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Table of Contents

Future documents to follow: • Ocean Worlds Missions Scenarios, Roadmaps & Technologies Goals, Objectives and Investigations for Ocean Worlds

A. Introduction.

The Roadmaps to Ocean Worlds team. At the OPAG meeting in February 2016, it was decided to form a Roadmap to Ocean Worlds group, with the following charter: • Identify and prioritize science objectives for Ocean Worlds (tied to the Decadal Survey) • Design roadmap to explore these worlds to address science objectives (Mission sequences, sustained exploration effort) • Assess where each Ocean World fits into the overall roadmap • Summarize broad mission concepts (Considering mission dependences & international cooperation) • Recommend technology development and detailed mission studies in support of the next decadal survey

The team is co-chaired by Terry Hurford and Amanda Hendrix, who organized a large team of individuals with expertise in the various related disciplines, including small bodies topics normally covered by SBAG, to provide inputs for this and future reports. The ROW team membership is detailed in Appendix D1.

Definition of Ocean World

Background, Philosophy and Major Finding There are several – if not many – ocean worlds or potential ocean worlds in our solar system, targets for future NASA missions in the quest for the understanding of the distribution of life in the solar system. This document lays out the science questions and investigations to be addressed for each of those targets. This document is designed to be [the first part of] a roadmap for charting the course to search for life at ocean worlds in our solar system. This document outlines the goals, objectives and investigations of NASA’s Ocean Worlds program as recommended by the Roadmaps to Ocean Worlds group chartered by OPAG with the support of SBAG.

In considering ocean worlds in the solar system, there are several candidate ocean worlds (exhibiting hints of possible oceans) and several confirmed ocean worlds, in addition to those worlds that may potentially harbor oceans. As a philosophy, the ROW team considers it critical to consider all of these worlds in order to understand the origin and development of oceans and life in different worlds: why does life originate and take hold at some ocean worlds and not others? Thus, the ROW team urges NASA to create a program that studies that spectrum of Ocean Worlds; if only 1-2 ocean worlds are explored and life is discovered (or not), we won’t fully understand the distribution of life and the repeatability of its occurrences in the solar system.

We have considered that Enceladus, Europa, Titan, Ganymede and Callisto have known subsurface oceans, as determined from measurements by the Galileo and Cassini spacecraft. These are confirmed ocean worlds. Titan is a unique case because of its surface liquids, which are relatively easily accessible, but any life they harbor is not water-based (unless hot spots at the bottoms of lakes have melted the crust so that liquid water is in contact with the non-aqueous lake liquids). Titan, Ganymede and Callisto’s subsurface oceans are expected to be covered by a relatively thick ice shell, with no surface evidence of the oceans as on Europa and Enceladus. Triton, Pluto and Ceres are considered to be possible ocean worlds based on hints from limited spacecraft coverage. For other bodies such as Miranda, our knowledge is limited enough that the presence of an ocean is uncertain but they are deemed credible possibilities.

The ROW team decided on an overarching goal for NASA’s Ocean Worlds program: Identify ocean worlds, evaluate their habitability, and search for life. This overarching goal has four underlying goals: 1) determine which bodies have oceans and understand how to determine whether other bodies host current oceans; 2) characterize the oceans; 3) characterize the habitability of the oceans; and 4) understand what kind of life could be present in these oceans and how to search for it.

The goals are described in detail, along with corresponding objectives and investigations, in sections below. Applications to specific prominent solar system targets are provided in the Appendices. Figure 1 demonstrates the state of knowledge of each objective, for potential target bodies. Goals and Objectives are linked to Decadal Survey goals in Table 1. Goals, Objectives and Investigations are listed in Table 2.

A major finding of this study is that in order to map out a coherent Ocean Worlds program, significant input is required from studies here on Earth: a rigorous R&A program is called for, to enable future Ocean Worlds missions to be thoughtfully planned. Research objectives/investigations involve questions that can be addressed here on earth – through modeling, field studies, lab work etc. so that spacecraft data can be undertaken and interpreted properly. The objectives laid out in this document cover both those that include measurements required to be made at the various target bodies, and measurements/studies that will need to be made here on Earth to prepare for those robotic measurements and to help in their interpretation. Thus the ROW team recommends a rigorous R&A program as part of the Ocean Worlds program. The Objectives that focus on R&A work, rather than on measurements made using spacecraft missions, are indicated in Table 2 and discussed in Sec. C. Table 1. Serena Diniega will work on this table. Mapping of Decadal Survey themes to Objectives shown in the Goals, Objectives, Investigations table (Table 2) (this is important in particular for the midterm review) Decadal Survey Cross-cutting science Relevant Ocean Worlds Objectives theme

Decadal Survey Satellites science theme Relevant Ocean Worlds Objectives

Decadal Survey small bodies/KBOs? Relevant Ocean Worlds Objectives science theme

Figure 1. Investigations Roadmap: demonstrating the state of knowledge for each (potential) target world; an alternate version could also be about the capability of each target to have that question answered. This is an idealized version; the real version will be included here.

B. Goals, Objectives, Investigations The overarching goal is to understand whether these worlds have oceans (when considering where to look for life, we need to know which targets have oceans), understand habitability within the oceans, and ultimately find whether the oceans harbor extant life.

Table 2. – See attachment

Goal I. Identify ocean worlds in the solar system. Before sending spacecraft to target bodies to search for life within the ocean, we must first determine the presence of an ocean. There are several questions that can be addressed in order to determine the presence of an ocean. For the known ocean worlds (Europa, Enceladus, Titan, Ganymede, Callisto), these questions have already been answered – or enough of the questions have been answered that the presence of an ocean is (reasonably) certain.

A. Is there a sufficient energy source to support a persistent ocean? A.1 Is there gravitational energy from a parent planet or satellite? A.2 Is there remnant radiogenic heating? A.3 Can the planet or satellite convert available tidal energy into heat? A.4 Is the planet or satellite’s orbital or rotational properties favorable to tidal dissipation?

Energy sources are perhaps the single most fundamental requirement for the maintenance of a present-day ocean on an otherwise frozen world. The identification of ocean worlds therefore requires identification of possible energy sources. Both radiogenic heating (e.g., for Ganymede, Callisto, and Titan) and tidal energy (e.g., for Europa, Enceladus) play a role in sustaining oceans. Available energy sources can be identified either through modeling or direct observation (or ideally a combination of the two). Theoretical modeling is an invaluable tool for predicting which bodies can sustain ocean worlds. Such models anticipated oceans on icy moons (e.g., Europa) long before such oceans were ever actually detected. However, modeling alone can lead to misleading results. In the case of Enceladus, theoretical models indicate that efficient conversion of tidal energy to heat should occur within the moon. Observations by the Cassini spacecraft have demonstrated that such heating does occur, but the moon emits more an order of magnitude more energy than theoretical models predicted. Identifying the sources of Enceladus’ energy is an open and active area of research.

Evaluating sources of energy requires addressing, at a minimum, the four sub- questions listed above. For the largest satellites, remnant radiogenic heating may be sufficient to maintain an internal ocean (A2), depending on the initial radiogenic content of the rock component, and the state of the overlying ice shell. For smaller bodies (e.g., Enceladus) dissipation of tidal energy is critical. Dissipation of tidal energy requires the presence of parent planet or satellite with sufficient gravitational energy to deform the body (A1). Pluto and Charon lack a source of such energy, so the energetics that permit a long-lived ocean are still in question. Additionally, the body’s orbit and/or rotation must be favorable to tidal dissipation, possibly through a high eccentricity (e.g., Europa), libration, or obliquity (e.g., Triton) (A4). However, these two requirements are insufficient to ensure internal oceans, as the planet or satellite must be able to convert available tidal energy to heat (A3). This is demonstrated by the satellite Mimas, which despite its high eccentricity dissipates little tidal energy, likely because it interior has remained cold since shortly after its formation. The complex feedback between the orbital/rotational evolution of potential ocean worlds (A4) and their dissipativeness (i.e., internal structure) (A3) requires careful theoretical modeling.

B. Are signatures of ongoing geologic activity (or current liquids) detected? B.1 Do signatures of geologic activity indicate the possible presence of a subsurface ocean? (surface hotspots, plumes, crater-free areas, volcanoes, tectonics) B.2 Does the body exhibit tidal and/or rotational evidence indicating the presence of a sub-surface ocean? B.3 Does the gravity and topography of the body indicate the presence of a sub-surface ocean? B.4 Are temporal changes observed at the body that would indicate the presence of a sub-surface ocean? B.5 Is there an atmosphere or exosphere that could be linked with the presence of a sub-surface ocean? B.6 Does the electromagnetic response of the body indicate the presence of a sub-surface ocean? B.7 Can the surface composition be linked with the presence of a sub- surface ocean? B.8 Is the signature of a surface liquid observed (e.g. specular reflection)?

Over the past three decades, numerous techniques have been developed for accessing whether subsurface oceans are present on an icy world. In some cases, investigation of a satellite’s surface is sufficient to infer the presence of an ocean below. Recent or ongoing geologic activity such as a young tectonized surface, hotspots, and plumes are indicative of a warm interior that can potential sustain an ocean (B1). For example, the plume of Enceladus, along with the young, crater-free terrain and warm fractures from which it emanates are strong indicators of an ocean, even in the absence of other geophysical data. Likewise, surface change (B4) would indicate ongoing geologic activity, again requiring warm interior. Surface composition (B7) can also indicate a subsurface ocean through the presence of chemical species originating in an ocean. In the case of Europa, consensus on the origin of identified chemical species remains elusive; however, improved spectral and spatial resolution data is likely to resolve the question. In rare cases where liquids may be present on a body’s surface rather than within its interior (e.g., Titan) liquids can be detected through their optical and radar properties, or by specular reflection (B8).

