In 2018, the Chinese lunar mission Chang'E‐4 (CE‐4; Wu et al., 2017 ) will explore the SPA basin. It will be the first in situ exploration of the farside of the Moon. The selected landing region for CE‐4 (45°S–46°S, 176.4°E—178.8°E; Wu et al., 2017 ) is located on the southern floor of the Von Kármán crater. In situ exploration within the Von Kármán landing region will bring unprecedented imaging, spectral, radar, and low‐frequency radio spectral data for the landing region, and it will greatly improve our understanding about the compositions of farside mare basalt, SPA compositional zones including SPA compositional anomaly and Mg‐pyroxene annulus, regolith evolution, and the lunar space environment. Indeed, the U.S. National Research Council ( 2007 ) has identified key scientific priorities for future lunar exploration that can be addressed from the Von Kármán crater, including the possibility to study the existence and extent of differentiation of the SPA melt sheet (Morrison, 1998 ; Nakamura et al., 2009 ; Vaughan & Head, 2014 ) and possible exposed upper mantle materials (Melosh et al., 2017 ; Moriarty & Pieters, 2018 ). In this study, we have (1) investigated the geological characteristics of the landing region using available remote sensing data including topography, high‐resolution imaging, and reflectance spectral data; (2) identified targets of high scientific interests; and (3) proposed testable hypothesis for the CE‐4 mission.

(a) Color‐coded Lunar Orbiter Laser Altimeter (LOLA) digital elevation model of the South Pole‐Aitken Basin (altitude scale on the left). The polygon shows the location of subsets b and c. Letter “A” indicates the location of the Antoniadi crater. (b) Global normalized reflectance (incidence = 30, emission = 0, phase = 30) of the 643 nm Lunar Reconnaissance Orbiter Camera (LROC) Wide‐Angle Camera (WAC) mosaic of the Von Kármán area. The rectangle shows the location of Figures 5 a and 8 a. The yellow arrows indicate the ray materials on the floor of Von Kármán crater, and they are converged to Finsen crater. (c) LROC WAC mosaic of the same area. The white box shows the location of subset d. (d) LRO WAC mosaic of Von Kármán crater. The white box shows the selected landing region, which shows the locations of Figures 2 a and 2 b, red box in Figure 3 a, black boxes in Figures 5 c and 5 d, 7 , and 9 a– 9 d. The red, green, blue and yellow rectangles show the locations of Figures 3 b– 3 d and 4 c, respectively. We informally named the unnamed impact crater (diameter= 3.6 km; central coordinates as 43.3°S 176.1°E) near the landing region as, which we use in italics to indicate its informal designation. North is toward the top in subsets b–d.

Von Kármán crater (diameter D = 186 km; central coordinates as 44.4°S, 176.2°E; Figure 1 a) lies within the Mg‐Pyroxene Annulus, just northwest of the SPACA terrain. This crater is a pre‐Nectarian crater (Losiak et al., 2009 ; Yingst et al., 2017 ), younger than the Von Kármán M basin ( D = 219 km; central coordinates as 49.4°S, 174.9°E), and older than the Leibnitz crater (diameter D = 236 km; central coordinates as 38.2°S 179.3°E: Figures 1 b and 1 c; Yingst et al., 2017 ). Mare basalts flooded parts of the Von Kármán crater floor during the Imbrian Period (Pasckert et al., 2018 ; Wilhelms et al., 1979 ; Yingst & Head, 1997 ; Yingst et al., 2017 ), but a portion of the central peak remains exposed near the center of the crater (Figure 1 d). Finsen crater (diameter D = 73 km; central coordinates as 42.3°S, 182.3°E), Von Kármán L crater (diameter D = 29 km; central coordinates as 47.8°S 177.9°E), and a similar‐sized unnamed crater nearby (we informally denote it as Von Kármán L' in this study) were formed subsequent to the Von Kármán crater (Wilhelms et al., 1979 ; Figure 1 c). Relatively higher albedo linear features on the mare basalt of Von Kármán crater converge toward the crater Finsen (Figure 1 b), which is located within the SPACA, suggesting that the SPACA material has been derived from Finsen crater.

