Mercurian Impact Craters: Implications for Polar Ground Ice

Icarus 141, 194–204 (1999) Article ID icar.1999.6165, available online at http://www.idealibrary.com on

Mercurian Impact Craters: Implications for Polar Ground Ice Nadine G. Barlow Department of Physics, University of Central Florida, Orlando, Florida 32816 E-mail: [email protected]

Ruth A. Allen Department of Geology, Dickinson College, Carlisle, Pennsylvania 17013

and Faith Vilas NASA Johnson Space Center, SN2, Houston, Texas 77058 Received October 24, 1996; revised March 19, 1999

Recent radar observations of Mercury have detected strong depolarized echoes from the north and south polar regions which have been interpreted by some as ice deposits in the floors of permanently shadowed impact craters. We have used the experience from Mars, where subsurface ice lowers the depth-to-diameter ratio (d/D) of impact craters, to test for subsurface ice deposits on Mercury. This analysis determines the d/D ratios for 170 impact craters in the Borealis (north polar), Tolstoj (equatorial), Kuiper (equatorial), and Bach (south polar) quadrangles of the planet. Possible effects from sun angle and terrain were eliminated. To test whether d/D differences could be detected at Mariner 10 resolutions (∼1 km/pixel), we perform a similar analysis using Mariner 9 images of Mars which have similar resolutions. We demonstrate that d/D differences due to terrain softening can be detected between craters in the martian polar regions and the equatorial regions at the Mariner 9 resolutions. Although our initial results indicate that the south polar Bach Quadrangle has a statistically lower d/D than the north polar (Borealis) or two equatorial (Tolstoj and Kuiper) quadrangles, further investigation reveals that this finding is most likely the result of the filtering which was applied to the images of the Bach quadrangle by JPL. Thus, no unequivocal evidence exists that the possible ice deposits in craters at Mercury’s north and south poles are the exposed portions of more extensive subsurface ice caps. Combined with the temporal constraint imposed by the fact that the proposed ice deposits are found only in USGS Class 4 craters, this suggests a large, rapidly emplaced exogenic source of water to Mercury during the Mansurian period. We suggest that the source was multiple impacts from a fragmented comet or a comet shower. °c 1999 Academic Press Key Words: Mercury; cratering; surfaces, planets; ices; Mars.

INTRODUCTION

Recent ground-based radar observations of Mercury have detected strong, highly depolarized echoes from the north and

south polar regions which have been interpreted as possible polar ice deposits (Slade et al. 1992, Harmon and Slade 1992, Butler et al. 1993, Harmon et al. 1994). The radar echoes that are found on the side of Mercury imaged by Mariner 10 have been correlated with relatively fresh impact craters (USGS grade 4) (Harmon et al. 1994). However, Mariner 10 images show no albedo changes between the visible portions of these craters and the surrounding terrain. This observational evidence suggests that the proposed ice deposits do not form continuous polar caps on the surface of Mercury but rather exist in the permanently shadowed floors of some high latitude impact craters. Ice deposits at Mercury’s poles were proposed prior to the bright radar signal discoveries (Thomas 1974, Kumar 1976, Gibson 1977). Theoretical studies indicate that surficial ice deposits can be stable throughout the lifetime of the planet in the permanently shadowed regions of craters poleward of about 83.5◦ latitude (Paige et al. 1992, 1993; Ingersoll et al. 1992), but are likely unstable across exposed surfaces regardless of latitude unless surface albedos are higher than observed (Paige et al. 1992, Butler et al. 1993). Based upon the strength of the radar reflectivity, Butler et al. (1993) proposed that either incomplete areal coverage occurred in the first radar images or a small (∼10-cm-thick) layer of dust or soil was emplaced over the ice deposits, or both. A small insulating layer of dust or soil also would prevent ice from sublimating off the surface. They argued that ice existing below such a small layer of particulate matter would need to be close to pure ice in order to be detected by the radar. The ice probably would have to be emplaced very rapidly, and Butler et al. suggested that volcanic outgassing of volatiles could be one endogenic source for the ice. Harmon et al.’s (1994) detection of bright radar echoes correlated with craters at the mercurian poles confirmed that the presence of the proposed ice deposits was not uniform across the surface

