There is a general consensus that cometesimals formed in an original reservoir, the primordial Kuiper-belt (PKB), between 20-40 au from the Sun (e.g., Nesvorný 2018). From there, they were scattered into the current trans-Neptunian region. The so-called scattered disk is the source reservoir for Jupiter family comets, JFCs (e.g., Duncan & Levison 1997). Crucially, the scattering phase of cometesimals must occur as the gas disk dissipates to prevent the damping of the scattered orbits. This is well described by the final phase of planetesimal-driven migration, particularly Neptune’s migration (e.g., Nesvorný et al. 2018). 

Current planetesimal formation models predict an initial Gaussian size distribution centred around 100 km (e.g., Polak & Klahr 2023). Therefore, the question arises if such an initial size frequency distribution (SFD) of the PKB can reproduce the observational constraints of the current Kuiper-belt. Indeed, we find in Bottke et al. 2023 that the collisional evolution of the PKB based on planetesimal formation models will evolve into a SFD that is consistent with (i) crater SFDs on icy satellites and KBOs and (ii) observed SFDs of populations derived from the PKB (e.g., Jupiter’s Trojans). The craters on icy satellites allow us to infer the SFD of the PKB population scattered onto planet-crossing orbits and those that went to the scattered disk (i.e., source of Centaurs/JFCs).

Because comets originate from that same PKB but are then scattered into the Kuiper-belt, they similarly go through the collisional evolution described above. This raises the question of whether the collisional grinding from 100 km cometesimals to the few km sized comets can preserve the primitive properties of comets (i.a., they contain highly volatile species such as CO and have very low densities). Here, we will present recent simulations showing the degree of compaction and heating due to the collisional evolution.

How to cite: Marschall, R., Morbidelli, A., Nesvorný, D., Bottke, W. F., and Vokrouhlický, D.: Comets are fragments: Determining the degree of processing during the collisional grinding of the primordial Kuiper-belt., Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-862, https://doi.org/10.5194/epsc2024-862, 2024.

Abstract

The super slow rotation of the main belt asteroid 253 Mathilde has puzzled scientists for over twenty years, since the first glimpse by the Near Earth Asteroid Rendezvous (NEAR) spacecraft in 1997 [1]. With a very long rotation period of 418 hours [2], Mathilde is one of the slowest rotators in the solar system. The YORP effect, which was previously suggested to explain the slow rotation of small asteroids [3], is unlikely to spin down a 53-km-diameter object in 4.5 Gyr. Alternatively, the impact mechanism could be an explanation. Mathilde has at least five giant craters that are comparable in size to its radius [4], where the related impacts must have delivered a fair amount of angular momentum to the body and possibly spin it down [5]. It is also surprising that all the giant craters are well preserved without destroying each other, possibly due to the highly porous nature [6]. However, how these giant cratering impacts interacted with the structure of Mathilde, and ultimately evolved it into a slowest rotator, remain unexplained by existing works.

To re-examine these questions, we performed smoothed particle hydrodynamics (SPH) simulations of the giant cratering impacts on a homogeneous porous target, implemented with the Drucker-Prager strength model, Tillotson EOS for basalt, and the P-α porosity model. The simulations include both phases of shock fragmentation and gravitational collapse, which lasts a few hours until the remnant body reaches a stable rotation. The angular momentum (AM) transfer efficiency ζ, the ratio between the rotation change of the target and AM of the impactor’s orbital motion, is then expressed as a function of the impact angle. To model the impact-induced spin history of Mathilde, we generated sequential/virtual impacts  in each of our 10⁵ Monte-Carlo tests (e.g., Figure 1), following the main belt size distribution and the intrinsic collision probability. The statistic results suggest a median spin change ∆ω of 0.6 rev/d, compared with the 2.2 rev/d in a perfectly inelastic case. Note that the spins of large asteroids (>50 km) peak at ~1–2 rev/d, which should represent the primordial spin at their formation. Assuming a initial spin of 1 rev/d, there are only 2% cases ending at < 0.5 rev/d, and 0.02% at < 0.1 rev/d. The slow rotating Mathilde could be in these rare cases, or formed at a much lower spin rate. Future works will investigate the effect of different porous structures like micro/marco-porosity or rubble piles, to help constrain the properties and impact process of Mathilde-like primitive asteroids.

Figure 1. Impact-induced spin history example, where a 53-km-diameter target is impacted by 75 impactors (>0.5 km) at different velocities and angles during 4.5 Gyr.

References

[1] J Veverka, P Thomas, A Harch, et al. NEAR’s flyby of 253 Mathilde: Images of a C asteroid. Science, 278(5346):2109–2114, 1997.
[2] Stefano Mottola, William D Sears, Anders Erikson, et al. The slow rotation of 253 Mathilde. Planetary and Space Science, 43(12):1609–1613, 1995.
[3] P Pravec, Alan W Harris, D Vokrouhlicky, et al. Spin rate distribution of small asteroids. Icarus, 197(2):497–504, 2008.
[4] AF Cheng and OS Barnouin-Jha. Giant craters on Mathilde. Icarus, 140(1):34–48, 1999.
[5] Masahisa Yanagisawa and Sunao Hasegawa. Momentum transfer in oblique impacts: Implications for asteroid rotations. Icarus, 146(1):270–288, 2000.
[6] Kevin R Housen, Keith A Holsapple, and Michael E Voss. Compaction as the origin of the unusual craters on the asteroid Mathilde. Nature, 402(6758):155–157, 1999.

How to cite: Jiao, Y., Asphaug, E., Cheng, B., and Baoyin, H.: Effect of Giant Cratering Impacts on the Slow Rotation of Asteroid Mathilde, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-287, https://doi.org/10.5194/epsc2024-287, 2024.

