Earth and Planetary Science Letters 294 (2010) 440–450
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Earth and Planetary Science Letters
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Mineralogy of recent volcanic plains in the Tharsis region, Mars, and implications for platy-ridged flow composition
N. Mangold a,b,⁎, D. Loizeau b,c, F. Poulet c, V. Ansan a, D. Baratoux d, S. LeMouelic a, J.M. Bardintzeff b, B. Platevoet b, M. Toplis d, P. Pinet d, Ph. Masson b, J.P. Bibring c, B. Gondet c, Y. Langevin c, G. Neukum e a Lab. Planétologie et Géodynamique de Nantes, CNRS/Univ. Nantes, France b Interactions et Dynamique des Environnements de Surface, UMR8148, CNRS/Univ. Paris Sud, France c Institut d'Astrophysique Spatiale, CNRS/Univ. Paris Sud, Orsay, France d Lab. Dynamique Terrestre et Planétaire, OMP, CNRS/Univ. Toulouse, France e Institut of Geosciences, Freie Universität, Berlin, Germany article info abstract
Article history: Volcanism on Mars occurred until recently, but the mineralogy of recent lava plains is poorly known because Accepted 28 July 2009 few regions display fresh outcrops devoid of dust. Using visible and near infrared data of the Mars Express Available online 26 August 2009 probe, two new volcanic plains in Noctis Labyrinthus have been identified, and the existence of a volcanic Editor: T. Spohn plain on the floor of Echus Chasma has been confirmed. Crater retention ages estimated for these three plains range between 50 and 100 My, corresponding to the Late Amazonian. These plains represent an excellent Keywords: opportunity to constrain the mineralogy of recent volcanic rocks. Results show that basaltic compositions Mars with plagioclase and high calcium pyroxene are predominant. The low olivine proportion suggests that the volcanism apparent fluidity of these flat plains is not related to magmas being ultramafic. In addition, a platy-ridged fl platy ows texture is observed in two of the studied regions. Our study shows, for the first time, that this texture is pyroxenes associated with volcanic rocks, and that these rocks are of typical basaltic mineralogy. Finally, these volcanic plains are located more than 1000 km east of previously known Late Amazonian volcanic centers of the Tharsis region, an observation to be taken into account when considering models of recent volcanism on Mars. © 2009 Elsevier B.V. All rights reserved.
1. Introduction plains, of which two are identified here for the first time using HRSC imagery. Volcanism on Mars is diverse and well developed, with evidence of High resolution imagery and topography can be combined with abundant plain-style volcanism and a variety of edifices. At the scale visible and near infrared spectral data to study both morphology and of Viking images, early investigators observed that the most recent composition of volcanic rocks on Mars (e.g., on Syrtis Major, Poulet et al., volcanism was located in distinct areas such as the Tharsis (including 2003). However, few spectral data exist for areas of recent volcanism, Olympus Mons) and Elysium regions (e.g., Greeley and Spudis, 1981). because such surfaces are often featureless due to widespread dust cover At that time, it was unclear if volcanic activity extended into the Late (e.g., Hamilton et al., 2003, Stockstill-Cahill et al., 2008). As described in Amazonian (b500 My). A better picture of recent volcanism is now this manuscript, three small regions (b50 km wide) were found to have given by the Mars Observer Camera (MOC) onboard Mars Global formed recently (b100 My). Two of these areas were found in Noctis Surveyor (Hartmann and Neukum, 2001; Berman and Hartmann, Labyrinthus, and one was found on the floor of Echus Chasma (Fig. 1). 2002) and by the High Resolution Stereo Camera (HRSC) onboard The resolution of OMEGA (Observatoire pour la Minéralogie, l'Eau, les Mars Express (Neukum et al., 2004; Chapman et al., 2007; Vaucher Glaces et l'Activité) spectral data provides the possibility to analyze the et al., 2009; Werner, 2009) (e.g., Arsia and Olympus Montes). Ages as composition of these recent plains in detail (Bibring et al., 2004). recent as several tens of millions of years have also been deduced for Plain-style volcanism, corresponding to broad lava flows without the Central Elysium Planitia, southeast of Elysium Mons (e.g., Vaucher obvious volcanic vents, is widespread on Mars (e.g., Hodges and et al., 2009). These young ages are of interest because they show that Moore, 1994). However, few mineralogical data exist with which to volcanism on Mars may still be active. This study focuses on three constrain their composition independently from rheological evidence suggesting a high fluidity. Additionally, platy-ridged flows that are common in Martian plains are the subject of debate due to their ⁎ Corresponding author. Lab. Planétologie et Géodynamique de Nantes, CNRS et Univ. Nantes, 2 rue de la Houssinière, 44322 NANTES Cedex, France. morphologic similarity with potential ice rafts, suggesting frozen seas E-mail address: [email protected] (N. Mangold). in the southern Elysium plains (Murray et al., 2005). Two of the three
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HRSC DEMs (Digital Elevation Models) were computed using the photogrammetric software developed at the DLR and the Technical University of Berlin (Scholten et al., 2005; Gwinner et al., 2005, Ansan et al., 2008). The image correlation was performed using a matching process at a different spatial grid size (Scholten et al., 2005). The height was calculated taking into account the Martian geoid, defined as the topographic reference for Martian heights, i.e. the areoid (Smith et al., 1999). In the Noctis Labyrinthus region, the DEMs were constructed using spatial gridding of ~30 m/pixel for orbits #1977 and #1999, and 100 m for orbit #2479. In the Echus Chasma region, only the MOLA (Mars Observer Laser Altimeter) DEM at 1/128 degree was used for the topography.