The farside South Pole‐Aitken (SPA) basin is the largest known impact structure on the Moon (Stuart‐Alexander, 1978 ; Wilhelms et al., 1987 ). Its geology provides insights into the composition of the lower crust and upper mantle, the impact flux in early lunar history, the nature and evolution of basin‐scale impact melt deposits, and the nature of large impact basins and their formation and modification processes (Garrick‐Bethell & Zuber, 2009 ; Head et al., 1993 ; Moriarty & Pieters, 2018 ; Potter et al., 2012 ; Spudis et al., 1994 ; Vaughan & Head, 2014 ). The SPA basin has been studied with spectral observations (e.g., Ohtake et al., 2014 ; Pieters et al., 2001 ) and recently been subdivided into four distinct compositional zones based on Moon Mineralogy Mapper (M 3 ) data (Moriarty & Pieters, 2018 ): (1) a central ~700‐km‐wide SPA compositional anomaly (SPACA), which exhibits a strong Ca‐pyroxene signature, which is different from typical mare basalts; (2) a Mg‐Pyroxene Annulus, which is characterized by Mg‐rich pyroxenes; (3) a Heterogeneous Annulus, which exhibits mixing of localized pyroxene‐rich units and feldspathic materials; and (4) the SPA Exterior, which is mafic‐free and dominated by feldspathic materials. Pyroxene compositions in both the Heterogeneous Annulus and Mg‐Pyroxene Annulus are similar. Moriarty and Pieters ( 2018 ) have indicated that the Mg‐pyroxene unit is beneath the SPACA Ca‐pyroxene unit from the stratigraphic relationships and central peak exposures among craters within SPACA. The material of Mg‐Pyroxene Annulus was the main material excavated and melted by the SPA‐forming event due to the relatively uniform composition and the great area, depth, and thickness of the Mg‐Pyroxene Annulus. The extremely large size of the SPA basin strongly suggests that it has excavated subcrustal and mantle material (e.g., Melosh et al., 2017 ); however, the lack of evidence for widespread dunite or olivine‐dominated mineral assemblages, and the dominance of Mg‐pyroxenes in the basin interior, suggests that lunar mantle must include a significant Mg‐pyroxene component (Moriarty & Pieters, 2018 ).

In order to distinguish geological units of different composition, age, or texture, a RGB composite mosaic of MI data was built (R = 750 nm/415 nm, G = 950 nm/750 nm, B = 415 nm/750 nm (e.g., Huang et al., 2011 ; Weitz & Head, 1999 ). Observed variations in this RGB composite that mimic the Clementine false color ratio are commonly due to variations of the titanium and iron contents and the maturity of the surface materials (Pieters et al., 1993 ). We marked locations representative of different spectral units in the MI RGB composite. The mineralogy of these locations was derived from spectroscopic data from Chandrayaan‐1 M 3 instrument (Pieters et al., 2009 ). M 3 is a visible to near‐infrared hyperspectral imager, with 85 spectral channels spanning from 430 to 3,000 nm. The data used in this study are archived in the Planetary Data System (PDS, version 1 of Level 2), radiometrically corrected (Green et al., 2011 ), geometrically corrected (Boardman et al., 2011 ), thermally corrected (Clark et al., 2011 ), and photometrically corrected (Besse et al., 2013 ). The optical period used in this study is the OP2C2, with a resolution of 280 m/pixel. We selected this optical period because it covers the entire landing site region. The continuum of the spectra was removed with the algorithm developed by Martinot et al. ( 2018 ), and a suite of band parameters were calculated. The continuum is modeled as linear segments connected to the original M 3 spectrum in points called tie points, defined in fixed intervals. The algorithm maximized the 1 and 2 μm absorption band areas and automatically extracted the continuum. After continuum removal, a suite of spectral parameters is calculated (e.g., band center position, band area, and band depth) and exported as maps.

