Stratigraphic techniques

Stratigraphy is the basic principle used by geologists to interpret the geologic record. Stratigraphic analysis was first developed in 1669 by Nicolaus Steno, who realized that recent geologic events deposit their materials atop older, pre-existing layers of rock and soil. This is the principle of superposition, where the oldest layers are the most deeply buried and the younger units are near the surface (Figure 5.2). By analyzing rocks in these stratigraphic sequences, geologists can determine the types and relative timing of the geologic processes that created the layers.

Superposition is often combined with the principle of transection to constrain the relative timing of geologic events. Transection states that younger features cut older features. Transection is often seen with tectonic, fluvial, and glacial processes, which incise canyons or channels into pre-existing (and thus older) terrain. Later events can infill these features - it is not uncommon to see evidence of an ancient channel which has been filled in with later sedimentary or igneous deposits when looking at a canyon wall or a highway roadcut.

Superposition and transection are the two major stratigraphic techniques used on Earth. On Earth we have the additional advantage of being able to analyze rocks in our laboratories. Igneous rocks often contain small amounts of radioactive elements that can be used to date the time of solidification of the rock (Section 2.2.2). The age

Figure 5.2 Superposition is clearly seen in the layers exposed in Arizona's Grand Canyon. The older layers underlie younger layers. (Image by author.)

obtained by geochronologic techniques is the absolute age of the rock. If no datable rocks are available, geologists can develop a relative chronology from the strati-graphic positions of the rock layers. Absolute chronologies therefore provide actual dates for the time at which events occurred while relative chronologies only tell you whether something is older or younger than something else.

A third stratigraphic technique is commonly employed on other solid-surfaced bodies in our Solar System. Since datable samples are limited or non-existent from these bodies, planetary geologists use the number of craters as a key to determining relative chronologies and approximate absolute chronologies. Crater density is the cumulative number of craters per unit area, typically given as the number of craters greater than or equal to some diameter per 106km2. The probability that a surface has experienced an impact increases with the age of that surface since impact cratering has occurred throughout Solar System history. Thus older units will have higher crater densities than younger terrains.

5.1.3 Crater statistical analysis

Age information from crater analysis is obtained through use of crater size-frequency distributions (SFDs) (Crater Analysis Techniques Working Group, 1979). SFDs compare the frequency of craters as a function of crater diameter. An initial assumption of SFD analysis is that such distributions approximate a power-law function for an incremental distribution:

where N is the number of craters of diameter D and larger, K is a constant that depends on the crater density, and a is the slope of the power-law function (also

Cumulative SFD Plot

Relative SFD Plot

Highlands - Plains

Relative SFD Plot

Highlands - Plains

Highlands Plains

Highlands Plains

Figure 5.3 Crater size-frequency distributions (SFD) are studied using cumulative and relative plots. (a) The cumulative SFD plots for craters on the martian highlands and northern plains are shown. (b) Compare the relative SFD plots to those in (a) for the same regions. Obvious differences in the shapes of the curves are readily apparent in the relative plots.

called the population index). The cumulative SFD is an example of an incremental distribution. Figure 5.3a shows a sample cumulative plot, where the number of craters of a certain diameter and larger per unit area (C) is plotted against diameter on a log-log plot. Error bars are calculated from

SFDs for craters larger than 5 to 8 km in diameter often show an approximately straight line with a — 2 slope, especially for younger surfaces. The negative slope indicates that smaller craters are more abundant than larger craters. The values of SFD for different surface units vary vertically on this plot depending on crater density - higher crater densities lie above lower crater densities. Thus, the relative positions of SFDs from two different terrains indicate that the surface corresponding to the higher curve is older than the surface displaying the lower curve. The cumulative plot suggests that SFDs for all terrain units can be approximated by a power-law function with a slope of approximately — 2 at all diameters. However, this may be a reflection of the plotting technique because the cumulative nature of this technique tends to smooth out frequency variations within a particular diameter range.

