Thursday, September 28, 1995 THE MARS PATHFINDER MISSION, PROJECT AND LANDING SITE 8:30 AM Golombek M. P.* The Mars Pathfinder Mission No abstract available. Spear A. J.* Status of the Mars Pathfinder Project No abstract available. Cook R.* Engineering Constraints on Pathfinder Landing No abstract available. Golombek M. P.* Mars Pathfinder Landing Site Selection No abstract available. Thursday, September 28, 1995 REGIONAL GEOLOGY OF CHRYSE PLANITIA 10:30 AM Tanaka K. L.* Regional Geology and Sedimentary Stratigraphy of Chryse Planitia, Mars Geologic/geomorphic mapping of Chryse Planitia and surrounding terrain [1] has revealed assemblages of landforms indicative of the basin's sedimentary stratigraphy. Here I refine the precision of relative ages of basin units and discuss progressive northeastward thickening of these units based on geomorphic indicators. Using crater densities in units of the geologic map of [1], Late Hesperian and Early Amazonian ages have been cited for deposits resulting from flooding into Chryse Planitia. These crater counts were limited to areas covered by higher-resolution images; area sizes range from 7 to 65 x 10^3 km^2. I have begun a more comprehensive study of crater densities for these same units (see Table 1); thus far I confirm, with much better precision, a Late Hesperian/Early Amazonian age for much larger areas of the units (>2 x 10^5 km^2). Additionally, these ages are based on craters >5 km diameter, which appear fully preserved. Craters <5 km have a somewhat lower density than expected relative to larger craters, which may be explained by obliteration by moderate erosion or deposition. In some areas of the basin, large, partly buried, and heavily modified crater rims outnumber large, relatively unmodified, superposed craters. The density of these large craters in places suggests that the underlying material may be much older than the basin sediments (Early Hesperian or older). In addition, I determine a Late Noachian/Early Hesperian crater density for unmodified ridged plains material (unit Hr) in western Chryse Planitia (Table 1). A few geomorphic clues indicate that sediments progressively thicken from southern and western Chryse Planitia to northeastern Chryse and southern Acidalia Planitiae (in the direction of fluvial flow, as indicated by teardrop-shaped bars), as follows: (1) channel bars and terraces become progressively thinner and disappear; (2) wrinkle ridges change from pristine to subdued to presumably buried in western Chryse (alternatively, they become lower in height and more subject to other degradational processes); (3) large, presedimentation crater rims form raised channel obstructions in Chryse Planitia (i.e., they are prows for many of the channel bars), then become subtle, mostly buried and modified in southern Acidalia, and finally disappear farther north; and (4) fractures and grooves are sparse and primarily align in the direction of flow in southernmost Acidalia Planitia (including within a large, degraded, infilled crater) and then become more polygonal in pattern farther north; some of the grooves have wide flat floors and have been interpreted to form by compaction of sediments hundreds of meters thick [2,3]. How thick are the sediments? Considering that channel terraces, wrinkle ridges, and degraded crater rims may be tens to hundreds of meters high [4], it appears that thicknesses of a few tens of meters may not be readily detected in most Viking images. As thickness approaches 100 m, landforms such as wrinkle ridges, smaller craters, and thinner channel bars may appear subdued. Those features would likely disappear for thicknesses of a few hundred meters, whereas large craters would appear buried and compaction fractures would become evident. As sediment thickness reaches several hundred meters, most large craters would disappear (given that large degraded martian crater rims seldom exceed 500 m in height in the planet's highlands [5], where they generally appear higher than in the northern plains). These geomorphic indicators observed in Chryse indicate progressive, northeastward thickening of basin sediments from on the order of 100 m near the mouths of Kasei and Simud/Tiu/Ares Vallis in western and southern Chryse Planitia to a few hundred meters into southern Acidalia Planitia. The Mars Pathfinder landing site (latitude 19.5 degrees N., longitude 32.8 degrees) lies on unit Hchl below the mouth of Ares Vallis. Here, the unit consists of a broad flood plain marked by bars and a few small knobs. The plain is likely made up of a thin veneer (tens of meters?) of flood deposits (none of the above sediment thickness indicators are observed). Table 1. Crater densities of basin units in Chryse and Acidalia Planitiae, Mars (after mapping of [1]). Cumulative no. of Map Unit Area of Count (km^) Craters per 10^6 km^2 Relative Age* >2 km >5 km AHcc 514,440 280+/-23 60+/-11 LH/EA AHcs 276,340 326+/-34 76+/-17 LH/EA AHcr 236,100 356+/-39 93+/-20 LH/EA Hchl 410,070 285+/-26 83+/-14 LH/EA Hr 230,240 747+/-57 230+/-32 LN/EH *Based on crater-density stratigraphy of [7]; N = Noachian, H = Hesperian, A = Amazonian, E = Early, L = Late. References: [1] Rotto S. and Tanaka K. L. USGS Map I-2441, in press. [2] Lucchitta B. K. et al. (1986) Proc. LPSC 17th, in JGR, 91, E166-E174. [3] McGill G. E. and Hills L. S. (1992) JGR, 97, 2633-2647. [4] Golombek M. P. et al. (1991) Proc. LPS, Vol. 21, 670-693. [5] Roth L. E. et al. (1989) Icarus, 79, 289-310. [6] Kargel J. S. et al. (1995) JGR, 100, 5351-5368. [7] Tanaka K. L. (1986) Proc. LPSC 17th, in JGR, 91, E139- E158. Rice Jr. J. W.* The Geologic Mapping of the Ares Vallis Region Geologic mapping is being conducted, at a scale of 1:500,000, on the Ares Vallis outflow region in southern Chryse Planitia. Ares Vallis is a 1500-km-long outflow channel that drains the cratered highlands to the south before debouching into the vast, topographically enclosed Chryse Basin. This area contains both erosional and depositional landforms (streamlined islands, flood-flow-modified impact craters and associated ejecta blankets, terraces, and knobby and etched terrains). The identification of deltaic deposits on Mars has been documented [1,2]. However, a more detailed account is being put forth for the depositional and erosional history of Ares Vallis based on the mapping being conducted. The region mapped contains the landing site (19.5 degrees N, 32.8 degrees W) for Mars Pathfinder. This site was chosen by Mars Pathfinder science personnel in June 1994. The Ares Vallis site was originally proposed by investigators J.W. Rice and R.O. Kuzmin et al. in April 1994 [3,4]. This mapping effort will aid in the analysis and interpretation of landforms and surface materials found at the landing site. The landing is scheduled for July 4, 1997. The 1:500,000-scale mapping of Ares Vallis is an order of magnitude better than previous maps of this region [5,6]. Viking orbiter images with resolutions on the order of 40 m/pixel are being used to define the various materials present in this region. Information on physical properties [7-10] are also being incorporated into this map. Finally, a more complete picture of the depositional history of the Chryse Basin is emerging and a facies model has been proposed [11] that should help explain both the Viking 1 and Pathfinder landing sites. References: [1] Rice J. W. and Scott D. H. (1993) LPI Tech. Rpt. 93-06, 45-46. [2] Rice J. W. and De Hon R. A. (1995) USGS Map I-2432. [3] Rice J. W. (1994) LPI Tech. Rpt. 94-04, 36. [4] Kuzmin R. O. et al. (1994) LPI Tech. Rpt. 94-04, 30-31. [5] Wilhelms D. E. (1976) USGS Map I-895. [6] Greeley R. et al. (1977) JGR, 82, 4093-4109. [7] Golombek M. P. et al. (1995) LPSC XXVI, 481-482. [8] Edgett K. S. (1995) LPSC XXVI, 353- 354. [9] Edgett K. S., this volume. [10] Parker T. J., this volume. [11] Rice J. W. and Edgett K. S., this volume. Crumpler L. S.* Geologic Mapping Traverse of the Highland to Lowland Transition in an Area Adjacent to the Mars Pathfinder Region Introduction: Mapping in the region between Central Chryse Planitia and Xanthe Terra is of relevance to understanding the Ares Vallis region because (1) it is the closest area to the proposed Pathfinder landing site for which regional 1:500,000-scale mapping and local ground-surface geologic characteristics have been determined (Fig. 1); (2) it traverses a region of the highland-lowland transition at a similar latitude and in an area of similar overall geologic units as the Pathfinder site; and (3) in contrast to the proposed landing site in Ares Vallis, it represents a portion of the Chryse Basin margin that is relatively little modified by outflow channels, and thus offers the opportunity to establish the nature of the regional surface in the transition as it might have appeared prior to Ares Vallis outwash. Thus, detailed mapping in this region is important to both the characterization of the highland-to-lowland transition in a region of extensive fluvial deposition and erosion, and to making educated inferences about the regional substrate geology of the Mars Pathfinder site. Objectives: Geotraverse mapping. This work builds on mapping recently completed in central Chryse Planitia in the vicinity of Mutch Memorial Station (MMS or VL-1 )from central Chryse Planitia southward onto the Noachian highlands (MTM 15047, 10047, and 05047 1:500,000-scale photomosaic sheets) in an effort to extend our detailed surface geologic knowledge outward from the Viking Lander 1 site, where we have actual ground truth, to surrounding geologic units. The approach taken by stacking three 1:500,000-scale quadrangles [1,2] is that of a geologic traverse along a relatively narrow corridor (transect or geotraverse) at a large map scale from a relatively young but typical surface in a lowland region to the complex and older surface of an adjacent highland. The goal is a geologic cross section across the lowland-to-highland boundary in an area where the transition is both topographically and geologically relatively gradual in comparison with many other localities around the margins of Chryse Planitia. Regional setting. The oldest map units are Noachian to Hesperian surface materials near the southern edge of Chryse Planitia. These continue northward where they are overlain in central Chryse Planitia, in the vicinity of MMS, by younger ridged plains that are interpreted to be outwash deposits from Maja Vallis to the west [1,2]. Hesperian ridged plains are among the oldest post-Noachian materials exposed in the highland-lowland boundary regions, and the large exposure here in southern Chryse Planitia affords an opportunity to assess some characteristics of this surface prior to outwash deposition and surface scouring. The highland-lowland boundary strikes approximately northwest- southeast across the center of the map area. At 1:500,000 scale the actual contact between the Hesperian ridged plains and the highlands material appears transitional in character, as numerous low hills or knobs, possibly residuals of cratered highland material, protrude through the ridged plains unit. Similar low knobs occur in the central basin east of MMS within the 1:500,000-scale quads connecting this area with the Pathfinder region. Numerous mare-type ridges, generally interpreted to be the result of small amounts of shortening and compression [5,6], occur on the Hesperian plains. Some arcuate arrangements of mare-type ridges, near the highland-lowland boundary within the Hesperian plains, are clearly superimposed on the buried rims of large highland craters. This observation implies that the highland surface may be preserved to some extent beneath the adjacent plains materials. Several sinuous-type channels (Hypanis and Nanedi Valles) trend southwest to northeast within the highland part of the proposed map area. High-resolution images across one of these affords the opportunity to map this channel in detail and an assessment may be made of the local evidence for the origin and modification of this type of channel. Bends in Mars sinuous channels are frequently characterized by circular or constant radius curvature in contrast to the more asymptotic curvature of river meanders on Earth. This might suggest that factors in addition to normal stream dynamics, such as variation in material properties resulting from the probable brecciated or cratered nature of the highlands, among other environmental influences (re-used lava channels), may have exerted a control on the sinuosity [7]. Summary: Several questions of regional, local, and topical significance to the Pathfinder site can be addressed through mapping the Xanthe Terra to Chryse Planitia traverse: What is the geologic history and stratigraphy of the transitional boundary between the highlands and lowlands throughout this region [8]? What is the origin of the numerous knobs within the transitional region; are they residual highland materials? What evidence is there for the origin and the nature of the emplacement of the Hesperian ridged plains? What are the gradients in thickness of the Hesperian ridged-plains material at the boundary, and what might these tell us about the underlying gradients of the highland surface? Are the surfaces of the intercrater highland plains the same material that forms the lowland plains? If not, why are the crater ages similar? What is the regional geologic section and how does it relate to the regional topographic characteristics? And finally, what is the evidence for the origin of the highland-lowland boundary in this region and can it test theories [3] of the origin of the global dichotomy? References: [1] Crumpler et al.(1995) 1:500,000 scale Misc. Inv. Geol. Map, USGS, in preparation. [2] Craddock R. A. et al. (1992) Proc. LPS, Vol. 23, 257-258. [3] McGill G. E. and Squyres S. W. (1991) Icarus, 93, 386-393. [4] Scott D. H. and Tanaka K. L. (1986) USGS, 1:15,000,000 Series, Atlas of Mars, Western Region, 1-1 802A. [5] Plescia J. B. and Golombek M. P. (1989) LPI Tech Report 89-04, 64-65. [6] Watters T. R. (1988) JGR, 93, 10236-10254. [7] Mars Channel Working Group (1983) GSA Bull., 94, 1035-1054; Rotto S. and Tanaka K. (1993) USGS Misc. Inv. Map 1-2441. Figure 1, which appears in the hard copy, shows the area previously mapped in central Chryse Planitia (solid boxes). The proposed Mars Pathfinder site is located several hundred kilometers to the east in a similar regional setting, but the Xanthe-Chryse transect may preserve substrate geologic characteristics that were originally present in the Ares Vallis area prior to outwash effects. Approximate location of highland-lowland boundary as mapped by Scott and Tanaka [4] is shown as a sinuous dashed line. Thursday, September 28, 1995 CHRYSE FLOODING 2:00 PM Rice Jr. J. W.