The detailed observations of the boulder populations on rubble-pile asteroid surfaces such as Ryugu, Bennu or Dimorphos provide insights into their physical properties and the evolution of their host asteroids [1]. Asteroid surfaces evolve through impacts and thermal processing – both mechanisms that may lead to boulder fragmentation (Figure 1). 

Figure 1: a) - c) Images acquired by the DART spacecraft approaching the surface of Dimorphos. Boulders which may have originated as a result of an impact event on the asteroid surface are highlighted in red (NASA/Johns Hopkins APL).

Here, we present numerical simulations using the Bern SPH code [2] to study the impact fragmentation of boulders on asteroid surfaces. Before applying our models to asteroid scales, we successfully validated them by comparing the results with boulder disruptions observed in laboratory impact experiments [3].

In scenarios of impacts on rubble-pile asteroids, our Bern SPH simulations include the displacement and late-stage evolution of the boulder fragments, using a recently developed approach [4,5]. This allows us to determine their final position on the asteroid surface or if they are ejected from the asteroid. 

Our initial results suggest that only for boulders with a relatively low bulk tensile strength (<≈ 10⁴-10⁵ Pa)  would the impact fragmentation allow the fragments to stay on the asteroid surface close to the impact point. For higher boulder tensile strength, the specific impact energies required for boulder fragmentation lead to fragment velocities above the escape speed of small asteroids (of the order of tens of cm/s). Our results thus provide constraints for the boulder properties and implications for the surface evolution of rubble-pile asteroids through impact processes. 

Figure 2: Bern SPH simulation of boulder fragmentation by impacts. The fate of the boulder fragments (final position on the surface or ejection) depends on the boulder strength, impact energy and surface gravity. 

Acknowledgements: S.D.R. and M.J. acknowledge support from the Swiss National Science Foundation (project number 200021_207359). R.-L.B. was funded by NASA New Frontiers Data Analysis Program grant number 80NSSC22K1035.

References:

[1] Ballouz, R-L et al. (2023). “Disrupted Boulders on the Surfaces of Near-Earth Asteroids Bennu, Ryugu, and Dimorphos”. LPI Contributions 2806, p. 2505.

[2] Jutzi, M., Benz, W. & Michel, P. (2008) Numerical simulations of impacts involving porous bodies: I. Implementing sub-resolution porosity in a 3D SPH Hydrocode. Icarus 198, 242–255.

[3] Cline II, CJ et al. (2023). “Using Controlled Impact Experiments to Understand the Effects of Cohesive Blocks on the Cratering Process”. Lunar and Planetary Science Conference.

[4] Jutzi, M., Raducan, S. D., Zhang, Y., Michel, P. & Arakawa, M. (2022). Constraining surface properties of asteroid (162173) Ryugu from numerical simulations of Hayabusa2 mission impact experiment. Nature Communications 13, 7134 

[5] Raducan, SD et al. (2024). “Physical properties of asteroid Dimorphos as derived from the DART impact”. Nature Astronomy 8, pp 445–455

How to cite: Della Moglie, P., Raducan, S. D., Jutzi, M., Ballouz, R. L., Cline, C. J., and Cintala, M. J.: SPH Simulations of Boulder Disruptions: from Laboratory Scale Impacts to Asteroids, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-803, https://doi.org/10.5194/epsc2024-803, 2024.

Introduction  Impact processes play a dominant role in the evolution of the solar system at nearly all size scales, from planet formation all the way to micrometeorite impacts on small bodies. The outcomes of collisions are often parameterized by Q, which is the specific impact energy. The parameter  Q^*_D denotes the catastrophic disruption threshold, which is defined as the specific energy required disperse half of the total mass involved in the collision, leaving remaining mass in the largest remnant. Many studies have attempted to determine  Q^*_D  numerically, demonstrating that it is a complicated function of the target’s physical and material properties (size, density, strength, internal structure, etc.) and impact conditions (impactor size, velocity, angle, etc.) (e.g., Benz and Asphaug 1999; Leinhardt and Stewart 2012; Jutzi 2015; Ballouz et al. 2015; Raducan et al. 2024, among others). Here, we introduce the concept of  Q^*_TD, which is the catastrophic disruption threshold for a natural satellite, accounting for tides from the central body. 
        We will present a suite of hundreds of simulations varying the impact conditions to derive a scaling relationship between the mass of the largest remnant, M_lr, and  Q^*_TD . We will show that the influence of an external tidal potential plays an important role in determining both the mass of the largest remnant and the size-frequency distribution of ejected fragments, even for targets well outside the primary’s Roche limit. Finally, we discuss the implications for this modified impact scaling law for small satellites such as Phobos, Deimos, and the inner satellites of the gas and ice giants.

Methods Most previous studies considered only non-rotating targets in inertial space, although pre-impact rotation can enhance mass loss (Ballouz et al. 2014) and the role of tides is known to play a significant role in some scenarios (Hyodo and Ohtsuki 2014). Accounting for tides in collisional disruptions means there are many more free parameters as the problem is no longer spherically symmetric. To keep the problem computationally feasible, we consider a target with a single set of physical and material properties. Then, we vary the impactor size, velocity, impact angle, impact direction, and the orbital distance of the target. We only consider relatively low-speed impacts (v ≪ 1 km s⁻¹), so both the collision and subsequent gravitational accumulation are modeled using pkdgrav, a N-body code that includes particle self-gravity and handles particle contacts using the soft-sphere discrete element method (Richardson et al. 2000; Schwartz et al. 2012; Zhang et al. 2017).