Not all ocean worlds reveal their present oceans in their surface characteristics. Ganymede and Callisto both have internal oceans, but their surfaces are currently inactive. For these worlds, and to confirm oceans on geologically active worlds, a number of geophysical measurements can be used to identify present-day oceans. In many cases, oceans can be revealed by the orbital and rotational state of a body (B2), if it can be measured carefully enough. For example, the magnitude of Enceladus’ physical libration requires the presence of a global ocean. Titan’s subsurface ocean is also revealed by differential rotation of its outer shell, which must be decoupled from its interior. For systems with a strong, inclined magnetic field (e.g., the Jupiter system), the electromagnetic response of the body provides a strong indication of an internal ocean (B6), as demonstrated for Europa, Ganymede, and surprisingly, Callisto. With sufficient flybys, gravity data can also indicate the presence of an ocean, especially when coupled with detailed topography (B3), as recently demonstrated for Dione.

Implicit in the discussion above is the necessity of spacecraft data to characterize the surfaces of bodies and acquire geophysical data. In some cases, ground-based data can aid our understanding. This is especially true for the monitoring of a satellite’s exosphere for potential plume activities (as in the case of Europa) (B5). However, unambiguously identifying ocean worlds requires detailed in-situ investigations.

C. How do materials behave under conditions relevant to any particular target body?

C.1. What are the phase relations of material composing ocean worlds at relevant pressures and temperatures?

C.2. What is the composition and chemical behavior of materials composing ocean worlds?

C.3. What are the rheologic mechanisms by which material deforms under conditions relevant to ocean worlds?

C.4. How does energy attenuation/dissipation occur under conditions relevant to ocean worlds? C.5 What are the thermophysical properties of material under conditions relevant to ocean worlds?

This is a “fundamentals” objective, discussed in Sec C.

Goal II. Characterize the ocean of each ocean world

A. Characterize the ocean: physical properties; ice shell thickness, depth of ocean, currents, gradients (T, salinity, pH, P, composition, radiation)

A.1 What is the thickness, composition (including the presence of any organics), porosity of the ice shell (crust) and how do these properties vary spatially and /or temporally? A.2 What is the thickness, salinity, density and composition of the ocean? How do these properties vary spatially and /or temporally? A.3 Characterize the seafloor

B. Characterize the ocean interfaces B.1. Characterize the seafloor; High-pressure ocean – silicate interaction B. 2. What is the thickness, composition (including the presence of any organics), porosity of the ice shell (crust) and how do these properties vary spatially and /or temporally? B.3. Characterize the ice-ocean interface

Add explanatory text from oceanography/interfaces group (German, Singer, Bowman, Rhoden, Schmidt)

Goal III. Characterize the habitability of each ocean world

A. What is the availability (type and magnitude/flux) of energy sources suitable for life, how does it vary throughout the ocean and time, and what processes control that distribution?

A.1 What environments possess redox disequilibria, in what forms, in what magnitude, how rapidly dissipated by abiotic reactions, and how rapidly replenished by local processes? A.2 (Where) is electromagnetic radiation available? In what wavelengths and intensity?

Life on Earth utilizes, as sources of energy, light within the visible to near-IR wavelength range and the Gibbs energy released in specific (mostly oxidation- reduction) chemical reactions. The present understanding of biological energy metabolism indicates that chemical energy sources must satisfy discrete minimum requirements for both Gibbs energy change (∆G) and power (flux of energy through time) in order to be useful. Light energy also must satisfy a discrete minimum requirement for flux (corresponding to light intensity; the requirement equivalent to ∆G is easily satisfied in any of the part of the wavelength range used by life). Additionally, the flux of energy constrains, in a direct relationship, both the maximum rate of new biomass synthesis (productivity) and the maximum quantity of standing biomass that can be sustained in steady state. That is, environments having greater energy flux can potentially support more abundant life, and might therefore be better targets for life detection.

Investigations B.1 and B.2 call for characterization of the availability of the two forms of energy known to be utilized by life on Earth (B1: chemical; B2: light). B1: Spacecraft observations that constrain the concentrations of redox-active species within the liquid environment will support calculation of Gibbs energy yields (∆G) associated with specific redox couples, and thereby identify metabolisms that satisfy the biological ∆G requirement. Assessment of energy flux will require spacecraft observations that constrain the rate of delivery of specific chemical species into the liquid environment – for example, delivery of oxidants into an ocean by overturn of surface ice and corresponding delivery of reductants by water-rock reaction. B2 requires observations or models that constrain the spectral character and intensity of light available within the liquid environment. In general, a kilometers thick ice cover will preclude solar influx, but transiently or locally thinner ice cover and black body radiation (e.g., from hydrothermal vents) might allow for some introduction of light into liquid water habitats. Moreover, Titan’s hydrocarbon lakes, should they prove to represent a solvent suitable for life, receive direct solar irradiation.

Discuss: II.A.1 is R&A; R&A + spacecraft measurements needed for II.A.2

B. What is the availability (chemical form and abundance) of the biogenic elements, how does it vary throughout the ocean and time, and what processes control that distribution?

B.1 What is the inventory of organic compounds, what are their sources and sinks, and what is their stability with respect to the local environment? B.2 What is the abundance and chemical form of nitrogen, oxygen, phosphorus, sulfur, and inorganic carbon, what are their sources and sinks, and are there processes of irreversible loss or sequestration relative to the liquid environment?

The biochemistry of life on Earth is built around a core of elements – C, H, N, O, P, and S – that are required by all known organisms, as well as a variety of other elements (e.g., specific transition metals) that are required by specific subsets of life. In Earthly environments where life’s requirements for water and energy are abundantly met – for example, the sunlit portions of Earth’s aquatic environments – the distribution of these elements (N, P, and Fe, in particular) can directly limit the abundance and productivity of life. Each of these elements can be incorporated into a diversity of chemical forms, some of which may be less accessible or inaccessible to biology.

Investigations B1 and B2 seek to establish the availability of these elements, in terms of both chemical form and abundance, and to characterize the processes that govern that availability. B1 emphasizes the importance of organic compounds, whose presence could be indicative of the stability of specific chemical bonds or functional groups in the host environment, or which might plausibly serve as feedstock for biochemistry. Spacecraft observations that constrain the abundance, molecular character, and distribution in chemical space of organics within the liquid environment will directly support this investigation. B2 targets the availability of the remaining biogenic elements and of inorganic carbon. Spacecraft observations that constrain the abundance and chemical form of compounds bearing these elements in the liquid environment will directly support this investigation.

Goal IV. Understand how life might exist at each ocean world and search for life

A. What are the potential biomarkers in each habitable niche? (determine what we’re looking for) A.1 What can we learn about life on ocean worlds from studying life on Earth? A.2 What niches for life are possible on ocean worlds? A.3 What can we learn about life by understanding the history of ocean worlds from their formation to the present? A.4 What should be our target indicators? (Life Detection Ladder) A. 5 How do we distinguish extant from extinct life in environments in which life might develop, and which timescales (e.g., for metabolism, reproduction, dormancy) matter?

Earth is the only planet where life is known to exist. Analogue studies of life in Earth’s oceans and other habitable niches provides our only anchor for extrapolations to other ocean-bearing worlds. Example of what can be learned from such studies include: (1) the range of physical (e.g. temperature, pressure, radiation levels) and chemical (e.g. pH, redox, salinity/water activity, major/trace elemental abundances) conditions that life tolerates; (2) whether known extremophiles exist at the temperatures, pressures, pH, radiation levels, etc. found on the ocean world; (3) whether there exist areas on Earth that do not support life (and why); (4) how long it takes ecosystems to colonize a given environment, or adapt/evolve to changing physicochemical conditions; (5) the metabolic diversity that might be expected in ocean world environments; and (6) the amount of potential biomass that could be sustained on a given ocean world.

Niches for life provide a solvent, energy, and nutrients for a sufficient amount of time. By definition, ocean worlds provide a solvent, and the materials planets form from is thought to contain nutrients. Bioavailable energy requires the co-location, in chemical disequilibrium, of electron donors and acceptors. Determining which metabolic strategies are possible, which dominate, and what their spatial and temporal distribution is involves finding out which electron donors (EDs) and acceptors (EAs) are present, what their abundances are, whether there exist mechanisms to bring them together, and the quantification of sources, sinks, and their variation in space and time.

Our current knowledge of biology is entirely based on life on Earth, but finding life elsewhere would provide other data points to infer universal properties of life. Key considerations include the emergence of life, how long a world has been inhabited, whether (and which) evolutionary pressures might have driven or prevented life’s diversification, and how geological and environmental changes could affect habitability over time.

Ideal indicators are specific to life (not found in abiotic systems), universal (not limited to life on Earth), and easy to detect. To date, we do not know of any single indicator that satisfies all three criteria. Darwinian evolution and reproduction may be the only specific and universal indicators, but we cannot measure them even for most of life on Earth, which we cannot culture. Growth (concurrent life stages) and activity (motility, feeding, biofilm formation) are specific, but not universal. Evidence for metabolism (isotopic fractionations from abiotic values, co-location of electron donors and acceptors) can be more easily detected, but its lack of specificity requires excellent contextual knowledge. Functional molecules (nucleic or amino acid polymers, pigments) are highly specific and easy to detect, but not universal. Potential biomolecule components (nucleic or amino acids, lipids, sugars), with structural preferences (non random chirality, carbon number, or trace element compositions) may be universal and easy to detect, but have low specificity.

Any extant life on ocean worlds is likely to be present as single-celled prokaryotes (and potentially eukaryotes). Micrometer-scale imaging can confirm the presence of cells irrespective of their chemistry, and fluorescence imaging with dyes specific to biochemicals (e.g. lipids, proteins, nucleic acids, sugars) can assess the co-location of chemicals relevant to life and cell-like features. The required resolution for such imaging should be the same on ocean worlds with liquid water as on Earth: the upper size limit for single celled organisms is set by diffusion rates similar in liquid water everywhere, and the lower limit is set by the need to contain a self-replicating genome and the chemistry required to manufacture cellular proteins.

As on the Viking missions, metabolism can be observed by supplying electron donors/acceptors and labeled nutrients to a sample, and monitoring abundances of labeled products over time. This approach would benefit from simultaneous micro- and macroscopic imaging of the reaction volume. Microscopic imaging could detect growth, reproduction, motility (easily distinguished from brownian motion), and taxis, and assess their correlation with chemical measurements of metabolic activity. Macroscopic imaging could perform the equivalent of colony counting. Rapid changes in activity can be stimulated by modest temperature changes or nutrient addition; although obvious changes, such as the regrowth of flagella from dormant bacteria, can take several hours.

