Figure 1 shows a subset of ultra-narrowband images reconstructed from the spectra that are characteristic of the dataset. The first row of images are from 1.55–1.57 μm, where CH4 absorption is weak, and primarily intensity from the surface is measured. The dominant surface feature is the bright feature located at ∼100°W.
Figure 1. Spectral image maps of Titan. Each image is integrated over 0.0009 μm around the central wavelength indicated. The sub-observer point is 111°W longitude and −23°N latitude, the phase angle is ∼6°, and the sub-solar point is East (to the right) of the sub-observer point.
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 The middle row of images in Figure 1 show that as CH4 opacity increases with wavelength, flux from higher altitude in the atmosphere is observed. At 1.6028 μm the surface dominates the signal, whereas at 1.6158 μm flux from a tropopause haze is observed (see below). At wavelengths longer than 1.6158 μm the bright surface feature is no longer visible. CH4 absorption at this wavelength is strong enough to completely attenuate the light from the sun before it reaches the surface. The observed signal is therefore from photons that are scattered from particulates in the atmosphere.
 The 1.6158 μm image shows that there is a localized enhancement of lower-atmospheric particulates near the southern (summer) pole. This enhancement occurs at the latitude where tropopause cirrus has been detected spectroscopically at 2.1 μm [Griffith et al., 1998; Brown et al., 2002] and has been directly imaged with a narrowband filter at 2.108 μm [Roe et al., 2002b] We show by modelling the observed spectra from this region that this haze is found at 30–50 km, near the tropopause at 42 km.
 At wavelengths longer than 1.6218 μm (bottom row of Figure 1), stratospheric particulates are measured. At these wavelengths the observed flux is diminished, signal-to-noise decreases, and limb brightening becomes more pronounced. The N-S stratospheric haze asymmetry imaged on the same night with an H1702 narrowband filter [Roe et al., 2002a] is seen in the 1.702 μm image and quantitatively modeled below. However, with the narrow bandpass and incomplete spatial coverage of these images we do not clearly observe the E-W asymmetry or polar collar seen by Roe et al. [2002a].
 The pathlength through Titan's atmosphere plays a critical role in determining the observed spectra. Therefore, we first compare spectra measured at the same Titan airmass – that is, points that are equidistant from the center of the disk. Figure 2a shows 8 spectra taken at a Titan airmass of A = 1.12 (0.5 Titan radii, RT, from disk center). The pixel-to-pixel noise of each spectrum is read-noise limited (∼30 DN, or in this case 0.003 I/F) and the similarity of the spectra longward of 1.63 μm is a reassuring consistency check of the data reduction, normalization, image reconstruction and centering.
Figure 2. Observed and modeled spectra from multiple locations at the same airmass on Titan's disk – the dashed vertical lines are the wavelength of the inset images and colored circles indicate the locations of the corresponding spectra. Surface albedo contrasts that are observed at 1.592 μm (A) do not significantly affect spectra from 1.61–1.66 μm. Enhanced tropopause haze at the southern limb (B) is probed from 1.61–1.62 μm. Panel C compares the observed spectra (points) with radiative transfer model (lines) that fit the tropopause haze enhancement. Altitude dependence of tropopause haze (D) is more sensitive to decrease (red dashed line) to 20 km than to increase to 70 km (red solid line).
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 Shortward of 1.60 μm the spectra diverge as they probe deeper into the troposphere. The greatest contrast in the spectra occurs at 1.592 μm, where the measurement is sensitive to surface albedo. Figure 2a shows that bright surface features show up in the spectra at 1.592 μm, and that they do not affect the spectra significantly at wavelengths longer than 1.610 μm. This qualitatively indicates that the spectra from 1.61 to 1.64 μm probe an altitude region from the troposphere to the lower stratosphere, and exclude the surface.
 The bright feature seen around the southern pole at 1.62 μm can be investigated by comparing spectra taken from around the limb (A = 2.35, RT = 0.9) of Titan (Figure 2b). A striking feature in these spectra is the increasing divergence at wavelengths shorter than 1.63 μm. The five spectra from the northern hemisphere overlap at 1.61 μm, and contrast with the three spectra from the southern hemisphere at this wavelength. This indicates a robust observational difference between the two sets of spectra.