Our simulations consider a spherical target, consisting of  randomly arranged particles with a size-frequency distribution following a truncated power law with an index of -3. The particles are cohesionless and have frictional angle of . Particles have a density of 2 g cm^-3, which gives the target a bulk density and radius of ρ_bulk ~1.35 g cm and R~10 km, respectively. The target has a surface escape velocity of v_esc ~9 m s⁻¹. The target is placed on a circular orbit around a Mars-mass primary at varying semimajor axes and allowed to settle into an equilibrium before the impact is performed. These simulations assume the target to be tidally locked to the primary, so the target’s spin rate varies along with its orbital distance. The target is slightly smaller in size and less dense than Phobos, however the results can be generalized to any small satellite around a planet.

Preliminary Results We show an example of a head-on collision with an impactor M_imp = 0.1M_targ at an impact velocity of v = 10v_esc ~ 90 m s⁻¹. Snapshots from this simulation at three orbital semimajor axes are shown in Fig. 1. These three semimajor axes correspond to aorb = [1.31, 1.44, 1.97]R_RRL, where R_RRL is the radius of the target’s rigid-body Roche limit. We see a sharp transition between the two innermost cases. When a_orb = 1.31R_RRL, the body first undergoes a catastrophic collision but the disturbs the satellite enough causing to to undergo a complete tidal disruption despite it technically being outside the Roche limit. When a_orb is increased slightly to 1.45R_RRL, the body is safe from disruption and the mass of the largest remnant is M_lr ~0.61M_tot, where M_tot is the combined mass of the target and impactor. As the orbital distance is increased further to a_orb = 1.97R_RRL, the target is able to retain more mass. We attribute this to two factors. Since the target starts in a tidally-locked configuration, a larger a_orb means a slower rotation and therefore a slightly weaker surface gravity. The second, and more important factor, is that the Hill sphere increases with a_orb, as does the escape velocity required to reach the Hill sphere, so less ejecta is able able to escape onto planetocentric orbits for the same impact.

Figure 1: Snapshots from the first 9 hours of three simulations with the same impact conditions at different orbital distances. The impactor is shown in green, traveling from left to right and the target is in white. From left to right, the semimajor axis of the target’s circumplanetary orbit corresponds to 1.31, 1.44 and
1.97 times the rigid-body Roche limit (R_RRL).

Acknowledgments  H.A. was supported by the French government, through the UCA J.E.D.I. Investments in the Future project managed by the National Research Agency (ANR) with the reference number ANR- 15-IDEX-01. P.M. acknowledges funding support from CNES.

References
Ballouz, R.-L. et al. ApJ 789, 158 (2014).
Ballouz, R.-L. et al. Planet. Space Sci. 107, 29–35 (2015).
Benz, W. & Asphaug, E. Icarus 142, 5–20 (1999).
Hyodo, R. & Ohtsuki, K. ApJ 787, 56 (2014).
Jutzi, M. Planetary and Space Science 107, 3–9 (2015).
Leinhardt, Z. M. & Stewart, S. T. ApJ 745, 79 (2012).
Raducan, S. D. et al. PSJ 5, 79 (2024).
Richardson, D. C. et al. Icarus 143, 45–59 (2000).
Schwartz, S. R., Richardson, D. C. & Michel, P. Granular Matter 14, 363–380 (2012).
Zhang, Y. et al. Icarus 294, 98–123 (2017).

How to cite: Agrusa, H. and Michel, P.: The influence of tides on catastrophic disruptions of close-in planetary satellites, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-524, https://doi.org/10.5194/epsc2024-524, 2024.

The hypothesis that the lunar core dynamo once generated a global magnetic field is widely accepted [1–2]. Paleomagnetic analysis of lunar samples shows that the lunar dynamo field operated from about 4.2 Ga to sometime between 1.92 and 0.8 Ga [3–5]. From the orbital observations, the strongest lunar magnetic crustal anomalies are found to concentrate on the lunar farside, and a prominent magnetic low correlates with the nearside Procellarum KREEP Terrane [6].

Impact cratering is a geological process that could either magnetize or demagnetize the lunar crust. Previous studies have investigated the magnetic signatures of impact basins hundreds of kilometers in size [7–9]. Among the largest multi-ring basins, only five Nectarian basins are unambiguously associated with central magnetic anomalies. For the smaller craters, previous studies have detected both magnetized [10] and demagnetized [11] signatures.

This study systematically analyzed the magnetic signatures of all lunar impact craters resolvable by the most recent magnetic field models. We used the locations and crater diameters from [12] which were based on optical images, and this database was supplemented by the peak-ring and multi-ring basins from [13] that were characterized using gravity data. We assigned the crater age using the updated version of the crater database of [14]. The surface magnetic field data from [15] was used as the primary model, and the model of [16] was then used to confirm the magnetic signature of the investigated crater.

The investigated craters in this study were classified into three classes: Magnetized (with central magnetic highs), demagnetized (with the magnetic low interior of the crater rim), and no signal (with no clear magnetic signal). The signal fidelity levels of the magnetized and demagnetized craters were further divided into three levels: Certain, probable, and possible. The two authors conducted the classification independently. To assess the likelihood of correct identification, we made use of two sets of synthetic magnetic field models generated by rotating the geographic coordinate system of the observed fields. We found more false identifications of magnetized craters than real observations when the crater diameter is smaller than 90 km. Therefore, we only report the results from the craters with diameters greater than 90 km.

In total, we analyzed 447 craters. Of these, only 26 and 42 were classified as magnetized and demagnetized craters by at least one of the analysts, respectively. If only considering the craters that both analysts agree upon, the numbers decrease to 10 and 30, respectively.

The debiased number of craters for a given classification was defined as the number of craters using the observed magnetic field models minus the number of craters using the synthetic maps in the given diameter or age class. For a few cases where the number of false identifications is larger than the number of using real detections, we simply set the debiased number as zero. The debiased percentage is simply the sum of the debiased numbers divided by the total number of craters in the given class.