On Earth, microevolution can be very rapid in single-celled organisms. Organisms with random mutations might develop into distinct populations with different nutrient preferences. Thus, evolution might be detected within days or weeks, by supplying a sample container heterogeneously with a variety of chemical stimuli, and measuring cell concentrations in different regions. Short of being able to induce evolution, a proxy is the observation of multiple species in a population, likely to arise in the presence of different habitable niches. Microbial diversity might be detected with short-term observations of differences in activity levels with different nutrient sources, motility styles (non-motile, motile, various swimming mechanisms), or even structures for organisms larger than a few microns.

B. How to search for and analyze data in different environments? B.1 How can we look for life on an ocean world remotely (from orbit or during a flyby)? B.2 How can we look for life on an ocean world in situ (landed, underwater, plume) investigations? B.3 How can we look for life on an ocean world with sample return science? B.4 Which science operational strategies should be used to detect life on ocean worlds?

Although remote detection of life seems challenging, the presence of life might be expressed at the surfaces of ocean worlds. Currently, we do not know what kind of evidence might be seen, at what abundances, or how this evidence might be modified by radiation processing and oxidation. It would also be useful to determine to what extent remote spectroscopy techniques could 3+ measure co-located electron donors (e.g. H2 ± CH4) and acceptors (e.g. O2, nitrate, Fe , CO2), and whether the geological context or other indicators allow the inference of how long this co- location has persisted over time. On airless worlds, measuring remotely the surface distribution of elements and looking for any deviations from background bulk concentrations could represent another possible avenue, but relevant spatial scales have not been constrained. As a final example, remote spectroscopy techniques could detect pigments and/or other specific biomolecules.

Possible in situ investigations for life could seek to determine whether the inventory of detected organic molecules differs from those expected to be synthesized by abiotic chemistry; look for morphological signatures that indicate microenvironments containing chemical gradients; active, sharp physicochemical gradients of metabolic interest in ice or water columns (e.g., gradients in pH, Eh, temperature); evidence for enzymatic catalysis in the formation pathways of detected organic molecules (e.g. stable isotopic fractionations of CHNOPS that differ from those of abiotic systems over the full range of plausible habitat temperatures); or chemical cycling (organic or inorganic) or electrochemical / electrical activity not explained by abiotic processes that could indicate metabolic activity. Planned in situ investigations could much benefit from lessons learned from previous in situ life-detection experiments, such as the Viking suite. In situ investigations can also search for extant preserved or dormant life in near-surface ice that results from upwelling through cracks; viable microorganisms have been found preserved in ancient ice on earth. Search for extant life in situ should also include surveys with optical microscopy - the temperatures and chemistry of environments where extant life is likely to form (liquid water environments), the sizes of organisms will be constrained by the same effects: chemical diffusion rates will set the upper limit, and the minimum amount of chemistry (information molecules and supporting chemistry) required for self-replication will set the lower limit. Fluorescent dyes that are specific to various classes of molecule without being single-molecule specific can be used to determine whether cell-like objects contain likely biotic chemistry, such as lipids, proteins, amino acids, and sugars. Various stimuli (chemical, photo, magneto, photo, thermo) can be applied to samples to induce changes in activity levels, such as growth, reproduction, or motility.

Possible investigations for life in the samples include determining (1) the distribution of detected organics, and whether it differs from that expected to be synthesized by abiotic chemistry; (2) signatures in the conformation (e.g., chirality) of specific organics that cannot be obtained abiotically; (3) isotopic compositions of inorganic compounds, as well as of H, C, N, O, and/or S in specific organic compounds, and any measurable deviations from expected abiotic compositions (4) direct observation of microorganisms. To inform these investigations, in situ contextual measurements at the time of collection are essential.

Key considerations for life detection include the choice of sampling location and sampled material (e.g. rocks, ice, water, soil, interface zones), contamination control, and meeting planetary protection requirements. Previous searches for life remotely, in situ, and in samples on Earth have taught us that such searches are highly path-dependent: decisions on which measurements to make depend on the results of previous measurements, and it is difficult to predict a priori which measurements will be needed. It would be helpful to identify relevant sets of complementary observations and possible decision trees for observation types. Also crucial is the ability to distinguish environments that are pre-biotic, host extant life, and post-biotic (harbored now extinct life). This requires reducing the possibility of false positives due to forward chemical and biological contamination (separately from planetary protection), by characterizing contamination signals and distinguishing them from indigenous signals. It also requires constraining the states in which we might we find evidence of life (e.g. live, dead, stasis/frozen, fossilized, chemical residue). Relevant planetary protection issues include quantifying any exchange (or lack thereof) of biological material between Earth and ocean worlds, and identifying any synergies between scientific and planetary protection priorities. limits of detection; context of any null result C. R&A Topics needed for Ocean Worlds.

Many of the objectives and investigations required to understand ocean worlds and any potential life cannot be accomplished using existing? spacecraft measurements, and/or we are not yet ready to make such measurements. More work is required here on Earth in order to undertake such spacecraft missions and/or to interpret the data. Some key areas of research fit well within NASA’s current R&A programs; these are important should continue, perhaps with an emphasis on Ocean Worlds. Some key areas of research do not fit in to current NASA R&A programs (e.g. oceanography/interfaces studies).

Needed because of long lead time on these projects. If we are arguing for a new program, we need to demonstrate here WHY.

Basic research is needed for:

GOAL I areas. Objective I.C. How do materials behave under conditions relevant to any particular target body?

A. How do materials behave under conditions relevant to any particular target body?

C.5. What are the phase relations of material composing ocean worlds at relevant pressures and temperatures?

C.6. What is the composition and chemical behavior of materials composing ocean worlds?

C.7. What are the rheologic mechanisms by which material deforms under conditions relevant to ocean worlds?

C.8. How does energy attenuation/dissipation occur under conditions relevant to ocean worlds? C.5 What are the thermophysical properties of material under conditions relevant to ocean worlds?

GOAL II areas

Interfaces with Earth oceanography people & collaborative studies (NOAA, NSF; the NASA Earth Science Oceans program – mostly remote sensing)

Studies related to ice shell – radar reflectivity etc in prep for Europa mission Studies related to seafloor/silicate interactions & high-pressure stuff (Habitable Worlds); thermodynamics/chemistry of water-rock interactions (HW/SSW); Reactions between water-rock (NSF & NASA)

Equipment: Simulation chambers- for creating relevant environments

GOAL III areas “Fundamentals of Habitability” questions:

1. Solvent 1.1. What solvents are suitable for life? 1.1.1. Are there limits (water activity, ionic strength, other) to the composition of aqueous environments that can support life? If so, what are they? 1.1.2. Are solvents other than water capable of supporting life? If so, what are they? 1.2. How do the temporal and spatial extents of liquid environments factor into habitability? How short/transient/small is insufficient? Can a liquid environment be too extensive in space and time to be habitable?

2. Energy 2.1. Can life take advantage of energy sources other than those known to support life on Earth (redox or visible-near IR light)? 2.2. When and how does energy availability constrain the type, diversity, and/or abundance of life? 2.3. Are there organismal, local, or planetary scale limits on how much energy is enough to support life? If so, what are they, and how do they depend on the physicochemical environment? 2.4. Are there upper limits to the amount of energy that can be constructively harnessed by life (vs. that same energy becoming destructive)?

3. Elemental and molecular raw materials 3.1. Can biochemistry be based on elements other than those utilized by life on Earth? 3.2. Are there lower limits on the abundance or constraints on the chemical form of elemental and molecular resources required for habitability?

4. Physicochemical environment 4.1. Can life transcend the physicochemical limits exhibited by life on Earth? Of particular importance for ocean worlds, what are the high pressure limits for life? 4.2. How do the physicochemical limits for life change under conditions of compound “extremes” or energy limitation?

5. Origin of Life 5.1. Under what physicochemical conditions, with what energy sources, and with what abundance and chemical form of elemental and molecular raw materials can life emerge? How does the probability of an origin of life vary within the range of permissive conditions? 5.2. How does the time scale required for life to emerge vary as a function of conditions within the permissive set? 5.3. Are there processes essential to the origin of life that definitively do not occur in sub-ice oceans?

Field work/analogue studies : oceans/ice shells/habitability (mostly currently through NSF); some through PSTAR, Exobio? Habitable Worlds

Note that some of these are extremely challenging questions that we cannot reasonably expect to answer in any short-term horizon, for example the question relating to “the probability” of the origin of life. The purpose in articulating them here is to create a complete list of open questions whose answers could strongly influence the way we choose to (prioritize a) search for life, whether or not we think those answers will be forthcoming in a meaningful time frame.

Note that many Goal III questions are covered in the Astrobiology program – need to make sure Ocean Worlds get enough emphasis.

Goal IV areas. What are the biomarkers & how do we detect them; Exobio, NAI; also basic biology/metabolism/conditions/survival (NSF?) Survival of amino acids/biomolecules under various conditions (e.g radiation) (Biosignatures working group of NAI)

Validation of biology technologies for application to ocean worlds

Disconnect between chemical detection & life detection – how to get past this; think of ecosystems & whole cells

Studies of non-water-based life (Habitable Worlds, Exobio)

Biopreservation potential for liquid h2o interacting with organics on ow surfaces (e.g. what to look for; cryobiology) Planetary protection. E.g. Would an Earth-based biomolecule definitely end up in an extraterrestrial ocean?

Note that some Goal IV topics are being addressed in COLDTech; this should continue.

D. Appendices. D1. ROW membership

List all names

D2. Enceladus

I. Is a subsurface ocean is expected based on the energy situation at Enceladus?

Enceladus’ present-day eccentricity, and thus its dissipation, is maintained by a 2:1 eccentricity-type mean motion resonance with Dione (Meyer & Wisdom 2008). In the particular case that Enceladus’s eccentricity remains constant, the heat production within Enceladus is inversely proportional to Saturn’s Q and is independent of Enceladus’s k2/Q (Meyer & Wisdom 2007). In this case, the eccentricity adjusts itself to produce the equilibrium heating rate, depending on the value of Enceladus’s k2/Q. The time-averaged Q of Saturn is generally assumed to exceed 18,000—if it were smaller, Mimas would have been inside Saturn’s Roche limit later than 4.5 Ga B.P. (Meyer & Wisdom 2007). If Q of Saturn > 18,000, then the upper bound on the time-averaged heat production within Enceladus is 1.1 GW (Meyer & Wisdom 2007). This is an order of magnitude smaller than the currently observed heat flow. (From Spencer and Nimmo, 2013).