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around the poles. Their results did not indicate if an insulation layer of material is necessary to explain the observations, and as such the existence of a particulate layer remains an open question. Other sources for the radar bright material at the poles have been proposed. Potter (1995) proposed that chemical sputtering of surface rocks by protons from the planet’s magnetosphere would be sufficient to produce water molecules, some of which could be trapped as surficial polar ice deposits. Killen et al. (1997) and Moses et al. (1999) investigated the possibility that the sources are emplaced volatiles from the impacts of external sources, including comets, asteroids, and interplanetary dust particles. Sprague et al. (1995) alternately proposed that the radar-bright spots are not water ice at all, but rather are sulfur deposits at the mercurian poles, a proposal that has recently been reexaminated by Butler (1997). In this paper we consider whether subsurface ice caps could exist at the poles of Mercury, exposed in the permanently shadowed regions of craters near the poles. There are a number of reasons why we could expect polar ice deposits on Mercury. Evidence continues to build that Mercury has experienced an active volcanic history. Radar signals from the unimaged side of Mercury show a large conical structure similar in relative size and shape to the large shield volcanoes seen on other terrestrial planets (Harmon 1997). Recalibration and new analyses of the Mariner 10 color images show color units correlated with geologic flow units suggesting a volcanic origin (Robinson and Lucey 1997), and possible pyroclastic centers (Robinson et al. 1998a). The global distributions of volcanic centers are nonrandom on large scales for the Moon and three other terrestrial planets. For Earth, Venus, and Mars, more than half of the population of source areas for volcanism occur in less than 30% of the planet’s surface area (Crumpler and Revenaugh 1997); mare deposits cover only 17% of the Moon’s surface area (Head 1976), and anomalous volcanic source areas on the Moon (e.g., Marius Hills and Aristarchus Plateau) are equally nonrandom (Guest and Murray 1976, Whitford-Stark and Head 1977, Head and Wilson 1992). Thus, conclusions about the presence and state of volcanism on Mercury based on observation of ∼45% of the planet’s surface could be unrepresentative of the true state of Mercury’s volcanic history. We conduct here the first experimental test of whether the radar bright signals at Mercury’s poles could be caused by water ice having an endogenic origin. Outgassing of volatiles as a result of volcanism could result in a trapped, clean subsurface ice layer at the poles. In addition, some percentage of volatiles outgassed across the rest of Mercury’s surface would migrate and be trapped in permanently shadowed regions at the poles. Volcanism on Mercury would have occurred early in the planet’s history after the interior likely had been heated by accretion, differentiation, electromagnetic induction by a T-Tauri-like solar phase, and/or decay of 26 Al during the first 106 –107 years of solar system existence. As discussed above, ice existing at Mercury’s poles should be stable across the planet’s lifetime.

Thus, any subsurface ice caused by outgassing due to volcanism likely will have existed over most of Mercury’s lifetime. Polar ice deposits are known to exist on two other terrestrial planets in the Solar System, Earth and Mars, and have been suggested at the polar regions of the Moon (Watson et al. 1961, Arnold 1979, Hodges 1980, Ingersoll et al. 1992, Nozette et al. 1996, Feldman et al. 1998). Mars in particular offers some interesting lessons in the effect of surface and near-surface ice on the morphologies of impact craters. Since most of Mercury is heavily cratered, such lessons from Mars are applicable to a search for subsurface ice on Mercury. Cosmochemical models of Mars suggest that the planet should be rich in volatiles, including H2 O. The low atmospheric pressure prevents liquid water from currently existing on the planet’s surface, but H2 O does exist in solid form in the polar caps and in vapor form in the martian atmosphere. In addition, the Viking, Mars Pathfinder, and Mars Global Surveyor missions provide abundant evidence that climatic conditions in the past have been much warmer and wetter, allowing rivers, lakes, and possibly oceans of liquid water to exist on the surface (Golombek et al. 1997, Smith et al. 1998, P. R. Christensen, press release of May 27, 1998). As climatic conditions changed and the environment became colder and drier, much of this water probably accumulated into subsurface reservoirs (Carr 1996). Geologic evidence of this subsurface H2 O includes small valley network systems formed by sapping (Baker 1982, Baker et al. 1992), outflow channels formed by vast quantities of water suddenly released from the substrate (Baker 1982), lobate ejecta flows around impact craters (Carr et al. 1977, Barlow and Bradley 1990), and softening of terrain features at high latitudes (Squyres and Carr 1986). Thermal models suggest that most of this subsurface H2 O is likely in the form of ice, although the proposed thermal gradient for Mars could allow liquid water to exist in brine reservoirs below 1 to 2 km depth in the equatorial regions (Fanale 1976, Squyres et al. 1992, Clifford 1993). The effects of solid versus liquid H2 O may be indicated by variations in the lobate ejecta morphology displayed by impact craters of various sizes at different latitudes—studies suggest that the smaller single lobe craters form by impact into ice while the larger multiple-lobe craters form by impact through ice and into water-rich reservoirs (Barlow and Bradley 1990). Another effect of subsurface H2 O on impact craters is the decrease in depth– diameter ratio (d/D) and the softening of sharp features such as crater rims with time (Fig. 1) (Cintala and Mouginis-Mark 1980, Squyres and Carr 1986). This “terrainsoftening” results from the lower strength properties of materials containing ice versus dry materials, which facilitate creep of the material. The change in crater depth with time due to such relaxation is given by (Melosh 1989) H (t) = H (0) e−t/τ , where H (t) is the crater depth at time t, H (0) is the initial crater depth, t is the elapsed time since crater formation, and τ is the