In Figure 1, we plot the average debiased percentages of craters with impact-related magnetized and demagnetized signatures as a function of diameter and age. When only considering the certain and probable fidelity levels, about 1%, 3%, and 14% of the complex craters (90–206 km), peak-ring basins (206–582 km), and multi-ring basins (582–1321 km) are found to have magnetized and demagnetized signatures, respectively. When considering all three signal fidelity levels, the results show the same trend of increasing with increasing diameter but show a higher percentage.

In terms of crater age, when only considering the highest two fidelity levels, we find that about 2% and 3% of craters with pre-Nectarian and Nectarian ages show magnetized signatures, respectively. None of the younger Imbrian, Eratosthenian, or Copernican periods show evidence of craters with magnetized signatures. For the demagnetized class, we find that about 0.3%, 3%, and 16% of pre-Nectarian, Nectarian, and Imbrian aged craters show demagnetization signatures, respectively. These results are compatible with a lunar dynamo operating during at least portions of the pre-Nectarian and Nectarian periods, and then either weakening or ceasing at the beginning of the Imbrian period.

Fig 1. Average debiased percentages of the two analysts for craters with magnetized and demagnetized signatures as a function of (a) diameter and (b) age using the surface magnetic field models. The number above each bar shows the average number of craters in the interval from the two analysts after debiasing.

Lastly, we placed constraints on the mechanisms of generating impact-related magnetic signatures. The excavation of crustal materials and thermal demagnetization can account for the magnetic lows within the crater rim. Shock demagnetization might account for the magnetic lows that extend beyond the crater if an ambient core-generated field was absent when the crater formed. In contrast, the magnetized signatures are likely to be the result of the heated materials cooling in the presence of an ambient magnetic field. The reasons for only a small number of craters with magnetic signatures could potentially be a dynamo that was episodic, frequently reversing, or unstable in intensity and direction. Specific impact conditions (e.g., impact angle and the iron-metal content of the impacting projectile) could perhaps also be required for generating a magnetic signature.

Reference [1] Wieczorek M. et al. (2022) Lunar magnetism, [2] Weiss B. and Tikoo S.  (2014), Science, 346, [3] Garrick-Bethell, I. et al. (2009) Science 323, 356–359, [4] Garrick-Bethell, I. et al. (2017) JGR-Planets, 122, 76–93, [5] Mighani S. et al. (2020) Sci. Adv., 6, [6] Wieczorek M. (2018) JGR-Planets, 123, 291–316, [7] Halekas J. et al. (2003) Meteorit. Planet. Sci. 38, 565–578, [8] Hood L. (2011)  Icarus, 211, 1109–1128, [9] Oliveira J. et al. (2017) JGR-Planets, 122, 2429–2444, [10] Halekas J. et al. (2002) GRL, 29, 23-1, [11] Arkani-Hamed J. and Boutin D. (2014) Icarus, 237, 262–277, [12] Robbins S. J. (2019) JGR-Planets, 124, 871–892, [13] Neumann G. (2015) Sci. Adv. 1, e1500852, [14] Losiak, A. et al. (2015) LPI crater database, [15] Tsunakawa H. et al. (2015) JGR-Planets, 120, 1160–1185, [16] Ravat D. et al. (2020) JGR- Planets, 125.

How to cite: Yang, X. and Wieczorek, M.: Magnetic signatures of lunar impact craters, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-229, https://doi.org/10.5194/epsc2024-229, 2024.

Introduction: The detection of water on the lunar surface by spacecraft flybys and ground observations has raised questions about its origins. While solar wind is likely an important source of H, leading to H₂O and -OH formation, other external processes such as interplanetary dust implantation, micrometeoroid impacts, and dielectric breakdown may also be important in the formation of water molecules. Whilst laboratory observations and studies are informative, they only report end result data, without providing a complete insight on the evolution of the surface composition. In addition, macroscopic simulation models using laboratory equipment are unable to simultaneously resolve chemical reactions at the atomic scale during collisions. In order to better understand water formation and retention, studies at the atomistic level need to be conducted. Molecular dynamics (MD) simulations can capture atomic bonds and molecule formation during different processes, allowing for studies of the atomistic processes underlying water production.

Huang et al. [1], used MD to study water production and retention after a micrometeoroid impact for near surface hydrogen implantation. Their simulations used a reactive force field (ReaxFF) in order to capture bond breaking and formation of H₂O molecules while testing different impact velocities for a 6nm impactor. Their results showed that the optimum velocity water production was 16 km/s while they also noted that even at the nanoscale molecules can still reach the Moons exosphere. Their work was limited to hydrogen at the surface only, focusing on water production during normal micrometeoroid impacts. However, the majority of solar wind protons are distributed throughout the substrate rather than concentrated at the near surface according to the site location on the Moon. Research is needed on the role of implantation depth and incidence angle on the water formation behavior during micrometeoroid impacts. Thus, in this study, we consider water formation and retention during micrometeoroid impacts for different realistic hydrogen implantation depth cases in a silica substrate. Using atomistic modelling we also examine the effect of different micrometeoroid sizes at normal and 30^o incident angles on water production and retention.

Methodology: We conducted MD simulations of micrometeoroid impacts on the lunar surface to better understand the underlying physics of water formation, retention. We used a reactive force field (ReaxFF) potential, originally parameterized to describe a silica-water system and previously used in micrometeoroid and H diffusion simulations in amorphous SiO₂ [1]-[4].

First, a 433 × 433 × 290 ų substrate with a free surface is created, representing a part of the outer surface of a lunar grain. Hydrogen atoms were then placed according to two different depth distribution cases. A near surface distribution is created by placing ~30000 H atom at the surface only and then impacting. A distribution with depth was created by running separate binary collision approximations following best practices for SW impacts [5] and then depositing H in the MD substrate according to these results. Micrometeoroids of 6nm and 9nm in diameter were simulated at velocities ranging from 12-20 km/s and at normal and 30^o degrees angles of incidence.