2.2. OMEGA/Mars Express data
OMEGA is a visible and near infrared (VNIR) hyperspectral imager providing three-dimensional data cubes with spatial samplings from a few kilometers to 300 m. For each pixel it provides spectra between 0.35 and 5.1 µm, using 352 contiguous spectral elements (spectels), 7–20 nm wide. The spectrometer consists of three detectors (from 0.35 to 1 µm, from 0.9 to 2.7 µm, and from 2.5 to 5.1 µm) (Bibring et al., 2004). For this study we used data recorded by the second detector where the sig- natures of the mafic minerals are well characterized. Pyroxene is detected from the presence of a broad band at 1.9 and 2.3 µm, res- pectively, for low-calcium pyroxene (LCP) and high calcium pyroxene (HCP). In our study, the detection of pyroxene is done with the spectral index used by Poulet et al. (2007) sensitive to the presence of both HCP and LCP. To account for instrumental and atmospheric biases, the detection is considered positive for values of the pyroxene spectral parameter larger than 1% (Poulet et al., 2007). The detection of olivine is based on the broad 1 µm band from the spectral parameter defined in Mustard et al. (2005), but no olivine has been detected by this method on the region studied. The depths of absorption bands were determined and mapped, allowing the presence of minerals to be identified in the top few micrometers of the Martian surface. Fig. 1. (top) Context map: MOLA shaded relief of the Tharsis plateau. (bottom) TES The identification of minerals from band depths does not, however, thermal inertia of the area in a extracted from Putzig et al. (2005). The three studied areas, NL1, NL2, and ECF, display inertia N300 J m− 2 K− 1 s− 1/2 (in yellow), while most provide a direct determination of all minerals. In addition, the spectral of the volcanoes and lava plains are buried beneath low inertia material (b200 J m− 2 parameter used for pyroxene is not designed to discriminate between − − K 1 s 1/2). the LCP end-member and the HCP end-member. Modeling of NIR spectra has been demonstrated to provide an accurate estimate for mineral abundances in mixtures of basaltic granular materials for a plains studied showed platy-ridged landforms with fresh outcrops wide range of particle sizes (Poulet and Erard, 2004). As described by enabling us to extract their mineralogy. Thus, our study focuses on the Poulet et al. (2009a), scattering models may be used to explore the morphology, mineralogy of recent volcanic plains that represent an parameter space and the influence of grain size of mineralogical excellent opportunity to constrain the composition of recent volcanic members, including spectrally neutral components. The model has rocks and understand their apparent fluidity, in order to provide two free parameters for each end-member: the average grain size and insights for the origin of recent volcanism on Mars. the relative abundance. Since aerosols influence spectral slopes, one additional free parameter is used to adjust the spectral slope con- fi – 2. Methods tinuum. The spectra are tted in the 0.99 2.49 µm wavelength range (OMEGA SWIR-C channel) using a simplex minimization algorithm. 2.1. HRSC/Mars Express data The synthetic spectrum is a nonlinear combination of the optical indices of the minerals selected to be part of the mixture, in proportion fi The HRSC camera acquired images in five panchromatic channels to their abundance. Upon obtaining a data spectrum t, the algorithm under different observation angles, as well as four color channels at a supplies the user with a model-derived spectrum as well as with a fi relatively high spatial resolution (Neukum and Jaumann, 2004). In percentage and grain size for each end-member used in the t. The our work we used only panchromatic nadir images, with a maximum principal mineralogical characteristics of the three areas studied were spatial resolution from 10 to 40 m pixel− 1, and two panchromatic determined using this method and are summarized in Table 1. stereoscopic images with a spatial resolution usually degraded compared to the nadir image. The coordinates of ortho-rectified Table 1 nadir images are defined in the planetocentric system of the Mars OMEGA derived modal mineralogy. IAU 2000 ellipsoid (Seidelmann et al., 2002). The following images HCP LCP LCP/HCP Neutral Olivine Dust RMS have been processed: #1977, #1999, #2402, #2479, and #3155 in Echus Chasma 25±6 6±3 23±13 47±10 b5 17±13 0.26 the Noctis region, with a spatial sampling of the nadir image at 15 m Noctis Labyrinthus 1 34±5 5±3 17±10 60±10 b5 b1 0.39 − − pixel 1 for #1977, #1999, and #3155, and 40 m pixel 1 for the Noctis Labyrinthus 2 30±5 8±3 29±15 48±7 b5 12±9 0.39 #2402 and #2479. Images of orbit #71 and #5379 are used in the Values are percentage (including the LCP/HCP ratio) and standard errors for − 1 Echus Chasma region with a spatial sampling at 30 m pixel . abundances indicate ±1σ. Neutral end-member corresponds to plagioclases. 442 N. Mangold et al. / Earth and Planetary Science Letters 294 (2010) 440–450
