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 Kepler is a 31 km diameter Copernican age complex impact crater located on the nearside maria of the Moon. We used Lunar Reconnaissance Orbiter imagery and topographic data in combination with Kaguya terrain camera and other image data sets to construct a new geomorphologic sketch map of the Kepler crater, with a focus on impact melt-rich lithologies. Most of the interior melt rocks are preserved in smooth and hummocky floor materials. Smaller volumes of impact melt were deposited in rim veneer, interior and exterior ponds, and lobe-like overlapping flows on the upper crater wall. Based on shadow lengths, typical flows of melt-rich material on crater walls and the western rim flank are ∼1–5 m thick, and have yield strengths of ∼1–10 kPa. The melt rock distribution is notably asymmetric, with interior and exterior melt-rich deposits concentrated north and west of the crater center. This melt distribution and the similarly asymmetric ray distribution imply a slightly less than 45° impact trajectory from the southeast. The exposed wall of Kepler displays distinct layering, with individual layers having typical thicknesses of ∼3–5 m. These are interpreted as flows of Procellarum mare basalts in the impact target. From the point of view of exploration, numerous fractures and pits in the melt-rich floor materials not only enable detailed studies of melt-related processes of impact crater formation, but also provide potential shelters for longer duration manned lunar missions.
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 New lunar data, spearheaded by those provided by the Lunar Reconnaissance Orbiter (LRO) spacecraft, present unprecedented possibilities for remote sensing studies of the Moon. High-resolution imagery and topographic data are particularly useful for studying the impact cratering processes. This type of remote sensing assessment is a mandatory precursor for the exploration of Moon that the National Research Council (of the National Academies of the USA, henceforth referred to as the NRC) outlined to resolve the highest science priorities for the Moon, the most important concept being the bombardment history (concept 1 [National Research Council (NRC), 2007]). In addition, the report stressed that the Moon is an accessible laboratory for studying impact cratering processes (concept 6 [NRC, 2007]) that will be applicable to all planetary surfaces. As natural probes of the crust, impact craters are also vital for unraveling the diversity of lunar crustal rocks (concept 3 [NRC, 2007]).
 Impact melt-rich rocks are crucial to establishing the flux and the types of projectiles that hit the Moon [e.g., Kring, 2009]. Their composition and distribution also provide key insights of the impact cratering process itself [e.g., Cintala and Grieve, 1998]. Thus, fresh complex craters with well-preserved melt-rich rocks are good targets for photogeologic analysis and later sampling by robotic and manned missions. Kepler crater (Figures 1–234) on the lunar nearside has been recommended as one of the sites for such studies [Wilhelms, 1993; Kring, 2009]. In this report, we characterize the distribution and diversity of impact melt-rich lithologies, place some constraints on their rheological properties, and provide general guidelines for robotic or manned scientific exploration of the Kepler area.
2. Geologic Background of the Kepler Area
 Photogeologically, Kepler is a prominent lunar complex impact crater with a final rim-to-rim diameter (D) of 31 km and an average rim-to-floor depth (d) of ∼2.7 km, with the highest parts of the northeastern rim crest rising ∼3.2 km above the crater floor (Figure 2). The crater is located at 8.1°N, 322°E on the lunar nearside, between Mare Insularum and Oceanus Procellarum (Figure 1). Very little photogeologic research has been carried out since the geologic mapping of the Kepler quadrant [Hackman, 1962]. Photogeologically Kepler deposits have only been divided into crater floor, slope, and rim materials [Hackman, 1962].
 The most extensively used data were the LRO's Narrow Angle Camera (NAC, resolution ∼0.5–1.2 m/pixel) and Wide-angle Camera (WAC, ∼60 m/px) monochrome and color images [Robinson et al., 2010] taken through 15 March 2011, and Kaguya Terrain Camera (TC, ∼7.4 m/px) evening and morning mosaics (Figure 3 [Haruyama et al., 2008]). Other data include Kaguya Multiband Imager (MI, ∼20 m/px; Haruyama et al., 2008]), Lunar Orbiter (LO, ∼60 m/px), Clementine UVVIS (Ultra-Violet – VISual, ∼100 m/px, 250 m/px for ray mapping), and SMART-1 AMIE (Small Missions for Advanced Research in Technology, Advanced Moon micro-Imager Experiment, ∼150 m/px) images, as well as oblique images from the Lunar Orbiter III and Apollo 12 missions. For ray mapping, Clementine iron and titanium maps were also utilized [Lucey et al., 2000]. Earth-based Consolidated Lunar Atlas (CLA) imagery [Kuiper et al., 1967] and Virtual Moon Atlas 5.1 software (C. Legrand and P. Chevalley, Virtual Moon Atlas, version Pro 5.1, 2010, available at http://www.ap-i.net/avl/en/download) were used for regional context.
 Elevations and slopes are derived from LRO Lunar Orbiter Laser Altimeter (LOLA) point data, with a nominal vertical accuracy of 10 cm and a spot size of ∼5 m [Smith et al., 2010]. LOLA makes five spot measurements in an X-pattern, with each orbit resulting in five near-equidistant parallel profiles spaced at ∼12 m [Smith et al., 2010]. In addition, the 512 px/degree (∼59 m/px at Kepler's latitude) LOLA digital elevation model (DEM) was used.
 LRO WAC and NAC images were radiometrically calibrated, map projected, and mosaicked with United States Geological Survey's Integrated Software for Imagers and Spectrometers (ISIS, version 3.2.1). Further processing and analysis was carried out with Adobe Photoshop CS5 Extended, ImageJ 1.44I, and ESRI ArcGIS 10 software. Ames Stereo Pipeline (version 1.0.2 [Broxton and Edwards, 2008; Moratto et al., 2010]) was used to analyze the morphology of the melt-rich lithologies on the eastern crater floor and rim in three dimensions, and to produce the perspective views (DEMs). A perspective view of the northern part of Kepler based on AMIE stereo imagery was also found useful [d'Angelo and Wöhler, 2008].
 Subroutine Adiabat_1ph (versions 2.0.1 and 3.1) of the MELTS algorithm [Ghiorso and Sack, 1995; Smith and Asimow, 2005] was used for the density and viscosity calculations. A temperature range of 1700–1125°C and an average lunar noritic composition, with comparisons to average lunar anorthosite and average Apollo 12 basalt compositions (Table 1) [Haskin and Warren, 1991] were assumed. MELTS calculations were carried out at 1 bar, which with respect to the rheologic properties of the melt can be considered a reasonable approximation for the near vacuum of the Moon. The viscosity calculation method incorporated in MELTS is that by Shaw , which is suitable for average norite, basalt, and anorthosite compositions.
 In this study, different geologic units (Figure 4) were mapped based on their morphology and location with respect to the crater. Their definitions, characteristics, and distribution are discussed below. The melt-rich units are smooth and hummocky floor materials, terrace ponds, wall lobes, rim veneer, and exterior ponds. This classification mostly follows the scheme by Howard and Wilshire . Melt-poor units we defined are the central uplift material, terrace zone, massive slumps, exposed wall, continuous and discontinuous ejecta deposits, and rays.
