Europlanet Science Congress 2020
Virtual meeting
21 September – 9 October 2020
Europlanet Science Congress 2020
Virtual meeting
21 September – 9 October 2020

Oral presentations and abstracts


The session will gather researchers of different communities for a better understanding of the evolution and properties of small bodies, in particular the parent bodies of meteorites.
It will address recent progresses made on physical and chemical properties of these objects, their interrelations and their evolutionary paths by observational, experimental, and theoretical approaches.
We welcome contributions on the studies of the processes on and the evolution of specific parent bodies of meteorites, investigations across the continuum of small bodies (comets, planetesimals, asteroids, dwarf planets) ranging from local and short-term to global and long-term (thermal and thermochemical) processes, studies of the surface dynamics on small bodies, studies of exogenous and endogenous driving forces of the processes involved, as well as statistical and numerical impact models for asteroids observed closely within recent missions (e.g., Hayabusa2, New Horizons, OSIRIS-REx).

Convener: Wladimir Neumann | Co-conveners: Sabrina Schwinger, Ottaviano Ruesch, Marco Ferrari

Session assets

Session summary

Chairperson: Wladimir Neumann, Sabrina Schwinger, Marco Ferrari, Ottaviano Ruesch
Processes: Impacts
Ronald-Louis Ballouz, Patrick Michel, Olivier Barnouin, Kevin Walsh, Martin Jutzi, Eri Tatsumi, Maria Antonella Barucci, Daniella DellaGiustina, Humberto Campins, Seiji Sugita, Seiichiro Watanabe, Hirdy Miyamoto, William Bottke, Harold Connolly, Makoto Yoshikawa, and Dante Lauretta

Disruption and Reaccumulation: 

Asteroids such as Ryugu and Bennu are likely fragments formed from a larger body that was disrupted in the main asteroid belt [1,2]. Numerical simulations of asteroid disruptions—including the fragmentation phase during which the asteroid is broken up into small pieces and the gravitational phase during which fragments may reaccumulate due to their mutual attractions—lead to a family of rubble piles over a range of sizes [3]. Considering microporous parent bodies of 100 km in diameter, we found that their disruption (Fig. 1) can lead to rubble piles with oblate spheroidal or top shapes [4]. Moreover, assuming that the parent body is hydrated, the various degrees of heating at impact can produce rubble piles with different level of hydration as a result of a single parent body disruption.

We proposed two scenarios where Ryugu and Bennu could originate from the same parent body. In scenario a, Ryugu and Bennu are composed from materials sourced from near the impact point and near its antipode, respectively. In scenario b, Ryugu and Bennu are composed from materials sourced from the parent-body center and near the impact point’s antipode, respectively. The detected signature of  exogeneous material introduces new complexities to the collisional origin of Ryugu and Bennu [5, 6].

Rubble Pile Contamination:

Due to the apparent spectral homogeneity observed on the surfaces of Bennu and Ryugu during the first observational campaigns, our simulations in [4] only considered the fate of material originating from the parent body, assumed to be homogenous in composition. However, subsequent spectral data from the OSIRIS-REx and Hayabusa2 missions show a small fraction of anhydrous silicate material on the surface of the two bodies [5, 6]. The presence of this material can be explained by retention of a projectile on either the parent body or on the rubble piles themselves after their formation.  However, projectile retention efficiencies for impacts of anhydrous silicates on hydrated minerals are poorly constrained [7, 8] for expected impact speeds in the main asteroid belt (~ 5km/s, [9]). Here, we investigate whether the family-forming catastrophic disruption can lead to the incorporation of impactor material in the reaccumulated family members, leading to the small fraction of apparently exogeneous material on their surface.

Figure 1 Outcome of a SPH simulation of the disrup-tion of a microporous 100-km-diameter parent body. Each particle is a fragment. Colors represent the vari-ous degrees of impact heating. This outcome is the starting point of the gravitational phase during which the fragments reaccumulate to form rubble piles.

Figure 1: Outcome of a SPH simulation of the disrup-tion of a microporous 100-km-diameter parent body. Each particle is a fragment. Colors represent the various degrees of impact heating. This outcome is the starting point of the gravitational phase during which the fragments reaccumulate to form rubble piles.