Fig. 2. OMEGA spectral mapping of pyroxenes in Noctis Labyrinthus. A map of HCP superimposed on a THEMIS daytime mosaic of the Noctis Labyrinthus region. Squares represent the two regions of interest NL1 (west) and NL2 (east). Colors represent band depths from 1% (dark blue) to 5% (red) for both maps. Mars Express orbits used are: #305, #331, #357, #431, #453, #486, #519, #1052, #1107, #1933, #1944, #1955, #2336, #2347, #2369, #2380, #2402, #2479, #3078, #4219, #4230, #4241, #4545, and #4556.
2.3. Other datasets Syria Planum is a volcanic plain with many small low shield volcanoes of Hesperian age (Baptista et al., 2008). To the east, West Mars Observer Camera (MOC) images are used to complement HRSC Tithonium Chasma is covered extensively by sand sheets and dunes nadir images at a better resolution, typically 2.8 m pixel−1 (Malin et al., that bury the canyon, and part of the Oudemans crater (the 100 km 1992). Nighttime and daytime infrared images of the Thermal Emission diameter crater to the southeast of the canyon). Therefore, none of Imaging System (THEMIS) onboard Mars Odyssey were used to provide these signatures have been found to be correlated with young a view of the surface properties (Christensen et al., 2003). Only the volcanic rocks. However, our interest focused on two spots (NL1 mosaics of thermal images are used at 230 m pixel−1.Wechosetousea and NL2 on Fig. 2) inside Noctis Labyrinthus (Fig. 2), where pyroxene color scale for these images and superposed them onto HRSC or THEMIS signatures appear to be pronounced and the thermal inertia is high. daytime images in order to observe where thermal images may be Spectra of both regions (Fig. 3) fit well with typical laboratory spectra indicative of the presence of indurated rock outcrops or coarse grains (as of pyroxenes, especially HCP, with broad 1 µm and 2.3 µm bands. seen in red for high nighttime temperatures) compared to more dusty, less indurated particles (as seen in blue for low night time tempera- 3.2. West canyon floor (NL1) tures). Values of thermal inertia for the regions of interest have been extracted from the TES (Thermal Emission Spectrometer) thermal The first canyon floor analyzed, NL1, is located at 99W, 7S (Fig. 4). inertia map at a resolution of 3 km (Putzig et al., 2005). It is about 50 km from west to east, 60 km from north to south, and 5 km deep below plateau level. The HRSC image shows a smooth floor 3. The floor of Noctis Labyrinthus Chasmata only interrupted by small buttes. The albedo of the floor is very low compared to walls and plateaus (Fig. 4a), and low albedo correlates 3.1. Regional context and the distribution of pyroxenes as seen with the high night time temperatures seen on the floor (Fig. 4d), by OMEGA corresponding to areas with more than 500 J m− 2 K− 1 s−1/2 on the TES thermal inertia map (Putzig et al., 2005)(Fig. 1b). Only the Pyroxene signatures are present throughout all of the southern uppermost part of the canyon walls (rock scarps) display nighttime part of Tharsis and in Valles Marineris (Fig. 2). Hence, these regions temperatures as high as the canyon floor (Fig. 4d). Pyroxene are devoid of dust as seen on the thermal inertia map (Fig. 1). signatures detected by OMEGA are found only on the canyon floor. However, most pyroxene signatures are observed either on dark sand The topography of the floor measured by MOLA data (Fig. 4b,c) dunes filling canyons floor, such as in Candor Chasma (Mangold et al., indicates that the floor is at a constant altitude of ~2250 m with 2008), or on canyon walls (e.g. Mustard et al., 2005). The source of limited variation (b50 m). All of the pyroxene signatures (Fig. 4e) are material is unknown for canyon floors. For canyon walls, rocks detected beneath this elevation contour. correspond to thick accumulations of lava flows of Hesperian or A closer look to the east of this canyon is possible with more Noachian age, which are beyond the scope of our study. By examining resolved HRSC data from orbit 1999 (Fig. 5). Here, the HRSC DEM the region of Noctis Labyrinthus (Fig. 2), we observe pronounced allows us to observe that all small buttes located on the canyon floors pyroxene signatures in the west of Syria Planum and West Tithonium are located above the 2250 m contour. This contour separates that Chasma. part of the floor with low albedo (0.15 at 1 µm OMEGA reflectance) N. Mangold et al. / Earth and Planetary Science Letters 294 (2010) 440–450 443
