al., 2006) and, properly calibrated, to infer the upper cloud and haze vertical structure using radiative transfer models (Sánchez-Lavega et al., 2006). The first published study of motions in Saturn s atmosphere using ISS images (Porco et al., 2005) showedtwo distinct patterns of zonal wind velocities at southern equatorial latitudes based on a scarce number of tracked features. One family of 14 low speed points was obtained using the MT2 methane band filter (727 nm) and another family of 25 higher speed points was obtained using the CB2 continuum filter (750 nm). Qualitatively, the different vertical penetration at /$ see front matter 2006 Elsevier Inc. All rights reserved. doi: /j.icarus

2 Saturn s equatorial jet 511 these two wavelengths in Saturn s atmosphere suggested that motions were sensed at two altitude levels, with winds decreasing with altitude, in agreement with the vertical wind shears derived above the clouds from temperature measurements (Flasar et al., 2005). No other study on the motions of vertical cloud structure at the equator using Cassini-ISS images has been published since then. Vasavada et al. (2006) studied the motions at non-equatorial latitudes, mainly the temperate latitudes, and Sánchez-Lavega et al. (2006) concentrated their study on the South Polar Region. The aim of this paper is to present a first consistent detailed study of the motions in the equatorial area during the first year of the Cassini mission ( ), based on cloud tracking in images obtained with filters that are sensitive to different altitude levels, and to quantify by means of radiative transfer models, the altitude location of the features used as tracers. We therefore present a study of the three-dimensional structure of the jet, in other words, the zonal wind meridional profile as a function of altitude and compare them with previous profiles retrieved from the Voyager mission, Hubble Space Telescope observations and ground-based data obtained since We discuss the implications of the observed temporal behaviour. 2. Observations and analysis Saturn s equatorial atmosphere was observed at high resolution by the Cassini Imaging Science Subsystem (ISS) (Porco et al., 2004) following the spacecraft orbital insertion on 1 July We have used in this study 76 Cassini-ISS images obtained between 8 10 May 2004, 6 10 September 2004 and February 2005 with filters that cover the wavelength range from 258 to 939 nm (see Fig. 1 in Sánchez-Lavega et al., 2006). These images allowed us to obtain precise tracking of cloud motions and to perform a spectral reflectivity analysis to determine the vertical cloud structure. The highresolution images were taken with the Narrow Angle Camera at a resolution between 15 and 50 km pixel 1. The images were first selected in pairs separated by about 10 h and then processed and navigated for latitude and longitude measurements of each individual pixel element in the disk, following procedures identical to those described in our previous paper (Sánchez-Lavega et al., 2006). Accurate navigation of the images requires fitting an ellipse to the planet s limb. Navigation of the May 2004 images was not difficult since the full disk was visible (May 2004). In the September 2004 series, Saturn disk was fragmented in four frames, with only one frame useful for cloud tracking at equatorial latitudes. Unfortunately the limb is barely visible in the frame so we used the adjacent frame, that includes mid- and polar latitudes (and the limb), as a reference to control the navigation and position of those features that are visible simultaneously in both frames. In February 2005 the images were taken with a different acquisition strategy scanning the equatorial region from east to west using twelve overlapping frames. The eastward frame shows a large portion of the limb and was used for an accurate navigation. The subsequent overlapping frames were navigated using as references those prominent features which are common to the adjacent frames. As a control, nearly simultaneous wide angle images of the whole planet were also navigated, checking the position of the largest target features. For the photometric study, we calibrated the May 2004 ISS images using the CISSCAL software (Porco et al., 2004) to get the intensity of the reflected solar radiation at each pixel on the planet disk, as a function of the scattering angles. We used the following Cassini filters for this part of the study: UV1 (264 nm), UV2 (306 nm), UV3 (343 nm), BL2 (440 nm), CB1 (619 nm), MT1 (619 nm), MT2 (727 nm), CB2 (752 nm), MT3 (890 nm) and CB3 (939 nm). Values of the methane absorption coefficients are showed in Table 2. In total we tracked 905 individual and unique cloud features, 614 in the CB2 and CB3 filters and 291 features in the MT2 and MT3 filters, spanning