We performed a series of numerical simulations of sub-catastrophic and catastrophic disruption of 1- and 100-km-diameter microporous asteroids. We account for both the parent body material and the projectile material in the subsequent gravitational phase when fragments re-accumulate to form the parent-body remnant and smaller rubble-pile family members. As in our previous works, the fragmentation phase was simulated using a Smoothed Particle Hydrodynamics (SPH) hydrocode, and the gravitational phase was computed using the N-body code pkdgrav, including the Soft-Sphere Discrete Element Method (SSDEM) [10]. We then track the surviving materials of both the projectile and the parent body, including their level of heating, as they reaccumulate. For each aggregate, we measure their shapes, the fractions of projectile and parent body materials that compose them, and their associated level of heating.

Projectile material was neglected in previous work because asteroid families appear spectrally homogeneous, suggesting that they are mostly made of the material of their parent body. The advanced observational capabilities of space missions enabled the discovery that this scenario may be more complex.


Observational analysis of exogenous material on Ryugu and Bennu provide constrains for our numerical simulations. In particular, the total volume and the spectral characteristics of the exogenous material can be measured [5,6,11]. The total volume bounds the required contamination efficiency and/or the total time needed to contaminate the parent body. The spectral analysis shows that Bennu hosts HED-like material whereas Ryugu has ordinary chondrite–like material.

This difference in the spectral signature of exogenous material may render scenario b (outlined above) invalid, as our preliminary calculations show that contamination on large 100-km parent bodies is likely only limited to its outer shell. Thus, it is difficult to form a 1st generation rubble-pile that has both: i) material from the parent body core, and ii) exogenous material that originated from the contamination of the original parent body’s outer shell. This scenario may be possible if the asteroid is a 2nd generation object, with its precursor being an approximately 20-km rubble-pile that incorporated material originating from both the  center and exterior of the parent body [12, 13]. Our numerical simulations will provide claraity on the feasibility of these various scenarios. Ultimately, analysis and comparison of the returned samples will provide clarity on the potential shared collisional origin of Ryugu and Bennu, and the prevalence of impact contamination in the Solar System. 


This material is based upon work supported by NASA under Contract NNM10AA11C issued through the New Frontiers Program. P.M. acknowledges support from the Centre National d’Études Spatiales and from the Academies of Excellence on Complex Systems and Space, Environment, Risk and Resilience of the Initiative d’EXcellence “Joint, Excellent, and Dynamic Initiative” (IDEX JEDI) of the Université Côte d’Azur. We are grateful to the entire OSIRIS-REx and Hayabusa2 teams for making the encounters with Bennu and Ryugu possible.


[1] Michel, P. et al. (2001) Science, 294, 1696–1700. [2] Walsh, K.J. (2018) ARA&A 56, 593. [3] Jutzi, M., et al. (2019) Icarus 317, 215. [4] Michel, P., Ballouz, R.-L. et al. (2020) Nature Comm. 11, 2655.  [5] DellaGiustina, D.N., et al. (2019) EPSC-DPS2019-1074.  [6] Sugimoto, C., et al. (2019) Asteroid Science in the Age of Hayabusa2 and OSIRIS-REx, 2051. [7] Avdellidou, C., et al. (2016) MNRAS, 456, 2957. [8] Daly, R.T., & Schultz, P. H. (2018) M&PS, 53, 1364. [9] Bottke, W.F., et al. (2005) Icarus, 179, 63.  [10] Ballouz, R.-L., et al. (2019) MNRAS 485, 697. [11] Campins, H., et al. (2020) EPSC. [12] Walsh, K.J., et al. (2020) LPSC 51, 2253.  [13] Sugita, et al. (2019) Science 364, 252.

How to cite: Ballouz, R.-L., Michel, P., Barnouin, O., Walsh, K., Jutzi, M., Tatsumi, E., Barucci, M. A., DellaGiustina, D., Campins, H., Sugita, S., Watanabe, S., Miyamoto, H., Bottke, W., Connolly, H., Yoshikawa, M., and Lauretta, D.: Modeling the contamination of Bennu and Ryugu through catastrophic disruption of their precursors , Europlanet Science Congress 2020, online, 21 Sep–9 Oct 2020, EPSC2020-510,, 2020.