The relative or R-plot technique overcomes the smoothing problem of cumulative plots. The R-plot is a differential plot of the form (Crater Analysis Techniques Working Group, 1979)

where dN is the number of craters with diameters in the range dD, D is the mean diameter of the range, C is a constant, and b is the differential population index. Because this equation is derived by differentiating the incremental form given in Eq. (5.1), b = a + 1. The R-plot differs from the cumulative plot in two ways: (1) it only includes the number of craters within a particular size range, thus frequency variations are not reduced as with the cumulative technique, and (2) it normalizes the curves to a common slope index of — 3 (corresponding to the cumulative slope index a = 2). The sizes of the diameter ranges (or bins), dD, are set at a/2 increments; hence dD may be 2km to 2 \/2km (= 2.8km), the next bin will be from 2.8 km to 2.8V2km (= 4.0km), etc. The relative technique plots the geometric mean diameter of the bin on the horizontal axis and the normalized parameter R on the vertical axis of a log-log plot. R is given by

where N is the number of craters with diameters in the range Da (lower diameter limit of bin) to Db (upper diameter limit, Db = Da a/2), D is the geometric mean of the diameter bin (D = ^jDaDb), and A is the area over which the craters are counted. SFDs following the —2 slope suggested from cumulative analysis will plot as a horizontal line on the R-plot, while SFDs following other slopes will plot as inclined lines. Error bars are calculated as

Figure 5.3b shows the R-plot of the same terrain units displayed in the adjacent cumulative plot. The lower curve (i.e., the younger terrain) approximates a horizontal line on the R-plot, corresponding to a power-law function of —2 cumulative slope. However, a single-sloped power-law function does not approximate the upper (older) curve. The upper curve is representative of the SFDs seen on heavily cratered regions of the Moon, Mercury, and Mars, while the lower curve is similar to that seen on the lunar maria and martian northern plains.

A third type of crater plot combines properties of the cumulative and relative plots. This is a log-incremental SFD, primarily used by W. Hartmann and colleagues (e.g., Hartmann, 2005). In this technique, the number of craters within a particular diameter bin per unit area is plotted on the ordinate while the diameter is plotted on the abscissa (Figure 5.4). This plot is similar to the R-plot except the results are not normalized to the —3 differential slope.

The cause of the multisloped highlands SFD curve is still debated. One suggestion is that this curve represents the SFD of terrains saturated with craters to o o

102 101 10° 10-1 10-2 10-3 10-4 10-5 10-6 10-7 10-8

Rbolide = 2.6, Schmidt-Housen scaling

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MOC R1001183 Detail MOC R1001183 THEMIS V02040007 Detail THEMIS V02040007 HRSC 2987 post surface HRSC 3042 post surface

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16 31 63 125 250 500 1 2 4 8 16 32 64 128 256 (m)(km)—

Diameter

16 31 63 125 250 500 1 2 4 8 16 32 64 128 256 (m)(km)—

Diameter

Figure 5.4 Isochron plots provide insights into the absolute ages of surface units based on crater density. This plot of craters in a region west of Tharsis indicates that the surface is approximately 100 Ma old. (Image courtesy of William Hartmann, Planetary Science Institute.)

(Hartmann, 1984, 1997). Saturation refers to surfaces which are covered with so many craters that any new impact erases an existing one, keeping the crater density approximately constant. No age information (other than "very old") can be obtained from saturated surfaces. While ancient surfaces may display crater saturation in some diameter ranges (particularly at the smaller crater diameters), other diameter ranges appear to be largely unaffected. Observational evidence against saturation is seen in places such as on the floor of the Moon's Orientale Basin and on the martian ridged plains. These terrains have lower crater density than the heavily cratered regions, but still display the multisloped curve shape (Barlow, 1988; Strom et al., 1992).