* Edgett K. S. Facies on Mars: A Model for the Chryse Basin Outflow Sediments Based on Analog Studies of Icelandic Sandar and the Ephrata Fan of the Channeled Scabland, Washington Nicholaus Steno introduced the concept of facies to the study of geology in 1669 [1]. Here, we describe "facies" in terms of laterally adjacent coeval sedimentary units. In this paper we extend the facies concept to the surface of Mars, specifically the basin of Chryse Planitia. This should aid in the interpretation of the Mars Pathfinder landing site. Mars Pathfinder is scheduled to land on the Ares Vallis floodplain on July 4, 1997. A sedimentary facies model has been developed for Chryse based on the combined analysis of geomorphology, albedo, and thermophysical properties. Seven major outflow channels debouch into the topographically enclosed Chryse Basin. We are developing a facies map to illustrate the areal extent of the sedimentary deposits exposed at the surface (those that result from the most recent outflow events in Chryse). A facies map indicates the distribution of different sedimentary facies over a geographic area for a specified moment or interval of geologic time [2]. Such maps are important for interpreting ancient depositional environments and paleogeography. The sandar of southern Iceland and the Ephrata Fan of the Channeled Scabland, Washington, are the analog surfaces used in this study. Sandar deposits (Icelandic glacial outwash fans/plains) cover extensive areas (up to 800 km^2) of the Icelandic coasts. Most of these deposits were laid down by jokulhlaups, which occur frequently and are caused by the abrupt drainage of ice-dammed lakes or subglacial volcanic eruptions [3,4]. The high current velocities and discharges associated with these outbursts are very significant factors for sediment transport along these fans. Discharges from these floods can be up to 100,000 m^3/s and can last for several days [4]. Sandar deposits display a succession of facies characterized by differences in fan gradient, clast size, bar morphology, and sedimentary structures [5]. These facies are described as follows: (1) "proximal facies," composed of steep gradients, large clasts, longitudinal bars, and few sinuous streams; (2) "mid-fan facies," typified by moderate gradients, clasts ranging from tens of centimeters to sand size, longitudinal bars, and numerous anabranching and bifurcating streams; (3) "distal facies," characterized by low gradients, mostly sand-sized material, linguoid bars, and branches that merge and form very wide shallow streams. Sandar are relevant to martian outflow channel deposits because they are the direct result of rapid sediment deposition from an expanding flow after leaving a constricted area. This is observed at the mouth of Maja Vallis, where the channel emerges from its constriction in the highlands and empties into the vast Chryse Basin [6]. Another important point to note is that discharge rates on sandar vary greatly. Discharges into Chryse were also probably variable due to the sudden and sporadic release of water associated with the production of the chaotic terrain. The other analog surface is the Ephrata Fan. This fan is an extensive, coarse gravel unit deposited subfluvially [7]. The Ephrata Fan was emplaced as flood waters emerged from the Soap Lake constriction, located at the mouth of the Lower Grand Coulee, and expanded outward into the Quincy Basin. Clast sizes on the fan surface range from house- sized boulders in the proximal fan zone to the sand of the Moses Lake dune field located in the fan's distal zone. We postulate that the Ephrata Fan in Quincy Basin is a very small analog/model of the outflow deposits in Chryse Basin. Rock abundance, thermal inertia, albedo, and geomorphic evidence indicate that the Chryse Basin follows the general sedimentary facies models of Icelandic sandar deposits and the Ephrata Fan. The transition to distal facies in Chryse Basin is most easy to recognize; it is indicated by the change in albedo from Chryse Planitia (0.21-0.26) to Acidalia Planitia (0.14-0.16). The lower albedo of Acidalia results from the presence of windblown sand [8], similar to the presence of eolian- reworked sand in the Moses Lake dune field on the Ephrata Fan. Viking 1 lander investigators in the 1970s commented on the relative paucity of sand at the landing site [e.g., 9]; we argue that this is because very little sand was deposited at that location; instead, the outflow floods washed most sand further out into Acidalia. Today, Acidalia is both sandy and rocky, probably owing to eolian erosion that has exposed buried rocks and possibly stripped away sand (perhaps transported to the north polar dune field). Both the Viking 1 and Mars Pathfinder landing sites occur in what we map as the mid-fan facies; these areas are rocky with little sand, and the rocks tend to trap windblown silt and dust. That the Viking 1 site might once have been inundated by floods is apparent, considering its proximity to the mouths of Maja and Kasei Valles. Our facies model for the last flood deposits in Chryse Basin is outlined as follows: (1) The proximal facies zone extends from chaotic terrain (channel sources) to the mouth of channels, a distance of ~1000 km; (2) mid-fan facies zone (transition zone) extends from channel mouths to about 800 km into Chryse bBsin; and (3) the distal facies zone begins at the albedo change between Chryse and Acidalia. References: [1] Walker R. G., ed. (1984) Facies Models, Geoscience Canada Reprnt. Ser. 1, 317. [2] Dott R. H. and Batten R. L. (1981) Evolution of the Earth, McGraw-Hill, 113 pp. [3] Krigstrom A. (1962) Geog. Ann., 44, 328-346. [4] Thorarinnsson S. (1960) Inter. Geol. Cong. 21st Session, 33-45. [5] Boothroyd J. and G. Ashley (1975) SEPM Spec. Pub., 23, 193-222. [6] Rice J. W. Jr. and DeHon R. A. (1995) USGS Map, I-2432. [7] Baker V. R. and Nummedal D. (1978) The Channeled Scabland, NASA, 186 pp. [8] Arvidson R. E. et al. (1989) JGR, 94, 1573-1587. [9] Sagan C. et al. (1977) JGR, 82, 4430-4438. Komatsu G.* Baker V. R. Catastrophic Paleoflooding at the Pathfinder Landing Site: Ares Vallis, Mars Paleodischarges for the martian outflow channels have been estimated by a number of researchers [1-3] using a modified Manning equation. This study updates our preliminary work [4]. The flow velocity (v) can be estimated by the Chezy equation, v = C(ds )^1/2 in which C is the Chezy coefficient; C = (2g/C(sub)f)^1/2 = 1/n(d)^1/6 (where g is gravity; C(sub)f is the friction coefficient; n is the Manning coefficient; d is depth); and s is energy slope. The Manning coefficient for Mars (n(sub)M) can be related to the empirical terrestrial Manning coefficient (n(sub)E), by the equation n(sub)M = n(sub)E(g(sub)E/g(sub)M)1/2 = 1.62n(sub)E (where g(sub)E is terrestrial gravity; g(sub)M is martian gravity). More realistically, the empirical Manning coefficient on Earth ranges over a factor of about 2 and, for our application, the influence on the final result is minimal. The Manning coefficient chosen for Mars (n(sub)M) paleoflooding is 0.0324 (n(sub)E = 0.02). We selected a reach where the channel is well defined and unusually deep. For simplicity, the 10 cross sections of the Ares Vallis are assumed to represent the paleogeometry of the channel at the time of flooding [4]. The slope between cross sections 1-8 is too small to measure, so we have assumed it to be 0.0001. Because high-water marks, such as trim lines and deposits, are not apparent on the available Viking imagery, we assumed that the water surface reached the rims of the channel. We also assumed that the flood did not overflow the rims of the channel. The peak discharge was calculated for each cross section, and we took the lowest peak discharge out of the 10 as the best estimate for the entire reach (Table 1). The resulting peak discharge is 0.57 x 10^9 m^3/s. This discharge rate is of the same order as the estimates for the Kasei Vallis [2]. For this discharge, the flow velocity ranges from tens of m/s to over 100 m/s. Froude numbers suggest that, at the steep section, the flow was supercritical and, at the less steep section, the flow was subcritical. We expect that the water may have incised the channel and, therefore, may not have filled it to the rim. In this case, the discharge could well have been much lower than the estimated peak. Calculations by [5] show that flow velocities of tens of m/s transport basalt fragments of one to several meters in diameter. A 100 m/s flow would transport basalt boulders larger than 10 m in diameter even by suspension. However, basalts are prominently jointed due to the cooling, and fragmentation during the transportation would cause considerable reduction of boulder sizes. The above calculation was applied to the deep section of the channel to estimate the peak discharge. The next U.S. Mars mission, Mars Pathfinder, has its primary candidate landing site located about 100 km to the north of the mouth of the Ares Vallis, which is one of several huge outflow channels debouching into the Chryse Planitia. At the landing site, the flood levels were estimated to be lower than the constricted section. This is the result of the pronounced expansion of the channel reach. However, the rich evidence of erosional landforms around the landing site suggests that even on the outwash fan deposited in this expansion, the peak flood power was still very high. The maximum flow velocity may still have been several tens of m/s. As the flow velocity decreases, the sorting of sediments occurs. The sediments transported through constricted reaches will settle on the outwash fan of the expanding reach. Hence, on the outwash fan, it is considered that the largest-sized boulders accumulated near the mouth, and the grain size should have decreased away from the mouth of the channel. As the flood level lowers toward the end of flood event, the sediment transport decreases. This leads to deposition of smaller-sized grains. As a result, the flood deposit stratigraphy probably displays an upward decrease of grain size. Moreover, the flooding may have occurred in multiple events, which would have caused redistribution of flood deposits. The preservation of the original flood deposits is also subject to modification by the other geological processes, including glaciation, impact cratering, and eolian processes. The size distribution of the boulders and smaller grains observed at the landing site depends on these factors as well as on the primary flood depositional processes. References: [1] Baker V. R. (1982) The Channels on Mars. [2] Robinson M. S. and Tanaka K. L. (1990) Geology, 18, 902-905. [3] De Hon R. A. and Pani E. A. (1993) JGR, 98, 9129-9138. [4] Komatsu G. et al. (1995) LPSC XXVI, 781-782. [5] Komar P. D. (1980) Icarus, 42, 317-329. Table 1. Flow Reconstruction using Manning Equation Section 10 Section 1 Manning Coefficient 0.032 0.032 Depth (m) 398.40 985 Slope 0.02 0.0001 Velocity (m/s) 148.91 25.43 Froude Number 5.46 0.49 Discharge (10^9 m^3/s) 0.57 0.57 Craddock R. A.* Tanaka K. L. Estimates of the Maximum and Minimum Flow Velocities of the Circum-Chryse Outflow Channels The Mars Pathfinder landing site was chosen in part because of its potential to offer investigators the opportunity to analyze a variety of material from different locations [1]. To know what we're getting out of the "grab bag" it is imperative that the detailed geology and hydraulic history of southern Chryse Planitia and the circum-Chryse outflow channel complex be understood ahead of time. Crude estimates of the maximum channel flow velocities can be made simply by knowing the depth and slopes of the outflow channels themselves. Although these characteristics have been derived in part by stereophotogrammetry [2], they are subject to a considerable amount of error, or ~+/- 1 km in the southern Chryse area [3]. Fortunately some Earth-based radar data exist that are both reasonably accurate and provide the spatial coverage necessary for determining the slopes of some of the channels [4,5]. Using these data, the bed shear stress of a flow, or the retarding stress at the base of a flow, tau(sub)b, can be estimated from the depth-slope formula tau(sub)b = rho-ghS (1) where rho is the density of the fluid, g is gravitational acceleration, h is the flow (or channel) depth, and S is the slope of the channel. This is equal to the bottom stress created by a flow, tau, where tau = rho C(sub)f u bar^2 (2) and C(sub)f is a dimensionless drag coefficient, and u bar is the mean flow velocity. Thus, the mean flow velocity for a channel can be calculated from u bar = (ghS/C(sub)f)^1/2 (3) The dimensionless drag coefficient can be adjusted for gravity by the expression C(sub)f = g((n^2)/(h^1/3)) (4) where n is the Manning roughness coefficient (units of s/m^(1/2)), which has been derived empirically from terrestrial observations. Application of an appropriate Manning roughness coefficient, n, to martian outflow channels is uncertain, so Robinson and Tanaka [6] used a range of values (0.015-0.035) in estimating the flow velocities of Kasei Valles. Because these values describe most environments free of vegetation, they appear to be reasonable values to apply to the circum-Chryse channels. Estimates of the mean flow velocities were calculated from this method (Table 1); however, at best these represent maximum values. Large-scale geologic mapping indicates that most channels were subjected to multiple episodes of flooding [7-9], which suggests that the channels may not have been completely full of water at any one time (i.e., bankfull discharge). This method is also not directly applicable to Simud and Tiu Valles because the Earth-based radar data indicate a positive downslope gradient [4,5], which may be due to modification (e.g., slack-water deposition) postdating channel formation. An alternative method for calculating lower channel flow velocities may be in the thermal inertia data made available by the Viking Infrared Thermal Mapper (IRTM) [10] and the Phobos Thermoskan [11] instruments. These data can be used to estimate the critical shear stress, tau(sub)cr, by assuming: (1) the effective particle size measured by the IRTM represents the median-sized bed material, D(sub)50; (2) the channel bed is planar; (3) the sorting coefficient (standard deviation) is 2.0 phi, implying that the material is poorly sorted, typical of most gravel-bed streams; (4) the calculated D(sub)84 particle size was the minimum-sized particle in motion at one time; (5) the density of the material is that of basalt (3.3 g/cm^3); and (6) the fluid that formed Shalbatana Vallis was water at 10 degrees C. The tacit assumption made is that the thermal inertia values measured in the channel represent unmodified channel materials. The validity of this assumption is discussed in general by Betts and Murray [11]. Shields [12] derived empirical relations for the dimensionless grain parameter, zeta*, and the dimensionless boundary shear stress, tau*. The dimensionless grain parameter is defined as zeta* = (D^3(rho(sub)s-rho)g)/(nu^2 rho) (5) where D is the particle diameter (in this case the D(sub)84 particle size derived from the thermal inertia data expressed in cm), rho(sub)s is the particle density (3.3 g/cm^3), rho is the fluid density (water at 10 degrees C or ~1.0 g/cm^3), g is the acceleration of gravity, and nu is the kinematic viscosity of the fluid (1.304 x 10^-2 cm^2/s for water at 10 degrees C). From the assumptions given, zeta* simplifies to zeta* = D^3 x 5.07 x 10^6 cm^-3 (6) From Shields' [12] curve, values for the dimensionless boundary shear stress, tau*, can be determined. For values of zeta* less than ~400, Shields extrapolated his curve. Although equation (7) shows that it is unlikely derived values of zeta* will be <400, White's [13] experimental data can be used to determine values for tau* in this range. The dimensionless boundary shear stress, tau*, is tau* = (tau(sub)cr)/((rho(sub)s-rho)gD) (7) where tau(sub)cr is the critical boundary shear stress needed to initiate sediment motion. This is assumed to be the bottom shear stress, tau(sub)b, during the waning stages of channel formation and can be used to estimate lower values of the channel velocities. The shear velocity, u* (expressed in cm/s), is u* = the square root of tau(sub)cr/rho = the square root of tau(sub)b/rho (8) By substituting equation (1) for tau(sub)b, equation (9) becomes u* = the square root of ghS (9) which also allows the depth (h) of the water in the channel during the low flood stages to be determined. This depth should be much less than the full depth of the channel. As Komar [14,15] notes, it is better to analyze the flow in terms of u* than u-bar due to the uncertainties in estimating reasonable values for C(sub)f. However, values of u-bar are more intuitive. These can be calculated from the relationship u bar = the square root of C(sub)f u* = g^(1/2)nh^(-1/6)u* (10) Of course, in order to estimate values of u-bar from the thermal- inertia-derived values of u*, reasonable values of the Manning coefficient, n, must be used. Obviously these should be the same range of values used to determine the channel flow velocities at bankful discharge (0.015-0.035). Table 2 lists the possible channel flow velocities and depths determined from the available thermal inertia data. They represent minimum estimates because the material contained on the surface of the channel floors, if unmodified, was probably emplaced during the waning stages of flooding. Actual channel flow velocities probably fall between the two values presented in Tables 1 and 2. Acknowledgements: This research is supported in part by NASA grant NAGW- 3365 (Smithsonian Institution). References: [1] Golombek M. P. (1995) LPS XXVI, 475-476. [2] U.S. Geological Survey (1991) Map I-2160 1:15M-scale. [3] Wu S. S. C. and Howington-Kraus A. (1988) NASA TM 4041, 551. [4] Downs G. S. et al. (1982) JGR, 87, 9747-9754. [5] Lucchitta B. K. and H. M. Ferguson (1983) JGR, 88, A553-A568. [6] Robinson M. S. and Tanaka K. L. (1990) Geology, 18, 902-905. [7] Chapman M. G. and Scott D. H. (1989) Proc. LPSC, 19th, 367-375. [8] DeHon R. A. and Pani E.A. (1992) Proc. LPS, Vol. 22, 63-71. [9] Craddock R. A. et al. (1993) LPS, XXIV, 335-336. [10] Craddock R. A. (1988) M.S. thesis, Arizona State Univ., 110 p. [11] Betts B. H. and Murray B.C. (1994) JGR, 99, 1983-1996. [12] Shields A. (1936) U.S. Dept. Agriculture, Soil Conservation Service Coop. Lab., Calif. Inst. Tech, 26. [13] White S. J. (1970) Nature, 228, 152-152. [14] Komar P. D. (1979) Icarus, 37, 156-181. [15] Komar P. D. (1980) Icarus, 42, 317-329. [16] Henry L. Y. and Zimbelman J.R. (1987) Papers Presented at the Third Annual Summer Intern Conf., LPI, 13-14. Table 1. Estimates of channel flow velocities determined from Earth- based radar derived measurements of channel depths and slopes. Thursday, September 28, 1995 LANDING ELLIPSE GEOLOGY AND SEDIMENTOLOGY 3:30 PM Kuzmin R. O.