Results: MD simulations showed that water generation during a micrometeoroid impact is highly dependent on impact velocity, incident angle, and initial depth distribution of H. In the majority of cases there was an increase in H₂O generation with impactor velocity. Interestingly, when the hydrogen is distributed throughout the depth (as opposed to at the near surface) there is a difference in both H₂O production rate and retention. For example, for a 9 nm impactor at 30 degrees there was in increase in retained water by 20% when H was distributed according to SW impacts. In contrast, for this same impactor there was a distinct drop in retained water by 24% when H was deposited at the surface only. We also observe more H₂O ejecta when compared to the distributed H cases. Therefore, whether micrometeorites lead to the formation or loss of H₂O strongly depends on the initial deposition profile of the implanted H. In addition, other ejected atoms and molecules such as H, H₂ and -OH leave the substrate during impact, these ejecta can potentially contribute to the lunar exosphere or be redistributed at nearby sites.

Results also indicate a dependence on the impact energy in the normal direction for the different cases simulated. For example, a 6nm micrometeoroid at 20 km/s and normal impact angle exhibits similar distribution behavior as a 9nm micrometeoroid at 12 km/s and normal impact angle. This indicates a dependence in impact energy as the two cases have close energies of 5 x 10^-14 J and 5.8 x 10^-14 J respectively. In addition, a decrease from initial concentration was observed at the near surface of the substrate, 0 - 40 Å in depth. This is likely due to the high impact energy experienced at the surface that raises the local temperature, breaking the chemical bonds of the surface water. Results point to the importance of considering both the impactor behavior (velocity/size/angle) along with the initial hydrogen deposition when studying water production. Those areas on the Moon which are directly exposed to SW and thus have H deposited at depth are more likely to retain produced water.

References:

[1] Huang, Z. et al. (2021), Geophysical Research 142 Letters, 48(15)

[2] Fogarty, J. C. et al., (2010) Journal of Chemical Physics Vol 132, Issue 17, p. 174704

[3] Morrissey, L. S. et al., (2022) Icarus Vol 379, p. 114979

[4] Sheikholeslam, S. A. et al., (2016) J Mater Chem C Mater Vol 4, Issue 34, pp. 8104–8110

[5] Morrissey L. S. et al. (2023) Planet. Sci. J. 4 67

How to cite: Georgiou, A., Huang, Z., Verkercke, S., Lewis, J., and Morrissey, L.: Water Formation During Micrometeoroid Impact on the Lunar Surface: A Molecular Dynamics Study, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-673, https://doi.org/10.5194/epsc2024-673, 2024.

Introduction

Impacts are one of the most destructive processes in the solar system, and impact craters have been identified on almost every type of solar system body. The lunar surface is covered in craters ranging from 2500 km in diameter, down to sub-millimetre scale. Lunar impact flashes (LIF) are caused by incandescence during an impact, and can be observed by ground based telescopes. Over 650 LIFs have been observed within literature [1,2,3], but despite this large volume of data, only 3 freshly formed craters with documented LIF have been located previous to this work. Such craters are important, as they serve as ground truth data for both the refinement of the luminous efficiency, η, typically taken as between 10^-2 and 10^-4, and for analysing which crater scaling law is most accurate at the scale of the observed craters.

Method

PyNAPLE is software we developed to locate the resultant crater within Lunar epoch, of an observed LIF [4]. Using the >650 LIFs available in literature, we applied constraints to filter out the unconfirmed events, and to prioritise the higher energy events, which had the highest probability of being detected with PyNAPLE. In total, this left 22 LIFs to be processed with PyNAPLE.

Results

After processing the 22 LIF events, there were sufficient LROC images to locate the freshly formed craters for six new events, as well as the three already identified within literature[4,5,6]. For one of these events, two candidate craters were found. Additionally, two unlinked craters were located during the search, however comparing the formation window of these craters to the database of LIFs, no candidate formation events were identified. A selection of six of these craters is shown in Fig. 1.

Figure 1: Six of the craters located from LIF observations by PyNAPLE.

Analysis & Discussion

For each of the 9 craters with known formation event, the likely parent meteoroid stream for each event can be obtained by comparing the LIF location to the meteoroid streams that were active and visible to the impacting location at the time of impact. Identification of the parent stream gives an approximate value for the impactors velocity, impacting angle, and projectile density.

Using the calibrated brightness of each flash, a value for the luminous energy, E_lum, can be obtained for each event. Using an estimate for η, the total kinetic energy of the impactor for each event can also be caluculated, KE = E_lum / η.

The crater scaling laws are several equations which all attempt to relate the kinetic energy of an impactor to the rim-to-rim diameter of the formed crater. As they were mostly derived from explosive tests, their accuracy at the <100m scale of these craters is unknown. Using two of the found craters, which share a parent meteoroid stream and therefore share similar properties, we can evaluate the accuracy of the most popular scaling laws, as shown in Fig. 2. From this comparison, as Shoemaker & Wolfe [7] most closely fits the data, we can conclude it is the most accurate at this range. This equation does not exactly fit, however, and while there are several factors that could contribute to this, such as the estimates for projectile density,target density, impactor velocity, and angle, the single most likely factor is the poorly constrained luminous efficiency.

Figure 2: The comparison of four popular scaling laws, using two of the located craters with same parent meteoroid stream as ground truth.

Under this assumption, a more accurate value for the luminous efficiency can be calculated from the observed craters. This can be done using a rearrangement of the crater scaling laws, to work out the required KE for the observed crater diameter, and subsequently using the observed luminous energy of the impact flash to calculate η = E_lum / KE.

