Impact Craters on Titan: Finalizing Titan's Crater Population

Western University Scholarship@Western

Electronic Thesis and Dissertation Repository

8-3-2018 2:00 PM

Joshua E. Hedgepeth The University of Western Ontario

Supervisor Neish, Catherine D. The University of Western Ontario

Graduate Program in Geophysics A thesis submitted in partial fulfillment of the equirr ements for the degree in Master of Science © Joshua E. Hedgepeth 2018

Follow this and additional works at: https://ir.lib.uwo.ca/etd

Part of the Geology Commons, Geomorphology Commons, and the Geophysics and Seismology Commons

Recommended Citation Hedgepeth, Joshua E., "Impact Craters on Titan: Finalizing Titan's Crater Population" (2018). Electronic Thesis and Dissertation Repository. 5618. https://ir.lib.uwo.ca/etd/5618

This Dissertation/Thesis is brought to you for free and open access by Scholarship@Western. It has been accepted for inclusion in Electronic Thesis and Dissertation Repository by an authorized administrator of Scholarship@Western. For more information, please contact [email protected]. Abstract

Titan is one of the most dynamic moons in the solar system. It is smaller than Earth and much colder, yet Titan is eerily similar to Earth, with rivers, rain, and seas, as well as sand seas that wrap around the equator. However, the rivers are made of hydrocarbons rather than water and the sand made of organics rather rock. We can use Titan’s impact craters to study how these processes modify the surface by comparing the craters depths, diameters and rim heights of Titan’s craters with fresh craters. Therefore, we have used the complete data set from NASA’s Cassini mission to update Titan’s crater population finding 30 new craters (90 total). We find that Titan’s craters are statistically shallower than those observed on similarly sized icy moons (e.g. Ganymede). We suggest this is due to sand and sediment infilling into the crater and fluvial erosion of the rims.

Keywords

Titan, Cassini-Huygens, Impact Cratering, Dragonfly, Geomorphology, Fluvial, Aeolian, Radar,

Co-Authorship Statement

Chapter 1 and 3 were written by J. Hedgepeth with edits and suggestions from Dr. C. Neish. Chapter 1 is a literature review of previous work. Chapter 3 is a review of the work presented here along with limitations, assumptions, and future work.

Chapter 2 is a modified article set to be submitted to Icarus for a publication tentatively titled “Impact Craters on Titan: Finalizing Titan's Crater Population”. Co-authors are Dr. Catherine Neish, Dr. Elizabeth Turtle, and Dr. Bryan Stiles and the Cassini RADAR Team. Dr. B. Stiles produced the SARTopo data. All RADAR data was collected by J. Hedgepeth, and all the data processing and results were produced by J. Hedgepeth. Dr. C. Neish, Dr. E. Turtle, and Dr. B. Stiles contributed to the interpretations and structure of the article as well as editorial suggestions.

Acknowledgments

Thank you to Dr. Catherine Neish for her great guidance and advice throughout this project. Thanks also for your patience and your encouragement for a healthy work life balance. Thank you to everyone in the Neish research group for all your support and aid. Thanks particularly to Alyssa Werynski, as the only other student in the group in the Titan field. Thanks to Dr Elizabeth Turtle for the guidance and for welcoming me into the wonderful Titan community. Thanks to Dr. Brittany Schmidt for suggesting I apply to Western, and thanks for your continual support as I grow as a young scientist. Without you I probably would not have made it this far. Thanks to my friends for the emotional support.

Thanks to Bryan Stiles for the awesome topography data that, while frustrating at times, has been the basis for the most interesting part of my research. Thank you to the Technologies for Exo-Planetary Science (TEPS) CREATE program for the funding and opportunities to broaden my background to include exoplanetary research and for the opportunities to network in this early stage of my career. Thank you to my committee for reviewing and critiquing my work. Thank you to my department and particularly to Dr. Gordon Osinski for two awesome classes that gave me the opportunity to study multiple terrestrial analogs that we see on Titan.

iii

Table of Contents

Abstract ...... i

Co-Authorship Statement ...... ii

Acknowledgments ...... iii

Table of Contents ...... iv

List of Tables ...... vii

List of Figures ...... viii

List of Appendices ...... xiii

Chapter 1 ...... 1

1 Introduction ...... 1

1.1 The Saturnian System ...... 1

1.2 Titan ...... 4

1.2.1 Titan’s Atmosphere and Origin ...... 6

1.2.2 Titan’s Surface ...... 10

1.3 Impact Cratering ...... 13

1.3.1 Formation ...... 13

1.3.2 Crater Morphology ...... 15

1.3.3 Crater Morphometry ...... 16

1.3.4 Titan Craters ...... 19

1.3.5 Crater Degradation ...... 23

1.3.6 Significance ...... 24

1.4 Radar Remote Sensing ...... 24

1.4.1 Surface Properties...... 25

1.4.2 Cassini RADAR and Related Data Products ...... 26

1.4.3 Cassini VIMS and ISS Instruments ...... 28

1.5 Summary ...... 29

1.6 References ...... 30

Chapter 2 ...... 41

2 Identifying Impact Craters on Titan and Constraining their Morphology (i.e. Diameter, Depths, and Rim Heights) ...... 41

2.1 Introduction ...... 41

2.2 Methodology ...... 43

2.2.1 Data Gathering and Processing ...... 43

2.2.2 Crater Mapping ...... 43

2.2.3 Crater Topography ...... 47

2.3 Results ...... 53

2.3.1 Final Crater Population ...... 53

2.3.2 Crater Counts ...... 61

2.3.3 Crater Topography ...... 63

2.4 Conclusion ...... 71

2.5 References ...... 73

Chapter 3 ...... 78

3 Conclusions ...... 78

3.1 Titan’s Craters ...... 78

3.2 Limitations ...... 80

3.2.1 Assumptions in Crater Mapping ...... 80

3.2.2 Radar Coverage ...... 81

3.2.3 Comparing Titan with Callisto and Ganymede ...... 82

3.2.4 Topography Data ...... 82

3.2.5 Methodological Assumptions ...... 83

3.2.6 Statistical Assumptions ...... 83

3.3 Future Work ...... 83

3.3.1 Landscape Evolution Models ...... 84

3.3.2 Comparing Radar Images and Erosion ...... 84

3.3.3 Updated ISS Map ...... 85

3.3.4 Updated SARTopo ...... 85

3.3.5 Measuring Stereo Rim Heights ...... 86

3.3.6 Dragonfly ...... 86

3.4 References ...... 87

Appendices ...... 90

Curriculum Vitae ...... 110

List of Tables

Table 2.1: The complete list of Titan's craters. Craters are classified by ID, ordered from largest to smallest diameter, with an asterisk indicating available SARTopo data. The crater names are defined by the publication they are from or the name they have been officially assigned. Unnamed craters are labeled with a W (Wood et al. 2010), NL (Neish and Lorenz 2012), and H2018 for those presented here...... 56

Table 2.2: Titan’s crater count corrected for incomplete coverage...... 62

Table 2.3: A list of Titan's craters with SARTopo and Stereo topography data available and their Rim Heights, Rim Depths, and Terrain Depths with the rim depths compared to relative depths (Rr)...... 65

vii

List of Figures

Figure 1.1: A modified image of the Saturnian System from NASA/JPL- Caltech. Permissions given from NASA...... 2

Figure 1.2: An image of the Cassini-Huygens spacecraft highlighting some of the major instruments on it, from NASA/JPL- Caltech. Permissions given from NASA...... 3

Figure 1.3: The timeline of the Cassini mission showing the number of flybys of each satellite by year. In 2017, the mission concluded with 22 flybys between Saturn and its rings (NASA/JPL-Caltech). Permissions given from NASA...... 4

Figure 1.4: Contrast enhanced image of Titan taken by Voyager 1 (Smith et al., 1981). Permissions given from Science...... 5

Figure 1.5: A temperature profile for Titan’s atmosphere with some of the major chemical processes observed within it. On the left are the depths over which each of the Cassini-Huygens instruments observed Titan (Hörst, 2017). Permissions given from Journal of Geophysical Research: Planets...... 6

Figure 1.6: A schematic view of Titan's methane cycle from Lunine and Atreya (2008). Timescales are rough estimates. Methane clathrates (green) or aquifers are a possible source of methane to maintain the atmosphere. It cycles from pole to pole to form the lakes (blue). As it cycles, humidity increases and allows for periodic rainfall (red), but methane is destroyed by radiation while organic hydrocarbons (purple) are formed and settle out to potentially form the dunes (orange). Permissions given from Nature Geoscience...... 7

Figure 1.7: A theoretical evolution of Titan's interior (b) and how it has affected the outgassing of methane throughout Titan’s history (a). The outgassing rate is controlled by the interior evolution (a). The final outgassing event is shown with 10% and 50% of the methane reservoir outgassed since the second outgassing event. This is modified from Tobie et al. (2006). Permissions given from Nature...... 8

Figure 1.8: Images of Titan's surface taken by the Huygens Descent Imager/ Spectral Radiometer highlighting the range of fluvial networks at the Huygens landing site (star) at

viii

600m elevation (C), and at the surface (D). Images A and B are ~8km and ~6km wide respectively. From Burr et al. (2013). Permissions given from Geologic Society of America...... 11

Figure 1.9: ISS Base map of Titan overlaid by radar-measured dune orientations (white vectors). (Lorenz and Radebaugh, 2009). Permissions given from Geophysical Research Letters...... 12

Figure 1.10: Schematic of the formation of simple (left) and complex (right) craters modified from Osinski et al. (2011). Permissions given from Earth and Planetary Science Letters. .... 14

Figure 1.11: Crater shape for a simple (A) and complex (B) shaped impact crater. Diameter becomes more difficult to constrain for complex (B) craters and especially for larger, multiringed craters. Modified from Osinski et al. (2011). Permissions given from Earth and Planetary Science Letters...... 16

Figure 1.12: Depth/diameter measurements for fresh impact craters on Ganymede. The thick lines are for lunar craters, and the thin lines are least-squares fits through the Ganymede data. Simple craters are solid dots, and complex craters are split into those with central peaks (open circles) and central pits and domes (crosses). The multiringed craters are shown with error bars. From Schenk (2002). Permissions given from Nature...... 18

Figure 1.13: Cassini RADAR images of craters on Titan having undergone different levels and types of degradation. (a) Crater #43 from Wood et al. (2010), D ~ 26 km. (b) Ksa, D ~ 39 km. (c) Momoy, D ~ 40 km. (d) Crater #49 from Wood et al. (2010), D ~ 60 km, (e) Soi, D ~ 78 km, (f), Sinlap, D ~ 82 km, (g) Hano, D ~ 100km, (h) Afekan, D ~ 115km, (i) Menerva, D ~ 425 km. Image from Neish et al. (2013). Permissions given from Icarus...... 19

Figure 1.14: Titan's crater population adjusted for limited coverage from Neish and Lorenz (2012) compared to past models (Artemieva and Lunine, 2005; Korycansky and Zahnle, 2005). Permissions given from Planetary and Space Science...... 20

Figure 1.15: Crater depths on Titan (●) compared to those on Ganymede (□, *) as a function of diameter. Squares are central peak craters and stars are central pits. These points are from

Bray et al. (2012) and the dashed trend line is from Schenk (2002). Plot from Neish et al., (2013). Permissions given from Icarus...... 23

Figure 1.16: Radar signals under different forms of backscatter. The roughness of the surface dictates the direction by which the radar signal will disperse. Smoother surfaces act like reflectors, directing the radar signal away from the sensor and producing minimal backscatter. Rough terrains scatter the signal in all directions, including back to the sensor increasing the power of the signal recorded. This is a general rule, but it is dependent on the incident angle (b) where all types of surface will reflect back more for lower incident angles (Neish and Carter, 2014). Permissions given from Elsevier...... 25

Figure 1.17: The types of radar data taken by Cassini. Image from Cassini RADAR Users Guide (Stofan et al., 2012). See text for full details. Permissions given from NASA...... 27

Figure 1.18: An overview of the SARTopo technique. Top left: A pixel with designated range (r) and Doppler has a varying gain (G) as a function of height (h). Top Right: The returned backscatter power (P) to the antennae, shown in black, has increased ranges for higher heights that can be defined as some normalized backscatter G(r,h). Bottom: The gain can be calculated for a given height, done over a range of heights, to find the gain profile that best correlates with the observed power values of P. 1 and 2 relates to the overlap swaths. See more details in text (Stiles et al., 2009). Permissions given from Icarus...... 27

Figure 2.1: New Crater #9 H2018 (Certainty=2, D=69km), in Cassini RADAR. The dark interior of the crater is a signature of the smooth sediment that fills the crater floor (circled in red). The rougher bright ejecta blanket and rim (circled in blue) is another signature characteristic used to identify a crater on Titan...... 45

Figure 2.2: Four example craters mapped in this work to illustrate how certainties are assigned. They are described by crater number (top left), name if applicable (top right), size (bottom left, km), and certainty 1-4 (bottom right). Momoy (Crater #12 H2018) has a distinct ejecta blanket and rim and a smooth, dark crater floor with a circular shape. Crater #9 H2018 looks similar, but its ejecta blanket is less clear and more degraded. Crater #2 H2018 clearly lacks any evidence of an ejecta blanket. Finally, crater #11 H2018 lacks an ejecta blanket, and its rim is

x difficult to differentiate from the rough surrounding terrain that likely existed prior to impact. Furthermore, the circular shape is questionable...... 46

Figure 2.3: Selk crater (Certainty=1, D=80km) observed in Cassini RADAR data. It has been mapped using the geo-referencing identifed by Wood et al. (2010) and Neish et al. (2012) with a black dashed circle. The current geo-position of Selk crater is mapped with a red circle. Overlain on the RADAR image is the available SARTopo elevation data measured in meters...... 47

Figure 2.4: A) Sinlap crater (Certainty=1, D=88km) and B) Soi crater (Certainty=1, D=84km) in Cassini RADAR with SARTopo data overlaid on it. The crater center and rims are identified with a black 'x'. Below is the topographic profile that goes through the crater with the regional slope subtracted out, plotted along the distance (km) of the profile. The rim to floor depth (dr) shown in black. The terrain to floor depth (dt) shown in green. Sinlap represents the topography measurements of a well-preserved crater compared to the heavily degraded Soi crater where dt~0km. More details in text. All of the crater topography data is shown in Appendix A. .... 49

Figure 2.5: Soi crater as observed through stereo data (overlaid on a RADAR image) (Kirk et al., 2012)...... 51

Figure 2.6: Thirty new craters have been identified and are shown in Cassini RADAR data. Craters are listed in descending diameter (D, bottom left). The certainty is denoted by C (bottom right). The Crater # for the H2018 study are shown in the top left. Official names are at the top right (where applicable)...... 54

Figure 2.7: A global Cassini RADAR mosaic of Titan mapping the 30 new craters (red) and the original 60 (yellow)...... 55

Figure 2.8: The correct crater count that Cassini RADAR would have detected based on its size given the total surface coverage through the end of the Cassini Mission (T126 flyby) compared to Neish and Lorenz (2012) (T65 flyby)...... 63

xi trend line that changes at ~35km, but we show it extended because the transition is gradual. b) Rim heights of Ganymede craters as measured by Bray et al. (2012) (data black diamonds, trend black dashed line) compared to those measured on Titan (red circles). As with Schenk (2002) rim to floor depths, the dashed lines are trend lines fit to the data over a range of crater diameters. However, the largest craters (~>50km) are shown as the upper and lower boundaries because the trend line breaks down. c) Crater depths, measured from local terrain of Ganymede craters (Bray et al., 2012 black diamonds) compared to Titan craters (red circles)...... 67

xii

List of Appendices

Appendix A: The SARTopo data for each crater studied overlain on a radar image of the region, plotted along latitude and longitude. The individual topographic profile(s) is/are identified on the radar image and ploted along the profile distance (km) with the regional slope subtracted out. Multiple profiles are numbered for reference...... 90

xiii 1

Chapter 1 1 Introduction

Saturn’s moon Titan is one of the most dynamic bodies in the solar system. Titan is unique because it has a thick atmosphere that extensively modifies the surface through erosional and depositional processes like those on Earth. These processes degrade the surface, and we observe this with impact craters on both Titan and Earth. The elevated crater rims erode while sand and sediment fill the crater floors. On Titan, the infill is made of fine grain organics that appear smooth, or radar dark. The elevated rims and surrounding ejecta blanket of excavated material is rough, appearing radar bright, until it becomes obscured through erosion and burial. Here we continue the broad study of impact craters on Titan by compiling a complete list of all the craters that were observed on Titan following the conclusion of the Cassini mission. We use the Cassini RADAR and Imaging Science Subsystem (ISS) data to identify and catalogue Titan’s impact craters, and characterize the morphometry of the craters (depth, diameter, and rim height) using the SARTopo dataset (Stiles et al., 2009). These observations will help to constrain the extent of the exogenic processes occurring on Titan and improve our understanding of how Titan’s surface is evolving.

1.1 The Saturnian System

Saturn has mesmerized astronomers since prehistoric times. Its size makes it one of the easiest planets to observe, yet it was not until 1610 (aided by the invention of the telescope) that Galileo Galilei was able to make the first view of its rings (Galilei, 1613; Galileo and Christoph, 2010). It was nearly 50 years later that Christiaan Huygens realized the rings were just that—rings—rather than moons (Huygens, 1659). In the early 1980s, the Voyager spacecraft gave us our first close look at Saturn (Smith et al., 1982, 1981). After Voyager, Saturn went from being a simple planet to being its own mini solar system (Figure 1.1), and it inspired NASA’s inevitable return to Saturn with the Cassini-Huygens mission.

Figure 1.1: A modified image of the Saturnian System from NASA/JPL- Caltech. Permissions given from NASA.

Cassini (Figure 1.2) launched in October of 1997 and arrived at the Saturnian system in July 2004. The primary objective was to understand Titan and the unique nature of its atmosphere (European Space Agency, 1988; Matson et al., 2003). Saturn (atmosphere, magnetosphere, etc.) was the second highest priority followed by its rings other moons. The Cassini orbiter was equipped with 12 instruments to study the Saturnian system and the Huygens probe was equipped with 6 instruments to study Titan specifically (European Space Agency, 1988; Lebreton and Matson, 2003). The primary imagers captured data in the infrared with the Imaging Science System (ISS), in visible and infrared with the Visible and Infrared Mapping Spectrometer (VIMS), and active microwave with the Radio Detection and Ranging (RADAR) instrument (Brown et al., 2004; Elachi et al., 2004; Porco et al., 2004). Cassini observed Saturn and its moons for nearly a decade and a half before it concluded its mission in September of 2017. In total, 127 flybys were done of Titan over that time (Figure 1.3).

Figure 1.2: An image of the Cassini-Huygens spacecraft highlighting some of the major instruments on it, from NASA/JPL- Caltech. Permissions given from NASA.

1.2 Titan

Titan was discovered by Huygens in 1656 (Huygens, 1656). In the mid-1900s, Kuiper (1944) used spectroscopy to confirm the existence of Titan’s atmosphere and the significant presence of methane. Scientists proceeded to consider what this meant for Titan’s surface and climate with some predicting a global ethane ocean and a strong greenhouse effect (Sagan, 1973; Lunine et al., 1983). During its flyby of the Saturn system, Voyager got the first close up look at Titan, but it could not see through Titan’s thick atmosphere (Figure 1.4; Smith et al., 1981).

Figure 1.4: Contrast enhanced image of Titan taken by Voyager 1 (Smith et al., 1981). Permissions given from Science.

Eventually, direct radar sensing would reveal Titan’s surface to be too bright to be covered by a global hydrocarbon ocean. Soon after, radiometry revealed the emissivity to be more consistent with an icy surface as well (Muhleman et al., 1990, 1995). Albedo maps from infrared imaging of the Hubble Space Telescope identified one of the first geologic features, that would later become known as Xanadu, from its abnormally high albedo (Smith et al., 1996). Arecibo radar images suggested the presence of localized bodies of liquid near the equator (Campbell et al., 2003). Cassini would later prove Campbell et al. wrong, that the liquids were at the poles not the equator, but these findings sparked scientists to consider the forms of landscape evolution Cassini-Huygens would observe. Most predictions still underestimated the level of degradation Cassini would observe because methane is far less effective at eroding water ice than water is of rock on Earth due to the lack of chemical weathering (Collins, 2005; Lorenz and Lunine, 1996).