A second possible explanation for the multisloped curve of the highlands is that erosion has preferentially destroyed the smaller craters, causing a change in slope at the lower-crater-diameter end. While erosion definitely affects the SFDs, it does not appear to be the total explanation for the frequency downturn in the 5 to ~70 km diameter range. Heavily cratered regions on the Moon, Mercury, and Mars all show the same multisloped curve, but the three bodies have experienced very different obliteration histories. In addition, the SFDs of relatively fresh impact craters (those still retaining ejecta blankets; by definition, fresh craters have not experienced much erosion) on heavily cratered martian terrains also show the downturn at smaller diameters (Barlow, 1990).

The different SFDs could also indicate two crater-production populations with different size-frequency distributions. The multisloped curve is seen on surfaces believed to date from the Late Heavy Bombardment (LHB) period (Section 2.1.2) while the flatter curve is seen on younger terrains. The current impactor population in the inner Solar System is derived primarily from asteroids with some contribution from comets (Bottke et al., 2002) and the crater SFD is consistent with the asteroid SFD (Bottke et al., 2005). LHB impactors likely originated from the planetesimal disk disrupted by the migration of the outer planets (Gomes et al., 2005), although a source in the main asteroid belt also has been proposed (Strom et al., 2005). If the SFDs reflect two different impactor populations, they imply that the LHB population was enriched in mid-sized and large impactors and depleted in smaller objects compared to the present-day asteroid SFD. The theory of different impactor populations indicates that not only can the SFD curves tell us the crater densities and thus the relative ages of the different terrain units but also whether the craters formed during the LHB or post-Heavy Bombardment periods. Such an analysis applied to Mars suggests that ~60% of the planet's surface dates from LHB while the remaining 40% formed more recently (Barlow, 1988).

Crater statistical techniques provide the best chronologic information when the number of craters per unit area is greatest, thereby reducing the uncertainties. For craters larger than ~5-km diameter, this requires performing the crater analysis over relatively large areas to obtain statistically viable results. Geologic units are seldom uniform over such large areas, resulting in the age information derived from such analyses being averages of several incorporated units. Thus, planetary geologists prefer using craters <5-km diameter, which are more abundant than the larger craters and allow statistically significant results to be obtained for regions of small areal extent (e.g., Hartmann et al., 1999; Neukum et al., 2004).

However, chronologic information derived from small-crater counts is affected by two major processes. Erosion, as noted previously, preferentially destroys smaller craters, leading to younger age estimates on surfaces where erosion has dominated. The other problem is secondary crater contamination, which may be responsible for a steepening of the SFD at <1-km diameter. Secondary craters are small craters formed from material ejected during formation of a larger (primary) crater. Secondary clusters and those close to the primary crater ("obvious secondaries")

can be easily identified and avoided, but distant secondaries can be difficult to distinguish from small primary craters. Inclusion of secondaries produces a higher crater density and thus indicates an older age for the surface. Some researchers argue that secondaries are a minor contributor to the SFD since fragmentation and grinding of ejected debris eliminates material in the smallest diameter ranges (Werner et al., 2006). Hartmann (2005) argues that the SFD is a combination of primary and distant secondaries which approximates the primary production curve over time. Clustering analysis suggests that the majority of small craters on Europa are secondaries, a result which is probably applicable to other Solar System bodies (Bierhaus et al., 2005). The 10-km-diameter martian crater Zunil displays thermally distinct secondary crater rays extending over 300 crater radii away from the primary crater (McEwen et al., 2005). The large number and wide distribution of Zunil secondaries supports the assertion that most martian craters <1-km diameter are probable secondaries. However, secondary crater production is enhanced on young lava plains and occurs only in association with larger craters in regions with thicker regolith (Hartmann and Barlow, 2006). This observation combined with the fact that secondary craters are concentrated in non-uniformly distributed rays indicates that secondary contamination is not evenly distributed across the entire planet. The combined problems of erosion and secondary craters lead to considerable uncertainty in chronologic information derived from small-crater analysis (McEwen and Bierhaus, 2006).

