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L. K. Tamppari, R. W. Zurek (JPL), D. A. Paige (UCLA)
Water ice clouds are known to form in the Martian atmosphere. One method for detecting these clouds has been to make use of the broadband 11 micron water ice absorption feature. In particular, this feature has been used in conjunction with the Viking IRTM channels at 11 and 20 microns (Hunt (1979), Christensen and Zurek (1984), and Tamppari et al., (1997)). Since the water ice absorption feature is deepest over the 11 micron IRTM channel, this channel's brightness temperature (T11) will be the combination of the surface and water ice cloud temperatures. Conversely, the 20 micron channel (T20) will only see a contribution from the surface. Previously, it had been assumed that a negative T11-T20 difference was a unique water ice signature.
However, we have found that under certain modelled conditions, both C02 ice and dust in the Martian atmosphere can produce a signature indistinguishable from water ice in the atmosphere. We have shown for a realistic Mars atmospheric mass that the amount of condensed CO2 required to cause the water ice-like signature is not physically possible. The dust case is not so straightforward. However, in the case of the Viking IRTM data, other channels can be employed to successfully detect water ice in the Martian atmosphere. This methodology and it's results will be presented. The Martian surface emissivity (Christensen, 1998) continues to play an important role in the determination of these atmospheric aeorsols. Implications for actual retrieval of atmospheric aeorsols, as well as their particle sizes and opacities, in light of these complications will be discussed.
Also presented will be a complete set of water ice cloud maps extending from late spring (Ls=80) in the first Viking Mars year through late summer (Ls=175) in the second Viking Mars year.