Investigation of Charon's Craters with Abrupt Terminus Ejecta, Comparisons with Other Icy Bodies, and Formation Implications

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Investigation of Charon's Craters With Abrupt Terminus Ejecta, Comparisons With Other Icy Bodies, and Formation Implications

Item Type Article

Authors Robbins, Stuart J.; Runyon, Kirby; Singer, Kelsi N.; Bray, Veronica J.; Beyer, Ross A.; Schenk, Paul; McKinnon, William B.; Grundy, William M.; Nimmo, Francis; Moore, Jeffrey M.; Spencer, John R.; White, Oliver L.; Binzel, Richard P.; Buie, Marc W.; Buratti, Bonnie J.; Cheng, Andrew F.; Linscott, Ivan R.; Reitsema, Harold J.; Reuter, Dennis C.; Showalter, Mark R.; Tyler, G. Len; Young, Leslie A.; Olkin, Catherine B.; Ennico, Kimberly S.; Weaver, Harold; Stern, S. Alan

Citation Investigation of Charon's Craters With Abrupt Terminus Ejecta, Comparisons With Other Icy Bodies, and Formation Implications 2018, 123 (1):20 Journal of Geophysical Research: Planets

DOI 10.1002/2017JE005287

Publisher AMER GEOPHYSICAL UNION

Journal Journal of Geophysical Research: Planets

Download date 02/10/2021 15:46:31

Item License http://rightsstatements.org/vocab/InC/1.0/

Version Final published version

Link to Item http://hdl.handle.net/10150/627125 PUBLICATIONS

Journal of Geophysical Research: Planets

RESEARCH ARTICLE Investigation of Charon’s Craters With Abrupt Terminus 10.1002/2017JE005287 Ejecta, Comparisons With Other Icy Bodies, Key Points: and Formation Implications • Report discovery of Charon abrupt termini ejecta craters Stuart J. Robbins1 , Kirby Runyon2 , Kelsi N. Singer1 , Veronica J. Bray3, Ross A. Beyer4,5 , • Describe properties of Charon abrupt 6 7 8 9 5 termini ejecta and their host craters in Paul Schenk , William B. McKinnon , William M. Grundy , Francis Nimmo , Jeffrey M. Moore , 1 5 10 1 11 context of other bodies John R. Spencer , Oliver L. White , Richard P. Binzel , Marc W. Buie , Bonnie J. Buratti , • Constrain formation mechanisms for Andrew F. Cheng2 , Ivan R. Linscott12, Harold J. Reitsema13 , Dennis C. Reuter14, these ejecta blankets Mark R. Showalter4 , G. Len Tyler12, Leslie A. Young1 , Catherine B. Olkin1 , Kimberly S. Ennico8 , Harold A. Weaver2 , and S. Alan Stern1 Supporting Information: • Supporting Information S1 1Southwest Research Institute, Boulder, CO, USA, 2Department of Earth and Planetary Sciences, The Johns Hopkins • Figure S1 University, Baltimore, MD, USA, 3Lunar and Planetary Laboratory, University of Arizona, Tucson, AZ, USA, 4Sagan Center, SETI • Table S1 Institute, Mountain View, CA, USA, 5Planetary Systems Branch, Space Sciences and Astrobiology Division, NASA Ames Research Center, Moffett Field, CA, USA, 6Lunar and Planetary Institute, Houston, TX, USA, 7Department of Earth and Correspondence to: 8 S. J. Robbins, Planetary Sciences at McDonnell Center for the Space Sciences, Washington University, St. Louis, MO, USA, Lowell [email protected] Observatory, Flagstaff, AZ, USA, 9Department of Earth and Planetary Sciences, University of California, Santa Cruz, CA, USA, 10Massachusetts Institute of Technology, Cambridge, MA, USA, 11NASA Jet Propulsion Laboratory, California Institute of Technology, Pasadena, CA, USA, 12Department of Electrical Engineering, Stanford University, Stanford, CA, USA, 13Ball Citation: 14 Robbins, S. J., Runyon, K., Singer, K. N., Aerospace (retired), Boulder, CO, USA, NASA Goddard Space Flight Center, Greenbelt, MD, USA Bray, V. J., Beyer, R. A., Schenk, P., … Stern, S. A. (2018). Investigation of Charon’s craters with abrupt terminus Abstract On the moon and other airless bodies, ballistically emplaced ejecta transitions from a thinning, ejecta, comparisons with other icy bodies, and formation implications. continuous inner deposit to become discontinuous beyond approximately one crater radius from the crater Journal of Geophysical Research: rim and can further break into discrete rays and secondary craters. In contrast, on Mars, ejecta often form Planets, 123,20–36. https://doi.org/ continuous, distinct, and sometimes thick deposits that transition to a low ridge or escarpment that may be 10.1002/2017JE005287 circular or lobate. The Martian ejecta type has been variously termed pancake, rampart, lobate, or layered, “ ” Received 15 FEB 2017 and in this work we refer to it as abrupt termini ejecta (ATE). Two main formation mechanisms have been Accepted 16 NOV 2017 proposed, one requiring interaction of the ejecta with the atmosphere and the other mobilization of Accepted article online 21 NOV 2017 near-surface volatiles. ATE morphologies are also unambiguously seen on Ganymede, Europa, Dione, and Published online 5 JAN 2018 Tethys, but they are not as common as on Mars. We have identified up to 38 craters on Charon that show signs of ATE, including possible distal ramparts and lobate margins. These ejecta show morphologic and morphometric similarities with other moons in the solar system, which are a subset of the properties observed on Mars. From comparison of these ejecta on Charon and other solar system bodies, we find the strongest support for subsurface volatile mobilization and ejecta fluidization as the main formation mechanism for the ATE, at least on airless, icy worlds. This conclusion comes from the bodies on which they are found, an apparent preference for certain terrains, and the observation that craters with ATE can be near to similarly sized craters that only have gradational ejecta.

1. Introduction

For the first several centuries of telescopic astronomy, impact craters were observed to either have simple, radial ejecta, or no obvious ejecta at all, for that was all that was observed on Earth’s moon (e.g., Gilbert, 1893). This type of ejecta typically has coarse, disordered texture closer to the rim that grades to finer material and then into discrete rays. Typically, the continuous portion of the ejecta extends approximately one crater radius from the rim. Mariner 9, the first spacecraft to orbit Mars, revealed a unique type of impact ejecta not yet observed elsewhere. These ejecta appeared to be smooth or hummocky, relatively thick compared with lunar ejecta, often showed a lobed or flower petal-like pattern, and the termini were typically ≈1–1.5 crater radii from the rim (e.g., Carr et al., 1977; McCauley, 1973). The ejecta deposits terminated in topographically distinct rampart edges (a ridge that forms a topographic high) or “pancake”-like edges (convex termini; first

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Figure 1. Examples of ATE craters found in the solar system where arrows on all bodies except Mars show margins of ejecta. (a) Mars (Mars Reconnaissance Orbiter CTX image P21_009381_2010_XN_21N080W, 6 m/pixel) with a 10-pixel Gaussian blur applied to downgrade it to roughly similar resolution as other examples, (b) Ganymede (Galileo image C055244/5113R, 160 m/pixel), (c) Europa (Galileo mosaic, 500 m/pixel), (d) Dione focused on Sagaris crater (Cassini ISS color mosaic, 400 m/pixel), (e) Tethys focused on Icarius crater (Cassini ISS color mosaic, 400 m/pixel), and (f) Charon focused on Nimo crater (New Horizons MVIC image mp1_0299180334_0x530_sci, 620 m/pixel); Nimo is the only crater that all of the ﬁrst three coauthors of this work agreed was a certain example of ATE (see section 2.2 and Figure S1 in the supporting information). In all images, north is up. For additional Charon examples, see Figure 3 and the supporting information.