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FIG. 1. Terrain softening on mars. These two images of martian impact craters show the effects of terrain softening on craters at high latitudes. (a) shows several craters in various stages of preservation near the martian equator. All however show fairly sharp, although degraded, appearance (Viking Orbiter Frame 441S01). (b) shows craters near the south polar region of Mars. Ridges and crater rims appear rounded and crater depths are shallower than those for similarly sized craters in the equatorial region. The shallow depths and rounded edges are common features of terrain softening, caused by subsurface creep of ice-rich materials (Viking Orbiter Frame 196S11). The largest crater in both images is about 15 km in diameter.

relaxation time, which is a function of crustal viscosity (η), crustal density (ρ), gravitational acceleration (g), and crater diameter (D): τ∼ = 8η/ρg D. The inclusion of H2 O into normally dry crustal material will decrease the viscosity of the material, leading to more rapid relaxation of the crater floor. This process, although requiring longer periods of time to occur than ejecta fluidization or channel formation, requires a much lower concentration of volatiles than most of the other geologic indicators of subsurface volatiles. Terrain softening appears to have occurred on Mars since impact craters at high latitudes (generally poleward of about ±40◦ latitude) show a decrease in d/D compared to craters at equatorial latitudes where ice/water is expected to be a lower percentage of the substrate (Jankowski and Squyres 1992). Simple craters at high latitudes are about 80% as deep as similarly sized craters at lower latitudes, while complex craters are about 50% as deep as their equatorial counterparts. The radar echoes suggesting possible ice deposits within impact craters at Mercury’s poles led us to question whether such ice might extend into subsurface polar ice deposits. The presence of ground ice would serve to illuminate the extent and formation of the radar bright echoes. Mercury shows no evidence of channels formed by either sapping or outbreak, so subsurface reservoirs of ground water are counterindicated. In addition, no lobate ejecta blankets have been observed surrounding mercu-

rian impact craters, indicating that volatiles do not constitute a large enough percentage of the subsurface materials to produce the fluidization process necessary to emplace the ejecta into lobate patterns. However, no quantitative study of possible changes in crater d/D with latitude which could reveal the effects of terrain softening has been undertaken for Mercury. As discussed above, the volcanic outgassing which could supply the volatile reservoirs at the polar regions occurred early in mercurian history and hence a sufficient period of time has elapsed during which we could potentially observe crater relaxation effects. We thus initiated a study to test this hypothesis and determine if the cratering record on Mercury reveals any indications of subsurface polar ice deposits. EXPERIMENTAL PROCEDURE AND INITIAL RESULTS

We selected four regions on Mercury for this study: the two polar regions (Borealis (H1) and Bach (H15)) where bright radar echoes have been observed, and two equatorial regions (Tolstoj (H8) and Kuiper (H6)) where neither surface nor subsurface ice is expected (Fig. 2). These four areas are representative of the terrain composing the mercurian surface as imaged by the Mariner 10 cameras. The surface of Mercury can be divided into several morphologic units, but primarily consists of heavily cratered terrain, intercrater plains, and smooth plains (Strom 1984, Spudis and Guest 1988). These three terrain units are well represented in the Borealis, Tolstoj, Kuiper, and Bach

FIG. 2.

Mercury map. This map shows the area of Mercury imaged by Mariner 10. Areas in this study are outlined in black.