Rutu Parekh, Katharina Otto, Ralf Jaumann, Klaus-Dieter Matz, Thomas Roatsch, Elke Kersten, Stephan Elgner, Katrin Krohn, and Carol Raymond


Ponded craters have been predominantly identified on small, dry planetary bodies like (433) Eros and Itokawa. We identified similar features on Vesta, where loose fragmented ponded materials are present on small crater floors. While the morphological details of the ponded features on Vesta and Eros/Itokawa are similar, their production mechanisms may vary, due to differences in gravity or the insolation environment Previous studies conducted on Vesta have provided evidence for volatile outgassing in some regions. In this study, we investigate the morphology of the ponded crater and possible involvement of volatiles outgassing and its interaction with surface material in producing ponded craters on Vesta.   

Ponded craters have widely received lime light due to its unusual characteristics on Eros revealed by the NEAR Shoemaker mission (Sears et al., 2015, Robinson et al., 2001,2002). In general, ponded craters show a smooth layer of fine-grained material with grain size less then cm (Robinson et al., 2001) partially covering topography of the crater floor (Figure 1). They may also possess varying sized boulders or unconsolidated material (Sears et al., 2015). The depth of the pond is about ~5% of the depth of the original crater (Robinson et al., 2001) on Eros. Other than the distinct morphological impression (smooth and flat floor), ponded crater regolith also shows sharp variation in the spectral signature (Robinson et al., 2001) which can be due to mineral heterogeneity (Robinson et al., 2001), space weathering (Sears et al., 2015, Heldmann et al., 2010, Robinson et al., 2001) or the difference in grain size between regolith and the surrounding region (Heldmann et al., 2010, Robinson et al., 2001).

Based on the evidences on Eros, the formation mechanisms of ponded craters include electrostatic levitation, seismic shaking and/or boulder comminution (Robinson et al., 2001). However, the effects of these mechanisms may vary on other dry planetary bodies with different compositions, gravity or insolation intensity. In our study, we characterize ponded craters on Vesta to understand their formation mechanisms and how interactions with the regolith may have influenced the generation of ponded craters.

Figure 1: A classic example of pond crater on Eros. The crater has diameter of ~0.09km (Robinson et al., 2001). The pond material has sharp boundary, low albedo and flat smooth surface which makes it easy to distinguish from the original carter floor. Image source: Robinson et al., 2001  

For the identification of ponded craters on Vesta, HAMO mosaics (~70 m/pixel) and LAMO mosaics (~20 m/pixel) provided by the NASA Dawn Mission were considered. To extract elevation information, a DTM of HAMO resolution was used (92 m/pixel) (Preusker et al., 2016) prepared from stereo-pairs.

So far, we have identified 10 ponded craters nearby the equator (0°-30°) on Vesta. Overall, the crater floor is fully or partially covered by fine and loose material. The usual diameter of ponds ranges from 0.9 -6.4 km within craters of 1.78- 8.43 km diameter. Most of the identified ponded craters, have clear flat floors in which the fine material is evenly distributed within the bowl-shaped depression (Figure 1) covering the original floor of the crater entirely. An example is given in Figure 2, located at 15°S,189°E. The carter has a diameter of ~8km. By fitting a polynomial shape (e.g. a parabola) to the crater walls, we estimated the original depth of the crater with ~0.66km. The ponded material has filled the original crater surface, producing a shallow depth crater. The material has filled ~0.31km of the crater, which means half of the original crater depth is infilled by the fine material. On Eros ponds have average infilling depths of ~10cm or 5% of the original depth (Robinson et al., 2001). The smaller infilling might be due to the fact that craters on Eros are significantly smaller in comparison with Vesta. However, it is unreasonable to draw any conclusions based on a single example. At the meeting we will present measurements of the rest of the identified ponded sites to understand the overall morphology and discuss formation mechanisms for ponds on Vesta based on our findings.