One of our best views of the surface came in January 2005 when the Huygens probe landed on the surface of Titan and saw a terrain covered with evidence of fluvial activity. It studied Titan’s atmosphere during its decent measuring aerosol and cloud properties, the chemistry and composition of the atmosphere and ionosphere, as well as the winds and temperatures (Lebreton et al., 2005). Similar studies were done to constrain the physical

6 properties and composition of the surface finding it to be similar to wet clay at least at the near surface (Zarnecki et al., 2005; Tomasko et al., 2005; Niemann et al., 2005).

1.2.1 Titan’s Atmosphere and Origin

While it was unable to image Titan’s surface, the Voyager spacecraft did perform measurements of Titan’s atmosphere (Lindal et al., 1983). It was found to consist mostly of nitrogen (98%) with a large percentage of methane (~2% in the stratosphere) (Coustenis, 2014; Tomasko et al., 2005; Lindal et al., 1983). Solar and other cosmic radiation is constantly ionizing and dissociating Titan’s atmosphere; these particles recombine and form heavy complex organic compounds (tholins) (Figure 1.5; Waite, 2005; Tomasko et al., 2005; Hörst, 2017). These are the source of Titan’s haze, layers of large molecules that scatter visible light and prevent it from reaching the surface.

Titan’s abundance of methane and organic compounds results in a complex cycle of rain, erosion, and deposition. Titan’s surface pressure is very similar to Earth’s (1.5 bars vs 1 bar), and with temperatures at 94K, methane can be stable as both a gas and liquid (Kouvaris and Flasar, 1991). Titan’s methane isn’t in an ocean like water on Earth, rather it is in large lakes, and the specific heat (the energy needed to raise a substance by 1K) is harder to overcome because of the lower solar flux (Lunine and Atreya, 2008). Nevertheless, rainfall and cloud events have been observed, if rarely (Turtle et al., 2011;

Porco et al., 2005; Griffith et al., 1998). Modeled rainfall requirements for the observed channels suggest that there should be twice as much methane vapor in the atmosphere than observed at the equator (Lunine and Atreya, 2008). Short term (100s of yrs) rainfall could be fueled by evaporation of the lakes, but methane is slowly being destroyed over a much longer lifespan of 10-100 Ma (Figure 1.6). Therefore, it has been suggested that outgassing event(s) have released methane from the interior in Titan’s past to resupply its atmospheric methane (Figure 1.7) (Choukroun et al., 2010; Choukroun and Sotin, 2012; Tobie et al., 2006).

Tobie et al. (2006) present a theoretical evolution of Titan’s interior that composes of three major outgassing events. The first event begins with the quick overturn of Titan’s

9 initial core. The outpouring of methane produces a thick layer of methane clathrate above the ocean because of how methane interacts with liquid water under Titan’s temperatures and pressures (Sloan, 1998; Tobie et al., 2006). Clathrates are ices with compounds that are trapped within the ice lattice, and have different rheological and thermal properties compared to ice (Durham et al., 2010). The low viscosity and low conductivity acts as an insulator, warming the ocean which releases methane. After differentiation, the silicate core begins to convect; the heat flux thins the clathrate layer with the buoyant methane accumulating at the base and escaping through cracks (Lunine and Stevenson, 1987; Tobie et al., 2006). As the interior cools, the liquid ocean begins to freeze to form a layer of ice I. The thick ice layer convects, forming warm plumes (likely from tidal dissipation; Sotin et al., 2002) that break through the clathrate layer making it unstable. Tobie et al. (2006) suggest the last stage occurred perhaps ~500 Ma. This model explores a range of parameters but finds that these changes only alter the length and intensity of these cycles. Furthermore, this model is consistent with isotopic signatures in the atmosphere.

Isotopic measurements help to verify the outgassing events but also constrain the origin and evolution of Titan’s volatiles. Evidence of 40Ar in Titan’s atmosphere reflects the decay of 40K which would have been sourced from rock-water interactions. Its existence in the atmosphere suggest potassium rich water magmas reached the surface through volcanism likely fueled by the outgassing (Niemann et al., 2005; Tobie et al., 2006; Waite, 2005). Another line of evidence comes from the lack of enrichment of heavier carbon isotopes in methane despite seeing it in nitrogen isotopes. The 15N/14N ratio is enhanced as heavier 15N sinks below the 14N which can escape more readily (Hidayat et al., 1998; Lunine et al., 1999). Inversely, the ratio of 13C/12C in hydrocarbons is closer to the terrestrial value suggesting it has not undergone the same escape enrichment. That is to say, the carbon in the hydrocarbons (i.e. methane) has been recently sourced (~1 Ga or less). Congruently, the deuterium in methane is lightly enriched (~1.5 times) which is still significantly lower than organic molecules in early solar nebula (i.e. the outer solar system) (van Dishoeck et al., 1993). This reflects the slightly warmer circum-Saturnian nebula that would have undergone more processing of the volatiles that likely sourced Titan’s reservoir (Lunine and Tittemore, 1993).

1.2.2 Titan’s Surface

The initial results from Cassini and the Huygens probe definitively proved that the surface of Titan is not covered with a global ocean of hydrocarbons. At the Huygens landing site, the surface was inferred to be largely a “bedrock” of water ice (with some impurities) (Tomasko et al., 2005). However, most of the planet is covered with a ~3cm layer of small organic haze particles (Atkinson et al., 2010; Janssen et al., 2016). Zarnecki et al. (2005) used the deacceleration of the probe to show that the surface was neither hard like solid ice nor soft like a fluffy blanket of snow (Seiff et al., 2005). Rather, it was shown to be more like wet clay. The subsurface was more difficult to constrain at the Huygens landing site, but it was best explained by damp sand that gets wetter with depth. It is important to recognize that this does not mean the entire surface is clay-like, but we can use this site as a form of ground truth for better constraining Titan’s entire surface using global remote sensing data sets to. However, this may speak to the existence of subsurface aquifers or methane layers, or it may be an indicator of the methane clathrates in the crust (Choukroun et al., 2010; Tobie et al., 2006).

Near the Huygens landing site, river channels flow from highlands to flatter lowlands (Figure 1.8; Tomasko et al., 2005). Tomasko et al. (2005) also showed that these highlands are rough (changes of 10s to nearly 100 meters), likely due to the complex channels cutting through them. The landing site is in the much flatter lowland region where a likely flood plain has left behind heavily eroded (rounded) pebbles of water ice on the order of 10s of cms (Lorenz, 2006; Tomasko et al., 2005). Towner et al. (2006) later used acoustics data from the Huygens descent to corroborate these results, showing that the region had large depressions in places but was smooth on the cm scale as you would expect from fluvial activity.

In fact, fluvial networks may cover more than 10% of Titan’s surface (Burr et al., 2009; Lorenz et al., 2008). Burr et al. (2013, 2009) detailed the full range of fluvial activity that was observed on Titan from narrow to wide to stubby channels with a range of bright and dark radar albedos. Recent results show that some channels are 100s of meters deep with liquid present (Poggiali et al., 2016). Overall, the existence of these features presents a contrasting scenario where Earth-like processes are occurring under completely different

11 physical conditions. Intriguingly, the erosion rates on Titan appear to be within the range of those seen on Earth, but these estimates are poorly constrained (Collins, 2005).

Figure 1.8: Images of Titan's surface taken by the Huygens Descent Imager/ Spectral Radiometer highlighting the range of fluvial networks at the Huygens landing site (star) at 600m elevation (C), and at the surface (D). Images A and B are ~8km and ~6km wide respectively. From Burr et al. (2013). Permissions given from Geologic Society of America.

Another important component of Titan’s surface likely comes from the organic particulates that fall out of Titan’s atmosphere and cover the surface (Janssen et al., 2016; Lorenz, 2006). These particles are very fine grained (

(Collins, 2005), nor are they large enough to create sand dunes because saltation (formation) on Titan would require particles 100s microns in size (Lunine and Lorenz, 2009). Nevertheless, we observe sand dunes on Titan made of what appear to be the same organics as those falling out of the atmosphere using compositional information from VIMS and radiometry data (Le Gall et al., 2011; Lorenz, 2006; Soderblom et al., 2007). These dunes reach across large distances, wrapping around Titan’s equator (Figure 1.9) (Barnes et al., 2015), so these dunes must be sourced from somewhere. Barnes et al. (2015) showed that this can happen through sintering, the fusing of particles through burial and lithification. However, other ways exist like fusing of particles suspended in liquid (flocculation) or through concentrated evaporites.

Figure 1.9: ISS Base map of Titan overlaid by radar-measured dune orientations (white vectors). (Lorenz and Radebaugh, 2009). Permissions given from Geophysical Research Letters.

Prior works have suggested sintering as the cause (Barnes et al., 2008) because large lakes of methane in the north and south poles provide the right environment for it to occur (Stofan et al., 2007). The existence of these lakes and paleolakes in the same region create a wet dry cycle that would be needed to sinter the fine haze particles (Grotzinger et al., 2013). It shows that sintering is a viable option, but the mechanical and thermal

13 properties of organics under Titan conditions is still poorly constrained (Barnes et al., 2008).

1.3 Impact Cratering 1.3.1 Formation

Impact craters form when large projectiles collide with a planetary body at hyper velocities (10s of km/s) (Gault et al., 1968). If the planetary body has an atmosphere, impact craters will only form if the projectiles are large enough (>10s of meters) to pass through it (French, 1998; Osinski and Pierazzo, 2012). These events impart high levels of energy in a very short amount of time, compressing and excavating the target material. The first stage in an impact cratering event is the compression stage. The compression stage is when the kinetic energy of the projectile is converted to internal energy (i.e. shock waves) that propagate through the surface (Grieve et al., 2014; Melosh, 1989, 2011). Impacts can generate hundreds of GPa of pressure that is converted to heat and transport energy. The region of highest shock vaporizes the bedrock and projectile, and transitions to melting, then brittle damage as the pressures diminish with distance from the impact site (Osinski et al., 2012).

Figure 1.10: Schematic of the formation of simple (left) and complex (right) craters modified from Osinski et al. (2011). Permissions given from Earth and Planetary Science Letters.

The shock wave degrades in strength as it expands outward, becoming a stress wave; this is because the energy covers a larger area (i.e. it is less concentrated) and is being expended through heat, phase change and material transport (Melosh, 1989; Osinski et al., 2012). In a matter of seconds, the excavation stage begins where target material is moved down and out, driven by the cratering flow-field (Figure 1.10; Osinski et al., 2012, 2011). The ejecta is formed by a mix of impact melt and breccias (excavated bedrock material) that is deposited up to 2 radii out, being thickest at the rims. These are rough and blocky

15 due to their clastic nature. This continues until there is not enough energy left to overcome the force of gravity keeping the material in the crater. The final stage of crater formation is the modification stage. Material stops flowing out of the crater and is gravitationally pulled back down (Kenkmann et al., 2012).

Once the formation is completed, the crater may continue to change overtime. Larger complex craters experience significant crater collapse, and weaker target rock (e.g. icy bodies) will viscously relax on the order of 10s to 100s Ma (Melosh, 1989). Viscous relaxation is like a glacier flowing downhill; gravity drives it down to an equilibrium state. This isn’t observed on rocky bodies because the viscosity is too high and even colder temperatures can sufficiently strengthen ice to significantly dampen the effect of relaxation. For example, on Titan there is insufficient energy or heat for much relaxation to occur (Schurmeier and Dombard, 2018). Sand deposition can produce added insulation and moderately enhance the process, but even then, the timescales are too long to compete with the much faster fluvial and depositional processes (Neish et al., 2016, 2013).

1.3.2 Crater Morphology

Crater morphology is primarily controlled by the crater size and the planetary properties (i.e. size, composition, structure) (Melosh, 1989). Craters form two well constrained shapes (Figure 1.11): 1) Simple (bowl like) craters with parabolic profiles. 2) Complex craters with central uplifts and significant wall slumping. However, it also forms a third, less constrained shape, 3) Multi-ringed basins which are the largest craters in the solar system. These craters are so large that the walls completely collapse erasing most evidence of a rim in favor of rings where the walls slid downward.

These shapes are idealized. Crater morphologies involve some variance, and transitional morphologies exist from one phase to another. However, large scale analysis of craters shows that impact crater morphologies follow distinct trends that are relatively similar for different planetary bodies. However, different target materials (e.g. rock vs ice) exhibit different depths and transitions between crater morphologies because of the material strength (Schenk, 2002). Even worlds of similar sizes and compositions can vary.

1.3.3 Crater Morphometry

The most descriptive characteristic of a crater is its diameter. The diameter of a crater can be used to predict the approximate depth and overall shape of a crater. The material properties of the planetary body also matter (see above), but the transition point from simple to complex to multiringed craters is unique for each body (Bray et al., 2008; Schenk, 2002, 1989). However, some planets are sufficiently similar to compare. Gravity is the driving factor behind the simple-complex transition diameter, but lithospheric structures

17 also influence this transition because they can be indicative of the thermal structure in the crust (Melosh, 1989; Schenk, 2002; Turtle and Pierazzo, 2001). The thermal structure controls how material will react to the force of a shock wave; warmer material tends to act more ductile (plastically deforming) than brittle (faulting) (Turcotte and Schubert, 2014). This is especially important on icy moons where temperature is the driving factor in viscous relaxation because it controls the viscosity of the material (Durham et al., 2010; Schurmeier and Dombard, 2018).

Therefore, it is important to accurately characterize the diameter of a crater to accurately constrain the global trend of crater morphologies with size, known as scaling laws (Figure 1.12). Schenk (2002) discusses how scaling laws change for the icy moons of Jupiter. Compared to the Moon, the simplest craters follow the same trend, but the transition to complex craters occurs much sooner because of the differences in material strength. For similar reasons, these icy moons reach a point where the complex craters (with central peaks) become central pit craters. It has been suggested that pits may form from weak ice collapse of the rapid release of volatiles (Carr et al., 1977; Passey and Shoemaker, 1982). However, the recent discovery of central pits on the Moon and Mercury suggests it does not require volatiles (ice) (Xiao and Komatsu, 2013), and new models suggests it may be melt drainage through fractures that occurs more easily with water than magma (Elder, 2015). The effect of size and strength on crater morphology is observed best with Europa which is embedded with significant tidal heating that thins its icy crust. Therefore, the transitions occur sooner, but worlds of similar size, thermal structure, and material strength often have very similar scaling laws (e.g. Ganymede and Callisto).

Titan, like Earth, undergoes a great deal of degradation (see Section 1.2 and 1.3.5), so the crater structure becomes harder to interpret. Terrestrial (and Titan) craters undergo significant degradation which make it difficult to find the crater rim even with topography (Figure 1.13). Without a population of fresh craters, it is difficult to develop a unique scaling law. Multiringed basins are similarly difficult to characterize even before erosion. In these cases, the best approach is to clearly identify the ambiguity required when interpreting a crater. An example of this is given in Schenk (2002), showing multiringed basins with error but not a trend line fit to it.

1.3.4 Titan Craters

There was limited information about Titan’s impact craters prior to the Cassini mission. Results leading up to Cassini’s arrival at Titan had only just begun to address the problem of a wet versus dry surface, and that made it difficult to consider the types of craters we would expect to see there. Engel et al. (1995) considered the effect of the atmosphere on crater production to try to constrain the history of Titan’s atmosphere using crater size distributions. Later models attempted to constrain cratering rates on Titan taking into account the influence of the atmosphere on crater production (Figure 1.14) (Artemieva and Lunine, 2005; Korycansky and Zahnle, 2005). They assumed most impactors are ecliptic comets (short period comets) and compared different sized impactors modeled after craters seen on Jupiter’s moons versus on Triton (Zahnle et al., 2003). A major factor in impactor flux is the gravity in the region (Shoemaker and Wolfe, 1982). This increases the speed and energy of incoming impactites. Saturn’s gravity is lower than Jupiter’s. It is also further from the sun, but that has a negligible effect on the impact flux. Perhaps more importantly, Titan’s neighbors are significantly smaller than Galilean siblings, so their

20 effect on the impact flux at Titan is likely minimal. Furthermore, Titan will likely have higher impact rates than the inner satellites because smaller orbits create faster rotations, making it harder to impact with. These models predicted a distinct drop-off in craters at diameters less than 20 km, but it was moderately dependent on the impactor sizes and how rich the impactor was in water ice (Korycansky and Zahnle, 2005).

Cratering rates on active worlds (e.g. Earth, Mars, Titan) where the impact crater population is significantly modified cannot be determined based solely on what we observe. It is hypothesized that cratering rates are split into two main “Populations” in the Saturn system. Some have suggested that these are a result of impacts from comets outside vs. inside the solar system while others suggest it is to do with impacts from large, accretionary material formed in the Saturn subnebula and later, smaller long-period comets

(Plescia and Boyce, 1985; Smith et al., 1981). Simulations have been done to model the orbits of comets to predict likely cratering rates (Bottke et al., 2002; Levison and Duncan, 1997) and combined with direct observations of impacts into giant planets like Jupiter (Zahnle et al., 1998). Zahnle et al., (2003) approximates the probability of impact as it relates to Jupiter’s impacting rate by using a Monte Carlo algorithm. Zahnle et al. (2001) developed this approach using comet planetary orbital dynamics to simulate the likelihood of impact. They considered two cases, one where Titan cratering rates are derived from observations of Triton and one derived from observations of Jupiter. Jupiter has a larger gravity effect on the acceleration of impactors, so it is likely to produce fewer smaller craters. Figure 14 (a) shows the case relative to Jupiter. The Triton estimate relies on smaller bodies from the Kuiper belt and as such would create more smaller craters. Artemieva and Lunine (2005) uses the modern flux of elliptical comets in the Kuiper belt to determine a rate, assuming a change in flux through time (Figure 14b). After counting the observed number of craters on Titan, we can predict the age using the rate per year (number of craters divided by the rate). Neish and Lorenz (2012) suggest a surface age of 200 Ma to 1 Ga, with the large variation due to differences in cratering rates and crater scaling laws used in the two models described above.

1.3.4.1 Craters Observed on Titan

As Cassini RADAR coverage of Titan grew, the population of known craters grew. The first crater assessment only reported three known craters based on RADAR coverage of 10% of the surface (Lorenz et al., 2007). Wood et al. (2010) updated the list with data from 2004 to 2007, creating a catalog of 49 craters, rating them from certain to nearly certain to probable. Identifying craters has proven difficult due to the limited (and low resolution) data on Titan, coupled with erosion and deposition altering the crater morphologies. The most recent catalog of Titan’s crater population had a total of 59 craters (Neish and Lorenz, 2012). Buratti et al. (2012) raised that to an even 60 when they identified a crater with the Cassini VIMS instrument that was not in any overlapping RADAR data. The age of the surface (crater retention rate) was judged to be between 0.2 and 1.0 Ga (Neish and Lorenz, 2012). Menrva (~425km) is the only very large crater (D>125km) on Titan (Neish et al., 2013), and thickening of the lithosphere would suggest

22 very large craters could not have formed when the crust was thinner (Tobie et al., 2006). However, craters of this size are rare, so rare that it is unlikely that another very large crater would have formed over the lifespan of the solar system (Neish and Lorenz, 2012).