The tidal heating of Enceladus could be higher if the current Q of Saturn is much smaller than 18,000, as suggested by a recent interpretation of astrometric data (Lainey et al. 2012). However, this scenario is somewhat unlikely as the astrometric analysis also shows rapid inward motion of Mimas, which is difficult to explain. A smaller Saturn Q is also difficult to reconcile with the current locations of the satellites unless the mid-sized satellites are significantly younger than the solar system (cf. Charnoz et al. 2011). A conductive or convective shell heated uniformly by 1.1 GW cannot sustain a global ocean over geological timescales (Roberts & Nimmo 2008). However, 1.1 GW of heating focused around the south polar region can sustain a regional sea indefinitely (Behounkova et al. 2012, Tobie et al. 2008). Thus, Enceladus could well have maintained a regional ocean for as long as the current e-resonance with Dione has existed, even though the current ocean seems to be global (Thomas et al. 2016). Similar regional oceans in other locations at earlier times may explain the older activity that has affected the leading and trailing hemispheres of Enceladus.

What evidence exists for a subsurface ocean, or any other type of subsurface liquid areas (i.e. maybe not a global ocean), or any surface liquids?

In July 2005, instruments on board the Cassini spacecraft detected endogenic heat and plume activity at Enceladus’ south pole associated with four prominent fractures (sulci) commonly referred to as tiger stripes (Spencer et al., 2006; Hansen et al., 2006; Porco et al., 2006; Spahn, 2006). Triangulation of Imaging Science Subsystem (ISS) images shows that these sulci are also the sources of the observed plumes (Spitale and Porco, 2007). From this point onwards a sub-surface ocean has been speculated to exist, since a subsurface, regional “sea” best explains the active south polar terrain’s topographic low (Collins & Goodman 2007).

The presence of subsurface liquid on Enceladus was strengthened by analysis of the chemical composition of the plume. It was found to be primarily composed of water vapor, with small amounts of carbon dioxide, methane, ammonia, methanol and many heavier hydrocarbons and organic molecules (see Table 1, Hansen et al. 2011; Waite et al. 2009, 2011). Ammonia and methanol are efficient antifreezes (e.g. Hogenboom et al.1997) and both have been detected in the plumes at the 1% and 0.02% level respectively (Waite et al. 2009). While methanol’s abundances are lower it has a lower vapor pressure than ammonia, meaning its concentration in the sub-surface liquid plume source may be higher, which would make the liquid more viscous than pure water. Higher concentrations of such antifreezes in the ocean could delay (but perhaps not prevent) complete freezing of a global ocean. Conversely, if tidal heating and liquid water are both confined to the south polar region, then steady-state situations permitting a regional sea and matching the observed heat output can be constructed (e.g., Tobie et al. 2008). The detection of Na in the erupted ice particles (Postberg et al. 2011) is most readily explained by the presence of liquid water that at some point was in contact with silicates (Zolotov 2007).

More recently analysis of gravity measurements support the existence of a sea 10 kilometres thick located beneath an ice crust 30 to 40 kilometres thick (Iess et al., 2015). However, careful analysis of Enceladus’ orbital librations suggests that the extent of subsurface liquid is global, even though the ocean is likely thicker below the South Pole (Thomas et al., 2016). II. What do we know about Enceladus’ ocean or liquids?

All of what we know about Enceladus’ sub-surface liquid is inferred from the composition of material ejected through its geysering activity. The composition of the plume, as sampled in situ by Cassini’s INMS instrument is given in Table 1. Results from Cassini’s CDA instrument showed that grains in Enceladus’ plumes primarily had one of three compositional types. One of the types, which occurs in ~6% of E-ring detections is particularly rich in sodium and potassium salts (0.5–2% by mass) (Postberg et al., 2011). It was observed that the composition of these salt- rich grains closely resembles the predicted composition of an Enceladean ocean that has been in prolonged contact with the rocky core of the moon (Zolotov, 2007). Due to this compelling compositional match, the grains are thought to form from frozen droplets present as spray over a liquid reservoir close to the surface (Postberg et al., 2011).

From the composition and size ranges of these particles (2 to 8 nm radius) it is inferred that these particles are produced by ongoing high-temperature hydrothermal reactions (>90 °C). Such activity is likely associated with global-scale geothermal activity, which quickly transports hydrothermal products from the ocean floor at a depth of at least 40 kilometres up to the plume of Enceladus (Figure 1, Hsu et al., 2015). Laboratory experiments support the hypothesis of high- temperature reactions occurring in Enceladus’ ocean to form silica nanoparticles and potentially explain the concentrations of ammonia and carbon dioxide in the plumes (Sekine et al., 2015).

Knowledge of the pH improves our understanding of geochemical processes in Enceladus’ ocean. Chemical models of ocean water on Enceladus, constructed using the plume composition, suggests that Enceladus’ ocean is a Na–Cl–CO3 solution with an alkaline pH of ∼ 11–12. This makes Enceladus’ ocean more analogous to soda lakes on the Earth than terrestrial seawater. The high pH is interpreted to be a key consequence of serpentinization of chondritic rock, as predicted by prior geochemical reaction path models; although degassing of CO2 from the ocean may also play a role depending on the efficiency of mixing processes in the ocean.

Serpentinization leads to the generation of H2, a geochemical fuel that can support both abiotic and biological synthesis of organic molecules such as those that have been detected in Enceladus’ plume (Glein et al., 2015).

Figure 1: A schematic of Enceladus’ interior (from Hsu et al., 2015).

III. What do we know about the availability of energy to support life?

Life as we know it requires the co-location of electron donors (reductants) and acceptors (oxidants) (McKay et al. 2008). Supply of reducing species to the ocean is expected to result from chemical reactions at water-rock interfaces, driven by heat flow within the silicate-bearing core and at the seafloor (Vance et al. 2007). Present- day radiogenic heating within the core is expected to generate approximately 0.3 GW (Porco et al., 2006; Schubert et al., 2007). It is expected that most of Enceladus’ heat is produced by tidal heating either in its core (especially if its unconsolidated, Roberts 2015), or within the ice shell, especially if the shell is at least locally decoupled from the underlying silicates by liquid water, Spencer and Nimmo, 2013). The amount of tidal energy Enceladus receives is not well constrained, primarily because of uncertainties in the Q of Saturn (see above). Furthermore, how tidal heat is transported is still poorly understood. However, as more thoroughly discussed above, the composition of Enceladus’ plumes points to high-temperature hydrothermal activity at the seafloor (>50 ºC). At the surface, Cassini/CIRS measurements have constrained Enceladus’ heat flow to be between 5.8 and 15.8 GW (Spencer et al., 2006; Howett et al., 2011). However, how this heat flow changes with depth is unclear. It has been proposed that the observed thermal signature is consistent with conductive heating of the surface by fractures with temperatures above 200 K (Abramov & Spencer 2009). It is not known how much of this heat is converted to chemical energy at the surface.

The supply of oxidants is less constrained. Oxidant sources likely include radiolytic products of ice in Enceladus’ near-surface, irradiated in the energetic environment generated by Saturn’s magnetic field, or of water in Enceladus’ interior caused by the decay of long-lived radionuclides. However, the magnitude of these sources is poorly known, with near-surface radiolysis thought to be lower than at Europa by a factor of a few (Hand and Carlson 2011).

For internal discussion: Specifically are radiolysis products produced from alkaline liquids, and does energy get produced this way? Could this provide energy away from the SPT? Are the products from this processes useful for astrobiology? - Explore this internally (Tom McCord, Bob Carlson, Kevin Hand)

What is known about the availability of biogenic elements at and inside Enceladus?

Again, our knowledge of biogenic elements comes from direct sampling of Enceladus’ plumes. Of the elements essential to terrestrial life (carbon, hydrogen, nitrogen, oxygen, phosphorus, sulfur), all but phosphorus (which is also likely to be present given its cosmochemical abundance) have been found in the Enceladus plume (Waite et al., 2009). The presence of nitrogen is particularly exciting, as Enceladus is the only place where nitrogen has been discovered in direct contact with a liquid reservoir.

IV. Given our current understanding of Enceladus’ habitability what biomarkers have been observed?

Nitrogen-bearing and organic species up to 100 amu (the upper limit of the Cassini/INMS’ range) have been observed in Enceladus’ plumes (Waite et al., 2009). Cassini/CAPS has detected compounds up to 10000 amu, suggesting that INMS detected fragments of these heavier molecules. How might we go about looking for and detecting life at this target?

Compelling evidence for the presence of life in Enceladus’ ocean could be obtained by directly observing organisms with microscopy. However, if life is present, the odds of sampling intact organisms are lower than those of sampling chemical compounds of biological origin (making up cells and/or resulting from metabolism). Measurements to detect such chemical signatures include sequencing of polymers, looking for chiral excesses, and determining the isotopic composition of major elements essential to life (H, C, N, O) in specific organics. Molecular structures and distribution would indicate how far toward life chemistry has progressed, while measurements on inorganic compounds (elemental and stable isotopic compositions, chemical forms) would help quantify the habitability of the ocean.

For such measurements, mass spectrometry techniques are best suited. However, many other techniques could be utilized too, including X-ray diffraction, IR, and UV sensing. These would allow determination of the mineralogy, crystal structure, and composition of Enceladean material at scales ranging from global to microscopic. Remote sensing techniques, particularly any that sense surface pigments, could be used to link in situ to global studies.

It is crucial to be able to assess and quantify any modifications made to material from Enceladus upon sampling. Such modifications can be due to sampling velocity (at km/s impacts, organics are destroyed, but small biopolymers encased in ice survive: see Aksyonov and Williams 2001; Gu et al. 1999; Burchell et al. 2014), interaction with the instrument (e.g. the titanium antechamber of Cassini’s INMS; Waite et al. 2013), and contamination by compounds onboard the spacecraft (Summons et al. 2014).