Performing this for each LIF linked impact crater, after outlier removal, produces an average value of η = 0.017. While this is slightly larger than the typically used values of between 10^-2 and 10^-4, the difference is close enough that this could be the result of small inaccuracies in the other parameters used, such as impact velocity, or projectile density. The inaccuary of the scaling law used would also effect this value; as more LIF-linked craters are obtained, and a greater statistical dataset is formed, a re-evaluation of the crater scaling laws, and possibly the derrivation of a new equation, is necessary for this work.

Bibliography

[1] Xilouris et al. (2018) A&A 619, A141 [2] Madiedo et al. (2015) Planetary and Space Science, 111, 105 [3] Suggs et al. (2014) Icarus, 238, 23 [4] Sheward et al. (2022). MNRAS, 514(3):4320–4328 [5] Robinson et al. (2015) Icarus , 252, 229 [6] Robinson, M., (2014) Another New Crater! Webpage: http://lroc.sese.asu.edu/posts/810 [7] Shoemaker & Wolfe (1982) Satellites of Jupiter, 277–339

How to cite: Sheward, D., Avdellidou, C., Delbo, M., and Cook, A.: The Resultant Craters from Lunar Impact Flashes, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-782, https://doi.org/10.5194/epsc2024-782, 2024.

Introduction: Fragmentation of rocky surfaces upon meteor impacts often generates fields of boulders around impact craters. Ejected boulders sometimes possess sufficient kinetic energy to create secondary craters, increasing overall crater density and introducing biases when estimating crater retention ages. Characterizing the kinetic energy of ejected boulders is required to correct such biases in age determinations [1]. However, the vast number of resolvable boulders around single impact craters has rendered morphometric studies of impact-generated boulders challenging. As a result, few boulder fields have been studied [2–11], impeding systematic analyses of the role of, e.g., impactor energy and target properties on spalled boulder properties. To enable more extensive analyses of boulder morphometrics, we developed BoulderNet [12], a machine learning-based algorithm that automatically detects the boulder outlines and characterizes their morphometrics from high-resolution satellite images. Here, we use BoulderNet around over 40 young and fresh impact structures on the lunar surface (i.e., without visible degradation) to better understand the role of the impactor energy and target properties in the boulder ejection process.

Methods: A couple of improvements to the previously published version of BoulderNet [12] were implemented. First, the model architecture was modified from Mask R-CNN [13] to YOLOv8 [14], leading to refined outline detections (thanks to the anchor-free nature of YOLOv8). Second, boulders around lunar cold spots and fresh impact craters (directly relevant to this study) were added to the training dataset. Overall, these changes resulted in the detection of smaller boulder sizes and an increase in recall and precision for most boulder sizes.

This updated version of BoulderNet was then used to investigate the youngest and freshest impact craters on the lunar surface - the so-called cold spots [15]. Cold spots smaller than 400 m in diameter were excluded from this study because most of their associated boulders are close to or below the limit of resolution of LRO NAC images [16]. A total of 42 cold spots were studied - 5 located in mare, 36 in highlands, and one on the floor of a larger impact crater - with diameters ranging from 420 to 2300 m. Boulder outlines were automatically detected within 1–4 radii away from the crater center. Boulder morphometrics (size, aspect ratio, orientation, spatial density, and location) were then automatically extracted and analyzed.

Fragmentation of rocky materials generates fragment-size populations that roughly follow a power-law distribution, N(>D)=C.D^-β, where N is the cumulative number of fragments with diameter > D, and β and C are constants. The slope parameter, β, varies with fragmentation process and history [3]. To further compare boulder populations around cold spots, we fit such a power law for each cold spot (including boulders with diameters greater than 2–4 m depending on image resolution). 

Results: As expected, more boulders are generated around larger impacts (Fig. 1). Furthermore, preliminary results suggest that the slope parameter, β, spans a wider range of values for smaller crater diameters (< 1000 m; β ~ -2.5 to -7.0), but clusters around -2.5 to -3.5 for larger craters (Fig. 2). Surprisingly, the correlation between the size of the largest boulders generated by a given impact and crater size is relatively weak (Fig. 3), especially in lunar highlands.

Fig 1. Number of boulders larger than 4 m as a function of primary crater diameter. Only boulders within 1 radius away from the crater rim are selected. Colors reflect terrain type. Symbol size reflects the resolution of NAC images from which detections were made (larger circle = coarser resolution). 

Fig 2. Slope parameter, β, of boulder populations around lunar cold spots as a function of primary crater diameter. Results from other studies are also shown for comparison. The range of catastrophic disruption for asteroid collisions is highlighted in gray.  

Fig 3. 99.9^th percentile of equivalent boulder diameter as a function of primary crater diameter.

Discussion and Conclusions: The general absence of correlation between the size of the largest boulders generated by a given impact and crater size hints at the importance of fragmentation history and the likely presence of pervasive fractures in the lunar crust, especially in highlands. This interpretation is possibly corroborated by our observation of larger variations in β for lunar craters smaller than a kilometer in diameter, which could be explained by more heterogeneous crustal materials in the first tenths of meters. We emphasize that the role of target lithology (mare vs. highlands) is difficult to assess due to the relatively small number of very fresh impact craters therein. To address this caveat, we will analyze boulders around fresh martian craters in similarly young lava flows.

References:

[1] Melosh (1984) Icarus 59.

[2] Shoemaker (1965) JPL Tech. Rept.

[3] Hartmann (1969) Icarus 10.

[4] Vickery (1986) Icarus 67.

[5] Bart & Melosh (2010) Icarus 209.

[6] Krishna et al. (2016) Icarus 264.

[7] Pajola et al. (2017) Icarus 296.

[8] Pajola et al. (2019) PSS 165.

[9] Watkins et al. (2019) JGR Planets 124.

[10] Pajola et al. (2021) Universe 7, 82.