MERCURY CRATERS AND POLAR GROUND ICE

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quadrangles used in this study. Crater size-frequency distribution studies indicate that the heavily cratered terrain is the oldest unit on Mercury, followed by the intercrater plains (Leake 1982, Strom and Neukum 1988). Both of these units are categorized as pre-Tolstojan in the stratigraphic sequence developed for Mercury, indicating formation ages prior to about 4.0×109 years (Spudis and Guest 1988). The smooth plains, particularly those surrounding the Caloris Basin, display lower crater densities and are thus younger, dating from the Calorian stratigraphic period around 3.9 × 109 years ago. For our purposes, we followed the example of Pike (1988) and divided the mercurian surface terrain into “smooth terrain” and “cratered terrain.” We chose to group the cratered terrain types (intercrater plains and heavily cratered terrain) in this manner and disregard any slight age differences that might exist in the cratered terrain. The “smooth terrain” contains only smooth plains material. We utilized the individual Mariner 10 photographic frames (not photomosaics) in this analysis. The impact craters used in this study are those that display fresh morphologic characteristics (i.e., sharp rims, no evidence of floor deposits, pristine to slightly eroded ejecta blankets). All of the craters in this study are superimposed on the underlying terrain units (“cratered terrain” and “smooth terrain”). Because of their fresh appearance, most craters in this study are expected to have formed during the Kuiperian (<1 Gyr old) or Mansurian (between 1 and 3.5 Gyr old) stratigraphic periods (Spudis and Guest 1988). The two polar regions selected for this study are the north polar Borealis Quadrangle and the south polar Bach Quadrangle (Davies et al. 1978). The Borealis Quadrangle (H1) covers the north polar region of Mercury above +65◦ latiude. The area from the eastern limit of visibility near 20◦ W longitude to about 100◦ W longitude is covered by the smooth plains materials of Borealis Planitia. The rest of the quadrangle is composed of heavily cratered terrain and intercrater plains. The Bach Quadrangle (H15) covers the south polar region of Mercury poleward of −65◦ latitude. The Bach Quadrangle is covered by heavily cratered uplands and intercrater plains of pre-Tolstojian age. We also examined the Tolstoj Quadrangle (H8) as representative of the mercurian equatorial region, an area on the planet subject to extreme temperatures. Because of the high surface temperatures encountered in this region, theoretical models predict that ice cannot exist on the surface or in the near-surface substrate (Paige et al. 1992, Butler et al. 1993, Ingersoll et al. 1992). The absence of radar bright echoes at the equator suggests that the radar bright material does not exist at this location on Mercury’s surface. The Tolstoj quadrangle is located to the southeast of the Caloris Basin and covers the area between ±25◦ and −25◦ latitude, 145◦ to 190◦ W longitude. Most of the northeastern portion of the quadrangle is covered by the smooth plains surrounding the Caloris Basin (Tir Planitia; Budh Planitia). The southwestern portion of the quadrangle consists primarily of intercrater plains, the exception being around Tolstoj Basin which is surrounded by plains characterized by