relative to surrounding bright terrain of smooth texture, containing pyroxene. These observations confirm that the whole floor is very flat at the elevation of 2200 to 2250 m over the entire canyon floor area. Close-ups of the canyon floor reveal further landforms not visible at a broader scale. First, a circular shape 1 km in diameter and about 100 m in depth is observed beyond the southern edge of the plain (Fig. 6a,b). It is breached in its northern part by a 200 m large, 30 m deep channel which joins the canyon floor 5 km to the north and vanishes inside the canyon floor shortly after intersecting the 2250 m elevation contour (Fig. 6b). In addition, most of the canyon floor is smooth at the HRSC scale with only a few sand dunes present in the central part. Nevertheless, the MOC close-up shows small impact craters and local outcrops, especially around unfilled pits (Fig. 6c), suggesting that sand exists but does not mantle all the surface. These images also display a lineation of NW–SE direction, resembling a buried dyke or a fracture zone, with collapse pits 100 m wide sitting on the fracture (Fig. 6c). Finally, the area located east of the canyon floor exhibits not only small buttes, but a series of pitted cones of about 100 m in diameter, which are observed at the contact of the 2250 m altitude contour (Fig. 6d). All these observations listed above indicate a volcanic origin for the canyon floor. By itself, the pyroxene-rich mineralogy coupled with the low albedo and the strong thermal inertia is good evidence for a volcanic origin (Fig. 4). Additionally, the combination of the deep circular landform together with a channel (Fig. 6) is similar to ter- restrial volcanic vents with lava channels emerging from them (e.g. pit crater with lava channel in the Snake River plain, Idaho, USA, Fig. 5G, in Hodges and Moore, 1994). Pitted fractures inside the canyon floor might correspond to other vents, which were partially filled by the lavas they extracted. The pitted cones to the east might be an Fig. 3. Spectra (average of 9 pixels) of the selected areas NL1, NL2 and ECF. Library additional source of lavas. Alternatively, they might correspond to reflectance spectra are relative reflectance. Library spectra correspond to diopside (HCP) rootless cones, also named pseudocraters. Rootless cones form at the and pigeonite (LCP) from the RELAB database (Reflectance Experiment LABoratory; http:// fl www.planetary.brown.edu/relabdata/). A spectrum from the dust region around NL1 is contact of lava ows and a volatile such as liquid water or water ice shown for comparison. (e.g. Lanagan et al., 2001). Their presence at the contact with the
Fig. 4. Close-up on NL1: (a) HRSC image #2379 inside the NL1 trough. (b) Same image as a with MOLA contours (250 m intervals). (c) MOLA DEM in between elevations 2200 m and 2300 m showing the very flat Chasma floor. (d) The HRSC image with THEMIS nighttime mosaic superimposed in color. Red colors represent high temperatures, blue colors represent low temperatures, and, therefore, finer material than in warmer locations. (e) An OMEGA map of pyroxene superimposed on the THEMIS daytime image. 444 N. Mangold et al. / Earth and Planetary Science Letters 294 (2010) 440–450
so, the fact that the albedo is slightly higher in the western part of the canyon floor suggests that the western part has been slightly mantled by a thin (b10 cm) coating of fresh dust hiding the rock surface at OMEGA wavelengths. An elliptical feature, marked by “E” is visible in the albedo (Fig. 7a), in the topography (Fig. 7c), and in the temperature maps (Fig. 7d). This landform is oriented NW–SE, is a few tens of meters deeper than the rest of the canyon floor, and could have served as a trap for recent dust deposition. A close-up on HRSC and MOC data shows a set of different textures (Fig. 8). First, the southern part exhibits a specific texture that differs from the smooth canyon floor. Here, plains appear as if it has been disrupted by liquefaction, or some kind of fluid activity. Indeed, small pieces of plains are found in the middle of this area and resemble solid plates floating over a fluid medium (Fig. 8a). The plates are still as smooth as the plains material, as seen on the MOC close-up (Fig. 8b). In between these plates, the texture is rough with a lot of fractures that do not cross the plates. Plates are up to 1 km long and are less than 30 m thick because the resolution of the HRSC DEM is insufficient to detect the plates. Outside this area, plains exhibit frequent small ridges (marked by “R” in Fig. 8c) several kilometers long and approximately 100 m wide. The ridges display a positive relief, as if they were formed by compression in soft material, but their curved shape without straight faults (Fig. 8c) does not plead for a tectonic origin by brittle faulting. In Fig. 5. (a) An HRSC close-up of orbit #1999 on the eastern side of the Chasma floor. (b) The HRSC DEM represented in between 2200 m and 2300 m of elevation, which shows Fig. 8c, left of the ridge, curvilinear shapes of material that appear to the flat floor to the west and the presence of hills above the floor to the east. canyon floor is consistent with such a hypothesis. The apparent size of these cones do match terrestrial cones as well as some other Martian examples (Lanagan et al., 2001; Meresse et al., 2008) and may indicate the presence of liquid water or water ice at the time of the volcanic activity. Finally, the constant altitude of the canyon floor suggests that the volcanic material is fluid enough to form a smooth plain which follows an equipotential surface. All this evidence favors formation as a lava plain, covering about 450 km2 if the 2250 m contour is taken as a reference.