The extent of erosion and infill in Titan’s craters is easily observed in the radar imagery (Figure 1.13). Smoother sediments have infilled crater centers and river channels cut through crater rims and walls (Neish et al., 2013). The limited topography data highlights how Titan’s craters are significantly shallower than those on other similarly sized icy bodies (Figure 1.15; Neish et al., 2013). We expect Titan’s craters to have the same morphology as those on Ganymede when they first form. Titan’s overall structure is like Ganymede’s. In a way, Ganymede’s craters are like a template for what Titan’s craters would look without the exogenic processes acting on them. Ganymede’s bulk density is 1940 kg/m3 with a radius 2634 km while Titan has a density of 1880 kg/m3 and a radius of 2575km (Buratti and Thomas, 2014). Not only that, Ganymede and Titan each have a thick ice crust (~100 km thick) covering a liquid water ocean and a deeper high pressure ice layer (Collins and Johnson, 2014; Mitri and Showman, 2008; Nimmo and Bills, 2010). Therefore, if we accept the assumption that the two are reasonably comparable, we can obtain quantitative constraints on how significantly erosion and infill affects the surface of Titan by comparing the Titan’s crater depths with fresh crater depths on Ganymede.

Figure 1.15: Crater depths on Titan (●) compared to those on Ganymede (□, *) as a function of diameter. Squares are central peak craters and stars are central pits. These points are from Bray et al. (2012) and the dashed trend line is from Schenk (2002). Plot from Neish et al., (2013). Permissions given from Icarus.

1.3.5 Crater Degradation

There are two primary modes of crater evolution. Evolution through endogenic processes such as viscous relaxation and tectonics is a long-term factor for icy bodies, but exogenic process such as erosion and deposition probably destroy Titan’s craters before relaxation can occur (Schurmeier and Dombard, 2018). Under the right conditions, however, viscous relaxation may play a moderate role in modifying the largest craters, especially if they are filled with a large amount of sand. Nonetheless, the main mechanism of landscape evolution of Titan’s craters is through infill of sand and to a lesser extent, fluvial erosion (Neish et al., 2016, 2013). The amount of organic infill varies depending on the crater location (Neish et al., 2015; Werynski et al., 2017). In the equatorial sand seas, craters are filled in by windblown sand, and sediment derived from the fluvial erosion of the rims. Infill will still occur in craters outside the equatorial sand seas through fluvial erosion, but the crater interiors here appear less organic rich in VIMS and radiometry data. For example, Soi crater (~78km) is heavily degraded, but it is less organic rich than other craters near the dunes (e.g. Shikoku) due to the relative lack of sand. Here, the organic

24 infill is a combination of organic-rich sediments deposited through fluvial erosion and infiltration by liquid methane.

The number of craters on Titan’s surface varies with latitude (Neish and Lorenz, 2014). The poles may have been saturated with liquid methane in the past, providing a marine environment in which impact craters do not form recognizable shapes. As a result, most of Titan’s craters are in the lower latitudes (Neish and Lorenz, 2014). However, there is also clear evidence of fluvial activity in the lower latitudes, suggesting fluvial processes play an important role here as well (Neish et al., 2016). In addition to distinct fluvial patterns in radar images, central uplifts are missing, suggesting fluvial modification by the moving of material from higher to lower terrain, or burial by infill (Neish et al., 2016)..

1.3.6 Significance

In this work, we continue the broad study of impact craters on Titan by updating Titan’s list of craters and updating the crater measurements at the end of the Cassini mission. Five years of new data is now available, and a systematic search for all of Titan’s craters has yet to be performed with the completed dataset. Therefore, we are creating a complete catalogue of all the craters on Titan and characterizing their morphology by measuring their crater depths and rim heights and then comparing them with fresh craters on Ganymede.

1.4 Radar Remote Sensing

Radar is an imaging technique that makes use of the transmission of radio waves that are returned as an echo to study the surfaces of different planets (Moreira et al., 2013; Neish and Carter, 2014). Radar is done actively (Figure 1.16), but microwave emissions can be used to passively study the surface emission (i.e. radiometry). The active nature of radar gives the user more control over the data being recorded and provides a variety of potential uses involving three basic techniques: 1) radar imagery (e.g. Cassini SAR data), 2) radar sounding (e.g. Carter et al., 2009), and 3) radar topography (e.g. Stiles et al., 2009). Synthetic aperture radar (SAR) images the surface at centimeter to decimeter scale wavelengths, while radar sounding emits a radio wave into the subsurface of a planet that

25 reflects the subsurface structure. Radar topography can be achieved through multiple means from measuring how long a signal takes to return (altimetry) to using stereo coverage through multiple fly overs.

1.4.1 Surface Properties

Radar data can provide information about surface properties that optical remote sensing does not (Neish and Carter, 2014). Typically, the radar albedo is most sensitive to surface roughness at the scale of the radar wavelength (typically centimeters to decimeters).

Smoother surfaces will have a more specular reflection that sends a signal away from the sensor producing a darker image while rougher surfaces produce more backscatter returning more signal to the receiver (Figure 1.16). However, albedo is also affected by surface properties such as the dielectric constant of the surface and the large-scale topography. We can begin to infer these properties using the radar imagery (Farr, 1993; Pettengill et al., 1997).

1.4.2 Cassini RADAR and Related Data Products

Cassini was equipped with a Ku-band (2.17 cm λ) radar instrument with 5 beams for collecting data (Elachi et al., 2004; Stofan et al., 2012). It was used to perform radiometry, scatterometry, altimetry and synthetic aperture radar (SAR) imaging (Figure 1.17). Radiometry occurred at the highest altitudes (100,000-30,000km) and produced polarized brightness temperatures that provide an array of information regarding Titan’s surface properties (e.g. temperature, dielectric constant, roughness at and below the surface, etc.). Scatterometry, or low-resolution altimetry (altitudes from 30,000-10,000km) gave a broad look at how backscattering varies over broad global terrains. Altimetry occurred at virtually the same altitudes as the SAR images (5,000-1,000km altitudes), providing topographic information to characterize Titan’s landforms and overall geoid (e.g. Mastrogiuseppe et al., 2014). The SAR mode captured the highest resolution images of the surface of Titan at 175 meters per pixel (Elachi et al., 2004; Lopes et al., 2010). These images were primarily sensitive to centimeter scale roughness. However, it did not monitor the signal polarization.

Figure 1.17: The types of radar data taken by Cassini. Image from Cassini RADAR Users Guide (Stofan et al., 2012). See text for full details. Permissions given from NASA.

In addition to altimetry, a relatively new technique, called SARTopo, was developed that provides topographic data with increased coverage than that available from the altimetry mode. Stiles et al. (2009) devised a way to back out the topography for most swatches of the SAR imagery (Figure 1.18). Each swatch is made with 5 antenna beams to gain a wider view of the surface. Where these beams overlap, the backscatter differences can be used to calculate the height given sufficient knowledge about how the spacecraft is operating. Alternatively, a height can be assumed to recalibrate the backscatter of two images to better align. Each SAR signal is sent in the form of 30-60 pulses. Each pulse is calibrated so they can be averaged into a single backscatter power section (Pi) across a set of range values (r). Then multiple normalized backscatter sections (G(r,h)) are produced at a range of ℎ = ±2 km over increments of 250m. A simple goodness-of-fit test for each G(r,h) will find which height value is the best fit.

The result is topography that is accurate to ~100m in the vertical direction with ~10km horizontal resolution. That is to say, the topography was averaged over a 10km width, but it was done in 300m increments to translate the topography in finer detail. Lorenz et al. (2011) compared Titan’s hypsometry (distribution of elevation) to other bodies in the solar system and found it to be much narrower with limited topographic changes suggestive of extensive erosion. Later, an interpolated global topographic map of Titan was made using a combination of SARTopo, stereo, and altimetry data, highlighting the limited coverage (~5% of the surface) of the SARTopo data set (Corlies et al., 2017; Lorenz et al., 2013). However, individual topographic profiles can be extremely helpful in constraining crater topography.

1.4.3 Cassini VIMS and ISS Instruments

Cassini RADAR is the primary data source used in this study. However, Cassini VIMS and ISS instruments offer complementary information about the surface. VIMS saw in the visual to near-infrared range (0.94-5.1 microns over 7 windows) with a variable resolution (Brown et al., 2004). Although it was acquired at lower resolution, VIMS has more coverage than Cassini RADAR and gives information about surface composition (e.g. Barnes et al., 2009). The ISS resolution is only ~4 km/pixel in the best cases, and it covers

29 a spectral range over visible to near infrared (0.2-1.1 microns) (Cassini ISS Team, 2015; Porco et al., 2004). It is similar to VIMS but was equipped with a series of settings and filters to enhance observations (e.g. Porco et al., 2005). The significantly lower resolutions make these datasets less ideal for crater mapping, but they each provide added context to what is observed in RADAR.

1.5 Summary

This thesis begins with a literature review of Titan, the missions that have visited it, and the process of impact cratering on Earth, Titan and other icy moons. The goal is to provide context to the work that we present in Chapter 2. In Chapter 2 we present the manuscript for “Impact Craters on Titan: A Final Assessment” where we have updated Titan’s list of impact craters, providing details about their morphology observed in radar images and the implications for Titan’s global crater population. Then we update past estimates of crater depths to quantify the level of degradation observed on Titan. There is current work to map the geology of Titan, including impact craters. However, that work is not a search to update and complete the crater population. This work is unique because there is no one else addressing this problem explicitly. We have expanded beyond past work by obtaining the first measurements of rim heights of impact craters on Titan. We conclude by discussing the implications of fluvial erosion and aeolian infill on Titan. Organics, mostly sourced from the dunes, infill craters and reduce their depth as methane fluvially erodes the raised rims. In Chapter 3, we summarize our findings and conclusions and discuss the limitations and assumptions of this work. We conclude by detailing the ways in which this work can be expanded or improved through future work.

1.6 References

Artemieva, N., Lunine, J.I., 2005. Impact cratering on Titan II. Global melt, escaping ejecta, and aqueous alteration of surface organics. Icarus 175, 522–533. https://doi.org/10.1016/j.icarus.2004.12.005

Atkinson, K.R., Zarnecki, J.C., Towner, M.C., Ringrose, T.J., Hagermann, A., Ball, A.J., Leese, M.R., Kargl, G., Paton, M.D., Lorenz, R.D., 2010. Penetrometry of granular and moist planetary surface materials: Application to the Huygens landing site on Titan. Icarus 210, 843–851. https://doi.org/10.1016/j.icarus.2010.07.019

Barnes, J.W., Brown, R.H., Soderblom, L., Sotin, C., Le Mouèlic, S., Rodriguez, S., Jaumann, R., Beyer, R.A., Buratti, B.J., Pitman, K., Baines, K.H., Clark, R., Nicholson, P., 2008. Spectroscopy, morphometry, and photoclinometry of Titan’s dunefields from Cassini/VIMS. Icarus 195, 400–414. https://doi.org/10.1016/j.icarus.2007.12.006

Barnes, J.W., Lorenz, R.D., Radebaugh, J., Hayes, A.G., Arnold, K., Chandler, C., 2015. Production and global transport of Titan’s sand particles. Planetary Science 4. https://doi.org/10.1186/s13535-015-0004-y

Barnes, J.W., Soderblom, J.M., Brown, R.H., Buratti, B.J., Sotin, C., Baines, K.H., Clark, R.N., Jaumann, R., McCord, T.B., Nelson, R., Le Mouélic, S., Rodriguez, S., Griffith, C., Penteado, P., Tosi, F., Pitman, K.M., Soderblom, L., Stephan, K., Hayne, P., Vixie, G., Bibring, J.-P., Bellucci, G., Capaccioni, F., Cerroni, P., Coradini, A., Cruikshank, D.P., Drossart, P., Formisano, V., Langevin, Y., Matson, D.L., Nicholson, P.D., Sicardy, B., 2009. VIMS spectral mapping observations of Titan during the Cassini prime mission. Planetary and Space Science 57, 1950–1962. https://doi.org/10.1016/j.pss.2009.04.013

Bottke, W.F., Morbidelli, A., Jedicke, R., Petit, J.., Levison, H., Michel, P., Metcalfe, T.S., 2002. Debiased orbital and absolute magnitude distribution of the near-Earth objects. Icarus 156, 399–433.

Bray, V.J., Collins, G.S., Morgan, J.V., Schenk, P.M., 2008. The effect of target properties on crater morphology: Comparison of central peak craters on the Moon and Ganymede. Meteoritics & Planetary Science 43, 1979–1992. https://doi.org/10.1111/j.1945-5100.2008.tb00656.x

Brown, R.H., Baines, K.H., Bellucci, G., Capaccioni, F., Cerroni, P., Clark, R.N., Coradini, A., Cruikshank, D.P., Drossart, P., Formisano, V., Jaumann, R., Langevin, Y., Matson, D.L., Mccord, T.B., Mennella, V., Miller, E., Nelson, R.M., Nicholson, P.D., Sicardy, B., Sotin, C., 2004. The Cassini Visual and INfrared Mapping Spectrometer (VIMS) Investigation. Space Science Reviews 115, 111–168.

Buratti, B.J., Sotin, C., Lawrence, K., Brown, R.H., Le Mouélic, S., Soderblom, J.M., Barnes, J., Clark, R.N., Baines, K.H., Nicholson, P.D., 2012. A newly discovered impact crater in Titan’s Senkyo: Cassini VIMS observations and comparison with other impact features. Planetary and Space Science 60, 18–25. https://doi.org/10.1016/j.pss.2011.05.004

Buratti, B.J., Thomas, P.C., 2014. Planetary Satellites, in: Spohn, T., Breuer, D., Johnson, T.V. (Eds.), Encyclopedia of the Solar System. Elsevier, Amsterdam ; Boston, pp. 759–777.

Burr, D.M., Jacobsen, R.E., Roth, D.L., Phillips, C.B., Mitchell, K.L., Viola, D., 2009. Fluvial network analysis on Titan: Evidence for subsurface structures and west-to- east wind flow, southwestern Xanadu. Geophysical Research Letters 36. https://doi.org/10.1029/2009GL040909

Burr, D.M., Taylor Perron, J., Lamb, M.P., Irwin, R.P., Collins, G.C., Howard, A.D., Sklar, L.S., Moore, J.M., Adamkovics, M., Baker, V.R., Drummond, S.A., Black, B.A., 2013. Fluvial features on Titan: Insights from morphology and modeling. Geological Society of America Bulletin 125, 299–321. https://doi.org/10.1130/B30612.1

Campbell, D.B., Black, G.J., Carter, L.M., Ostro, S.J., 2003. Radar Evidence for Liquid Surfaces on Titan. Science 302, 431–434. https://doi.org/10.1126/science.1088969

Carr, M.H., Crumpler, L.S., Cutts, J.A., Greeley, R., Guest, J.E., Masursky, H., 1977. Martian Impact Craters and Emplacement of Ejecta by Surface Flow. Journal of Geophysical Research 82, 4055–4065.

Carter, L.M., Campbell, B.A., Watters, T.R., Phillips, R.J., Putzig, N.E., Safaeinili, A., Plaut, J.J., Okubo, C.H., Egan, A.F., Seu, R., Biccari, D., Orosei, R., 2009. Shallow radar (SHARAD) sounding observations of the Medusae Fossae Formation, Mars. Icarus 199, 295–302. https://doi.org/10.1016/j.icarus.2008.10.007

Cassini ISS Team, 2015. Map of Titan, Cassini ISS Global Mosaic.

Choukroun, M., Grasset, O., Tobie, G., Sotin, C., 2010. Stability of methane clathrate hydrates under pressure: Influence on outgassing processes of methane on Titan. Icarus 205, 581–593. https://doi.org/10.1016/j.icarus.2009.08.011

Choukroun, M., Sotin, C., 2012. Is Titan’s shape caused by its meteorology and carbon cycle? Geophysical Research Letters 39, n/a-n/a. https://doi.org/10.1029/2011GL050747

Collins, G., Johnson, Torrence V., 2014. Ganymede and Callisto, in: Spohn, T., Breuer, D., Johnson, T. V. (Eds.), Encyclopedia of the Solar System. Elsevier, Amsterdam ; Boston, pp. 813–829.

Collins, G.C., 2005. Relative rates of fluvial bedrock incision on Titan and Earth. Geophysical Research Letters 32, n/a-n/a. https://doi.org/10.1029/2005GL024551

Corlies, P., Hayes, A.G., Birch, S.P.D., Lorenz, R., Stiles, B.W., Kirk, R., Poggiali, V., Zebker, H., Iess, L., 2017. Titan’s Topography and Shape at the End of the Cassini Mission. Geophysical Research Letters 44, 11,754-11,761. https://doi.org/10.1002/2017GL075518

Coustenis, A., 2014. Titan, in: Spohn, T., Breuer, D., Johnson, T.V. (Eds.), Encyclopedia of the Solar System. Elsevier, Amsterdam ; Boston, pp. 883–905.

Durham, W.B., Prieto-Ballesteros, O., Goldsby, D.L., Kargel, J.S., 2010. Rheological and Thermal Properties of Icy Materials. Space Science Reviews 153, 273–298. https://doi.org/10.1007/s11214-009-9619-1

Elachi, C., Allison, M.D., Borgarelli, L., Encrenaz, P., Im, E., Janssen, M.A., Johnson, W.T.K., Kirk, R.L., Lorenz, R.D., Lunine, J.I., Muhleman, D.O., Ostro, S.J., Picardi, G., Posa, F., Rapley, C.G., Roth, L.E., Seu, R., Soderblom, L.A., Vetrella, S., Wall, S.D., Wood, C.A., Zebker, H.A., 2004. Radar: The Cassini Titan Radar Mapper. Space Science Reviews 115, 71–110.

Elder, C.M., 2015. The Effects of Melt on Impact Craters on Icy Satellites and on the Dynamics of Io’s Interior (PhD). The University of Arizona, Arizona.

Engel, S., Lunine, J.I., Hartmann, W.K., 1995. Cratering on Titan and implications for Titan’s atmospheric history. Planetary and Space Science 43, 1059–1066. https://doi.org/10.1016/0032-0633(95)00044-6

European Space Agency, 1988. Cassini Saturn Orbiter and Titan Probe, Report on the Phase A Study,. ESA.

Farr, T.G., 1993. Radar Interactions with Geologic Surfaces, in: Guide to Magellan Image Interpretation, JPL Publication. NASA and Jet Propulsion Lab, p. 150.

French, B.M., 1998. Traces of Catastrophe: A Handbook of Shock-Metamorphic Effects in Terrestrial Meteorite Impact Structures (Technical Report No. LPI-Contrib- 954). Lunar and Planetary Institute, Houston, TX.

Galilei, G., 1613. Istoria e Dimostrazioni intorno alle Macchie Solari (Letters on Sunspots). Accademia dei Lincei.

Galileo, G., Christoph, S., 2010. On Sunspots. University of Chicago Press.

Gault, D.E., Quaide, W.L., Oberbeck, V.R., 1968. Impact Cratering mechanics and structures, in: Shock Metamorphism of Natural Materials. Mono Book Corp., Baltimore, MD, pp. 87–99.

Grieve, R.A.F., Osinski, Gordon R., Tornabene, L.L., 2014. Planetary Impacts, in: Spohn, T., Breuer, D., Johnson, T.V. (Eds.), Encyclopedia of the Solar System. Elsevier, Amsterdam ; Boston, pp. 83–99.

Griffith, C.A., Owen, T., Miller, G.A., Geballe, 1998. Transient clouds in Titan’s lower atmosphere. Nature, Letters to Nature 395, 4.

Grotzinger, J.P., Hayes, A.G., Lamb, M.P., McLennan, S.M., 2013. Sedimentary Processes on Earth, Mars, Titan, and Venus, in: Comparative Climatology of Terrestrial Planets. University of Arizona Press. https://doi.org/10.2458/azu_uapress_9780816530595-ch18

Hidayat, T., Marten, A., Bezard, B., Gautier, D., Owen, T.C., Paubert, G., 1998. Millimeter and submillimeter heterodyne observations of Titan: The vertical profile of carbon monoxide in its stratosphere. Icarus 133, 109–133.

Hörst, S.M., 2017. Titan’s atmosphere and climate. Journal of Geophysical Research: Planets 122, 432–482. https://doi.org/10.1002/2016JE005240

Huygens, C., 1659. Systema Saturnium (The System of Saturn). Adriaan Vlacq, The Hague.

Huygens, C., 1656. De Saturni luna observatio nova. In Opera varia.