Figure 2: Example of a ponded crater on Vesta. (a) The ponded material is covering the original surface within central crater region and exhibits a flat and smooth texture. (b) The elevation profile of the crater in (a) highlighting the flat ponded material in red. The ponded material infilled the deeper parts of the crater depression, masking the original shape (dash-dotted line) and generating a flat floor.


[1] Heldamann et al., 2010, Icarus, 206, 685-690. [2] Robinson et al., 2001, Nature, 413,396-400. [3] Robinson et al., 2002, Meteorit. Planet. Sci., 37 1651-1684. [4] Sears et al., 2015, PSS, 117, 106-118.


How to cite: Parekh, R., Otto, K., Jaumann, R., Matz, K.-D., Roatsch, T., Kersten, E., Elgner, S., Krohn, K., and Raymond, C.: Ponded craters on Vesta, Europlanet Science Congress 2020, online, 21 Sep–9 Oct 2020, EPSC2020-656,, 2020.

Stamatios Xydous, Angeliki Papoutsa, Ioannis Baziotis, Jinping Hu, Chi Ma, and Paul Asimow


Sodic plagioclase is common in Earth’s crust and in many differentiated and undifferentiated meteorites. Under high temperature (HT) and high pressure (HP) conditions in asteroidal collisions, sodic plagioclase may transform into either hollandite-structured lingunite [1] or the recently discovered albitic jadeite [2]. When stoichiometric jadeite forms by decomposition of albite, the excess silica forms an SiO2 polymorph, often stishovite [3]. Albitic jadeite, by contrast, a Na-rich analogue of tissintite [2], is super-silicic, vacancy-rich pyroxene with excess Si coordinated in the octahedral M1 site. Searching for albitic jadeite alongside other P-sensitive mineral assemblages is therefore potentially important for expanding the list of pressure constraints available for impact events.

We report preliminary results on the occurrence of albitic jadeite within shock veins in the L6 ordinary chondrites Ozerki and Chug-Chug-011 (Fig. 1). Ozerki (fell 21st June 2018 in Russia) is moderately shocked (S4/5) and un-weathered (W0); it was recovered quickly (25th June 2018) after its fall. Chug-Chug-011 is a find, recovered in 2018 in Antofagasta, Chile; it is weakly shocked (S2), with minor weathering (W1).

Materials and Methods

Polished thin sections of Ozerki and Chug-Chug-011 were carefully examined for shock indicators and HP polymorphs, with intensive focus on the melt veins (MVs). We used optical microscopy, a JEOL JSM-IT300LV scanning electron microscope, a JEOL JXA 8900 electron probe micro-analyzer, and a dispersive confocal Renishaw inVia Reflex Raman microscope (514 nm laser).

Petrography & mineral chemistry

The thin section of Ozerki displays two discrete areas (Fig. 1A); light-colored chondritic and dark-colored impact melt-rich area. We focused on a network of shock veins intruding the light-colored area. The MVs are dark, variable width (40-850 μm), and locally disrupted by angular to sub-rounded clasts. Clasts are more abundant in wider MVs; jigsaw-fit breccia textures are widespread. Clasts, mostly silicate, concentrate in the center of each MV, whereas the margins are rich in metallic segregations and sulfides.

In Chug-Chug-011, three different MVs (~100 μm wide) crosscut the meteorite matrix (Fig. 1B). Elongated silicate clasts oriented parallel to the veins are common in their central domains.

In Ozerki, albitic jadeite forms acicular to dendritic crystallites aggregates (≤ 2 μm) associated with feldspathic glass (Fig. 2A). In Chug-Chug-011, albitic jadeite is found within a composite clast: low Ca-pyroxene surrounds sodic plagioclase (Fig. 2B). Crystallites near the core of the plagioclase show brighter backscatter than those near the rim.

Albitic jadeite in Ozerki yields an empirical formula (Na0.70Ca0.15K0.050.14)(Al0.82Si0.10Fe0.04)Si2O6 whereas that from Chug-Chug-011 is variable: (Na0.57-0.64Ca0.07-0.07K0.03-0.05Mg0.01-0.070.16-0.29)(Al0.78-0.86Si0.10-0.18Fe0-0.05Mg0-0.13)Si2O6, with Ca# [100×Ca/(Ca+Na)] from 10 to 13.