along the latitude band from 8 N to 40 S (all latitudes in this paper are planetographic). The higher latitude cut at 8 N is imposed by the ring obstruction when projected on the planetary disk. The lower limit 40 S is where the wind reaches velocity zero. These two filter pairs combination is well suited to detect the clouds with good contrast at different altitude levels in the equator (Fig. 1). In the latitude band from 8 N to about 13 S, features present in the CB2 CB3 wavelengths are absent in the MT2 MT3 ones and show a different morphology, suggesting they are separated in altitude since both filter couples sense different altitude levels. Features in the CB2 CB3 filters tend to be isolated spots, with a bright contrast relative to background clouds and hazes. Sometimes, complexes clusters of brighter ( white ) spots (feature C in Fig. 1) with typical sizes km are observed in rapid evolution at latitudes 0 5 S. These are the most conspicuous cloud structures seen at these wavelengths. No trace of them is apparently detected in the MT2 MT3 filter couple. These two filters show, on the contrary, different cloud textures. In the northern equatorial part at 5 N, large elongated and narrow filaments spanning km beginning from a round spot run nearly parallel along the equator (features A at 5 N in Fig. 1). The band between latitudes 3 8 S, shows in MT2 MT3 wavelengths a cluster of short filaments ( 2000 km in length) tilted in the northwest to southeast direction by relative to the equator (features A at 5 S). The apparently decoupled morphology of cloud features at two levels strongly suggests a different dynamical mechanism in the origin of features in this latitude range. At the other wavelengths we see the same structures as at MT2 MT3 or CB2 CB3 but with smaller contrast, so they add nothing to the cloud track job. For example, the UV filters (mainly the UV3) show at most the features seen in MT2 MT3 but under a contrast reversal ( white structures are seen black ). In MT1 the contrast strongly decreases so we barely see some of the MT2 MT3 features. The same applies to the CB1 filter when compared to CB2 CB3. On the other hand, the broadband continuum filters (BL1 and GRN, blue and green, respectively) show the same features as in MT2 MT3 but at a lower contrast. Finally, the IR3 filter shows a mixture of features, some as in MT2 MT3 and others as in CB2 CB3 since

3 512 A. Sánchez-Lavega et al. / Icarus 187 (2007) Fig. 1. Map projection of the equatorial region of Saturn showing the cloud morphology as seen by Cassini-ISS on September 9, 2004: (a) image N obtained with the MT3 filter (890 nm) sensitive to high altitudes (individual features denoted as A); (b) image N obtained with the CB3 filter (938 nm) sensitive to deeper clouds (individual and distinct features labelled C). Latitudes are planetocentric. this broadband filter covers a large wavelength range encompassing their wavelengths. All this indicates qualitatively that the variability of the features contrast with wavelength is a consequence of their altitude location in the atmosphere within the hazes in the equatorial area and that only two families of features are available for cloud tracking. 3. Results: The equatorial jet velocities In Fig. 2 we show the individual cloud speeds measured at the two filter combinations as a function of latitude, tracing the meridional profile of zonal winds at two altitude levels. The zonal velocities are calculated relative to the Voyager radio-rotation period System III reference frame with a period τ III = 10 h 39 min 22.4 s (Desch and Kaiser, 1981). However, it should be noted that Ulysses (Galopeau and Lecacheux, 2000) and Cassini radio data (Gurnett et al., 2005) indicate lower radio-rotation periods (in the range 10 h 46 min, Sánchez- Lavega, 2005), in agreement with direct measurements of the magnetic field rotation (10 h 47 min, Giampieri et al., 2006). Until a definitive rotation period is proposed by the International Astronomical Union, we keep System III as a reference since it allows us making a comparison with previous studies on Saturn s atmospheric motions. Larger rotation periods would