 A 2-stream approximation solves the radiative transfer equation for plane-parallel layers [McKay et al., 1989]. Line-by-line techniques are used to calculate the absorption coefficients from the methane line parameters compiled in the GEISA database [Husson et al., 1991]. Particles are assumed to scatter as spherical particles, and the Voyager thermal profile is adopted [Lellouch et al., 1989]. The modelled spectra cut off at 1.617 μm as limited by laboratory data for the CH4 spectrum's pressure and temperature dependence. This limitation sets the depth into the troposphere that is accurately modelled to 30 km. Improved laboratory data for weak CH4 transitions in this region would allow analysis of the existing data to probe deeper into the troposphere. Currently, the rich 1.6 μm spectra corresponding to regions below 30 km, where the intriguing cloud and tropospheric haze formation occur, can only be interpreted qualitatively.
 The initial estimate of the haze extinction vertical profile follows a scale height of 21 km. A rudimentary χ2 minimization technique is used to systematically vary the relevant atmospheric parameters – the haze densities in the 9 uppermost layers, the surface albedo, and the CH4 mixing ratio – while determining the mean deviation between the observed spectrum and model output. We fit a series of values (∼10) for each parameter while holding all others at their initial value. The parameter adjustment that yields the smallest χ2 is then updated, and the processes is repeated until the mean squared deviation between the model and the data is below the observation uncertainty (∼20–30 iterations).
 Since multiple solutions for the haze extinction profile give similar values for the minimum χ2 – the extinction may be shown to be anti-correlated in neighboring layers of the model – we select for those extinction profiles that follow the atmospheric scale height with our initial estimate. Profiles that deviate significantly from an exponential decrease with altitude are rejected as unphysical. Changes of a few percent from the best fit optical depths produce modelled spectra outside of the pixel-to-pixel noise of the observations, so uncertainty in the optical depth is estimated to be ∼20%, primarily from the uncertainty in the photometry.
 The extinction profiles of the haze that fit the spectra in Figure 2c are shown in Table 1. The optical depth in the southern polar region from 30–40 km (τS = 0.100) is almost double the corresponding region from the northern hemisphere (τN = 0.062). At higher altitude in the stratosphere (90–150 km) the haze optical depth is greater in the North than in the South, which is in agreement with narrowband images at 1.702 μm [Roe et al., 2002a]. The total haze optical depths in both the North and South agree with values derived from 2 μm images [Gibbard et al., 2004].
Table 1. Comparison of Altitude Dependence of Stratospheric Haze Layers in the Northern Mid Latitudes and Near the Southern Pole.
|Stratospheric Haze Altitude Profiles|
|Altitude (km)a||Pressure (mbar)||Temp (K)||25°N||80°S|
|Optical Depth||Extinction (km−1)||Optical Depth||Extinction (km−1)|
| || ||Sum||0.23|| ||0.30|| |
 The best fit profiles follow the atmospheric number density scale height [Yelle et al., 1997], h = 18 km from 0–60 km and h = 35 km from 80–200 km. This is in general consistent with current microphysical models of particle settling [McKay et al., 2001; Luz et al., 2003] and consistent with extinction profiles determined from modelling 0.89–0.95 μm narrowband HST images [Young et al., 2002].
 Sensitivity tests for the altitude of the lower-stratospheric haze layers, in particular the altitude of the enhancement in the southern hemisphere, are shown in Figure 2d. Systematically shifting the haze opacity parameters down by one layer clearly under-predicts the signal at 1.62 μm and the slope of the spectrum as it increases toward shorter wavelengths. Exchanging the values for haze opacity at 30 and 40 km – representing a haze enhancement higher in the atmosphere – is a much more subtle effect, but nonetheless produces a spectrum that is outside of the observational error. This indicates that while the observed haze may extend slightly below the tropopause, is does not extend deep into the troposphere, and is more likely to be distributed about the tropopause (at 42 km) and above.
 The distinct drop in haze opacity below 30 km observed by Young et al.  is not readily tested with the analysis presented here. While clearing of aerosols by rainout may occur below 30 km, minor changes to the extinction profile (Figure 2d) produce results outside of the observational uncertainty, suggesting that the atmosphere at the south pole is not cleared up to 80 km [McKay et al., 2001].
 It's not clear what concentrates Titan's particles in Titan's southern tropopause. The particles that we witness lie above the daily cumuli that appear at Titan's south pole. These particles reside at altitudes similar to those observed at 2 μm [Bouchez, 2003] and at 0.7–1.0 μm (C. A. Griffith et al., Imaging temporal changes on Titan, submitted to Icarus, 2004). Potentially they are related to the discrete clouds and form from the same updrafts that may form the lower clouds. Another possibility is that they represent haze concentrated dynamically by Titan's circulation [Rannou et al., 2003]. Alternatively we may be observing the condensation of ethane on settling particles, which occurs above that of methane [Barth and Toon, 2003, 2004].