Pyroxene Raman spectroscopy

Raman spectra of the albitic jadeite in Ozerki display five distinct peaks at 376, 526, 698, 986 and 1036 cm-1 (Fig. 3A). In Chug-Chug-011, the predominant peak is at 698 cm-1, but there is a noteworthy 1016 cm-1 peak in addition to the “typical jadeite” 1038 peak. This may be associated either with a diopside-related structure or another high-P clinopyroxene (Fig. 3B).

Discussion and Conclusions

In Ozerki, albitic jadeite was found in the middle of ~70 μm and ~300 μm wide MVs. The presence of equant idiomorphic crystals with 120° triple junctions suggests that these MVs reached peak HT above the liquidus of the matrix. From such conditions, a ~300 μm wide vein surrounded by cold matrix conductively cools and solidifies in ~6.5 ms, which is an upper limit for growth time of minerals in the MV. Albitic jadeite is less dense than lingunite, implying formation from sodic plagioclase at lower pressures. The absence of lingunite suggests maximum pressures below 21 GPa. According to experiments [4] in jadeite-rich compositions (Jd70-80), jadeite + stishovite + garnet is stable at 13.5-21.5 GPa. However, the absence of stishovite and garnet in our MV may only reflect sluggish nucleation of these phases rather than an insufficient peak P<13.5 GPa [5]. The presence of albitic jadeite, by itself, therefore yields only an upper limit and not a fully quantitative P constraint.

In Chug-Chug-011, high-pressure Na-clinopyroxene [(Na0.49Ca0.15K0.03Mg0.240.09)(Al0.62Si0.04Fe0.13Mg0.21)Si2O6] is enclosed in a melt pocket included in pyroxene that is in turn entrained in a MV. The bright crystallites near the center of the pocket yield compositions and spectra similar to the HP-sodic clinopyroxene identified by [6]. The backscatter-dark crystallites closer to the pocket margins better match albitic jadeite. Neither phase is yet calibrated for shock pressure. However, the presence of a mixed xieite-chromite spectrum at the rim of another MV in the section suggests higher P conditions, 18-23 GPa (Fig. 3C). The same MV shows minor wadsleyite peaks near its center, requiring gradients over space or time in preserved P and T conditions across the MV.


This research received support from European Social Funds and the Greek State (call code EDBM103).


[1] Gillet, P., et al. 2000. Science, 287(5458), 1633-1636; [2] Ma, C., et al. 2020. 51st LPSC, #1712; [3] Liu, L.G. 1978.EPSL, 37(3), 438-444; [4] Bobrov, A.V. et al. 2008. GCA, 72, 2392-2408, 2008; [5] Kubo, T., et al. 2009. Nature Geoscience, 3, 41-45, 2009; [6] Baziotis, I., et al. 2018. Scientific Reports, 88, 9851, 2018.

Fig. 1: Transmitted-light mosaics of (A) Ozerki and (B) Chug-Chug-011; rectangles indicate the areas hosting HP polymorphs (Figs. 2, 3).


Fig. 2: A) BSE image of MV in Ozerki showing albitic jadeite crystals (spectra #31 and #51 in Fig. 3A) in a partly crystallized melt area. B) BSE image of MV in Chug-Chug-011 with albitic jadeite (spectra in Fig. 3B). Bright core near C3 may be HP sodic clinopyroxene (see text).


Fig. 3: A) Ozerki Raman spectra: typical jadeite peaks at ~698, 986, and ~1036 cm–1 in spectra #3_31, #3_51. Spectrum #5_11 shows the 698 cm-1 peak but the two higher wavenumber peaks are not clearly resolved. B) Chug-Chug-011 Raman data: jadeite peak at 698 cm-1 is apparent. The peak at 960 cm-1 in spectrum C12 is apatite. C) Chug-Chug-011 MVA spectra: rim point MVA3_91 is mixture of chromite and xieite with olivine. Center point MVA1_151 shows wadsleyite peaks at 720 and 915 cm-1.