radial lineations and grooves. This latter plains unit, called the Tolstoj Basin Rim Material, is interpreted to be ejecta from the Tolstoj Basin (Schaber and McCauley 1980). The Tolstoj Basin and Rim Material define the Tolstojan stratigraphic unit, emplaced around 3.9– 4.0 × 109 years ago (Spudis and Guest 1988). In addition to their locations, these quadrangles were chosen for this study based on the spatial resolution and Sun angles (the angle formed by the Sun and a normal to the planet’s surface at a given location) of the Mariner 10 images. Our study includes determination of crater depth from shadow length estimates, so our analysis was restricted to craters with Sun angles between 55◦ and 89◦ . This limitation allowed us to accurately measure only those craters in latitudes equatorward of ±85◦ , since Sun angles for craters at higher latitudes were too low and shadows cast by the crater rims extended beyond the floor up onto the opposing crater wall. Thus, we cover a small portion of the region near Mercury’s poles that theoretically could have ice in permanently shadowed regions of craters and presumably could be affected by subsurface ice. Most of the craters correlated with the radar bright echoes exist above the ±85◦ latitude or in unimaged areas of the mercurian surface. However, our study was able to include three of the craters suggested to contain ice based on the radar reflectivities (crater W in the Borealis Quadrangle, and craters G and V in the Bach Quadrangle) (Harmon et al. 1994, Robinson et al. 1998b). Shadow length measurements and rim-to-rim crater diameters were obtained by digitization of Mariner 10 images with 1 km/pixel resolution or better. Crater depths were determined using the shadow length measurements and the approximate Sun angles, following the method described by Clow and Pike (in Pike 1988). Values of d/D were calculated for all individual craters in this analysis. Measured craters were mapped on a shaded relief map and divided into simple and complex craters, based on the morphologic and morphometric measurements by Pike (1988), who found the simple-to-complex transition diameter for mercurian craters is 10.3 ± 4 km. We related the mean d/D ratio to the change in latitude by looking at craters within 10◦ latitude bins (Table I). We also looked for changes in d/D within the quadrangles for both the simple and complex craters. Results were analyzed using a Student-t distribution statistical test. We test the accuracy of the measurements by repeating the procedure with a “control” equatorial region, the Kuiper Quadrangle (H6). Craters within Bach, Borealis, and Tolstoj had similar lighting while those in the Kuiper Quadrangle average lower Sun angles. These control measurements allowed us to compare measurements of fresh craters on similar terrain (i.e., cratered terrain) under different lighting conditions, testing whether the extreme Sun angles would affect the resulting d/D measurements. The Kuiper Quadrangle is located between +25◦ and −25◦ latitude, 10◦ to 70◦ W longitude, on the eastern limit of Mariner 10 coverage. It is entirely covered by the pre-Tolstojianaged heavily cratered terrain and intercrater plains (DeHon et al.

MERCURY CRATERS AND POLAR GROUND ICE

TABLE I Average Crater d/D by Latitude Average d/D Quadrangle Latitude range

Simple

Complex

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ANALYSIS OF MARTIAN CRATERS

Number of craters Simple

Complex

9 16

2 9

Borealis Borealis

75◦ –80◦ N 65◦ –75◦ N

0.09 ± 0.03 0.09 ± 0.02 0.10 ± 0.05 0.09 ± 0.05

Tolstoj Tolstoj Tolstoj Tolstoj Tolstoj

15◦ –25◦ N 5◦ –15◦ N 5◦ S–5◦ N 5◦ –15◦ S 15◦ –25◦ S

0.09 ± 0.03 0.10 ± 0.06 — 0.10 ± 0.04 0.14 ± 0.08

0.09 ± 0.02 0.06 ± 0.02 0.06 ± ∞ 0.18 ± 0.10 0.10 ± 0.09

5 3 0 4 9

3 3 1 4 3

Kuiper Kuiper Kuiper Kuiper Kuiper

15◦ –25◦ N 5◦ –15◦ N 5◦ S–5◦ N 5◦ –15◦ S 15◦ –25◦ S

0.09 ± 0.03 0.07 ± 0.04 0.09 ± 0.02 0.06 ± 0.01 0.11 ± 0.10

0.16 ± ∞ 0.09 ± 0.04 0.09 ± 0.02 0.03 ± ∞ 0.08 ± 0.04

5 7 6 2 2

1 6 2 1 4

Bach Bach

65◦ –75◦ S 75◦ –85◦ S

0.07 ± 0.04 0.05 ± 0.02 0.06 ± 0.04 0.03 ± 0.02

33 8

12 10

1981). Table I shows the mean d/D values obtained for 10◦ latitude bins for the Kuiper and Tolstoj Quadrangles. No significant difference between these values exists, suggesting that the Sun angle differences are not important contributors to any observed differences in crater d/D between mercurian quadrangles. The Mercury study areas were selected primarily for their range in latitudes and Sun angle; geologic terrain units were a secondary consideration. However, terrain does often affect the d/D for impact craters (Pike 1988; Barlow 1993, 1999), so we investigated whether any statistically significant changes in our d/D values could be attributed to differences in geologic units (i.e., cratered terrain versus smooth plains). We found that there was a d/D difference between cratered terrain and smooth plains (in Borealis, simple craters on smooth plains had an average d/D of 0.15 compared to the 0.09 value found in cratered terrain; in Tolstoj, simple craters on smooth plains showed an average d/D of 0.16 compared to 0.11 in cratered terrain). Pike (1988) also found differences in the characteristics of simple and complex craters between the different terrains. Because cratered terrain is widespread in all four quadrangles studied here, we concentrated our study only on craters within the cratered terrain. A total of 170 fresh craters having diameters ranging from 0.7 to 32.0 km were identified and measured. Simple craters constituted 109 of the 170 craters while complex craters account for the remaining 61. We found a large deviation in the d/D values among individual craters in each quadrangle. We attribute this result to the low image resolution of the Mariner 10 FDS images (generally between 1 and 1.5 km/pixel). We wondered whether the relaxation in crater depth expected from subsurface ice deposits could actually be detected with these resolutions, so we tested the validity of our results by returning to the experiences with Mars.