Janssen, M.A., Le Gall, A., Lopes, R.M., Lorenz, R.D., Malaska, M.J., Hayes, A.G., Neish, C.D., Solomonidou, A., Mitchell, K.L., Radebaugh, J., Keihm, S.J., Choukroun, M., Leyrat, C., Encrenaz, P.J., Mastrogiuseppe, M., 2016. Titan’s surface at 2.18-cm wavelength imaged by the Cassini RADAR radiometer: Results and interpretations through the first ten years of observation. Icarus 270, 443–459. https://doi.org/10.1016/j.icarus.2015.09.027

Kenkmann, T., Collins, G.S., Wünnemann, K., 2012. The Modification Stage of Crater Formation, in: Osinski, G.R., Pierazzo, E. (Eds.), Impact Cratering. John Wiley & Sons, Ltd, Chichester, UK, pp. 60–75. https://doi.org/10.1002/9781118447307.ch5

Korycansky, D.G., Zahnle, K.J., 2005. Modeling crater populations on Venus and Titan. Planetary and Space Science 53, 695–710. https://doi.org/10.1016/j.pss.2005.03.002

Kouvaris, L.C., Flasar, F.M., 1991. Phase equilibrium of methane and nitrogen at low temperatures: Application to Titan. Icarus 91, 112–124. https://doi.org/10.1016/0019-1035(91)90131-C

Kuiper, G.P., 1944. Titan: a Satellite with an Atmosphere. Astrophysical Journal 100, 378.

Le Gall, A., Janssen, M.A., Wye, L.C., Hayes, A.G., Radebaugh, J., Savage, C., Zebker, H., Lorenz, R.D., Lunine, J.I., Kirk, R.L., Lopes, R.M.C., Wall, S., Callahan, P., Stofan, E.R., Farr, T., 2011. Cassini SAR, radiometry, scatterometry and altimetry observations of Titan’s dune fields. Icarus 213, 608–624. https://doi.org/10.1016/j.icarus.2011.03.026

Lebreton, J.-P., Matson, D.L., 2003. The Huygens Probe: Science, Payload and Mission Overview, in: Russell, C.T. (Ed.), The Cassini-Huygens Mission. Springer Netherlands, Dordrecht, pp. 59–100. https://doi.org/10.1007/978-94-017-3251- 2_2

Lebreton, J.-P., Witasse, O., Sollazzo, C., Blancquaert, T., Couzin, P., Schipper, A.-M., Jones, J.B., Matson, D.L., Gurvits, L.I., Atkinson, D.H., Kazeminejad, B., Pérez- Ayúcar, M., 2005. An overview of the descent and landing of the Huygens probe on Titan. Nature 438, 758–764. https://doi.org/10.1038/nature04347

Levison, H., Duncan, M.J., 1997. From the Kuiper belt to Jupiter-family comets: The spatial distribution of ecliptic comets. Icarus 127, 13–32.

Lindal, G.F., Wood, G.E., Hotz, H.B., Sweetnam, D.N., Eshleman, V.R., Tyler, G.L., 1983. The atmosphere of Titan: An analysis of the Voyager 1 radio occultation measurements. Icarus 53, 348–363. https://doi.org/10.1016/0019-1035(83)90155- 0

Lorenz, R.D., 2006. What Titan is Really Like : In-Situ Measurements of the Titan Environment by the Huygens Probe. IEEE, pp. 1–10. https://doi.org/10.1109/AERO.2006.1655736

Lorenz, R.D., Lunine, J.I., 1996. Erosion on Titan: Past and Present. Icarus 122, 79–91. https://doi.org/10.1006/icar.1996.0110

Lorenz, R.D., Mitchell, K.L., Kirk, R.L., Hayes, A.G., Aharonson, O., Zebker, H.A., Paillou, P., Radebaugh, J., Lunine, J.I., Janssen, M.A., Wall, S.D., Lopes, R.M., Stiles, B., Ostro, S., Mitri, G., Stofan, E.R., 2008. Titan’s inventory of organic surface materials. Geophysical Research Letters 35. https://doi.org/10.1029/2007GL032118

Lorenz, R.D., Radebaugh, J., 2009. Global pattern of Titan’s dunes: Radar survey from the Cassini prime mission: PATTERN OF TITAN’S DUNES. Geophysical Research Letters 36, n/a-n/a. https://doi.org/10.1029/2008GL036850

Lorenz, R.D., Stiles, B.W., Aharonson, O., Lucas, A., Hayes, A.G., Kirk, R.L., Zebker, H.A., Turtle, E.P., Neish, C.D., Stofan, E.R., Barnes, J.W., 2013. A global topographic map of Titan. Icarus 225, 367–377. https://doi.org/10.1016/j.icarus.2013.04.002

Lorenz, R.D., Turtle, E.P., Stiles, B., Le Gall, A., Hayes, A., Aharonson, O., Wood, C.A., Stofan, E., Kirk, R., 2011. Hypsometry of Titan. Icarus 211, 699–706. https://doi.org/10.1016/j.icarus.2010.10.002

Lorenz, R.D., Wood, C.A., Lunine, J.I., Wall, S.D., Lopes, R.M., Mitchell, K.L., Paganelli, F., Anderson, Y.Z., Wye, L., Tsai, C., Zebker, H., Stofan, E.R., 2007. Titan’s young surface: Initial impact crater survey by Cassini RADAR and model comparison. Geophysical Research Letters 34. https://doi.org/10.1029/2006GL028971

Lunine, J., Stevenson, D.J., 1987. Clathrate and ammonia hydrates at high pressure: Application to the origin of methane on Titan. Icarus 70, 61–77. https://doi.org/10.1016/0019-1035(87)90075-3.

Lunine, J., Tittemore, W.C., 1993. Origins of Outer-Planet Satellites. Protostars and Planets III 1149–1176.

Lunine, J., Yung, Y.L., Lorenz, R., 1999. On the Volatile Inventory of Titan from Isotopic Abundances in Nitrogen and Methane. Planetary and Space Science 47, 1291–1303.

Lunine, J.I., Atreya, S.K., 2008. The methane cycle on Titan. Nature Geoscience 1, 159– 164. https://doi.org/10.1038/ngeo125

Lunine, J.I., Lorenz, R.D., 2009. Rivers, Lakes, Dunes, and Rain: Crustal Processes in Titan’s Methane Cycle. Annual Review of Earth and Planetary Sciences 37, 299– 320. https://doi.org/10.1146/annurev.earth.031208.100142

Lunine, J.I., Stevenson, D.J., Yung, Y.L., 1983. Ethane Ocean on Titan. Science, New Series 222, 1229–1230.

Mastrogiuseppe, M., Poggiali, V., Hayes, A., Lorenz, R., Lunine, J., Picardi, G., Seu, R., Flamini, E., Mitri, G., Notarnicola, C., Paillou, P., Zebker, H., 2014. The bathymetry of a Titan sea. Geophysical Research Letters 41, 1432–1437. https://doi.org/10.1002/2013GL058618

Matson, D.L., Spilker, L.J., Lebreton, J.-P., 2003. The Cassini/Huygens Mission to the Saturnian System, in: Russell, C.T. (Ed.), The Cassini-Huygens Mission. Springer Netherlands, Dordrecht, pp. 1–58. https://doi.org/10.1007/978-94-017-3251-2_1

Melosh, H.J., 1989. Impact Cratering: A Geologic Process, Oxford monographs on geoloy and geophysics. Oxford University Press.

Melosh, J.H., 2011. Planetary Surface Processes, Cambridge Planetary Science. Cambridge University Press.

Mitri, G., Showman, A.P., 2008. Thermal convection in ice-I shells of Titan and Enceladus. Icarus 193, 387–396. https://doi.org/10.1016/j.icarus.2007.07.016

Moreira, A., Prats-Iraola, P., Younis, M., Krieger, G., Hajnsek, I., Papathanassiou, K.P., 2013. A tutorial on synthetic aperture radar. IEEE Geoscience and Remote Sensing Magazine 1, 6–43. https://doi.org/10.1109/MGRS.2013.2248301

Muhleman, D.O., Grossman, A.W., Butler, B.J., 1995. Radar Investigation of Mars, Mercury, and Titan. Annual Review of Earth and Planetary Sciences 23, 337–374. https://doi.org/10.1146/annurev.ea.23.050195.002005

Muhleman, D.O., Grossman, A.W., Butler, B.J., Slade, M.A., 1990. Radar Reflectivity of Titan. Science, New Series 248, 975–980.

Neish, C.D., Barnes, J.W., Sotin, C., MacKenzie, S., Soderblom, J.M., Le Mouélic, S., Kirk, R.L., Stiles, B.W., Malaska, M.J., Le Gall, A., Brown, R.H., Baines, K.H., Buratti, B., Clark, R.N., Nicholson, P.D., 2015. Spectral properties of Titan’s impact craters imply chemical weathering of its surface. Geophysical Research Letters 42, 3746–3754. https://doi.org/10.1002/2015GL063824

Neish, C.D., Carter, L.M., 2014. Planetary Radar, in: Spohn, T., Breuer, D., Johnson, T.V. (Eds.), Encyclopedia of the Solar System. Elsevier, Amsterdam ; Boston, pp. 759–777.

Neish, C.D., Kirk, R.L., Lorenz, R.D., Bray, V.J., Schenk, P., Stiles, B.W., Turtle, E., Mitchell, K., Hayes, A., 2013. Crater topography on Titan: Implications for landscape evolution. Icarus 223, 82–90. https://doi.org/10.1016/j.icarus.2012.11.030

Neish, C.D., Lorenz, R.D., 2014. Elevation distribution of Titan’s craters suggests extensive wetlands. Icarus 228, 27–34. https://doi.org/10.1016/j.icarus.2013.09.024

Neish, C.D., Lorenz, R.D., 2012. Titan’s global crater population: A new assessment. Planetary and Space Science 60, 26–33. https://doi.org/10.1016/j.pss.2011.02.016

Neish, C.D., Molaro, J.L., Lora, J.M., Howard, A.D., Kirk, R.L., Schenk, P., Bray, V.J., Lorenz, R.D., 2016. Fluvial erosion as a mechanism for crater modification on Titan. Icarus 270, 114–129. https://doi.org/10.1016/j.icarus.2015.07.022

Niemann, H.B., Atreya, S.K., Bauer, S.J., Carignan, G.R., Demick, J.E., Frost, R.L., Gautier, D., Haberman, J.A., Harpold, D.N., Hunten, D.M., Israel, G., Lunine, J.I., Kasprzak, W.T., Owen, T.C., Paulkovich, M., Raulin, F., Raaen, E., Way, S.H., 2005. The abundances of constituents of Titan’s atmosphere from the GCMS instrument on the Huygens probe. Nature 438, 779–784. https://doi.org/10.1038/nature04122

Nimmo, F., Bills, B.G., 2010. Shell thickness variations and the long-wavelength topography of Titan. Icarus 208, 896–904. https://doi.org/10.1016/j.icarus.2010.02.020

Osinski, G.R., Grieve, R.A.F., Tornabene, L.L., 2012. Excavation and Impact Ejecta Emplacement, in: Osinski, G.R., Pierazzo, E. (Eds.), Impact Cratering. John Wiley & Sons, Ltd, Chichester, UK, pp. 43–59. https://doi.org/10.1002/9781118447307.ch4

Osinski, G.R., Pierazzo, E., 2012. Impact Cratering: Processes and Products, in: Osinski, G.R., Pierazzo, E. (Eds.), Impact Cratering. John Wiley & Sons, Ltd, Chichester, UK, pp. 1–20. https://doi.org/10.1002/9781118447307.ch1

Osinski, G.R., Tornabene, L.L., Grieve, R.A.F., 2011. Impact ejecta emplacement on terrestrial planets. Earth and Planetary Science Letters 310, 167–181. https://doi.org/10.1016/j.epsl.2011.08.012

Passey, Q.R., Shoemaker, E.M., 1982. Craters and Basins on Ganymede and Callisto- Morphological Indicators of Crustal Evolution. Satellites of Jupiter 379–434.

Pettengill, G.H., Campbell, B.A., Campbell, D.B., Simpson, R.A., 1997. Surface Scattering and Dielectric Properties, in: Venus II : Geology, Geophysics, Atmosphere, and Solar Wind Environment. Unviersity of Arizona Press, pp. 527– 546.

Plescia, J.B., Boyce, J.M., 1985. Impact cratering history of the Saturnian satellites. Journal of Geophysical Research: Solid Earth 90, 2029–2037.

Poggiali, V., Mastrogiuseppe, M., Hayes, A.G., Seu, R., Birch, S.P.D., Lorenz, R., Grima, C., Hofgartner, J.D., 2016. Liquid-filled canyons on Titan. Geophysical Research Letters 43, 7887–7894. https://doi.org/10.1002/2016GL069679

Porco, C.C., Baker, E., Barbara, J., Beurle, K., Brahic, A., Burns, J.A., Charnoz, S., Cooper, N., Dawson, D.D., Del Genio, A.D., Denk, T., Dones, L., Dyudina, U., Evans, M.W., Fussner, S., Giese, B., Grazier, K., Helfenstein, P., Ingersoll, A.P., Jacobson, R.A., Johnson, T.V., McEwen, A., Murray, C.D., Neukum, G., Owen, W.M., Perry, J., Roatsch, T., Spitale, J., Squyres, S., Thomas, P., Tiscareno, M., Turtle, E.P., Vasavada, A.R., Veverka, J., Wagner, R., West, R., 2005. Imaging of Titan from the Cassini spacecraft. Nature 434, 159–168. https://doi.org/10.1038/nature03436

Porco, C.C., West, Robert A., R.A., Squyres, S., McEwen, A., Thomas, P.C., Murray, C.D., Delgenio, A., Ingersoll, A.P., Johnson, T.V., Neukum, G., Veverka, J., Dones, L., Brahic, A., Burns, J.A., Haemmerle, V., Knowles, B., Dawson, D., Roatsch, T., Beurle, K., Owen, W.M., 2004. Cassini Imaging Science: Instrument Characteristics and Anticipated Scientific Investigations at Saturn. Space Science Reviews 115, 363–497.

Sagan, C., 1973. The greenhouse of Titan. Icarus 18, 649–656. https://doi.org/10.1016/0019-1035(73)90068-7

Schenk, P.M., 2002. Thickness constraints on the icy shells of the galilean satellites from a comparison of crater shapes 417, 3.

Schenk, P.M., 1989. Crater formation and modification on the icy satellites of Uranus and Saturn: Depth/diameter and central peak occurrence. Journal of Geophysical Research: Solid Earth 94, 3813–3832. https://doi.org/10.1029/JB094iB04p03813

Schurmeier, L.R., Dombard, A.J., 2018. Crater relaxation on Titan aided by low thermal conductivity sand infill. Icarus 305, 314–323. https://doi.org/10.1016/j.icarus.2017.10.034

Seiff, A., Stoker, C.R., Young, R.E., Mihalov, J.D., McKay, C.P., Lorenz, R.D., 2005. Determination of physical properties of a planetary surface by measuring the deceleration of a probe upon impact: Application to Titan. Planetary and Space Science 53, 594–600. https://doi.org/10.1016/j.pss.2004.05.002

Shoemaker, E.M., Wolfe, R.F., 1982. Cratering Time Scales for the Galilean Satellites, in: Satellites of Jupiter. University of Arizona Press, pp. 277–339.

Sloan, E.D., 1998. Clathrate Hydrates of Natural Gases, 2nd ed. Marcel Dekker, New York.

Smith, B.A., Soderblom, L., Batson, R., Bridges, P., Inge, J., Masursky, H., Shoemaker, E., Beebe, R., Boyce, J., Briggs, G., Bunker, A., Collins, S., Hansen, C.J., Johnson, T.V., Mitchell, J.L., Terrile, R.J., Cook, A.F., Cuzzi, J., Pollack, J.B., Danielson, G.E., Ingersoll, A.P., Davies, M.E., Hunt, G.E., Morrison, D., Owen, T., Sagan, C., Veverka, J., Strom, R., Suomi, V.E., 1982. A New Look at the Saturn System: The Voyager 2 Images. Science, New Series 215, 504–537.

Smith, B.A., Soderblom, L., Beebe, R., Boyce, J., Briggs, G., Bunker, A., Collins, S.A., Hansen, C.J., Johnson, T.V., Mitchell, J.L., Terrile, R.J., Carr, M., Cook, A.F., Cuzzi, J., Pollack, J.B., Danielson, G.E., Ingersoll, A., Davies, M.E., Hunt, G.E., Masursky, H., Shoemaker, E., Morrison, D., Owen, T., Sagan, C., Veverka, J., Strom, R., Suomi, V.E., 1981. Encounter with Saturn: Voyager 1 Imaging Science Results. Science, New Series 212, 163–191. https://doi.org/10.1126/science.212.4491.163

Smith, P.H., Lemmon, M.T., Lorenz, R.D., Sromovsky, L.A., Caldwell, J.J., Allison, M.D., 1996. Titan’s Surface, Revealed by HST Imaging. Icarus 119, 336–349. https://doi.org/10.1006/icar.1996.0023

Soderblom, L.A., Kirk, R.L., Lunine, J.I., Anderson, J.A., Baines, K.H., Barnes, J.W., Barrett, J.M., Brown, R.H., Buratti, B.J., Clark, R.N., Cruikshank, D.P., Elachi, C., Janssen, M.A., Jaumann, R., Karkoschka, E., Mouélic, S.L., Lopes, R.M., Lorenz, R.D., McCord, T.B., Nicholson, P.D., Radebaugh, J., Rizk, B., Sotin, C., Stofan, E.R., Sucharski, T.L., Tomasko, M.G., Wall, S.D., 2007. Correlations between Cassini VIMS spectra and RADAR SAR images: Implications for Titan’s surface composition and the character of the Huygens Probe Landing Site.

Planetary and Space Science 55, 2025–2036. https://doi.org/10.1016/j.pss.2007.04.014

Sotin, C., Head, J.W., Tobie, G., 2002. Europa: Tidal Heating of Upwelling Thermaplumes and the origin of Lenticulae and chaos melting. Geophysics Research Letters 29, 1233. https://doi.org/10.1029/2001GL013844

Stiles, B.W., Hensley, S., Gim, Y., Bates, D.M., Kirk, R.L., Hayes, A., Radebaugh, J., Lorenz, R.D., Mitchell, K.L., Callahan, P.S., Zebker, H., Johnson, W.T.K., Wall, S.D., Lunine, J.I., Wood, C.A., Janssen, M., Pelletier, F., West, R.D., Veeramacheneni, C., 2009. Determining Titan surface topography from Cassini SAR data. Icarus 202, 584–598. https://doi.org/10.1016/j.icarus.2009.03.032

Stofan, E.R., Elachi, C., Lunine, J.I., Lorenz, R.D., Stiles, B., Mitchell, K.L., Ostro, S., Soderblom, L., Wood, C., Zebker, H., Wall, S., Janssen, M., Kirk, R., Lopes, R., Paganelli, F., Radebaugh, J., Wye, L., Anderson, Y., Allison, M., Boehmer, R., Callahan, P., Encrenaz, P., Flamini, E., Francescetti, G., Gim, Y., Hamilton, G., Hensley, S., Johnson, W.T.K., Kelleher, K., Muhleman, D., Paillou, P., Picardi, G., Posa, F., Roth, L., Seu, R., Shaffer, S., Vetrella, S., West, R., 2007. The lakes of Titan. Nature 445, 61–64. https://doi.org/10.1038/nature05438

Stofan, E.R., Wall, S.D., Stiles, B.W., Kirk, R.L., West, R.D., Callahan, P.S., 2012. Cassini RADAR Users Guide. NASA and Jet Propulsion Lab.