4 Saturn s equatorial jet 513 Fig. 2. The zonal wind velocity profile as measured from cloud tracking on images obtained from two different missions: (a) continuous dark line, Voyager profile for (Sánchez-Lavega et al., 2000) corresponding to the altitude range 360 ± 140 mbar (Pérez-Hoyos and Sánchez-Lavega, 2006); (b) red dots and its averaged curve correspond to Cassini-ISS CB2 CB3 filter (939 nm) and, blue dots and its averaged curve correspond to the MT2 MT3 filter (889 nm), both from May 2004 to February (For interpretation of the references to colour in this figure legend, the reader is referred to the web version of this article.) give lower zonal wind speeds. The zonal wind speed (u N )ina reference frame with rotation period τ N is given by u N = u III 2πR(ϕ)cos ϕ [ (1/τ III ) (1/τ N ) ] u III being the zonal wind speed in System III and R(ϕ) the planetary radius at latitude ϕ. A fit to the binned profiles using boxes with 0.5 width in latitude (except in few latitudes where 1 boxes were used due to the lack of data), gives an average velocity of 263 m s 1 in MT2 MT3 wavelengths for latitudes 8 Nto15 S and 364 m s 1 in CB2 CB3 wavelengths for latitudes 2 Nto 12 S. The dispersion in wind speed ranges from 1 to 25 m s 1 and is due to the contribution of the measurement errors (typically 1 5 m s 1 ) with the same sources as explained in Sánchez-Lavega et al. (2006), to local motions and probably to altitude differences in the tracers coupled to the vertical wind shears (see below). In Table 1 we give the binned-averaged values of the zonal wind speeds in these two filters which correspond to the fits in Fig. 2. There are several characteristics of the two profiles that merit a comment. The averaged MT2 MT3 profile (high altitude) shows a small latitudinal shear in the wind velocity from latitudes 8 Nto15 S. The lack of cloud features with enough contrast to be tracked at the equator (from 3 Sto3 N) prevents us from measuring the full profile but the fit to the points close to this latitude suggests that at the equator the winds converges toward a minimum speed of 225 m s 1. The CB2 CB3 profile shows between latitudes 5 S and 10 N (jet peak) a more undulating pattern with some individual points reaching maximum speeds of 400 m s 1. The averaged profile shows a strong minimum between 5 N and 8 N where the speeds are similar to those in the MT2 MT3 filters at these latitudes ( 275 m s 1 ). It must be noted that these measurements are difficult due to the blockage of the atmosphere by the rings, and future measurements under a better viewing angle should confirm this point. The two profiles merge at latitude 16 N and seem to follow a similar trend at higher latitudes. Our results agree with the 39 points presented by Porco et al. (2005) in the equatorial latitudes in CB2 and MT2. In addition, our measurements obtained with the Hubble Space Telescope (Sánchez-Lavega et al., 2003, 2004) fit well the Cassini MT2 MT3 profile. 4. Vertical cloud structure in the equatorial region In order to retrieve the altitude location of the tracers used to determine the wind profile, we have used intensity calibrated ISS images from 10 May 2004 that give the reflected solar ra-

5 514 A. Sánchez-Lavega et al. / Icarus 187 (2007) Table 1 Averaged mean zonal flow on Saturn s equatorial region as a function of latitude as determined from Cassini-ISS images at two altitude levels in Zonal winds in CB2 CB3 filters Latitude graphic (deg) ±Latitude graphic (deg) U (m s 1 ) ± U (m s 1 ) Zonal winds in MT2 MT3 filters Latitude graphic (deg) ±Latitude graphic (deg) U (m s 1 ) ± U (m s 1 )

6 Saturn s equatorial jet 515 diation as a function of wavelength and scattering angles. We interpret the reflectivity measurements using a radiative transfer code that reproduces the sunlight absorption and scattering by both gas and haze-cloud particles in Saturn s atmosphere (Acarreta and Sánchez-Lavega, 1999; Pérez-Hoyos et al., 2005; Sánchez-Lavega et al., 2006). We have analysed the spectral reflectivity at four latitudes (3.5 N, 0,10 and 20 S), selected according to the cloud morphology and wavelength dependent wind profiles. Fitting procedures were similar to that presented in Pérez-Hoyos et al. (2005) and Sánchez-Lavega et al. (2006), with the following few modifications: (1) The tropospheric haze values tested were adapted to those expected for the equatorial region (top pressure values ranging from 30 to 100 mbar, bottom values from 100 mbar to the ammonia condensation level and optical thickness from 0 to 100). (2) We also checked three phase functions for the tropospheric haze (Mie, double Henyey Greenstein and isotropic), and (3) we varied the number of layers (from 2 to 4) in order to find a better fit. The models that best fit the observations require two detached aerosol layers at all the latitudes explored as typically found in other recent models of the atmosphere (Pérez- Hoyos et al., 2005; Karkochska and Tomasko, 2005). In Fig. 3 we show an example of our model observations fit for the latitude 3.5 N and in Table 3 we give the details of the Fig. 3. Comparison between the observations of the spectral reflectivity (I/F) (dots) and our vertical structure atmospheric model (continuous line), as a function of the longitude in the planet. The data and model are for latitude 3.5 N and the wavelength can be identified by the filter name (Porco et al., 2004; Sánchez-Lavega et al., 2006).