4.1. Distribution and Characteristics of Impact Melt-Rich Lithologies
 As expected, the largest continuous concentrations of melt-rich material can be found on the floor of Kepler. Photogeologically, the melt-rich material forms two distinct units that grade into each other. The smooth floor material unit (Figure 5) forms a level, smooth, fractured, and fairly dark unit mainly in the northern and western parts of the crater floor, covering altogether about 35 km2 (Figure 4). Where stereo imagery is available (mainly the eastern part of the crater), the fractures in most cases are associated with topographic gradients (Figure 5). This relationship applies to all generally more or less level melt-rich units.
 The hummocky floor material unit (Figures 5–7) is more extensive than the smooth floor unit, covering ∼78 km2 (Figure 4). It is reminiscent of the smooth floor unit, but as the name suggests, the surface of the unit is not level. Hummocks of higher albedo material are typically tens to a few hundred meters in diameter, and the material between these individual hummocks is similar to the smooth floor material. However, stereo imagery reveals the hummocky floor unit has also longer wavelength topographic irregularity, quite distinct from the smooth floor unit. In many places, best seen in the eastern part of the crater, there is a distinct topographic bench (Figure 7) near the contact between the hummocky floor and the lowermost part of the terrace zone, with heavy fracturing of the melt-rich floor material where the topographic gradient is highest.
 In addition to fractures, collapse pits some tens of meters in diameter are seen in the hummocky floor unit. Prominent ledges in some of the fractures and pits offer possibilities for estimating the thickness of the visible melt-rich material. An apparently deep pit (Figure 5) near the southern edge of the floor shows signs of granular flows on its eastern wall, implying the maximum slope is at the angle of repose, or ∼31° [e.g., Quaide and Oberbeck, 1968]. Assuming the wall slope is 31°, the shadow length gives a maximum depth of ∼30 m for the pit. Shadow measurements of ledges of melt-like material in other pits and fractures, where slopes give the impression of being shallower than the angle of repose, imply that the thickness of the directly observable melt-rich material in both hummocky and smooth floor units is commonly on the order of 5–10 m.
 Assuming the smooth floor unit has 80 volume-% melt and a thickness of 10 m, and the hummocky floor material is 5 m thick and composed of 50 vol.-% melt, the total combined observed melt rock volume is only ∼0.5 km3. However, as the observed ledges form only the uppermost crust of the melt-rich units, and some melt also must have flowed down to the fractures in the crater floor, the amount of impact melt derived from this rough approximation is a minimum estimate.
Terrace ponds are isolated and relatively small patches of melt-rich material on top of the terrace zone and massive slump material. Definite flow features related to terrace ponds are rarely observed (although Kaguya TC images show flows from a pond on the southeastern terrace zone to the crater floor). However, there are characteristics indirectly indicating the flow of material, such as narrow radial depressions (Figure 8) interpreted as erosional channels caused by impact melt-rich material flowing downward and then pooling in local depressions (similar to, e.g., Aristarchus [Mustard et al., 2011] or Thales [Bray et al., 2010]). In the geomorphologic sketch map (Figure 4) only the larger fractured pools are included, but probably many of the smaller patches of dark material are similar ponds of melt-rich rock.
Wall lobes are lobate, mutually overlapping sheets of smooth, usually fairly dark material on the upper crater wall (Figure 9). In the current (until 15 March 2011) NAC coverage, they are best seen on the western wall, and are much harder to distinguish on other parts of the crater wall. Shadow length measurements imply that their thickness varies from ∼0.8 m to perhaps slightly more than 5 m. Due to their mutually overlapping layered nature, the length of the individual lobes is hard to define, but seems to be typically a couple of hundred meters. Instead of individual lobes, the wall lobes should be considered as a zone of the upper crater wall, where this characteristic flow morphology of melt-rich material occurs. It is noteworthy that the “stubby” and sheet-like nature of the wall lobes clearly distinguishes them from long, narrow, digitate clastic debris flows that commonly occur on the crater wall, partially covering some of the wall lobes.
Rim veneer covers most of the circumference of Kepler (Figure 4). It is smooth but typically highly fractured, and gives the impression of forming a thin coating on the rim (Figure 10). The veneer often seems to be the source for the wall lobes, and it grades into the exterior ponds, although distinct flow channels (Figure 10) are rarely seen. In the geomorphologic sketch map (Figure 4) ponded material that appears to be continuous with the veneer has been included in the rim veneer unit. The leveed flow in Figure 10 exhibits flow fronts (black-in-red arrows) on the side of the main flow, where the flow was diverted toward a small local depression. The flow and the associated pond are also fractured perpendicular to the flow direction.
 The larger areas mapped as rim veneer on the southwest, southeast, and north sides of the crater (dotted in Figure 4) are difficult to map, as the surface is covered with many small patches of slightly darker, apparently ponded but poorly delineated material. There are some indications of flow, but the characteristic fracturing or ledges of the rim veneer are generally lacking. We interpret these areas as being covered by a very thin veneer of impact melt-rich material. Thus, although the mapped areas are large (and some data imply that, e.g., in the northeast they should perhaps be even larger), the volume of the melt rock in these deposits is very small.
Exterior ponds (Figure 10) with their generally level, slightly fractured dark surfaces are similar in appearance to the terrace ponds, and many of them are closely related to the rim veneer. The most striking feature of the exterior ponds is their non-symmetric distribution around Kepler (Figure 4). There are large ponds north of the crater and numerous small ponds close to the northwest rim crest. Some small and poorly defined ponds are present south and particularly southeast of the crater (dotted in Figure 4), but in general they are rare or lacking on the eastern and southern to southwestern sides of the crater. As the larger ponds can readily be identified in Kaguya TC and MI or LRO WAC imagery, the asymmetric pond distribution is not due to uneven high-resolution NAC coverage or unfavorable illumination geometry. Instead, it appears to be a true phenomenon related to the initial distribution of the melt-rich material.
4.2. Melt-Poor Interior Deposits
 Geologic units with a low apparent melt content were also mapped. These include the central uplift, terrace zone, massive slumps, and the exposed wall. All of these, and particularly the central uplift and massive slump units, include small deposits of melt-rich material that, mostly due to their limited size and very poorly discernible boundaries, are not depicted in Figure 4.
 The central uplift material in Kepler does not consist of a single peak, but instead forms a cluster of rounded hills with typical slopes of ∼15°, although steeper slopes occur locally. The most prominent part of the central peak cluster is slightly north from the crater center, and the majority of the central uplift forms a crudely crescent-shaped structure open to the southeast (Figure 4). The highest of the hills crossed by LOLA profiles rise ∼550 m above the surrounding crater floor. This closely matches the predictions from the scaling relationships by Wood  and Hale and Grieve , but is ∼65 m higher than Wood and Andersson's  and ∼150 m lower than Pike's  scaling laws predict. The surface of central uplift material has a higher albedo than the surrounding materials and typically has more boulders on the surface than the surrounding crater floor (Figure 6). No layering or other structures can be seen in the central uplift.
 Ledges of material with melt-rich characteristics can be seen onlapping the lower parts of the uplift, and melt lithologies are also present in some of the troughs between the hills, making the transition to floor material sometimes difficult to discern. The easternmost separated and fairly low (∼85 m above the surrounding floor) parts of the central uplift (Figures 3, 4 and 6) may instead be extremely large bedrock clasts embedded in the hummocky floor material.