3.3. East canyon floor
The second canyon floor analyzed, NL2, is located at 96.2W, 7.2S (Fig. 7). It is about 50 km in length from west to east, 40 km in length from north to south, and 6 km in depth. The HRSC image shows a smooth floor with small buttes in the eastern and southern part. The albedo of the floor is low when compared with the walls and plateaus (0.11 at 1 µm OMEGA reflectance), and displays a gradation through the floor from the west to the east, the eastern region being darker than the western region (Fig. 7a). Low albedo correlates with the high night time temperatures seen on the canyon floor (Fig. 7d), reaching up to 600 J m− 2 K− 1 s− 1/2 on TES data in the darkest areas (Putzig et al., 2005)(Fig. 1b). Here too, only the uppermost part of canyon walls (namely the rock scarps) display temperatures as high as the canyon floor. Pyroxene signatures detected by OMEGA are found mainly on the eastern part of the canyon floor (Fig. 7e). The topography of the floor as measured by HRSC data (Fig. 7b,c) indicates that the floor is constrained between the altitudes of 1850 and 2000 m. The 2000 m contour follows the smooth texture of the canyon floor very well. We observe that the transition between high temperatures (red) and low temperatures (blue to yellow) at the southern and northern edge of the canyon floor correlates with the 2000 m altitude level (Fig. 7d). Dust, as seen from blue colors (low Fig. 6. (a) HRSC orbit #1999 close-up of the southern part of the Chasma. (b) The same night temperatures) and high albedos in the HRSC image, mantles as a with HRSC DEM contours superimposed. This area displays a central pit to the south most of the canyon walls and debris aprons. As dust deposition is a from which emerges a sinuous channel ending in the Chasma floor. (c) MOC image #R1800704 in the central area showing two elliptical pits formed along a fracture zone. gradual process at the regional scale, this observation likely illustrates (d) An HRSC close-up on the eastern margin of the Chasma floor showing a series of fl that dust was deposited before the canyon oor was resurfaced, and small (b500 m) cones in the center, and round hills that emerge from the floor of the that dust was buried beneath the high thermal inertia material. Even Chasma by a few tens of meters at the maximum. N. Mangold et al. / Earth and Planetary Science Letters 294 (2010) 440–450 445
Fig. 7. Close-up on the NL2 trough: (a) HRSC image #1977 inside the NL1 trough. (b) The same as a with HRSC DEM contours (250 m intervals). Note the presence of irregular contours for the 2000 m elevation contour, which is due to the DEM accuracy (b30 m) relative to the flatness of the area. (c) An HRSC DEM in between elevations 1860 and 2080 m showing the very flat Chasma floor with a possible elliptical trough in the middle (identified by E in a). (d) The HRSC image with the THEMIS nighttime mosaic superimposed in color. (e) An OMEGA map of pyroxene superimposed on the THEMIS daytime image. have been deformed in a ductile way can be observed. Small impact spectra. It most likely corresponds to rocks mantled by eolian deposits craters can also be observed in this image. such as bright dust. The window of low albedo rocks therefore signals a In summary, this volcanic plain has a relatively high nighttime location where the spectral composition can be explored. The high temperature with rough textures of bedrock exposure with low albedo (0.19) compared to the Noctis plains (below 0.15) could sug- albedos and pyroxene signatures on fresh outcrops. These character- gest that a thin dust mantle (b100 µm thick) is present, but not in istics argue in favor of a volcanic origin such as lava flood plains. The sufficient proportion to mask the detection of pyroxene. plate shape, size and width are consistent with those frequently The overall topography of the Echus Chasma floor is flat, sug- present in Martian lava flood plains, usually referred to as platy- gesting that a low viscosity fluid filled the trough (Fig. 9c). At the scale ridged lava flows (e.g., Keszthelyi et al., 2000). Other landforms such of MOC images, surface morphology is complex and includes pressure as ridges and curved textures are also consistent with viscous lava ridges, straight fractures, and rafted plates (Fig. 9a,b). Individual flows, despite the fact that this type of ridge does not seem frequent in plates are several tens of meters to several kilometers wide, relatively other Martian flood plains. The total surface of this volcanic plain is smooth, and separated by rougher and brighter terrains. The smooth about 800 km2. A difference between the NL2 and NL1 plains is that plates often seem to fit together like a jigsaw puzzle. Therefore, these there is no obvious vent found in, or along the plain NL2. Nevertheless, plates can be interpreted to be pieces of a ruptured crust that have it is possible that the elliptical trough (E) corresponds to a lava vent rafted over a fluid medium. Pressure ridges are interpreted to exist at that has been entirely buried beneath its flows. the edges of plates, since these plates are rammed into each other during emplacement. These characteristics are typical of platy-ridged 4. Echus Chasma floor flows, as described by Keszthelyi et al. (2000). In summary, this plain displays mineralogical features and morphologic patterns that