Tobie, G., Lunine, J.I., Sotin, C., 2006. Episodic outgassing as the origin of atmospheric methane on Titan. Nature 440, 61–64. https://doi.org/10.1038/nature04497

Tomasko, M.G., Archinal, B., Becker, T., Bézard, B., Bushroe, M., Combes, M., Cook, D., Coustenis, A., de Bergh, C., Dafoe, L.E., Doose, L., Douté, S., Eibl, A., Engel, S., Gliem, F., Grieger, B., Holso, K., Howington-Kraus, E., Karkoschka, E., Keller, H.U., Kirk, R., Kramm, R., Küppers, M., Lanagan, P., Lellouch, E., Lemmon, M., Lunine, J., McFarlane, E., Moores, J., Prout, G.M., Rizk, B., Rosiek, M., Rueffer, P., Schröder, S.E., Schmitt, B., See, C., Smith, P., Soderblom, L., Thomas, N., West, R., 2005. Rain, winds and haze during the Huygens probe’s descent to Titan’s surface. Nature 438, 765–778. https://doi.org/10.1038/nature04126

Towner, M.C., Garry, J.R.C., Lorenz, R.D., Hagermann, A., Hathi, B., Svedhem, H., Clark, B.C., Leese, M.R., Zarnecki, J.C., 2006. Physical properties of Titan’s surface at the Huygens landing site from the Surface Science Package Acoustic Properties sensor (API-S). Icarus 185, 457–465. https://doi.org/10.1016/j.icarus.2006.07.013

Turcotte, D.L., Schubert, G., 2014. Geodynamics, Third edition. ed. Cambridge University Press, Cambridge, United Kingdom.

Turtle, E.P., Perry, J.E., Hayes, A.G., Lorenz, R.D., Barnes, J.W., McEwen, A.S., West, R.A., Del Genio, A.D., Barbara, J.M., Lunine, J.I., Schaller, E.L., Ray, T.L.,

Lopes, R.M.C., Stofan, E.R., 2011. Rapid and Extensive Surface Changes Near Titan’s Equator: Evidence of April Showers. Science 331, 1414–1417. https://doi.org/10.1126/science.1201063

Turtle, E.P., Pierazzo, E., 2001. Thickness of a Europan Ice Shell from Impact Crater Simulations. Science 294, 1326–1328. https://doi.org/10.1126/science.1062492 van Dishoeck, E.F., Jansen, D.J., Phillips, T.G., 1993. Submillimeter Observations of the Shocked Molecular Gas Associated with the Supernova Remnant IC 443. Astronomy and Astrophysics 279, 541–566.

Waite, J.H., 2005. Ion Neutral Mass Spectrometer Results from the First Flyby of Titan. Science 308, 982–986. https://doi.org/10.1126/science.1110652

Werynski, A., Neish, C.D., Le Gall, A., Janssen, M.A., 2017. Compositional Variations of Titan’s Impact Craters Indicates Active Surface Erosion.

Wood, C.A., Lorenz, R., Kirk, R., Lopes, R., Mitchell, K., Stofan, E., 2010. Impact craters on Titan. Icarus 206, 334–344. https://doi.org/10.1016/j.icarus.2009.08.021

Xiao, Z., Komatsu, G., 2013. Impact Craters with Ejecta Flows and Central Pits on Mercury. Planetary and Space Science 82, 62–78.

Zahnle, K., Dones, L., Levison, H.F., 1998. Cratering Rates on the Galilean Satellites. Icarus 136, 202–222. https://doi.org/10.1006/icar.1998.6015

Zahnle, K., Schenk, P., Levison, H., Dones, L., 2003. Cratering rates in the outer Solar System. Icarus 163, 263–289. https://doi.org/10.1016/S0019-1035(03)00048-4

Zahnle, K.J., Schenk, P.M., Sobieszczyk, D., Levison, H.F., 2001. Differential cratering of synchronously rotating satellites by ecliptic comets. Icarus 153, 111–129.

Zarnecki, J.C., Leese, M.R., Hathi, B., Ball, A.J., Hagermann, A., Towner, M.C., Lorenz, R.D., McDonnell, J.A.M., Green, S.F., Patel, M.R., Ringrose, T.J., Rosenberg, P.D., Atkinson, K.R., Paton, M.D., Banaszkiewicz, M., Clark, B.C., Ferri, F., Fulchignoni, M., Ghafoor, N.A.L., Kargl, G., Svedhem, H., Delderfield, J., Grande, M., Parker, D.J., Challenor, P.G., Geake, J.E., 2005. A soft solid surface on Titan as revealed by the Huygens Surface Science Package. Nature 438, 792– 795. https://doi.org/10.1038/nature04211

Chapter 2 2 Identifying Impact Craters on Titan and Constraining their Morphology (i.e. Diameter, Depths, and Rim Heights) 2.1 Introduction

The Cassini-Huygens spacecraft orbited the Saturnian system from 2004 to 2017. It uncovered unprecedented information about Saturn, its rings and moons. We obtained new data about its atmosphere and how it is interconnected with its surface, creating extensive erosional and depositional processes that alter the surface (Burr et al., 2013; Hörst, 2017; Lorenz et al., 2008a; Lorenz, 2006; Lorenz and Radebaugh, 2009). During the Cassini mission, the surface of Titan was observed by three instruments: the Imaging Science System (ISS) at ~4 km/pixel, the Visible and Infrared Mapping Spectrometer (VIMS) at ~500 m/pixel, and the Radio Detection and Ranging (RADAR) instrument at ~350 m/pixel (Brown et al., 2004; Elachi et al., 2004; Porco et al., 2004). ISS studied Titan’s clouds and haze and was used to perform large scale geologic studies of the surface. VIMS also studied atmospheric processes and helped constrain mineralogical compositions of Titan’s surface. RADAR studied surface compositions, measured the topography of the surface and lakes, and obtained the highest resolution images of Titan’s surface (Corlies et al., 2017; Janssen et al., 2016; Le Gall et al., 2011). By the end of the Cassini mission, Titan was revealed to be a diverse geologic landscape with dynamic surface processes that modify the surface.

Crater morphologies are one way of constraining how Titan’s surface is changing (Neish et al., 2016, 2013). Impact craters form distinct morphologies that are fairly well constrained (Bray et al., 2012; Melosh, 1989; Schenk, 2002). The morphologies (depth, diameter, rim heights) of these fresh craters provide quantitative constraints with which to assess how degraded Titan’s craters are. Neish et al. (2013) used Ganymede and Callisto as analogues for fresh Titan craters because these moons are about the same size and density as Titan, with similar gravities. Seven craters (the total population of craters with topography at the time) were used in that study. Neish et al. (2013) showed that Titan’s craters were shallower than craters on Ganymede; the freshest craters on Titan came close

42 to, but never exceeded, the depths observed on Ganymede. This quantitatively showed that Titan was actively changing, and Titan’s crater population reflected this.

However, the earlier crater assessments had already provided a good overview of the state of Titan’s crater population (Lorenz et al., 2007; Neish and Lorenz, 2012; Wood et al., 2010). For example, we had already learned that erosion and burial was likely the dominant mode of crater degradation. Dark crater floors were indicative of fine grained sand particles that appear smooth in radar. There was also ample evidence that fluvial erosion was taking place (Collins, 2005; Lorenz et al., 2008). Furthermore, Titan’s low population of craters was consistent with a heavily modified surface that was rather young (~0.2-1.0 Gyr) (Neish and Lorenz, 2012). The last assessment of Titan’s crater population identified a total of 60 craters from a combination of RADAR and VIMS data through 2010 (Buratti et al., 2012; Lorenz et al., 2007; Neish and Lorenz, 2012; Wood et al., 2010). However, there have been substantial updates to the data since these studies were completed.

Cassini ended its mission with ~69% of Titan’s surface mapped in Cassini RADAR. This is more than double the amount of coverage (~33%) available in the last assessment of Titan’s craters (Neish and Lorenz, 2012). Furthermore, the older data has been updated with newer, slightly different, georeferencing that has shifted the position of previously known features. The amount of topography has grown but still only covers ~9% of Titan’s surface, with ~5% coming from SARTopo with improved height values (Corlies et al., 2017; Stiles et al., 2009). ~2% of the topography comes from stereo methods (Kirk et al., 2012), but this has gone virtually unchanged since Neish et al. (2013) did their study of Titan’s crater topography. Nevertheless, seven years of new data needs to be formally reviewed and mapped to finalize Titan’s crater population. Furthermore, changes to the existing dataset need to be identified to ensure Titan’s crater population is accurately archived. Finally, improvements in the coverage and quality of SARTopo data on Titan has improved and with it we can better constrain the surface processes modifying Titan. A larger population of impact crater topographies is available that can better represent the state of Titan’s craters and improve the measurements with better calibrated data. Therefore, there is much to be gained from an updated look at Titan’s crater population.

In this work, we seek to provide a final look at Titan’s crater population at the end of the Cassini mission. We identify the location of all previously known craters and make note of any new craters identified in the complete data set. Furthermore, we analyze Titan’s crater morphologies to constrain the extent of erosion and deposition on Titan. We identify all craters with topographic data and compare the depths of Titan’s craters with similarly sized craters on the icy moons Ganymede and Callisto. We also consider rim depths to study the effect of fluvial erosion independently of sand and sediment infill (which drives down the crater depths). Our goal is for this work to serve as a final review of Titan’s crater population until further missions provide more data to use (e.g. Turtle et al., 2018a).

2.2 Methodology 2.2.1 Data Gathering and Processing

NASA missions archive their data on the Planetary Data System (PDS). The PDS is publicly available, and it is located at pds.nasa.gov. For this study, we used the radar data from the Cassini RADAR instrument located in the Cartography and Imaging Sciences node. All of the RADAR data is available on the website, but we used the SAR imagery stored in the Basic Image Data Record (BIDR). The data has been processed into .IMG files for each swath, or flyby of the surface. These files are further divided by resolution ranging from 8 pixels per degree (ppd) to 256 ppd. Each swath is georeferenced and can be studied individually, but the images must be projected for larger, global, comparisons.

Image processing was done using a publicly available software ISIS (Integrated Software for Imagers and Spectrometers). The software is provided by the United States Geological Survey (USGS). Each image was reprojected using a simple cylindrical global map projection. The individual swaths were studied within ISIS or by importing them into commercially available GIS software (e.g. ArcGIS). Mosaicking of individual images was done within ISIS.

2.2.2 Crater Mapping

We constructed a global RADAR mosaic using the Synthetic Aperture Radar (SAR) and High-Altitude SAR (HiSAR) image swaths through the last close flyby of Titan, T126.

Images had resolutions that ranged between 64ppd and 256ppd. Before mosaicking, each file was manually evaluated, and those with the highest resolution and relative quality were stacked on top. We also utilized a global ISS mosaic constructed at ~11ppd (4km per pixel) (Cassini ISS Team, 2015). This mosaic is considerably lower in resolution than the RADAR mosaic, but the increased coverage can reveal craters that were not observed by the RADAR instrument.

We use a simple approach for identifying crater candidates. Craters are often circular, but there are a great number of suspiciously circular features on Titan that may have been formed through other geologic processes. We therefore focused our search on circular features with rough (bright) ejecta surrounding smooth (dark) interiors (Figure 2.1). Crater ejecta is commonly radar bright due to the extreme blockiness of the ejected material – see Thompson et al. (1981) for examples from the Moon. Most craters on Titan have been infilled by fine grained sediments, leading to a smooth crater floor. We continue the estimation of certainty from Wood et al. (2010) rating craters from Certain (C1), Nearly Certain (C2), Probable (C3), and a new fourth category of Possible (C4). These classifications roughly relate to how many lines of evidence it has (circularity, ejecta blanket, rims, and dark crater floors), but it also considers a judgement of how confident we are in each line of evidence. Below, we present four examples of craters with certainties from C1 to C4. We identify which lines of evidence are met for each (Figure 2.2). The highest certainty crater (Momoy) shows evidence for all four requirements (circular, ejecta, rim, and crater floor). As we look at Crater #9 H2018, the ejecta blanket becomes less distinct, so we lower its certainty. The craters of certainty C3 and C4 lack ejecta blankets and the crater rims are difficult to differentiate from the rough terrain. There is not a hard line between each certainty factor; these represent overall assessments of what evidence of being a crater exists and how compelling each piece of evidence is. In addition, where SARTopo data exists over a putative crater, we used that data set to differentiate a circular mound from the circular depression of a crater. Most SARTopo data only exist for previously identified craters (Neish et al., 2013; Neish and Lorenz, 2012; Wood et al., 2010), so this is not as useful at identifying new craters.

In addition to identifying new craters, we also reanalyzed previously known craters to identify any errors in the existing dataset (Figure 2.3). We determined that the geo- referencing utilized in Wood et al. (2010) and Neish and Lorenz (2012) no longer matches

47 the existing geo-referencing assigned to the RADAR and SARTopo data; changes in the dataset have shifted the crater centers westward. These types of changes highlight the importance of systematically identifying each crater, including the known population, to ensure they are accurately and fully characterized with the Cassini dataset now complete.

2.2.3 Crater Topography

The topographic dataset is far less extensive than the radar and ISS coverage. Only 14 unique craters have usable topographic data available; 12 are covered in SARTopo, 8 have associated RADAR stereo pairs, and none are covered by the RADAR altimeter. As a result, six craters have overlapping SARTopo and stereo data. For this study, we

48 exclusively analyze the SARTopo data (Figure 2.4) because past works have already derived crater depths from the existing stereo data (Figure 2.5; Neish et al., 2018, 2016, 2013). This data set consists of two-dimensional profiles that cross over the crater in a somewhat random location. The heights are measured in meters around a sphere of radius 2575km because Titan’s geoid was not effectively constrained when the technique was first developed (Stiles et al., 2009). The geoid has been constrained to vary between ~ ± 20km around the reference sphere from the poles to the equator (Iess et al., 2010).

SARTopo data is not publicly available on the PDS. It was provided by Dr. Bryan Stiles to all RADAR team members, including co-author Dr. Catherine Neish (Stiles et al., 2009). The data exists in the form of .csv files of tabulated data with a file for each Titan flyby. I uploaded the data into MATLAB and created a global dataset to compare topographies across different swaths. Each crater was investigated to find all the craters that had SARTopo coverage. Then I created and used a series of functions and scripts to determine the morphometry of each crater with SARTopo coverage. We use the topographic profiles available for each crater to derive the crater diameter and crater depths relative to the local topography and relative to the rims; as a result, we are also able to derive the first ever measurements of rim heights on Titan (Figure 2.4; Appendix A). The rim positions are identified using the peak in topography at the edge of the crater, and the crater floors are identified as the lowest point along the topographic profile in the crater. With these points defined, we can derive each of the desired crater measurements. Crater diameter is defined as the distance from rim to rim (Turtle et al., 2005). Erosion may give an apparent diameter that is slightly larger than when it formed, but it is likely not more than the level of error in our results. The limited topography rarely goes directly through the center of the crater. Therefore, more steps must be taken to derive the actual diameter (Hedgepeth et al., 2018). The simplest approach is to measure the distance from each rim to the center of the crater that was found through radar mapping. The diameter is the average of these distances, but the error in this estimate comes from the assumption that the crater is perfectly circular, where the deviation from the average is the observed error. These errors can be ~20% at times, and it is important to be aware of these when doing diameter dependent studies (e.g. depth to diameter comparisons; Bray et al., 2012; Neish et al., 2013).

(km) of the profile. The rim to floor depth (dr) shown in black. The terrain to floor depth (dt) shown in green. Sinlap represents the topography measurements of a well- preserved crater compared to the heavily degraded Soi crater where dt~0km. More details in text. All of the crater topography data is shown in Appendix A.

Depth to diameter studies of Titan’s craters have shown that they are shallower than crater depths on similar icy moons (Neish et al., 2013). We are updating these results with the new data available to us but are also taking into consideration the effects of resolution. SARTopo averages ~10km of topography, but it does so over ~300m increments (Stiles et al., 2009). The horizontal resolution of topography on other icy moons (0.3-1.5km/pixel) is much higher than in SARTopo. In addition, topography on Ganymede and Callisto was derived from a combination of stereo photogrammetry and shadow length measurements,

50 in contrast to the SARTopo technique which is derived from the calibration of the RADAR data. These differences in topography production may cause small, steep features like crater rims to be artificially reduced in height in the SARTopo data set because they are being averaged with the lower elevations around it. Comparing rim to floor crater depths on Titan to those on similar icy moons gives a quantitative constraint on the degradation of Titan’s craters, but we do a similar analysis of the depth measured from the local terrain to the floor. We can identify the rims in SARTopo; the error would only lower the highest peak, not erase the rims entirely. Therefore, by removing the rim heights, we can constrain the level of degradation independent of erosion of the rim and, by extension, remove the bias in the depth measurements. We also compare the rim to floor depth results to those derived from Titan stereo topography, which uses more comparable techniques at similar resolution (~1.4km/pixel; Kirk et al., 2012) and is available for several of the craters that also have SARTopo. This will help us to determine whether previous results are a robust result, or the cause of differences in the data sets. If the effect is negligible, we also have reliably derived rim heights that can constrain the amount of fluvial erosion independent of depositional processes because the rim heights are not affected by infill. This is especially useful because our measurements would be the first rim heights measured on Titan.

Figure 2.5: Soi crater as observed through stereo data (overlaid on a RADAR image) (Kirk et al., 2012).

To obtain measurements, we follow the methods of Bray et al. (2012, 2008) to ensure our measurements are comparable to Ganymede craters. We begin by subtracting out the regional slopes. In an ideal situation, this produces topographic profiles with the terrain elevation at ~0m and the crater parameters are the large-scale fluctuations that are left in the topography as seen in basic crater models (Figure 2.4). A linear fit is assigned to the topography between 1.5푟 to 5.0푟 from each rim, where 푟 is the radius of the crater. We exclude heights within the rim of the crater to avoid the topography of the crater because it is the anomaly that we are trying to detect. That is to say, we model the shape of the landscape to subtract it out, leaving only the crater shape that has deformed it. We exclude height measurements for another 1.5푟 outside the rims to avoid topography that is biased upward by the ejecta blanket for the same reason. We perform a linear fit because we find this best describes the large-scale changes of the local topography in this way; attempts to fit to higher order equations (order 2, 3 and 4) produce fits that are essentially still linear.

When we subtract out the regional slope, the terrain becomes centered at ~0m. The rim heights (퐻푟) are how high the rims are elevated over the local terrain (~0km),

퐻푟푖 = 퐻푟푖푚 − 퐻푡푒푟푟푎푖푛 (2.2.3.1) where 퐻푟푖푚 is the SARTopo measurement at the rim, 퐻푡푒푟푟푎푖푛 is 0 with the slope removed, and i is for a specific profile on a single side of the crater (Figure 2.4). This measurement is done on the left and right of the profile then averaged to derive a final rim height that accounts for fluctuations in crater topography. This is repeated if more profiles are available and the total average is used. The error in these results is dependent on the intrinsic error reported in each measurement along the SARTopo profile. To obtain the bounds of the error, we compare the upper and lower limits of each rim height measurement (whether in a single profile or across multiple profiles) obtained by adding and subtracting the intrinsic error for each point. Then the highest and lowest possible heights are used as the error bounds.

We use the same steps used to derive the average rim heights to find the average depths, and we determine the maximum and minimum errors allowed, again using the intrinsic errors in each height measurement. The terrain to floor depth (푑푡) (Figure 2.4) is the change in elevation from the terrain (퐻푡푒푟푟푎푖푛~ 0 km) to the floor (퐻푓푙표표푟),

푑푡푖 = 퐻푡푒푟푟푎푖푛 − 퐻푓푙표표푟 (2.2.3.2)

The rim to floor depth (푑푟) (Figure 2.4) is the change in elevation from the rim (퐻푟푖푚) to the floor (퐻푓푙표표푟),

푑푟푖 = 퐻푟푖푚 − 퐻푓푙표표푟 (2.2.3.3) where 퐻푓푙표표푟 is the SARTopo height measurement at the lowest point of the crater floor.

푑푟 is the traditional way of describing crater depths (e.g. Bray et al., 2012). Our approach simply breaks it down into its main components: rim measurements and interior depth. Rim to floor depths are just the combination of these two measurements (i.e. 푑푟 = 푑푡 + 퐻푟), where 퐻푟 is modified by erosion and 푑푡 is modified by infill and deposition. Our work is

53 the first to propose this new measurement of terrain depth, not only on Titan, but as a way of constraining crater morphometry.