 The collapse of the wall of Kepler resulted in two types of materials, classified as terrace zone material of moderate albedo and massive slump material. However, in most cases the difference between the two is not clear-cut. In general, the surface expression of the terrace zone material is smoother than that of the fairly rugged massive slumps. Importantly, the terrace zone material hosts linear troughs that run roughly parallel to the crater rim, interpreted to be normal faults responsible for the terrace- or stair-like appearance of some parts of the crater wall (Figures 3 and 8), typical of a crater modification zone. Terrace ponds of melt-rich material can be found in many of the troughs (Figure 8). The blocks defining those troughs appear to have moved down fault planes as coherent blocks, rather than as rock avalanches, which is typical for the massive slump material. A sketch of the main faults can be seen in Figure 3. In addition to the faults and the terrace ponds, other prominent features of the terrace zone are narrow radial grooves running between terraces, or from a terrace to the crater floor (Figure 8).
 The faults are mainly present on the western half of the crater wall (Figure 3). Based on LOLA point data, approximate minimum vertical displacements (throws) of these terrace-forming normal faults vary from ∼1000–1600 m for the fault closest to the final crater rim, whereas the inner faults have vertical displacements of only ∼40–60 m. The southeastern wall has few faults, and it seems that massive slumping covers at least some probable faults on the eastern and northwest terrace zones. The extensive image coverage shows that the observed east/west difference in the fault distribution is not influenced by lighting geometry, but it is a true distinction in the structure of the crater wall.
 The formation of massive slumps is a major contributor to the scalloped, partially polygonal plan view of Kepler (see Öhman  for a discussion). Terrace ponds and other smaller and poorly defined (hence, not depicted in Figure 4) occurrences of melt-rich material are in many places present on top of the massive slump material. The massive slump material is best exposed on the north-northeast lower wall of the crater (Figure 3), whereas the slump deposits on the southern wall are minor, and their nature can properly be seen only in images having a low solar angle. The massive slump material has an irregular, sometimes grooved surface, and particularly on the lower north-northeast wall it can be seen to form a large bulge, extending to the crater floor. On the south-southwest wall of the crater, the upper wall hosts a small bulge of similarly textured massive slump material, below which is the terrace zone. However, below the terrace zone there is another slump-like feature, which protrudes far onto the crater floor and central uplift region. Despite some morphological differences, this material has been classified as massive slump material in Figure 4.
 Relatively high-albedo material comprising the exposed wall unit occurs above the terrace zone or the massive slump material. The best exposures of this material are on the northern and particularly the northwestern wall (Figure 11). Although spatially fairly limited, a characteristic feature of this unit is the presence of distinct horizontal layering (Figure 11; also observed by Zanetti et al. [2011a]), traces of which can be seen throughout the upper crater wall. This layering is visible in a Kaguya TC morning image (Figure 3) as well as in Apollo 12 oblique imagery, but NAC images provide the most detailed view. The layers are visible partially due to their tendency to form steeper parts of the wall, thus casting shadows. However, based on scarce LOLA point data, even the steepest slopes are only ∼40°–47°. Another factor in the visibility of layers is their color. Particularly in smaller exposures on the southern and northeastern walls, the darker color of the layers is emphasized by debris flows of dark material originating from the layered rocks. Yet, some of the outcrops of layered rocks have a brighter albedo than the surrounding crater wall. A more detailed assessment about the origin of this variation, and the relative importance of factors like composition, maturity, and illumination geometry lies beyond the scope of this study.
 The most prominent layered section on the northwestern wall is ∼100 m high, the best exposed part being ∼65 m (Figure 11). Less striking layers are present ∼400 m higher, close to the rim crest (Figure 11). The elevation of the layers varies somewhat in different parts of the crater wall, but the best exposed sections in the northwest are ∼760–860 m below the lunar reference elevation, i.e., several hundred meters higher than the surrounding maria (Figure 2). The exposed section of the layers in the northern and northwestern wall is ∼2500 m in length, with ∼1200 m of practically continuous outcrop. Individual layers have a maximum thickness of ∼3–5 m. There are also a few boulders in different parts of the terrace zone which display weak traces of probable layering.
 The upper boundary of the exposed wall material is often difficult to decipher, as the illumination geometry strongly affects differentiating the wall lobes from the exposed wall (particularly other than the distinctly layered parts). It should also be noted that not all of the material classified as exposed wall material in Figure 4 is literally exposed bedrock. Commonly it is covered by what appear to be a thin scree deposit and long, narrow, digitate clastic debris flows forming talus deposits. In addition, debris fans deposited on the units below are common on the lowermost part of the exposed wall.
4.3. Ray and Ejecta Distribution
 A sketch map of the distribution of continuous and discontinuous ejecta deposits, as well as rays around Kepler is provided in Figure 12. The continuous and discontinuous ejecta deposits (Figures 4 and 12) are classified as being melt-poor, as the melt-rich proximal ejecta deposits are described separately. The continuous ejecta grades into discontinuous ejecta, and the discontinuous ejecta grades into rays. The continuous ejecta near the crater rim often forms narrow linear ridges, oriented transverse with respect to the crater radial. Farther out these grade into larger dune-like features and radial deposits [e.g., Oberbeck, 1975; Settle et al., 1979]. The difference between continuous and discontinuous ejecta deposits used in this study is the notable presence of secondary craters, marking the start of discontinuous deposits. There is not a distinct boundary between the continuous and the discontinuous deposits, but in general the continuous deposits are symmetrically distributed around the crater, extending to slightly less than a crater diameter from the rim. However, the discontinuous ejecta blanket is wider to the east (∼2.5 crater diameters from the rim) than to west and north (∼1.5 crater diameters from the rim). This is opposite of the distribution of rays (see below). From a stratigraphic point of view it is worth noting that based on NAC imagery, the continuous ejecta deposits cover crater Kepler F (Figure 2) ∼15 km west of Kepler, and the discontinuous deposits cover Kepler A (Figure 2) ∼42 km southeast of Kepler.
Rays are defined here as bright linear albedo features radial to Kepler, located beyond the discontinuous ejecta (Figure 12). Thus, in the sketch map the rays are shown only beyond the discontinuous ejecta blanket (similar to the work of Hackman ), although in some cases the most prominent rays can be traced on the discontinuous ejecta blanket as well. In Clementine titanium and iron maps the rays, as well as the more proximal ejecta, are seen as features darker than the surrounding maria, i.e., the material ejected from Kepler (as well as from Copernicus and Aristarchus) has a lower titanium and iron abundance than the maria surface. Contrary to the mapped discontinuous ejecta distribution, Kepler's most prominent and longest rays are concentrated on the northern and northwestern sides of the crater. In a sector from west-southwest to northwest the rays reach a distance of at least ∼460 km from the center of Kepler, whereas south and east of Kepler they extend only ∼200–275 km.
4.4. Physical Properties of Melt-Rich Lithologies
 We estimated the physical properties of the impact melt, particularly its yield strength (τy) and bulk effective viscosity (ηe). The yield strength (in Pa) estimates were based on shadow length measurements (in meters) of the wall lobe thicknesses [Hulme, 1974; Moore et al., 1978] using equation
In equations (1)–(3), ρ is the density of the melt (taken, based on noritic composition and MELTS calculations, to be 2665 kg/m3 at 1200°C), g is the surface acceleration of gravity (1.62 m/s2), θ is the downhill slope in degrees, H is the thickness of the lobe (in meters), wf is the width of the flow, and wb is the width of the levee (∼17 m in the most prominent case).