1.1. Terminology The terminology for the impact crater ejecta on Mars that differ from the radial patterns seen on the Moon or Mercury has varied considerably throughout the literature, with Barlow et al. (2000) identifying no less than two dozen different terms. Barlow et al. (2000) recommended the term “layered ejecta” with additional morphologic descriptors indicating the number of apparent “layers,” the sinuosity or circularity of the termini, and the shape of the termini (pancake or rampart). However, this terminology has been criticized as the term itself implies a layer or layering that is not necessarily observed. In the outer solar system, the term pancake is often used to describe this class of features, but given its connotation from the Barlow et al. (2000) work to describe only a subclass of ejecta morphologies, pancake has led to additional confusion. In this work, we propose a new term: abrupt terminus ejecta or “ATE” for short. The most distinguishing feature of these ejecta on any body is an abrupt, well-defined terminus, in contrast with simple radial ejecta that grade from continuous to discontinuous and have been observed on the Moon for centuries. There may be additional facies interior to the abrupt terminus, there may be an additional radial ejecta deposit overtop or below and that continues beyond the abrupt terminus, the abrupt terminus may not be primary morphology, and the abrupt terminus may not have a single cause. But, morphologically, it is the abrupt, well-defined termini that have distinguished these ejecta since the first examples were identified on Mars. As with the Barlow et al. (2000) work, additional modifiers could be attached to the ATE term, such as a rampart versus convex (pancake) terminus or a measurement of its circularity. Similarly, a measure of how far the ejecta travels from the crater rim could be used to distinguish ATE on Mars, icy satellites, or other inner solar system bodies. For example, Copernicus crater on the Moon shows ATE, but the ATE extends ~0.3–0.4·R from the rim while the continuous radial component extends significantly farther; in contrast, ATE on Mars typically

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Table 1 extend ~1.0–1.5·R from the rim, though they can extend much farther Average CEER Calculated for Average and Maximum Distances From (see Table 1). Therefore, while the addition of another term has the the Crater Rim From This Work, Reprocessed From Robbins and Hynek (2012a) (Mars) potential to add confusion, we think that this new term is nongenetic, can accommodate additional modifiers to capture subtypes, and Average of average- Average of maximum- captures within it the most basic feature that distinguishes these Body distance CEER distance CEER ejecta from more typical radial types. Mars Equatorial ≈1.2 ± 0.4 Equatorial ≈2.0 ± 0.7 ≈ ≈ South Polar 1.6 ± 1.1 South Polar 2.5 ± 1.8 1.2. Existing Formation Models North Polar ≈1.8 ± 0.7 North Polar ≈2.8 ± 1.2 Ganymede 0.9 ± 0.2 1.2 ± 0.3 ATE on Mars were originally thought to be the result of wind erosion Ganymede* 0.8 (McCauley, 1973). Later, Viking Orbiter images showed this interpretation Europa 1.2 ± 0.5 1.9 ± 0.9 to be unlikely (Woronow & Mutch, 1980), and the two current dominant Europa* 0.8 Dione 0.8 ± 0.0 1.2 ± 0.1 formation hypotheses are they form either by impact into and melting Charon 0.8 ± 0.2 1.4 ± 0.4 or vaporization of subsurface volatiles (e.g., Carr et al., 1977; Wohletz & Sheridan, 1983) or by ballistic ejecta entrainment within atmospheric Note. The asterisk indicates data from Boyce et al. (2010). Note that the values in this table are overall averages, and a weak diameter dependence vortices generated by interactions between the advancing ejecta was found (Figure 5), so the significance of these averages is subject to curtain and the atmosphere (while shear, enabled by water, could interpretation. These values are for ejecta morphologies with a single contribute to additional mobility of the inner facies that become facies only. ground-hugging after emplacement) (e.g., Barnouin-Jha & Schultz, 1998; Schultz, 1992a; Schultz & Gault, 1979). While the volatile fluidization hypothesis has gained significant traction, and there are numerous variations from its basic premise including a ground-hugging flow versus vaporized volatiles interacting with the ejecta curtain (Boyce & Mouginis-Mark, 2006; Osinski, Tornabene, & Grieve, 2011; Senft & Stewart, 2008; Weiss & Head, 2013, 2014), there is no consensus that it is the sole—or even dominant—mechanism for formation of ATE. Additionally, it has been suggested that a hybrid or composite of the two mechanisms may be required (Barlow, 2005; Komatsu et al., 2007), and Schultz and others classify their original work (see above) as a hybrid approach. Hybrid models for Martian ATE have been proposed because certain morphologic and morphometric features are better explained by each end-member model. As examples, some ATE on Mars appear to flow around topographic highs, which is more consistent with a ground-hugging flow; and some ATE morphologies on Mars are well correlated with near-surface ice, which is more consistent with the broad class of subsurface volatiles models; but the large distances that ATE travels for certain subtypes are more consistent with atmospheric entrainment models. Evidence bolstering the volatile fluidization hypothesis was the discovery of the same basic morphology of distinct ejecta facies on some icy satellites (in contrast with radial, gradational lunar ejecta), including Ganymede (e.g., Figure 1b) and Europa (e.g., Figure 1c) around Jupiter (e.g., Horner & Greeley, 1982; Lucchita, 1980; Moore et al., 1998; Strom, Woronow, & Gurnis, 1981), and our own Saturnian satellite survey has revealed up to three examples of ATE on Dione (e.g., Figure 1d) and one on Tethys (Figure 1e). Recent work has suggested that these ejecta may form on Mercury (e.g., Xiao & Komatsu, 2013) and while many look like landslides (which may or may not involve volatiles) (Evans & Degraff, 2003; Xiao & Komatsu, 2013), at least one (around Hokusai crater) appears to be similar to Martian types (Barnouin, Ernst, & Susorney, 2015). A third formation mechanism—a volatile-rich impactor—has been suggested for Hokusai’s ejecta. Dry granular flow has also been proposed (e.g., Barnouin-Jha, 2005; Wada & Barnouin-Jha, 2006) and is a potential mechanism to explain the best example of a similar landslide-like feature on Earth’s moon, over part of Tsiolkovsky crater’s ejecta.

