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Central Peak Crater


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DefinitionComplex crater with a single central uplift, a tight cluster of peaks, or a tightly spaced ring-like arrangement of peaks (e.g., Baker et al. 2011).CategoryA type of complex crater.DescriptionThe central peak is the simplest interior feature of complex craters. Many central peak craters have scalloped rims, terraced inner walls, and hummocky floors, on both rocky and icy bodies. These are inferred to represent failure by slumping and mass wasting of materials onto the floor (Greeley et al. 2000). The central peak itself can be a simple peak at or near the center of the crater floor, or can be composed of multiple uplift segments.MorphometryCentral peak diameter and height increase proportionally with crater rim crest diameter (Hale and Head 1979 and references therein). The top of the central peak is generally below the rim and the surrounding terrain (Öhman 2009 and references therein) (Fig. 1), although central peaks in the largest craters can reach and exceed the surroundi ...
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Central Peak Basin
Central Peak Crater
Veronica J. Bray
, Teemu O
Henrik Hargitai
Lunar and Planetary Laboratory, University of
Arizona, Tucson, AZ, USA
Arctic Planetary Science Institute, Rovaniemi,
NASA Ames Research Center / NPP,
Moffett Field, CA, USA
Complex crater with a single central uplift, a tight
cluster of peaks, or a tightly spaced ring-like
arrangement of peaks (e.g., Baker et al. 2011).
A type of complex crater.
The central peak is the simplest interior feature of
complex craters. Many central peak craters have
scalloped rims, terraced inner walls, and hum-
mocky floors, on both rocky and icy bodies.
These are inferred to represent failure by
slumping and mass wasting of materials onto
the floor (Greeley et al. 2000). The central peak
itself can be a simple peak at or near the center of
the crater floor, or can be composed of multiple
uplift segments.
Central Peak Crater 249
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Central peak diameter and height increase pro-
portionally with crater rim crest diameter (Hale
and Head 1979 and references therein). The top
of the central peak is generally below the rim and
the surrounding terrain (O
¨hman 2009 and refer-
ences therein) (Fig. 1), although central peaks in
the largest craters can reach and exceed the sur-
rounding terrain level (e.g., Pythagoras, a 144 km
diameter crater, Fig. 2). On some icy bodies, peak
height can exceed rim height (e.g., ‘Gula’ on
Central Peak Crater, Fig. 1 Morphometric parameters
and features of fresh central peak craters (Melosh 1989;
Turtle et al. 2005; Bray et al. 2008). Drim-to-rim diame-
ter, measured from rim crest to rim crest, D
diameter of
the flat inner floor, W
the width of the terraced zone, D
central peak diameter “measured at the contact with the
surrounding crater floor (i.e., measured at the top of the
allochthonous crater-fill breccias and melt rocks)” (Turtle
et al. 2005). Replaced by central uplift diameter for eroded
craters. dcrater depth measured from the maximum rim
elevation to the lowest point on the crater floor, h
height: the height of the crater rim above the average
surrounding terrain level, h
central peak height: maxi-
mum elevation of the central peak summit above the crater
floor. (1) Pre-impact terrain, (2) crater wall, terraced zone
if terraced, (3) central peak, (4) flat crater floor, (5) near-
rim ejecta blanket, (6) crater rim crest
Central Peak Crater,
Fig. 2 The 144 km
diameter Pythagoras crater
on the Moon. A SW to NE
topographic profile was
taken from the Global
Lunar Digital Terrain
Model (100 m resolution,
Scholten et al. 2012)to
show the large central peak
that rises higher than the
surrounding ground level.
Lunar Reconnaissance
Orbiter (NASA/GSFC/
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Ganymede). At lunar crater, diameters of approx-
imately 80 km central peak heights decrease. This
is also reflected in decreasing volume measure-
ments of central peaks (e.g., Hale and Grieve
1982; Bray et al., 2012) and is thought to mark
the beginning of the transition in crater morphol-
ogy from central peaks to peak-rings. Central
peak morphometric parameters include peak
diameter, peak height, peak area, volume, trend
(direction or elongation), offset (of the largest
feature of the peak from the crater center), and
azimuth (direction of offset) (Allen 1975).