 Yield strength measurements based on flow thicknesses (equation (1)) were made from the upper western wall, where wall lobes are ubiquitous and favorably illuminated in NAC M104769976RE (1.2 m/px; Figure 9). The slope of the upper wall is ∼29°. The measured lobe thicknesses are generally on the order of a few meters, varying from ∼0.8 m to possibly slightly more than 5 m.
 Levee widths were measured only from one channel close to the lobes (visible on the same NAC image, as well as on M148414768RE (0.6 m/px)), but situated on the outside crater rim flank (Figure 10). The slope of this part of the rim flank is ∼4°–5°.
 Given the assumptions and measurements described above, equation (1) implies that the yield strength of the wall lobes is typically ∼4–5 kPa (Figure 13), varying from ∼2 to ∼10 kPa. For the same lobes, equation (2) gives unrealistically low strengths. This is due to the very large width of the wall lobes, i.e., equation (2) is applicable to more channelized flows, not to a wide zone of flowing melt. The yield strength of the leveed flow is somewhat lower than that of the wall lobes, being ∼1 kPa, in agreement with the melt flow morphology. The wall lobes with higher yield strength are relatively short (typically on the order of a 100 m, although their layered nature makes it difficult to define the starting points of individual lobes) and wide, whereas the leveed flow is longer (altogether ∼1700 m) and narrow. Both seem to have their origin in the melt layer deposited on the rim, now seen as thin veneer covering most of the rim crest (Figure 4). The change in temperature or composition has only a relatively minor effect on the yield strength (or the density) of the liquid phase (Figure 13).
 Yield strength enables estimation of the flow thicknesses for cases where it cannot be measured based on shadow lengths. According to Hulme , the flow thickness (in meters) at the surface of the centerline of the flow (H0) is given by
and the critical thickness (Hc), which defines the minimum thickness (in meters) of the flow for downhill movement to occur, by
where θG is the slope expressed as a gradient (i.e., the difference in elevation divided by the distance, θG being equal to tan θ). The width of the leveed flow (Figure 10) on the western flank is not easy to measure, but it is ∼100 m. Assuming a yield strength of 1 kPa, according to equation (4) the thickness of the flow at the centerline is ∼5 m. Conversely, assuming the leveed flow is ∼5 m thick, equation (2) gives a yield strength of ∼1 kPa for a ∼100 m wide flow. For the critical thickness (Hc), assuming a measured slope (θ) of 5° and therefore a gradient (θG) of 0.875 for the leveed flow (wb = 17 m), equations (5) and (6) give approximately Hc = 2.6 m and Hc = 3.0 m, respectively. The good agreement between the results from equations (4), (5), and (6) implies that at least for this leveed flow, a thickness on the order of ∼5 m and yield strength of ∼1 kPa are reasonable.
 The viscosities calculated with MELTS subroutine Adiabat 1_ph assume that no clasts or crystals are entrained in the melt. However, impact melt rocks typically consist of melt matrix with a variable amount of relic lithic clasts, and crystals growing from the cooling melt. Therefore, we calculated the effective viscosity (ηe) of the bulk melt by applying equations relating the change in relative viscosity to particle (clast or crystal) concentration [e.g., Petford, 2009]. In addition to particle concentration, particle shapes also have a significant effect on the viscosity [e.g., Mueller et al., 2011]. However, compared to the large uncertainties in estimating the concentration of particles and their effect on the impact melt viscosity, the effect of particle shapes is probably minor, and was not considered in the following calculations. The equations for effective viscosity we applied are
 As discussed by Stickel and Powell  and Petford , for example, different effective viscosity equations generally agree with each other in low particle concentrations, but they differ substantially at higher concentrations. The equations used here were chosen to portray those uncertainties, thus providing a reasonable estimate for the possible range of effective viscosities encountered in Kepler impact melts.
 The results (Figure 14) imply that a superheated (∼1700°C) melt without clasts has a viscosity of ∼5 Poise ( = 0.5 Pa·s), whereas melt entraining 30 vol.-% clasts has a higher viscosity of ∼19 Poise (equation (7)). At a lower melt temperature (∼1275°C) closer to the liquidus, the viscosities are ∼120 and ∼430 Poise for a clast-free melt and a melt including 30 vol.-% clasts, respectively. At high temperatures, equations (7) (E–R) and (8) (G–N–A) produce similar results in the 30 vol.-% clasts case, whereas at lower temperatures, equation (9) (K–D) is a closer match to equation (8) (G–N–A). For 50 vol.-% clast contents, equation (8) (G–N–A) gives very high viscosities, particularly for temperatures below the liquidus. For comparison and setting reasonable minimum and maximum viscosity estimates for pure melt in Kepler, Figure 14 also includes clast-free viscosity curves for the average composition of basalts from the Apollo 12 landing site, as well as for average anorthosite (Table 1) [Haskin and Warren, 1991]. These compositions can be considered the end-members of the possible target lithologies, and, therefore, also of the impact melt.
where Q = 6·10−4. For the leveed flow (τy = ∼1 kPa; Figure 10), equation (10) results in a viscosity estimate of ∼9.5·104 Poise (9500 Pa·s), and for the wall lobes (τy = ∼4.5 kPa; Figure 9) ∼3.5·106 Poise (3.5·105 Pa·s). In our model calculations with a noritic composition (Figure 14), such viscosity values are applicable to melts below the liquidus with clast contents of ∼30–55 vol.-% (equation (9), K–D). It is worth noting, however, that the E–R equation (7) does not predict such high viscosities at all, and using the G–N–A equation (8) with ∼50 vol.-% clasts contents, these viscosity values are encountered even above the liquidus.
 In addition to effective viscosity and yield strength, we estimated the flow velocities of the impact melt-rich lithologies. The mean flow velocity (U) of a laminar flow of a Newtonian liquid is given by the Jeffreys  equation
where f is a friction factor, taken by Wilson and Mouginis-Mark  to be ∼0.01. Note that in a turbulent flow the viscosity is irrelevant unless it is significantly higher than about 10000 Poise [Wilson and Mouginis-Mark, 2001]. In the present study, the limit of significantly higher viscosities than 10000 Poise is generally reached at temperatures below ∼1250°C, i.e., in melts that are starting to crystallize, although this is strongly dependent on the clast content and particularly on the equation used (Figure 14).
 As an example, we assume a 2 m thick flow on a 25° slope (similar to the wall lobes on the upper part of the crater wall; Figure 9) with 30 vol.-% clasts at a temperature of 1200°C. Such a melt is starting to crystallize and has a viscosity of ∼2200 Poise (equation (7) and Figure 14). If it behaved as a Newtonian fluid and flowed in a laminar fashion (equation (11)), it would have a velocity of ∼11 m/s (∼40 km/h). However, impact melts are Bingham fluids, i.e., a yield stress must be applied before any flow can take place. Depending on, for example, local topography, impact melts also at least partially are likely to be turbulent, so a velocity of ∼17 m/s (∼60 km/h; equation (12)) is a better estimate. In this case (near liquidus), equations (11) and (12) are in a good agreement, although at higher temperatures (and, therefore, lower viscosities) the velocity estimates provided by the two equations can differ by two orders of magnitude.