1.3. Importance of These Ejecta on Charon Three of the six bodies that are now known to show ATE that extend to ~1·R have of order 1–10 craters with these ejecta (Tethys, Dione, and Europa), Ganymede has on order 100, but on Mars, approximately one third of craters D ≥ 3 km (tens of thousands) show ATE (Robbins & Hynek, 2012a). We observed 38 examples of ATE on Charon (none on Pluto), comprising up to approximately one third of Charonian craters D ≳ 15 km that could be observed over ≈45% of the surface, representing a relative population closer to Mars than any other observed body in the solar system. A study of these craters and their ejecta could help constrain formation models by cataloging and better understanding the subtypes of ejecta that occur on airless bodies, such as Charon, and comparing them among each other and with Mars.

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Figure 2. Basemap of Charon. The magenta contours bracket favorable illumination for topographically identiﬁable ejecta, the cyan contours are ≤0.75 km/pixel, and the yellow contours at 65° are emission angle. The intersection of these three regions is unshaded, which indicates where we should reasonably be able to identify any ATE-bearing craters with diameters D ≥ 5 km. Additional dark blue outlines indicate higher resolution areas. The black/green dashed outline is the approximate boundary of Vulcan Planum (southern margin based on terminator of high-resolution imagery). The green markers are craters that display certain ATE morphology from both Robbins and Runyon, the white markers are craters that display probable ATE morphology by both Robbins and Runyon (or certain by one and probable by the other), and the grey are the remaining possible ATE. The letters indicate craters that are shown in Figure 3.

After discussing methods (section 2), the craters and ejecta are described on Charon (section 3), compared with other bodies (section 4), and the bodies on which ATE are found are compared with those on which they are not (section 5). Section 6 summarizes the implications of these craters on Charon and broadly addresses their constraints on different classes of formation models, and we brieﬂy summarize the broader implications in section 7. The crater catalog is included in the supporting information along with images of all 38 possible examples of these craters on Charon. Note that this paper uses informal nomenclature on Charon that has not yet been approved by the International Astronomical Union.

2. Methods 2.1. Image Base Image data used in this study come from the LOng-Range Reconnaissance Imager (LORRI) (Cheng et al., 2008) and Multi-spectral Visible Imaging Camera (MVIC) (Reuter et al., 2008) aboard the New Horizons spacecraft (Fountain et al., 2008; Stern, 2008). We used individual images as well as basemaps produced from the LORRI and MVIC panchromatic images in this work (Schenk et al., 2016). New Horizons returned coverage of ≈80% of Charon’s surface, but south of ≈30–40°S is in permanent shadow and will not be completely illu- minated until 2114 Common Era (Zangari, 2015). Approximately the same coverage was returned for Pluto. Only images within a few hours of closest approach are suitable for identifying ATE, and therefore, this study is limited to the encounter hemisphere. During that time, Charon rotated very little, and the fixed Sun poses an additional constraint on identifying ATE, which require morning/evening Sun angles, like other topographic features. Almost the entire encounter hemisphere was imaged at better than ≈0.75 km/pixel, which would allow for the reasonable detection of all D ≥ 5 km ATE craters (the craters subtending >6 pixels, and their ejecta >10 pixels), though detecting ejecta margins may be more difficult (see section 6.1). The basemap, and the intersection of favorable Sun (≈45°–85°), pixel scale (≲0.75 km/pixel), and emission angle (≲65°), is shown in Figure 2, which also shows outlines for sub-0.75 km/pixel regions. These three constraints to crater identification are empirical: the pixel scale described above, the Sun angle based on experience and some studies (Ostrach et al., 2011; Wilcox et al., 2005), and the emission angle based on experience.

2.2. Ejecta Identification Craters on Charon were identified through standard morphologic means, discussed in detail in Robbins et al. (2017). Ejecta classification was less straightforward and sometimes difficult due to the relatively restricted

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Figure 3. Examples of ATE on Charon. (top) Panchromatic images; the scale bar is 20 km. The arrows are representative: the solid arrows point to what we interpret as deﬁnite ejecta borders; the open arrows point to what some have interpreted as potential borders; the red are ATE boundaries, the yellow and green to what some might interpret as external ejecta facies; the magenta indicate potential facies inside the ATE. (bottom) Topography based on stereogrammetry and photoclinometry. (a) Spock crater is emplaced on terrain modiﬁed by tectonics. The ejecta is preserved enough to trace its perimeter. Several montes are to the south, located ≈24.5–26.0 km and ≈43–58 km from the rim, covering 25° and 33–70° of azimuth. The montes could be interpreted to approximately follow Spock’s rim, but the authors’ consensus is that these are not additional ejecta facies. Topography possibly indicates a thickening toward the distal margins. (b) Revati crater has linear disruptions to the southwest. It displays a single nearly continuous ejecta facie, but interior to it in at least one area is an additional facie and some features exterior to the continuous ejecta could be interpreted as additional ejecta (though we do not). Revati’s topography indicates a likely thickening at the distal margins. (c) Siddhartha crater imaged near the terminator that displays a hummocky texture that, toward the north- northeast, blends with the surrounding surface.

image data for Charon (section 2.1 and Figure 2). Robbins first identified all craters he thought could possibly have ejecta. He classified the ejecta into “certain,”“probable,”“possible,” and “radial.” The craters were numbered at random and sent to Runyon and Singer, who followed the same system while adding “none” for no ejecta around the crater. The definitions for this subjective classification are as follows: 1. Certain: ATE is obvious, ejecta subjectively appears to have an abrupt terminus rather than grade to the surrounding surface, and it may have ramparts over some regions. 2. Probable: ATE is less obvious but could be distinct for some azimuths; there is a small, subjective possibility that the ejecta may grade from continuous to discontinuous ejecta rather than have an abrupt terminus. 3. Possible: Could be ATE, could be clumpy lunar appearing ejecta, could be preexisting surface features, perhaps limited by Sun angle, or smear. Subjective “50/50” chance. 4. Radial: There may be ejecta (>50% subjective chance), and if there is, 90–100% subjective chance it is lunar appearing. We made this determination by first examining what appears to be the ejecta itself: Does there appear to be a quasi-circular, continuous ATE at an approximately consistent distance from the crater rim? If this feature appears but is not continuous, could the discontinuity be disruption by other features (e.g., in Figure 3a; Spock crater has tectonic disruption of its ejecta to the north)? If this feature is present and continuous, or disruption could be plausibly explained by preexisting features or subsequent modification, then the

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ejecta was generally interpreted as more likely to be ATE. If not, then the ejecta was less likely to be ATE. Figure S1 in the supporting information shows all possible ATE examples and gives the certainty from Robbins, Runyon, and Singer that these ejecta are ATE. In this analysis, Robbins was the most liberal with 14 certain, 13 probable, and 10 possible. Runyon had 12 certain, 13 probable, 11 possible, and 2 that Robbins identified as possible he considered to be lunar-like ejecta (one that he identified as possible, Robbins considered to be lunar-like). Singer had 1 certain, 4 probable, and 8 possible with the remainder having no obvious ejecta, or surrounding features were considered to be background or preexisting properties of the surrounding surface. For context, both Robbins and Runyon have previously studied these ejecta in detail on Mars. The analysis within this manuscript uses a combination of Robbins’ and Runyon’s classifications for its conclusions; if we used the most conservative classifications, no analyses could be conducted. Therefore, we present our findings, state here and in Figure S1 in the supporting information where we have differences as authors, provide our conclusions based on a subset of our findings, and let the reader draw their own interpretations.