* Greeley R. Mars-Pathfinder: Geology of the Landing Site Ellipse Nineteen years after the last successful landing on Mars by Viking, the Mars-Pathfinder mission will provide new data for the martian surface and atmosphere. An important advance will be the inclusion of a microrover for sampling and imaging surface rocks [1]. The Mars-Pathfinder landing site must satisfy engineering constraints such as topographic level, surface roughness, and latitude to provide maximum solar power for the lander and rover. For scientific goals, the landing site should be in a region that has a wide variety of rock types in a small area. From terrestrial experience [2], such a site is typical for channel and canyon deltas that drain highland or mountain areas of various ages and rock types. For example, the Furnace Creek alluvial fan in Death Valley contains basalt, rhyolite, dionte, quartzite, limestone, and gabbro within 1 m radius of a simulated landing site. These rocks were transported from the surrounding mountains by intermittent streams. There is a good probability that a similar depositional regime can be found on Mars in deltas of the large outflow valleys. Ares and Tiu Valles are among the largest outflow channels on Mars. Their watersheds include various ancient Noachian highland materials and Hesperian-age ridged plains. Consequently, the delta-fluvial deposits of these channels, selected as the nominal site for Pathfinder, should include a wide variety of rock types and ages. Analysis of Viking images (30 m/pixel) for the nominal landing site ellipse on the delta deposits of Ares-Tiu Valles shows that the surface is mostly smooth and slightly undulating, resembling the surface of the Viking l landing site. Crater counts of the landing site ellipse suggest a late Hesperian age. The surface of the site is complicated locally by streamlined islands that have terraces. About eight terrace levels are recognized on the flanks of the streamlined islands. The terraces could result from erosion of layered highlands or from fluvial deposition during different stages of flooding. The ellipse area also includes knobs and fluidized ejecta from large impact craters. Scabby or etched terrain occupies the western edge of the landing site ellipse. This terrain may have resulted from fluvial plucking and subsequent eolian deflation. The main trend of the etched terrain is parallel with Tiu Vallis, suggesting an origin related to the formation of the channel. To the north, the etched terrain grades into incipient chaotic terrain, indicating that the etching process may have been enhanced by ground- sapping processes [5]. Large blocks ranging from the limit of resolution to 1 km are found in the northwestern part of the ellipse. Some blocks may be remnants of crater rims and highland material; others could be remnants of eroded fluvial deposits. Craters and blocks constitute about 4% of the surface area [4]. Locally, many small impact craters (100-200 m in diameter) have dark halos, suggesting that they excavated dark albedo subsurface material. Albedo variations on the plains of the landing site ellipse may be due to variations in thickness of sand and dust. Typical values of the thermal inertia in the area range from 9.8-12.9 [6] to 8.4-10.9 [7] for fine-component material. This suggests possible particle sizes of 500 to 2300 micrometers [8], typical sizes for medium and very coarse sand. Thermal remote sensing data of Mars [9] suggest a rock abundance of the landing site to be similar to the Viking 1 landing site, except for the fine-component material. The current modification of the landing site by eolian process is demonstrated by crater streaks and fine lineations on the surface. The Ares-Tiu Valles delta is characterized with weak radar echoes [10], high values of fine-component thermal inertia, and moderate rock abundances. We suggest that a reasonable explanation for this discrepancy is that sand-gravel-rocky surface material is interbedded with finer deposits and scattered rock fragments, and that the thickness of the fine component is less than the penetration depth of the radar. On the Earth, such a surface structure is typical for some fluvial plains and desert pavements. References: [1] Golombek M. P. (1995) LPSC XXVI, 475-476. [2] Greeley R. and Kuzmin R. O. (1994) Mars Pathfinder Landing Site Warkshop, Houston., 33-34. [3] Kuzmin R. O. et al. (1994) Mars Pathfinder Landing Site Workshop, Houston, 35-36. [4] Golombek M. P. (1995) LPSC XXVI, 481-482. [5] Carr M. H. et al. (1976) Science, 193, 766-776. [6] Christensen P. R. and Moore H. (1992) in Mars, 687-729, Univ. of Arizona, Tucson. [7] Christensen P. R. and Malin M. (1993) LPS XXIV, 285. [8] Edgett K. S. and Christensen P. R. (1994) JGR, 99, 1997-2018. [9] Christensen P. R.(1986) Icarus, 68, 217. [10] Tyler G. L. et al. (1976) Science, 193, 812. Parker T. J.* Viking Stereo of the Ares Vallis Site: Sedimentological Implications The Ares Vallis site has one very important advantage over other potential grab-bag sites on Mars that might be accessible to the spacecraft. Because this site was one of the first choices for Viking Lander 1, excellent, same-orbit stereo images were acquired early in the Viking mission. These stereo pairs are on the order of 40-m/pixel resolution with a separation angle between looks of about 48 degrees, corresponding to about a 40-m vertical precision for topographic measurements based on parallax displacement. Subtle topography on the plains is most easily viewed in stereo, and is indispensable to studying the geology of the landing site and surrounding region. In addition, it will be extremely useful in determining the exact location of the Pathfinder lander on the martian surface after the landing, by potentially providing a correlation between objects viewed on the horizon with a three-dimensional aerial view at least several months earlier than Mars Global Surveyor images could "find" the lander. The Pathfinder landing ellipse was placed within a plains region beyond the mouth of Ares Vallis to avoid large topographic hazards. But this plains region is not without its interesting, and in some cases very problematic, landforms. The ellipse contains the following features: (1) primary impact craters; (2) small secondary impact craters; (3) streamlined islands and (4) longitudinal grooves [e.g., 1]; (5) "scabland" or "etched" terrain [e.g., 1], mostly outside ellipse to west); (6) pancakelike shields (7), dikelike structures [e.g., 2,3], and (8) knobs or "buttes" [e.g., 3]; and (9) a previously undetected, subtle undulating or hummocky topography to the plains surface. The nature of these landforms has important bearing on what we can anticipate Pathfinder may see once the first images are transmitted to Earth. But why is there even the least confusion about the nature of the landing site? After all, it lies just beyond the mouth of one of the largest outflow channels on Mars, orders of magnitude larger than the Channeled Scablands on Earth. Shouldn't we expect a fluvial sedimentary deposit? Probably. But a number of things could have happened to the site to change that in the more than 1.8 Ga since the latest Ares Vallis flood, including (1) eolian reworking or burial; (2) permafrost modification; (3) burial by lacustrine or marine sediment; (4) desiccation; and (5) burial by volcanic plains. Points (2) and (4) would have little effect on the lithology at the lander scale, since they involve the least modification of the postflood surface. Points (1), (3) and (5), on the other hand, could drastically affect the kinds of materials Pathfinder would examine. Thankfully, due to the relatively high thermal inertias (see discussions elsewhere in this volume) any eolian or fine lacustrine blanket is not likely to be thick enough to preclude access to surface rocks. To further complicate the picture, even without subsequent modifiers, fluvial processes may leave behind (1) fluvial deposits; (2) relatively unmodified, pre-existing terrain (sediment "bypassing" or throughput); or (3) an eroded, pre-existing bedrock surface. The differences between these three are determined by the channel's energy at this location relative to reaches both upstream and downstream. High-energy flows will transport sediment through the site or scour the existing bedrock surface, whereas a drop in transport energies beyond the mouth of the channel will result in sediment deposition. Again, based on these considerations, we should probably expect a fluvial sedimentary deposit. What about the problematic landforms? How does the range of interpretations for these features affect what Pathfinder is likely to find? Taking these one at a time and based on a preliminary examination of the stereo pairs, I offer the following suggestions: (1) and (2): Large craters are relatively few within the landing ellipse. Small secondaries are much more abundant, and occur in a number of large clusters across the ellipse. Small craters probably won't pose a serious landing hazard. Landing near, even within, a crater may increase the grab-bag potential of the specific landing site by providing blocks of ejected material. (3) and (4): Streamlined islands and longitudinal grooves scattered within and around the landing ellipse suggest a primary fluvial plains surface not buried by thick eolian, lacustrine, or volcanic deposits. (5): "Scabland" terrain to the west of the ellipse appears to be due to eolian deflation after, rather than fluvial plucking during, the flood. Etching of the plains from around a few of the pancakelike shields and exhumation of several dikes appears to have occurred in this area. In addition, at least one small crater's ejecta blanket appears to have locally armored the plains surface, protecting it from the etching process. Other occurrences in this region of less-pronounced "scabland" terrain are more intriguing. One notable example in Viking images 004a19 and 004a79 (stereo pair) has an almost parabolic shape oriented ~45 degrees with respect to local channel scour, and therefore probably unrelated to it. (6) and (7): The pancakelike shields and dikes exhibit no evidence of either channel scour or streamlining, further indicating they formed after the flood. Though there is no reason volcanic landforms couldn't be found in a fluvial setting (they often are on Earth), these occur exclusively on the plains surface in this region, never on the flanks or "tails" of streamlined islands. This might suggest that they are related in some way to the plains "deposits," perhaps as pseudovolcanic sedimentary structures, such as sand volcanos, mud laccoliths, and clastic dikes. Though there are no terrestrial analogs to these features at such large scales, the martian floods were certainly capable of rapidly emplacing thick, wet sedimentary deposits that could have experienced dewatering on a grand scale after cessation of the flood. (8): The knobs may be remnants of resistant material left behind after the flood. However, most show no streamlining, whereas several similar- size streamlined knobs can be found throughout the southern Chryse Basin. In addition, many clusters of these knobs can be found, particularly in the "shadows" of flow obstructions, similar to occurrences of large boulders in the scablands, but on a scale too large for suspension, or even bedload transport. Alternatively, they could be similar to "rafted" blocks in terrestrial floods. Most notable of these are fragments of concrete from the Saint Francis Dam in southern California, which are on the order of thousands of tons in mass and were rafted more than a kilometer downstream when the 47 x 10^6 m^3 reservoir failed catastrophically in 1928. (9): The gently undulating surface of the plains within the ellipse (most pronounced near its center) became apparent when I contrast- enhanced the Viking images to bring out subtle details within the plains (typically saturating sunward and shaded slopes of knobs and craters). This surface almost appears rippled, though the ripples are poorly organized and trend with their crests parallel to the flow direction. The wavelengths are on the order of a few kilometers and amplitudes are up to a few tens of meters. They may be hummocks or lobes of sediment deposited by the flood. References: [1] Baker V. R. (1973) GSA Spec. Paper 144, 79. [2] Greeley R. et al. (1977) JGR, 82, 4093. [3] Hodges C. A. and Moore H. J. (1994) USGS Prof. Pap. 1534, 194. Thursday, September 28, 1995 POSTER SESSION 5:00 PM Golombek M. P. Parker T. J. Moore H. J. Slade M. A. Jurgens R. F. Mitchell D. L. Characteristics of the Mars Pathfinder Landing Site The preliminary landing site selected for Mars Pathfinder and a candidate for final validation is at the mouth of Ares Vallis in southeastern Chryse Planitia (19.5 degrees N, 32.8 degrees W) [1]. Ares Vallis is a large outflow channel that drained the highlands to the southeast. The region contains large streamlined islands of older plateau materials and smooth outflow deposits in the channels. Many of the streamlined islands have terraces that may represent layering, downcutting stages during flooding, or both. Linear features extending downstream from many of these islands may be longitudinal grooves [2] or other primary flow features, indicating a surface composed of materials deposited by the floods. Potential source materials for the outflow deposits include ancient Noachian crustal units (Npl(sub)1, Npl(sub)2), Hesperian Ridged Plains (Hr), and a variety of reworked channel materials (Hcht, Hch, Hchp). The 100 x 200 km landing ellipse is located on a broadly undulating, level surface between (1) streamlined islands and knobby terrain to the east, (2) large streamlined islands to the south, (3) large fresh impact craters to the north, and (4) scabby or etched terrain to the west (Fig. 1). Scabby or etched terrain is appears rough in high-resolution images (~40 m/pixel), with 10-30-m-high scarps that may have resulted from fluvial plucking or eolian deflation [3]. Portions of the landing site are peppered with secondary craters with dark rims of either excavated or partly buried low-albedo material; morphologies of the secondaries indicate a primary or source crater to the south. Crater counts in the landing ellipse indicate a Late Hesperian age [4] with a -2 power-law size-frequency distribution and 2419 craters >=1 km diameter, 445 craters >=2 km diameter, 64 craters >=5 km diameter, all normalized per million km^2. The following numbers of craters lie within the 15,700 km^2 landing ellipse: 49 <0.5 km in diameter, 38 between 2 and 3 km in diameter, and part of one 10-km-diameter crater. There is about a 3% chance of landing within a crater at this site. Within the landing ellipse there are approximately 275 small hills that range from 60 m (which is close to the resolution limit) to 7 km in diameter, although most are <1 km in diameter. Some streamlined islands have concentrations of hills on their downstream sides, which suggests that they were carried as bedload during flooding and deposited where flow velocities decreased; flow reconstruction calculations, however, suggest a maximum boulder of only 10 m diameter could have been carried [2]. Cumulative frequencies of hills within the landing ellipse roughly fit a -1.6 power-law size-frequency distribution between diameters of 0.2 and 1 km. There are 828 hills >=1 km diameter, 382 hills >=2 km diameter, and 64 hills >=5 km diameter (all normalized per million km^2). The actual number of hills in the landing ellipse >=0.25 km diameter is 168, with 62 hills >=0.5 km diameter, 13 hills >=1 km diameter, 6 hills >=2 km diameter, and 1 hill >=5 km diameter. Photoclinometry (symmetric method) was used to estimate the heights and slopes of 12 hills. Results indicate most hills have overall slopes of about 10 degrees, with maximum local slopes up to 25 degrees. A general relationship exists between hill diameter and height h = -33.2 + 0.15d, where h is height and d is diameter (r = 0.86). There is about a 1% chance of landing on a hill at this site. Albedo, color, thermal inertia, and rock abundance suggest that the Ares Vallis landing site shares many of the same characteristics as the Viking landing sites. A combined albedo, thermal inertia (reported in 10^-3 cgs units), and rock abundance dataset kindly provided by P. Christensen [5] was used in our evaluation of the landing site. Albedo varies from 0.19 to 0.23 (1 degree bin data [6]), thermal inertia varies from 9.8 to 12.9 (0.5 degree bin data [7]), fine-component thermal inertia varies from 8.4 to 10.9 (1 degree bin data [8]), and rock abundance varies from 18% to 25% (1 degree bin data [9]) over the Pathfinder landing ellipse. For comparison, the Viking 1 and 2 landing sites have values of rock abundances of 16% and 23%, albedos of 0.25, thermal inertias of 8.4 and 7.8, and fine-component thermal inertias of 7.1 and 6.2 respectively. Red (0.155-0.187) and violet (0.058-0.079) radiances of the Ares landing site using Viking orbiter frames in the 344S series yield an average red-to-violet ratio of 2.3 (range: 2.05- 2.95); Viking lander 1 red (0.165-0.184) and violet (0.062-0.071) radiances yield an average red-to-violet ratio of 2.6 (range: 2.4-2.8), derived from the same orbiter images. The lower albedo, lower red-to- violet ratios, greater rock abundances, and higher thermal inertias for the Ares Vallis site suggest a slightly rockier and less dusty surface than the Viking landing sites. Surface materials at the Ares Vallis site should be similar to those at the Viking landing sites [10]. Color radiances and their ratios suggest a variety of materials that include cohesive soillike materials, dust, and coated and uncoated rocks [11]. Rather high fine-component thermal inertias also suggest that cohesive soillike materials, compatible with successful landing and roving, dominate the surface, with less low- cohesion, low-strength drift material prevalent at the Viking lander 1 site. Rock abundance estimates suggest ample rocks are present and available for analysis with the imaging system and the rover-mounted alpha proton X-ray spectrometer. In contrast with the weak radar echoes received during 3.5-cm wavelength observations for the Viking mission in 1976 [12], strong radar echoes were received from the site during 3.5-cm wavelength Goldstone observations in early 1995. Preliminary analysis of selected delay-doppler echoes with subradar points between 19.7 degrees and 20.2 degrees N and 31.1 degrees and 34.6 degrees W yields RMS slopes of 4.6 degrees +/- 0.7 degrees and normal reflectivities of 0.057 +/- 0.008. Continuous-wave (CW) echoes with subradar points along 18.7 degrees N between 31.9 degrees and 34.6 degrees W yield an RMS slope of 6.4 degrees +/- 0.7 degrees and the following cross sections: total polarized - 0.101, quasispecular - 0.048, polarized diffuse - 0.053, and depolarized - 0.020. A conservative estimate for the experimental uncertainty in these CW cross sections is ~25%. Preliminary comparisons with CW observations at the same wavelength in the southern hemisphere imply that the surface of the Ares site may be rougher at a scale of 0.4-10 m (RMS slope 6.4 degrees +/- 0.7 degrees) and has reflectivities (about 0.058 when integrated from 0 degrees to 30 degrees [13]) that are comparable to or smaller than averages in the south. At these southern latitudes [14], average RMS slopes are 4.04 degrees +/- 1.47 degrees (1988 data) and 4.25 degrees +/- 0.71 degrees (1990 data) with average reflectivities of 0.0603 +/- 0.0296 (1988 data) and 0.1062 +/- 0.0175 (1990 data). Diffuse echo strengths and their ratio are more or less normal at the Ares site; spectra show little or no unusual structure. Delay-doppler reflectivities (0.057 +/- 0.008) are consistent with dry soillike materials with poorly constrained bulk densities of 1.3 +/- 0.3 g/cm^3 [15]. Bulk densities such as these indicate a surface that is consistent with our interpretations of the remote sensing data described above--namely a surface compatible with successful landing and roving, with less of the drift material present at the Viking lander 1 site. The average radar reflectivities also suggest a surface that will adequately reflect radar altimeter transmissions during descent of the Pathfinder spacecraft used for firing the solid rockets and inflating the airbags. The elevation of the landing site appears to be well below the 0 km elevation required to provide sufficient atmosphere for the flight system parachute. The USGS topographic map [16] lists the Ares site at about -2.0 km elevation (relative to the reference surface), which is also the elevation listed for the Viking 1 site. Earth-based radar tracks at 22.71 degrees N obtained on January 20, 1980, and 22.89 degrees N latitude obtained on January 15, 1980, cross the longitudes of both the Ares landing site and the Viking Lander 1 site (22 degrees N, 46.5 degrees W), yielding elevations of -1.7 km and -1.8 km vs. -1.6 and -1.8 km respectively. Other radar tracks at latitudes of 21.59 degrees N obtained on February 16, 1980, and 21.3 degrees N obtained on February 22, 1980, suggest an elevation of -2.0 km for the Ares landing site. Because of uncertainties in relating an elevation, with respect to some reference martian figure, to atmospheric pressure (which is what is important for operation of the parachute), we have simply assumed that the elevation of the Ares site is the same as Viking Lander 1 and determined the atmospheric surface pressure from the extremely repeatable surface pressure measurements for the appropriate landing season and day (6.85 mbar). References: [1] Golombek, this volume. [2] Komatsu et al. (1995) LSC XXVI, 779; LSC XXVI, 781. [3] Greeley et al. (1977) JGR, 82, 4093. [4] Tanaka (1986) Proc. LPSC, in JGR 81, E139. [5] Christensen and Edgett (1994) LPI Tech Rpt. 94-04, 19. [6] Pleskot and Minor (1981) Icarus, 45, 447. [7] Christensen and Malin (1993) LPS XXIV, 285. [8] Christensen and Moore (1992) in Mars, 687-729, Univ. of Arizona, Tucson. [9] Christensen (1986) Icarus, 68, 217. [10] Moore et al. (1987) USGS Prof. Paper, 1389, 222. [11] Arvidson et al. (1989) JGR, 94, 1573; Dale-Bannister (1988) LPS XIX, 239. [12] Tyler et al. (1976) Science, 193, 812; Downs G. S. et al. (1978) Icarus, 33, 441-453. [13] Moore and Thompson (1991) Proc. LPS, Vol. 21, 457. [14] Thompson et al. (1992) LPS XXIII, 1431-1432. [15] Moore and Jakosky (1989) Icarus, 81, 164-184. [16] USGS, 1989, Misc. Inv. Map I-2030. Edgett K. S. Viking IRTM Observations of the Anticipated Mars Pathfinder Landing Site at Ares Vallis Albedo, rock abundance, and thermal inertia derived from Viking Infrared Thermal Mapper (IRTM) observations provide some insight as to the nature of the surface materials that occur at the Ares Vallis landing site [1,2]. Observations of the Ares Vallis site are compared here with the Viking 1 and 2 lander sites, which offer some measure of "ground truth." Moderate resolution (30-60-km-sized areas) observations are the best available for the Ares Vallis site, and include an IRTM-derived albedo map (1 degree latitude/longitude bins) compiled by L. K. Pleskot [3], thermal inertia (0.5 degrees bins) from Christensen and Malin [4,5], and rock abundance and fine component thermal inertia maps (1 degree bins) derived from IRTM data by Christensen [6]. In the landing ellipse, albedo ranges from about 0.19 to 0.23 and is generally lower at the eastern end of the landing ellipse, higher toward the west. The comparable albedo at the Viking lander sites is about 0.25 at both. Thermal inertia, computed using the Kieffer thermal model [7] with the "2% assumption" (wherein the atmospheric contribution to downgoing radiation is 2% of the maximum solar insolation) varies between 410 (9.8) and 540 (12.9) J m^(-2) s^(-0.5) K^(-1) (10^(-3) cal cm^(-2) s^(- 0.5) K^(-1)). (Note that both S.I. units and commonly-used "Kieffer" units for thermal inertia are given here; the latter is in parentheses). In general, thermal inertia is higher at the eastern end of the landing ellipse and lower toward the west. For comparison, the Kieffer model thermal inertia of the Viking 1 site is about 360 (8.5) and for Viking 2 is about 330 (7.9). Rock abundance, a parameter derived from thermal inertia and differences in temperature at 7, 9, 11, and 20 micrometers, ranges from about 25% at the east end of the ellipse down to about 18% at the west end. The uncertainty here is on the order of 5% to 10% rocks [6]. Modeled rock abundance at the Viking 1 site is about 15 +/- 5%; at Viking 2 it is 20 +/- 10% [8]. The rock abundances estimated for the Ares Vallis landing ellipse are similar to the rock abundances of the two Viking lander sites. The corresponding fine component (Kieffer model) thermal inertia, a byproduct of the rock abundance modeling, is about 350 (8.4) to 460 (10.9) from west to east in the Ares landing ellipse, as opposed to 300 (7.1) and 260 (6.2) for the Viking 1 and 2 sites, respectively. A search was conducted for high-resolution (2-5-km-sized areas) IRTM observations of areas within the Ares Vallis landing ellipse. One Viking orbiter track meeting appropriate search criteria (spacecraft range <= 2500 km, emission angle <= 60 degrees, L(sub)s = 350 degrees-115 degrees, Hour 0-6) was found to pass about 200 km to the north of the landing ellipse; the thermal inertias there were consistent with the moderate-resolution results, but allowed a more detailed map along the orbiter's flight path. Unfortunately, no such IRTM data were found to pass through the Ares landing ellipse. One daytime (10-14 H) high- resolution IRTM track from Viking 2 was found, but accurate computation of thermal inertia is problematical for daytime data. In general, this track indicates thermal inertias similar to those in moderate resolution. Lately there has been considerable discussion about the uncertainty in thermal inertia derived under the relatively dusty atmosphere of the Viking era [9-13]. Hayashi et al. [10], using the Haberle-Jakosky coupled surface-atmosphere model approach [9], note that the Kieffer model thermal inertias for the Viking 1 and Viking 2 sites are about 60- 80 (1.4-1.9) thermal inertia units too high. The corresponding range of thermal inertias for the Ares Vallis site would also drop by about 100 (2.4) units, thus ranging from about 310 (7.4) in the west to 440 (10.5) in the east. The general trend from higher thermal inertia in the east to lower in the west remains unchanged. From the conclusions of Hayashi et al. [10], it seems that the fine-component thermal inertias (by- product of rock abundance) for the Ares landing ellipse would drop by about 90 (2.2) units, to range from 260 (6.2) in the west to 370 (8.8) in the east. The Ares Vallis landing site has the potential for being somewhat different than the two Viking lander sites. However, the differences might turn out to be as subtle as the differences observed when one compares the two Viking sites. The Ares site in general is not radically different from the previous sites; this may turn out to be helpful for reinterpretation of the geology of the Viking sites. In general, the Ares site is about as rocky as the two Viking sites, but the somewhat lower albedo and higher fine-component thermal inertias suggests there might be more sand (or at least, less dust) at the Ares site. The fine- component thermal inertias suggest effective particle sizes in the medium to medium-coarse sand range (300-600 micrometers) throughout the region; this assessment is consistent with new thermal conductivity results from Presley [14]. The east-to-west variation in thermophysical properties might indicate that there are coarser deposits of sediment at the eastern end of the landing ellipse. It seems likely that the landing site will look less like the Viking 1 and 2 sites if Mars Pathfinder touches down at the eastern end of its landing ellipse. The implications for eolian features that might occur at the landing site are discussed in a subsequent abstract [15]. References: [1] Golombek M. P. et al. (1995) LPS XXVI, 481-482. [2] Edgett K. S. (1995) LPS XXVI, 353-354. [3] Christensen P. R. (1988) JGR, 93, 7611-7624. [4] Christensen P. R. and Malin M. C. (1988) LPS XIX, 180-181. [5] Christensen P. R. and Malin M. C. (1993) LPS XXIV, 285-286. [6] Christensen P. R. (1986) Icarus, 68, 217-238. [7] Kieffer H. H. et al. (1977) JGR, 82, 4249-4291. [8] Christensen P. R. (1982) JGR, 87, 9985-9998. [9] Haberle R. M. and Jakosky B. M. (1991) Icarus, 90, 187- 204. [10] Hayashi J. N. et al. (1995) JGR, 100, 5277-5284. [11] Paige D. A. et al. (1994) JGR, 99, 25959-25991. [12] Edgett K. S. and Christensen P. R. (1994) JGR, 99, 1997-2018. [13] Bridges N. T. (1994) GRL, 21, 785- 788. [14] Presley M. A. (1995) Ph.D. dissertation, Arizona State Univ., Tempe, 237 pp. [15] Edgett K. S., this volume. Kochemasov G. G. Possibility of Highly Contrasting Rock Types at Martian Highland/Lowland Contact "Stream sediment sampling" was proposed in 1979 as a rational tool for collecting and studying various rock fragments on the martian surface (9th Gagarin reading on aeronautics and aviation) and was again discussed in 1988 (18th Gagarin reading) [1]. This idea was based on the author's experience in stream sediment, heavy fraction, and rock fragment sampling as a geological prospecting tool in various African and Asian environments. A particular parallel was drawn between the martian environment and that of mountain deserts of northern Africa (Anti-Atlas) where eolian contamination is rather pronounced and which has to be borne in mind during the martian rock-sampling mission. Experiments