The sampling platform could either be a Saturn or Enceladus orbiter, a spacecraft able to hover in the plumes, or a lander in the plume fallout zone. With these architectures, sampling velocities would decrease from km/s, to hundreds of m/s (plume ejection velocities), to zero velocity for surface material. A lander could allow subsurface sampling but would be unable to sample several locations, unless provided with roving or hopping capabilities. These architectures would be suitable for remote, in situ, or sample return investigations. Remote and in situ measurements are essential in order to provide context in the search for life. Sample return enables path-dependent analyses that can adapt to unexpected findings, analyses with instruments that cannot be flown, and sample archival for analysis with technology that postdates a mission by decades.

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D4. Titan

D5. Triton

D6. Ganymede & Callisto

D7. Ceres & small bodies Introduction: Ceres is the largest object in the asteroid belt (mean radius 469.7 ± 0.2 km) and is the only dwarf planet in the inner solar system (semi- major axis ~2.8 AU). More data are available about Ceres than most small bodies, because the Dawn spacecraft orbited Ceres from March 2015 (Russell et al., 2016) and observed Ceres with its Framing Camera (FC) (Sierks et al., 2011), Visible-Infrared Mapping Spectrometer (VIR) (De Sanctis et al., 2011), Gamma Ray and Neutron Detector (GRaND) (Prettyman et al., 2011) and radio-science experiment (Konopliv et al., 2011). Pre-Dawn thermal evolution models and telescopic observations were summarized by McCord et al. (2011). Our knowledge of Ceres is rapidly changing because the Dawn spacecraft is still acquiring data at Ceres. Thus, the following represents the state of knowledge as of July 2016.

A former global subsurface ocean on Ceres: Thermal evolution models of Ceres, using shape data derived from telescopic observations (e.g. Thomas et al., 2005; Parker et al., 2006; Drummond et al., 2014), indicate that following Ceres’ accretion, the dwarf planet could have differentiated into a silicate core and a water-rich outer layer. These models show that Ceres could have harbored a global subsurface ocean for several hundred million years after its formation (e.g. Castillo-Rogez and McCord, 2010; McCord and Sotin, 2005). Dawn data broadly agree with these models: the degree-2 gravity field and inferred moment of inertia indicate that Ceres is close to hydrostatic equilibrium and exhibits some central condensation, although the degree of differentiation is lower than expected (Park et al., 2016). However, Ceres’ surface retains ~16 km of relief and numerous impact craters, which suggests that the outer layer is stiffer than previously thought, and contains only ≤40% water ice by volume (Bland et al., 2016). It could be that the water- ice- rich top of the frozen global ocean was eroded by impact cratering and sublimation. In this case, only the base layer would remain, forming Ceres’ present surface and outer layer of (possibly stratified) salts, silicates and ≤40 vol.% water ice (Castillo- Rogez et al., 2016). Alternatively, a mud layer, with salts and perhaps clathrates, could be obtained if fine grains of rock (perhaps akin to carbonaceous chondrite matrix) remained suspended in the hydrosphere by convection in liquid and were not excluded from ice upon refreezing (Travis et al., 2015; Neveu and Desch, 2015). The freezing of a subsurface global ocean and associated tensile stresses could have fractured Ceres’ near- surface, providing a formation scenario for observed large- scale linear features (average length of ~200 km and maximum width of ~11km) called the Samhain Catenae (Scully et al., 2016). In addition, there is evidence for small- scale regions of water ice on Ceres’ surface and in the outer layer, consistent with the proposition that that outer layer contains ≤40 vol.% water ice. Sparse patches of water ice are associated with slumping regions that likely recently exposed ice from the near- subsurface (Combe et al., 2016). Ceres’ surface also displays lobate flows of similar morphology to water- ice- rich flows on Earth, Mars, and icy satellites. These flows are more numerous towards the cooler poles where water ice should be stable closer to the surface (Schmidt et al., 2016; Buczkowski et al., 2016; Hayne et al., 2015).

Possible composition of Ceres’ ocean: Ceres’ global near-infrared spectrum, as observed by VIR, is best fit by a mixture of ammonia-bearing phyllosilicates, magnesium-bearing phyllosilicates, carbonates, and a dark component (De Sanctis et al., 2015). The uniform composition of the phyllosilicates suggests aqueous alteration was global (Ammannito et al., 2016). This composition can be reproduced qualitatively with geochemical models of the interaction of liquid water, possibly bearing carbon and nitrogen-rich volatiles, with silicates and organics of chondritic composition, at temperatures below 50ºC and pressures encountered from Ceres’ center to near-surface (Castillo-Rogez et al., in prep). The fluids in equilibrium with these rock assemblages are rich in carbonates, ammonia/ammonium, sodium, and possibly chloride if concentrated upon freezing. These modeled, qualitative ocean compositions are a close match to those observed at Ceres’ bright spots (De Sanctis et al., 2016). Presumably the salinity of Ceres’ ocean would have increased as it froze, precluding a quantitative assessment of solute concentrations through time.

Evidence of geologically recent or current subsurface liquid: Ahuna Mons, a ~4 km high and ~21 km long mountain, is interpreted as a viscous cryo- volcanic dome formed by the ascent from depth of cryo-magma. This cryo- magma is interpreted to consist of low-eutectic salt melt, ice, and silicate solid, which may have incorporated liquid water as brines in the melt and/or as water ice in the solid component (Ruesch et al., 2016). Moreover, the bright spots observed in the 92-km-diameter Occator crater mainly consist of sodium carbonates, interpreted as the residue of crystallized brines extruded onto the floor of Occator from depth (De Sanctis et al., 2016). Current or recent liquid brines may result from transient, near-surface heating by impacts (De Sanctis et al., 2016). Alternatively, present near-surface conditions may allow for the current presence of small amounts of brines, although the formation of transient liquid brines from impact heating is more likely .

Summary: Ceres could have supported a global subsurface ocean during its early history. This ocean would have frozen within the first few hundred million years of Ceres’ evolution, unless convective mixing in a muddy interior kept heat from the decay of radionuclides well distributed. Impact cratering and sublimation could have eroded the water-ice-rich top of the frozen global ocean, leaving exposed today the bottom portion, which consists of a mixture of silicates, salts and water ice. Dawn observations suggest possible current reservoirs of liquid brines, or transient near- surface brines resulting from impact-induced heating. The composition of these brines, rich in carbonate, ammonium, and sodium, may mirror the dominant solutes in the former ocean.

References: Ammannito, E., et al. 2016. Distribution of Phyllosilicates on Ceres. Lunar and Planetary Science Conference XXXXVII. Abstract 3020. Bland, M. T., et al. 2016. Composition and structure of the shallow subsurface of Ceres as revealed by crater morphology. Nature Geoscience, in press. Buczkowski, D. L., et al. 2016. The Geomorphology of Ceres. Geological Society Of America Annual Meeting. Abstract 282816. Castillo-Rogez, J. C., et al. 2016. Loss of Ceres’ Icy Shell from Impacts: Assessment and Implications. Lunar and Planetary Science Conference XXXXVII. Abstract 3012. Castillo-Rogez, J. C., and McCord, T. C. 2010. Ceres’ evolution and present state constrained by shape data. Icarus 205, 443-459.

Combe, J.-P., et al. 2016. Detection of H2O-rich materials on Ceres by the Dawn mission. Lunar and Planetary Science Conference XXXXVII. Abstract 1820. De Sanctis, M. C., et al. 2011. The VIR spectrometer. Space Science Reviews 163, 329-369. De Sanctis, M. C., et al. 2015. Ammoniated phyllosilicates with a likely outer solar system origin on (1) Ceres. Nature 528, 241-243. De Sanctis, M. C., et al. 2016. Bright carbonate deposits as evidence of aqueous alteration on Ceres. Nature, in press. Drummond, J. D., et al. 2014. Dwarf planet Ceres: Ellipsoid dimensions and rotational pole from Keck and VLT adaptive optics images. Icarus 236, 28- 37. Hayne, P. O., and Aharonson, O. 2015. Thermal stability of ice on Ceres with Rough topography. Journal of Geophysical Research Planets 120, 1567-1584. Konopliv, A. S., et al. 2011. The Dawn gravity investigation at Vesta and Ceres. Space Science Reviews 163, 461-486. McCord, T. B., and Sotin, C. 2005. Ceres: evolution and current state. Journal of Geophysical Research 110, E05009. McCord, T. B., Castillo-Rogez, J., and Rivkin, A. Ceres: Its Origin, Evolution and Structure and Dawn’s Potential Contribution. Space Science Reviews 163, 63-76. Neveu, M., and Desch, S. J. 2015. Geochemistry, thermal evolution, and cryovolcanism on Ceres with a muddy ice mantle. Geophysical Research Letters 42, 10197-10206. Park, R. S., et al. 2016. Gravity science investigation of Ceres from Dawn. Lunar And Planetary Science Conference XXXXVII. Abstract 1781. Parker, J. W., et al. 2006. Ceres: High-resolution imaging with HST and the determination of physical properties. Advances in Space Research 38, 2039-2042. Prettyman, T. H., et al. 2011. Dawn’s Gamma Ray and Neutron Detector. Space Science Reviews 163, 371-459. Ruesch, O., et al. 2016. Ahuna Mons: A geologically-young extrusive dome on Ceres. Lunar and Planetary Science Conference XXXXVII. Abstract 2279. Russell, C. T., et al. 2016. Dawn Arrives at Ceres: Exploration of a Small Volatile-rich World. Lunar and Planetary Science Conference XXXXVII. Abstract 1275. Schmidt, B. E., et al. 2016. Ground ice on Ceres? Lunar and Planetary Science Conference XXXXVII. Abstract 2677. Scully, J. E. C., et al. 2016. Implications for the geologic evolution of Ceres, Derived from global geologic mapping of linear features. Lunar and Planetary Science Conference XXXXVII. Abstract 1618. Sierks, H., et al. 2011. The Dawn Framing Camera. Space Science Reviews 163, 263-327. Thomas, C., et al. 2005. Differentiation of the asteroid Ceres as revealed by its Shape. Nature 437, 224-226. Travis, B. J., et al. 2015. Unconsolidated Ceres Model has a Warm Convecting Rocky Core and a Convecting Mud Ocean. Lunar and Planetary Science Conference XXXXVI. Abstract 2360.