Relative ages obtained from crater statistical analysis can be used to estimate absolute ages of surface units. Radiometric dating of lunar samples returned by the Apollo and Luna missions allows correlation of crystallization ages with the crater density of the region surrounding the sample site. The resulting crater density versus age plot is the lunar crater chronology (LCC) graph (Figure 2.2). The LCC can be used to estimate absolute ages of non-sampled regions of the Moon based on the crater density of those regions.

The LCC can be extrapolated to other Solar System bodies to obtain absolute age estimates provided the following three conditions are met: (1) the population of impacting bodies is the same between the Moon and body of interest; (2) the timing of the LHB between the Moon and other body is known; and (3) the impact flux for the body relative to the Moon is known. The similarity of the SFDs for the Moon and Mars indicates that condition 1 is met. Recent numerical modeling suggesting that the LHB resulted from disruption of a planetesimal disk by migration of the outer planets (Gomes et al., 2005) indicates that the beginning and end of LHB were approximately simultaneous throughout the inner Solar System, satisfying condition 2. The Mars/Moon impact rate ratio can be estimated from SFDs of the impacting objects and knowledge of impact probability for asteroids/comets crossing the orbits of the Moon and Mars (Ivanov, 2001). The ratio of the impact rate of projectiles of a certain size per unit area compared to that impact rate on the Moon is called the bolide ratio (Rb). The current Mars/Moon Rb value is 4.93, but because of orbital eccentricity variations experienced by Mars (Section 7.4.2), the value can be as low as 2.58 (Ivanov, 2006).

The LCC can be extrapolated to Mars by utilizing the Mars/Moon impact rate ratio and scaling relationships for converting impactor size to final crater diameter (Section 5.3.1). Hartmann (2005) has produced isochron plots for Mars which allow one to estimate the formation age of a surface from crater counts (Figure 5.4). Variations in ages associated with different diameter ranges provide insights into processes such as erosion and possible secondary crater contamination.

5.2 Martian geologic periods

Martian history is divided into periods based on stratigraphic relationships and occurrence of certain geologic processes (Figure 5.5) (Tanaka et al., 1992). The oldest period is the Noachian, named after the Noachis region in the southern highlands. Noachian terrains formed during the LHB and are thus very heavily cratered surfaces. The range of degradation associated with craters formed during this period suggests high erosion rates due to geologic processes such as rainfall and fluvial erosion (Craddock and Howard, 2002). The Noachian period is divided into early (prior to ~3.95 Ga ago), middle (between 3.95 and 3.8 Ga ago), and late (3.83.7 Ga ago) periods based on crater densities (Hartmann and Neukum, 2001). The Hesperian Period (named after Hesperia Planum) represents middle martian history and is characterized by volcanic extrusions creating the ridged plains and a decline in the high impact rates that occurred during the LHB. It is subdivided into early (3.7-3.6 Ga ago) and late (~3.6-3.0 Ga ago) periods. The Amazonian Period (from Amazonis Planitia) covers the planet's most recent history during which erosion rates have been low and volcanic activity has been concentrated in the Tharsis and Elysium regions. The Amazonian Period is divided into early (~3.0-1.8 Ga ago), middle (~ 1.8-0.5 Ga ago), and late (~0.5 Ga ago to present) periods.

5.3 Geologic processes

5.3.1 Impact cratering

Impact craters are the most common geologic features observed on most planetary surfaces. The term crater, Latin for cup, was introduced by Galileo in 1610 to describe the approximately circular depressions observed on the lunar surface. Most early scientists believed that the lunar craters were of volcanic origin, primarily because of their circular appearance and terrestrial experience with volcanic craters

Hesperian Noachian

Figure 5.5 Stratigraphie techniques provide insights into the ages of surface units. This map shows the distribution of units formed during the Noachian, Hesperian, and Amazonian periods. (Image courtesy of Trent Hare, USGS.)