How to cite: Xydous, S., Papoutsa, A., Baziotis, I., Hu, J., Ma, C., and Asimow, P.: High-pressure clinopyroxene formation in L6 chondrites (Ozerki, Chug-Chug-011): Implications for impact processes, Europlanet Science Congress 2020, online, 21 Sep–9 Oct 2020, EPSC2020-932,, 2020.

Processes: Erosion / Space Weathering
Leon Hicks, John Bridges, Takaaki Noguchi, and Jack Piercy


Space weathering due to the bombardment of electrons and solar wind upon the exposed lunar surface shows as an apparent spectral darkening and reddening in ground-based and lunar-orbital observations [1].

Space weathered rims have been observed on soil surface samples, returned by the Apollo landings [1,2], featuring amorphized material and nanophase Fe metal (npFe⁰) particles formed due to the implantation of solar wind H⁺ ions reducing the host grain mineral oxides to form metal [2,3]. Oxidation of these Fe particles has also been shown, and a suggested correlation between oxidation and lunar soil maturity [3].

In this study, we investigate Fe-redox changes in the space weathered rims of Apollo 17 lunar surface soil samples, using TEM and X-ray nanoprobe Fe-K XANES.

Samples and Methods

The lunar sample number is 78481,29, collected from the top 1 cm of trench soils during the Apollo 17 lunar landing [4]. Observing an abundance of blisters (a feature of space weathering [1,5]) on the grain surfaces indicated space weathering to be investigated, and FIB-SEM lift-out sections were extracted from at least three sample grains.

The three grains consisted of two augite pyroxenes, En₈₁Fs₁₆ (#A17-3) and En₈₅Fs₁₂ (#A17-5), and one olivine Fa₃₉ (#A17-6), with partially amorphised space weathered rims measuring up to ~100 nm thick featured in all three (e.g. Figure 1a), observed in HAADF-STEM imaging using a JEOL JEM-ARM300CF at ePSIC in the Diamond Light Source synchrotron facility. A deposition of tungsten (W-dep) on the grain surface was used during FIB lift-out preparation which can be seen in Figure 1a.

Fe-redox was analysed using the I-14 X-ray Nanoprobe Beamline at Diamond, similarly to previously investigated Itokawa asteroid samples [6,7,8]. A series of XRF maps are obtained over the Fe-K XAS energy range 7000 and 7300 eV, with a higher resolution of energy increments over the XANES features (~7100-7150 eV). Using Mantis 2.3.02 [9] to process the XANES mapping, spectra can be isolated for the space weathered zone (SWzone), the grain ‘Host’ mineralogy, and the W-dep. All measured spectra are normalized in Athena 0.8.056 [10].


Nano-grains measuring ~2-3 nm in diameter within the partially amorphised space weathered zones (see Figure 1b), confirmed to be Fe metal when measuring lattice fringe spacings. Figure 1c shows a nano-grain with lattice spacings measuring ~2.06 Å, observed in the #A17-3 augite sample. Other lattice spacing measurements of up to ~2.10 Å in each of the three lunar grains confirms the presence of Fe metal, similar to previous studies of Itokawa samples [5].

In the Fe-K XANES analysis, the 1s→3d pre-edge centroid positions are defined by the intensity-weighted average across baseline subtracted peaks (see Figure 2). A shift in the 1s→3d centroid position from the host mineral to the SWzone suggests Fe redox variation, where an positive shift in energy position is an increase in ferric (Fe³⁺) content, based on a ferric-ferrous ratio (Fe³⁺/ΣFe) defined by Fe-K XANES measurements of standard reference minerals. The largest variation observed between SWzone and the respective substrate host mineral is a positive shift of 0.23 ±0.06 eV, equivalent to a ferric increase of Fe³⁺/ΣFe = 0.14 ±0.03 from host to SWzone. There is a consistent positive shifting in the 1s→3d centroid energy positions observed in all three lunar samples, with average ferric increases of: Fe³⁺/ΣFe = 0.08 ±0.03 (#A17-3); Fe³⁺/ΣFe = 0.11 ±0.03 (#A17-5); and Fe³⁺/ΣFe = 0.10 ±0.03 (#A17-6).


