Mantle dynamics of Western Pacific and East Asia: Insight from seismic tomography and mineral physics

Recent results of high-resolution seismic tomography and mineral physics experiments are used to study mantle dynamics of Western Pacific and East Asia. The most important processes in subduction zones are the shallow and deep slab dehydration and the convective circulation (corner flow) processes in the mantle wedge. The combination of the two processes may have caused the back-arc spreading in the Lau basin, affected the morphology of the subducting Philippine Sea slab and its seismicity under southwest Japan, and contributed to the formation of the continental rift system and intraplate volcanism in Northeast Asia, which are clearly visible in our tomographic images. Slow anomalies are also found in the mantle under the subducting Pacific slab, which may represent (a) small mantle plumes, (b) upwellings associated with the slab collapsing down to the lower mantle, or (c) sub-slab dehydration associated with deep earthquakes caused by the reactivation of large faults preserved in the slab. Combining tomographic images and earthquake hypocenters with phase diagrams in the systems of peridotite + water, we proposed a petrologic model for arc volcanism. Arc magmas are caused by the dehydration reactions of hydrated slab peridotite that supply water-rich fluids to the mantle wedge and cause partial melting of the convecting mantle wedge. A large amount of fluids can be released from hydrated MORB at depths shallower than 55 km, which move upwards to hydrate the wedge corner under the fore-arc, and never drag down to the deeper mantle along the slab surface. Slab dehydration reactions at 120 km depth are the antigorite-related 5 reactions which supply water-rich fluids for forming the volcanic front. Phase A and Mg-surssasite breakdown reactions at 200 and 300 km depths below 700 °C cause the second and third arcs, respectively. Moreover, the dehydration reactions of super-hydrous phase B, phases D and E at 500–660 km depths cause the fluid transportation to the mantle boundary layer (MBL) (410–660 km depth). The stagnant slabs extend from Japan to Beijing, China for over 1000 km long, indicating that the arc–trench system covers the entire region from the Japan trench to East Asia. We propose a big mantle wedge (BMW) model herein, where hydrous plumes originating from 410 km depth cause a series of intra-continental hot regions. Fluids derived from MBL accumulated by the double-sided subduction zones, rather than the India–Asia collision and the subsequent indentation into Asia, are the major cause for the active tectonics and mantle dynamics in this broad region.     

Deep versus shallow origin of gravity anomalies, topography and volcanism on Earth, Venus and Mars

The relation between gravity anomalies, topography and volcanism can yield important insights about the internal dynamics of planets. From the power spectra of gravity and topography on Earth, Venus and Mars we infer that gravity anomalies have likely predominantly sources below the lithosphere up to about spherical harmonic degree l=30 for Earth, 40 for Venus and 5 for Mars. To interpret the low-degree part of the gravity spectrum in terms of possible sublithospheric density anomalies we derive radial mantle viscosity profiles consistent with mineral physics. For these viscosity profiles we then compute gravity and topography kernels, which indicate how much gravity anomaly and how much topography is caused by a density anomaly at a given depth. With these kernels, we firstly compute an expected gravity-topography ratio. Good agreement with the observed ratio indicates that for Venus, in contrast to Earth and Mars, long-wavelength topography is largely dynamically supported from the sublithospheric mantle. Secondly, we combine an empirical power spectrum of density anomalies inferred from seismic tomography in Earth’s mantle with gravity kernels to model the gravity power spectrum. We find a good match between modeled and observed gravity power spectrum for all three planets, except for 2?l?4 on Venus. Density anomalies in the Venusian mantle for these low degrees thus appear to be very small. We combine gravity kernels and the gravity field to derive radially averaged density anomaly models for the Martian and Venusian mantles. Gravity kernels for l?5 are very small on Venus below ≈800 km depth. Thus our inferences on Venusian mantle density are basically restricted to the upper 800 km. On Mars, gravity anomalies for 2?l?5 may originate from density anomalies anywhere within its mantle. For Mars as for Earth, inferred density anomalies are dominated by l=2 structure, but we cannot infer whether there are features in the lowermost mantle of Mars that correspond to Earth’s Large Low Shear Velocity Provinces (LLSVPs). We find that volcanism on Mars tends to occur primarily in regions above inferred low mantle density, but our model cannot distinguish whether or not there is a Martian analog for the finding that Earth’s Large Igneous Provinces mainly originate above the margins of LLSVPs.     