(1) Central peaks may be single massifs, ridges,
or various types of clusters of peaks (Beer
and M
adler 1837, 130}77). Classification of
the central peaks by Hale and Head (1979)
included simple and complex types and arcu-
ate, symmetric, and linear types. The geom-
etry or the morphologic complexity of the
peak does not appear to be related to the
size of the crater (O
¨hman 2009 and refer-
ences therein) (Fig. 3).
(2) On low gravity bodies, broad central massifs
are produced in small bowl-shaped craters/
basins (complex crater, low gravity)
(Schenk et al. 2012) Fig. 4. These are not
created via the same mechanism as typical
central peaks.
Central Peak Crater, Fig. 3 Central peak craters. (a)
75 km diameter King crater, Moon displaying a lobster-
claw-like central peak (El-Baz 1978) 5.0N 120.5E.
LROC WAC mosaics. (NASA/GSFC/ASU); (b)93km
diameter Icarus crater, Moon. It has an unusually high
central peak (Allen 1975), almost reaching the crater rim
(Scholten et al. 2012) 5.3S 173.2W. LO-I-033 M
(NASA); (c) 36 km diameter terraced crater on Mercury
at 54.33S 311.2E. MESSENGER Image ID: 699956
(NASA/Johns Hopkins University Applied Physics Labo-
ratory/Carnegie Institution of Washington); (d)37km
diameter Saskia crater, Venus. Magellan radar mosaic
P36711 (NASA/JPL)
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The central peak is a result of uplift of the mate-
rial that originally underlay the transient cavity’s
central region (Kuiper 1954) during the modifi-
cation stage of impact crater formation. Uplift is
associated with the response of the target material
to the unloading by the rarefaction wave, as well
as with the converging of inwards and down-
wards collapsing material from the transient cra-
ter walls (Fig. 4). This brings subsurface rocks to
the surface (e.g., Grieve and Pilkington 1996).
This is in contrast to the central mounds of
small simple craters which are made up of debris
slumped in from the crater walls.
The existence of central peaks in impact cra-
ters prompted discussion of how the surface of
a planet or moon can act in a fluid-like manner
during impact crater formation. It is thought that
a transient weakening mechanism is required,
which allows the fast flow of the crater material
to form a central peak, followed by the return of
the normal material strength, which then main-
tains the uplift as part of the final crater morphol-
ogy (e.g., Melosh 1989).
Surface Units
(1) Central peak.
(2) Flat floor (annular basin, ring depression;
circular trough; rim syncline; annular
trough): the ring plains between crater rim
and central peak. The central peak is
surrounded by an annulus of fragment-
containing impact-melt sheet. Breccias com-
prised of different rock types (polymict brec-
cias) from various locations are found at the
base of the annular plain, over the fractured
bedrock. The central uplift exposes originally
deep-seated, highly disturbed shocked rocks.
(3) The rim area is structurally displaced, ter-
raced, and pervasively fractured (crater
rim). The underlying bedrock of the terrace
is covered by impact melt and ejecta; most of
the latter are highly mixed and moderately
shocked. Bodies of ejected melt tend to pool
in surface depressions on top of the breccias
¨rz et al. 1991).
Prominent Examples
Type example: Tycho, Moon. On the icy moons:
Melanthius, Tethys; Herschel, Mimas (Fig. 5).
Regional variations (see also: complex
Craters with central peaks appear in differ-
ently sized craters on differently sized planets.
The size of transition depends on local gravita-
tional acceleration and target material character-
istics (Pike 1980).
Moon: Central peaks are found with increas-
ing frequency in craters between 17 and 35 km
diameter on the Moon. Lunar craters from 35 to
about 170 km in diameter possess a central peak.
Central Peak Crater, Fig. 4 Central peak crater forma-
tion (After Fig. 3.10 from French (1998), modified)
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The height of the peak reaches its maximum in
80 km diameter craters (Hale and Head 1979
and references therein). Craters smaller than
40 km in diameter in maria have a higher abun-
dance of central peaks than in highlands (O
2009 and references therein).