Echus Chasma is the source area of the Kasei Vallis outflow corresponds to a volcanic filling. The rough texture observed on channels that dissect volcanic plateaus from the equator to 20N, west MOC images, the lack of sand dunes, and the relatively high thermal of the large volcanoes of Tharsis (Fig. 1). The Echus-Kasei valley is inertia favor a rocky bedrock locally exposed without a thick dust partially filled by lava flows originating from the west, from the main mantle. The area of the plains defined by the OMEGA signature is volcanic center (e.g. Scott and Tanaka, 1986). While no vent is approximately 2000 km2. observed in our studied area, volcanic fissures exist on the floor of the valley north of the area studied, suggesting that local emissions from 5. Determination of ages and compositions the Chasma floor are also possible (Chapman et al., 2010-this issue). The studied area is a zone with two types of material, a high 5.1. Ages thermal inertia, low albedo region, in the middle of plains that are generally of low thermal inertia and high albedo (Fig. 9). Pyroxene is Ages are obtained counting craters in the three regions studied detected where the albedo is low (0.19 in 1 µm OMEGA reflectance) according to the diagram and methods described in Hartmann and and the inertia high (Fig. 9d, e). The morphology observed in these Neukum (2001). This diagram divides the distribution range (crater plains shows a lack of sand dunes and a rough surface at MOC scale diameter) into log intervals incremented in square root 2. Using studies suggesting the surface is relatively free of aeolian material (Fig. 9a,b,f). from the Moon, meteoroid distributions and impact frequency, Typical spectra of this region have a well defined pyroxene signature, predicted “isochrons,” or crater size-frequency distributions, have as visible in Fig. 3. In contrast, the high albedo region in the surround- been derived for well-preserved surfaces of various ages, such as 1 Ga, ings consists of low inertia material that appears featureless in OMEGA 100 Ma, 10 Ma, and so on. Determined ages are valid for fresh surfaces 446 N. Mangold et al. / Earth and Planetary Science Letters 294 (2010) 440–450
not affected by erosion or deposition. Crater counts should follow isochrons when the surface is fresh, while crater counts crossing isochrons indicate a surface that has been subsequently modified. Counts were done in a GIS software using sinusoidal projections of HRSC and MOC images for measuring crater diameters and the area of the unit. NL1 crater counts are determined for over 250 km2 on the eastern part of the plain (Fig. 5). We limit counts to HRSC data due to the paucity of images with better resolution in this case. An age following approximately the 100 My isochron can be retrieved graphically. A model age of 77±21 My is determined using the more precise mean square root method proposed by Vaucher et al. (2009). The younger age with this method is due to the fact that smaller craters are more nu- merous, therefore statistically more reliable in the mean square method. NL2 crater counts are determined over an area of 800 km2 from HRSC images, and over an area of 210 km2 from the three MOC images crossing this part of the plain. The plot reveals a good consistency between craters counted in HRSC and MOC images (Fig. 10, circles). The isochron for this plot corresponds to an age of 50 to 100 My. This age is well constrained from the large range of craters used. A precise model age of 66±5 My is determined using the method of Vaucher et al., 2009. Echus Chasma Floor crater counts were derived from the relevant HRSC mosaic over an area of 1600 km2. The crater counts are close to the 100 My isochron (Fig. 10, triangles). The slight shift to a younger age for smaller craters could be related to the statistically low number of large craters, or to erosion/deposition processes that influenced part of the canyon floor, thus removing some small craters. Despite this slight difference in the isochron slope, ages in the range 50–100 My provide reasonable fits. A model age of 52±8 My is determined using the mean root squares method (Vaucher et al., 2009), assuming the smaller bin is not affected by more recent obliteration. This age is consistent with that of about 70 My established using cumulative methods (Chapman et al., 2007, 2010-this issue). Plots of crater density follow isochrons well enough to confirm that the ages correspond to the formation of the plains, or the plots would have crossed isochrons without following them. One should remember that the absolute ages have uncertainties of a factor 3 to 4 (e.g. Hartmann and Neukum, 2001). In addition, the correction factor due to a decreasing impact flux through time might underestimate these ages by a factor of roughly 3 (Quantin et al., 2007). Even with these limitations, the three studied plains are still extremely young relative to the majority of the Martian surface (300 My is an age well into in the Late Amazonian epoch). The estimated ages can be compared to other studies of volcanic centers on Mars. These ages are among the youngest found on Mars for volcanic plains, and are of the same order as those proposed for the calderas of Arsia Mons or Olympus Mons in the Tharsis region (e.g., Neukum et al., 2004) and volcanic plains in Central Elysium Planitia (Vaucher et al., 2009).