Atmospheric control of the cooling rate of impact melts and cryolavas on Titan’s surface

As on Earth, Titan’s atmosphere plays a major role in the cooling of heated surfaces. We have assessed the mechanisms by which Titan’s atmosphere, dominantly N₂ at a surface pressure of 1.5 × 10⁵ Pa, cools a warm or heated surface. These heated areas can be caused by impacts generating melt sheets and (possibly) by endogenic processes emplacing cryolavas (a low-temperature liquid that freezes on the surface). We find that for a cooling cryolava flow, lava lake, or impact melt body, heat loss is mainly driven by atmospheric convection. Radiative heat loss, a dominant heat loss mechanism with terrestrial silicate lava flows, plays only a minor role on Titan. Long-term cooling and solidification are dependent on melt sheet or flow thickness, and also local climate, because persistent winds will speed cooling. Relatively rapid cooling caused by winds reduces the detectability of these thermal events by instruments measuring surface thermal emission. Because surface temperature drops by ≈50% within ≈1 day of emplacement, fresh flows or impact melt may be difficult to detect via thermal emission unless an active eruption is directly observed. Cooling of flow or impact melt surfaces are orders of magnitude faster on Titan than on airless moons (e.g., Enceladus or Europa).Although upper surfaces cool fast, the internal cooling and solidification process is relatively slow. Cryolava flow lengths are, therefore, more likely to be volume (effusion) limited, rather than cooling-limited. More detailed modeling awaits constraints on the thermophysical properties of the likely cryomagmas and surface materials.     

Hydrovolcanic features on Mars: Preliminary observations from the first Mars year of HiRISE imaging

We provide an overview of features indicative of the interaction between water and lava and/or magma on Mars as seen by the High Resolution Imaging Science Experiment (HiRISE) camera during the Primary Science Phase of the Mars Reconnaissance Orbiter (MRO) mission. The ability to confidently resolve meter-scale features from orbit has been extremely useful in the study of the most pristine examples. In particular, HiRISE has allowed the documentation of previously undescribed features associated with phreatovolcanic cones (formed by the interaction of lava and groundwater) on rapidly emplaced flood lavas. These include “moats” and “wakes” that indicate that the lava crust was thin and mobile, respectively [Jaeger, W.L., Keszthelyi, L.P., McEwen, A.S., Dundas, C.M., Russel, P.S., 2007. Science 317, 1709-1711]. HiRISE has also discovered entablature-style jointing in lavas that is indicative of water-cooling [Milazzo, M.P., Keszthelyi, L.P., Jaeger, W.L., Rosiek, M., Mattson, S., Verba, C., Beyer, R.A., Geissler, P.E., McEwen, A.S., and the HiRISE Team, 2009. Geology 37, 171-174]. Other observations strongly support the idea of extensive volcanic mudflows (lahars). Evidence for other forms of hydrovolcanism, including glaciovolcanic interactions, is more equivocal. This is largely because most older and high-latitude terrains have been extensively modified, masking any earlier 1-10 m scale features. Much like terrestrial fieldwork, the prerequisite for making full use of HiRISE’s capabilities is finding good outcrops.     

Resolving the Elysium Controversy: An open invitation to explain the evidence

The origin of Mars’ Elysium-Amazonis plain has been the subject of debate for three decades, with a flood-volcanic origin the majority view today. Stratigraphical observations indicate that this view, based mainly on form analogy with Earth, is in error. A general unwillingness to engage with these observations in published work is detrimental not only to wider perceptions of planetary geology, often viewed as little more than the science of resemblance as a consequence, but also to superposed theories of surface chronology and source of the young martian meteorites, which are characterised by inconsistency and paradox as a result. Adherents to the volcanic hypothesis must convincingly explain how the observed field-relations are consistent with a lava surface across this region, or accept that these deposits are not what they currently believe them to be.     