Unlike Mercury, Mars has been imaged by a number of different spacecraft at varying resolutions and seeing conditions. Virtually all recent research about Mars has been based on the Viking images, which were obtained during relatively clear atmospheric conditions and at fairly high resolutions (200 m/pixel for the entire planet, and better than 10 m/pixel resolution in some areas). The documentation of latitudinal variations in crater d/D ratios was obtained using Viking data (Squyres and Carr 1986, Jankowski and Squyres 1992). The earlier Mariner 9 images of Mars, however, provide a better match in the quality and resolution (approximately 1 km/pixel) to the images we used in our Mercury study. Thus, knowing the expected outcome of comparing craters in polar and equatorial regions on Mars, we partially repeated the procedure for calculating d/D ratios on Mars using Mariner 9 images so as to determine if lower resolution images would still show the effect. Two areas were selected on Mars to conduct the comparison study of image resolution: an area near the south pole and an equatorial region (Fig. 3). The south pole study area is located between −50◦ and −75◦ latitude and westward from 340◦ W longitude to 40◦ W. Sun angles range from 54◦ to 78◦ for this area. This study area includes portions of the Argyre (MC-26), Noachis (MC-27), and Australe (MC-30) Quadrangles. Most of the terrain covered by this study area consists of old Noachian-aged (older than 3.9 × 109 years) cratered plateau materials (Hodges 1980, Peterson 1977, Condit and Soderblom 1978, Scott and Tanaka 1986, Greeley and Guest 1987, Tanaka and Scott 1987), with Argyre crater rim material dominating the western edge of the area. The second area studied on Mars lies between +25◦ and −25◦ latitude and extends westward from 345◦ W longitude to 40◦ W. Sun angles for this area range from 49◦ to 70◦ . This study area covers portions of the Oxia Palus (MC-11), Arabia (MC-12), Margaritifer Sinus (MC-19), and Sinus Sabaeus (MC-20) quadrangles. This area also lies within the heavily cratered southern highlands of Mars and is dominated by the same Noachian-aged cratered plateau materials seen in the south polar study area (Wilhelms 1976, King 1977, Saunders 1979, Moore 1980, Scott and Tanaka 1986, Greeley and Guest 1987, Tanaka and Scott 1987). Thus both study regions on Mars display similar terrain units, and terrain effects on crater d/D can be eliminated. We measured 131 craters (81 simple, 40 complex) on Mars for these heavily cratered regions near the South Pole and equator. The results closely matched the findings from studies using Viking data in that there is a statistically significant difference in d/D between the polar (d/D = 0.14 ± 0.02 for simple; d/D = 0.10 ± 0.02 for complex) and equatorial (d/D = 0.20 ± 0.02 for simple; d/D = 0.16 ± 0.02 for complex) regions on Mars, with craters at higher latitudes displaying lower d/D values due to terrain softening. The positive results of this test give us confidence that similar variations can be identified on Mercury at equivalent Mariner 10 resolutions.

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FIG. 3. Mars map. This map shows the area of Mars imaged by the Viking Orbiters. The Mariner 9 areas used in this study are outlined in black.

STATISTICAL COMPARISONS OF d/D VALUES

Table II summarizes the number of craters, average d/D values, and standard deviation for the four mercurian quadrangles used in this study. Detailed information about every crater in this analysis is available from the authors on request. The majority of the craters are excavating to about 1 km depth, with only a few larger craters extending down as deep as 2 to 3 km. Hence 1 km is the depth to which we test the proposal of subsurface ice caps. Table II shows that craters within the Borealis, Tolstoj, and Kuiper quadrangles display similar d/D ratios, but craters in the south polar Bach Quadrangle have a lower d/D.

We performed a Student-t distribution test on these results to determine if the craters in Bach were statistically different from the craters in the other quadrangles. The Student-t test was selected for this analysis since we are testing two sample populations to determine if the means of the two populations are identical (Devore 1987). We compared each quadrangle with the other three to determine if the d/D values were statistically identical at the 95% confidence interval. The results of the Student-t test, shown in Table III, indicate that the crater d/D for Borealis, Tolstoj, and Kuiper are statistically identical, but that the d/D of simple craters in Bach is statistically different from the other three quadrangles.