7 516 A. Sánchez-Lavega et al. / Icarus 187 (2007) Table 2 Methane absorption coefficients Filter K (1/km-am) ϖ 0 UV ± 0.1 UV ± 0.1 UV ± 0.1 BL ± 0.05 CB ± MT ± MT ± 0.01 CB ± MT ± CB ± aerosol structure of these two hazes. There is an uppermost stratospheric haze extending between the pressure levels 20 and 50 mbar formed by small particles with a mean radius of 0.15 µm with refractive indexes m r = 1.43 and m i ranging from 0.01 (UV) to (IR) (Pérez-Hoyos et al., 2005; Pérez-Hoyos and Sánchez-Lavega, 2006). This layer has an optical depth of less than 0.2 at wavelengths above 300 nm and no detached features are detected in it. The model gives the following optical depth dependence with wavelength for the stratospheric haze: τ str = exp [ λ(nm) ]. Below, a tropospheric haze-cloud layer extends between 60 mbar (the tropopause level) and the base of the ammonia cloud ( 1.5 bar). It is in this vertically extended layer where detached cloud features form and where winds are measured at different altitudes. Since Cassini observations were performed at a high phase angle (α 90 ), greater than those restricted to α<6 for ground-based and Hubble Telescope images, it is possible to get complementary information on the phase function. We found that a two-term Henyey Greenstein phase function with enough backscattering (function parameters f = , g 1 = 0.1, g 2 = 0.3) or an isotropic phase function, both with the same single scattering albedo ϖ 0 decreasing in the ultraviolet-blue spectral range (see Table 2), can reproduce the observations. Our modelling rules out Mie scatterers for the tropospheric haze (radius µm). The haze is very thick in the equator (3.5 N and 0 ) with optical thickness of 30 decreasing to 7 at latitude 20 S where the two wind profiles converge. We use this model to locate our tracers in altitude. The uncertainty in the location of this layer s bottom level is much higher than the one of the top pressure level. Increasing the optical depth of the tropospheric haze allows putting lower its bottom part, still keeping within the uncertainties of the reflectivity measurements. For the optical depth uncertainty given in Table 3 at each latitude band, we get the bottom of this layer to be between 1.0 and 1.4 bar. The MT2 and MT3 filters are centred in methane absorption bands and are sensitive to the higher atmospheric levels (optical depth 1 2 reached at levels 60 and 100 mbar). According to our model interpretation, they show features by the reflectivity contrasts in the top level of the tropospheric haze produced by small altitude differences in the top level of the haze (Sánchez- Lavega et al., 2004). We quantify the altitude differences to be 10 km to reproduce the observed reflectivity. This is the case of the features type A shown in Fig. 1 and placed at latitudes 5 N 0 and 3 8 S. These features are also seen in the ultraviolet filters ( nm) by a contrast reversal relative to methane filters, due to a combination of Rayleigh scattering by the gas and absorption by the haze particles. On the other hand, to precisely retrieve the altitude location of the bright cloud structures seen in CB2 CB3 (feature C in Fig. 1), we have calculated the change of their contrast relative to surroundings as a function of altitude when placing them within the tropospheric haze as in Pérez-Hoyos and Sánchez- Lavega (2006). This is shown in Fig. 4 for the different wavelengths we used, comparing them with the observed values. The CB2 and CB3 filters have the deepest penetration since no significant absorption and scattering by the gases occur at their effective wavelengths. According to Fig. 4, the bright clouds type C are placed deep at 700 mbar, in the region of the ammonia cloud whose base is predicted to be located at 1.5 bar. The filters MT1 (619 nm, methane) and CB1 (619 nm) are sensitive to intermediate atmospheric levels ( mbar). Features seen in the MT2 MT3 filters are barely seen in MT1. The same occurs with CB2 CB3 and the CB1 filter that shows the deep features with a very low contrast. No features up to a 1% contrast are detected between 100 mbar and 600 mbar levels with MT1 and CB1. We conclude that the haze is homogeneous in particle distribution