in the Anti-Atlas have shown that significant eolian contamination exists in fine (<0.5 mm) dry mountain alluvial fractions. Hence, relatively large rock and mineral fragments are more safe for "on-the-spot" study of a catchment area and preparing a return collection. The majority of planetologists believe, based on remote spectral studies, that the difference between low and highland rock types is not very great (as the difference between the fresh and weathered basalts - palagonites [2]). We think that this conclusion is controversial, considering, for example, an enormous albedo difference between Arabia Terra and Syrtis Major Planitia, not occasionally likely the highest difference in the inner solar system. It was recently shown that there is regular change of crucial planetary crust characteristics with increasing solar distance ([3] and ref. herein). This resonance behavior is related to wave tectonics that considers the interference of lithospheric (geospheric) stationary waves warping planetary spheres in four directions (orthogonal and diagonal) and having lengths proportional to the planetary orbital periods (Mercury pi R/16, Venus pi R/6, Earth pi R/6, Earth pi R/4, Mars pi R/2, where R is a planet's radius). This "intricate weaving" produces "rounded" tectonic blocks and surface relief, both increasing with the solar distance. Subsided (oceanic) and uplifted (continental) segments of the planetary crusts, composed of relatively dense and light materials (principle of block angular momentum conservation govern this behavior) tend to have density contrasts growing in the same direction (Fig. 1). Tracing the chemistry change of basaltic plains is most reliable as their soils were studied directly on Earth, Venus, and Mars. Iron content, and hence density of mare basalts, correlates with the planetary relief amplitude or the amplitude of producing its lithospheric wave (e.g., the deeper primary Pacific Ocean has more Fe-rich tholeiites than the shallower secondary Indian and Atlantic Oceans, which helps us to understand the governing principle). Basalt contents of the Earth's primary Pacific depression are (in weight %) Fe/Si 0.38 and Fe/Mg 1.89 [4]; martian basalts, respectively, 0.64 and 2.53 [4]; venusian ones 0.31 and 1.10 [5]; and mercurian ones 0.16 and 0.32. For Mercury we took into account a mean estimate of Fe content in the surface rocks (5%), the high Mg content of its mantle, and the closeness of its crust composition to anorthosites with small albedo contrast between "mare" and "highlands" [4]. Figure 1 shows ratios of the above Fe parameters compared to the terrestrial ones taken as 1 (solid line - relief, dash line - Fe/Si, dots - Fe/Mg). Highland compositions: andesitic terrestrial [6] changes to alkali basaltic venusian [5]. The composition of the highlands regions of Mercury (bright cratered plains) is taken to be somewhat more anorthitic or less dense (enstatite anorthosite) than that of the "smooth dark plains." Decreasing highland densities with increasing solar distance predicts the "lightest" continents on Mars, which is supported by the very sharp albedo contrast between "cratered old terrains" (bright areas) and plains (dark areas), indications of viscous magmas, and low- density rocks (gravity data [7]). "White rock" [8] with its very high (resembling ice) albedo could be albitite--a light acid variety of plagioclasite. High SiO2 in this rock and in the bulk highland rock, equal to the SiO2 content of albite (70%), follows from 60% SiO2 in the bulk crust (= martian dust enriched in feldspar [4]) and 45% SiO2 in the lowland basalts covering 1/3-2/5 of Mars' surface. Albitite could be magmatic or metasomatic in origin. The formation of this acid Na-rich plagioclasite is consistent with high pressure caused by warping Mars lithospheric waves ( high and anisotropic pressure squeezing the planet). As the lowland rocks recede from the Sun they become more Fe rich and dense (anorthite enstatitite 2.93 g/cm^3; Mg basalt 2.95; tholeiite 3.0; Fe basalt 3.1); the highland rocks, inversely, become less dense (enstatite anorthosite 2.90 g/cm^3; alkali basalt 2.85; andesite 2.75; albitite 2.65). The density contrast between the highland and lowland rocks increases: 0.03; 0.1; 0.25; 0.45 g/cm^3 (Fig. 1, dot-dash line, reduced to the terrestrial contrast 0.25 g/cm^3), correlating with the relief range. Such regularity is caused by the action of Le Chatelier rule, according to which equilibrium disturbance brings forces creating obstacles to it: increasing surface warping (relief range) brings increasing density contrast between lowland and highland rocks. This tends to level angular momenta of rising and falling blocks. We suggest that at the low/highland contact in the Ares Vallis outflow area the Pathfinder could encounter rocks of the Fe-tholeiite family mixed with light (not dense) rocks rich in Na such as albitites and syenites (some resemblance with mangerite-anorthosite and anorthosite- granite formations of the Earth). References: [1] Kochemasov G. G. (1989) Gagarin Reading on Aeronautics and Aviation, 1988, Moscow, Nauka, 275 (in Russian). [2] Kieffer H. et al (1992) Mars, Univ. of Arizona, Tucson. [3] Kochemasov G. G. (1994) 20th Russian-American Microsymposium on Planetology, Abstracts, Moscow, Vernadsky Inst., 46-47. [4] Basaltic Volcanism on the Terrestrial Planets (1981) Pergamon, 1286 pp. [5] Basilevsky A. T. et al. (1992) JGR, 97, 16315-16335. [6] Taylor S. R. and McLennan S.M. (1985) The Continental Crust, Blackwell, Oxford, 312 pp. [7] Phillips R. J. et al (1973) JGR, 78, 4815-4820. [8] Ward A. W. (1979) JGR, 84, 8147-8166. Figure 1 appears here in the hard copy. Friday, September 29, 1995 WHAT WILL PATHFINDER FIND? 8:30 AM Tanaka K. L.* Potential Source Rocks of Sedimentary Deposits at the Pathfinder Landing Site No abstract available. Rice Jr. J. W.* Edgett K. S.* A Sojourner's Prospectus: Provenance of Flood-Transported Clasts at the Mars Pathfinder Landing Site Introduction: Mars Pathfinder and its microrover, recently named Sojourner (after the 19th Century civil rights figure, Sojourner Truth), will examine rock and soil composition at the Ares Vallis landing site, centered at 19.5 degrees N, 32.8 degrees W. The Ares site was chosen because it may provide a "grab bag" of flood-transported rock types [1]. This paper discusses terrestrial examples of deposits and processes in catastrophic flood and periglacial settings on Earth, and how these lend credence to the conclusion that the Ares site might provide a rich sampling of diverse materials. The purpose is to generate discussion on the nature of the Ares "grab bag." What Rocks Might Have Been Sampled?: The units cut by Ares Vallis, according to Scott and Tanaka [2] are "Hesperian ridged plains material" and two types of "Noachian plateau cratered material." Tiu Vallis cuts the same units. An older, Noachian proto-Ares Vallis cuts similar units and may have stretched from the Nereidum Montes through Uzboi, Ladon, and Margaritifer Valles to Chryse [3]. The Hesperian ridged plains are likely dominated by mafic or ultramafic lavas, perhaps interbedded with sediments and volcanic ash. The Noachian units likely include impact breccias interbedded with lavas, pyroclasts, and eolian and lacustrine sediment. Hydrothermal minerals/rocks might be among the materials at the Ares site [4], as well as "hardpan" or other low-temperature diagenetic debris [5]. The only clue we have to the composition of Noachian rocks is the ALH 84001 orthopyroxenite meteorite, which might be a ~4 Ga sample from the martian highlands [e.g., 6]. Missoula Flood Deposits in Washington and Oregon: The Missoula flood deposits suggest how a diversity of lithologies can be deposited in a location like the Ares site. The late-Wisconsin sediments in the Quincy Basin, Washington, came mainly from floods that poured through the Grand Coulee (see field guide in this volume, and references therein). Most of the larger clasts (cobbles, boulders) were deposited in the Ephrata Fan. Observations made in June 1995 (with M. Golombek [7]) show that about 95% of the rocks on the surface of the fan are basalts, ~5% are granodiorites and <<1% are others (metamorphic and sedimentary rocks [8]). The two dominant rock types outcrop in the Grand Coulee [9]. The sand dunes south of the Ephrata Fan have a wider variety of flood- deposited clasts: basalt (~55%), metamorphic rock (~6%) quartz (~30%), and other minerals [10] derived from rocks that outcrop as far away as Idaho. At the Ares landing site, we expect that the smaller clasts are most likely to include rock fragments that were carried the greatest distances by water. There is another way for a flood to carry large clasts great distances. The Missoula flood deposits include ice-rafted debris. The most spectacular ice-rafted features occur in the Willamette Valley, Oregon [11]. The boulders in Willamette Valley include metamorphic, plutonic, and volcanic rocks transported from outcrops occurring from Lake Pend Oreille, Idaho, down to Portland, Oregon [11]. Amazingly, one ice-rafted boulder was the Willamette Meteorite, a 14,000 kg Ni-Fe that was originally incorporated into the Cordilleran ice sheet, transported down the Purcell Trench glacial lobe, then brought to Oregon by a flood-borne ice raft [12]. We think that some of the large (~1-km) knobs (which lack streamlining) at the Ares landing site could be ice-rafted blocks. High-Latitude Fluvial Flood Processes on Earth: Rivers in high northern latitudes provide additional insight as to mechanisms that transport and deposit clasts from a variety of settings. In addition to ice rafting, whole, thick sections of frozen rock and soil may be transported by water following bank collapse. Bank collapse can occur all along the course of a fluvial system, such as the Colville River in Alaska, and is usually initiated by thermo-erosional undercutting of an ice-rich bank by the warmer water in the stream [13,14]. Some high-latitude fluvial systems discharge into ice-covered seas; the ice cover allows sediment to be transported great distances (tens of kilometers on Earth). This process is well documented at the Colville River delta, where springtime flood waters meet the still-frozen Beaufort Sea [15]. Although controversial, this process could have occurred at the Ares site if an ice-covered sea was ever present. Diversity of Rock Types at Viking 1 Lander Site?: The Ares landing site is thought to be a "grab bag," with rocks and perhaps intact soil units that have been transported by floods. The Ares site may offer new insight into the Viking 1 lander site, the geologic history of which remains controversial. Here we note that the rocks at the Viking 1 lander site appear to have a variety of lithologies [e.g., 16]. Perhaps what will be found at the Ares Vallis site will help us reinterpret what was observed at the Viking 1 site two decades ago. References: [1] Golombek M. P. et al., this volume. [2] Scott D. H. and K. L. Tanaka (1986) USGS Map I-1802-A. [3] Parker T. J. (1989) LPS XX, 826-827. [4] Farmer J. D. et al. (1995) LPS XXVI, 393-394. [5] Treiman A. H., this volume. [6] McSween H. Y. Jr. (1994) Meteoritics, 29, 757- 779. [7] Golombek M. and D. Rapp, this volume. [8] Grolier M. J. and J. W. Bingham (1971) USGS Map I-589. [9] Stoffel K. L. et al. (1991) Wash. Div. Geol. Earth Res. Map GM-39. [10] Petrone A. (1970) M.S. thesis, Wash. State Univ., Pullman. [11] Allison I. S. (1935) GSA Bull., 46, 605-632. [12] Allen J. E. and Burns M. (1986) Cataclysms on the Columbia, Timber Press, Portland, Oregon. [13] Williams P. J. and Smith M. W. (1989) The Frozen Earth, Cambridge Univ., New York. [14] Walker H. J. (1973) Artic Inst. North America Tech. Paper, 25, 49-92. [15] Reimnitz E. and Bruder K. F. (1972) GSA Bull., 83, 861-866. [16] Sharp R. P. and Malin M. C. (1984) GSA Bull., 95, 1398-1412. Treiman A. H.* Hardpan and Other Diagenetic 'Rock' in the Catchment of Ares Valles and Surrounding Areas The Mars Pathfinder landing site, in the outflow plains of Ares Vallis, will presumably present a "grab-bag" of many rock types transported from upstream. Based on the geology of the Valles Marineris (just southwest of Ares Vallis), the Ares Vallis catchment could include hardpan and other diagenetic "rocks" lithified by low-temperature aqueous solutions. The chemical effects of diagenesis may be detectable with the M/P APX instrument; if present, but unrecognized, diagenetic effects could lead to erroneous geochemical inferences. Extensive Hardpan: The "cliffs" at the wall tops in Valles Marineris represent an extensive hardpan horizon [1] (Fig. 1). The cliffs represent a relatively resistant packet of layers: an upper darker layer (~50 m thick); a central brighter layer (~250 m thick); and a lower dark layer (~100 m thick). The layer packet is exposed on chasma walls from western Noctis Labyrinthus (110 degrees W, VO 423A75) to eastern Eos Chasma (31 degrees W, VO 964A22), and in continuous reaches of 200-300 km of wall in Hebes, Ius, Eos, and Coprates chasmae. The packet is visible on all planar chasma walls (fault or landslide) where Viking Orbiter or Mariner 9 image resolutions are <150 m, and rarely on spur- and-gully walls or in images of poorer resolution. The packet appears at the same stratigraphic position (<50 m from the wall tops) and thickness (400+/-200 m) over this huge area, independent of elevation (3-10 km), age of adjoining plateau surface (Noachian through upper Hesperian), and geological structures (e.g., impact craters, Fig. 2). From its continuity and consistent stratigraphy, the layer packet appears to be a single geologic feature, exposed discontinuously over the whole Valles Marineris (4000 km E-W, 800 km N-S). Age relations between geology of the plateaus and the layering packet suggest that it did not form by deposition. If the layers were deposited in Noachian times, they should dip beneath the resurfacing materials of the Hesperian age ridged plains of Lunae and Syria Planae. However, the layers remain at the wall tops. If the layers were deposited in Hesperian times, they should not be present under Noachian age plateaus; they are present. Nor can a depositional origin explain the continuity of the packet beneath impact craters (Fig. 2). However, a diagenetic origin for the layering, alteration of pre- existing material in situ, can explain its structure transgressions, stratigraphic transgression, and lateral extent. Diagenetic effects controlled by proximity to the ground surface would appear independent of surface elevation, and transgress pre-existing features, including the boundary between Noachian and Hesperian deposits and the structures and lithologic changes beneath and inside impact craters [2]. Given similar compositions and permeabilities to fluids, diagenetic transformations should affect all earlier materials, regardless of their ages. And diagenetic layers can be of regional extent, like hardpans or calcretes on Earth. Diagenesis has been invoked to explain rock shapes at Viking lander sites and clodding of near-surface sediment [3,4], but not for the extent or thickness of the upper-wall layers in Valles Marineris. Similar upper-wall layering is visible in the Ares Vallis catchment, although neither exposures nor image resolutions are as good as for Valles Marineris. The layering packet appears present at Hydaspis Chaos (1.6 degrees N, 29.6 degrees W; VO 083A34) and Iani Chaos (2.6 degrees S, 48.59 degrees W; VO 405B09), and in some crater walls (6.5 degrees N, 19.9 degrees W; VO 745A31). Hardpan material from these layers could have been transported to the M/P landing site. Hardened Fault Zones: Some graben-bounding faults on the plateaus adjacent to Valles Marineris appear to continue as ridges on the walls of chasmae. These ridges are most prominent at the boundary between Ophir and western Candor Chasmae (Candor Labes, 4.5 degrees S, 72 degrees W), and are now under study. The ridges are probably relict fault zones, cemented and hardened by precipitates from aqueous solutions (e.g., [5]). Conclusion: Rocks and regolith in the Ares Vallis catchment may have been chemically and structurally affected by aqueous solutions under diagenetic conditions. On Earth, hardpans can be cemented by Ca carbonate, Ca sulfate, salts, iron oxide hydroxides, clays, and silica. The chemical signatures of the first three cements may be directly detectable with the M/P APX instrument (as C, S, and Cl). Chemical signatures of the last three may not be detectable directly, and could lead to geochemical errors. For instance, an ironstone hardpan developed in basalt could masquerade as a ferroan ultramafic lava. References: [1] Treiman A. H. et al. (1995?) JGR. [2] Melosh (1989) Impact Cratering. [3] Fuller and Hargraves (1978) Icarus, 34, 614. [4] Jakosky and Christensen (1986) JGR, 91, 3457. [5] Mozely and Goodwin (1995) Geology, 23, 539. Figure 1, which appears in the hard copy, shows wall-top layering, south-east wall of Hebes chasma (arrow) near top of landslide scarp (~2 km tall, ~30 degrees slope angle). VO 915A06 (2.2 degrees S, 73.5 degrees W), scale bar 10 km. Figure 2, which appears in the hard copy, shows the south wall of Gangis Chasma, VO 014A30 (centered 9.7 degrees S, 44.6 degrees W), oblique view south, scale bar 10 km E-W at wall. Layering (arrows) undeflected from plains to crater floor under crater rim. An extra layer to west thickens and deepens beneath crater. Mellon M. T.