[11] Mistick et al. (2022) Icarus 376.

[12] Prieur et al. 2023, JGR Planets 128.

[13] He et al. arXiv:1703.06870.

[14] Ultralytics: YOLOv8 (2024), https://github.com/ultralytics/ultralytics.

[15] Williams et al. (2018) JGR Planets 123.

[16] Robinson et al. (2010) Space Sc. Rev. 150.

How to cite: Prieur, N. C., Xiao, Z., Kerner, H., Werner, S., and Lapôtre, M.: Systematic Analysis of Boulder Populations around Lunar Cold Spots, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-1252, https://doi.org/10.5194/epsc2024-1252, 2024.

Introduction

The contribution of secondary craters to the size-frequency distribution of impact craters on planetary surfaces has been debated controversially [1-3], in particular as they have the potential to affect the dating of planetary surfaces via crater counting. The identification of secondary craters is straightforward in proximity to the source crater. Fragments impacting near the primary crater have relatively low velocities and produce irregular and non-circular shaped craters that are shallower than fresh primaries [4]. Here we present the first field of secondary craters found on Earth that allows to study secondary crater formation in detail.

Fig.1 DEM map of southeastern Wyoming, USA, showing the locations of the secondary craters and the trajectory fans.

The secondary crater field in Wyoming

In 2018 a field of small impact structures was discovered in Wyoming in tilted Permian strata of the Rocky Mountains Front Range system [5] (Fig. 1). The confirmation of the impact origin of these craters was based on the documentation of shock features in quartz grains. At first, [5] interpreted this cluster of craters as a crater strewn field that was formed by the atmospheric fragmentation of a single meteoroid. The maximum theoretical spreading of such strewn fields is restricted and should not exceed more than one kilometer perpendicular to the trajectory [6]. Since the discovery of the first craters, additional craters have been identified in the same strata in an area that measures  90 by 40 km (Fig.1) [8] and excludes a meteoroid-break-up scenario. Instead it was shown that the craters represent secondary craters.

Secondary crater morphologies and trajectory reconstruction

The diameters of the 31 confirmed and 188 potential crater structures range in size from 10 m to almost 80 m. Many craters are circular, but their ellipticity can reach up to 1.7 (Fig. 2). All craters are exposed at the same stratigraphic level so are assumed to be of equivalent, 280 Myr age. The freshest structures contain steep crater walls, raised rims with overturned ejecta flaps and remains of the proximal ejecta blankets. Estimates of depth-diameter ratios are 0.1 and less. We observe irregular crater clusters and crater chains, where craters overlap [7]. Occasionally relics of herringbone patterns were observed (Fig. 2). We used the orientation of the long axis of elliptical craters and the crater chain alignments for trajectory reconstruction. Criteria for up-range and downrange distinction included overturned ejecta flap downrange and V-shaped herringbone patterns (Fig. 2). The craters define fan-like corridors of trajectories for each crater field (Fig. 1). Tracing back the trajectories allowed the location of the primary crater to be estimated in the area of the intersection of the corridors, centered at 41°28’N and 103°59’W [7]. All discovered secondary craters occur at a distance of 150+/-50 km to the proposed primary crater. We calculated the ballistic paths of ejecta taking into account aerodynamic drag. Trajectories of ejecta with 1, 2 and 4 m radius were modeled with ejection angles ranging from 30 to 60 degrees and initial speeds of 1, 2 and 4 km/s.  Modeling showed that the impacts occurred at around 700-1000 m/s with impact energies of about 12 to 400 GJ. Such impacts are capable of generating craters of 8–55 m in diameter and may generate small volumes, where shock pressures are sufficient to form shock microstructures in quartz [7].

Fig.2 Drone image of a the southern part of the Sheep Mountain secondary crater field.

The primary crater

The possible location of the primary crater is situated in the Northern Denver basin, where 280 Myr old strata are deeply buried beneath younger beds. Ejecta scaling suggests that the primary crater may have a diameter of 50-65 km. We will analyze newly released high-resolution Bouguer gravity data to constrain location and size of the possible crater. The borehole I-35 Hawk Fee, situated in the area of interest, shows some breccia layers at 3023-3066 m depth at the respective stratigraphic level, but shock features could not unequivocally found until now.

A chain of primary craters in Wyoming?

Recently, a new impact structure has been discovered in the Bighorn Basin of NW Wyoming, named Jake Seller Draw impact structure [8] some 300 km NW of the secondary crater field. The 4.3-km-diameter structure was recognized as a seismic disturbance at a depth of ∼6.5 km making it the most deeply buried impact structure known on Earth to date. Shock features were detected from boreholes drilled into the center of the structure and its ejecta. Seismo-stratigraphy and drilling showed that the crater also formed 280 m.y. ago. The coincident stratigraphic age of the Jake Seller Draw impact structure with the ages of the Wyoming crater field, its SE-NW alignment with the fan of secondary craters and the proposed source crater of the Wyoming crater field, sug­gest a causal relationship between the Jake Seller Draw structure and the Wyoming crater field [8]. In addition to that, the buried 7 km diameter Cloud Creek impact crater [9] lies exactly on the same trajectory some 110 km NW of the secondary crater field (Fig.3). Its published age apparently rules out a connection to the other craters, but in a personal communication with the first author of [9], the age of Cloud Creek was presented as not robust. We propose that the Wyoming impact event comprised of three asteroids and impacted the Earth along a SE to NW impact trajectory. The largest of the primary craters formed the secondary crater field that is preserved in downrange direction.