in this altitude range and no detached features are seen within it. 5. Discussion The picture that emerges from our wind and reflectivity measurements in the equatorial area ( 5 Nto12 S) is that we are sensing distinct cloud structures at two different levels 60 mbar and mbar where winds blow on the average at 263 and 364 m s 1, respectively. Fig. 5 summarises schematically the situation. Features type A highly placed are drawn as a wavy pattern, representing the altitude differences Table 3 Vertical cloud structure in Saturn s equatorial region Latitude P top stratosphere (mbar) P bottom stratosphere (mbar) P top troposphere (mbar) P bottom troposphere (bar) Troposphere optical depth τ trop Stratosphere particle radius a (µm) 3.5 N 20± ± ± ± ± ± ± ± ± ± ± ± S 20± ± 5 50± ± ± ± S 20± ± ± ± 0.2 7± ± 0.05

8 Saturn s equatorial jet 517 Fig. 4. Altitude levels of the individual cloud tracers placed within the tropospheric haze. The lines show the calculated dependence of the reflectivity contrast with the altitude for a bright cloud feature when embedded within the tropospheric haze at various wavelengths: red (filter CB3), black (filter CB2), green (filter GRN), blue (filter CB1) and blue dots (filter MT1). Circles mark the observed contrast for features type C at each wavelength. The vertical line at contrast 1% marks the detection limit. No features are identified above this limit in the MT1 images. Our models put these deep features at 700±100 mbar. Horizontal error bars are derived from the dynamic range of the images, vertical error bars are pressure error corresponding to the contrast error. (For interpretation of the references to colour in this figure legend, the reader is referred to the web version of this article.) in the top of the tropospheric haze layer. The region named B corresponds to intermediate altitudes where no individual clouds elements are detected. Finally, features C are detached bright clouds located deeper in the ammonia main cloud deck. The dual aspect of the meridional wind profile (the two separated branches with no intermediate wind speeds between them as shown in Fig. 2), is well explained by the concentration of cloud tracers at levels mbar (detected in UV and MT filters) and at the mbar (detected in the CB filters) with a lack of tracers in between ( mbar altitude range). Obviously, small differences in altitude of the tracers and the existence of a vertical wind shear can explain the velocity dispersion seen when averaging both profiles (Table 1 and Fig. 2). Convergence of the two branches of the wind profile occurs at latitudes 16 S where our photometric models indicate that clouds form at a similar level within the tropospheric haze (pressures 200 mbar, Sánchez-Lavega et al., 2004). The same occurs at about 5 N where accurate photometry cannot be performed due to ring projection. The vertical wind shear from the averaged wind speeds is u/ z 40 m s 1 per atmospheric scale height (H = 54 km in Saturn s equator), i.e., winds increase with depth by 100 m s 1 in 142 km altitude difference. This vertical wind shear is similar to that measured in the equatorial region of Jupiter by the Galileo probe (Atkinson et al., 1998) Fig. 5. Scheme of the vertical cloud structure in Saturn s equatorial region and retrieved zonal wind speeds. The upper stratospheric haze is located above the tropopause (level P 0 ). Altitude differences at the top of this haze layer (drawn as a wavy pattern) are seen as contrasted features (type A) in MT2 MT3 filters (see Fig. 1). The level P l = 600 ± 100 mbar marks the visibility limit with the MT1 filter, and the level P c = 700 ± 100 mbar, marks the location of the deep clouds in CB1 CB2 CB3 filters (features named C, see Fig. 1). No distinct clouds are detected in between levels P 0 and P l (region named B). The crosses mark the eastward zonal jet velocity at the two altitude levels where clouds are detected and tracked. The total optical depth (τ ) of the tropospheric haze is indicated, together with the location of the base of the ammonia cloud layer. and agrees with the range of Cassini-CIRS derived vertical wind shears of 52 m s 1 /H at 156 mbar, although it is greater than the 26 