2.3 Results 2.3.1 Final Crater Population

Thirty new craters on Titan have been discovered (Figure 2.6) in this work (H2018). Nearly half of the new craters were discovered using HISAR images; none were discovered in the ISS mosaic. It is more difficult to confidently identify craters in the lower resolution HiSAR data set, so most of these craters were classified as C3 (probable) or C4 (possible). The two largest new craters are both 69km in diameter and the smallest is 7km. Most of the craters have a dark interior and a bright, rough rim and ejecta blanket. In some cases, the ejecta blends with the rim or the rim does not seem to be surrounded by an ejecta blanket. Crater #1 (C3) doesn’t exhibit a clear ejecta blanket, nor has it retained a rough rim around the entire crater. Crater #4 (C3) is one of the few cases that has a radar bright interior; this may be a function of its location in the radar bright labyrinth terrains (Lopes et al., 2016; Malaska et al., 2017, 2014). Crater #8 (C3) is not completely mapped in SAR, and its rim is less distinct. However, the interior is moderately darker than the surroundings with areas of moderately brighter ground that may be an ejecta blanket. Crater #11 (C4), #19 (C3), and #20 (C3) show a distinctly darker interior with a possible bright ejecta blanket. The ejecta blanket is more akin to labyrinth terrain which distorts the shape of the dark terrain as well as making it slightly less circular than expected. Crater #14 (C3) lacks a very distinct rim, but there is a hint of an ejecta blanket surrounding a distinctly darker circular interior. Crater #17 highlights the difficulty of identifying smaller (<30km) craters in the HiSAR data set. Crater #18 (C3) is distorted by encroaching sand dunes from the east but still retains a bright rim and ejecta to the west. Craters #22 (C4), #23 (C3), #27 (C3), and #30 (C4) all lack a clear ejecta blanket and only show light evidence of a bright rough rim.

Combining these new craters with the previous data set, Titan’s final crater population is 90 impact craters identified on the surface (50% larger than the previous data set). Most of the craters lie at equatorial latitudes. Nearly 85% of Titan’s craters are within ±50o of the equator, and even when the planet is split into equal surface areas (±30o), more than half (65%) of the crater population remain (Figure 2.7). This may be a result of the increased fluvial and lacustrine activity in the higher latitudes (Burr et al., 2013; Neish et al., 2016; Neish and Lorenz, 2014). The entire population has been updated to reflect their center latitude and longitude and diameter (Table 2.1).

Figure 2.7: A global Cassini RADAR mosaic of Titan mapping the 30 new craters (red) and the original 60 (yellow).

ID Crater Name Longitude Latitude Diameter Certainty RADAR (km) Swath

1* Menrva 86.97 19.98 400 1 T003, T077, T108

2* Forseti 10.74 25.76 140 1 T023, T113

3* Afekan 200.27 26.00 115 1 T043, T083

4 Paxsi 341.53 5.51 115 2 T104

5* Hano 344.98 40.51 105 1 T016, T084, T104

6* Sinlap 15.98 11.51 88 1 T003, T108, T113

7* Soi 140.97 24.29 84 1 T016, T055, T056, T120

8* Selk 199.05 6.92 84 1 T036, T095, T098, T120, T121

9* Crater #1 240.33 31.94 70 3 T084, T104 H2018

10 Crater #2 129.24 14.49 69 3 T120 H2018

11 Guabonito 151.71 -11.32 66 2 T013, T048

12 Crater #26 W 85.99 -8.14 62 2 T013, T029

13 Crater #3 138.44 1.84 62 2 T091 H2018

14 Crater #4 110.83 -6.16 60 3 T113 H2018

ID Crater Name Longitude Latitude Diameter Certainty RADAR (km) Swath

15 Nath 7.77 -30.61 60 2 T050

16 Crater #49 W 188.91 -10.92 60 3 T008, T036, T041, T121

17 Crater #6 347.75 1.48 57 2 T104 H2018

18 Crater #7 164.40 31.33 55 3 T120 H2018

19* Ksa 65.34 13.72 45 1 T017, T077, T113, T083

20* Shikoku 165.05 -7.80 41 1 T013, T048, T121

21 Crater #47 W 184.47 -7.25 41 3 T008, T120, T121

22 Crater #25 W 88.67 -11.54 40 2 T003, T013, T113

23 Santorini 147.55 2.15 40 1 T044, T056

Crater #8 293.05 52.09 40 3 T084, T104, 24 H2018 T016

Crater #9 25

H2018 205.01 -14.69 38 2 T113

26 Crater #10 340.81 -8.75 35 3 T104 H2018

27 Crater #45 W 18.50 8.10 34 3 T003, T121

28 Crater #11 353.88 54.20 33 4 T083, T098 H2018

29 Crater #23 W 49.80 47.60 33 2 Ta

30 Crater #5 NL 140.54 12.13 33 2 T056

31 Momoy 44.58 11.66 32 1 T017, T077, T113

ID Crater Name Longitude Latitude Diameter Certainty RADAR (km) Swath

32 Crater #13 153.97 19.84 32 3 T049 H2018

33 Crater #43 W 50.18 10.61 32 3 T017, T077, T083

34 Crater #14 188.43 21.18 31 3 T098 H2018

35 Crater #46 W 40.43 -49.60 30 3 T36

36 Crater #44 W 62.08 12.46 29 3 T017, T077

37 Beag 169.47 -34.74 28 1 T098

38 Crater #16 149.14 -68.63 27 4 T058, T077 H2018

40 Crater #22 W 29.96 11.35 26 2 T029

41* Crater #10 129.69 -30.58 25 3 T041 NL

42 Crater #17 348.98 -6.32 27 3 T023, T104, H2018 T113, T120

43 Crater #18 207.82 -13.10 25 4 T058, T064, H2018 T095

44 Crater #4 NL 141.78 11.01 22 2 T056, T020

45 Crater #42 W 166.86 -10.56 22 3 T013, T121

46 Crater #21 W 84.27 -10.53 20 2 T013, T041, T113

47 Crater #19 280.44 0.02 19 3 T071 H2018

48* Crater #3 NL 150.67 -16.33 19 2 T048, T058, T057

39 Crater #20 127.04 -3.12 19 3 T058, T057, H2018 T020

ID Crater Name Longitude Latitude Diameter Certainty RADAR (km) Swath

49 Crater #41 W 260.19 29.95 18 3 T043, T083

50 Crater #9 NL 162.65 19.15 18 3 T057, T056

51 Crater #21 121.49 -15.59 18 3 T021 H2018

52 Crater #22 62.47 -9.77 18 2 T113 H2018

53 Crater #20 W 77.90 -7.70 17 2 T013

54 Crater #40 W 67.33 -12.41 17 3 T013, T113, T121

55 Crater #23 170.68 21.76 17 4 T056, T098 H2018

56 Crater #24 353.82 36.55 17 3 T084 H2018

57 Crater #17 W 75.93 -13.20 16 2 T013

58 Crater #7 NL 160.15 -38.27 16 3 T057, T120

59 Crater #13 W 309.90 83.30 15 2 T029

60 Crater #15 W 250.12 39.36 15 2 T021, T084, T092

61 Crater #19 W 74.88 -13.25 15 3 T013

62 Crater #39 W 129.70 33.40 14 3 Ta

63 Crater #18 W 249.26 38.16 12 2 T084, T092

64 Crater #25 175.58 -6.68 12 3 T048, T077 H2018

65 Crater #38 W 24.51 -57.10 11 3 T039

66 Crater #8 NL 194.79 0.06 11 3 T036, T061, T098, T121

ID Crater Name Longitude Latitude Diameter Certainty RADAR (km) Swath

67 Crater #48 W 173.40 22.70 10.5 3 T021, T029, T083, T120

68 Crater #16 W 239.00 82.50 10 2 T029

69 Crater #37 W 349.40 41.80 10 3 T018

70 Crater #35 W 266.93 23.88 9 3 T098

71 Crater #36 W 105.20 64.40 9 3 T056, T098

72 Crater #26 170.97 -37.79 9 3 T021, T049 H2018

73 Crater #27 171.35 -39.96 9 3 T018 H2018

74 Crater #28 73.33 12.31 9 4 T104 H2018

75 Crater #33 W 324.50 83.90 8 3 T71

76 Crater #34 W 228.20 36.40 8 3 T025

77 Crater #9 W 254.00 39.00 8 3 T030

78 Crater #29 235.25 -60.12 8 3 T021, T084, H2018 T091

79 Crater #7 W 39.42 -18.92 8 3 T025

80 Crater #14 W 349.10 44.00 7 2 T018

81 Crater #30 104.62 67.45 7 4 T043, T113, H2018 T020, T083, T098

82 Crater #11 W 343.30 25.28 6 3 T016

83 Crater #32 W 229.00 83.00 6 3 T025

84 Crater #12 W 12.91 -54.48 5 2 T007

85 Crater #10 W 43.40 18.58 4 2 T003

ID Crater Name Longitude Latitude Diameter Certainty RADAR (km) Swath

86 Crater #8 W 17.12 14.04 4 2 T003

87 Crater #30 W 43.15 -18.89 4 3 T025

88 Crater #6 W 38.40 -17.10 3 2 T025

89 Crater #29 W 261.14 30.50 3 3 T021

90 Crater #31 W 45.90 -28.60 3 3 T025 2.3.2 Crater Counts

We have produced a final crater count, corrected for the incomplete coverage of Titan by Cassini RADAR, following the same approach as Neish and Lorenz (2012). We use a Monte-Carlo approach to determine the probability of detecting craters of different sizes on Titan (e.g. Neish and Lorenz 2012). We begin by constructing two 360×180 arrays where each cell represents a 1°×1° area of Titan’s surface. The first array represents the SAR coverage where each “active” (1) cell represented where there was coverage and all others are “inactive” (0) where there is no coverage. The second array represents the random position of a crater of a specific size. A random center latitude and longitude is assigned, and each cell along the rim and within the interior of the crater is mapped as "active" (1) with all cells outside the crater as "inactive" (0). We determine in which cells the crater is active by finding the distance between each cell and the assigned center latitude and longitude. All cells with a distance ≤ the radius of the crater is considered “active.” Finally, if any cell is active in both arrays, then the crater is classified as “discovered”. As 푛 → ∞, where n is the number of times the experiment is done, the frequency of discovering the crater approaches the probability of discovering it. We repeat this process 1,000 times for diameters between 30-1200 km in bins of √2 km, consistent with previous Titan crater assessments (Artemieva and Lunine, 2005; Korycansky and Zahnle, 2005a; Neish and Lorenz, 2012).

When Neish and Lorenz (2012) accessed Titan’s crater population only ~33% of the surface was mapped in SAR. With ~69% of the surface now mapped, the probability of

62 missing a crater has decreased. The effect is dampened with increasing crater diameter as the probability of detecting larger craters approaches 100%. The crater count is adjusted to included possibly undetected craters,

−1 퐶 = 퐶푎푐푡푢푎푙 ∗ 푃 (2.3.2.1) where 퐶 is the corrected crater count, 퐶푎푐푡푢푎푙 is the actual crater count, and 푃 is the probability of being detected (Table 2.2).

Table 2.2: Titan’s crater count corrected for incomplete coverage.

Correct Number of Probability of Diameter (km) Number of Craters Detection Craters

2√2 − 4 3 68.9 4.4

4 − 4√2 4 68.9 5.8

4√2 − 8 4 68.9 5.8

8 − 8√2 15 68.9 21.8

8√2 − 16 6 68.9 8.7

16 − 16√2 15 68.9 21.8

16√2 − 32 10 68.9 14.5

32 − 32√2 15 68.8 21.8

32√2 − 64 7 73.3 9.5

64 − 64√2 6 76.8 7.8

64√2 − 128 4 80.1 5

256√2 − 512 1 95.3 1

We find that the final crater count follows the same general trend as previous assessments (Figure 2.8). There are fewer smaller craters than originally thought (Neish and Lorenz, 2012). However, Neish and Lorenz (2012) noted they found there was an excess of smaller craters than originally predicted (Artemieva and Lunine, 2005; Korycansky and Zahnle, 2005b), and we find it is more consistent with original predictions. Smaller impactors (that produce craters with diameters < 20 km) on Titan are preferentially broken apart by the atmosphere. The general trend remains, and it suggests the crater retention age derived by Neish and Lorenz (2012) remains a plausible estimation for Titan (i.e., 0.2 to 1.0 Ga).

2.3.3 Crater Topography

We find that only ~30 craters (new and old) have SARTopo data near (≤ 3푟) the crater, and only 15 of these pass over the crater itself. Three of those are too noisy or are missing too much of the interior to be interpreted. All three of these craters were < 20 km in diameter, and as we approach these smaller crater sizes, it is more difficult for a profile with such coarse resolution to be usable. This leaves 12 craters with high enough quality SARTopo data to interpret. Stereo results have been reported for 8 craters, only two of

64 which are unique from the 12 observed in SARTopo. Here, we update and finalize the crater depths using the SARTopo data, and we compare this to previously measured stereo data to judge the efficacy of each method (Table 2.3). Six of the eight craters with stereo data have depths that match up with SARTopo results within errors, while two are significantly less. Neish et al. (2018) explore the possible reasons for this marked difference, suggesting that a lack of features in the interior may bias the interpolation to the rim height thereby decreasing the measured depth of the crater. Given the similarities in depths between SARTopo and stereo, we therefore consider it reasonable to use SARTopo data to compare with the depths of craters other icy moons, substituting stereo where SARTopo is not available.

Titan’s craters can be compared to similarly sized craters on Ganymede and Callisto. This is because Ganymede, Callisto, and Titan are all about the same size with similar gravities, densities, and interior structures. Impact craters follow specific morphologies depending on the size of the impact and the body in question (Schenk, 2002). For example, Europa transitions into complex craters sooner than its larger siblings (Ganymede and Callisto) because of its thin crust. We consider these good analogues for Titan craters, but we also acknowledge the potential for clathrates or a subsurface methane reservoir to modify Titan’s crater trends. These overall effects are difficult to predict, and no detailed modeling of these effects has yet occurred. Clathrates would strengthen the crust (Durham et al., 2010) possibly producing craters with smaller diameters than those on Ganymede. On the other hand, wetlands or marine environments may form craters with minimal surface topography (Neish and Lorenz, 2014).

Table 2.3: A list of Titan's craters with SARTopo and Stereo topography data available and their Rim Heights, Rim Depths, and Terrain Depths with the rim depths compared to relative depths (Rr).

Rim Rim Terrain Diameter Relative Relative Crater Heights Depth Depth a b Technique (km) Depth Rr Depth Rr 푯풓 (km) 풅풓(km) 풅풕(km)

+25 +95 +130 +170 +0.10 +0.13 1 Menrva 400−30 150−120 430−115 280−210 0.64−0.11 0.54−0.14 SARTopo

+40 +125 +145 +185 +0.10 +0.20 1 Forseti 125−40 285−115 405−185 120−120 0.66−0.12 0.56−0.16 SARTopo

+10 +60 +0.05 +0.07 2 140−10 - 180−60 - 0.85−0.05 0.80−0.07 Stereo

+5 +115 +160 +75 +0.14 +0.18 1 Afekan 115−5 235−120 530−165 295−80 0.56−0.13 0.43−0.16 SARTopo

+10 +75 +100 +135 +0.07 +0.09 1 Hano 105−20 230−105 400−80 170−145 0.67−0.08 0.57−0.11 SARTopo

+5 2 100−5 - ~0 - ~1 ~1 Stereo

+1 +175 +255 +175 +0.21 +0.26 1 Sinlap 88−1 375−150 790−245 410−145 0.31−0.22 0.17−0.27 SARTopo

+2 +100 +0.09 +0.10 4 82−2 - 700−100 - 0.38−0.09 0.30−0.10 Stereo

+15 +45 +80 +95 +0.07 +0.09 1 Soi 85−15 220−45 220−80 0−0 0.81−0.07 0.77−0.08 SARTopo

+2 +120 +0.10 +0.11 3 78−2 - 240−120 - 0.78−0.10 0.76−0.11 Stereo

+2 +50 +125 +155 +0.09 +0.11 1 Selk 84−2 280−54 470−105 190−130 0. 58−0.11 0.51−0.13 SARTopo Crater #1 70+25 200+125 290+165 90+105 0.73+0.15 0.70−0.16 SARTopo1 H2018 −25 −110 −155 −75 −0.16 +0.17

+1 +60 +90 +65 +0.09 +0.07 1 Shikoku 41−1 255−65 300−80 45−45 0. 66−0.10 0.72−0.08 SARTopo

+5 +80 +0.09 +0.07 3 42−5 - 340−80 - 0.61−0.09 0.60−0.07 Stereo

+2 +50 +85 +75 +0.11 +0.09 1 Ksa 45−2 395−60 795−100 400−90 0.12−0.10 0.24−0.08 SARTopo

+2 +175 +0.20 +0.15 4 39−2 - 750−175 - 0.13−0.20 0.34−0.15 Stereo

+1 +100 +0.11 +0.09 4 Momoy 40−1 - 680−100 - 0.22−0.11 0.40−0.09 Stereo

+5 +70 +0.08 +0.06 3 Santorini 40−5 - 340−70 - 0.61−0.08 0.70−0.06 Stereo Crater #10 25+2 0+20 +45 140+65 0.86+0.06 +0.05 SARTopo1 NL −2 −0 105−45 −60 −0.06 0.88−0.05

Crater #3 +5 +100 +170 +155 +0.26 +0.21 1 20−5 130−105 250−175 120−120 0.63−0.25 0.69−0.21 SARTopo NL 1This work, 2Neish et al. (2018), 3Neish et al. (2016), 4Neish et al. (2013), aBray et al. (2012), Schenk (2002)b

We chose to compare Titan’s craters with those of Ganymede rather than Callisto because there is more data available. We use the Ganymede crater parameters from Bray et al. (2012) and Schenk (2002) to compare with those of Titan (Figure 2.9). Comparing the crater depths between the two moons allows us to constrain the level of degradation occurring on Titan’s craters. We also present relative depths (R), measured from rim to 푑 floor, where 푅 = 1 − 퐺푎푛푦푚푒푑푒 and d is the depth of the crater (Table 2.3). We use the rim 푑푇푖푡푎푛 to floor measurements because these have been characterized with trend lines for ‘Ganymede crater depths. The freshest craters will have 푅 → 0 because the Titan depths approach the Ganymede depths. Qualitatively, the depths of craters on Titan appear shallower than the depths of fresh craters on Ganymede (Figure 2.9). The smaller crater depths on Titan suggest infill by aeolian and fluvial erosion, as first proposed by Neish et al. (2013). With relative depths, we begin to quantify this relationship, and we continue by statistically testing what we think we are seeing.

Figure 2.9: a) Crater depths, measured from rim to floor, of Ganymede craters (Bray et al., 2012 black diamonds; Schenk 2002 dashed black line) and Titan crater depths as measured by SARTopo (red circles) and stereo data (blue squares). The Schenk (2002) data is in the form a trend line that changes at ~35km, but we show it extended because the transition is gradual. b) Rim heights of Ganymede craters as measured by Bray et al. (2012) (data black diamonds, trend black dashed line) compared to those measured on Titan (red circles). As with Schenk (2002) rim to floor depths, the dashed lines are trend lines fit to the data over a range of crater diameters. However, the largest craters (~>50km) are shown as the upper and lower boundaries because the trend line breaks down. c) Crater depths, measured from local terrain of Ganymede craters (Bray et al., 2012 black diamonds) compared to Titan craters (red circles).

2.3.3.1 Crater Depth We wish to test the probability that Titan depths (rim to floor) can be compared to fresh Ganymede craters that have been altered by infill and erosion (Figure 2.9a). We used the goodness-of-fit Kolmogorov-Smirnov (KS) test to judge whether the two depth distributions come from the same population or not (e.g. Tornabene et al., 2018) . The technique finds the maximum difference (KS) between the observed and modeled distribution functions, 퐹푛(푥) and 퐹(푥) respectively (Trauth 2015):

퐾푆 = max|퐹푛(푥) − 퐹(푥)| (2.3.3.1.1) where n is the number of measurements.