2.3. Identification of Additional Abrupt Terminus Morphology Our proposed ATE nomenclature easily permits additional morphologic descriptors, and in past research these descriptors have sometimes proven useful in various analyses. First is the number of facies interior to the final one. Second is whether the ejecta appears to be a contiguous feature throughout, or whether it appears broken into numerous azimuthally separated petal-like structures. Both of these morphologies are typically easy to observe visually. An additional morphology is whether the ATE is a rampart—characterized by a topographic rise relative to the rest of the ejecta deposit—or not (often termed pancake). Possible ramparts on different ejecta deposits were identified visually as what appears to be a highlight on the Sun-facing side and shadow or darker area on the anti-sunside. While this technique is not possible for every ejecta deposit, nor is it possible for the entire perimeter of the ejecta, it can generally be used to distinguish a rampart from preexisting undulating terrain or large blocks. This distinction is made by ensuring that what appears to be a rampart is continuous for a “reasonable” portion of the ejecta perimeter. For example, Spock’s ejecta appears to have a rampart from approximately SSE through SSW, or about 45° of azimuth around the crater. There are small potential rampart areas elsewhere along the perimeter, but they are not continuous enough to be a convincing rampart. Similarly, Figure 3b shows what may look like a rampart around the south clockwise to the west margin, but it is less conclusive; however, the topography data indicate elevated margins, and so it was marked as a likely rampart. Importantly, the determination here of ramparts is generally made at the edge of resolution, for we are using pixel scales of 160 m or coarser; as such, we may be overinterpreting these data, and so the classification of ramparts is not used in any of our implications sections.

2.4. Measurement of Abrupt Terminus Morphometry Abrupt terminus ejecta have two morphometric properties that are typically measured. The first is the lobateness (or circularity), which is the ratio of the ATE perimeter to the circumference of a circle with the same area as the ATE, often symbolized as Γ. Lobateness was measured using the same vertex spacing of 500 m for all bodies presented in this study, which is the same as Robbins and Hynek (2012a), so that the data could be directly compared (using different vertex spacing can affect the results due to the fractal nature of perimeters (Mandelbrot, 1967)). The ATE perimeter could not be traced in all cases due to superposed features and/or preexisting topography, such as Revati crater in Figure 3b. However, in cases of very small disruptions such as a relatively small, superposed crater a few pixels across, a simple visual interpo- lation was done. The lobateness of ejecta may be related to flow instabilities, either in the framework of atmospheric vortices (Barnouin-Jha and Schultz, 1998) or fluidized ejecta (Baratoux et al., 2002). The second common morphometric measurement is how far the continuous ejecta facies extend from the crater rim. Distance from the rim is called “ejecta runout distance” in Martian studies, a term we make more generic here as “continuous ejecta extent” (“CEE”). Continuous ejecta extent from the rim divided by crater radius is “ejecta mobility ratio” in the literature (e.g., Mouginis-Mark, 1979), though here we use “continuous ejecta extent ratio” (“CEER”) to maintain a nongenetic term. CEER is often measured for the farthest distance any component of the ejecta extends from the rim, but Robbins and Hynek (2012a, 2012b) used average ejecta extent in this calculation because they considered it more representative of the overall dynamics.

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For compatibility, we calculated average and maximum CEE and CEER and reanalyzed all data from Robbins and Hynek (2012a) to produce these four measurements. Average CEE was calculated as the radius of a circle with the same area as the ejecta plus crater, minus the radius of the crater; maximum CEE was calculated as the maximum distance from the crater rim of any point traced along the ATE perimeter.

2.5. Crater Size-Frequency Distributions An analysis in this work compares crater size-frequency distributions (SFD). This analysis was done in two ways. The ﬁrst used recommendations of the Crater Analysis Techniques Working Group (1979) to bin data and use Poisson (N1/2) error bars. Both cumulative and relative plots are used, the cumulative being where each smaller diameter bin is the sum of its data and all data larger than it, and the relative is a differential SFD normalized to a À3 slope power law. The second type of SFD shown relies on a new technique proposed in Robbins et al. (2018), which uses a kernel density estimator (treats each crater as a Gaussian with μ equal to the diameter measured and σ equal to 10%·D (informed by variability studies)). The uncertainty is based on a modiﬁed bootstrap with replacement that samples the original crater diameters where there is “dense” data (the difference in two crater diameters in a sorted list is less than the sum of their σ) and the SFD where data are not dense. The data are displayed as both cumulative and relative versions from this method, and the graph includes a rug plot, which is a small mark along the x axis at each originally measured crater diameter.

2.6. Topography New Horizons obtained extensive stereo coverage of the encounter hemispheres of both Pluto and Charon, and established stereogrammetric methods that use parallax and feature-matching were used to derive terrain from those images (Schenk, Wilson, & Davies, 2004). The stereogrammetric method has an approximate ground scale of ≈4.4 km and vertical precision ≈0.1–0.25 km (Cook et al., 1996; Schenk et al., 2016). Stereogrammetry was augmented by photoclinometry for a 300 m/pixel product for Charon; long- wavelength errors in the photoclinometry were corrected by tying it to the stereogrammetric product through a method similar to that of Schenk (1989). There are no signiﬁcant, short-wavelength albedo differences of terrain in the region of Charon where most ATE reside (see section 3.1) (Buratti et al., 2017) and so there should not be any large errors affecting the photoclinometry.

2.7. Central Tendency Ellipses An analysis in this work relies on central tendency or standard deviational ellipses. This is a common technique that computes the minimum-area ellipse that bounds x% of the data. These were constructed for the inner 68% of the data (1σ assuming a Gaussian distribution) and can be used to make any trends (or lack- there-of) appear more clearly.

2.8. Crater Preservation In some of the discussion, we use the idea of a well-preserved impact crater. While recent studies have shown that crater preservation is a difﬁcult metric to tally and requires hectometer-scale resolution (e.g., Piatek et al., 2016, 2017; Tornabene et al., 2016), the fact that the craters still possess an ejecta deposit, are not buried, show “sharp” features at the pixel scales available, and are not overprinted with comparable-sized craters can be reasonable proxies for good preservation. It is craters that fall in this category to which we refer when discussing such a classiﬁcation.

3. Observations of Craters With Abrupt Terminus Ejecta on Charon 3.1. Geographic Distribution Figure 2 shows the locations of all 38 potential ATE craters. All but one of the 38 is within the intersection of the three empirical observing constraints (if detection is extended to the terminator). This region of favorable observing conditions is dominated by Vulcan Planum, which could be a selection effect (see section 5). The spatial density of the ATE within Vulcan Planum is 70 ± 12 per 106 km2 versus 18 ± 11 per 106 km2 (three craters) in the remaining favorably observed region (assuming Poisson N1/2 uncertainty). If using just craters Robbins and Runyon marked as certain, the spatial density is 20 ± 6 (10 craters) on Vulcan Planum versus 0.