 To first order, the individual wall lobes may have been formed within tens of seconds from the time the melt was emplaced on the rim. The leveed flow would have initially formed in a few minutes, but the flow likely continued for a slightly longer time to produce the exterior pond at the end of the flow. As can be seen in Figure 10, the melt flowed in several episodes, as there are several flow fronts on the side of the main flow where a small local depression diverted part of the flow. Similarly, the wall lobes overlap each other (Figure 9), indicating that the formation of the entire zone of wall lobes was a multiphase process.
5.1. The Layered Material and Mare Basalt Thickness
 Traces of layering can be seen around most of the periphery of Kepler, but the layers are most pronounced on the northwestern upper crater wall (Figure 11). The most prominent layered section is ∼100 m thick, but layered material is also present closer to the crater rim crest (Figure 11), ∼400 m above the lowermost layers. We tentatively interpret these layers as individual mare basalt flows, uplifted by the rim formation by several hundred meters (Figure 2). Hence, ∼400 m represents the total minimum thickness of the mare basalts. This agrees reasonably well with the estimates by Hörz  and Heather and Dunkin .
5.2. The Volume of Impact Melt
 As Kepler is a fresh, essentially undegraded crater, the observed distribution of the main impact melt-rich units closely reflects their original distribution. Thus, the ejected impact melt was never a homogenous sheet, but instead formed discrete pools in the ejecta; these pools were mainly concentrated on the downrange side. After the deposition, minor flowing and ponding of melt-rich material in local depressions took place. However, the volume of the impact melt-rich lithologies in and around Kepler is very poorly constrained. Approximately 0.5 km3 can be regarded as the minimum estimate of melt-rich material within the smooth and hummocky floor units. According to Cintala and Grieve , about 45% of the melt produced should be retained in the crater during the formation of a 25 km transient cavity (assuming an anorthosite target and a chondritic projectile). Thus, based on direct measurements and estimates from imagery, ∼1.1 km3 (i.e., ∼0.5 km3 or ∼45% interior, and ∼0.6 km3 or ∼55% exterior) of impact melt is the minimum estimate of the total melt produced in the Kepler impact.
 This value can be compared to predictions from different scaling relationships. Assuming the most probable impact angle of 45° [Gilbert, 1893; Shoemaker, 1962], which for the Kepler case also has some observational support (see below), a 1900 m diameter projectile with a density of 3000 kg/m3 impacting at 17 km/s, a basaltic target composition with a density of 2800 kg/m3, and a transient cavity diameter of ∼25 km, scaling laws by Abramov et al.  suggests that ∼16 km3 of melt should have been produced. If an anorthosite target is assumed, the melt volume is ∼36 km3 [Abramov et al., 2012]. This is in good agreement with a scaling law prediction of ∼34 km3 by Cintala and Grieve . If 16–36 km3 of melt was produced and ∼7.2–16.2 km3 (45%) of it remained inside the crater, a ∼65–145 m thick layer of clast-free melt, or ∼130–290 m of melt with 50 vol.-% clasts would cover the 113 km2 now occupied by the smooth and hummocky floor units. Therefore, the melt-rich lithologies on the floor of Kepler are at least 5–10 m thick (observations of ledges), and depending on their clast contents, may approach a thickness of ∼65–290 m (scaling law predictions).
 After their emplacement, the melt-rich materials on the crater floor subsided. This is indicated by a topographic bench with heavy fracturing at the location of the highest flexure, best seen between the eastern hummocky floor and the terrace zone (Figure 7). This is similar to, for example, the floor of Aristarchus impact crater, or the Kilauea Iki volcanic crater in Hawaii [Strom and Fielder, 1970]. In addition to cooling contraction, the subsidence was probably mainly the result of compaction and material flowing into the empty spaces particularly within the hummocky floor unit.
5.3. The Origin and Physical Properties of the Impact Melt
 As mentioned above, ∼400–500 m seems to be a reasonable estimate for the total basalt thickness in the Kepler region. Given Kepler's proximity to the Imbrium Basin (Figure 1) and the presence of up to 700 m high ridges mapped as Imbrium ejecta (Figure 2; [Hackman, 1962; Wilhelms and McCauley, 1971]), Imbrium ejecta is another constituent of the Kepler impact melt rocks. According to McGetchin et al. [1973, equation 3], the thickness of Imbrium ejecta in the Kepler area is expected to be ∼380–660 m (see Kring  for a description of these uncertainties), in a reasonable agreement with the height of the ridges. Thus, as a first order estimate, the mare basalts and Imbrium ejecta together constitute perhaps 1 ± 0.2 km of the target stratigraphy.
 Assuming Kepler had a transient cavity diameter of ∼25 km [e.g., Collins et al., 2005] and other assumptions as above, the depth of melting likely reached ∼4 km [Cintala and Grieve, 1998]. Based on a scaling law by Cintala and Grieve , the stratigraphic uplift in Kepler is expected to be ∼3 km, somewhat less than the minimum depth of the origin of central peaks established by the ∼4 km estimate for the depth of melting. Taken together with the lack of olivine spectral signature [Le Mouélic et al., 1999], these estimates imply that olivine-rich material is not present in Kepler at the depth of ∼3–4 km, from where a substantial portion of the melt originated [e.g., Cintala and Grieve, 1998]. Although the relatively thin mare basalt layers must have contributed to the composition of the impact melt, they most likely were not the main source. Hence, the composition of the Kepler impact melt is probably not more mafic than that of the mare basalts. On the other hand, given the presence of pre-impact mare basalts, the interpreted noritic spectral signature of the crater wall [Le Mouélic et al., 1999], and the largely noritic composition of Imbrium ejecta, the probable Kepler impact melt composition cannot be approximated as pure anorthosite either. Given these boundary conditions, our best current estimate is that the Kepler impact melt composition is probably noritic. Therefore, the viscosity and yield strength of the (clast-free) Kepler impact melt also must lie between those of basalt and anorthosite, and are most likely reasonably well approximated by average norite composition, as in Figures 13 and 14.
 We envision Apollo impact melt breccia samples 14303, 14306 and 15405, which have noritic melt compositions and variable clast/melt-ratios [Wiik et al., 1973;Ryder and Bower, 1976; Simonds et al., 1977; Ryder and Spudis, 1987], to be reasonable conceptual analogs for the impact melt-rich rocks in Kepler. Sample 15405 has flow banding [e.g., Ryder and Bower, 1976] that might be representative of textures found in melt rock deposits that occur along Kepler's crater wall, along the crater floor, and in flows streaming outward from the crater rim (Figure 10). Samples 14303 and 14306 are without any significant melt matrix, so they are potential textural and compositional counterparts to material along the base of impact melt rock units throughout the crater, or interspersed with melt-bearing units in the hummocky floor material (Figures 5–67). We also anticipate textural analogs to Apollo samples 68815 (a vesicular glassy polymict impact melt breccia) and 68415 (a clast-poor impact melt breccia) [Ryder and Norman, 1980] in the central part of the crater and potentially along the crater walls where flows are observed (Figures 8–9). Preferential ejection of melt spherules from the uppermost mare lithologies might distribute basaltic melt spherules, like sample 60095, far beyond the crater rim [Ryder and Norman, 1980; see also Stöffler et al., 2002].