The absolute model age of the landing region is derived from crater statistics. Since obvious secondary craters (i.e., secondaries) that occur in chains and clusters have almost covered the entire landing region (see section 3.2 ), we employed two independent methods to estimate the model age. (1) The mare surface on the floor of Von Kármán was interpreted to be emplaced by one episode of basaltic volcanism based on the uniform reflectance spectral characteristics (Yingst & Head, 1999 ). Therefore, the model age for the mare surface is derived via craters that are substantially larger than the largest obvious secondaries (diameter D > ~2 km; see section 3.1 ) on the LROC WAC mosaic. (2) There are a few subareas, including two within the proposed landing region of CE‐4, where craters larger than ~300–400 m are not heavily modified by obvious secondaries. The model ages for these subareas are derived by studying the size‐frequency distributions of craters larger than 300–400 m in diameter using Kaguya TC mosaics. During analysis of the crater statistics, the CraterTools toolbox in ArcMap was employed to collect the craters by fitting circles based on three points on the crater rims. The crater chronology and production functions proposed by Neukum et al. ( 2001 ) were used to derive the model ages based on the Poisson Possibility analyses method advocated by Michael et al. ( 2016 ).

The thickness of ejecta deposits (i.e., the mixture of ejecta and excavated local material) that were transported from craters outside of the landing region are estimated using the empirical scaling laws estimated by Xie and Zhu ( 2016 ). Local geological context suggests that the Finsen, Von Kármán L, and Von Kármán L' craters have contributed most of the ejecta deposits within the proposed landing region (McGetchin et al., 1973 ; Oberbeck et al., 1975 ; Petro & Pieters, 2006 ; Sharpton, 2014 ). We have a detailed description of the model in section 4.4 .

We estimated the thickness of the regolith within the proposed landing area using NAC images that have incidence angles less than 55° following the method described in Quaide and Oberbeck ( 1968 ). Quaide and Oberbeck ( 1968 ) found that relatively fresh concentric craters with diameters less than 250 m could be used to estimate the thickness of the regolith with the equation: thickness = ( k − D F / D A ) * D A * tan ( α )/2, where D A is the rim‐to‐rim diameter of a crater, D F is the diameter of the inner concentric ring, k is an empirically constant, and α is the angle of repose of materials on the surface of the Moon. The angle of repose of lunar regolith ( α ) is 31°, so the corresponding k is 0.86 and the corresponding slope of inner walls of fresh craters is 31° ± 2° (Quaide & Oberbeck, 1968 ). It is considered more robust to identify concentric craters using images with smaller incidence angle; otherwise, the flat bottomed craters and concentric craters could be misrecognized as normal craters and flat‐bottomed craters, respectively (Fa et al., 2014 ). Thus, we conservatively chose these LROC NAC images with incidence angles less than 55°. Only relatively well‐preserved craters with diameters less than 250 m were used to estimate the regolith thickness (Oberbeck & Quaide, 1967 ) due to the difficulties in diameter identification and measurements using degraded craters (Soderblom, 1970 ).

We analyzed the local geology using the Lunar Reconnaissance Orbiter Camera (LROC) Wide Angle Camera (WAC) mosaic (100 m/pixel) for regional context (Robinson et al., 2010 ), Kaguya Multiband Imager (MI) 750 nm reflectance mosaic (14 m/pixel) for albedo variations (Ohtake et al., 2008 ), Terrain Camera (TC) morning mosaic (7 m/pixel) for local context (Kato et al., 2010 ), and LROC Narrow Angle Camera (NAC) images (0.5–1.6 m/pixel) for small geological feature identification. Topographic analyses were performed using the merged digital elevation model (DEM) that is derived from the LRO Lunar Orbiter Laser Altimeter (LOLA; Smith et al., 2010 ) and Kaguya TC data (SLDEM: 59 m/pix; Barker et al., 2016 ). We produced a slope map for the landing region using SLDEM at a 59 m scale.

Absolute model ages derived from crater statistics for the mare unit located within the floor of Von Kármán crater. (a) Counting areas on the crater floor based on Kaguya TC mosaic. The mare units on the crater floor were classified as a single geological unit based on recent geological mapping (Yingst et al.,). The blue area shows the counting area where only craters larger than 2 km in diameter are used in deriving the model age. The locations of the three subareas are shown in yellow polygons. These regions were selected as they are less affected by obvious secondaries in comparison to the rest of the crater floor. (b) Model age for the mare unit derived from probability analysis (Michael et al.,). (c) Detailed surface morphology of the subarea 1 shown in Kaguya TC mosaic. This area is located within the selected landing area, where the population of obvious secondary craters is smaller. Degraded secondary craters within this subarea are less than 400 m in diameter. (d) Model age for the subarea 1 on the mare surface derived from probability analysis (Michael et al.,). (e) Context of subarea 2 shown by Kaguya TC mosaic and the location of the counting area (blue polygon). (f) Model ages derived for the crater population in subarea 2. (g) Context of subarea 3 shown by Kaguya TC mosaic and the counting area (blue polygon). (h) Model ages derived for the crater population in subarea 3. North is to the top in (a), (c), (e), and (g).