Mars: Central peak craters occur between 5.0
and 139.8 km (Barlow 2010). The peak basal
diameters are 40–50 % of the crater diameter in
fresh craters. Central peak transition occurs at
5.2 km in the southern highlands of Mars, at
8.4 km in the northern plains, and at 11 km at
high latitudes poleward of 40in the zone of
water-saturated surface materials. There are
only two examples found in polar terrain. Central
peak may not form in these zones until signifi-
cantly larger diameters are reached, or it does
form but quickly collapses (Robbins and Hynek
2012). Hydrothermal alteration may have
occurred in several Martian impact craters’ cen-
tral uplift materials and may have produced
phyllosilicate assemblages there due to the heat
released by the impact and the fragmentation that
occurred during the uplift process (Barnhart
et al. 2010).
Icy moons: Transition from simple bowl-
shaped craters to central peak craters occurs at
smaller diameters on the icy satellites than on
rocky bodies of similar gravity (Schenk 1989).
Below crater diameters of 12 km, central peak
craters on Ganymede and simple craters on the
Moon have similar rim heights, indicating com-
parable amounts of rim collapse. This suggests
that the formation of central peaks at smaller
crater diameters on Ganymede than the Moon
(1.9 km on Ganymede compared to 25 km on
the Moon) is dominated by enhanced central floor
uplift rather than rim collapse. Central peak
diameters on Ganymede are typically one-third
of the crater diameter (Bray et al. 2008). On
Callisto, central peak craters range from about
5 to 40 km in diameter (Greeley et al. 2000).
Unlike on the Jovian moons, central peaks on
icy Saturnian satellites often exceed the elevation
of the background terrain and rise above the cra-
ter rim. These unusually high peaks are proposed
to be the products of concentrated floor uplift due
to high central temperatures (Dombard
et al. 2007).
Low gravity (0.2–0.3 m/s
) midsize icy satel-
lites and Vesta: see complex crater, low
History of Investigation
Central peaks were named centralgebirge (central
mountain) by Schro
¨ter (1791). Beer and M
Central Peak Crater, Fig. 5 130 km diameter Herschel Crater on the 381 415 km diameter Mimas at 0N, 245E, in
orthographic and perspective views. PIA12739, Cassini (NASA/JPL/Space Science Institute)
Central Peak Crater 253
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described “centralberge” as structures that “lie
often exactly in the middle of walled plains,
ring mountains, with no connection to the wall”
(Beer and M
adler 1837, 129}76). They distin-
guished the following peak classes (Fig. 6):
(1) centralketten ([linear] central chains) (e.g.,
Humboldt) (Beer and M
adler 1837, 130}77),
(2) centrale massengebirge (central mass moun-
tains: steep, with multiple peaks. When seen at
terminator, they look as lightspots in the dark
areas) (e.g., Theophilus, Petavius), (3) einzelne
centralberge (individual central mountains) (e.g.,
Copernicus, Gassendi), and (4) centrale piks
(central peaks: a single individual, sharp-pointed
peak (e.g., Alphonsus).
Kuiper (1954) noted that “among the most
puzzling lunar phenomena are the central peaks
found in a fraction of the craters. Some writers
have regarded them as volcanoes which depos-
ited the crater walls, thought to consist of ashes;
others, as a rebound action of the soil after the
impact or even as the “stuck” impactor itself.”
Kuiper suggested a rebound origin, during which
“the impact caused local melting or fissures
which tapped a large reservoir of lava farther
down” and concluded that “the central peaks
might consist of extrusive, igneous rock.” In
Kuiper’s model, central peak formation depends
on the presence of a lava reservoir available
before impacts; therefore, central peak could
only be formed in the so-called crust-melting
era, i.e., “shortly before, during, and shortly
after the formation of the maria” but not in the
“pre- and post-melting eras.” “It is not unlikely
that in marginal cases the heat of impact contrib-
uted to the availability of a lava reservoir; but the
relations stated on the presence and absence of
peaks indicate that the impacts were not the prin-
cipal cause of the lava.”
In modern studies, the use of the term
“rebound” is not recommended, as it implies
that elastic forces are the main cause of the uplift.