 Earlier estimates of the viscosity of lunar impact melts seem to be scarce. Hulme and Fielder  assumed the viscosities for Copernicus and Tycho melt flows to be 107 Poise and for Aristarchus 3000 Poise. These values are based on the assumption that they are similar to the viscosities of terrestrial lava flows having comparable yield strengths. Hulme and Fielder  do not provide a temperature for their viscosity values, but they may be referring to liquidus temperatures. If this is the case, the viscosity of Copernicus and Tycho flows is about four orders of magnitude higher than our estimate for Kepler, whereas the viscosity of Aristarchus flows is very close to that of Kepler (assuming 30 vol.-% clasts and K–D equation (9)). Our viscosity estimates based on yield strength (equation (10)) [Moore and Ackerman, 1989; Wilson and Head, 2003] fit the model calculations with MELTS (equations (8) and (9) and Figure 14), thus supporting their validity.
 Our best estimate for typical yield strength of the Kepler impact melt forming the wall lobes is ∼4–5 kPa (varying from ∼2 to ∼10 kPa), and ∼1 kPa for the leveed flow (Table 2). The flow fronts associated with the leveed flow, the curving of the leveed flow, and the presence of the leveed flow on the melt pond where melt-rich material had pooled before the formation of the levees, clearly indicate that this melt-rich material was not emplaced in its current location ballistically. Instead, it was deposited in a late-stage process reminiscent of terrestrial lava flows, with measured and estimated rheologic properties similar to other lunar impact melts and terrestrial analogs (e.g., Kilauea basalts [Moore et al., 1978]).
Table 2. Calculated Yield Strengths of Lunar Impact Melt Flows
 In general, our map of ejecta distribution (Figure 12) coincides remarkably well with that given by Hackman . Nevertheless, there are some differences. Crater Kepler F, located ∼15 km west of Kepler's rim (Figure 2), was mapped by Hackman  as having rays on top of Kepler's continuous and discontinuous ejecta, thus making Kepler F younger than Kepler. However, modern imagery does not show any traces of these rays, and ejecta from Kepler can be seen deposited on the wall of Kepler F. Some images, particularly from AMIE, show traces of a deposit that partially coincides with Hackman's  Kepler F ray distribution. We interpret this as continuous ejecta from Kepler F, partially seen as a topographic feature blanketed by Kepler ejecta, which according to McGetchin et al.  should be about 20 m thick at Kepler F.
 Our map of the continuous ejecta (Figures 4 and 12) is quite symmetric, extending slightly less than one crater diameter from the rim in every direction, whereas in Hackman's  map the continuous ejecta extends notably farther to the east than to other directions. Both mapping efforts, although ours should be regarded as a sketch, agree that the discontinuous ejecta (Figure 12) is distributed farthest in the east. This directional pattern contrasts with the ray and melt-rich material distribution (see below).
 The rays of Kepler are clearly not distributed symmetrically around the crater (Figure 12), which has been previously noted [e.g., Elger, 1895]. This has been partially attributed to “shielding” effects by topographic highs, briefly discussed by Fielder , and also mentioned by Baldwin . Based on terrestrial imagery, one of the conspicuous gaps in the most prominent rays was interpreted to have its vertex at a southwest–northeast oriented ridge [Fielder, 1961] of rugged material at the edge of the discontinuous ejecta blanket 45 km west of Kepler's rim, rising about 650–700 m above the mare surface (Figure 2). However, faint rays are present up to ∼35 km west of the ridge before they disappear, as was also noted by Fielder . Other roughly similar apparent gaps in the ray pattern can be seen to be associated with smaller ridges ∼65–90 km west-northwest from Kepler. Somewhat lower rugged ridges ∼60 km north of Kepler have a generally ∼north–south (or Kepler-radial) orientation (Figure 2), which Fielder  suggested was why no gaps in the ray pattern could be seen.
 However, many major gaps in the ray pattern have no apparent association with any of the topographic highs, and Clementine UVVIS imagery reveals that in addition to the weak rays being present immediately downrange (west) of the major ridge, the prominent rays west of Kepler that drew Fielder's  attention actually continue across the ridge, and can be traced back to what is mapped (both in Figure 12 and by Hackman ) as the discontinuous ejecta blanket. Similar observations can be made with respect to the other rugged ridges as well. Thus, it seems unlikely that the ridges around Kepler would have blocked the ray-forming fragments, resulting in the asymmetries in the ray distribution. Some topographic blocking may have occurred, but the apparent correlation of the ridges and gaps in particularly lower resolution imagery seems mainly coincidental. In addition to impact angle (see below), asymmetries in the ejecta distribution are largely controlled by pre-existing target inhomogeneities, as was shown by Shoemaker [1960b, 1962]. This is one plausible explanation for the asymmetry of the Kepler ray system.
 A possible contributing factor in the apparent ray asymmetry is the effect of rays from other Copernican craters. Rays from Tycho are parallel to Kepler rays in the southeastern part, and some of the Aristarchus rays parallel Kepler's rays in the northwestern part. Even in these areas, however, in most cases the rays can be reliably assigned to Kepler, Aristarchus, or Tycho. By far the most significant “alien” ray contribution comes from Copernicus. East of Kepler, on Mare Insularum, in some cases it is impossible to assign a certain origin for a weak ray segment. Figure 12 includes only those rays where a genetic relationship with Kepler could be determined with reasonable certainty. However, even in the eastern areas, the overall effect of rays from other craters than Kepler can be assumed to be fairly low.
 The factors described above are not enough to account for the entire asymmetry. In the zone from west-southwest to northwest, the rays extend at least to about 460 km from the center of Kepler, whereas they reach only ∼200–275 km in the south and east. The most prominent rays are concentrated in the western and northwestern sectors (Figure 12). We believe that the most plausible explanation for most of the observed ray asymmetry around Kepler is an oblique impact, as was mentioned in passing by Gault and Wedekind . Based on the ray distribution, the direction of the impact was somewhere from the southeastern sector.
 No large-scale experimental data exists about the distribution of distal ejecta (rays) in oblique impacts, but numerical modeling of the Ries-Steinheim impacts and the formation of the moldavite strewn field suggests that in impact events occurring at angles of 30°–50° from the horizontal, distal ejecta forms a fan-shaped field with an apical angle of ∼75° of the fan [Stöffler et al., 2002]. As the Kepler rays are not so tightly clustered but merely strongly asymmetrically distributed, 30°–50° is probably the lowermost possible impact angle. This estimate also agrees with the generally symmetric distribution of the proximal ejecta and the fairly circular (scalloped or somewhat polygonal; see Öhman  for a review on crater morphology) plan view and rim topography of Kepler, as well as the presence of smooth floor material on the inferred downrange side of the crater floor [Gault and Wedekind, 1978; Schultz and Anderson, 1996; Herrick and Forsberg-Taylor, 2003].
 However, the mapped shape of the discontinuous ejecta blanket (or the continuous blanket as mapped by Hackman ) depicted in Figure 12 does not fit the impact direction inferred from the ray and melt rock distributions. This is possibly due to observational bias. The surface material near the eastern edge of the discontinuous ejecta blanket is dominantly not mare, but more rugged material with higher albedo interpreted as Imbrium ejecta [Hackman, 1962; Wilhelms and McCauley, 1971]. Thus, clear-cut rays are more difficult to detect, but secondaries and scouring of the surface, typical for the discontinuous ejecta blanket, can be seen. If the target material was darker mare, rays probably would be most dominant ejecta feature, and therefore the area would be classified as rays.