5.2. Composition
The mineralogy derived from spectra is shown in Table 1. Plagioclases abundance, inferred by OMEGA is included in a mineral group referred to as “neutral components,” for which the abundance is the combined abundance of spectrally featureless phases in the near infrared (plagioclases, high silica phases and quartz). Plagioclases are expected to be predominant in this group for smooth lava plains. Grain sizes were determined in models from a few tens to a few hundreds of micrometers, except for dust which was below 10 μm. Results indicate that plagioclase, HCP and LCP are the dominant minerals for the three regions listed in Table 1. A contribution of dust Fig. 8. (a) An HRSC close-up (#1977) on the southern part of the smooth plain NL2. is required in two regions. Olivine was not required in the fitof Notice the disruption of the surface in the center relative to the surrounding flat area. spectra. This shows that olivine was below the detection limit (~5% for (b) MOC image #R0800697 of this disrupted area. Three plates N500 m wide are visible grain size of 100 µm, or a slightly higher proportion for grains below with their smooth texture. They are separated by rougher terrains. (c) MOC image #R0601602 of the smooth plain. R represents a sinuous ridge. The terrain in the center 10 µm). The derived abundances of plagioclases and pyroxenes (40 to and in the left part of the image displays fine texture with flow patterns. 60% plagioclases, 30 to 40% pyroxenes) are typical of basaltic N. Mangold et al. / Earth and Planetary Science Letters 294 (2010) 440–450 447
Fig. 9. (a) MOC R0100171 of the Chasma floor. (b) MOC #E1701699 of the Chasma floor. Both close-ups show ridges and plates suggesting a solidified crust floating over a fluid medium. (c) The THEMIS daytime mosaic with MOLA contours superimposed (250 m intervals). (d) The THEMIS daytime mosaic with the THEMIS nighttime mosaic superimposed in color showing warmer locations at night. (e) An OMEGA pyroxene map superimposed on the THEMIS daytime mosaic. (f) An HRSC mosaic (orbit #71 and #5379) of Echus Chasma floor. Mars Express orbits used for OMEGA data processing are #581 and #1257.
composition. The three regions also display a very low LCP/HCP ratio, of 1:1 to 1:4 (Poulet et al., 2009b), whereas NL1 reaches a ratio of 1:6 some of the lowest ratios known to date on Mars. Indeed, most (17% in Table 1). Such characteristics were interpreted to support the Noachian crustal rocks or Hesperian aged plains have LCP/HCP ratios idea that there is a change in the LCP/HCP ratio through Martian 448 N. Mangold et al. / Earth and Planetary Science Letters 294 (2010) 440–450
thus showing that the origin of the volcanism is local, below the canyon floor. As the origin of the volcanism is not the focus of our study, we do not enter into in-depth discussions on that question, but the observed volcanic activity, well to the east of the Tharsis area, should be taken into account in geophysical models attempting to explain recent volcanism on Mars (e.g., Schumacher and Breuer, 2007). Possible explanations include the strong faulting in this region, especially in Noctis Labyrinthus, compared to regions devoid of widespread faulting, and, the thinner crust on the floor of Chasmata than on the plateau (N5 km of difference of elevation). Both of these features may enable magmatic activity to reach the surface.