The null hypothesis states that the observed measurements are described with the same distribution as the model. The null hypothesis is rejected if the calculated KS is more than the critical KS value or the p value is less than the significance factor where p is the probability of observing a statistic this extreme if the null hypothesis is true (O’Connor and Kleyner, 2012).

Regarding Titan, the null hypothesis is that Titan’s craters are fresh, with similar depths to Ganymede. The average crater depths (rim to floor) for Titan give a KS = 0.65

(KScritical=0.449) and a p=0.0153%. Even when we use the upper limits of Titan’s depths, we find KS = 0.447 and p=2.3%. These results suggest we can reject the null hypothesis that Titan’s crater depths follow the same distribution as fresh Ganymede crater depths. Even in the “worst case” scenario, there is only a 2% probability of observing a statistic this extreme if the craters come from the same population.

Therefore, we deem it likely that Titan’s craters represent a distinct population of craters from Ganymede. On average, Titan’s rim to floor crater depths are 49% more shallow. The most degraded craters are the smallest (10

69 by filling the crater floor. However, craters between 30km to 90km are less infilled, and this suggest something else is at play. If we assume constant infill for the craters, it may suggest that craters of this size are younger craters. The large error in diameter measurements may also make it difficult to compare for such a small population of craters (12 total). Nevertheless, we compare rim heights to see if the pattern holds.

However, first we must consider whether SARTopo measurements are reliable indicators of rim heights. The lower resolution of SARTopo could play a significant role in what differences we can observe in the data. The topography on Ganymede has much better resolution (0.05-1.0 km, Schenk and Ridolfi, 2002) than that of Titan (~10km, repeated over 300m increments, in SARTopo) (Stiles et al., 2009). The lower resolution results in averaging longer wavelength signals while rims peak over shorter wavelengths. The resultant averaging of topography in SARTopo may artificially lower the rim heights observed on Titan, such that the crater depths would be decreased. However, the Titan stereo data has resolutions more comparable to Ganymede (~1.4 km, (Kirk et al., 2012)). Given that Titan stereo and SARTopo measurements appear very comparable, we suspect that the effect of resolution is minimal. A KS test on the two populations give a p=96.9% with a KS=0.208 (KScritical= 0.410). Therefore, the null hypothesis (that the two represent the same population) cannot be disregarded. In fact, the two populations match remarkably well, so SARTopo topography is expected to be as good an estimate of the height of the rim peaks on Titan as the topography on Ganymede.

2.3.3.2 Rim Heights

Measuring rim heights provides constraints on fluvial erosion, given aeolian erosion of the rims is unlikely (Collins, 2005; Forsberg-Taylor and Howard, 2004). Aeolian processes tend to fill in basins without abrading or covering up topographic highs. Therefore, we measure the depth from the local terrain to constrains the crater infill, and we use the rim heights to quantify the fluvial erosion (Figure 2.9b and 2.9c).

Titan’s rim heights are not as a high as the rim heights observed on Ganymede. This suggests that not only does infill play a significant role but fluvial erosion of the rims does too (Neish et al., 2016). A KS-test gives similar results to the KS test of Titan’s crater

depths (KS=0.606, KScritical=0.449 and p=0.0549%). Therefore, we reject the null hypothesis that the rim heights on Titan follow the same distribution as Ganymede. We find that Titan’s rims are, on average, even more degraded than the depths (54%). We see the same pattern with crater size as with the crater depths. We see the largest craters are the most degraded (72%). Without the smallest craters, the rim heights follow the expected pattern; the largest craters take longer to fill and should experience more fluvial erosion before they are infilled. The anomaly of the <30km craters may be a limitation of our method which calibrates for local topography, and as we reach smaller craters, the fluctuations in local topography become more exaggerated and difficult to account for. However, it is possible that this is not the case, and smaller craters experience these processes differently, so we continue by comparing the depth without the rims (terrain to floor) for a closer look.

The rim to floor depths are affected by infill (changes in 푑푡) and erosion to the rim

(changes in 퐻푟). By subtracting out the rim heights, we can observe how much of the degradation is caused by infill. However, the infill may be aeolian or from fluvial deposition. First, we perform a KS test and find that the Titan and Ganymede populations are the most similar between the three measurements (p=9%, KS=0.37 and

KScritical=0.449). It is still a very low probability of these two populations being the same (9%), so it is unlikely. On average, these depths are 41% less shallow than those on Ganymede. The smallest craters are much shallower (69%). We hypothesize that the smallest craters have infilled to an ~equilibrium state where erosion of the rim has begun to erase away the final remnants of the crater. Modeling suggests erosion and infill should be largest for small craters with the rate of infill decreasing with time (Forsberg-Taylor and Howard, 2004). However, we do not observe this beyond 30km. Craters become more degraded and infilled for larger diameters. Medium sized craters (30km-90km) craters still exhibit less degradation; infill has only made the craters 34-36% less shallow. The rims exhibit the same pattern but with more degradation, being 45-58% lower than on Ganymede. Tthe largest craters are 58% infilled and have 62% degraded rims. The fact that larger craters are more degraded than the medium sized craters, from both infill and erosion, is counter intuitive if we are to believe infill degrades smaller crater more quickly (Neish et al., 2013). However, it may just support the idea that the largest craters are

71 generally older, so they would be exposed to degradation for a longer period. This suggests the medium sized craters are younger, and in fact, these craters also the freshest in RADAR and relative depths. Alternatively, it may suggest significant changes in infill and erosion rates across Titan’s surface.

2.4 Conclusion

In this work, we have produced the final post-Cassini crater population for Titan. We found that Titan’s crater population is 50% larger than previously thought. However, adjusting for the increased coverage shows that the population is in line with previous assessments, with slightly fewer smaller craters (≤ 8 km). Our results are consistent with previously reported crater retention ages between 200 Ma to 1.0 Ga. An updated ISS map is to be released by the end of 2018 that has remarkably increased resolution. Our results suggest there are still a few craters left unseen, and this may help to better constrain Titan’s crater population.

We then used the data available to further characterize 12 of Titan’s craters using SARTopo data. We demonstrated that these results are comparable to stereo topography data in most cases. In cases where they are not, stereo results in lower depths, possibly due to bias in how the stereo processing interpolates the original images. Comparing these results with Ganymede shows that Titan’s craters are 50% shallower, driven by a mix of infill (41% shallower) and erosion (54% more degraded rims). This suggests fluvial erosion may play a much larger role than previously thought and that weathering rates likely vary across Titan (Werynski et al., 2017). There is ample evidence of fluvial activity all over Titan’s surface to support this. However, there are craters that contradict this. For example, Soi (D=85km) is one of the most degraded craters on Titan (R=0.81, compared to Bray et al., 2012), yet its rims are only ~100m more degraded than one of the freshest craters on Titan, Sinlap (D=88km, R=0.31). Neish et al., (2015) suggested Soi’s rims should be moderately eroded compared to Sinlap with more aeolian infill and fluvial deposition in the floor, and this is consistent with our findings.

We were also concerned that lower rim heights may be artificially lower due to low resolution averaging, but comparisons of stereo and SARTopo show similar results (Figure

2.9). Stereo data has more comparable resolution to that the topography of Ganymede; this suggests lower resolution does not significantly affect the result given that SARTopo and stereo results are generally comparable.

Future missions may help constrain our results and test the efficacy of the SARtopo data. There is a need to ground truth what we have observed in Titan’s impact craters. Dragonfly is a rotorcraft lander in the New Frontiers 4 Phase A competition. This mission would study Titan’s chemistry and geology (Turtle et al., 2018). If selected, it would investigate the transition from pre-biotic to biotic in the origin of life, and one possible candidate for investigation would be Titan’s impact craters as a source for prolonged liquid water on the surface (Neish et al., 2018). Therefore, we may receive unprecedented data about some of Titan’s impact craters, which would contribute significantly in constraining the processes and morphologies that shape its surface.

2.5 References

Bray, V.J., Schenk, P.M., Jay Melosh, H., Morgan, J.V., Collins, G.S., 2012. Ganymede crater dimensions – Implications for central peak and central pit formation and development. Icarus 217, 115–129. https://doi.org/10.1016/j.icarus.2011.10.004

Cassini ISS Team, 2015. Map of Titan, Cassini ISS Global Mosaic.

Collins, G.C., 2005. Relative rates of fluvial bedrock incision on Titan and Earth. Geophysical Research Letters 32, n/a-n/a. https://doi.org/10.1029/2005GL024551

Forsberg-Taylor, N.K., Howard, A.D., 2004. Crater degradation in the Martian highlands: Morphometric analysis of the Sinus Sabaeus region and simulation modeling suggest fluvial processes. Journal of Geophysical Research 109. https://doi.org/10.1029/2004JE002242

Hedgepeth, J.E., Neish, C.D., Turtle, E.P., Stiles, B.W., 2018. Impact Craters on Titan: Finalizing Titan’s Crater Population, in: Titan Is Terrific. Presented at the 49th Lunar and Planetary Science Conference, Lunar and Planetary Institute, Woodlands, Texas.

Hörst, S.M., 2017. Titan’s atmosphere and climate. Journal of Geophysical Research: Planets 122, 432–482. https://doi.org/10.1002/2016JE005240

Iess, L., Rappaport, N.J., Jacobson, R.A., Racioppa, P., Stevenson, D.J., Tortora, P., Armstrong, J.W., Asmar, S.W., 2010. Gravity Field, Shape, and Moment of Inertia of Titan. Science 327, 1367–1369. https://doi.org/10.1126/science.1182583

Kirk, R.L., Howington-Kraus, E., Redding, B., Callahan, P.S., Hayes, A.G., Le Gall, A., Lopes, R.M.C., Lorenz, R.D., Lucas, A., Mitchell, K.L., Neish, C.D., Aharonson, O., Radebaugh, J., Stiles, B.W., Stofan, E.R., Wall, S.D., Wood, C.A., 2012. Topographic Mapping of Titan: Latest Results. Presented at the 43rd Lunar and Planetary Science Conference, Lunar and Planetary Institute, Woodlands, Texas.

Korycansky, D.G., Zahnle, K.J., 2005a. Modeling crater populations on Venus and Titan. Planetary and Space Science 53, 695–710. https://doi.org/10.1016/j.pss.2005.03.002

Korycansky, D.G., Zahnle, K.J., 2005b. Modeling crater populations on Venus and Titan. Planetary and Space Science 53, 695–710. https://doi.org/10.1016/j.pss.2005.03.002

observations of Titan’s dune fields. Icarus 213, 608–624. https://doi.org/10.1016/j.icarus.2011.03.026

Lopes, R.M.C., Malaska, M.J., Solomonidou, A., Le Gall, A., Janssen, M.A., Neish, C.D., Turtle, E.P., Birch, S.P.D., Hayes, A.G., Radebaugh, J., Coustenis, A., Schoenfeld, A., Stiles, B.W., Kirk, R.L., Mitchell, K.L., Stofan, E.R., Lawrence, K.J., 2016. Nature, distribution, and origin of Titan’s Undifferentiated Plains. Icarus 270, 162–182. https://doi.org/10.1016/j.icarus.2015.11.034

Lorenz, R.D., 2006. The Sand Seas of Titan: Cassini RADAR Observations of Longitudinal Dunes. Science 312, 724–727. https://doi.org/10.1126/science.1123257

Lorenz, R.D., Lopes, R.M., Paganelli, F., Lunine, J.I., Kirk, R.L., Mitchell, K.L., Soderblom, L.A., Stofan, E.R., Ori, G., Myers, M., Miyamoto, H., Radebaugh, J., Stiles, B., Wall, S.D., Wood, C.A., 2008. Fluvial channels on Titan: Initial Cassini RADAR observations. Planetary and Space Science 56, 1132–1144. https://doi.org/10.1016/j.pss.2008.02.009

Lorenz, R.D., Radebaugh, J., 2009. Global pattern of Titan’s dunes: Radar survey from the Cassini prime mission: Pattern of Titan’s Dunes. Geophysical Research Letters 36. https://doi.org/10.1029/2008GL036850

Malaska, M.J., Lopes, R., Mitchell, K.L., Radebaugh, J., Verlander, T., Schoenfeld, A., 2017. Classification of Labyrinth Terrains on Titan. Presented at the 48th Lunar and Planetary Science Conference, Lunar and Planetary Institute, Woodlands, Texas.

Malaska, M.J., Radebaugh, J., Lopes, R., Mitchell, K.L., Hayes, A.G., Le Gall, A., Turtle, E.P., Solomonidou, A., Lorenz, R., 2014. Labyrinth Terrain on Titan, in: Ice in the Solar System: Mars to Pluto and Icy Worlds in Between. Presented at the 2014 Geological Society of America Annual Meeting, Geological Society of America, Vancouver, BC.

Melosh, H.J., 1989. Impact Cratering: A Geologic Process, Oxford monographs on geoloy and geophysics. Oxford University Press.

Neish, C.D., Lorenz, R.D., 2014. Elevation distribution of Titan’s craters suggests extensive wetlands. Icarus 228, 27–34. https://doi.org/10.1016/j.icarus.2013.09.024

Neish, C.D., Lorenz, R.D., 2012. Titan’s global crater population: A new assessment. Planetary and Space Science 60, 26–33. https://doi.org/10.1016/j.pss.2011.02.016

Neish, C.D., Lorenz, R.D., Turtle, E.P., Barnes, J.W., Trainer, M.G., Stiles, B., Kirk, R., Hibbitts, C.A., Malaska, M.J., 2018. Strategies for Detecting Biological Molecules on Titan. Astrobiology 18, 571–585. https://doi.org/10.1089/ast.2017.1758

O’Connor, P.D., Kleyner, A., 2012. Pratical Reliability Engineering, 5th ed. John Wiley & Sons, Ltd.

Schenk, P.M., 2002. Thickness constraints on the icy shells of the galilean satellites from a comparison of crater shapes 417, 3.

Schenk, P.M., Ridolfi, F.J., 2002. Morphology and scaling of ejecta deposits on icy satellites. Geophysical Research Letters 29. https://doi.org/10.1029/2001GL013512

Thompson, T.W., Zisk, S.H., Shorthill, R.W., Schultz, P.H., Cutts, J.A., 1981. Lunar craters with radar bright ejecta. Icarus 46, 201–225. https://doi.org/10.1016/0019- 1035(81)90209-8

Tornabene, L.L., Watters, W.A., Osinski, G.R., Boyce, J.M., Harrison, T.N., Ling, V., McEwen, A.S., 2018. A depth versus diameter scaling relationship for the best- preserved melt-bearing complex craters on Mars. Icarus 299, 68–83. https://doi.org/10.1016/j.icarus.2017.07.003

Turtle, E.P., Barnes, J.W., Trainer, M.G., Lorenz, R.D., Hibbard, K.E., Adams, D.S., Bedini, P.D., Brinckerhoff, W.B., Cable, M.L., Ernst, C., Freissinet, C., Hand, K., Hayes, A.G., Hörst, S.M., Johnson, J.R., Karkoschka, E., Langelaan, J.W., Lawrence, D.J., Le Gall, A., Lora, J.M., MacKenzie, S.M., McKay, C.P., Neish, C.D., Newman, C.E., Palacios, J., Panning, M.P., Parsons, A.M., Peplowski, P.N., Radebaugh, J., Rafkin, S.C.R., Ravine, M.A., Schmitz, S., Soderblom, J.M., Sotzen, K.S., Stickle, A.M., Stofan, E.R., Tokano, T., Wilson, C., Yingst, R.A., Zacny, K., 2018. Dragonfly: In situ exploration of Titan’s Organic Chemistry and Habitability. Presented at the 49th Lunar and Planetary Science Conference 2018, Lunar and Planetary Institute, Woodlands, Texas.

Turtle, E.P., Pierazzo, E., Collins, G.S., Osinski, G.R., Melosh, H.J., Morgan, J.V., Reimold, W.U., 2005. Impact structures: What does crater diameter mean?, in: Special Paper 384: Large Meteorite Impacts III. Geological Society of America, pp. 1–24. https://doi.org/10.1130/0-8137-2384-1.1

Werynski, A., Neish, C.D., Le Gall, A., Janssen, M.A., 2017. Compositional Variations of Titan’s Impact Craters Indicates Active Surface Erosion.

Wood, C.A., Lorenz, R., Kirk, R., Lopes, R., Mitchell, K., Stofan, E., 2010. Impact craters on Titan. Icarus 206, 334–344. https://doi.org/10.1016/j.icarus.2009.08.021

Chapter 3 3 Conclusions

Titan is a dynamic world eerily like Earth yet still fundamentally different. It’s less than a quarter the size of Earth and nearly 200K colder. As a result, Titan’s surface is covered by a thick icy crust rather than a rocky crust. It has an atmosphere ~1.5 times the pressure of Earth’s and consists mostly of nitrogen and methane. When exposed to solar radiation, these ionize and dissociate the nitrogen and methane. These particles recombine over time to create heavier organic molecules which form the thick haze that makes Titan virtually impenetrable by visible light. These “tholins” slowly fall out and coat the surface with an organic layer. We see global sand dunes made of organics across the equator, but whether these are made from the haze particles is unclear because the grains involved are at least an order of magnitude different in diameter. The organics are not the only thing to fall from the atmosphere; a hydrological cycle with methane as the working fluid can modify the surface. Rivers like those we see on Earth cover the surface. Some of these rivers flow into lakes or, in some cases, seas of liquid methane and ethane. This diverse geologic landscape is uncannily similar to Earth, and this motivates us to investigate how these terrestrial processes operate in such a different environment. In this work, we have reviewed Titan’s crater population to obtain a broad look at the global effect of these processes by assessing how the degree to which Titan’s crater population has been altered by exogenic processes. We used NASA’s Cassini Mission data to identify the craters on Titan’s surface and to constrain their morphology (depth, diameter, and rim heights). We use previously published work on impact craters on icy moons similar in size to Titan to understand what Titan craters would look like without these processes. We use similar icy moons as a template for how impact craters form, and by comparing crater depths and rim heights we can constrain exactly how much degradation has affected the morphology of Titan’s craters.

3.1 Titan’s Craters

We discovered 30 new craters on Titan in this work. This brought the entire population to 90 craters, 50% higher than previous counts based on an incomplete data set. Most of the

79 population is found in the equatorial and mid-latitudes (85% in ±50°). This may be representative of the increased fluvial and lacustrine activity in the higher latitudes, or the presence of a former polar ocean (Neish et al., 2016; Neish and Lorenz, 2014). However, we do caution that in the northern latitudes it is difficult to differentiate between small round lakes and impact craters. We judge that most of these lakes are not impact craters, because it is unlikely that such a high density of craters would impact in such a small area. Furthermore, we are unable to distinguish evidence of ejecta blankets, and while some lakes exhibit elevated rims, they are not rough or bright in radar as with other moons and can be explained through other means (Birch et al., 2018).

The final crater population is consistent with estimates from the last global survey after we adjust for limited coverage (Neish and Lorenz, 2012). The distribution of crater density by size drops for larger craters (+100km) with Menrva standing out as a very large crater (~400km) that is likely to form approximately once in the lifetime of the solar system. The atmosphere disrupts smaller impactors, making it more difficult for smaller craters (~20km) to form (Korycansky and Zahnle, 2005). The atmosphere also erodes smaller craters more quickly than larger ones, making it difficult to identify them after they have been altered (Neish et al., 2016). Our new age estimate for Titan’s surface is consistent with ages less than a billion years and perhaps as low as a few hundred million years.

Titan’s crater depths and rim heights appear not only different but largely degraded compared to the craters on Ganymede. On average, Titan’s craters are ~50% shallower than Ganymede’s. The shallowest craters on Titan are its smallest (10km-30km) at 73% less than those on Ganymede. There is evidence that erosion and infill should affect smaller craters the most (Forsberg-Taylor and Howard, 2004; Neish et al., 2016), but that is not the case for Titan’s medium sized craters which appear ~20% fresher than its largest. This may be indicative of localized erosion and deposition, or it may suggest that Titan’s oldest craters are older.

