Meteorite impact craters represent important geological features for revealing the near-surface layers of a planetary surface. In the case of Mars, this characteristic was proposed as a useful method to study spatial variations of such attributes as the distribution of sub-surface volatiles, and heat flow. Using the Planetary Image Cartography System (PICS) software, a quantitative analysis was completed of the geometry of fresh impact craters in the Hesperia Planum region of Mars, where an uniform target material and optimum viewing geometry make possible an analysis of target effects over a large geographic region. Because of the morphologic similarity to the lunar maria, it is likely that Hesperia Planum comprises a series of flood lavas that partially infilled topographic depressions within the Martian highlands. Measurements of partially buried crater rims suggest that the lava flows within Hesperia Planum are between 200-400 m thick

Hayashi, J. N.

Mouginis-Mark, P. J.

English

NASA Technical Reports Server

N92-10873. 
DEPTH/DIAMETER RELATIONSHIPS OF FRESH CRATERS WITHIN
HESPERIA PLANUM, MARS.
P. J. Mouginis-Mark & J. N. Hayashi, Planetary Geosciences Division, Dept. Geology &
Geophysics, SOEST, University of Hawaii, Honolulu, HI 96822.
INTRODUCTION
Meteorite impact craters represent important geological features for revealing the
near-surface layers of a planetary surface. In the case of Mars, this characteristic has been
proposed as a useful method to study spatial variations of such attributes as the distribution of
subsurface volatiles (1, 2) and heat flow (3). Using the Planetary Image Cartography System
(PICS) software, we have completed a quantitative anal_is of the geometry of fresh impact
craters in the Hesperia Planum region of Mars (-23 - 38 S, 233 - 248ow), where a uniform
target material and optimum viewing geometry make possible an analysis of target effects over a
large geographic region. Because of the morphologic similarity to the lunar maria, it is likely.that
Hesperia Planum comprises a series of flood lavas that partially infdled topo.graphic depressions
within the Martian highlands (4). Measurements of partially buried crater nms suggest that the
lava flows within Hesperia Planum are between 200 - 400 meters thick (5).
NEW DEFI'H/DIAMETER DATA
Our measurements were made from Viking Orbiter images obtained during orbits 417S,
418S and 419S. These images provide an unusually constant range of illumination geometries
(incidence angle = 63 - 80 °) and spatial resolution (88 - 96 m/pixel) for a large geographic area,
and to our knowledge represent the best data set that exists for the quantitative analysis of the
ridged plains materials on Mars. Our sample contains 109 craters in the diameter range 1.97 -
15.44 kin. Of these craters, 57 are morphologically very fresh, possessing complete rims, well
preserved ejecta blankets that completely surround the parent crater, have radial striations upon
the ejecta lobes or sharp distal ramparts, and have no superposed impact craters (cf. 6, 7).
Our depth/Diameter (d/D) data for the 57 pristine Hesperia Planum craters show
significant scatter, but particularly for craters < 5 km dia. (Fig. la, lb). We have divided the
population of fresh craters into two categories, denoted as "shallow" and "deep" (the "deep"
craters are -50 - 100% deeper than the "shallow" fresh craters of the same diameter). Once the
spatial distribution of these craters is investigated (Fig. lc), differences in their distribution is
evident. "Deep" craters axe in general located in the northern part of I-Iesperia Planum that is
investigated here, while "shallow" craters are more often found in the south. Clearly this
relationship is non-exclusive, but the spatial pattern does appear to be a real phenomenon.
We have also investigated the possible existence of the equivalent "deep" and "shallow"
craters for fresh craters larger than 6 km in diameter. From the d/D plot (Fig. lb), it is less clear
that two populations of large craters exist, so that we have separated craters that have the highest
d/D ratios from those with lower ratios but the same diameter. When the spatial distribution of
these two classes of large craters are compared (Fig. ld), we see a less uniform trend. From
Fig. ld, we therefore infer that whatever property of Hesperia Planum is responsible for
affecting craters smaller than 6 km diameter, this attribute does not affect larger craters.
DISCUSSION
Two mechanisms may explain this poleward shallowing of fresh craters: 1) spatial
variations in the thickness of the lava flows that comprise the surface material of Hesperia
Planum, and 2) latitudinal variations in the amount of volatiles in the near-surface (top 500 m?)
of the target. The thickness of the Hesperia Planum lava flows is believed to vary from more
than 400 m in the northwest of our study area to less than 200 m to the southeast (5), and the
underlying materials are most likely to be heavily cratered terrain similar to the materials that
outcrop around the perimeter of Hesperia Planum (4). In order for target properties to affect the
geometry of the crater, we assume that the crater cavity must be excavated to the depth of the
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interface between the lava flows and the underlying materials. The alternative method for
modifying small crater geometry in Hesperia Planum is one that involves a spatially-variable
concentration of volatiles in the top few hundred meters of the target at the time that the craters
were formed. This idea is not new, indeed it was the preferred mechanism proposed by Cintala
and Mouginis-Mark (8) to explain the variations in crater depth/diameter ratios. Fanale et al. (9)
have modeled the situation where volatiles were driven towards the poles over Martian history,