A minor Fe-redox variation of up to Fe³⁺/ΣFe ~0.14 ±0.03 has been observed in these lunar grains, similar to the space weathered rims of asteroid Itokawa grains with increased ferric contents ranging Fe³⁺/ΣFe ~0.02-0.14 ±0.03 [6,7,8]. The ferric increase is in the dominant silicate phase of the amorphized rims, as Fe metal appears to have had no influence on the Fe-K XAS spectra. Additionally, the npFe⁰ particles show no oxidation present, confirmed by the Fe metal lattice spacings, unlike the oxidised Fe particles observed in previous space weathered lunar regolith samples [3].

The minor increase in ferric material is likely the result of solar wind H⁺ implantation on the ferrous silicate grain surfaces, segregating the Fe to form Fe metal. The resulting water vapour then oxidises some of the remaining Fe to Fe³⁺ in the silicate phase.

It has been proposed that the spectral reddening observed on surfaces of airless bodies is due to the npFe⁰ particles formed in space weathering [1], and spectral reflectance models have also shown that Fe³⁺ can increase reddening too [3]. Continued investigations into the mineralogical complexities associated with space weathering could reveal more spectral effects. Further samples to be investigated will include new Apollo samples, as well as potential space weathered material collected by asteroid missions Hayabusa2 and OSIRIS-REx.


The three Apollo lunar surface grains feature space weathered rims up to ~100 nm thick. The npFe⁰ particles measured lattice spacings of ~2.06-2.10 Å, suggesting they had not been oxidised. However, Fe-K XANES analyses suggest a minor oxidation of up to Fe³⁺/ΣFe ~0.14 ±0.03 occurring in the dominant silicate phase of the space weathered rims. These results are consistent with previously analysed space weathered asteroid Itokawa samples, suggesting Fe-redox variations in silicate material exposed on the surfaces of airless planetary bodies may be a key part of space weathering effects.


[1] Pieters C. A. and Noble S. K. (2016) JGR: Planets, 121, 1865-1884. [2] Hapke B. (2001) JGR, 106, 10039-10073. [3] Thompson M. S. et al. (2016) Meteorit. Planet. Sci., 51, 1082-1095. [4] Butler P. (1973) MSC 03211 Curator’s Catalog. pp 447. [5] Noguchi T. et al. (2014) Meteorit. Planet. Sci., 49, 188–214. [6] Hicks L. J. et al. (2019) LPSC L, Abstract #1805. [7] Hicks L. J. et al. (2019) 82nd Annual Meeting of The Meteoritical Society, Abstract #6330. [8] Hicks L. J. et al. (2019) 2nd British Planetary Science Conference, Abstract. [9] Lerotic M. et al. (2014) J. Synchrotron Radiat., 21, 1206–1212. [10] Ravel B. and Newville M. (2005) J. Synchrotron Radiat., 12, 537–541.

How to cite: Hicks, L., Bridges, J., Noguchi, T., and Piercy, J.: Fe-K XANES and HR-TEM analyses of Apollo lunar grain space weathered surfaces, Europlanet Science Congress 2020, online, 21 Sep–9 Oct 2020, EPSC2020-715,, 2020.

Processes: Chemical Alteration
Angeliki Papoutsa, Stamatios Xydous, and Ioannis Baziotis


Calcium-aluminum-rich inclusions (CAI) in carbonaceous chondrites are considered to be of the most primitive objects, formed during the nebula phase of the Solar System [1].  As such, they are of paramount importance, since they preserve record of early solar processes and conditions during their formation, reflected in a series of mineral assemblages [2][3]. We present preliminary textural and chemical results of our petrological study of CAIs from the CM2 chondrite AMU 17290.

Materials and Methods

AMU 17290 (section #9) is the only meteorite recovered to date from Amundsen Glacier, during the Antarctic Search for Meteorites (ANSMET) expedition in 2017-2018, and was classified as a CM2 carbonaceous chondrite. The studied sample is a polished section, that has been carefully examined for its texture, mineralogy and mineral chemistry by using optical microscopy (reflected light), Scanning Electron Microscopy (JEOL JSM-IT300LV at the American School of Classical Studies in Athens), and electron microprobe microanalysis (JEOL JXA-8900 Superprobe at the Agricultural University of Athens).