Obvious secondary crater chains have dominated the entire mare unit on the floor of Von Kármán (Figure 3 a). Most of the secondaries are less than 2 km in diameter (section 3.2 ). Therefore, we studied the size‐frequency distribution of craters larger than 2 km in diameter to estimate the model age of the mare unit. The absolute model age derived by probability analyses (Michael et al., 2016 ) is 3.6 (+0.09, −0.2) Ga, which falls in the Imbrian period. This result is consistent with the recent geological mapping (Yingst et al., 2017 ), and crater statistics (Haruyama et al., 2009 ). We also selected three subregions on the mare units that have been less affected by large secondary craters in order to verify the model age. Figure 8 a shows the locations of the counting areas. The surface morphology of the subregions shows that craters larger than ~300–400 m are not obvious secondary craters (e.g., Figures 8 c, 8 e, and 8 g), and model ages derived from the crater counts are identical with those estimated from craters larger than 2 km in diameter.

We analyzed relatively fresh concentric impact craters using LROC NAC images with incidence angle less than 55° to estimate the regolith thickness using the method of Quaide and Oberbeck ( 1968 ). These images have solar angles greater than 35°, which are larger than the repose angle of the regolith (31°). Therefore, we can clearly determine the morphology of the impact craters with these NAC images (Fa et al., 2014 ). The estimated thickness of the regolith in this area varies from ~2.5 to 7.5 m (Figure 7 ). It appears that the regolith in the northeastern portion of the region is thicker than that of the southwestern portion, which is consistent with a larger contribution of ejecta deposits (Figure 10 ), as well as a strong gardening effect of secondary crater chains (Figure 2 a) formed as a result of the Finsen crater‐forming event.

(a) M750 nm reflectance mosaic of Von Kármán crater (stretch: 0.044–0.082). High albedo linear features oriented WSW‐ENE are visible on the crater floor, and they can be traced back to Finsen crater in Figure 1 b. The location of this subset and (b) is indicated in Figure 1 b. (b) Mcolor composite of Von Kármán crater (same region as a). R = 2 μm band center (2,000–2,500); G = 2 μm band depth (0.04–0.17); B = reflectance at 1580 nm (0.085–0.15). Stretch values are indicated in brackets. LCP‐bearing material appears in light blue, and HCP‐bearing material appears green. Spectral parameters were calculated using the method described in Martinot et al. (). (c) MI 750 nm reflectance mosaic of the landing area. Relatively high albedo features are mostly associated with ejecta deposits of small fresh craters, and impact rays from the Finsen crater, whereas low albedo features are mostly associated with mare basalts and ejecta deposits of large craters. The landing area is indicated as a black rectangle, which is indicated in Figure 1 d. (d) RGB composite of MI data. R: 750 nm/415 nm (1.797–1.925), G: 950 nm/750 nm (0.877–1.034), B: 415 nm/750 nm (0.515–0.564). Stretch values are indicated in brackets. The landing area is indicated as a black rectangle, which is indicated in Figure 1 d. (e) Continuum‐removed Mspectra of locations (1 to 7) in Figure 5 d. Each spectrum corresponds to a single pixel. Spectra were processed with the method described in Martinot et al. (). A pigeonite (LCP) and an augite (HCP) spectrum from the RELAB database are displayed above the observed spectra for comparison (respective RELAB‐ID: DL‐CMP‐008 and AG‐TJM‐010). The shaded areas represent the diversity of values of the 1 and 2 μm bands.