Although a complete mechanical understanding
of the central peak formation is still lacking,
elastic rebound is known to have a very minor
effect, except perhaps in some special circum-
stances (Melosh and Ivanov 1999). “Uplift” of
the sub-surface material is preferred to
Central Peak Crater,
Fig. 6 Variations of
central structure
morphology (Beer and
adler 1837). (a) Central
peak chain: Humboldt
(chain length: 100 km,
27.2S 80.9E), (b) central
complex mountain:
Theophilus (complex
diameter: 35 km, 11.4S
26.4E), (c) central
individual peaks:
Copernicus (peak cluster
diameter: 25 km, 9.7N
20.0W), (d) central peak:
Alphonsus (peak diameter:
8 km; 13.4S 2.8W).
LROC WAC mosaics
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The progressive development of central peaks
to peak rings was identified by Hartmann and
Wood (1971) when lunar far side images became
available, which showed different unflooded
basin morphologies (Fig. 7).
For Venus, analysis (Herrick and Sharpton
2000) implied that many central peaks of Venu-
sian impact craters are perhaps unusually high:
they were thought to reach the elevation of the
surrounding terrain and sometimes to surpass it.
For some cases, this may be true, but later studies
have shown that the most prominent feature of
the Venusian central peak elevation data is large
scatter (Herrick and Rumpf 2011). New, higher-
resolution topographic dataset of Venus is
required to fully resolve this question.
Related Terms
Landforms derived from different processes:
Central Dome Crater
Central Peak Crater, Fig. 7 Sequence of crater forms
from the appearance to the disappearance of central peak.
Central peaks show progressively more complex and
extended forms with increasing diameters. Generalized
profiles are also shown (Hartmann and Wood 1971). (a)
No central peak (Censorinus, D =4 km, 0.4S 32.7E), (b)
central peak (Lansberg, D =40 km, 0.3S 26.6W), (c)
central peak complex (Bullialdus, D =59 km, 20.7S
22.2W), (d) ringed peak cluster (Gassendi, D =
110 km, 17.5S 39.9W), (e) protobasin (Compton, D =
162 km, 55.3N 103.8E), (f) no central peak, peak ring
¨dinger, D =312 km, 75.0S 132.4E). LROC WAC
mosaics (NASA/GSFC/ASU)
Central Peak Crater 255
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Central Mound Crater
Landforms derived from the impact processes
with progressively increasing diameter and rela-
tively larger rings on rocky bodies:
Central peak crater (from craters with a single
central peak to craters with ring-like central peak)
(ringed peak-cluster basin and protobasin)
peak-ring crater, multiring basin.
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The high level of endogenic geological activity makes the terrestrial record of impact difficult to read. In their largely uneroded states, terrestrial impact structures have the basic so-called simple and complex forms observed on other planetary bodies, but few of them have morphometric parameters, such as apparent and true depth and stratigraphic uplift, that can be defined. Erosion severely affects such parameters, and can even result in a positive topographic form due to differential erosion. The principal criterion for the recognition of terrestrial impact structures is, therefore, not their form, but the occurrence of shock-metamorphic effects. In addition to a characteristic geological signature, terrestrial impact structures have characteristic geophysical signatures. The most common is a Bouguer gravity low, which extends out to the rim. The magnetic signature can be more varied but generally corresponds to a subdued low. The geophysical, geological, and morphological characteristics at terrestrial impact structures are summarised in tabular form as an aid to the recognition of additional structures.
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New analysis of central pit, central peak, and elliptical craters finds that pit craters are concentrated on volcanic units, pit and peak diameters display a linear relation with crater diameter, and elliptical craters show no obvious orientation changes with time.
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The detailed morphology of impact craters is now believed to be mainly caused by the collapse of a geometrically simple, bowl-shaped "transient crater." The transient crater forms immediately after the impact. In small craters, those less than approximately 15 km diameter on the Moon, the steepest part of the rim collapses into the crater bowl to produce a lens of broken rock in an otherwise unmodified transient crater. Such craters are called "simple" and have a depth-to-diameter ratio near 1:5. Large craters collapse more spectacularly, giving rise to central peaks, wall terraces, and internal rings in still larger craters. These are called "complex" craters. The transition between simple and complex craters depends on 1/g, suggesting that the collapse occurs when a strength threshold is exceeded. The apparent strength, however, is very low: only a few bars, and with little or no internal friction. This behavior requires a mechanism for tem-porary strength degradation in the rocks surrounding the impact site. Several models for this process, including acoustic fluidization and shock weakening, have been considered by recent investigations. Acoustic fluidization, in partic-ular, appears to produce results in good agreement with observations, although better understanding is still needed.