6.2. The composition of platy-ridged flows
The existence of platy-ridged flows with rafted slabs and pressure ridges is well known on Mars (i.e., Keszthelyi et al., 2000). However, this type of pattern is not frequent for terrestrial volcanism, and has been discussed almost exclusively for flows in Iceland (Keszthelyi et al., 2004, Haack et al., 2006). Therefore, not all workers agree that these features on Mars are lava flows. Mudflows and pack ice (icebergs floating over water) have a very similar shape, with rafted slabs of ice over water or rocks over mud. Such a scenario has been proposed for platy flows seen in Kasei Vallis by Woodworth-Lynas and Guigné (2003) and Williams and Malin (2004). In this respect we note that Fig. 10. Crater counts for the three studied plains presented on Hartmann isochrons. “n” Kasei Vallis originates in Echus Chasma where our platy flows are σ is the total number of craters counted for each image. Error bars (1 ) only included observed. In addition, Central Elysium Planitia plains have been when larger than circles, triangles and squares. Bin divisions in the horizontal axis fi include 500 m, 707 m, 1 km, 1.4 km, 2 km, 2.8 km, and so on. The vertical axis gives the proposed to have been lled by a frozen sea displaying rafted plates density of craters (number of craters/km2 in each bin). (Murray et al., 2005). The problem of these contradictory interpreta- tions is that they are based on geomorphic landforms of very similar shape. No spectral study of platy-ridged flows has been performed to geological history, with an enrichment in HCP in recent lava flows discriminate these hypotheses, principally because all of these regions (Poulet et al., 2009b). are blanketed by dust. However, in the regions studied here, we observe two plains that display platy-ridged textures: the Eastern Noctis Chasma (NL2) and Echus Chasma (Figs. 8 and 9 respectively). 6. Discussion Both of these areas display a mineralogy of basaltic lava flows with predominant plagioclases and pyroxenes. While we cannot confirm 6.1. The origin of recent volcanism inside the Tharsis area that all platy-ridge flows are magmatic in origin, our data provide the first unambiguous evidence that such morphological features on Mars Recent volcanism in the Tharsis region is generally considered to may indeed be associated with basaltic lava flows. be limited to the main volcanic edifices, such as Arsia or Olympus Models of platy-ridged lava flows generally show a viscosity in the Mons and surrounding plains. In our study, we have found two pre- range of 100–1000 Pa s. (Keszthelyi et al., 2000). Viscosities of below viously undescribed plains b100 My old. These plains were previously 103 Pa s, and yield strengths less than 200 Pa are found in the Central mapped as undifferentiated floor material of Early to Middle Elysium Planitia (Vaucher et al., 2009), a location where extensive Amazonian age and of diverse nature (Scott and Tanaka, 1986). No platy flows have also been reported (Keszthelyi et al., 2000). recent volcanism has ever been reported inside the Valles Marineris– Unusually low viscosity or low yield strengths for silicate magma Noctis Labyrinthus troughs, despite the large amount of Hesperian age may be related to different parameters, including a high temperature volcanism surrounding these areas (e.g., Baptista et al., 2008). In or a low proportion of Si. The mean compositions of the two plains are addition, Echus Chasma floor has been confirmed to be covered by that of classical basalts, and not that of an ultrabasic rock. Although volcanic terrains of similar ages. The location of the three studied the mineralogy of platy-ridged flows has been established at the regions, much further to the East than the main volcanic centers of surface of the flows, this does not give the mineralogy at the bottom of Tharsis, is a feature of interest in itself to understand the evolution of flows. A downward segregation of olivine may have taken place, thus Valles Marineris Chasmata. being a possible explanation for the lack of olivine detection in the The Echus Chasma floor could have been filled from lava flows OMEGA spectra. However, the fact that olivine is systematically coming from the western plateau of the Echus-Kasei system of troughs, absent, or at least in limited proportion in small grains, even in eroded as suggested by Viking based mapping (Scott and Tanaka, 1986). Such parts of NL2, and that the detection of pyroxene, another dense an observation would suggest that the volcanic centers are actually on mineral, is clear, are both arguments that plead against a loss of the Tharsis plateau, close to the volcanoes. Nevertheless, fractures and olivine from the surface. Thus, the lack of detection of olivine suggests landforms observed on the bottom of Echus Chasma north of the that an ultrabasic composition was not necessary to explain the lavas studied area suggest that lavas were also extruded from fissures on the fluidity. A similar conclusion has been drawn from Icelandic examples canyon bottom (see Chapman et al., 2007, 2010-this issue for more where platy flows correspond to tholeiitic basalts (Keszthelyi et al., details), not only from the western plateaus. This suggests that local 2004). Thus, mechanisms such as deep crustal melting, dissolved magmatic sources also exist below the canyon floor. The two Chasmata volatiles (Vaucher et al., 2009), or high effusion rates (Keszthelyi et al., of Noctis Labyrinthus were not filled from the plateaus since there is no 2000) are necessary to explain the fluidity of platy-ridged flows rather sign of lavas flowing into the canyons. NL1 displays volcanic vents, than an ultramafic composition. N. Mangold et al. / Earth and Planetary Science Letters 294 (2010) 440–450 449
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