Variations in composition are identified in the reflectance spectra of the proposed landing area (Figure 5 ). Spectrally, the entire region is dominated by pyroxene signatures (Figures 5 b, 5 d, and 5 e). Pyroxene reflectance spectra are characterized by the presence of diagnostic, broad absorption bands located around 1 and 2 μm, shifting toward longer wavelengths with increasing iron or calcium content (e.g., Klima et al., 2007 ). The 1 and 2 μm absorption band center positions, displayed on the parameter maps, show that the mare itself is rather homogenous (Figure 5 b). Minor variations of the 1 and 2 μm absorption band center positions of pyroxene spectra are observed within impact crater ejecta located on the mare floor, suggesting variations in chemistry with depth (Figure 5 b). On the MI color composite (Figure 5 d), relatively fresh small craters (diameter ~66 to 324 m; Figure 6 a) show blue‐toned ejecta (e.g., site 1 in Figure 5 d) and higher albedo in the MI 750 nm reflectance data (Figure 5 c). Their reflectance spectra are consistent with that of Low‐Calcium Pyroxene (LCP)‐bearing material (Figure 5 e). The orange‐toned ejecta (e.g., site 2 in Figure 5 d) are related to larger craters (diameter 252–950 m; Figure 6 b), with spectra consistent with Higher‐Calcium Pyroxene (HCP)‐bearing materials (Figure 5 e). The spectra of the ejecta of Ba Jie crater are similar to the spectra of the orange‐toned ejecta (sites 4, 5, 6, and 7 in Figure 5 d). The ejecta on the rim of Ba Jie crater presents HCP signatures, but with larger spectral contrast (site 5 in Figure 5 d). The mare itself (site 3 in Figure 5 d) has a spectral signature consistent with these HCP‐rich materials with weaker absorption bands.

(a) LRO NAC image (M1100060445RE) of a region west of the landing region. The red arrows indicate the position of asymmetric sinuous ridges. (b) LRO NAC image (M1183658592RE) of an area in the western part of the landing region. The blue arrows indicate the position of sinuous troughs. (c) The spatial relationships between sinuous ridges and sinuous troughs observed in the vicinity Ba Jie crater. The red lines indicate sinuous ridges, and the blue lines indicate sinuous troughs. North is toward the top in all the panels.

Notably, we have identified extensive sinuous ridges within the landing region, which are different from wrinkle ridges (e.g., Binder, 1982 ) and degraded secondary crater chains (e.g., Lucchitta, 1977 ) on the Moon. The ridges are asymmetric in shape and extend in a sinusoidal‐like shape (Figure 4 a): the widths of the ridges are tens of meters. Sinuous troughs are formed between the ridges (Figure 4 b), and both the troughs and ridges are spatially associated with Ba Jie crater, which is to the west of the landing region (Figure 4 c).

(a) TC morning mosaic shows secondary craters occurrence all over the landing region (indicated by the red box). The red box is the same location of the white box in Figure 1 d. (b) NE‐SW trending linear features formed by Finsen, which are interpreted as highly degraded secondary craters with diameters >500 m. (c) N‐S trending secondary craters (> 1 km) overlapping the previous NE‐SW trending linear features formed by Finsen. The source crater is most likely Antoniadi crater. (d) NE‐SW relatively fresh secondary craters superposing the NE‐SW trending linear features. The source craters are most likely the Von Kármán L and Von Kármán L' craters. North is up in all the panels. The locations of (b)–(d) are indicated in Figure 1 d.

Obvious secondary craters are widespread within the entire landing region (Figure 3 a). With irregular planar morphology, secondaries are recognized by their spatial occurrences in chains and/or clusters that exhibit herringbone‐shaped morphology. Secondaries within a given chain and cluster have similar preservation states and their elevated crater rims all point in the same direction. Based on the preferential spatial orientation of secondaries within the landing region (Figures 1 d and 3 ), we have identified at least four sets of various‐sized secondary craters that have different preservation states. The NE‐SW trending secondary craters (Figure 3 a) are larger than 500 m in diameter, and they converge toward the Finsen crater (Figures 1 b and 3 a). These secondaries are heavily degraded because their rims now occur as subparallel ridges (Figure 3 b) and the original shallow cavity has been gradually filled by mass wasting and subsequent ejecta deposition. The second set of secondaries generally trend north to south and are larger than 1 km in diameter. The secondaries overlap those formed by Finsen and have a better preservation state (Figures 3 c and S1 ). The steepened crater walls are located at the southern part of the secondaries, suggesting that the source crater is located to the south of Von Kármán. The minimum diameter of the parent crater should be at least 20 km considering that the maximum ratio between continuous secondaries and primaries is ~5% on terrestrial planets (Melosh, 1989 ) and that distant secondaries are much smaller than those on continuous secondaries facies (Xiao, 2016 ). Tracing southward along a great ellipse circle in ArcMap, we noticed that the Antoniadi crater ( D = 138 km; central coordinates as 69.3°S, 186.9°E) is the most likely source crater that fulfills the above criteria (Figure S1 ). Judging by the same criteria, the freshest secondaries within the proposed landing region are the east‐northeast and south‐southwest trending secondaries that are located on top of the secondaries formed by Finsen and Antoniadi. These secondaries (e.g., Figure 3 d) are much smaller, and they are ~250–500 m in diameter (Figure 3 d). Tracing along the secondaries, we find that the Von Kármán L and Von Kármán L' craters to the south of the landing region are the possible source craters (Figure S2 ).