Midsized and Ceres-like Asteroids (Pallas, Hygiea) Besides Ceres the asteroid belt counts a number of large asteroids, the majority of which is believed to be volatile-rich based on their spectral classes and densities. Notable asteroids include Hygeia (XXkm), Pallas (XXkm), and 24 Themis (XXkm). The latter is the largest remnant of a ~390 km parent of a >2000 daughter family created upon catastrophic disruption (ref). 24 Themis’ surface exhibits the signature of ice (ref) although this interpretation has been debated (ref). In any case the Themis family hosts most of the main belt comets whose activity is attributed to water sublimation (Hsieh et al.) Thermal evolution models by Castillo-Rogez and Schmidt (2010) indicate that Themis could differentiate a rocky core and volatile-rich shell but not preserve a deep ocean until present. It is important to note that while the Themis daughters display B and C spectral types they lack evidence for hydrated material on their surfaces (Takir and Emery 2012). Hence it is unclear whether their parent experienced global melting and hydrothermal activity triggered by 26Al decay, as is generally expected for objects formed in the main belt.

Hygiea’s spectral properties are very similar to Ceres (Takir and Emery 2012), in particular with the presence of the 3.06 micron band associated with the presence of ammonium, likely in the form of ammoniated silicates (De Sanctis et al. 2015). Smaller than Ceres, Hygiea’s interior is less likely to have preserved an ocean until present, although the presence of ammonia could promote remnant liquid pockets depending on Hygiea’s cooling history.

Pallas is notable as the largest B-type body in the solar system (true?). The origin of the blue slope in the visible is unknown but it has been suggeted that it is evidence for outgassing (ref) or for the intimate mixture of ice and hydrated silicates (ref). Interior evolution models by Schmidt and Castillo- Rogez (2012) suggest that Pallas’ interior could differentiate a rocky core an volatile-rich shell if the body formed early enough to accrete 26Al. Schmidt and Castillo-Rogez (2012) theorized that the large basin observed in Pallas’ south hemisphere is the signature of a large impact that removed a significant fraction of a volatile-rich shell, exhuming oceanic material as suggested by UV observations (ref). It is unlikely for Pallas to host a deep ocean at present, although these models did not account for the potential presence of antifreeze volatiles. Indeed Pallas’ Tisserand parameter is akin to comets, which is strong suggestion that this large body has migrated from colder regions, in the outer solar system. However an alternative explanation for its high inclination is that it has been pumped up by solar resonances.

The asteroid belt also includes a large number of volatile-rich asteroids in the 100-200 km size range. Most of them belong to the spectral class Ch per the presence of a marked absorption feature at ~0.65 micron that has been attributed to the mineral cronstedtite and is thus evidence for past aqueous alteration (McAdam et al. 2015). These objects are however too small for the long-term preservation of a deep liquid layer as long-lived radioisotope decay is the only source of long-term heat within these objects (ref). That class also includes large C- asteroids that share similar properties with Ceres in the infrared (Takir and Emery 2012, Takir et al. 2015), and especially the 3.06 micron band attributed to ammonium (De Sanctis et al. 2015; Berg et al. 201X). If some ammonia remained present in solution then it could have acted as antifreeze. However the expected abundance of that volatile based on cosmochemical models (ref) is too small for the amount of remaining liquid to be of significant with respect to ROW. Nevertheless these bodies have been recognized as “abodes for life” (Abramov and Mojzsis) owing to their potential for hosting hydrothermal environments in their early history that could have been the sites of advanced organic chemistry.

Summary: The potential for long-term preservation of a deep ocean in 400+km asteroids directly scales from Ceres’ own potential. Hence future missions should aim at assessing the presence and quantifying the extent of liquid in Ceres’ interior as a reference for better understanding the class of wet asteroids as a whole. Asteroids in the 100-200 km range are believed to have hosted hydrothermal environments in the past but are unlikely to contain large bodies of water at present.

Trojan Asteroids (Michael, Micah) The ~10^6 Trojan asteroids located at Jupiter’s L4 and L5 Lagrangian points have sizes ranging up to 240 km. Several scenarios have been suggested for the emplacement of the Trojan asteroids (see Morbidelli et al. xx for a review, or Levinson et al for Trojans specifically). They may be remnants of an original population of planetesimals in Jupiter’s region or captured during a late migration event, such as described in the Nice model (Gomes et al. 2005). The spectral similarities of Trojan asteroids with P- and D-type asteroids and short-period comets suggests a genetic relationship between these populations. A relationship to Kuiper belt objects has also been recently suggested (Wong and Brown, in review at AJ). In summary Trojan asteroids are likely to be rich in ice and other volatiles, as well as organics, although these materials have proved elusive in telescopic observations. Objects bigger than ~100 km could have undergone partial melt and hosted a temporary ocean if they formed early enough to accrete short-lived radioisotopes. However, the low densities measured for a couple of large Trojan asteroids (~1g/cm3 for both 617 Patroclus and 624 Hektor) indicate that these objects have preserved a bulk porosity fraction of up to 30% (Marchis et al. 2015), or that they accreted much less rock than suggested by cosmochemical models (Castillo-Rogez et al. 2012). Hence the possibility for the formation of a deep ocean at some point in Trojan history either from short- or long-lived radioisotope decay heat is low. One cannot exclude that the two large binary asteroids, Patroclus and Hektor, were subject to tidal dissipation but orbit circularization should have occurred on a short timescales (ref). Long-term evolution is primarily via collisions; they may produce temporary and local hydrothermal systems, although telescope observations have not detected any evidence for hydrated material at the surface of these objects.

Summary: observational evidence to date and thermal modeling suggest that large Trojan asteroids cannot host deep oceans at present and the prospect for large bodies of water to be present in these objects in the past is low. D8. Pluto & KBOs I. The existence of an ocean on Pluto, and any other KBO, is based on theoretical modeling that computes the heat flow from the original accretion, impacts, and most importantly, radionuclides in its interior. The abundance of radioisotopes is assumed from its rock fraction. Models for Pluto show a liquid ocean for some cases. KBOs can have a substantial rock fraction (e.g., 2/3rds by mass for Pluto, and likely larger for Eris), but the types and physical distribution of radioisotopes (and other materials) will affect energy production over time.

Another potential source of energy is tidal heating in the case of KBOs with moons. The Pluto–Charon system is tidally evolved, but other systems may still have tidal heating.

There are a number of geologic features on Pluto that suggest the existence (or past existence) of a subsurface ocean, or at least a partially fluid mantle. The orientation of Sputnik Planitia on the anti-Charon point of Pluto may indicate reorientation and tidal wander of the shell facilitated by a subsurface ocean. Pluto has almost no equatorial flattening, indicating it remained warm enough at least through its early history to not retain or “freeze in” an equatorial bulge. Extensional tectonic deformation on both Pluto and Charon implies at least partial freezing of a subsurface ocean. Some of the young terrains on Pluto may have been resurfaced by icy volcanism, indicating at least somewhat mobile material was present and perhaps some melt at depth. In addition to water ice, Pluto has other volatile ices that might melt in the subsurface (N2, CH4, CO). Features like the dendritic-style patterns on Pluto’s surface may have been sculpted by some kind of liquid (potentially under glaciers). Liquid is not currently stable on Pluto’s surface, but earlier epochs of higher atmospheric pressure are possible. The Kuiper Belt is relatively unexplored, so both interior modeling and rudimentary astronomical observations can provide evidence for subsurface oceans. Besides measuring the density of an object (possible only for the case of a multi-body system), the measurement of size is key, as the geometric albedo of the surface can then be known. Thermal emission measurements are an indirect way of measuring the size of KBOs, and can be done for only the largest objects. Another, more direct technique is stellar occultations, which can provide accurate measurements of the occulting body’s diameter with several observation sites on Earth. Highly reflective surfaces, such as those found on Enceladus and Sputnik Planum, and to a lesser extent on Europa and Triton, are indications of active geologic processes that may originate in subsurface oceans

Another key measurement is rotational lightcurves, which pinpoint areas of activity. In retrospect, Sputnik Planitia first revealed itself in the large amplitude of Pluto’s lightcurve.

III. Radiogenic heating would be a source of heating on KBOs like Pluto and Eris because of their relatively large, rocky interiors. We do not know the state of chemical disequilibrium inside these bodies.

Theoretical work (Simonelli et al., 1989) showed that most of the carbon in the Kuiper Belt would be in the form of hydrocarbons. The reddish surface color and the photochemical processes in Pluto’s haze layer strongly suggests the presence of complex hydrocarbons, and the New Horizons data is currently being analyzed to give an inventory of these molecules.

Estimates of the total surface tectonic deformation plausibly attributable to the phase change of a subsurface ocean could place constraints on the total ocean volume within Pluto. Further, analysis of the tectonic structures themselves (amount of strain, spacing, etc.) would provide information on the local thickness of the ice shell, and thus some insight into the depth of the ocean. These findings could then be compared with thermal evolution models for Pluto.

IV. The biomarkers for KBOs would be similar to other icy worlds. In addition to remote sensing techniques, if young surface areas or plumes exist and appear to be bringing material up from below, a sampling mission may be able to conduct compositional tests. D9. Other Satellites Mimas Mimas is the midsized satellite closest to Saturn, and the smallest one, with a mean radius of 198.2 km. Its density of ~1150 kg/m3 corresponds to a rock volume fraction of about 8%, which means Mimas is most likely a heat starved body. Close to its primary, this moon could be subject to intensive tidal dissipation provided that its ice could become warm in the first place. Mimas’ currently high eccentricity of ~2% however conflicts with the idea that Mimas is dissipative at present and in the recent past (MAKE A NOTE ABOUT RESONANCES). Mimas’ large libration amplitude (Tajeddine et al., 2014) has been attributed to the decoupling of an ice shell from the deep interior via a liquid layer, or by the presence of a highly non-hydrostatic core. The latter explanation is more likely especially as there are increasing observations and models in favor of a formation of Saturn’s midsized moons in the rings (Charnoz et al. 2010; Cuk et al., 2016). In this context Mimas would be the youngest moon to have emerged from the rings. The heat budget of satellites formed this way is not well modeled (we don’t know when the satellite seeds formed in the rings), but their long-lived radioisotope makeup is likely to be significantly depressed with respect to a scenario where the satellites formed in Saturn’s subnebula. In conclusion it is unlikely that Mimas held an ocean in recent time or even at any time in the course of its history.