Petrography and mineral chemistry

AMU 17290 presents a high matrix percentage that covers ~75 vol.% of the section, while chondrules (excluding CAIs) comprise ~20 vol.%. The chondrules are predominantly porphyritic olivine (PO) and porphyritic olivine-pyroxene (POP) type, and are surrounded by fine-grained accretion rims.  Olivine has nearly pure forsterite composition (Fo99). Pyroxene, on the other hand, from the chondrules, span it`s composition between augite (Wo38En60Fs2) and enstatite (Wo4En93F3). Chondrule fragments, are entrained in the matrix, and unlike chondrules, they do not preserve their spherical shape. Secondary gypsum appears in the matrix, and together with serpentine and other fine-grained phyllosilicates, represent alteration products.

Refractory inclusions are CAIs and amoeboid olivine assemblages (AOAs). As expected in CM chondrites, CAIs in AMU 17290 are less than 1mm in diameter, and classified as fluffy type A. AMU 17290, in particular, contains CAIs of three primary mineralogical compositions: a) spinel-perovskite (Fig. 1A), b) spinel-diopside-perovskite (Fig. 1B), and c) spinel-diopside-hibonite-perovskite (Fig. 1C,D). Most inclusions, are surrounded by fine-grained phyllosilicate rims, and abundant cronstedtite, and tochilinite (Figs. 1 and 2), while secondary calcite has been identified, as well. Even though, the mineralogy of CAIs resembles type A inclusions, melilite, a typical phase in such CAIs, is noticeably absent.

Formation of CAIs and early aqueous alteration

The CAIs are considered to have formed during the first stages of the solar nebula [1]. Considering an equilibrium condensation model, coexistence of hibonite and perovskite, points out to an estimated formation temperature of ~1751 K (at 10-3 bar), where hibonite first condenses, followed by perovskite [3]. However, both hibonite and perovskite are found as inclusions within spinel (Fig. 1C,D), and therefore a paragenetic sequence cannot be accurately inferred.

The absence of melilite, in type A CAIs has been previously attributed to early stages of aqueous alteration [5,6]. Another hypothesis, considering the limited occurrence of hibonite in AMU, is that condensation of these CAIs may have occurred near the thermal boundary between hibonite and before perovskite. If limited condensation of hibonite occurred in these CAIs, there was consequently limited interaction with it and the solar nebula, which is assumed to form melilite [7]. Nonetheless, petrographic evidence suggests an effect of early aqueous alteration of CAIs, that presumably occurred while on the parent body. The presence of cronstedtite in CAIs, has been associated with the early dissolution of melilite by fluids with high Fe/Si ratio at low temperatures (50-120°C) [8,9]. Such fluids, may be derived by reaction of water and dissolving Fe-rich phases, such as metallic Fe-Ni, and Fe-silicates, which occur in the matrices of the meteorites [9]. A rather large spherical, porous, spinel-perovskite-rich CAI (200 μm in diameter), shows a two-layer rim: a) a cronstedtite-tochilinite rim in direct contact with the spinel, and b) an outer Al-diopside rim in contact with cronstedtite-tochilinite (Fig. 2). Tochilinite in that layer occurs as intergrowths with cronstedtite. Although the presence of tochilinite suggests an early aqueous alteration at a temperature range higher than that of cronstedtite (120-160°C), the coexistence of these two secondary phases has been experimentally placed at 80-120°C, in neutral-alkaline conditions [9]. Secondary calcite occurs in the interior of the inclusion, within the spinel, suggesting localized Ca-mobility along with Fe and Si that formed cronstedtite and tochilinite. Therefore, a hypothesis under investigation, is that the combined occurrence of secondary calcite, cronstedtite and Al-diopside in absence of melilite, may result from the dissolution of the latter. Such an alteration may have been instigated by Fe-rich fluids, during early alteration on the parent body, under low temperatures.


This project is supported by European Social Funds and the Greek State (call code EDBM103). US Antarctic meteorite samples are recovered by the Antarctic Search for Meteorites (ANSMET) program which has been funded by NSF and NASA, and characterized and curated by the Department of Mineral Sciences of the Smithsonian Institution and Astromaterials Curation Office at NASA Johnson Space Center.


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