(a) Colorized SLDEM topography overlain in transparency over the TC morning mosaic of the landing region. The black arrow points to the mound with the highest altitude in this region. The white arrows indicate NE‐SW linear features with elevated topography. (b) Slope map of the landing region calculated with a baseline of 59 m. The location of (a) and (b) is indicated in Figure 1 d.

The selected landing region is located on the mare units within Von Kármán crater. The average elevation of this area is about −5926 m, with a standard deviation of 20.4 m. The elevation ranges over about 321 m. The highest geological feature in the region is the mound located near the north boundary (Figure 2 a). The proposed landing region is generally flat at a scale length of 59 m, as nearly 99% of the area has a slope less than 15°, and the slopes of about 94% of the area are less than 5° at a 59 m scale length (Figure 2 b). Local steep slopes (>15°) are mostly associated with craters larger than ~1 km diameter (Figure 2 b). The northeastern and southwestern portions of the landing area are lower in elevation compared with the northwestern and southeastern portions (Figure 2 a). NE‐SW linear features with elevated topography are visible across the landing region (Figure 2 ) and correspond to the ejecta deposits and secondary craters that originated from the Finsen crater (Figure 1 b).

4 Discussion

4.1 Context of CE‐4 Mission CE‐4 is scheduled to launch in 2018 and will be the first lunar farside in situ exploration mission. The CE‐4 mission will be carried out in two phases. First, a relay satellite with two microsatellites will be launched by a CZ‐4C rocket from Xichang, China. The relay satellite will be sent to the Earth‐Moon Lagrange Point 2. A Dutch‐made low‐frequency radio spectrometer (0.1–80 MHz) is carried by the relay satellite to perform space physics measurements. In addition, the relay satellite will carry several laser reflectors for assisting orbital determination. This mission will also include two microsatellites that will orbit the Moon, and they will be equipped with very‐long‐baseline interferometry instruments to conduct radio science experiments. One of the microsatellites will have a visible wavelength microcamera contributed from Saudi Arabia. Six months after the launch, the second part of the CE‐4 mission, which is composed of a lander and a rover, will be launched by a CZ‐3B rocket launched from Xichang, China. Since both the lander and the rover were designed as a backup for the Chang'E‐3 (CE‐3) mission, some of the science payloads on CE‐4 are similar to those on CE‐3 (Jia et al., 2018), which include a landing camera, a terrain camera, a panorama camera on the lander and a visible/near infrared imaging spectrometer (He et al., 2014), and two ground penetrating radars (Fang et al., 2014) on the rover. Additional instruments on the lander (Jia et al., 2018) include (1) a low‐frequency radio spectrometer (0.1–40 MHz) to perform joint space physics observations with the low‐frequency radio spectrometer on the relay satellite, (2) a German lunar neutron and radiation dose detector to explore the farside surface radioactive environment, and (3) a lunar microecosystem for astrobiology experiments and public outreach. A new instrument on the rover (Jia et al., 2018) is the Swedish neutral atom detector designed to study the interaction between the solar wind and lunar surface materials. However, the CE‐4 will not be equipped with the alpha particle X‐ray spectrometer that was previously used by CE‐3 to detect elemental abundances within surface material (Jia et al., 2018; Wu et al., 2017). Still the lander and rover together should be able to perform imaging, spectral, radar, and low‐frequency radio spectral measurements.