Hesperian Noachian

Robinson Projection Center Longitude 0

4,000 I km

Figure 5.5 Stratigraphie techniques provide insights into the ages of surface units. This map shows the distribution of units formed during the Noachian, Hesperian, and Amazonian periods. (Image courtesy of Trent Hare, USGS.)

(calderas). Acceptance of the idea that craters could result from collisions of large chunks of space debris with a planetary surface was gained only after (1) laboratory experiments showed that circular craters could be produced during high-velocity (hypervelocity) impacts (Ives, 1919; Gault et al., 1968), (2) nuclear and chemical explosions provided insights into the physics of impact events (cf. Roddy et al., 1977), and (3) identification of shock metamorphic features in rocks surrounding Meteor Crater and Ries Crater (Shoemaker and Chao, 1962; Shoemaker, 1963).

Impact crater formation is divided into three stages (Gault et al., 1968; Melosh, 1989): contact/compression, excavation, and modification. The projectile first encounters the surface during the contact/compression stage, generating shock waves which propagate through both the target and the projectile. The Hugoniot equations, derived from conservation of mass, momentum, and energy, describe the material characteristics on either side of the shock front:

The subscript 0 refers to the material properties - pressure P, density q, specific (per unit mass) internal energy E, and specific volume V ( = 1/q) - before passage of the shock front, while those parameters without a subscript refer to the compressed material (after passage of the shock wave) (Figure 5.6). The shock wave moves with velocity U and, after the shock front has passed, the compressed material has a particle velocity of up. The reference frame when using the Hugoniot equations is the rest frame of the uncompressed material (up0 = 0).

The shock wave pressure declines exponentially with distance from the impact site (Figure 5.7) and the shock wave turns into a seismic (elastic) wave once the pressure drops below 1 to 2 GPa. As the shock wave encounters a free surface (either the surface of the target material or the back surface of the projectile), it is reflected as a rarefaction or release wave back into the material. As rarefaction waves pass through the highly shocked material, they unload some of the pressure, resulting in melting and vaporization of some of the material. The contact and compression stage ends when the rarefaction wave engulfs and destroys the projectile. The projectile usually has traveled a distance equivalent to its diameter during the contact/compression stage, corresponding to only a few seconds or less.

Surface

Shock wave

Figure 5.6 The passage of the impact-induced shock wave alters characteristics of the rock. Before the shock wave passes, the material is characterized by density (q0), pressure (P0), and specific internal energy (E0). The shock wave travels with velocity U. Once the shock front passes, the rock acquires new values of density (q), pressure (P), and specific internal energy (E), as well as a particle velocity (mp). Relationships between these before-and-after values are obtained from the Hugoniot equations.

Projectile 10 km/s

Ejecta

Figure 5.6 The passage of the impact-induced shock wave alters characteristics of the rock. Before the shock wave passes, the material is characterized by density (q0), pressure (P0), and specific internal energy (E0). The shock wave travels with velocity U. Once the shock front passes, the rock acquires new values of density (q), pressure (P), and specific internal energy (E), as well as a particle velocity (mp). Relationships between these before-and-after values are obtained from the Hugoniot equations.

Projectile 10 km/s

Ejecta

Shock metamorphism zone

Shock pressure contours

Shock metamorphism zone

Fracturing and brecciation

Shock pressure contours

Fracturing and brecciation

1 GPa

Figure 5.7 Shock-induced pressure declines with distance from the impact site. Near the impact site, pressures are high enough to induce melting and vaporization. Further from the impact, the passage of the shock wave produces fracturing of the rock. (Reprinted by permission from Lunar and Planetary Institute, French [1998], Copyright 1998.)