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3.2. Size Distribution The diameter range of the identified ATE craters is ≈7.5–66 km, and Figure 4 shows the SFD with different thresholds of certainty (see sections 2.2 and 2.5). Regardless of certainty of the ejecta classification (between classifications by Robbins and Runyon), these observations hold: (1) ATE-bearing craters compared with others in Vulcan Planum follow the same SFD for D ≳ 20–30 km; and (2) ATE-bearing craters D ≲ 20–30 km do not follow the same SFD as the rest of Vulcan Planum, instead showing a marked decrease in number. Whether this may be a detection bias is discussed in section 6.

3.3. Morphology Three primary morphologies were discussed in section 2.3. The first is whether there exist additional facies between the crater rim and term- inal margin, and none were convincingly found on Charon. The next classification is whether the ejecta appears to be primarily contiguous throughout or whether it appears separated into distinct, petal-like lobes. None of the latter were observed on Charon. The third classification is whether the terminus is rampart or pancake shaped, where rampart was indicated if any reasonable portion of the terminus showed a relative topographic high. There are craters that we are confident show ramparts at least over part of their distal margins (e.g., Figure 3a, which shows Spock crater), given the caveat in sections 2.3 and 6.1. Of all potential ejecta that could be classified (30 of 38), two-thirds are pancake and one-third have at least some regions that show potential ramparts, such as Spock (fractions are used to avoid the appearance of two significant figure precision). This classification shifts when moving to only ejecta that are most certain, changing to two-fifth pancake and three-fifth rampart. The shift is because ejecta that are most distinct have the most obvious boundary between the ejecta and surrounding surface, and this boundary is sharpest when marked by a rampart. While the identification of at least some of these potential ramparts is generally convincing visually, their width tends to be a few pixels regardless of the pixel scale. This phenomenon is troubling, and so we Figure 4. Size-frequency distributions (SFDs) of craters in Vulcan Planum leave an in-depth study and comparison of ramparts to future work. (see Figure 2) and subsets of Charonian ATE craters in this study. Three versions of the subset are included: craters both Robbins and Runyon 3.4. Morphometry marked as certain; marked as certain or probable; or marked as possible, probable, or certain. The grey gradient indicates D ≤ 5 km, the approximate Two primary morphometries were discussed in section 2.4. The first is limit of crater detection completeness in Vulcan Planum based on all lobateness, Γ, a measure of the circularity or sinuosity of the ejecta ter- available images (see section 6.1 for why ATE detection limit may be larger). minus (a perfect circle is Γ = 1.0). In Barlow et al. (2000), the ejecta was (top row) Cumulative SFDs. (bottom row) Relative (“R”) SFDs. (left column) considered circular if Γ < 1.5 and sinuous if Γ > 1.5. For measurable Traditional binned SFDs (Crater Analysis Techniques Working Group, 1979). (right column) New versions of the SFDs recommended by Robbins et al. Charonian ejecta, 88% (14 of 16) are circular and 12% are sinuous (the (2018). Axes aspect ratio is 1:1 such that log10 horizontal is the same as log10 two sinuous being marked as certain or probable by both Robbins and vertical; completeness limit is estimated based on the downturn in the dif- Runyon). The full range is Γ = 1.07–1.60. Because decreasing vertex spa- ferential SFD for Vulcan Planum (Robbins et al., 2017). The 2σ confidence cing results in an increased perimeter length, it would be expected that band estimated from bootstrap sampled from a hybrid of the crater distri- Γ could decrease slightly with diameter as there is less ejecta with smal- bution and original diameters 10,000 times. Area for Vulcan Planum is 4.9·105 km2, area for ATE is 6.5·105 km2. ler craters. However, there is no trend with location, and there is only a weak correlation between Γ and D (Figure 5, top). The second morphometry is the absolute and relative continuous ejecta extent (CEE and CEER), the averages of which are shown in Table 1 with other solar system bodies. It shows that the average CEER is 0.8 ± 0.2 (μ ± σ), while the maximum CEER is 1.4 ± 0.4. Geographically, larger CEE and CEER are in the northeastern portion of Vulcan Planum, but there is no strong trend.

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Figure 5. Morphometric trends of ATE craters on six solar system bodies shown in two different formats. (left column) Scatterplot that excludes Mars data for readability. (right column) Ellipses that represent the inner 68% of the data (section 2.7); the ellipses visually reduce the data set for easier inspection of trends and clustering. (top row) Ejecta lobateness versus crater diameter showing that lobateness increases on all bodies with increasing crater diameter; some of this trend could be a resolution effect (see section 2.4). (middle and bottom) CEER (section 2.4) versus crater diameter. (middle) Maximum ejecta distance from the crater rim. (bottom) Average ejecta extent from the crater rim. CEER ‡ generally decreases with increasing diameter at similar rates on all bodies except for Europa. Crater data from Robbins and Hynek (2012a) instead of this study; note that they only calculated values for a few craters D < 3 km, but D ≥ 3km should be a complete census. Because there are so many more smaller than larger craters, the central tendency ellipses are pulled to smaller diameters. *Data from Boyce et al. (2010) for Lonar crater.

4. Comparison of Abrupt Terminus Ejecta on Charon With Other Bodies ATE with termini that extend to ~1·R have been unambiguously observed on Charon, Dione, Tethys, Europa, Ganymede, and Mars. In this section, we describe the Charonian observations in context with these other solar system bodies.

4.1. Geographic Distribution On Charon, ATE are predominantly found in the Vulcan Planum region. ATE appearing preferentially on a speciﬁc terrain is common. On Mars, they are found preferentially on speciﬁc terrains and are not uniformly distributed (e.g., Barlow, 1988; Barlow & Bradley, 1990; Barlow & Perez, 2003; Robbins & Hynek, 2012b). For example, ATE with one interior facie (“DLE” from Barlow et al., 2000) are predominantly found in middle to high northern latitudes, while those with multiple inner facies (“MLE”) are often found emplaced on volcanic terrain and along the dichotomy boundary. Ganymede’s ATE are found exclusively on grooved terrain (Horner & Greeley, 1982). In contrast, ATE on Europa, Dione, and Tethys are not preferentially found on