 Taken together, the observations of the asymmetric distribution of interior and exterior melt rock deposits and rays, and the symmetric continuous ejecta blanket and the crater rim topography and morphology imply that the impact direction was from the southeastern sector, but the angle of impact was not very shallow, i.e., not notably less than 45°.
5.5. Comparison With Other Impact Craters
 There are very few well-preserved craters on Earth to evaluate the details of impact melt distribution and emplacement. The best terrestrial analog for Kepler crater is the Ries crater, a ∼24 km diameter complex crater. Like Kepler, it was produced from a layered target, albeit one of layered sediments over a crystalline basement [e.g., Pohl et al., 1977; Stöffler et al., 2002; Kring, 2005, and references therein] rather than layered mare basalt flows over basin ejecta. As discussed above, both craters display asymmetric distribution of distal ejecta (Figure 12). At Ries, melt-bearing clastic polymict breccias (i.e., suevite) were deposited in the modification zone and appear to have ponded between blocks [e.g., Pohl et al., 1977; Kring, 2005]. There are, however, only a few known melt rock ponds several tens of meters in diameter that have been found at the Ries on the downrange inner rim slope of the crater [e.g., Stöffler et al., 2002]. The well-known impact melt rock occurrence at Polsingen displays an elongate shape, implying an origin as a melt flow on the crater wall. However, the evidence of impact melt rock distribution and the types of deposits it forms is patchy due to erosion, vegetation, and human occupation. The clarity of constraints on the interpretations provided by Kepler is dramatic.
 The third highest lunar science priority of the NRC  is to calibrate the impact flux to the Moon, and Kepler is a key site for doing so [Kring, 2009]. Kepler is also the smallest Copernican complex crater recently recommended as an exploration target [Kring, 2009], and thus the most “compact” impact site where exploration with major science return can be expected. From the mission planning perspective, Kepler's location has some notable advantages for possible robotic and/or manned in situ exploration. The nearside location enables continuous communication, and the proximity to the equator is beneficial for a simple and relatively low delta-v (i.e., change in velocity) trajectory [Patapoff, 1967]. However, Kepler's 14 Earth-day period of sunlight presents temperature and illumination limitations that have to be taken into account. From the scientific point of view, however, the most crucial practical aspect is the slopes. In previous recommendations for future mobility capabilities, Lunar Exploration Science Working Group (LExSWG)  recommended rovers be able to traverse up to 25° slopes. A preliminary robotic asset used at Meteor Crater [Kring et al., 2007] ascended up to a 40° slope. A crew rover is still in development. The chassis was initially designed to accommodate up to 15° slopes in analog sites on Earth [Harrison et al., 2008]. When outfitted with a complete cabin, simulating a small pressurized rover, the Lunar Electric Rover (LER, or Space Exploration Vehicle, SEV) climbed 18°–20° slopes on cinder-covered volcanic vents, suggesting a 25° slope capability on the Moon is reasonable. Using that constraint, we consider two types of exploration options at Kepler crater.
 First, we consider two alternatives for a manned rover mission landing inside Kepler, with a maximum traverse distance of 10 km from the landing point (10 km being the walk-back distance). Landing area A is on the northern smooth floor material just north of the highest central peaks, close to 8.27°N, 322.04°E (Figure 15a). This area has a low fracture and hummock density, making landing and particularly traversing easier and safer than on much of the southern crater floor (see, e.g., Figure 5), and giving a potentially high science return. In the immediate vicinity of the landing area A is, for example, a ∼800 m long discontinuous fracture (Figure 15a) and pit system in the melt-rich floor material. This section provides ample possibilities for studying the fundamental questions related to, for example, impact melt crystallization processes (NRC  goal 6a) and the deposition of crater-fill materials. Clast-poor impact melt rocks, most likely to be discovered in the smooth floor unit, also offer the best possibilities for radiometric dating, providing crucial constraints on the Copernican cratering flux (NRC  goal 1d).
 Landing area A also has some collapse pits where one side of the pit has an impact melt rock ledge, but the other side is seemingly shallow, thus enabling rover studies of the pit interiors and potential caves within the melt-rich crater floor material. More pits are commonly found in the southern hummocky floor (Figures 5 and 15c). In addition to their scientific value, caves in the melt-rich rock would be highly useful sites for future exploration as natural shelters. From landing area A, also the largest crater (D = 500–550 m, located on the southwestern floor at 8.08°N, 321.84°E) formed in the melt-rich floor materials in Kepler can be reached, providing excavated materials from the depth of ∼90 m [Croft, 1980].
 The importance of ledges, cliffs, and freshly excavated material is emphasized by the continuous regolith production on the Moon. Assuming the age of Kepler crater is ∼100–800 Ma [Basilevsky et al., 1977; König et al., 1977; Wilhelms, 1987], and that during that time the regolith accumulation rate has been ∼1 mm/My [Quaide and Oberbeck, 1975], regolith thickness in Kepler would be on the order of 10–80 cm, thus significantly hampering the possibilities of sampling actual bedrock. However, locally steep slopes (i.e., over ∼31° or the angle of repose [Quaide and Oberbeck, 1968]), like the walls of fractures or pits in the smooth floor material (Figures 5 and 15), provide access to true outcrops of bedrock without a significant regolith cover.
 Unlike the prominent central peaks of larger complex craters such as Copernicus or Tycho, the central uplift complex of Kepler is formed by fairly gently sloping hills. Thus, some of the central peaks (Figures 6 and 15b) of Kepler can be traversed (slopes ∼10°–21°). The central uplift provides opportunities to sample the upper crust of the Moon to a depth of about 3–4 km. As Kepler is in the central PKT [Jolliff et al., 2000], sampling the central uplift will improve our understanding of the extent and composition of KREEP-rich lithologies (NRC  goal 3a), and quantify the complexity of the lunar crust (NRC  goal 3d). In addition, the mechanics of central uplift formation can be studied in detail. However, central uplift traverses are complicated by the numerous large boulders (and presumably finer-grained loose material) covering much of the central peaks (Figure 6; see also Figure 15b). Conversely, boulders that have rolled down the slopes provide sampling opportunities with some stratigraphic control.
 Much of the terrace zone and the massive slumps can also be reached after landing on the northern crater floor in area A, as the slopes are generally 20°–25°. This enables structural studies of terrace formation and the general crater modification mechanics (closely related to NRC  goal 6c). However, the upper part of the crater wall is fairly steep with average slopes of 30°–40°, the steepest slopes being caused by the exposed layered material (Figure 11). Thus, if a rover landed on the floor of Kepler, getting out of the crater will likely be impossible, and therefore in situ analysis of the upper exposed wall materials could not be performed. Despite this, sampling of the layered portions with some stratigraphic control is possible, because debris flows originating from the layered materials reach the terrace zone. If the layers are individual basalt flows, this enables achieving several of the goals of the NRC  concept 5 concerning lunar volcanism.