* Ground Ice at the Mars Pathfinder Landing Site Water on Mars is of central interest to many areas of martian research, ranging from climate studies to geologic evolution. Ice in the martian subsurface (ground ice) is proving to be an important and dynamic reservoir. Its presence during recent geologic history may have had some influence on the geomorphic character of the Pathfinder landing site. The presence and implications of ground ice will be discussed. At present epoch ground ice is not likely to be stable at the low latitude of the intended Pathfinder landing site. In addition to low latitude, the relatively high average thermal inertia and low albedo typical of this region produce surface and subsurface temperatures much in excess of the atmospheric frost point. (The frost point is the temperature at which atmospheric water vapor would condense and, below which subsurface ice is stable with respect to sublimation and loss to the atmosphere). With the exception of higher long-term atmospheric water abundances than have been observed (raising the frost point) and small localized regions of unusually low thermal inertia (lowering mean ground temperatures), ground ice present in the near-surface regolith would rapidly sublime, diffuse to the surface, and be lost to the atmosphere. In the recent geologic past, however, the situation was probably very different. Large oscillations in the martian orbit would have caused considerable changes to the pattern of insolation and to the martian climate as a whole. Primarily, an increase in the obliquity (tilt of the spin axis relative to the orbital plane) would increase the amount of solar energy being deposited in the polar regions, while simultaneously decreasing that in the equatorial and midlatitude regions. An increase in polar insolation would have increased the rate at which water sublimes from the polar caps during the summer season and generally increased the atmospheric water content and the frost point temperature. Similarly a decrease in equatorial solar heating would have lowered the regolith surface and subsurface temperatures, allowing ground ice to become stable (and rapidly condense from atmospheric vapor) in regions in which it was previously unstable. This would have been the case for the intended Pathfinder landing site. Such episodes of stability have occurred periodically throughout history when the obliquity reaches moderately high values, the last episode occurring mearly 500,000 yr ago. On Earth ground ice present in regions of permafrost influences the geomorphic character of the surface, producing such periglacial landforms as ice-wedge polygons, thermokarst pits, and solifluction lobes. The episodic presence of ground ice at the Mars Pathfinder landing site may have had a similar impact. It is possible that the formation of thermal-tension fractures in the ice-cemented martian permafrost may have produced polygonal terrain on a scale similar to terrestrial counterparts. Polygonal terrain has been suggested in the Viking lander 2 images and may be observed in Pathfinder images. Although such polygons will not be active at the Pathfinder landing site in the present absence of ground ice, abundant ground ice at high obliquity would facilitate their development, which may in turn leave a geomorphic signature until the present epoch. In addition, periodic inflation and deflation of the ground due to "frost heave," which may be associated with the episodic condensation and sublimation of ground ice, might have produced thermokarst topography (differential collapse related to differences in regolith structure and composition). The Pathfinder landing site may exhibit a geologic character, in part due to the influences of ground ice. Certainly, other processes are in force, such as eolian and impact. These processes will compete with periglacial processes, particularly since the last occurrence of ground ice, making periglacial landforms difficult to identify. Despite this difficulty, the impact of ground ice may be evident at the lander site. Moore H. J.* What Will Pathfinder See and Do on Mars? Experiences of the Viking landers are excellent guides for Pathfinder. Lander 1 operated for nearly four martian years after landing on July 20, 1976. Like the Viking landers, Pathfinder will observe materials, sample them, deform and disturb them, estimate their physical properties, witness their response to martian winds, and measure things. Major differences between the Viking landers and Pathfinder include the cameras and analyses sampling techniques. Cameras on both spacecraft have similar resolutions and stereometric capabilities, but spectral capabilities are better for Pathfinder. Samplers on the landers acquired and delivered samples to analytical instruments aboard the landers, but the Pathfinder microrover will carry an analytical instrument to soillike materials and rocks. Things that Pathfinder might see and do are illustrated below with Viking lander observations and results. Bright drifts, deflated by wind, are striking features seen in Lander 1 images. These drifts, with cross-laminations that dip north-eastward [1], are superposed on a rocky substrate. In the distance, a couple of undeflated duneforms rest on the rocky substrate. The substrate includes a cohesive soillike material atop and admixed with rocks and rock fragments ejected from nearby impact craters; there may be outcrops of bedrock [2]. Rocks are striking features seen in the Lander 2 images [3]. Among the rocks are crusty to cloddy soillike materials and thin deposits that appear to be fine grained--like drift material. Estimates of the mechanical properties of soillike materials were derived from the responses of the materials to the footpads and sampler [4].In order of increasing strength, the sampled materials are (1) drift (Lander 1), (2) crusty to cloddy (Lander 2), and (3) blocky (Lander 1). Drift material is fine grained with grain sizes near 0.1 to 10 micrometers [5,6]. Friction angles estimated from deformations in front of the sampler during trenching are near 18 degrees and imply that drift material is porous with a low bulk density; its bulk density in the X- ray fluorescence spectrometer (XRFS) analysis chamber is 1100 +/- 150 kg/m^3 [7]. Cohesions range from 0 to 3.7 kPa. Disruption of crusty to cloddy materials produces crusts and fines in places and more or less equidimensional prismatic clods about 0.04 m across and less in other places. Local surfaces are littered with small clodlets (a few millimeters across), but surfaces of crusts are also seen at the surface. Individual mineral grains are in the 0.1 to 10 micrometer range [6,8]. The friction angle is about 34 degrees and compatible with moderately dense soillike materials. Cohesions range from 0 to 3.2 kPa. Disruption of blocky material produces strong prismatic clods. Smooth, tamped, and compressed surfaces in some sample trenches of blocky material argue for the presence of particles generally finer than 50 micrometers, but large comminutor motor currents suggest that significant amounts of strong millimeter-sized objects (such as rocks) are present. Some small unweathered rocks may be imbedded in blocky material because several centimeter-sized objects deposited on the XRFS funnel are dark gray with color reflectances consistent with mafic igneous rocks [9]. Blocky material is the strongest of the three soillike materials, chiefly because of its cohesion, which is near 5 kPa. Its friction angle is about 31 degrees. Friction angles and cohesions of blocky material are compatible with moderately dense loess on Earth. Chemical compositions of the soillike materials determined by the XRFS are remarkably similar at both sites [10] although the sites are separated by about 6500 km. The mineralogy of the surface materials are unknown, but the surface materials may be weathering or alteration products of mafic igneous rocks or smectite clays. Palagonite was suggested as an analog for the soillike materials because of its spectral properties [11], but palagonite does not reproduce the results of the Viking Labeled Release experiments [11]. The weight percents of SO3 and Cl in soillike materials correlate directly with their visual classification and relative mechanical strengths of the materials. The gas chromatograph mass spectrometer detected water and CO2 in the soillike materials [12]. It appears probable that adsorbed water and CO2 are present in the soillike materials. Little is known about the rocks at both sites because they were never sampled by Viking. Textures and appearances of the rocks vary [13]. Most of the rocks have irregular shapes but a few appear to be rectangular prisms. Surfaces of most rocks are pitted, others are smooth and unpitted, and others are knobby. The surface sampler did not chip, scratch, or spall surfaces of the rocks, so they do not have thick, weak, punky rinds [14]. The rocks are probably like terrestrial rocks with bulk densities near 2600 kg/m^3, cohesions near 10^3-10^4 kPa, and friction angles in the range of 40 degrees-60 degrees. Many surfaces of rocks have red colorations, while others are decidedly less red and nearly gray [15]. Red surfaces are consistent with coatings of a palagonite-like material on unaltered rocks produced by eolian deposition, weathering, or both. Gray surfaces suggest that the underlying rock is mafic. Hopefully, Pathfinder will obtain chemical compositions of both red and dark gray rock surfaces. Modifications of the surface by natural processes were mild. Two small slope failures occurred on steep slopes of drifts at the Lander 1 site. A few to tens of micrometers of bright red dust were deposited from great dust storms in the falls of the first and fourth years and the winter of the first year. Mild winds from local dust storms reworked thin layers of dust during the second and third years [16]. At the Lander 1 site, a local dust storm occurred during the late winter of the first year. Late in the winter of the third year, a local dust storm eroded trenches, conical piles, and other surface materials around Lander 1 [16]. Finally, a great dust storm was in progress in the fall of the fourth and final martian year [17]. Surface phenomena were different at Lander 2 because ices and dusts were deposited during the first two great dust storms [18]. During the second winter, ices and dusts were again deposited on the surface but no great dust storm was in progress. The ices evaporated and a few micrometers of dust remained behind. No strong winds capable of eroding the surface or the conical piles of loose materials placed on and among rocks by the sampler ever occurred while Lander 2 observed Mars. References: [1] Mutch et al. (1976) Science, 194, 87. [2] Binder et al. (1977) JGR, 82, 4439. [3] Mutch et al. (1977) JGR, 82, 4452. [4] Moore et al. (1982) JGR, 87, 10043; (1987) USGS Prof. Paper 1389, 222. [5] Ballou et al. (1978) Nature, 271, 644. [6] Moore and Jakosky (1989) Icarus, 81, 164. [7] Clark et al. (1977) JGR, 82, 4577. [8] Oyama and Berdahl (1977) JGR, 82, 4669. [9] Dale-Bannister et al. (1988) LPSC XIX, 239. [10] Toulmin et al. (1977) JGR, 82, 4625; Clark et al. (1982) JGR, 87, 10059. [11] Banin et al. (1992) in Mars, p. 594. [12] Arvidson et al. (1989) Rev. Geophys. Space Phys., 27, 39. [13] Garvin et al. (1981) Moon Planets, 24, 355; Sharp and Malin (1984) GSA Bull., 95, 1398. [14] Moore et al. (1978) USGS Prof. Paper, 1081, 21. [15] Guinness et al. (1987) Proc. LPSC, in JGR, 92, E575. [16] Arvidson et al. (1983) Science, 222, 463; Moore (1985) Proc. LPSC, in JGR, 90, D163. [17] Tillman (1988) JGR, 93, 9433. [18] Jones et al. (1979) Science, 204, 799; Wall (1981) Icarus, 47, 173; Svitek and Murray (1990) JGR, 95, 1495. Edgett K. S.* After the Flood: A Preview on Eolian Features at the Mars Pathfinder Landing Site in Ares Vallis After the last floods poured through Ares Vallis, wind was the main sedimentary agent at work in the Mars Pathfinder landing site region. Based upon field observations of the Ephrata Fan, a major flood deposit in the Channeled Scabland of Washington, it appears likely that wind should have reworked the fine sediment that was left after the Ares floods, while other sediments, especially airborne dust, were deposited in the area. On the Ephrata Fan, flood-deposited sand was reworked by wind to form dunes [1], and where sand was not present, windborne silt and volcanic ash accumulated between rocks [2]. Relative to the Viking landers, Mars Pathfinder is very well suited to study eolian features. The origin and physical nature of some eolian landforms at the Viking lander sites remains unresolved. For example, the particle size of sediments in the drifts at the Viking 1 site is undetermined [3,4]. Because of its stereo/multispectral imaging [5] and microrover capabilities, Mars Pathfinder might resolve similar puzzles in Ares Vallis. If Viking 1 had had the microrover, it could have driven out to the nearby drifts and determined (1) if the drifts have cross beds, (2) if the drifts are cohesive and/or cemented sand or dust, (3) if features identified as sand ripples [4] were such, and (4) if all the drifts are dark with bright red coatings, as suggested by the presence of one low-albedo drift [4]. If Viking 2 had had a microrover, its tiny cameras could have been used to determine whether ripples in nearby troughs [6] consisted of granules (2-4 mm-sizes) [4]. Finally, if the Vikings had had Mars Pathfinder's multispectral imaging ability [5], some information about the mineralogy of eolian deposits could have been obtained. Albedo, thermal inertia, and rock abundance offer clues to the nature and distribution of eolian debris [7]. In a global context, the Ares landing site has an intermediate albedo and thermal inertia [8-10]; it is not dust covered like Arabia, nor is it sandy like the north polar sand sea. Rock abundance at the Ares site is similar to the Viking 1 site, but the albedo is slightly lower and the thermal inertia is slightly higher [9,10]. Thermal and albedo properties of the landing ellipse change from east to west, with darker surface materials (probably sand) occurring in the northeast near the margin of Acidalia [10]. The