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TABLE II Characteristics of Craters by Quadrangle Quadrangle

Number of craters

Mean d/D

Standard deviation

Borealis Simple Complex

25 11

0.10 0.09

0.045 0.046

Tolstoj Simple Complex

21 14

0.11 0.11

0.060 0.075

Kuiper Simple Complex

22 14

0.09 0.09

0.034 0.043

Bach Simple Complex

41 22

0.06 0.04

0.034 0.018

Note. Statistics include only craters on the cratered terrain.

We addressed several aspects of image acquisition and processing which we believed might contribute to this difference in d/D. The images utilized in our analysis of the Borealis, Tolstoj, and Kuiper quadrangles are primarily from the first Mariner 10 pass during March 1974 (M1 encounter). The signal-to-noise ratio of these data was good and many of the images were recorded on the tape recorder and transmitted back to Earth at low data rate. However, during the six hours of closest approach, images were transmitted to Earth in real-time at a rate of 117.6 kbit/s (Soha et al. 1975). These are the images used in our analysis. Some images of Bach were also acquired during this pass, but the images are of very low resolution. Bach was imaged at resolutions comparable to those that we used in the other three quadrangles during the second Mariner 10 pass by Mercury in September 1974 (M2 encounter), but by this time the tape recorder had failed so all the data had to be sent in real time at 117.6 kbit/s (Soha et al. 1975). Some images of the equatorial quadrangles were also acquired at high enough resolution during this encounter to be useful in our analysis. TABLE III Results of Student-t Test Quadrangles compared

Calculated t-value

Comparison t-value

Result

Tolstoj–Borealis—Simple Tolstoj–Borealis—Complex Tolstoj–Bach—Simple Tolstoj–Bach—Complex Tolstoj–Kuiper—Simple Tolstoj–Kuiper—Complex Borealis–Bach—Simple Borealis–Bach—Complex Kuiper–Bach—Simple Kuiper–Bach—Complex Borealis–Kuiper—Simple Borealis–Kuiper—Complex

0.61 0.78 4.20 4.22 1.35 0.87 4.09 4.52 3.34 4.86 0.85 0.00

2.01 2.01 2.00 2.45 2.01 2.06 1.99 2.04 1.99 2.03 2.01 2.07

Same Same Different Different Same Same Different Different Different Different Same Same

We compared images of 22 craters that were obtained during both the M1 and M2 encounters. As before, we digitized the crater diameters and shadow lengths and computed the crater depths from the shadow length and Sun angle data. The measured diameter of each crater in this comparison was identical between the two encounters within the measurement limits. When we compared the d/D values for identical craters between the M1 and M2 encounters, we found that they were statistically identical at the 95% confidence interval. Hence errors introduced in the acquisition and processing of images acquired during the two different encounters cannot explain the lower d/D values found in Bach. Images downlinked to the Earth in real time were filtered after being radiometrically decalibrated (Soha et al. 1975) in order to correct single-pixel errors introduced by the faster data transmission. All of the Bach Quadrangle images used in our analysis were filtered. In Borealis, only 3 simple craters out of the total of 25 were on images that were filtered, while in Tolstoj 10 craters were measured on filtered images and 11 were not. We compared the d/D values between filtered and nonfiltered craters in Tolstoj, where the number of craters in each category is comparable. The average simple crater d/D for the filtered craters was 0.08 ± 0.04, while for nonfiltered images the mean d/D was 0.14 ± 0.06. Performing a Student-t test on these results, we found that the means of the two groups are statistically different at the 95% confidence level (computed t = 2.746; comparison t = 2.093). This indicates that images processed using the JPL filter are displaying lower crater d/D than images processed without the filter. Since the images used in the analysis of the Bach quadrangle were all filtered while most of those in the other three quadrangles were not, we believe that this is the major contributor to the difference in d/D values that we have obtained and that no evidence exists for terrain softening at either pole of Mercury. DISCUSSION