We have generated a new, 384,343-entry global crater database of Mars, statistically complete for craters with diameters D ≥ 1 km. In this release, the database contains detailed morphologic and morphometric data for craters D ≥ 3 km (future releases will extend these to smaller diameters). With detailed topographic data for the largest crater database to-date, we analyzed crater depth-to-diameter ratios for simple and complex morphologies across various terrains and for the planet as a whole and investigated the simple-to-complex morphology transition. Our results are similar to those in the published literature, but we found a substantial terrain dependence of the simple-to-complex transition that occurs at ˜11-km-diameter craters at high latitudes. This suggests a model that requires melting of volatiles during high-latitude crater formation that fill the crater during the modification phase but will still support the simple morphology to larger diameters. We also use this database to reexamine previously observed distributions and patterns to show its fidelity and to further explore other global relationships of fresh craters, those with central peaks, pits, and summit pits. We present the global distribution of craters with different types of ejecta and morphometric properties. Overall, this database is shown to be comparable to previous databases where there is overlap and to be useful in extending prior work into new regimes.
We derived near-global lunar topography from stereo image data acquired by the Wide-angle Camera (WAC) of the Lunar Reconnaissance Orbiter Camera (LROC) system. From polar orbit tracks, the LROC WAC provides image data with a mean ground resolution at nadir of 75 m/pixel with substantial cross-track stereo overlap. WAC stereo images from the one-year nominal mission and the first months of the science mission phase are combined to produce a near-global digital terrain model (DTM) with a pixel spacing of 100 m, the Global Lunar DTM 100 m, or “GLD100.” It covers 79°S to 79°N latitudes, 98.2% of the entire lunar surface. We compare the GLD100 with results from previous stereo and altimetry-based products, particularly with the Lunar Orbiter Laser Altimeter (LOLA) altimetry, which is the current topographic reference for the Moon. We describe typical characteristics of the GLD100 and, based upon the comparison to the LOLA data set, assess its vertical and lateral resolution and accuracy. We conclude that the introduced first version of the stereo-based GLD100 is a valuable topographic representation of the lunar surface, complementary to the LOLA altimetry data set. Further improvements can be expected from continuative investigations.
We present three investigations that use the Venusian impact crater population to constrain the planet's resurfacing history. We evaluate stereo-derived topography for 91 Venusian craters that have a diameter (D) greater than 15 km. Craters with radar-bright floors have greater rim-floor depths and rim heights than craters with radar-dark floors. For the bright-floored craters rim-floor depths are dbf = 0.483 D0.165 and rim heights are rhbf = 0.056 D0.483. Trends for dark-floored craters are ddf = 0.424 D0.108 and rhdf = 0.181 D−0.025. For a 60 km crater, this represents differences of 290 m in rim-floor depth and 240 m in rim height. We interpret these results to indicate that dark-floored craters have experienced postimpact volcanic embayment and filling. We examine the population of craters with D > 20 km that have radar-dark halos surrounding their continuous ejecta (114 craters). We find that a portion of the halo has been removed for almost all dark-floored craters, consistent with our interpretation that dark-floored craters have been affected by postimpact volcanism. Finally, we assessed geologic histories of 12 large impact structures with stereo coverage. All but one of these structures has experienced postimpact volcanism or tectonic deformation, often in multiple episodes. In summary, widespread volcanic and tectonic activity occurred throughout the time period of emplacement of the crater population. Postulated resurfacing histories that consider the majority of craters to be at the top of the stratigraphic column are invalid, and the mean surface age of Venus is young (∼150 My).