The excavation stage opens the crater. As the shock wave encounters the free surface, part of it is converted into kinetic energy and the rest is reflected as a rarefaction wave. The rarefaction waves will fracture the target material as long as the stress of the rarefaction wave (a tensional wave) exceeds the mechanical strength of the rock. The portion of the shock wave converted to kinetic energy will accelerate fragmented material outward. Some of this material is accelerated up and out

Projectile

Uplifted TC rim

Projectile

Uplifted TC rim

Figure 5.8 The transient cavity (TC) forms during the excavation stage, as the shock wave is converted into kinetic energy. Material can be either displaced downward, or up and out of the crater. The material ejecta outward forms the raised rim and ejecta blanket. (Reprinted by permission from Lunar and Planetary Institute, French [1998], Copyright 1998.)

Figure 5.8 The transient cavity (TC) forms during the excavation stage, as the shock wave is converted into kinetic energy. Material can be either displaced downward, or up and out of the crater. The material ejecta outward forms the raised rim and ejecta blanket. (Reprinted by permission from Lunar and Planetary Institute, French [1998], Copyright 1998.)

of the crater, forming the ejecta blanket and part of the uplifted rim (the rest of which results from structural uplift from the shock and rarefaction waves) (Figure 5.8). The rest of the material is displaced downward, resulting in a bowl-shaped cavity called the transient crater. The transient crater has the greatest depth (d) that the crater will ever have and is between a third and a quarter of the transient crater diameter (Dt).

Scaling relationships are used to extrapolate results from laboratory experiments to impact events (Melosh, 1989). The value of Dt is related to energy of the impact (W), size of the projectile (L), acceleration of gravity (g), densities of the projectile and target (qp and qt, respectively), and impact angle (measured from horizontal surface) (h) through

Dt = 1. 8q0-1V-33g-o-22La 13 W0 ' 22 (sin h)033 . (5 .9)

Smaller craters tend to be bowl-shaped depressions, approximating the shape and depth of the transient crater. These craters are called simple craters (Figure 5.9a). Depth (d) of a simple crater is related to its rim diameter (Dr) by d « y. (5 .10)

Larger craters display more complicated morphologies and are called complex craters (Figure 5.9b). Complex craters are shallower compared to their size than simple craters, typically about one-tenth of the rim diameter. The diameter where craters transition from simple to complex structures is proportional to g-1, although target characteristics also contribute. The simple-to-complex transition diameter

Figure 5.9 Impact craters display morphologic differences as diameter increases. (a) The smallest craters display a bowl-shaped appearance and are called simple craters. This 2-km-diameter simple crater on Mars is located near 34.4°N 241.2°E. (MOC image MOC2-1274, courtesy of NASA/MSSS.) (b) Larger craters, called complex craters, display more complicated morphologies, including central peak structures, shallower floors, and terraced walls. This complex crater is 27 km in diameter and located at 28.6°N 207.0°E. (THEMIS image V17916017, NASA/ ASU.)

Figure 5.9 Impact craters display morphologic differences as diameter increases. (a) The smallest craters display a bowl-shaped appearance and are called simple craters. This 2-km-diameter simple crater on Mars is located near 34.4°N 241.2°E. (MOC image MOC2-1274, courtesy of NASA/MSSS.) (b) Larger craters, called complex craters, display more complicated morphologies, including central peak structures, shallower floors, and terraced walls. This complex crater is 27 km in diameter and located at 28.6°N 207.0°E. (THEMIS image V17916017, NASA/ ASU.)

(DSC) on Mars is predicted to occur near 10 km, but observations place it closer to 8 km (Garvin and Frawley, 1998). The lower DSC probably results from substantial quantities of ice in the near-surface material reducing the strength of the target material. Variations in morphometric properties such a crater depth, volume, and rim height with latitude also are attributed to higher concentrations of subsurface ice near the poles (Garvin et al., 2000a). Table 5.1 provides morphometric relationships for fresh martian impact craters.