4.2 Geological Features Several features of interest have been identified within the landing region, including farside mare basalts, sinuous linear features, and ejecta rays. Nonnearside mare basalts will be investigated in situ for the first time at Von Kármán and may bring new clues about the farside and SPA volcanism. Spectral data show HCP signatures associated with the mare unit, which is consistent with most lunar mare (e.g., Staid et al., 2011). However, in contrast to the CE‐3 landing site, olivine has not yet been detected at Von Kármán (Ling et al., 2015; Zhang et al., 2015), which suggests that the farside mare might be of slightly different composition. The mare unit is homogeneous and likely to represent a single eruptive episode during the lunar peak volcanic period that occurred in the Late Imbrian (3.80–3.20 Ga; e.g., Yingst & Head, 1999). A single eruptive episode was one in which the mechanism of emplacement did not significantly vary over the period of activity, meaning that the rock unit left in place should share generally similar morphological and compositional characteristics, making it a viable geologic unit. Sinuous linear features, ridges and troughs, have been identified in the vicinity of Ba Jie crater. They are distinct from lunar wrinkle ridges, but their origin remains controversial. For example, Oberbeck (1975) proposed these to be ejecta deposits typical of small craters in which the ejecta is emplaced at relatively low velocity, and Atwood‐Stone et al. (2016) suggested that these structures are likely the results of Kelvin‐Helmholtz instabilities within the ejecta flows. The key information to understand their formation mechanism is the subsurface structure of these features, that is, the depth and structural disturbance of the ejecta deposits of the Ba Jie crater, and whether or not fractures deeper than Ba Jie's ejecta exist. The ground penetrating radars onboard CE‐4 could reveal the subsurface structure of these linear features and provide clues to their possible formation mechanism. The secondary craters produced by the Finsen crater‐forming event, and associated with relatively high albedo linear features, are heavily degraded but also of interest (Figures 1b and 3). Hawke et al. (2004) proposed four different mechanisms to explain the formation of impact rays: (1) immature primary ejecta emplacement, (2) secondary craters immature local material deposition, (3) the action of debris surges downrange of secondary clusters, and (4) immature interior. Reflectance spectra for the landing region show that the rays formed by Finsen appear to have distinct composition compared with the buried mare basalts (Figure 5b), indicating that the rays are composed of primary ejecta deposits. With the aid of the CE‐4 cameras and radar system, the thickness and spatial distribution of impact ejecta from the Finsen crater could be more well constrained. This would be the first in situ constraint for the subsurface structure of compositional rays, which will serve as an observational basis for understanding the efficiency of material transport by impact cratering on the Moon.