These depths, measured from rim to floor, are affected by fluvial erosion of the rim, and infill of the crater floor by sand and sediment. Aeolian infilling by sand does not impact the height of the rim in any significant way (Forsberg-Taylor and Howard, 2004).

Therefore, to constrain the amount of fluvial erosion occurring on Titan, we measured the rim heights of the crater. When we break it down to degradation by erosion (rim heights) and infill (terrain to floor depths) we find that erosion plays a larger role than infill; there is a 41% drop in terrain depths and a 54% drop in rim heights. This suggests erosion may play a bigger role than previously thought (Neish et al., 2016, 2013). In fact, infill appears to correlate with higher erosion, suggesting fluvial deposition may contribute more than aeolian infill. Given the distribution of impact crater depths, fluvial erosion has been assumed to play a secondary role to aeolian infilling (Neish et al., 2013). However, we can now use rim height to determine the extent of fluvial erosion in any individual crater. For example, Sinlap (D=88km) and Ksa (45km) are the freshest craters (R=0.31 and R=0.12 respectively, compared to Bray et al., 2012) and exhibit the highest rims observed on Titan. In contrast, Soi (D=85km) is one of the most degraded craters (R=0.81, compared to Bray et al., 2012) with rims only ~100m more degraded than Sinlap’s. Previous studies have suggested Soi’s rims should be moderately eroded compared to Sinlap but its floor significantly more buried by a mix of aeolian infill and fluvial deposition (Neish et al., 2015). We find this is the case, with the bulk of its degradation coming from infill and deposition which virtually flattens it with respect to the terrain.

3.2 Limitations

This work has updated Titan’s crater population and provided new estimates on Titan’s crater parameters (i.e. diameter, depth and rim heights). In doing so, we have better constrained how the different erosional processes (aeolian infill and fluvial erosion) modify Titan’s impact craters. However, it is important to recognize that the data sets we are working with are limited in size and quality. Here we present the assumptions and limitations that went into producing this work.

3.2.1 Assumptions in Crater Mapping

When searching for crater signatures it is not enough to look for suspiciously circular features (SCF). There are distinct signatures that we searched for to identify crater. Well preserved impact craters have distinct shapes and formation features that are consistent regardless of the target body (Pike, 1977; Schenk, 2002). Therefore, we expect

Titan’s craters to act the same way. It is after extensive degradation and infill that impact craters become indistinguishable from SCFs. Smaller craters tend to be more degraded and limited resolution can prevent us from distinguishing them as distinctly as larger craters. Crater ejecta appears to be the first feature to fade, as it is thinner and easier to degrade than the crater rim. This makes it difficult to confidently identify craters in many cases; it is possible there are many more degraded craters on Titan that are unrecognizable from orbit. Furthermore, evidence of large outgassing events present the possibility of large marine regions; marine impacts do not form the types of impact structures that can be easily identified (Neish and Lorenz, 2014). In addition, there are features on Titan with rough round rims and dark interiors that may not be impact craters. In the higher latitudes, methane lakes exhibit elevated rims with floors filled with liquid methane or ethane. Therefore, we are unable to confidently separate any lakes formed by impact cratering from the vast majority of lakes, which are likely formed by other processes (Birch et al., 2018, 2017).

3.2.2 Radar Coverage

Roughly 75% of Titan’s surface has been mapped by Cassini’s high-resolution RADAR instrument. ISS covered the entire surface over a range of resolutions, but these are all well below the resolution of the RADAR instrument. Still, we used ISS data to search in areas of Titan where there is no RADAR data. Past works identified some craters using ISS and VIMS, but we did not find any that are not currently mapped in radar. Our Monte-Carlo approach in determining the crater population does not consider ISS data (Cassini ISS Team, 2015; Porco et al., 2005). This would change the probability of detecting the largest craters on Titan because because our ability to detect them is partly dependent on the resolution of the surface, so with complete coverage the probability of detecting a crater above the resolution threshold approaches 100% (not accounting for distortion by erosion and infill). We don’t expect the adjusted counts to be significantly affected by this given how quickly the probability of detection approaches 100%,.

3.2.3 Comparing Titan with Callisto and Ganymede

An important assumption in our work is that fresh Ganymede craters are analogous to fresh Titan craters. There is a strong basis for this assumption; Ganymede and Titan share several key characteristics. They are each about the same size and density. However, Titan is covered with organics and hydrocarbons (Hörst, 2017; Soderblom et al., 2007). The mixing of impurities into the ice may affect the physical properties of the surface. For example, the methane clathrates in ice is expected to increase its viscosity (Durham et al., 2010). Evidence of outgassing events supports the presence of these impurities, at least in the past (Choukroun et al., 2010; Neish and Lorenz, 2014; Tobie et al., 2006). The effect of these impurities on the physical properties are poorly constrained, and further investigations are required to fully understand their effect on Titan’s crust and impact cratering process.

3.2.4 Topography Data

It is important to recognize that we are limited by a very small dataset to infer trends (14 craters with stereo and SARtopo data). That said, the Cassini extended mission allowed the team to focus on collecting topography data for the strongest crater candidates. Thus, our results represent a lower limit on the amount of erosion that has taken place on Titan, because the more degraded craters were less likely to be targeted for topography data. There are also some concerns about the quality of the topography data itself. SARTopo reports the height every 300m using the average over 10km at each step. This spatial resolution is an order of magnitude lower than that available for Ganymede, so it is possible that the sharpest peaks (e.g. rims) are artificially averaged down to significantly lower heights. However, the Titan stereo data has similar resolution to the Ganymede topography data, and in most cases, it appears to match with SARTopo where the both exist (with two notable exceptions, where the stereo produced lower heights than the SARTopo). We therefore judge that the quality of the SARTopo data is sufficient for comparison of crater depths between Titan and Ganymede.

3.2.5 Methodological Assumptions

We have repeated the steps of analyzing crater depths, diameters and rim heights performed by Bray et al. (2012). We subtract out any slope in the local terrain by fitting a line to terrain 5 radii out from the crater, excluding the crater itself (±1.5 radii from the crater center). We assumed a linear fit because all higher order fits were essentially linear. With the local slope subtracted, crater floor depths and rim heights were measured from the now zeroed terrain height. However, Bray et al. (2012) did not include terrains that were heavily disrupted by the impact crater being studied. Unfortunately, the nature of Titan is that it is rare to find topography that has not been significantly modified. Understanding this modification is the motivation of making these measurements. Therefore, it is important to recognize that because of this modification, the measurements we are taking may be skewed in ways that do not occur on Ganymede. This is the cause of the high level of uncertainty in our results, and it may be the cause of the ambiguity we observe between the different average depths and rim heights.

3.2.6 Statistical Assumptions

The goodness of fit test we use, the KS-Test, is valid even for the small populations used in this work. Nevertheless, a higher sample size would increase the quality of the results by better representing the entirety of Titan’s impact craters. It is worth noting that these tests are being run for N=12 for Titan and N=66 for Ganymede. Even within Ganymede’s 66 craters, anomalies exist (Figure 2.9).

3.3 Future Work

There are several ways to improve this work. This work has increased the population of known craters on Titan by 50%. We have improved our understanding of how degraded Titan’s craters are by updating previous works and contributing brand new results to study. These results present the opportunities for further study of Titan’s surface and ways of overcoming the limitations associated with our work.

3.3.1 Landscape Evolution Models

One obvious step forward is to use our results to improve landscape evolution models. Broadly speaking, our results constrain the level of infill and erosion by comparing Titan’s crater depths and crater rim heights to those of Ganymede. Past landscape evolution models have demonstrated that fluvial erosion is a viable candidate for crater modification on Titan (Neish et al., 2016). Kinser (2016) used a similar approach to test the effect of different parameters (e.g. erosion rate and grain size) and put constraints on how fluvial erosion is occurring on Titan by modeling the erosion and tracking the relative crater depths over time. Her goal was to identify which scenarios produced results that match the craters we see on Titan given the time constraints defined by Neish et al. (2016). Crater depth is dependent on both rim erosion and infill by deposition and aeolian processes, but a similar approach can be used to model the change in rim heights (and potentially central peak heights).

Comparing rim heights would test a parameter independent of infill such that we can constrain the level of erosion without the infill to bias the erosion levels. This would 1) achieve better constraints on the different fluvial erosion parameters and 2) provide estimates on how much infilling has occurred by deposition versus aeolian infill. Recent efforts have constrained how material is transported across Titan’s surface (Malaska et al., 2016). Prior works suggested infill was the main modifier of Titan’s craters (Neish et al., 2013). These results would help constrain the sediment flux near craters. Lastly, we can assess how significantly the thermal insulation from the infill will effect viscous relaxation of the crater (Schurmeier and Dombard, 2018).

3.3.2 Comparing Radar Images and Erosion

Titan has 90 impact craters that we know of, yet only 12 of these have known depths and rim heights. To understand the level of degradation of the other 78 impact craters, we could visually quantify the level of erosion, whether by number or lengths of rivers or radar albedo (brightness), then compare this with the level of degradation measured by the crater depths and rim heights. We could build a relationship between the erosion and radar to constrain the depths of the other craters. We could also look at impact craters on Earth or

Mars, and conduct a similar comparison, to understand erosion of impact craters more generally where more data is available as well as the ability to ground truth our results.

3.3.3 Updated ISS Map

The purpose of this work was to finalize Titan’s crater population using the now complete Cassini dataset including ISS global coverage. We are limited by what we can observe through ISS because it is of lower resolution than RADAR, and only the largest craters are visible in this data set. However, ongoing work by members of the Cassini team has produced a significantly improved ISS map of Titan (Turtle et al., 2018b). The data is not yet released to the public, but it may reveal the existence of more craters in regions without RADAR coverage. This will improve our coverage and test our predictions Titan’s crater population using the Monte Carlo method.

3.3.4 Updated SARTopo

Ongoing efforts in the Titan community are in the works to improve the coverage of SARTopo (Mastrogiuseppe, 2018). Mastrogiuseppe (2018) proposes that wherever we have radar coverage we can have full coverage topography (rather than thin lines) by splitting the bins of the bursts of signals that are taken of the surface for each image pixel. Theoretically, halving the bins would halve the quality of the data and the resolution, but he seems to think this can be done with much less loss of quality. While this work is ongoing, it is promising. It would drastically improve our ability to constrain Titan’s crater morphologies far beyond what we have done here. A larger population of fresh craters could be measured to obtain a trend of crater depth to diameters similar to those done on other icy moons (Bray et al., 2012; Schenk, 2002). This could be used to test whether fresh Titan craters on Titan really do follow similar trends to Ganymede and Callisto, assuming enough fresh craters exist. Furthermore, a larger collection of crater parameters can be measured for better statistical analysis, and better averaging of individual craters could be taken rather than a single or two cross-sections. Lastly, it would allow for more landscape evolution modeling of entire craters (e.g. Neish et al., 2016).

3.3.5 Measuring Stereo Rim Heights

This work updated Titan’s crater population using new data from Cassini. In this work, we updated SARTopo measurements using new SARTopo data. We did not perform measurements of stereo data because previous work has already been done and no new data was available. However, we choose to remove rim heights on Titan as a way of removing possible bias in the SARTopo measurements (from lower resolution). Once we demonstrated that bias is probably negligible, we were able to present the first rim height measurements on Titan from any dataset. However, measuring the rim heights from stereo data would further test that hypothesis by comparing the results with SARTopo, and it could be done using methods consistent with Bray et al. (2012) methods of taking multiple cross sections across the crater. This would expand the list of craters we have rim height measurements for. With such a small population, every possible point is useful for interpreting the surface. Lastly, these results can be compared to SARtopo results by testing the effect of resolution when measuring a simulated cratered shape to better constrain the error even if it is minimal per our results (e.g. Stewart and Valiant, 2006).

3.3.6 Dragonfly

Dragonfly is one of two missions now being considered in NASA’s New Frontiers program (Turtle et al., 2018a). Dragonfly is a proposed rotorcraft that would go to Titan and explore the surface chemistry, geology, interior, and atmosphere. The objective is similar to a Mars rover but designed to take advantage of Titan’s thick atmosphere. The primary goal of the mission is to determine Titan’s potential for life. It includes a range of instruments that would also aid in interpreting the geology of its surface, and this level of scrutiny would help us constrain the processes acting on any impact crater visited by the mission. This type of in-situ test would give us an entirely new level of insight into impact craters on Titan (Neish et al., 2018).

There is no shortage of opportunities moving forward with these new discoveries we present here. Cassini may have ended, but our work on Titan has only just begun.

3.4 References

Birch, S.P.D., Hayes, A.G., Dietrich, W.E., Howard, A.D., Bristow, C.S., Malaska, M.J., Moore, J.M., Mastrogiuseppe, M., Hofgartner, J.D., Williams, D.A., White, O.L., Soderblom, J.M., Barnes, J.W., Turtle, E.P., Lunine, J.I., Wood, C.A., Neish, C.D., Kirk, R.L., Stofan, E.R., Lorenz, R.D., Lopes, R.M.C., 2017. Geomorphologic mapping of titan’s polar terrains: Constraining surface processes and landscape evolution. Icarus 282, 214–236. https://doi.org/10.1016/j.icarus.2016.08.003

Birch, S.P.D., Hayes, A.G., Hofgartner, J.D., 2018. The Raised Rims of Titan’s Small Lakes. Presented at the 49th Lunar and Planetary Science Conference, Lunar and Planetary Institute, Woodlands, Texas.

Cassini ISS Team, 2015. Map of Titan, Cassini ISS Global Mosaic. https://astrogeology.usgs.gov/search/map/Titan/Cassini/Global- Mosaic/Titan_ISS_P19658_Mosaic_Global_4km

Hörst, S.M., 2017. Titan’s atmosphere and climate. Journal of Geophysical Research: Planets 122, 432–482. https://doi.org/10.1002/2016JE005240

Kinser, R.M., 2016. Geological Conditions Required for the Fluvial Erosion of Titan’s Impact Craters. Forida Institute of Technology, Melbournem, FL. Thesis.

Korycansky, D.G., Zahnle, K.J., 2005. Modeling crater populations on Venus and Titan. Planetary and Space Science 53, 695–710. https://doi.org/10.1016/j.pss.2005.03.002

Malaska, M.J., Lopes, R.M., Hayes, A.G., Radebaugh, J., Lorenz, R.D., Turtle, E.P., 2016. Material transport map of Titan: The fate of dunes. Icarus 270, 183–196. https://doi.org/10.1016/j.icarus.2015.09.029

Mastrogiuseppe, M., 2018. Recent Results from T126 Data Analysis and T108 Update.

Neish, C.D., Lorenz, R.D., 2014. Elevation distribution of Titan’s craters suggests extensive wetlands. Icarus 228, 27–34. https://doi.org/10.1016/j.icarus.2013.09.024

Neish, C.D., Lorenz, R.D., 2012. Titan’s global crater population: A new assessment. Planetary and Space Science 60, 26–33. https://doi.org/10.1016/j.pss.2011.02.016

Pike, R.J., 1977. Size-dependence in the shape of fresh impact craters on the moon., in: Roddy, D.H., Pepin, R.O., Merrill, R.B. (Eds.), Impact and Explosion Cratering. Pergamon Press, New York, pp. 489–509.

Schenk, P.M., 2002. Thickness constraints on the icy shells of the galilean satellites from a comparison of crater shapes 417, 3.

Schurmeier, L.R., Dombard, A.J., 2018. Crater relaxation on Titan aided by low thermal conductivity sand infill. Icarus 305, 314–323. https://doi.org/10.1016/j.icarus.2017.10.034

Stewart, S.T., Valiant, G.J., 2006. Martian subsurface properties and crater formation processes inferred from fresh impact crater geometries. Meteoritics & Planetary Science 41, 1509–1537. https://doi.org/10.1111/j.1945-5100.2006.tb00433.x

Tobie, G., Lunine, J.I., Sotin, C., 2006. Episodic outgassing as the origin of atmospheric methane on Titan. Nature 440, 61–64. https://doi.org/10.1038/nature04497

Turtle, E.P., Perry, J.E., Barbara, J.M., Del Genio, A.D., Sotin, C., Rodriguez, S., Lora, J.M., Faulk, S., West, R.A., Karkoschka, E., McEwen, A.S., Mastrogiuseppe, M., Hofgartner, J.D., Corlies, P., Kelland, J., Hayes, A.G., Pitesky, J., Ray, T.L., 2018b. Titan Insights During the Final Year of the Cassini Mission. Presented at the 49th Lunar and Planetary Science Conference, Lunar and Planetary Institute, Woodlands, Texas.

Appendices

+ퟐퟓ A1: Menrva crater (푫 = ퟒퟎퟎ−ퟑퟎ풌풎) topography that was averaged for profiles 1-4.

A2: Forseti (푫 = ퟏퟐퟓ ± ퟒퟎ풌풎) topography that was measured using a single profile.

A3: Afekan crater (푫 = ퟏퟏퟓ ± ퟓ풌풎) topography that was averaged for profiles 1-3.

+ퟏퟎ A4: Hano crater (푫 = ퟏퟎퟓ−ퟐퟎ풌풎) topography that was averaged for profiles 1-3.

A5: Sinlap crater (푫 = ퟖퟖ ± ퟏ풌풎) topography that was averaged using a single profile.

A6: Soi crater (푫 = ퟖퟓ ± ퟏퟓ풌풎) topography that was averaged using a single profile.

100

A7: Selk crater (푫 = ퟖퟒ ± ퟐ풌풎) topography that was averaged for profile 1, and profiles 2-4 are observing the crater rim.

101

102

103

A8: Crater #1 H2018 (푫 = ퟕퟎ ± ퟐퟓ풌풎) topography that was averaged using a single profile.

104

A9: Shikoku crater (푫 = ퟒퟏ ± ퟏ풌풎) topography that was averaged for profiles 1-3.

105

106

A10: Ksa crater (푫 = ퟒퟓ ± ퟐ풌풎) topography that was averaged for profile 3. While closest to the edge, it shows the freshest profile of the crater. Profiles 1-2 shown for comparison.

107

108

A11: Crater #10 NL (푫 = ퟐퟓ ± ퟐ풌풎) topography that was averaged using a single profile.

109

A12: Crater #3 NL (푫 = ퟐퟎ ± ퟓ풌풎) topography that was averaged using a single profile.

110

Curriculum Vitae Josh Hedgepeth Email: [email protected], Phone: 404-713-2463

Post- Secondary Education 2018 – 2022 University of Western Ontario Accepted into, PhD. in Geophysics and Planetary Science, Start September 2018

2016 – 2018 University of Western Ontario Candidate for MSc. in Geophysics and Planetary Science, Graduation August 2018

Professional Experience

September 2016 – July 2018 Graduate Research and Teaching Assistant Department of Earth Sciences, University of Western Ontario

January 2015 – July 2016 Undergraduate Research and Teaching Assistant Department of Earth and Atmospheric Sciences, Georgia Institute of Technology

Honors and Awards

President's Undergraduate Research Awards January 2016 The Effect of Existing Geologic Features on the Formation of Europa’s Chaotic Terrains Total budget: $1,500

Technologies for Exo-planetary Sciences CREATE Program 2017-2019 Impact Craters on Titan: Finalizing Titan’s Crater Population Yearly budget: $8,000

Publications in Progress

Hedgepeth, J.E, Walker, C., Schmidt, B.E. (2018) Constraining the Yearly Stresses in Helheim Glacier though Crevasse Mapping. in prep.

Hedgepeth, J.E, Schmidt, B.E. (2018) Relating the Extent of Resurfacing Events in Europa’s Chaotic Terrains to Local Geology. in prep.

Hedgepeth, J.E., Neish, C.D., Turtle, E.P., Stiles, B.W. (2018) Impact Craters on Titan: Finalizing Titan’s Crater Population in prep