Material ejected during crater formation forms the ejecta blanket surrounding the crater. The ejecta blanket does not contain material from the entire cavity. The depth of excavation (dex) for the ejected material is approximately one-third of the transient crater depth:

Ejecta blankets are divided into continuous and discontinuous sections. The continuous ejecta blanket abuts the crater rim and is composed of a continuous blanket of debris typically extending between one and three crater radii from the rim. The discontinuous ejecta blanket extends beyond the continuous ejecta blanket in discrete streaks radial to the crater. Secondary craters populate most of the discontinuous ejecta blanket (Figure 5.10).

Table 5.1 Impact crater morphometric parameters as function of crater diameter (D)

Parameter

Simple craters

Complex craters

Central peak diameter (Dcp) —

Inner cavity wall slope (s) s = 28.40D

= 0.36D049

Figure 5.10 Ejecta blankets can be divided into continuous and discontinuous parts. This image of the ejecta blanket associated with a 28.3-km-diameter crater located at 23.2°N 207.8°E shows the layered continuous ejecta blanket extending outward from the rim. Beyond the edge of the layered ejecta blanket, small secondary craters are seen, which constitute the discontinuous ejecta blanket. (THEMIS image V01990003, NASA/ASU.)

Most fresh martian impact craters are surrounded by a layered ejecta blanket with one (single layer), two (double layer), or more than two (multiple layer) ejecta layers (Figure 5.11) (Barlow et al., 2000). This ejecta morphology is distinct from the radial morphology seen around craters on dry bodies like the Moon. Two models have been proposed to explain this layered ejecta morphology: (1) impact into and vaporization of subsurface volatiles (Carr et al., 1977; Stewart et al., 2001), and (2) interaction of the ejecta curtain with the martian atmosphere (Schultz, 1992; Barnouin-Jha et al., 1999a, b). While the atmosphere plays some role, most of the evidence suggests that subsurface volatiles are the dominant contributor in the formation of the layered ejecta morphologies (see review in Barlow, 2005).

Figure 5.11 Layered ejecta blankets are divided into single layer (SLE), double layer (DLE), and multiple layer ejecta (MLE) based on the observed number of ejecta layers. (a) SLE craters display one complete ejecta layer around the crater, like this 9.9-km-diameter crater located at 19.54°S 277.01°E. (THEMIS image V18738005.) (b) DLE craters, like this 12.0-km-diameter example at 55.59°N 268.68°E, display two complete ejecta layers. (THEMIS image I13646007.) (c) MLE craters display three or more partial or complete ejecta layers. This MLE crater is 29.4 km in diameter and is located at 23.15°S 281.35°E. (THEMIS image I06994003.) (NASA/ASU.)

Figure 5.11 Layered ejecta blankets are divided into single layer (SLE), double layer (DLE), and multiple layer ejecta (MLE) based on the observed number of ejecta layers. (a) SLE craters display one complete ejecta layer around the crater, like this 9.9-km-diameter crater located at 19.54°S 277.01°E. (THEMIS image V18738005.) (b) DLE craters, like this 12.0-km-diameter example at 55.59°N 268.68°E, display two complete ejecta layers. (THEMIS image I13646007.) (c) MLE craters display three or more partial or complete ejecta layers. This MLE crater is 29.4 km in diameter and is located at 23.15°S 281.35°E. (THEMIS image I06994003.) (NASA/ASU.)

Table 5.2 Diameter ranges of specific interior morphologies

Interior feature

Diameter range

Central peaks

—6-175 km

Central pits

5-60 km

Peak ring basins

—50-500 km

Multi-ring basins

>500 km

Source: Data from Barlow (1988).

Source: Data from Barlow (1988).