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specific terrain types, though this could be because craters, in general, are only found on one terrain type on Europa, and the relative homogeneity of the other surfaces and small ATE numbers (for Dione and Tethys). 4.2. Size Distribution On Charon, ATE followed the typical SFD of Vulcan Planum’s other craters for D ≳ 20 km, and there was a relative paucity of ATE around craters D ≲ 20 km. This trend of a diameter threshold is also observed on Mars where craters with MLE peak in number D ≈ 20 km, DLE at D ≈ 15 km, and SLE at D ≈ 5 km (Robbins & Hynek, 2012a, 2012b). (Without full databases of Ganymedian and Europan craters, a similar comparison cannot be made; Dione and Tethys have too few craters for a meaningful comparison.) 4.3. Morphology For the first morphology discussed in section 3.3—facies interior to the ATE—Charon is the same as both Dione and Tethys with no examples of inner facies visible in available data. It is similar to Europa in that we identified only one Europan example of inner facies (though 1 of 9 is ~10% of the total). These bodies are less similar to Ganymede, where we identified at least two ejecta with inner facies and Boyce et al. (2010) measured six. Finally, Mars is unique among the six bodies: of craters D ≥ 3 km that have ATE, 78% have no inner facies, 18% have a single inner facie, and 4% have more than one (Robbins & Hynek, 2012b). We think that this is a true difference rather than a selection bias, for we would expect ~25 with one inner facies and ~5 with multiple on Ganymede if the trend were to hold, and ~7 and ~1–2 on Charon, respectively. The second morphology—the shape of the ejecta deposit—is similar on Charon to the ATE on the four other icy moons. It is dissimilar to Mars where the ejecta display a broad array of shapes. The third morphology— whether the ejecta ends in a topographic rampart—is difficult to assess because of the relatively coarse elevation data in the outer solar system. Boyce et al. (2010) classified several ATE on both Ganymede and Europa as having ramparts, but we are not as certain about those features. Therefore, we refrain from making comparisons and conclusions based on the presence or absence of ramparts. 4.4. Morphometry On Charon, the first morphometric property of lobateness revealed few ejecta Γ > 1.5. For direct comparison, we independently mapped the ATE on Dione (3 of 3), Europa (9 of 9 including one with inner facies), and Ganymede (94 of 137 including 2 with inner facies). Ejecta with Γ > 1.5 was found to be similarly rare on Ganymede (1 ATE) and Europa (0 ATE) both from our measurements and Boyce et al. (2010) (Figure 5). Mars is significantly different, for ≈85% of D ≥ 15 km craters have ATE with Γ > 1.5 (Robbins & Hynek, 2012b). There is no obvious geomorphologic province trend for lobateness on any body. In more detail and without the artificial Γ = 1.5 threshold, Figure 5 demonstrates that Charonian ATE tend to have greater Γ than the other icy bodies we examined except Dione. Charon’s ejecta follow the trend of Europa’s if the Europan ejecta were present at larger diameters and/or Charonian ejecta were present at smaller diameters. However, they are markedly more sinuous than Ganymedian ATE. Mars’ ATE, in contrast, has significantly larger Γ both on average and as a function of diameter. Ganymede’s and Europa’s also vary weakly with diameter (similar to findings by Neal & Barlow, 2003), while Mars’ vary significantly and Charon’s Γ versus D trend is intermediate. The second morphometry, CEER, is compared to other bodies in Table 1 and Figure 5; the table includes Boyce et al. (2010) values. While Boyce et al. (2010) found lower values than we, the trend for the CEER to be less for Ganymede and Europa than Mars is similar (Neal & Barlow, 2003, and Barlow, 2006, data are also reported in Boyce et al., 2010). Our CEER data versus D are shown in Figure 5 (middle and bottom). Fortunately, the average versus maximum ejecta extent in the definition for CEER do not show significantly different results—this similarity indicates that either definition for CEE and CEER does not matter in this appli- cation and type of study. For average CEER, Charon is most similar to Dione and second-closest to Ganymede (Table 1). For maximum CEER, Charon is most similar to Ganymede and second-closest to Dione, indicating consistent results between the two methods. In contrast, Mars’ equatorial region is similar to Europa, but Mars’ polar latitudes (poleward of ±40° latitude) are dissimilar from all other bodies (Table 1). We observe a small but robust dependence on crater diameter for CEER based on the major axis direction of our central tendency ellipses for all bodies; this trend is different from Boyce et al. (2010), who found no D-dependence on any body. A possible explanation for the difference between this study and Boyce et al. (2010) is that they limited their survey to craters D > 10 km except for ≈6 data points, and the bounding ellipses they illustrate were manually drawn rather than calculated based on an objective algorithm.

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Figure 6. Examples of craters on Charon in the Vulcan Planum region that (a) show no obvious ejecta, (b) show lunar-like ejecta (cyan arrows), and (c) show ATE (red arrows) near to a comparable sized crater that does not (from Figure 6a). It is possible that the crater in Figure 6a has lunar-like ejecta, and Figures 6a and 6b have ATE, but they may be hidden by resolution, lighting, and/or signal-to-noise of the images. Figure 6a is from a 400 m/pixel basemap, but the crater is also in the 160 m/pixel closest approach sequence and also shows no ejecta—the closest approach imagery is not shown here because of the comparatively low signal-to-noise, and they cannot be easily combined because of the slightly different solar incidence angles and parallax.

5. Discussion of the Presence of Abrupt Terminus Ejecta Versus Absence of It Here we work from Charon toward the Sun and discuss bodies that do not have this ejecta and include some interpretation. It is always difficult to prove a negative, but sufficient imagery exists for many bodies that allow such a comparison. On Charon in Vulcan Planum, ATE do not surround otherwise well-preserved craters, and some have what appear to be preserved Moon-like radial ejecta (Figure 6). Some of these craters are similar in size and location to nearby ATE-bearing craters. This juxtaposition is observed on all other bodies where ATE is present, where otherwise practically identical, adjacent craters have different ejecta. A caveat is that while we demonstrated in section 3.1 that the abundance of ATE on Vulcan Planum does not appear to be an observational bias, it is possible that the relative smoothness of Vulcan Planum could enhance ATE detection relative to other areas of Charon. ATE were not found on Pluto despite an image campaign that yielded similar illumination and emission angle but pixel scales up to 2× better. While it is possible that ATE are present on other terrains New Horizons did not image on such a geologically diverse world as Pluto (Moore et al., 2016; Stern et al., 2015), we consider this an unlikely explanation. It is more probable that (1) they do not form on Pluto due to differences in the upper crust between the two bodies; (2) the absence of smooth topography on Pluto’s older terrains impedes ejecta flow, preventing them from forming; or (3) the ejecta do form on Pluto (by whatever mechanism) but geolo- gic processes are so rapid the ejecta were removed. It is hard to state whether ATE can be ruled out on Uranian or Neptunian satellites due to the relatively poor coverage and large pixel scales. Around Saturn, we surveyed every spherical moon (Mimas, Enceladus, Tethys, Dione, Rhea, Titan, and Iapetus) and only observed three and one ATE on Dione and Tethys, respectively; we observed two or three potential deposits on Titan, but interpretation is difficult without corre- sponding visible-light imagery, and we cannot be certain whether they are the same type of feature so we omit Titan as a constraint. Around Jupiter, we observed of order 100 ATE on Ganymede and order 10 on Europa. This difference is despite significantly better coverage of Europa (15.6% of Europa was imaged at ≤500 m/pixel versus 1.7% of Ganymede), although Europa has relatively few craters and only 35% the surface area of Ganymede. Despite Ganymede having more ATE, researchers have proposed that their exclusivity to certain terrains could be a selection effect caused by ejecta flow disruption or difficulty identifying the ejecta on older, more cratered and topographically complex terrain (Horner & Greeley, 1982; Wada & Barnouin-Jha, 2006). No ATE were found on Callisto. Completing the observed icy bodies, Ceres is icy (e.g., Combe et al., 2016; Küppers et al., 2014; Lebofsky et al., 1981) and yet full coverage at high resolution shows only possible evidence of ATE (e.g., Hughson et al., 2017) though significant numbers of impact-generated landslides exist (Schmidt et al., 2017).