m s 1 /H derived at 500 mbar (Flasar et al., 2005; Fletcher et al., 2005). Regarding our previous Hubble Space Telescope (HST) wind measurements for (Sánchez- Lavega et al., 2003, 2004), we see that they are in excellent agreement with the lower speed branch profile presented here, indicating that the HST tracers were located in the upper troposphere (Pérez-Hoyos and Sánchez-Lavega, 2006). With respect to temporal changes in the equatorial winds, we recently showed from an altitude study of the prominent plumes used as tracers in Voyager images in that the high speed winds were not placed deeper than the Cassini features tracked in the CB filters (Pérez-Hoyos and Sánchez- Lavega, 2006). A comparison of Cassini CBs and Voyager wind profiles (Fig. 2) suggests that a real change has occurred in the equatorial winds between latitudes 3 S and 8 N, the only latitude band with common data from both missions and where comparison is possible. The wind speed change could amount to 50 m s 1 at latitude 0 but could be as large as 180 m s 1 at 6 N(Fig. 2). Notoriously this was the latitude where the Great White Spot, a large storm system, erupted in September 1990 (Sánchez-Lavega et al., 1991; Barnet et al., 1992a; Beebe et al., 1992), deeply modifying the equatorial cloud structure and probably inducing subsequent large scale disturbances as those observed in (Sánchez-Lavega et al., 1996, 1999). In addition, this region is subject to periodic ring

9 518 A. Sánchez-Lavega et al. / Icarus 187 (2007) shadowing and radiative heating variations that could also intervene in circulation changes (Barnet et al., 1992b). As an alternative explanation to the temporal variability in wind speed, but in disagreement with our photometric analysis (Pérez-Hoyos and Sánchez-Lavega, 2006), the Voyager tracers could have been placed deeper in the atmosphere, at an altitude where the wind speeds are high. Using the vertical wind shear of 40 m s 1 /H, a linear extrapolation indicates that the Voyager winds at latitude 0 should be placed at 2 bar (at the bottom of the ammonia cloud base formation). However at latitude 6 N, the large wind difference of 180 m s 1 would need a huge vertical wind shear of 180 m s 1 /H to keep the tracers and winds no deeper than the 2 bar. If we extrapolate the measured wind shear of 40 m s 1 /H downwards, clouds would be located at latitude 0 at a non-realistic level of 10 bar. Another possibility to explain the wind differences is that the Voyager plumes used as altitude tracers of the flow had their cloud tops highly placed, at 360 mbar according to our photometric model (Pérez-Hoyos and Sánchez-Lavega, 2006), but that they were in fact deeply rooted, moving themselves with a higher flow velocity at their feet, located perhaps at the level of the ammonia cloud base ( 1.5 bar). This deep flow dragging the plumes should be the one measured using the smaller tracers present in Voyager s green and orange filters but not in the violet filter (Sánchez-Lavega et al., 2000). The plumes should extend vertically more than ln(0.36/1.4) 1.4 scale heights with their lifetime (observed to be several days) being limited by the flow vertical and meridional wind shears. Future additional measurements of Cassini-ISS images and studies with the other instruments should give precise information on the 3D temperature structure, aerosol properties and their vertical distribution, and chemical abundances and distributions of the different species (including aerosols) in the equatorial region, as well as their temporal variability. Cloud motions from the near infrared imaging with VIMS in the 4 5 micron window, with deeper penetration than in the visible, and the determination of the precise altitude location of the tracers at these wavelengths (Baines et al., 2005), will complete the full 3D structure of the winds Saturn s upper troposphere. Acknowledgments This work has been supported by the Plan Nacional de Astronomía y Astrofísica (MEC) AYA and Grupos UPV 15946/2004. R. Hueso was supported by the Ramón y Cajal Program (MEC). This research made use of the public Cassini images available at NASA Planetary Data System. 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