Tethys Tethys is Saturn’s 5th largest satellite with a mean radius of 533 km (Thomas et al., 2007). Its mean density is 973 kg/m3 (Thomas et al., 2007) suggesting it is mostly ice. Depending on the assumed densities of rock and ice, the rock fraction ranges from 6% by mass (Thomas et al., 2007) to 12.1% by mass (Multhaup and Spohn, 2007). If porosity is accounted for, the rock fraction may be as high as 14% (Castillo-Rogez in prep.). Nimmo et al. (2011) found that Tethys’ degree 2 shape is not consistent with hydrostatic equilibrium. Due to its small size and low rock fraction, it is unlikely that heating from accretion or radiogenic decay was significant. Tethys eccentricity is indistinguishable from zero, so it is not currently experiencing any tidal heating. Thermal evolution models (excluding tidal heating) predict that if convection ever occurred, it ceased early in solar system history (Multhaup and Spohn, 2007).

One of Tethys’ most prominent features is the large Ithaca Chasma, a rift that is about 1000 km long and 2-3 km deep. It is 100 km wide in the north and becomes two narrower branches in the south. Its flanks are raised by up to 6 km above its surroundings (Giese et al., 2007). Crater counting suggests it formed approximately 4 Gyr ago (Giese et al., 2007). Previous work has suggested that the formation of Ithaca Chasma may be related to the 400 km Odysseus impact basin or expansion due to the freezing of Tethys’ interior (Smith et al. 1982). More recent work involving crater counts suggests that Ithaca Chasma pre-dates Odysseus (Giese et al., 2007). However, the formation of Ithaca Chasma is still debated. Flexural modeling of Ithaca Chasma suggests that the surface heat flux was 18-30 mW/m2 when it formed (Giese et al., 2007). This may imply a past history of tidal heating and a higher eccentricity for Tethys, which could have occurred if Tethys passed through a 3:2 resonance with Dione (Chen and Nimmo, 2008). An ocean would not be required to produce the necessary heat flux, but if the eccentricity remained high for a prolonged period of time, for example excited by a resonance, then the ice layer would start to melt which would increase tidal heat production and further thin the ice shell (Chen and Nimmo, 2008). New estimates of the thermal conductivity profile which include the effects of porosity yield a heat flow an order of magnitude less than previously assumed and suggest that tidal dissipation is not required to explain the flexure observed at Ithaca Chasma. This implies that an ocean would not be required at any time in Tethys’ history (Castillo-Rogez, in prep.).

Dione Dione is one of Saturn’s larger mid-sized icy satellites, radius of~561 km, and orbits at a distance of ~377000 km with an eccentricity of 0.0022 (bulk density: ). The surface of Dione contains both heavily cratered and heavily tectonized regions. Young fault scarps known as the “wispy terrains” that first discovered by the Voyager missions (Smith et al., 1981; Plescia, 1983) and have been better imaged by the Cassini mission (e.g. Wagner et al., 2009). Unfortunately, the age of these features is not well constrained (Kirchoff and Schenk, 2015) but their crosscutting relationships suggest they are among the youngest features on the surface (Martin et al., 2015). The wispy terrain structures appear to be extensional in origin (Jaumann et al., 2009) but the moon also contains ridges that point to a compressional history (Collins et al., 2009). Some researchers have suggested that Dione is active and contributing to Saturn’s E-ring (e.g. Khurana et al, 2007) and plasma to Saturn’s magnetosphere (Burch et al, 2007). However, no plume activity has been observed at Dione (Buratti et al, 2011). The surface is composed mostly of water ice (Hussmann et al., 2015) and a thin coating of “dark material” that is similar to dark material found on Phoebe, Hyperion, and Iapetus (Clark et al., 2008). Dione’s interior is likely differentiated into a rocky core and icy lithosphere (Thomas, 2010) where the rocky core composes about 48% of the mass, higher than most other icy satellites of Saturn, the exception being Enceladus (Hussmann et al. 2015). The ice shell likely experienced convection in the past but it is uncertain if such convection continues today (Zhang and Nimmo, 2009) and thus, little is known about the possible transport of material between the surface and deeper layers. Similar to other icy satellites, the heating to melt any subsurface ice into a global ocean would have to come predominantly from tidal heating. Geologic evidence for a present-day subsurface liquid layer is inconclusive; however, the large ridges (Hammond et al., 2013) and other terrains suggest at least a time in the satellite’s past when it did have a subsurface liquid layer.

Rhea Rhea is the largest inner satellite of Saturn, i.e., within Titan’s orbit, with a mean radius of ~764 km. Its mean density of 1236 kg/m3 corresponds to a rock volume fraction of 0.2. Tidal heating is not a major heat source in Rhea (i.e., compared against heat transfer) owing to the moon’s relatively large semi-major axis. However, crater morphology reveals that Rhea benefited from enhanced heat flux with respect to long-lived radioisotope decay production (White et al. 2013). This could point to a resonance event. It is also possible that Rhea formed close to Saturn, possibly from ring material (Charnoz et al. 2010) and was subject to significant tidal dissipation in its early history. Gravity measurements obtained from multiple flybys by Cassini-Huygens revealed that Rhea is not in hydrostatic equilibrium, which precludes inferences on its interior structure. The departure from hydrostaticity is attributed to a large core oblateness that may be the signature of an origin in the rings (Charnoz et al. 2010), also found at Mimas (Tajeddine et al. 2014) and Enceladus (McKinnon 2013). However, in its current location Rhea’s heat budget is limited to long-lived radioisotope decay. Early models suggested that the latter heat source was sufficient to promote the formation of a deep ocean in Rhea (Hussmann et al. 2006). However, this feature would have most likely been short-lived and is not expected at present.

Iapetus Located at about 60 Saturn’s radii, Iapetus is the farthest regular satellite of Saturn. Its mean radius is about 735 km. It is slightly denser than water ice, at 1088 kg/m3. This corresponds to a volume fraction of rock of 0.15. The low silicate content combined with weak tidal forcing from Saturn imply a very small heat budget for this satellite. Iapetus’ very old surface uncovered by the Cassini-Huygens spacecraft is consistent with thermal models that predicted internal temperatures could not reach the water ice melting point. If Iapetus accreted ammonia, as suggested by cosmochemical models), then its interior could have been subject to partial melting and separation of the rock from the volatile phase (Castillo-Rogez et al. 2005; Robuchon et al.). That state could have lasted a few tens of My; otherwise Iapetus likely remained frozen for most of its history. This scenario is supported by the analysis of crater morphology (i.e., relaxation) showing that the main heat source throughout Iapetus’ history has been long-lived radioisotope decay (White et al. 2013). Magnetometer observations by Cassini-Huygens did catch a magnetic anomaly that could be explained by interaction between Iapetus’ intrinsic magnetic field and solar wind. The origin of the magnetic field (remnant or induced in a deep ocean) cannot be further constrained with available data (Leisner, et al., 2008). The very low likelihood that Iapetus could harbor a deep ocean today does not by itself justify a mission to explore this object. However, future missions to the Saturnian system should consider including Iapetus in their science planning, if possible, to obtain additional magnetic field measurements at this object. In conclusion it is unlikely that Iapetus held an ocean in recent time or even at any time in the course of its history.

Miranda Miranda has only been visited by one spacecraft: Voyager 2. Due to the local season during the Voyager encounter, only the southern hemisphere has been imaged. However, Voyager made a relatively close pass by Miranda, which allowed for relatively high resolution of this small, icy moon. Its density is about 1200 kg/m3, indicating a likely silicate component. Hussman et al. (2006) model an ice shell about 130 km thick around a 100 km radius core. The surface of Miranda is quite dramatic, showing a very disrupted and fractured nature. Most prominent are three banded regions, often called coronae, which have many similarities to banded regions on Ganymede, Enceladus, and Ariel (Pappalardo, 1994). Two main models of formation have been proposed for the coronae: the first involves massive disruption of Miranda, perhaps by an impact, and subsequent reassembly. Denser, silicate pieces that would have been moved closer to the surface through this process then sunk towards the core, causing the coronae to form above them (Janes and Melosh 1988). This mechanism would lead to contractional surface tectonic features, however, analysis of the surface features appears more consistent with extensional surface tectonics (Schenk, 1991; Pappalardo et al., 1997). Thus, the majority of works focus on formation of the coronae as a surface expression of extension above buoyant diapirs (Croft and Soderblom,1991; Greenberg et al., 1991, 2014; Pappalardo et al., 1997). More recently, Hammond and Barr (2014) considered convection-driven resurfacing to form the extensional features of the coronae. Their global modelling was able to reproduce the locations of the coronae. A significant challenge in explaining any extensional mechanism for formation of the coronae is supplying sufficient heat into the body. Hussman et al. (2006) performed interior modeling for many icy moons and found no plausible scenario for Miranda to have sufficient radiogenic heat to produce such diapirs. Additionally, they find the current eccentricity is too small to support sufficient tidal heating. However, Beddingfield, Burr, and Emery (2015) have recently re- evaluated the nature of faults on Miranda and found evidence consistent with a previous epoch of tidal heating, such as would be caused by orbital resonances with other moons (eg. Collins et al. 2009 and references within). Hussman et al. did not predict a subsurface ocean at Miranda, however, they also did not predict an ocean at Enceladus, which is believed to have a global subsurface ocean today. Miranda’s surface features are consistent with significant past tidal heating. Did this involve an ocean? If so, how long did it persist, and does it persist today?

Ariel Ariel is the fourth largest Uranian moon, radius 579 km, and the second farthest major satellite from Uranus, tidally locked in an orbit with a period of 2.52 days, semi-major axis 7.29 R U (or 191,000 km), and eccentricity 0.0017 (Tittemore and Wisdom, 1990). Ariel has a bulk density of 1660 kg/m3, similar to those of Titan and Oberon and consistent with a rock:ice mass ratio of ~0.6 (Smith et al., 1986; Hussmann et al., 2006). Densities and rock mass fractions of the majority of the large Uranian satellites are higher than for their Saturnian counterparts (e.g., Brown et al., 1991).