low-albedo Acidalia Planitia is thought to be rocky with sand in the form of sand sheets or drifts [7,8]. The dark sands of Acidalia are probably mobile and moving slowly into the northeast end of the Ares Vallis landing ellipse. Mars Pathfinder's wind sock experiment and meteorological station should provide information about eolian events that occur during the mission. The July 1997 landing corresponds to L(sub)s 141 degrees, or mid- Northern Summer. Northern Summer should be the least windy season, according to GCM work [11]. Little eolian activity is to be expected during the 30-day primary mission; the strongest winds (needed to move sediment) typically occur during the seasons that have the strongest annual winds (late Northern Autumn through Winter) [see 12]. No eolian dunes will be found at the Ares landing site; none are observed in Viking orbiter images. It is possible, however, that dark sand might be accumulated in eolian drifts and/or granule ripples, particularly in the eastern half of the landing ellipse. The two fields of small "crater clusters" [13] within the landing ellipse have dark material surrounding them. This dark material must be sandy to have maintained a low albedo over time, and might be (1) sand-sized impact glass [14] or (2) an indicator that the craters penetrated to a lens of flood-deposited sand. Like on the Ephrata Fan of Washington, the Ares site is probably located in a rocky or gravelly facies [15]. The site has probably accumulated some airborne dust, forming a discontinuous mantle between rocks like at the Viking 1 site. The light-toned "etched terrain" [16] just outside the southwest end of the landing ellipse might be eolian-scoured, similar to features Sharp [17] described elsewhere on Mars. If so, then the southwestern part of the landing ellipse might have small eolian deflation pits and/or remnant knobs and mesas. Other features, like ventifacts or pitted rocks, might be found, but their presence is difficult to predict. The Mars Pathfinder landing site in Ares Vallis is not likely to be a vigorously active eolian environment, but may be more active than the Viking lander sites. The Mars Pathfinder site is probably most similar to the Viking 1 site, although with somewhat more windblown sand. The new capabilities of Mars Pathfinder will allow more detailed investigation of eolian features, giving new clues about the particle sizes and compositions of eolian sand and dust. In turn, Mars Pathfinder offers a chance to re-interpret the geology of the Viking lander sites, provided that there are features similar to those seen by Viking 1 or 2. References: [1] Petrone A. (1970) M.S. thesis, Washington State Univ., Pullman, Washington. [2] Edgett K. S. field notes, 1990-1995. [3] Mutch T. A. et al. (1976) Science, 193, 791-801. [4] Sharp R. P. and Malin M. C. (1984) GSA Bull., 95, 1398-1412. [5] Reid R. J. and Singer R. B. (1995) LPS XXVI, 1155-1156. [6] Mutch T. A. et al. (1977) JGR, 82, 4452- 4467, Fig. 6a. [7] Christensen P. R. and Moore H. J. (1992) in Mars (H. H. Kieffer et al., eds.), pp. 686-729, Univ. of Arizona, Tucson. [8] Arvidson R. E. et al. (1989) JGR, 94, 1573-1587. [9] Golombek M. P. et al. (1995) LPS XXVI, 481-482. [10] Edgett K. S., this volume. [11] Greeley et al. (1993) JGR, 98, 3183-3196. [12] Edgett K. S. and Blumberg D. G. (1994) Icarus, 112, 448-464. [13] Greeley R. et al. (1977) JGR, 82, 4093-4109. [14] Schultz P. (1994) Eos Suppl., 45, 406. [15] Rice J. W. and Edgett K. S., this volume. [16] Parker T. J., this volume. [17] Sharp R. P. (1973) JGR, 78, 4222-4230. Smith P. H.* A Targeting Strategy for Ensuring a Hillside View at Ares Valles An inspection of the Mars Pathfinder landing site at Ares Vallis with the stereo frames provided by the Viking Orbiter cameras shows a number of large-scale features that may be of considerable interest if they are within range of the IMP camera. This paper will characterize the heights of the features resolved by Viking and make estimates of how they might appear to the IMP camera at various distances, with the aim of adjusting the landing ellipse to maximize the probability of seeing one or more hills. Certainly, the most distinctive features within the landing site are the stream-flow islands, which have a tear-drop shape and are sometimes associated with an impact crater. Close inspection of orbiter pictures shows that the slopes of these features are layered, but little can be learned about this layering since it shows up at the limit of resolution. Possible explanations include topography such as sedimentary layers, ancient shorelines, erosion from multiple stream flows, or lava flows. It would be unfortunate to miss seeing such an interesting structure because the actual landing site happens to be 20 km away, below the local horizon. Naturally, one cannot prevent limiting the horizon if Pathfinder lands in a ditch or local depression. However, barring bad luck of that sort it is possible to improve the chances of seeing interesting steep terrain by choosing a landing ellipse within the vicinity of the Ares Vallis delta that guarantees large vertical features are within every view circle of the landing ellipse. A view circle can be defined by the distance that IMP (the Pathfinder camera) can resolve a steep slope with better resolution than the Viking orbiter pictures; with a 1-mrad/pixel resolving power, this distance is about 5 km and gives a pixel size of 5 x 5 m. Note that at this distance the surface curvature will only hide the lower 1 m of the feature. To estimate the density of sites, I divide the landing area by gridlines 10 km apart, twice the radius of the viewing circle. The major tall features within the area are apportioned to a gridbox and are counted. The best landing ellipse is the one that gives the highest probability of putting an observable tall feature within every grid box contained in the landing ellipse. Movements of the ellipse that serve to maximize the chance of viewing tall structures will not take away from the grab-bag nature of the site nor hinder any of the other goals of the mission. Besides searching for ancient shorelines and layering features, there are other important goals attached to seeing local hills besides studying the hills for themselves. First, the landing site appears to have a high enough density of tall objects that we are likely to be able to see two or more separate hills. This leads to the ability to exactly locate our landing site within the Viking pictures using simple triangulation from our relative coordinate system. It may be difficult for the MOC camera aboard the Mars Global Surveyor to identify our precise location on a rather featureless terrain. And second, the importance of a dramatic view should not be minimized. A 300-m-tall object seen at 3 km distance will subtend about 100 pixels and would be at least 300 pixels wide; a full color image of this hill would be very exciting. Tucsonans know that the beauty of the desert is revealed by the surrounding mountains. Friday, September 29, 1995 HAZARDS: ROCKS AND ROUGHNESS 1:30 PM Golombek M.* Rapp D. Size-Frequency Distributions of Rocks on Mars Predicting the size-frequency distribution of rocks at different locations on Mars is difficult owing to the limited dataset (ground truth from only two sites at the surface), but is critical for determining potential landing hazards for future Mars landers. In this abstract we (1) review rock frequency data at the two Viking landing sites and a variety of sites on Earth, with special reference to larger rocks that could be hazardous to a lander, (2) describe the data in terms of simple mathematical expressions, and (3) provide a means of extrapolating the data to any location on Mars from relationships between the rock frequency curves and remote sensing data. We used rock lengths, widths, and heights carefully measured from the stereo Viking landing images by Henry Moore and co-workers [1] consisting of a total of 421 and 486 rocks in areas of 83.7 m^2 and 83.76 m^2 at the Viking 1 and 2 sites (VL 1, VL 2), respectively. The rock data plotted in either cumulative number per square meter or cumulative fractional area vs. diameter have similar shapes at both Viking sites, displaying a convex up-curved shape on log-log plots that can be fit well with simple exponential functions. The rock data do not appear linear on log-log plots, so that power-law functions (commonly used to fit crater size-frequency data) overestimate the frequency and fractional area covered by both large-diameter and small-diameter rocks (Fig. 1). Similar shaped size-frequency distributions of rocks are found at a wide variety of rocky surfaces on the Earth (Fig. 1). Data collected by Malin [2] for (1) Icelandic catastrophic outflow deposits, (2) Antarctic dry valley wall talus, and (3) Hawaiian volcanic ejecta, as well as data we have collected from (4) Mars Hill, (5) an abandoned and washed alluvial fan in Death Valley, (6) a presently active alluvial fan on the eastern side of the Avawatz Mountains in the Mojave Desert, (7) two eroded and mass-wasted volcanic surfaces (basalt and tuff breccia) in the Eastern Mojave Desert (Goldstone), (8)catastrophic outflow deposits of the Ephrata fan in Washington state, and (9) a giant boulder field in the Leaf Basin of northern Quebec, in which boulders are transported downslope and washed in an intertidal zone [3], all show curved convex up size-frequency rock distributions on a log-log plot. Data from these sites have all been fit reasonably well with simple exponential functions, which describe both the precipitous drop-off in rocks with large diameters as well as the shallowing in cumulative number or area of rocks at small diameters. The VL 2 site is believed to be ejecta from the nearby crater Mie [4], whereas VL 1 is believed to be a partially covered and eroded lava flow surface, possibly with some local crater ejecta and flood deposits [5]. As a result, they appear to have formed by very different geologic processes, yet the shape of the rock size-frequency distributions at both sites are the same. The sites on Earth include alluvial fan water- rich debris flows (active and abandoned), catastrophic flood deposits, eroded volcanic surfaces, volcanic ejecta, and talus slopes, yet the rock size-frequency distributions are all similarly shaped. All sites show a precipitous fall-off in number or fractional area of rocks at large diameters, which may have something to do with the dearth of large coherent blocks of material and the inability of geologic processes to transport such large blocks without breaking them into smaller pieces. The consistency of the size-frequency rock distributions found on Earth and the two Viking landing sites suggests that similar shaped rock size- frequency distributions are applicable to other areas on Mars. A combined fit to both VL cumulative fractional area of rocks vs. diameter data was made with a general exponential function of the form F(sub)k(D) = k exp{-q(k) D}, in which F(sub)k(D) is the cumulative fractional area covered by rocks of a given diameter or larger, k is the total area covered by rocks at the site, and q(k) = (0.571 + 0.492/k). Simple linear height vs. diameter relationships, related to k, H = (0.25 + 1.4 k) D, were also derived from H/D ratios of ~3/8 and ~1/2 at VL 1 and 2 respectively, which suggest that less rocky areas on Mars have rocks with lower H/D ratios than more rocky areas. Height was then substituted into the general exponential function derived for diameter, which yielded F(sub)k(H) = k exp{-p(k) H} and p(k) = (0.571 + 0.492/k)/{0.25 + 1.4 k}, which describes the cumulative fractional area of rocks vs. height for any given total rock coverage. Viking thermal inertia measurements and models developed by Christensen [6] have been used to estimate the fractional surface area covered by high thermal inertia rocks greater than about 10 cm diameter vs. smaller particles, such as sand and dust, with low thermal inertia for 1 degree latitude by 1 degree longitude remotely sensed areas on Mars. Because the cumulative fractional area covered by rocks of 10 cm diameter and larger is fairly close to the total rock coverage, it can be used as the pre-exponential constant k in the general exponential function fit to the VL rock data to describe the cumulative fractional area vs. diameter or height at any location on Mars. This calculation is conceptually equivalent to Christensen and Malin's [7] suggestion that rock abundances on Mars reflect the thickness of mantling fine material. In this simple model, the maximum rock abundances (~30%) occur in areas with no mantling sand or dust, and less rocky areas (down to 2%) are mantled by progressively greater thicknesses of dust (up to 1 m thick). The exponential curves in Fig. 1 show these distributions in terms of cumulative area vs. diameter for any value of rock abundance, and the equations derived above show the decrease in H/D for less rocky areas. Results indicate that most of Mars is rather benign with regard to hazards from landing on large rocks. Roughly 50% of Mars has rocks covering only 8% or less of its exposed surface [6]. For total rock coverage of 8% analogous to VL 1, about 0.2% of the surface is covered by 20 cm or higher rocks. A surface covered with 12% rocks has only 1% of its surface area covered by rocks higher than 20 cm. The Mars Pathfinder lander airbag system is being designed to accommodate landing on 0.5 m high boulders. Such a landing system could land on a surface covered by about 20% rocks, similar to VL 2, with 1% of the surface covered by rocks of 0.5 m or higher. Surfaces with 20% or fewer rocks account for over 90% of the surface of Mars, so that such a landing system could be sent to all but the rockiest 10% of Mars with a low probability of landing on a >=0.5-m-high rock. The Ares Vallis landing site being considered for Mars Pathfinder has total rock abundances of ~20% [8], indicating a low probability of failure due to landing on large rocks. References: [1] Moore and Keller (1990) NASA TM-4210, 160; 1991, NASA TM-4300, 533. [2] Malin (1989) NASA TM-4130, 363. [3] Lauriol and Gray (1980) G. Soc. Can., Pap. 80-10, 281. [4] Mutch et al. (1977) JGR, 82, 4452. [5] Binder et al. (1977) JGR, 82, 4439. [6] Christensen (1986) Icarus, 68, 217. [7] Christensen and Malin (1993) LPSC XXIV, 285. [8] Golombek et al., this volume. Figure 1, which appears in the hard copy, shows cumulative fractional area covered by rocks vs. diameter for each rock at VL 1 and 2, Mars Hill sites (MH), Ephrata fan sites (EF), and highly binned data (only 4- 5 data points) from a variety of surfaces on the Earth. Power-law distribution suggested for VL 2 [1] also shown. Solid curves are the rock distributions predicted for various rock abundances (2%-30%) on Mars derived from a combined exponential fit to VL 1 and 2. Sites on the Ephrata fan, a depositional fan in the Quincy Basin of the Channeled Scabland, and potentially analogous to the Pathfinder landing site in Ares Vallis, Mars, include one with extreme rock coverage (70%), near the field trip stop in Rocky Ford Creek and another with 2% rock coverage. Harmon J. K.