The excavation stage ends when the transient crater reaches its maximum size. For craters whose final shape and size are dictated by gravity rather than the strength of the target material, the time for transient crater formation is related to Dt and g (Melosh, 1989):

The modification stage extends from the end of transient crater formation until the crater is completely destroyed by subsequent geologic processes. During the modification stage, simple craters usually undergo some sliding of debris from the crater walls but most of the changes to simple craters result from infilling by other geologic processes such as volcanic, eolian, and fluvial activity. Complex craters display a wider variety of internal structures. As with ejecta structures, interior features tend to transition to different types as a function of crater size and location on the planet (Table 5.2). Martian complex craters often display a peak or pit in the center of the crater. Central peaks (Figure 5.12a) result from uplift of the floor after passage of the shock wave, with the uplift freezing into place to form the central peak. Larger craters can display a mountainous ring (peak ring) rather than a central peak complex (Melosh, 1989). The peak ring structure (Figure 5.12b) likely results from collapse of the central peak and formation of a "ripple" which freezes into place. The largest craters, such as Hellas and Argyre, may be multi-ring structures, although their outer rings are not obvious. Central pits are seen on Mars and on many of the icy moons in the outer Solar System. Martian central pits can occur either directly on the crater floor or on top of a central rise (Barlow, 2006) (Figure 5.12c). Central pits likely form from vaporization of ice under the central part of the crater floor, resulting from shock heating of the material and explosive release of the resulting gases (Wood etal., 1978; Pierazzo etal., 2005). The walls of complex craters typically have slopes greater than the angle of repose when the crater first forms, resulting in collapse of the walls to form terraces (Figure 5.12a).

Martian impact craters display a wide range of ejecta and interior morphologies, indicating complexities in target properties during crater formation and modification

Figure 5.12 Interior morphologies vary with crater size and location. (a) Central peaks, such as the one visible in this 30.1-km-diameter crater at 8.30°N 302.54°E, are common features in complex impact craters. The collapse of the crater walls to form wall terraces can also be observed. (THEMIS image V18886012, NASA/ ASU.) (b) Central peaks often expand into peak rings for larger craters. A peak ring is clearly visible in 224-km-diameter Lyot Crater, located at50.38°N 29.33°E. (Viking mosaic image, NASA/JPL.) (c) Central pits are common in impact craters on Mars and on icy moons of the outer Solar System. This floor pit occurs in a 27.9-km-diameter crater located at 4.30°N 294.05°E. (THEMIS image V17526014, NASA/ASU.)

Figure 5.12 Interior morphologies vary with crater size and location. (a) Central peaks, such as the one visible in this 30.1-km-diameter crater at 8.30°N 302.54°E, are common features in complex impact craters. The collapse of the crater walls to form wall terraces can also be observed. (THEMIS image V18886012, NASA/ ASU.) (b) Central peaks often expand into peak rings for larger craters. A peak ring is clearly visible in 224-km-diameter Lyot Crater, located at50.38°N 29.33°E. (Viking mosaic image, NASA/JPL.) (c) Central pits are common in impact craters on Mars and on icy moons of the outer Solar System. This floor pit occurs in a 27.9-km-diameter crater located at 4.30°N 294.05°E. (THEMIS image V17526014, NASA/ASU.)

Figure 5.13 Pedestal craters are perched above the surrounding terrain by removal of fine-grained material by eolian and/or sublimation processes. These two pedestal craters are each about 1.5 km in diameter and are located near 52.0°N 150.5°E. (THEMIS image V11541006, NASA/ASU.)

processes operating after formation. For example, many small craters at high latitudes are elevated above the surrounding terrain (Figure 5.13). These "pedestal craters" formed when surrounding material was removed either by eolian deflation or ice sublimation (Barlow, 2006). Erosional processes have removed the elevated rims and infilled the floors of many older craters, particularly those found on the ancient Noachian surfaces. Computer simulations of the topography produced by different erosional processes have been compared to the topographic profiles across eroded craters. The results suggest that rainfall was a primary agent of erosion during the Noachian period (Craddock and Howard, 2002). Crater degradation in the Hesperian and Amazonian periods has been dominated by eolian and volcanic activity, with only localized fluvial/glacial erosion (Bibring et al., 2006).

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