and the trends revealed for the small craters (Fig. lc) would support this poleward migration
model. Our current analysis of the Hesperia Planum craters may also show that for craters < 5
km in diameter, there is a gradual transition in the amount of volatiles present within the target.
In particular, if a spatial variation in the amount of volatiles were the controlling factor in crater
geometry, then we would expect to see a relationship between the crater's latitude and d/D.
A further test of the latitudinal variation in d/D ratio is to consider the geometry of other
craters in Hesperia Planum that have well preserved rim crests but lack the continuous ejecta
blankets that were a feature of the pristine craters. We have also measured the depth and
diameter of 48 additional craters smaller than 6.0 km diameter that fall into this category, and
have separated these craters into those that have high and low d/D ratios, using the same
subdivisions that were applied to the pristine craters with ejecta blankets. When the spatial
distribution of these craters without ejecta blankets is investigated, the same latitudinal variation
can be seen as for the pristine craters: craters with large d/D values are typically equatorward of
30os, while low d/D craters are found at all latitudes. From inspection of De Hon's (5)
thickness map, the ridged plains at certain latitudes in this area are interpreted to vary from < 200
m in the east to > 400 m in the west. Indeed, it is apparent that all but two of the 24 craters with
a large d/D ratio are equatorward of 30°S, but are formed in parts of Hesperia Planum that have
estimated thicknesses between < 200 to > 400 m. Thus we feel that this east-west variation in
plains thickness (which is probably the entire range of plains thicknesses within our study area)
was not the primary physical characteristic of the target that caused some craters to be unusually
deep or shallow. We therefore believe that the spatial variation in the amount of volatiles at
shallow depth is the most reasonable explanation for the observed spatial distribution in crater
geometries.
CONCLUSIONS
Our analysis of the d/D information has revealed that craters < 6.0 km dia. were most
likely affected by a latitudinally variable layer of volatiles. Craters at higher latitudes are
shallower than craters of the same diameter that are closer to the equator. This relationship
appears to be true both for craters with well preserved ejecta blankets and for deeper craters that
lack ejecta but have well preserved rims. While we are unable to quantify the change in the
amount or depth of burial of the volatiles that we propose to have been responsible for the
shallowing of these craters, our observations add to the current body of evidence that volatiles
existed within the shallow Martian crust. Furthermore, these volatiles must have been present
quite recently because they affected some of the freshest impact craters, which we infer to be
quite young (<1 byrs?) due to the lack of erosion of their ejecta blankets and rim crests.
The Viking Orbiter data that we used are particularly well suited to this form of digital
analysis, due to their uniform spatial resolution, large area of coverage of a single geologic unit,
and the small range of illumination geometries; the use of the digital Viking Orbiter data for other
areas of Mars may thus permit the qualitative inter-comparison of volatile concentrations based
on the geometry of fresh meteorite craters. The depth of the largest crater with a high d/D ratio in
any area may define the maximum thickness of the volatile-depleted layer during the time interval
over which the craters formed. In addition, once the Mars Observer Laser Altimeter (MOLA)
topographic data become available, MOLA measurements may also be useful for investigating the
former distribution of volatiles on Mars.
This research was funded by NASA Planetary Geology Grant NAGW-437.
371
REFERENCES
1) Mouginis-Mark,P. (1979). J. Geophys. Res., 84, 8011 - 8022. 2) Barlow, N. & T.
Bradley (1990). Icarus, 87, 156 - 179. 3) Boyce, J. M. (1979). Rpts. Planet. Geol. Prog.
1978 - 1979, NASA TM-80339, 114 - 118. 4) Greeley, R. & J. Guest (1987). U.S.
Geological Survey, Miscellaneous Map 1-1802-B. 5) De Hon, R. A. (1985). Rpt. Plan.
Geol. Prog., 1984, NASA TM-87563, 242 - 244. 6) Pike, R. J. (1980). Proc. Lunar Planet.
Sci. Conf. 11th, 2159 - 2189. 7) Mouginis-Mark, P. (1981). Icarus, 45, 60 - 76. 8)
Cintala, M. & P. Mouginis-Mark (1980). Geophys. Res. Lttrs., 7, 329 - 332. 9) Fanale, F.
P., J. R. Salvail, A. P. Zent & S. E. Postawko (1986). Icarus, 67, 1 - 18.
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Fig. 1 A) d/D plot showing craters smaller than 6.0 km dia. Crater population has been
subdivided into two classes of craters that appear to be "deep" or "shallow" for each diameter
B) d/D plot of pristine craters larger than 6.0 km in dia. Note-that the scales used here are line_
to expand the scatter of data points. Craters with the largest d/D ratio are called "deep" craters,
while smaller d/D values are associated with "shallow" craters. (2) Spatial distribution of the
small fresh craters identified in Fig. l a. Note that shallow craters are preferentially located
poleward of 30°S, while "deep" craters are predominately equatorward of this latitude. D)
Spatial distribution of deep and shallow fresh craters in the diameter range 6.0 - 15.44 km, as
identified in Fig. lb. Note that there does not appear to be any systematic pattern to the
distribution of these two classes of crater, suggesting that the excavation depth exceeds the
thickness of the near-surface volatile layer.
372


Depth/diameter relationships of fresh craters within Hesperia Planum, Mars

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