As with the other Uranian satellites, less than half of Ariel's surface was observed during the Voyager 2 encounter in 1986; only ~35% was imaged with sufficient resolution and illumination for analysis of geological features (Plescia et al., 1987). Imaging covers only the southern hemisphere, from the South Pole to 0°-30°S latitude, with the highest resolution at longitudes ranging from the sub-Uranian to the trailing side of the moon (~240°E - ~30°E). Voyager's closest approach to Ariel was 127,000 km and the highest resolution images acquired of the satellite were 1.3 km/line-pair (Croft and Soderblom, 1991). Cratering statistics show a paucity of large craters, the largest observed is 85 km in diameter (Plescia et al., 1987; Strom, 1987; Smith et al., 1986). This distribution suggests complete resurfacing or modification of at least the observed area after accretion by tectonism, cryovolcanism, viscous relaxation, mantling or burial by infalling material, or some combination thereof (Croft and Soderblom, 1991; McKinnon et al., 1991; Plescia et al., 1987; Smith et al., 1986).

Ariel's surface has undergone widespread faulting, with a system of fractures and graben 15-50 km wide cutting its surface and topographic relief of up to 4 km (Brown et al., 1991). The extensional tectonic regime is potentially the result of freezing of a liquid water ocean (Smith et al. 1986), although formation of a subsurface ocean would likely have required more than radiogenic heating alone (e.g., Plescia, et al., 1987). The tectonic activity appears to have been accompanied by extensive emplacement of viscous flows (Croft and Soderblom, 1991; Schenk, 1991; Jankowski and Squyres 1988; Stevenson and Lunine, 1986; Smith et al., 1986), suggestive of extrusive cryovolcanism, that emplaced a series of smooth units that embay and partially bury craters, surround nunataks of cratered plains materials, and fill graben floors with convex deposits bounded by troughs 1-2 km deep (Jankowski and Squyres, 1988; Ruzicka, 1988) along the graben walls and often exhibiting medial troughs.

Analysis of Hussmann et al. (2006) did not suggest an interior ocean; however, the extensional past tectonic regime and potential surface flows of viscous material could be indicative of the past presence of a liquid water ocean and/or subsurface liquid reservoirs, potentially the result of heating during passage through tidal resonances with other satellites (Tittemore, 1990; Tittemore and Wisdom, 1990). The heat flow required by the elastic thickness of the lithosphere at the time of the tectonic activity is higher than can be explained by the likely resonances (Peterson et al., 2015).

Ariel has the highest disk-averaged albedo, 0.4 (Smith et al. 1986), of the Uranian satellites and a maximum normal albedo of 0.55 is seen for younger craters and ejecta (Veverka et al., 1991). It has a mean surface T 60K (Castillo and Lunine, 2012). CO 2 ice present on the surface could be the result of outgassing, delivery during impacts, and/or radiolysis by magnetospheric charged particle bombardment of dark presumably carbonaceous material from the rings (Grundy et al., 2006).

The diverse Uranian satellite system provides key constraints on tidal evolution processes, the potential for creating and maintaining interior oceans, and the conditions that determine whether oceans form and whether they can be detected. Of the Uranian satellites, Ariel seems to offer the highest potential for geologically recent endogenic activity (Castillo-Rogez and Turtle, 2012).

Theoretical models do not predict an ocean at present; however, the surface is young and exhibits what appears to be cryovolcanic features and should therefore be investigated as a possible ocean world.

Umbriel Umbriel is Uranus’ 3rd mid-sized icy satellite (moving away from the planet). It radius of 584.7 km (Thomas, 1988) is similar to that of its inner neighbor Ariel, and smaller than the two outer satellites Titania and Oberon. Umbriel’s density is 1.39 g cm-3 (Jacobson et al. 1992), which is the second lowest in the system of five mid-sized satellites (Miranda’s is lower), suggesting an ice- rich and possibly porous body. Available data sets are limited. Umbriel was imaged by Voyager 2 in 1986, but only with a resolution of ~10 km per line pair. Due to the flyby geometry, only 20% of the surface is covered at sufficient resolution for geologic mapping (Plescia, 1987). Ground based telescopic data provides additional information, but is generally limited to its surface composition (e.g., Grundy et al. 2006).

Umbriel’s surface is dominated by water ice, although CO2 has also been detected (Grundy et al. 2006). It’s albedo, which is similar to that of Callisto, is lowest among the Uranian mid-sized satellites (Smith et al., 1986). In contrast to the other Uranian satellites, the surface is uniformly dark, with only a few high albedo regions associated with impact craters (e.g., Wunda and Vuver). The dark surface may result from exogenic deposition of dark dust (like Iapetus) (Smith et al., 1986) or endogenic resurfacing (Croft and Soderblom, 1991).

Geologically, Umbriel’s surface is dominated by impact craters (Smith et al., 1986, Plescia, 1987). Crater counts of the mapped region of Umbriel indicate that the surface is ancient (Plescia 1987), although counts by Strom (1987) suggest that the surface is not primordial. There is no unequivocal evidence for geologic activity on Umbriel. Croft and Soderblom (1991) report numerous lineaments and troughs, including a set of horst and graben. They also interpret a large depression identified in limb profiles as tectonic, although the feature is perhaps more likely a large basin (Thomas, 1988). All of the features identified by Croft and Soderblom (1991) have also been interpreted as arcuate crater rims seen at low resolution (Smith et al., 1986, Plescia, 1987). The extent of tectonic deformation is therefore unclear. Subtle albedo patterns observed by Helfenstein et al. (1989) have also been offered as evidence for endogenic (cryovolcanic) resurfacing.

Unlike the Jupiter and Saturn systems, the Uranian satellites are not currently in resonance so no dissipation of tidal energy is likely at present (see, e.g., Peale, 1999). However, the evolution of the Uranian system, which includes chaotic behavior, is complex, and Umbriel may have participated in several resonances in its past (Tittemore and Wisdom, 1988, 1990). None of these resonances are thought to significantly excite Umbriel’s eccentricity (Tittemore and Wisdom 1988, 1990), so an energy source to drive geologic activity appears to be lacking. It should be noted, however, that current studies have not included dissipation and interior evolution within the satellites themselves, and the coupled orbital-thermal evolution of the system should be reevaluated in the light of current understanding of satellite interiors [Peale 1999]. Furthermore, accretional energy alone may be sufficient to differentiate the satellite if ammonia is present (Squyres et al., 1988).

Evaluation as an Ocean World: Given the limited, low resolution imaging data available, Umbriel is a planetary Rorschach test: the observations and inferences made regarding it appear to reflect the propensities of the investigator, rather than Umbriel itself. To first order, Umbriel is a low-albedo, heavily cratered satellite expressing limited geologic activity, and no apparent source of tidal energy now or in its past. As such it may be an unlikely body to contain a subsurface ocean (Hussmann et al., 2006). However, some investigators have argued for extensive tectonic and cryovolcanic activity, which might imply the existence of an ocean at some point in Umbriel’s past or present. Furthermore, we should always remember the lesson of Mariner 4: don’t judge a body until you’ve seen its entire surface. We simply don’t have enough data to fairly assess the potential for oceans on Umbriel.

Titania and Oberon Titania and Oberon are the largest and most massive moons of Uranus. Based on their relatively high density of ~1600 - 1700 kg/m3 (Jacobson et al., 1992), Hussmann et al. (2006) estimated a rock fraction of ~50 - 60%. In addition to H2O ice, CO2 has been detected in near-infrared spectra of the surfaces of Titania, Oberon and Ariel (Grundy et al., 2006). The source of CO 2 on the surfaces of these moons is not currently known, and possible explanations include radiolytic or photolytic production of CO2 from carbonaceous surface material of endogenous or exogenous origin, as well as outgassing of volatiles from the interior. If Titania and Oberon are differentiated and their interiors contain an antifreeze agent, a liquid water ocean may presently exist at the core-mantle boundary. Hussmann et al. (2006) predicted a present-day liquid water ocean at the core-mantle boundary of Titania and Oberon if the primordial composition included even a small amount of initial NH3 – on the order of a percent. If this is the case, these oceans are predicted to underlie an ice shell with a thickness of > 100 km. However, as higher order components of the gravity fields of these moons are not known, an undifferentiated interior consisting of a homogenous rock-ice mixture cannot presently be ruled out for either moon (Hussmann et al., 2015). Based on crater size distributions, Plescia, (1987) estimated that the surface of Titania is younger than that of Oberon and Umbriel, indicating that some sort of global resurfacing process has taken place, likely sometime early in Titania’s history. As a side note here, it’s worth pointing out that Plescia, (1987) also found that Ariel and regions of Miranda to have the youngest surfaces of the large Uranian satellites. The surface of Titania is marked by a number of tectonic features in the form of extensive faulting and parallel scarps, forming large canyon-like features, the largest of which (Messina Chasma) extends for 1500 km (Smith et al., 1986). The major, rift canyons include both NW-S and NE-SW components and appear to post-date almost all other structures on the surface (Croft, 1989). Several regions on the surface (notably adjacent to several faults and near the large Ursula impact crater) contain terrain that appears to be smooth. This may possibly indicate a surface expression of cryovolcanic activity (more recent than the global resurfacing), or alternatively, may be due to blanketing by impact ejecta (Plescia, 1987).

Oberon is the most heavily cratered of the major Uranian moons and the surface crater density is near saturation, indicating an ancient surface (Plescia, 1987). Like Titania, the surface of Oberon exhibits major tectonic features in the form of an extensive network of rift canyons, thought to be caused by a global expansion of ~0.5 % (Croft, 1989). The tectonic activity creating the rift canyons appears to have occurred in two phases, as both degraded and fresh canyon morphologies are observed (Croft, 1989). Patches of very dark material has been observed in the floor of several of the large impact craters on Oberon, and it has been proposed that this may be some expression of cryovolcanism (Smith et al., 1986).

Theoretical models and the tectonically deformed surfaces of the satellites suggest that oceans could be present on Titania and Oberon. There is therefore, good reason to think that these targets may be potential ocean worlds and they should be explored further.

References:

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