Using radargrammetry we have created high-resolution topographic maps of 74 Venusian craters, including all bright-floored craters over 12 km in diameter covered by Magellan stereo imagery. Our trend for rim-floor depths RF as a function of diameter D for bright-floored craters in the volcanic plains is RF = (0.345±0.05) D0.235±0.05, and for dark-floored craters in the plains, RF = (0.5±0.1) D0.0035±0.07. Rim heights RH for bright-floored craters in the plains are RH = (0.06±0.03) D0.4±0.2 (D>15 km), and for dark-floored craters in the plains, RH = (0.027±0.015) D0.4±0.2. These trends indicate that bright-floored craters 30 km in diameter are, on average, deeper than dark-floored craters by 180 m from rim to floor and have a 140 m higher rim, and at 90 km in diameter they are 380 m deeper from rim to floor and have a 220 m higher rim. The bright- and dark-floored populations are different at a 99% confidence level for both rim-floor depths and rim heights. The interpretation most consistent with our data and previous work by others is that Venusian craters with radar-dark floors have been partially filled and had their ejecta blankets embayed by regional-scale lava flooding. When the topography is examined in conjunction with the imagery, it is clear that many dark-floored craters have been surrounded by lavas that rose nearly to the crater rim even though a substantial portion of the crater's ejecta blanket was retained. Most previous analyses of the resurfacing history of Venus have relied on past interpretations that only a small percentage of Venusian craters are embayed by exterior volcanism. Because most craters on Venus have dark floors, our data indicate that the majority of Venusian craters have been surrounded and partially filled by postimpact lavas, and consequently, those previous analyses may have significantly underestimated the amount of volcanism on the Venusian surface over the past few hundred million years. Rim-floor depths for Venusian craters are consistent with the inverse gravity trend observed for the terrestrial planets, and they are ∼50% deeper than current estimates for complex craters on the Earth. Unlike the other terrestrial planets, neither terrain-floor depths nor central structure heights increase with increasing crater diameter. An interesting trend for which we have no explanation is that on Venus, the Moon, Mars, and Ganymede, central peaks generally rise to within a constant elevation relative to the surrounding terrain, but that elevation is lower on the Moon and Mars than on Venus and Ganymede.
Several of the icy satellites of Saturn (e.g., Dione, Mimas, Tethys) possess large craters (greater than 100 km in diameter) with very prominent central peaks. These peaks often exceed the elevation of the background terrain and rise above the crater rim. It has been thought that these peaks denote something peculiar associated with large crater formation on mid-size icy satellites, but we show that these peaks are instead the product of post-impact relaxation of the crater topography. Unlike previous relaxation studies, ours is the first to include remnant heat from the impact. We constrain the post-impact temperature field beneath a crater using hydrocode simulations. This temperature field is then used as input into a thermomechanical finite-element simulation, where we not only track the relaxing topography but also the diffusing thermal anomaly. Because the post-impact temperatures are highest in the central portion of the crater floor, a transient phase of rapid relaxation as the thermal anomaly dissipates results in concentrated floor-uplift, thereby explaining these ginormous central peaks. The background heat flow needs to be high enough to accommodate this deformation, which allows estimation of the thermal state of these satellites when these craters formed. Similarly extreme peaks are not found on the icy Galilean satellites of Jupiter because central peaks in large craters are replaced by other complex crater morphologies such as central pits.
Abstract— We examine the morphology of central peak craters on the Moon and Ganymede in order to investigate differences in the near-surface properties of these bodies. We have extracted topographic profiles across craters on Ganymede using Galileo images, and use these data to compile scaling trends. Comparisons between lunar and Ganymede craters show that crater depth, wall slope and amount of central uplift are all affected by material properties. We observe no major differences between similar-sized craters in the dark and bright terrain of Ganymede, suggesting that dark terrain does not contain enough silicate material to significantly increase the strength of the surface ice. Below crater diameters of ˜12 km, central peak craters on Ganymede and simple craters on the Moon have similar rim heights, indicating comparable amounts of rim collapse. This suggests that the formation of central peaks at smaller crater diameters on Ganymede than the Moon is dominated by enhanced central floor uplift rather than rim collapse. Crater wall slope trends are similar on the Moon and Ganymede, indicating that there is a similar trend in material weakening with increasing crater size, and possibly that the mechanism of weakening during impact is analogous in icy and rocky targets. We have run a suite of numerical models to simulate the formation of central peak craters on Ganymede and the Moon. Our modeling shows that the same styles of strength model can be applied to ice and rock, and that the strength model parameters do not differ significantly between materials.