4.3 Stratigraphy of the Landing Region Impact craters are probes of local stratigraphy (Melosh, 1989). Variations of reflectance spectra of crater ejecta deposits indicate vertical variations in composition/mineralogy in the landing area (section 3.3). We used the geometric relationship between the diameter of impact craters and the depth from which the ejecta was excavated to reconstruct the regional stratigraphy (Figure 9). The maximum depth of excavation is approximately 1/10 of the transient crater diameter, which equals to 0.84 times final crater rim‐to‐rim diameter for simple craters (Melosh, 1989). Therefore, we used this relationship for the diameter measured on the image to calculate the excavation depth of each of the simple craters located upon the mare unit. Note that most of the small craters less than 2 km in diameter are probably secondary craters; thus, they have smaller excavation depths than similar‐sized primary craters (McEwen & Bierhaus, 2006; Oberbeck, 1975), indicating that the calculated excavation depths are the upper bound for the actual values. Figure 9 Open in figure viewer PowerPoint Proposed stratigraphy of the landing region. The thickness of the layers is not to scale. The regolith constitutes the uppermost layer of the reconstructed stratigraphic column (A in Figure 9), with a thickness of ~2.5 to 7.5 m (the regolith in the northeastern part is thicker than in the southwestern part in the landing region; Figure 7). The main uncertainties in the regolith thickness estimates come from LROC NAC images with the necessary illumination geometry (incidence angle less than 55°), the determination of rim‐to‐rim diameter of craters, and the limitations of the estimation method (Quaide & Oberbeck, 1968). Beneath the regolith is the Finsen LCP‐bearing ejecta layer (B in Figure 9; which might be discontinuous or variable in thickness). Represented in blue tones in the MI color composite (Figure 5d), this layer is excavated by craters ~66 to 336 m in diameter (Figure 6a). The majority of these craters is 96 to 156 m in diameter (Figure 6a), suggesting that the LCP‐bearing material is at least ~8 to 13 m deep. The HCP‐bearing layer (D + E + F in Figure 9) is exposed in the orange‐toned ejecta in the MI color composite (Figure 5d) of larger craters with diameters ranging from 268 to 988 m (Figure 6b) as well as Ba Jie crater. The majority of orange toned ejecta craters fall in the 388 to 628 m range in diameter (Figure 6b), indicating that the HCP‐bearing materials (D) are at least ~33 to 53 m deep. Ba Jie crater is ~3.7 km in diameter, suggesting that local minimal depth of the HCP‐bearing materials (and hence mare unit) is greater than 310 m. The spectrum extracted from the site 5 of Ba Jie crater's ejecta has a deeper absorption band than the spectra of sites 6 and 7 (Figure 5e). This absorption band depth difference could indicate that the material of site 5 is slightly distinct compared to material of sites 6 and 7, possibly more enriched in HCP or with a different grain size or texture (or less mature). Alternatively, the ejecta could be thicker at site 5 and less mixed with the underlying likely space‐weathered layer, resulting in a more intense signature. Spectral variations within the ejecta of Ba Jie crater (Figures 5d and 5e) could hint at a subtle vertical compositional difference in the layer of HCP‐bearing materials. The material located at site 5 (layer F) are from deeper portions of the preimpact stratigraphy than the material located at sites 6 and 7, due to the fact that the deeper‐seated material tends to be ejected closer to the crater rim (Stöffler et al., 1975). Therefore, there is probably a layer of enriched HCP‐bearing material (F) under the HCP‐bearing‐material layer, and could imply a possible paleo‐regolith layer (E) between layers D and F if there were at least two episodes of basalt emplacement. The paleo‐regolith thickness (if such a layer exists) could be studied with the radar instrument onboard the CE‐4 rover, similarly to the detections made at the CE‐3 landing site (Xiao et al., 2015). It appears reasonable to speculate that there could be somewhere in the stratigraphy a layer of mixed LCP and HCP‐bearing material (C) due to collision of these two types of materials. However, we are not able to constrain the thickness of this mixed layer. The main uncertainties of the reconstructed stratigraphy come from ejecta of unidentified impact craters, mixing of ejecta of local materials, and products of uncertain geological events between stratigraphic layers. We propose a stratigraphic column with several layers beneath the HCP‐bearing‐material layer/mare unit based on the regional setting and previous geological maps (Wilhelms et al., 1979; Yingst et al., 2017). However, the thickness of these layers cannot be constrained using available data and previous mapping results. We expect a layer G made of ejected material from Imbrian‐aged craters (e.g., Alder) resurfacing HCP‐bearing basalts. Layer B to G were formed during the Imbrian epoch. Ejecta of Leibnitz crater (H) occurred as the layer beneath, and this impact event occurred in Nectarian. Then there should be a layer (I) of breccia from Von Kármán crater forming event lying above the target materials (J) of Von Kármán crater. Layer I and J are of pre‐Nectarian age (Figure 9), likely to be part of the SPA basin Mg‐Pyroxene Annulus (Moriarty & Pieters, 2018). The two ground penetrating radars onboard the CE‐4 rover will be able to reveal the subsurface structure of the landing area and test the stratigraphy predicted in this study. In a manner similar to the CE‐3 Yutu rover, the radar system of the CE‐4 rover has two frequency channels with different penetrating depths and vertical resolutions: Channel 1 has a frequency of 40–80 MHz, whereas Channel 2 has a frequency of 250–750 MHz (Jia et al., 2018). The radar system on the CE‐3 mission demonstrated that the Channel 2 radar could detect details of the subsurface structures up to a depth of ~12 m, and the Channel 1 radar could reveal subsurface structures up to ~400 m (Xiao et al., 2015). Therefore, the ground penetrating radars Channel 1 can detect layers A, B, C, D, E, and F, and it could detect layer G, H, I, and J depending on their thickness. Channel 2 could detect detailed structures within layers A and B, and the upper portion of the layer C.