* Campbell B. A. Radar Scattering Characteristics of Ares Vallis and Environs from Arecibo Observations Quasispecular radar echoes can provide estimates of surface roughness (rms slope theta(sub)r) and the dielectric constant along the sub-Earth track on Mars. Such measurements were used for landing hazard assessment by the Viking 1 lander site selection team [1,2]. The radar theta(sub)r measurements were partly responsible for the decision to reject the original A1 site in favor of the smoother A1WNW site. Since the Mars Pathfinder lander site lies very close to the rejected VL1-A1 site, it makes sense to take a fresh look at the radar characteristics of this area. The sub-Earth track crossed the Pathfinder site during the most recent Mars opposition in early 1995, and preliminary results from Goldstone 3.5-cm ranging observations indicate high radar roughness (theta(sub)r > 10 degrees) over the landing site [3]. Although telescope upgrading work prevented us from making similar observations at Arecibo in 1995, we do have some Arecibo 12.6-cm radar data from earlier oppositions covering portions of the Ares Vallis region. Here we present some results from those observations. In Fig. 1 we show theta(sub)r values estimated from Arecibo ranging observations at 20.6 degrees N in 1980 (filled boxes) and at 23.2 degrees N in 1982 (open boxes); each of these points is obtained by fitting a single delay template to the region within 4 degrees of the sub-Earth point for a given 30-s data block. Also shown are theta(sub)r estimates from CW observations taken at 23.1 degrees N in 1976 (solid line) and at 22.5 degrees N in 1967 (dashed line); both these lines are adapted from Fig. 1 of Tyler et. al. [2]. These data all agree in showing that the Chryse plains are rough compared to the ridged plains and cratered plains to the east and west. Superimposed on this general trend are more localized roughness variations that can only be properly analyzed by making template fits to narrow Doppler slices. For example, the delay-Doppler array in Fig. 2 shows a "hole" at the central leading edge that corresponds to relatively weak echoes from the main channel of Ares Vallis at 23.2 degrees N, 31 degrees W. Template fits to this channel feature give theta(sub)r ~ 5 degrees for an assumed reflectivity of rho(sub)o = 0.07, indicating that at this latitude Ares Vallis is rougher than the terrain to the east, but not particularly rough in an absolute sense. In Fig. 3 we show a delay-Doppler plot taken at a sub- Earth point of 20.6 degrees N, 36.5 degrees W, closer to the Pathfinder landing site. The single delay template fit to this entire array gives theta(sub)r = 7.2 degrees and rho(sub)o = 0.068, values typical of Chryse in general. The righthand edge of this plot corresponds to the point where the sub-Earth track grazes the north edge of the Pathfinder landing site ellipse, and the weakness of the echoes at this point in the plot are consistent with some increase in roughness as one approaches the lander site from the west at this latitude. A template fit to a Doppler slice centered at 32.8 degrees gives theta(sub)r = 8.5 degrees for an assumed rho(sub)o = 0.068. This is significantly rougher than the northern branch of the Ares Vallis channel at 23.2 degrees N. It is smoother than the theta(sub)r > 10 degrees reported by the JPL group from their track across the lander site, but we do not consider this difference as very significant given that (1) the Arecibo scan is a degree north of the lander site, (2) the observations were made at a longer wavelength, and (3) we have had to assume a rho(sub)o value because we do not have sufficient longitude coverage at this latitude to do the same "Downs-style" scattering analysis as was done on the 1995 Goldstone data by the JPL group. In addition to these quasispecular results, we have coverage of the Chryse area in depolarized reflectivity maps from 1990 random-code observations. Depolarized enhancements are indicative of surface roughness at wavelength (decimeter) scales, i.e., scales smaller than those influencing theta(sub)r. The strongest such enhancements are found on the Tharsis volcanos and flows and in the Elysium Basin and outflow channel. In the Chryse channel region the strongest enhancements are found in Maja Vallis and the plateau bordering Simud Vallis. We find no strong feature at the Pathfinder lander site, although we do see a modest enhancement within the Ares Vallis channel south and east of the lander site. This indicates that the Ares Vallis channel is rockier than the cratered and ridged plains through which it flows, but does not have the sort of chaotic texture typical of the major volcanic provinces. We will discuss results from the depolarized maps in more detail at the workshop. References: [1] Masursky H. and Crabil N. L. (1976) Science, 193, 809- 812. [2] Tyler G. L. et al. (1976) Science, 193, 812-815. [3] M. Slade, private communication. Figures 1-3 appear in the hard copy. Slade M. A.* Mitchell D. L. Jurgens R. F. Dual Polarization Continuous Wave (CW) Observations from the Ares Vallis Site and Environs During the 1995 Mars Opposition, CW radar observations were performed over the Ares Vallis landing site in order to help assess the type of terrain that the Pathfinder spacecraft and rover would encounter. A total of nine transmit-receive cycles, or runs, were performed on 17 April, 1995. For each run, a ~100% circularly polarized signal was transmitted toward Mars, and echoes were recorded simultaneously in both the same sense of circular polarization as transmitted (the SC sense) and the opposite (OC) sense. The helicity of circular polarization is reversed upon reflection from a surface that is smooth on all scales within about an order of magnitude of the wavelength, but SC echo power can arise from single backscattering from a rough surface, from multiple scattering, or from subsurface refraction. The circular polarization ratio is thus a measure of near-surface complexity, or "roughness," at scales near the 3.5-cm observing wavelength. Figure 1 shows the OC Doppler spectrum obtained on April 17, 1995, when the mean subradar longitude and latitude were 32.7 degrees W. and 18.7 degrees N. The OC echo has been decomposed into two components: (1) a quasispecular echo (heavy solid line) originating from a small region, typically a few degrees in radius, centered on the subradar point and (2) a diffuse echo (heavy dashed line) reflected from the entire visible disk. The quasispecular component is modeled with a Hagfors law, yielding a C parameter of 80, corresponding to an rms surface slope (on scales much larger than the observing wavelength) of 6.4 degrees. The total OC radar albedo is 0.101, and the quasispecular portion of the OC radar albedo is 0.048. Figure 2 places these results into the context of the day's observations, showing the total OC radar albedo and the circular polarization ratio for this and the other runs of April 17. The implications of this part of the 1995 Mars radar dataset will be discussed at the meeting. Haldemann A. F. C.* Muhleman D. O. Jurgens R. F. Slade M. A. Assessment of Pathfinder Landing Site with Goldstone Radar Ranging and Goldstone-VLA Dual-Polarization Imaging The preliminary landing site for the Mars Pathfinder Lander/Rover has been chosen in southeastern Chryse Planitia (19.5 degrees N, 32.8 degrees W) [1]. The region is believed to be relatively smooth. Its location at the mouth of the Ares Vallis outflow will hopefully provide a variety of sampling opportunities for the Rover experiments. A more detailed review of the landing site characteristics can be found elsewhere in this volume [2]. The Ares Vallis site had been rejected as a Viking Lander site, however, due to Goldstone CW radar echoes with low signal-to-noise (SNR). We will present an updated assessment based on two types of recent radar observations: radar ranging with the Goldstone DSN antenna and radar imaging with the Goldstone-VLA combination. Both experiments were at X-band (3.5 cm). Goldstone radar ranging measurements were performed during Mars' most recent opposition in the winter of 1994-95, when the planet was at a distance of about 0.6 AU. Five radar tracks of interest for the primary landing site were recorded and are listed in Table 1. Three of the tracks pass over the site itself or its immediate vicinity. The Goldstone radar system is now 12 dB more sensitive than it was at the time of the Viking landing site assessment, and so we expect that a correspondingly more detailed assessment can be carried out. The data are in the form of delay-Doppler views, which fall along a subradar track on the planet. Views are collected during 10-min receive cycles separated by 10-min radar transmit cycles. Each view has about a 4 s integration time, and are summed in groups of 4 to improve SNR. The Doppler frequency resolution provides a longitudinal resolution of about 4 km, while the delay resolution makes for a latitudinal resolution of about 100 km. The data is of good quality, and should certainly provide consistent topographic profiles, which will be compared to the Mars digital elevation model [2]. The topographic profiles will be used to assess the kilometer-scale roughness in the landing site region. The delay-Doppler data can be fit with radar models, in particular the Hagfors scattering law, to extract reflectivity and the Hagfors C parameter, a proxy for rms slope. The reflectivity relates to the dielectric properties of the upper meter, and contain information about the bulk density of the material and some information about chemistry, e.g., mafic vs. silicic material. It will be interesting to find out how well these material property assessments are borne out when the Lander and Rover return results. The C parameter can be used in a comparative manner to assess surface texture or roughness, both on the scale of the wavelength and on scales large compared to the wavelength. The texture may be expressed on the surface, and thus is evident in visual images, or may be covered by up to a meter of soil and appear smooth in visual images. Our investigation will make comparisons of the radar results with Viking imagery, and these analyses may drive some choices of model parameters to aid the fitting routines. Indeed, the study will require careful analysis of the applicability of the Hagfors scattering model, and the limits of its interpretation (see for example [3]). The second dataset at our disposal was produced by interferometric imaging at the VLA of radar echoes from Mars. This is a huge dataset, which first observed the Stealth region on Mars, an essentially radar- absorbing region to the southwest of Tharsis [4]. These data were also used to observe the radar properties of the martian poles [5]. The resolution of these images is only of the order the landing site ellipse, but the image data cover a much larger range of incidence angles than the delay-Doppler data, and will thus be extremely useful to constrain the models used to fit the delay-Doppler data. Further, the Goldstone-VLA image data are in both circular polarizations. Polarization characteristics are a vital but complex part of the radar experiments. The details of the surface morphology are encoded into the polarization properties through multiple scattering, critical angle internal reflections, locally coherent scattering processes, and effects we have yet to imagine. Particular polarization signatures have been exploited for Mars [4], and will further provide constraints on the interpretation of the regional variation in wavelength-scale roughness. The results will be presented at the workshop. REFERENCES : [1] Golombek, this volume. [2] USGS (1993) Mars Global Topography CD-ROM, v.7. [3] McCollum and Jakosky (1993) JGR, 98, 1173- 1184. [4] Muhleman D. O. et al. (1991) Science, 253, 1508-1513. [5] Butler B. J. et al. (1994) LPSC XXV. Table 1. Dates and locations of radar tracks on Mars to be used in the radar ranging part of the study. Date Latitude Longitude (deg.) Locations (deg.) Begin End 18 Dec. 1994 21.8 56.5 94.2 Kasei, NE Tharsis 25 Dec. 21.7 330.8 25.1 Arabia, Estern Chryse 27 Jan. 1995 19.9 336.2 76.2 Arabia, Chryse, Kasei 29 Jan. 19.8 322.9 48.7 Arabia, Ares and Tiu 30 Jan. 19.7 321.2 39.9 ibid. Shepard M. K.* Guinness E. A. Arvidson R. E. Spatial Statistical Analysis of the Ares Vallis Region from Viking Orbiter and Worst-Case Scenario for Subpixel-Scale Roughness One of the immediate concerns of the Pathfinder mission team is the surface roughness of the proposed landing site. Is it a reasonably safe area to place a soft lander? This question is a difficult one to answer directly since Viking Orbiter images of the proposed site have minimum pixel resolutions of 40 m. Boulders or precipices much smaller than this could cause a mission catastrophe. The objective of our research is to infer some quantitative measure of the subpixel-scale roughness in this region using techniques of spatial statistics. In an area of homogeneous terrain, i.e., terrain of one geologic provenance, one can expect some statistical behavior of the topography as a whole. When this area is imaged by a satellite sensor, the terrain is essentially broken into equal-sized areas or pixels. A pixel represents the average reflectance behavior of all the area within its purview. The reflectance of each pixel is a function of the surface composition, the topography, and shadows cast by the topography. In the simplest case, the reflectance of a pixel is a linear sum of the reflectance of the surface material and the reflectance of shadows, weighted by their respective fraction of the pixel area. Assume a surface of single composition, illuminated by the Sun at an incidence angle >45 degrees. If the topography of the surface in question has elements with vertical scales comparable to the pixel in size, those elements will dominate the pixel reflectance because of the shadows present. Furthermore, the larger the topographic elements are with respect to the pixel size, the greater the variance in the number of those elements and shadows present. As a whole, pixels in this area will show a relatively large degree of variance in reflectance because some will have large numbers of shadows and others will not. Furthermore, the greater the contrast between illuminated surface and shadow, the greater the variance. Conversely, if the topography is dominated by elements with scales much smaller than the pixel size, then any given pixel will have (statistically) the same number of these elements as every other pixel. Therefore, as a whole, the pixels in this area will have less variance in reflectance. In areas where the surface is not of a single composition, there will be variance from pixel to pixel caused by compositional heterogeneity. The amount of variance caused by compositional heterogeneities can be determined by examining the surface when the Sun is directly overhead. At this point, there will be no shadows cast by topography and the entire variance in the scene can be attributed to compositional variation. This variance can then be "subtracted" from the variance in the scene at higher Sun angles, where the variance is a combination of compositional variations and topographic shadows, thereby leaving the variance caused by the topography alone. The scale of subpixel roughness can be quantified by utilizing a combination of analytic methods and forward models of computer synthesized surfaces. This methodology has been developed and tested successfully on an Earth analog site [1]. Unfortunately, all of the high-resolution (40 m/pixel) Viking Orbiter scenes of the Ares Vallis region have Sun angles of ~60 degrees. It is therefore impossible to estimate the variance in the images caused by compositional variances. Many (if not all) of the scenes are also plagued by a large atmospheric optical depth that significantly reduces the contrast between shadows and illuminated surface. However, by assuming that the surface is compositionally homogeneous, i.e., no variance due to compositional heterogeneity, a worst-case scenario for subpixel-scale roughness can be developed. We will present the results of this worst-case scenario, as well as other information gleaned from an analysis of the differences in variance behavior observed in different spectral windows (images acquired through the clear, violet, red, and green filters). References: [1] Shepard M. K. et al. (1993) LPSC XXIV, 1293-1294.