The results of this study conclude that, within the accuracy possible from the Mariner 10 imagery and compensating for the differences introduced by the high-pass filtering, no variation in crater depth–diameter ratios which can be attributed to subsurface ice caps can be detected in the polar regions of Mercury. This observation only rules out the presence of enough subsurface ice to have caused terrain softening. Our results support the theory that, if the bright radar reflectivities are the result of ice, the ice is concentrated in only small deposits within the permanently shadowed floors of high latitude impact craters and does not extend into a large subsurface ice deposit. One question that previous theories have not fully address is why the radar bright signals are only located in the interiors of the USGS Class 4 craters and not the older Classes 1–3 or younger Class 5 craters. The Class 4 craters date to the Mansurian chronostratigraphic system of Mercury’s history, occurring roughly 3.0–3.5 Gyr ago (Spudis and Guest 1988). This

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isolation of the bright radar signals in these craters places potential constraints on the relative age and duration of the event or events that created the material which is the source of the radar bright signals. Many studies have shown that ice at the poles of Mercury, especially in the permanently shadowed floor regions of larger craters, could theoretically be maintained over long time periods by one or more of a variety of steady-state sources and sinks (Paige et al. 1992, 1993; Ingersoll et al. 1992). If the source is endogenic (i.e., volcanism), then the layer of relatively clean ice was emplaced prior to the USGS Class 4 craters being created. If this emplacement occurred close in time to the creation of the Class 4 craters, perhaps only those craters would have excavated into the subsurface ice-rich layer and perferentially collected the ice along their floors. Another possibility is that the source of the ice occurred even further back in Mercury’s chronology and was exposed in the older, more degraded craters present on Mercury’s surface (any or all of the USGS Classes 1–3) as well as the Class 4 craters. Material formed by processes occurring between the formation of the Class 3 and Class 4 craters, such as ejecta from younger crater formation or depositions from nearby micrometeorite impacts, then buried the ice in the Class 1–3 craters. However, neither of these theories explain why the younger USGS Class 5 craters do not expose the subsurface ice layer. Likewise, exogenic sources of water should collect in the permanently shadowed regions of craters beyond only the Class 4 craters if a steady-state or slower process created the water ice. One possible explaination for why the Class 1–3 craters do not show the radar signals is that infilling of crater floors by volcanism and/or ejecta blanketing could either bury any ice deposits or, together with rim degradation, keep the crater too shallow for permanently shadowed regions to occur on the floor. Similarly, the steeper walls of younger Class 5 craters may reflect and reemit more radiation to their interiors which would increase the temperature and limit the amount of ice present (Paige et al. 1992). Between these two extremes are the Class 4 craters, which should be able to retain water ice in their permanently shadowed floors. However, based on the results of this study, we propose another possible explaination. The temporal constraint imposed by the presence of this water ice only in USGS Class 4 craters, combined with our evidence that no widespread subsurface ice layer exists, supports a large, rapidly emplaced exogenic source of water to Mercury during or near the Mansurian period. This event was likely large enough to have involved the entire planet, as water must have been deposited at both poles. The event also was likely a limited event—a more common solar system phenomenon would have created polar ice deposits throughout more of Mercury’s history. We suggest that multiple impacts from one fragmented comet, such as the multiple impacts from Comet Shoemaker– Levy 9 into Jupiter’s atmosphere in 1994, could provide the ice necessary to create the cleaner ice deposits distributed across Mercury’s polar regions. Alternatively, impacts from a comet shower could distribute water ice across the surface of the planet

with perhaps a better latitudinal distribution than a fragmented comet could provide. We note that the relative velocities of Mercury and any impacting comets must be very low for water to be retained at the poles (Butler et al. 1993), as theoretical studies indicate that comet impacts to the Earth would lose most of their volatiles during high-velocity collisions (Chyba 1987). Swirl patterns near craters on Mercury have been attributed to local regolith scouring and melting by higher density gas/dust zones (10 − 10 g cm − 3 ) within the inner comae of comets, suggesting observational proof that comets have impacted the mercurian surface (Schultz 1988, Schultz and Srnka 1980). Once the ice is deposited by comet impact, steady-state source and sink processes would be expected to maintain the ice deposits in the polar regions. ACKNOWLEDGMENTS The authors thank the Lunar and Planetary Institute for access to equipment and Mariner 10 data, Debra Rueb for photographic reproductions, and Fred H¨orz for instrumental support. We thank M. G. Lancaster and two anonymous reviewers for constructive comments which have greatly improved the manuscript. Support for R.A.A. came from the 1995 NASA Johnson Space Center Intern Program and F.V. acknowledges the NASA Planetary Astronomy Program.

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