 As an alternative, we also briefly examine the possibilities presented by a slightly higher-risk landing in area B (8.16°N, 321.98°E), which is in a relatively flat part of the hummocky floor, and largely surrounded by the main central uplift complex (Figure 15b). Essentially all the science goals that can be reached by landing in area A are also available from area B, but the latter option provides some additional benefits. This is because the entire floor of Kepler is within a 10 km radius. However, traverses from area B are more complex due to hummocky floor material with large heavily fractured areas. As an example, area B enables sampling of a large layered boulder on the southern terrace zone at about 7.83°N, 322.00°E (Figure 15e). An unusual, ∼770 m wide radial graben on the slope of the terrace zone (at 7.85°N, 312.13°E) would also be within reach. The western terrace zone (at about 8.10°N, 321.76°E; Figure 15d) enables studying a section of the lower crater wall and the terrace formation, as well as small overlapping layers of impact melt-rich material.
 In our second exploration scenario, crew and a lunar rover are deployed outside Kepler. In this case, part of the impact melt rock studies are still possible due to the presence of extensive exterior ponds (particularly north of Kepler; Figure 4), rim veneer, and some flows (Figure 10). These are easily reached with LERs, as the rim flank slope is shallow (<10°).
 Landing outside the crater offers possibilities not readily available inside the crater. Kepler is surrounded by mare units of several age groups, including the youngest mare surfaces so far identified on the Moon [Hiesinger et al., 2003; Morota et al., 2011]. Therefore, NRC  goals 5a, 5b, and 5d regarding the origin, variability, and the ages of the mare basalts, as well as the evolution of volcanism can be addressed. The Kepler target stratigraphy can be studied by traversing the proximal ejecta, as ejecta blankets preserve the original stratigraphy of the target, albeit inverted [e.g., Shoemaker, 1960a, 1963] and mixed with local material [e.g., Oberbeck, 1975]. Craters Kepler A, B, C, and F (Figure 2) have diameters of ∼7–11 km, and, therefore, excavation depths on the order of 1 km [Croft, 1980]. Hence, these craters, although older than Kepler (at least Kepler A and Kepler F), provide additional probes of mare units. Kepler is also in the heart of the PKT [Jolliff et al., 2000] and may be located on a large shield volcano [Spudis et al., 2011; see also McGovern and Litherland, 2011], both worthy and accessible targets for detailed sampling, addressing NRC  concepts 2 and 5 about the lunar interior and volcanism, respectively.
 If the 10 km walk-back restriction is relaxed, other potential research targets outside Kepler include the massifs interpreted to be Imbrium ejecta [Hackman, 1962; Wilhelms and McCauley, 1971], the closest one being ∼20 km north-northwest of Kepler (Figures 2 and 4). These massifs provide an opportunity to refine the pre-existing age estimates for one of the most significant basin-forming impacts in lunar history [e.g., Wilhelms, 1987], thus addressing the most important goal set by the NRC . In addition, ray materials from both Copernicus and Aristarchus reach Kepler, and although these may prove to be difficult to identify on the lunar surface, sampling is strongly encouraged, as successful dating of these materials along with Kepler melt rocks will greatly improve our understanding of the Copernican cratering flux (NRC  goal 1d).
 Farther from Kepler, a possible isolated volcanic dome ∼40 km northwest of Kepler [Lena and Wöhler, 2009] is a potential mission objective. In the future, with more enhanced traversing capabilities, another recommendable exploration area is the volcanic Milichius dome field ∼175 km northeast of Kepler [e.g., Baldwin, 1949; Wilhelms, 1987; Wood, 2003; see also McGovern and Litherland, 2011]. It should also be noted that the goals of concepts 7 and 8 of the NRC  regarding regolith processes and the lunar atmosphere and dust environment, respectively, can be addressed by a mission landing either inside or outside Kepler crater.
 To summarize, if a manned or a robotic lunar mission explores the Kepler region, landing inside the crater will enable fundamental studies of impact melts and cratering mechanics, but landing outside and close to the rim of Kepler will provide more data on endogenic processes. Integrating the new LRO NAC and WAC imagery and LOLA topography with older data sets provides excellent tools for planning either mission.
6. Summary and Conclusions
 New geomorphologic sketch mapping of the Kepler crater (Figure 4) reveals that impact melt-rich material is present in a number of different morphologies both inside and outside the crater. Most of the photogeologically observable melt rock forms smooth and hummocky floor units, but their thickness is poorly constrained. Terrace ponds overlie both the terrace zone and the massive slump materials. The upper crater wall hosts a zone of mutually overlapping wall lobes of melt rock, having typical thicknesses from less than a meter up to perhaps five meters (Figure 9). Exterior melt-rich deposits include thin and poorly delineated rim veneer, which is closely related to the interior wall lobes and exterior ponds closest to the rim (Figure 10). However, most of the exterior ponds are farther beyond the rim crest, and heavily concentrated on the northern and northwestern sides of the crater (Figure 4).
 The yield strengths of the wall lobes are typically on the order of 4–5 kPa (Figure 13), comparable to other lunar craters (Table 2) and, for example, Kilauea basalts in Hawaii [Moore et al., 1978]. We also estimated the effective viscosity of the Kepler impact melts, assuming a noritic melt composition and different particle concentrations. Our best current estimates (Figure 14) for the viscosity of Kepler impact melt at the liquidus (∼1275°C) is ∼2000–15000 Poise (200–1500 Pa·s), but we acknowledge that the currently limited understanding of the effect of particle shapes and concentration on the effective viscosity of magmas [e.g., Petford, 2009] makes this an order of magnitude estimate at best.
 The upper northeastern crater wall displays distinct layering ∼100 m thick with individual layers having a maximum thickness of ∼3–5 m (Figure 11). Altogether layered material can be seen to span ∼400 m in elevation. We interpret these layers as mare basalt flows, and ∼400 m as the total minimum thickness of the eastern Procellarum mare basalts.
 Both interior and exterior melt rock deposits (Figure 4), as well as rays (Figure 12), are distributed asymmetrically around Kepler. Our observations imply an oblique impact from the southeastern sector, with an impact angle of ∼45° (but not substantially lower).
 The Kepler area presents interesting exploration opportunities for a rover mission. In addition to duration and evolution of lunar volcanism (related to maria, domes, and possibly a shield volcano) and the origin of the PKT, these questions include, for example, Copernican cratering flux, age of the Imbrium impact, and various cratering processes.
 A. Nahm (UTEP) deserves our gratitude for her most valuable comments on an earlier version of the manuscript, as well as for help provided during the whole project. B. Fessler (LPI) is thanked for his expert assistance with ISIS and ArcGIS, and K. Joy (LPI/NASA JSC), G. Kramer (LPI) and J. Rapp (LPI/NASA JSC) for helpful discussions. L. Gaddis and the ISIS team (USGS), as well as the NASA Ames Stereo Pipeline team deserve our thanks for providing the vital software, and for their useful comments on running them. The efforts of the science teams involved in the production of all the data sets used, particularly the LROC, Kaguya TC and MI, and LOLA teams, is gratefully acknowledged. J. Spray, an anonymous reviewer, and associate editor M. Wieczorek are thanked for their constructive reviews. This research was funded by NASA Lunar Science Institute contract NNA09DB33A (PI David A. Kring). This is LPI Contribution 1651.