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Finally, ATE are absent on all observed small, relatively dry, rocky bodies (Phobos, Deimos, asteroids). For large, relatively dry rocky bodies, we have already addressed Mars. Some researchers have proposed that ATE exist on Earth around several craters (e.g., Meteor Crater (e.g., Grant & Schultz, 1993), Ries (e.g., Kenkmann & Schönian, 2006; Osinski et al., 2011), and Tsenkher (e.g., Komatsu et al., 2006)), but because of erosion, we do not include Earth in this study. There are some examples on the Moon, such as Copernicus, but the ATE extends significantly less than on the icy bodies and Mars. Here we reiterate that we are studying crater ejecta deposits that occur regardless of preexisting topography. Therefore, while the lunar Tsiolkovsky crater also shows ATE, it is into a preexisting topographic low, and it occurs only over a small arc of the ejecta perimeter. This morphology is similar to the landslides on Mercury (Xiao & Komatsu, 2013) and Ceres (Schmidt et al., 2017). While radar data exist for Venus’ surface and many ejecta appear to be ATE, interpretation of ejecta morphology is similarly difficult as with Titan and so they are not used as a constraint here, though other work has considered Venus as an end-member for atmosphere versus atmosphereless ATE study (e.g., Barlow, 2005; Schultz, 1992b). Finally, there is one possible ATE discovered to date on Mercury around Hokusai crater (Barnouin et al., 2015), though the ATE extends to ~0.6–0.7·R, significantly less than ATE discussed on bodies in other sections of this work.

6. Interpretation and Implications for Formation In this section, we bring together the different observations of, and lack of, ATE on various solar system bodies. ATE may not have a single cause across all bodies, and they may not even have a single cause on a single body, but given the above data, several different observations can be made in the context of different classes of proposed formation mechanisms. To reiterate, we focus here on ATE that extend ≳1·R from the crater rim and we specifically exclude Venus, Earth, and Titan. 6.1. Formationally Ambiguous, but Important Observation Figure 4 shows that ATE form preferentially on Charon around larger diameter craters and may have a cutoff, diameters below which they do not form. This finding on Charon is potentially biased: Despite more than sufficient 155–650 m/pixel coverage through Vulcan Planum to detect all craters D ≳ 5–6 km (and ejecta that is twice as large), the spatial scales may not be enough to identify ejecta margins, a requirement for identifying ATE. If we were to treat the finding as reality, it would not be surprising because the same phenomenon is observed on Mars. Restriction to larger craters would likely indicate either (a) relatively more melting and vaporization by larger/faster (higher energy) impactors is required to form these ejecta types and/or (b) ATE requires excavation further into the surface of the body to reach material needed to form the ejecta. While this restriction may seem to exclude an atmospheric formation mechanism, one could propose a sce- nario where a larger impact event produces a denser local atmosphere (or produces any atmosphere for atmosphereless bodies) that supports the formation of ATE (e.g., Horner & Greeley, 1982). Therefore, while we think that this observation is important if it is real, it does not help to differentiate between the predomi- nant proposed formation mechanisms. 6.2. Erosional Features On Mars, features known as pedestal craters are now generally interpreted to be erosional features where overlying ejecta has cemented the underlying material and the surrounding surface has eroded away (e.g., Barlow, Boyce, & Cornwall, 2014; Kadish, Barlow, & Head, 2009; Kadish & Head, 2011; Schultz & Lutz, 1988; Wrobel, Schultz, & Crawford, 2006). The pancake Martian ATE are also thought by several researchers (e.g., Barlow, 2015; Costard, 1989; Schultz, 1992a) to be erosional remnants of other ATE morphologies, an important interpretation since most ATE on icy bodies appear to have this type of terminus rather than ramparts. However, erosion cannot be the mechanism for ATE types on Charon because Charon’s current surface material (water- and ammonia-ice) is stable at Charonian temperatures over the lifetime of the solar system (Fray & Schmitt, 2009). 6.3. Surface Gravity The icy Galilean satellites have similar surface gravity yet ATE do not exist on Callisto; similarly, Charon’s surface gravity is almost identical to Ceres, yet Charon unambiguously has ATE but craters on Ceres do not. Surface gravity does not control morphometry given (1) the greater similarity between Europa and Charon than between Europa and Ganymede for lobateness (Figure 5) (despite Europa’s surface gravity being

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closer to Ganymede), (2) there is no grouping by surface gravity in CEE versus diameter (Figure 5), (3) nor does the average CEER trend with it (Table 1). Therefore, there is no obvious trend with or restriction by surface gravity for any form of ATE so long as ejecta can form on the body (e.g., very small bodies lack the surface gravity to retain ejected material).

6.4. Impact Velocity There is no dependence on solar distance nor planetocentric distance (for giant planet satellites) with whether ATE occur or not, despite the order of magnitude variations across the solar system. Therefore, there is no linear dependence on impact velocity for ATE formation.

6.5. Impactor Composition Impactor composition is rarely considered in ATE formation, and our data do not support it as a substantial component. If volatiles are required in the impactor, that would require an icy impactor, which are likely only ~5% of inner solar system impactors (e.g., Bottke et al., 2002; Granvik et al., 2016; Strom et al., 2005) but are the majority of outer solar system impactors (e.g., Zahnle et al., 2003). Therefore, if the impactor composition was a primary driver, there should be a preference for ATE around most craters on most or all bodies in the outer solar system and little in the inner solar system. While impactor composition may still be a factor (e.g., the existence of ATE with ramparts around Hokusai on Mercury (Barnouin et al., 2015)), it is unlikely to be a signiﬁcant driver given present-day knowledge about impactor populations throughout the solar system, although uncertainty remains (e.g., Nesvorný, Roig, & Bottke, 2017; Steel, 2014).