The morphologies of volcanic landforms at Central Elysium Planitia: Evidence for recent and fluid lavas on Mars

This paper focuses on physical parameters (flow rates and rheological properties) of lava flows observed in the Central Elysium Planitia (CEP) region of Mars. The flows are modeled as Newtonian fluids, using the Jeffrey's equation and the concept of Graetz number, or alternatively as Bingham fluids. In addition to these approaches, a theoretical model of the shape of shield volcanoes based on the solution for the porous flow of an unconfined aquifer is applied to 5 shields, providing independent quantifications of rheological variations between the shields. This analysis indicates that of the five volcanoes studied, two are partially buried by lava postdating their formation, a result which has been confirmed independently in one case by high resolution images. Our observations reveal that two types of lava flows may be found in the CEP region. The first group is composed of large lava flows with viscosities around ∼2.5×¹⁰⁵ Pa s or yield strengths ranging from 100 to 500 Pa. The second group includes small lava flows of the shield volcanoes and large leveed lava channels on the plains with viscosities below 10³ Pa s, or yield strengths less than 200 Pa. When compared with other volcanic regions on Mars investigated with similar approaches, these latter values are, at present, the lowest inferred for martian lava flows. Several hypotheses for the formation of these lavas are discussed in the context of CEP given that low viscosity can be the result of (1) high temperature, (2) low crystal content, (3) low Si abundance of the liquid phase, and/or (4) the presence of dissolved volatiles. Two scenarios are considered. In the first one, it is demonstrated that low viscosity lavas (of low silica content) can be produced in the context proposed by Schumacher and Breuer [Schumacher, S., Breuer, D., 2007. Geophys. Res. Lett. 34. L12202] for recent volcanism. However, geochemical maps derived from GRS measurements do not provide support for anomalously low silica concentrations in this region. In the second scenario, a water-rich magma is proposed, although arguments in favor of a water-rich mantle source below the CEP are not available at the present time.     

Depth profiles of venusian sinuous rilles and valley networks

More than 200 venusian channels and valleys have been mapped based on analyses of Magellan SAR images. Sinuous rilles are the most abundant channels among six types of venusian channels, and they are widely distributed on Venus. Morphological characteristics of venusian sinuous rilles include sinuous narrowing reaches, source depressions, and length of several 10s to a few 100s of km. This type of channels is known to exist on the Moon and possibly on Mars. Valley networks on Venus often occur in the vicinity of or in connection to sinuous rilles. Cross-sectional morphologies of sinuous rilles and valley networks are of special importance in discussing their formation processes both qualitatively and quantitatively. We reconstructed cross-sectional profiles of 6 sinuous rilles and 2 valley networks using a new radar clinometric method. Reconstructed cross-sections revealed that floors of the channels and valleys are clearly lower than the surrounding plains. This finding implies that the sinuous rilles and the valley networks have erosional origins. Longitudinal depth profiles of the sinuous rilles show distinct decreasing trends toward the termini. Such decreasing trends of depths are qualitatively in agreement with theoretical models and laboratory experiments of thermal erosion. In order to verify this assertion quantitatively, we conduct simple 1-dimensional model calculations under the assumption that both channel-forming lavas and ground substrate are tholeiitic basalt. For initial lava thicknesses in the range 2-6 m, the model calculations yield good matches to the depth profiles. Estimated duration of lava effusion ranges from several months to a few years. These numerical results support thermal erosion of the sinuous rilles but do not necessarily exclude contributions from mechanical erosion processes. Valley networks seem to have formed under a strong structural control in comparison to sinuous rilles. The valleys vary widely in characteristics of the depth profile and flow directions relative to surface slopes. Therefore valley networks appear to have originated from diverse formation mechanisms.     

The frequency-area distribution of volcanic units on Venus: Implications for planetary resurfacing

The areas of volcanic units on Venus have been measured on the 1:5000000 geological maps published by NASA/USGS. These data were used to obtain a frequency-area distribution. The cumulative frequency-area distribution of 1544 specific occurrence of units cover six orders of magnitude from the largest unit (30 × 10⁶ km²) to the smallest (20 km²). The probability distribution function has been calculated. The medium and large volcanic units correlate well with a power-law (fractal) relation for the dependence of frequency on area with a slope of −1.83. There are fewer small units than the expected values provided by the power-law relation. Our measurements cover 21.02% of the planetary surface, 3.59% of the study area was found to be tessera terrain and is excluded from this study of volcanism. The measurements were restricted to areas where geological maps have been published. The analysis was performed on two independent areas of the planet, with a complete coverage of published maps. In both areas the largest volcanic unit covers a significant portion of the surface (58.75% and 63.64%, respectively). For the total measured volcanic units (excluding tessera), these two largest units (that could correspond to the same unit or not) cover the 61.18% and they are stratigraphically superimposed on older volcanic units which cover 3.37% of the area. The remaining area (35.45%) is occupied by younger volcanic units stratigraphically superimposed on the large volcanic unit(s). These results are based on the independent mapping of a large number of geologists with different ideas about the geodynamical evolution of Venus and different criteria for geological mapping. Despite this fact, the presence of these very large units is incompatible with the equilibrium resurfacing models, because their generation at different ages would destroy the crater randomness. Our frequency-area distribution of the mapped volcanic units supports a catastrophic resurfacing due to the emplacement of the largest unit(s) followed by a decay of volcanism. Our data for the frequency-area distribution of volcanic units provide new support for catastrophic resurfacing models. It is difficult to make our observations compatible with equilibrium, steady-state resurfacing models.     