6.6. Atmosphere, Exclusively Some of the original atmospheric models (e.g., Schultz, 1992a) used atmospheric effects to explain the flow-like pattern of inner and outer facies. The researchers were able to produce a contiguous ejecta rampart superpos- ing ballistically emplaced deposits under 0.06–0.3 bar atmospheric pressure, and they were able to produce ejecta flow lobes under 0.3–0.7 bar without any volatile component (laboratory conditions were scaled to reasonable values for Mars through tables and equations given in Schultz, 1992a), and the work acknowledged the potential role of volatiles. In the same and previous work (e.g., Schultz & Gault, 1979), it was also acknowledged that postemplacement mobility of the inner facies could be enhanced by the presence of a lubricant. Absence of ATE on at least one body that has an atmosphere indicates that an atmosphere does not guaran- tee their formation; we refer to Pluto, though its surface pressure is 3–60 μbar (Gladstone et al., 2016) while Mars’ current pressure averages 6 mbar and Earth’s 1 bar, indicating perhaps a minimum atmospheric pressure (which would control wind speeds) and density (which would control particle entrainment) may be required for some of ATE subtypes (though local gas pressure would be significantly increased immediately following an impact event). Given the different fractions and de facto existence of some morphologies of ATE on Mars versus other bodies, it is likely that atmospheric interaction at a minimum assists formation of some of the ATE morphologies which better fits the originally proposed atmospheric model; Schultz (1992a) suggested that shifts in atmospheric pressure or ejecta size by a factor of 5 could affect the observed morphologic differences. This is supported by larger lobateness, CEER, and more inner facies of Martian ATE compared with any other body. However, the fact that one can observe what appear to be two craters of similar diameters and similar preservation states close to each other, and one of those craters may have ATE and the other not, would tend to indicate that local variations in the upper crust may exert more control over whether ATE form. This is an observation made by past researchers (e.g., Barlow, 2005; Schultz, 1992a) for Mars that we extend to Charon (and observed on the other icy bodies where ATE are found).

6.7. Subsurface Volatile Fluidization and Ground-Hugging Flows Based on their existence on atmosphereless bodies and the substantial amounts of ices on the surface of most bodies with ATE, surface or near-surface properties likely contribute to ATE formation. Characteristics of planetary surfaces easily vary on scales of kilometers, which can explain why adjacent craters of similar size and preservation can have one with ATE and the other with non-ATE ejecta. Surface properties can explain why these ejecta are observed on some bodies and not others (though the exact conditions for formation remain unknown), why they occur on bodies without atmospheres, why they occur with varying frequency on different bodies at similar solar distances, and why they are sometimes correlated with certain terrains on different solar system bodies (with the caveat of certain observational or preservation

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biases discussed in section 5). Nontraditional volatiles could explain ATE on Mercury, too, with sulfur play- ing the part of water-ice elsewhere. However, an upper crust rich in volatiles is clearly not enough for ATE formation: Enceladus, Europa, and Dione have icy surfaces and are thought to have subsurface oceans (though evidence is much stronger on Enceladus and Europa), yet ATE are not found on Enceladus and are rare on Europa and Dione. This is despite the ice shells of Enceladus and Europa likely being of comparable thickness (e.g., Nimmo & Manga, 2009; Spencer et al., 2009) though Dione’s ice shell is likely at least twice as thick (Beuthe, Rivoldini, & Trinh, 2016). Most other satellites of Saturn are also thought to be icy, yet ATE are rare. Similarly, Ganymede and Europa have ATE but Callisto does not, despite all three bodies being icy, atmosphereless, and likely having subsurface oceans (e.g., Carr et al., 1998; Kivelson et al., 2000; Schubert et al., 2004).

6.8. Hybrid Both Barlow (2005) and Komatsu et al. (2007) suggested a hybrid mechanism between atmospheric interaction and subsurface volatile ﬂuidization on Mars that is a plausible mechanism for, at a minimum, enhancing formation of this ejecta type on icy bodies. A hybrid model could also apply on Charon: We speculate that if the interior of the body releases a gaseous volatile such as ammonia that seeps upward (Grundy et al., 2016) and becomes cold-trapped, it could be immediately released upon crater formation and form a transient, localized atmosphere that could enhance ATE formation. The presence of ammonia on Charon’s surface could also enhance ﬂow: At impact speeds of ~2 km/s, ice-ice impacts will bring ammonia hydrate or dihy- drate partially to the molten regime (~176–188 K (Croft, Lunine, & Kargel, 1988; Hogenboom et al., 1997; Kargel, 1992)) that can lubricate the ejecta emplacement. The same basic idea is often thought to occur on Mars with water-ice, which is what many researchers have invoked to explain why there is a bias for ATE— especially ATE with inner facies—to occur on Mars at high northern latitudes (e.g., Barlow & Perez, 2003; Robbins & Hynek, 2012b) where Mars Odyssey’s Gamma Ray Spectrometer data have been interpreted to show enhancements of near-surface water-ice (Feldman et al., 2004; Wilson et al., 2018).

7. Conclusion Putting the observations and interpretations together, our findings agree with past work showing that conditions to form the basic class of abrupt terminus ejecta exist on multiple bodies in the solar system, but Mars possesses a unique set of conditions that either (a) modifies the post-ballistic emplacement of ejecta or (b) there are varied mechanisms to form the ejecta across the solar system that produce what is generally interpreted as the same class of features. This latter possibility we consider to be more likely to explain some of the ATE seen on bodies such as Mercury or the Moon where the relative rarity of ATE suggests more rare conditions could be involved to cause their formation, such as a comet impactor in the inner solar system or enhancement of dry flows (e.g., Wada & Barnouin-Jha, 2006). We think that this is a robust result despite the tens of thousands of examples of ATE on Mars versus a total ~150 examples on other bodies. We think this primarily because of the distinct differences in morphometric (Table 1, Figure 5) and morphologic properties and the relatively large fraction of ATE with inner facies on Mars compared with other bodies. Future work and models should consider and be consistent with these observations. A unified model for for- Acknowledgments mation of ATE—if one can be made—would also need to explain why they are not observed on certain Support for the authors was made bodies that seem like they should contain these features based on their similarities with other bodies that possible through NASA’s New Horizons mission within the New Frontiers pro- do. While the addition of Charon is an important set of data to this overall problem of planetary geology, gram. S.J. Robbins was additionally how the ejecta form and why they form on certain bodies, locations on those bodies, and with certain proper- ’ supported by NASA s Mars Data Analysis ties are still inconclusive. Program award NNX15AM48G. The authors thank N.G. Barlow for nomenclature assistance, W.F. Bottke for sug- gestions on inner solar system impactor References populations, and R.R. Herrick for discus- Baratoux, D., Delacourt, C., Allemand, P. (2002). An instability mechanism in the formation of the Martian lobate craters and the implications sions on Venus radar data. The authors for the rheology of ejecta. Geophysical Research Letters, 29, 1210. https://doi.org/10.1029/2001GL013779 thank P. Schultz, P. Mouginis-Mark, G.R. Barlow, N. G. (1988). Crater size-frequency distributions and a revised Martian relative chronology. Icarus, 75(2), 285–305. https://doi.org/ Osinski, and N.G. Barlow for their helpful 10.1016/0019-1035(88)90006-1 reviews. All image data used in this work Barlow, N. G. (2005). A review of Martian impact crater ejecta structures and their implications for target properties. In Special Paper 384: Large are available via NASA’s PDS’ Small Meteorite Impacts III (Vol. 384, pp. 433–442). Boulder, CO: Geological Society of America. https://doi.org/10.1130/0-8137-2384-1.433 Bodies Node; the derived crater data are Barlow, N. G. (2006). Impact craters in the northern hemisphere of Mars: Layered ejecta and central pit characteristics. Meteoritics and available in the supporting information. Planetary Science, 41(10), 1425–1436. https://doi.org/10.1111/j.1945-5100.2006.tb00427.x

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