Venus astra/novae: Estimates of the absolute time duration of their activity

In order to assess the age relations between astra/novae (features with extensive radial fracture-graben systems) and their surroundings, and to determine the duration of their activity, we undertook a photogeologic analysis of Magellan images of 78 astra, 49 dark-parabola craters and 114 clear-halo craters. For seven of these 78 astra it was found that the astrum-forming faulting started before or close to the time of emplacement of regional plains and extended into the second part of post-regional-plains time. Because the mean age of the regional plains is close to the mean surface age of Venus (which is estimated to be ∼750 m.y), this means that the duration of activity of these seven astra was several hundred millions of years. This is longer than the duration of activity of ongoing mantle plumes on Earth, but shorter than the duration of activity of the plume feeding the martian volcano Olympus Mons. The basic morphologic characteristics of these seven astra, as well as their age relations with other geologic units, are similar to those of the majority of other astra; therefore, such a long duration could also be typical of some other astra. We confirm the two-phase (pre- and post-regional-plains) evolution of astrum-forming faulting suggested in previous studies. For the first phase we show evidence for purely tectonic faulting caused by the diapiric rise of a mantle plume. For the second phase we find evidence supporting the interpretation of other studies that the observed faults resulted from subsurface dike intrusions produced by magmatism associated with the plume. We also found that faulting during the second phase was not instantaneous but distributed over a long period of time.     

A combined spectrophotometric and morphometric study of the lunar mare dome fields near Cauchy, Arago, Hortensius, and Milichius

In this study we examine the spectral and morphometric properties of the four important lunar mare dome fields near Cauchy, Arago, Hortensius, and Milichius. We utilize Clementine UV-vis multispectral data to examine the soil composition of the mare domes while employing telescopic CCD imagery to compute digital elevation maps in order to determine their morphometric properties, especially flank slope, height, and edifice volume. After reviewing previous attempts to determine topographic data for lunar domes, we propose an image-based 3D reconstruction approach which is based on a combination of photoclinometry and shape from shading. Accordingly, we devise a classification scheme for lunar mare domes which is based on a principal component analysis of the determined spectral and morphometric features. For the effusive mare domes of the examined fields we establish four classes, two of which are further divided into two subclasses, respectively, where each class represents distinct combinations of spectral and morphometric dome properties. As a general trend, shallow and steep domes formed out of low-TiO₂ basalts are observed in the Hortensius and Milichius dome fields, while the domes near Cauchy and Arago that consist of high-TiO₂ basalts are all very shallow. The intrusive domes of our data set cover a wide continuous range of spectral and morphometric quantities, generally characterized by larger diameters and shallower flank slopes than effusive domes. A comparison to effusive and intrusive mare domes in other lunar regions, highland domes, and lunar cones has shown that the examined four mare dome fields display such a richness in spectral properties and 3D dome shape that the established representation remains valid in a more global context. Furthermore, we estimate the physical parameters of dome formation for the examined domes based on a rheologic model. Each class of effusive domes defined in terms of spectral and morphometric properties is characterized by its specific range of values for lava viscosity, effusion rate, and duration of the effusion process. For our data set we report lava viscosities between about 10² and , effusion rates between 25 and , and durations of the effusion process between three weeks and 18 years. Lava viscosity decreases with increasing R₄₁₅/R₇₅₀ spectral ratio and thus TiO₂ content; however, the correlation is not strong, implying an important influence of further parameters like effusion temperature on lava viscosity.