Fig. 3 Map of the primary and secondary impact structures in Wyoming

References

[1] Ivanov,B.A. 2006, Icarus, 183, 504–507. [2] McEwen,A.S. & Bierhaus,E.B. 2006, AREPS, 34, 535–567. [3] Zanetti,M. et al. 2017, Icarus, 298, 64–77. [4] Pike, R.J. & Wilhelms, D.E. 1978, Proceedings, LPSC 9, 907–909. [5] Kenkmann, T. et al. 2018, Scientific Reports, 8, 13246. [6] Artemieva,N.A. & Shuvalov,V.V. 2001, JGR 106, 3297–3309. [7] Kenkmann,T., et al. 2022, GSA Bull. https://doi.org/10.1130/B36196.1. [8] Sturm,S., et al. 2024, GSA Bull. https://doi.org/10.1130/B37164.1. [9] Stone,D.S. & Therriault,A.M. 2003, MAPS, 38, 445–455.

How to cite: Kenkmann, T., Sturm, S., Müller, L., Fraser, A., Cook, D., Sundell, K., and Rae, A. S. P.: Secondary cratering: a case study on Earth , Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-1067, https://doi.org/10.5194/epsc2024-1067, 2024.

Introduction: Reconstruction of paleoenvironments, especially aquatic, is important in search for potential Life habitats, e.g. on Mars. Remote sensing, preferably combined with rovers, give information on telltale geochemistry and landforms. However, this may be hampered by extensive surface erosion. Impact craters offer efficient sediment traps relatively protected from erosion. In aquatic (“marine”) environments, water may rush back into the crater during early modification generating “resurge deposits” [1]. Resurge deposits in drill cores from several impact craters show a direct relation between average clast frequency per meter (‹N›), event magnitude (i.e., projectile diameter, d) and target water depth (H) as ‹N›=-15(d/H) +100 for a “low” (e.g. moat) location, or  ‹N›=-13(d/H) +150 for a “high”, more turbulent position (e.g., near rim) [2;3]. This implies that any of these factors can be calculated if the other two are known. This was applied on Rochechouart impact structure that was debated if marine-target or not due to lack of marine sediments in the area [4]. Here, we study drill cores from the interior of the today burried Decorah crater, Iowa (43°18′ 49″N, 91°46′19″W), to learn more about its paleoenvironment.

The 5.6km Decorah crater seemingly lacks a central uplift expected for its size and target, and is suggested to be marine-target [5]. Target rocks were Upper Cambrian and Lower Ordovician sandstones and dolomite [5;6]. Earliest post-impact infill by marine Winneshiek Shale places the impact in Darriwilian[7]. Decorah is extensively drilled, but only two yielded cores useful for this study; the ~33m “H2” core (0.32km inside the eastern rim), and the ~28m “CS1” core (approx. halfway between south-western rim and crater center) [cf. 5].

Methods:  We log the polymict breccias of the 5cm in diameter H2 and CS1 cores as their observed grading [5] suggests resurge deposits [cf. 1]. At the time of writing, only H2 is fully logged and interpreted. CS1 will be presented at the conference. The logged section of H2 spans 9.6m from the bottom of the core until the clasts sizes become to small (<5mm) for the method, i.e., the line-logging technique previously used for Lockne, Tvären, Chesapeake Bay, Wetumpka, Flynn Creek, Chicxulub, and Rochechouart [cf 1;4;8;9;10;11]. For most of these craters it was possible to work directly on the cores, but here we have, similarly to the Chicxulub study [8], used digital core photos and the software JMicroVision 1.2.7. The core recovery in the logged intervals of H2 and CS1 was >98%. Size sorting calculates as the standard deviation of the clast size per length unit (here half a meter). Roundness is estimated with a grain shape comparator [cf. 12]. Matrix- or clast support of each clast is based on contact with adjacent clasts and plotted as a ratio per length unit. Alltogether, plotted values show relative variations indicating trends, not absolute values. In addition, clast colors and textures were noted to enable an association with the target stratigraphy [e.g., 4; 8]

Results and discussion: A selection of the 416 clasts examined in H2 is shown in Fig 1. The plots in Fig. 2 support the normally graded appearance of the breccia as noted by [5]. When comparing with logs from aforementioned craters, where cores have reached through resurge deposits into underlying slump and avalance breccias, it is evident that H2 ends within resurge deposits.  Nevertheless, the cored interval shows similar trends as several of the other craters. We primarily compare with Rochechouart. There, the sequence is subdivided into 6 intervals [see fig. 4 in 4]. Especially intervals 3–6 in Rochechouart show similarities to H2, whereas intervals 1-2 likely were not cored at Decorah, but would be expected as they represent the inevitable initial stages of the resurge. Intervals 3 and 4 at Rochechouart (“outwards passage of anti-resurge” followed by development of “body of standing water”) are characterized by slight upwards increase in clast frequency (until 32m in H2) followed by slight decrease (until 30.25m in H2). In the same interval at Rochechouart, the clast size and size sorting remained stable (up to 30.25m in H2), but accompanied by an increase in clast angularity (up to 30.25m in H2). With interval 5 at Rochechouart, a new pulse in transport energy caused a strong increase in clast size and drop in size sorting, as well as slight increase in roundness (29.25–30.25m in H2). This is then in Rochechouart followed by interval 6 that includes an increase in clast frequency and size sorting, and a generally normal grading towards the top (29.2m and upwards in H2). There is also a general upwards decrease in roundness similar to H2. This interval is interpreted to represent settling of material in a now almost water-filled crater, with seiches causing minor repeated beds (e.g., at 26m and 27m in H2)

H2 shows an obvious change in clast lithologies at interval 29.25–30.25m (Fig. 2). The “white” and“dark green” fragments are followed by “dark brown-red”, “dark grey” and “light brown”, whereas “terracotta colored”, “light tan-grey” and “light grey” appear throughout the logged sequence, possibly as the basement clasts do at Rochechouart.  

Two stratigraphic intervals can be correlated with clast types; The Upper Cambrian Lone Rock Formation, which contains glauconitic and feldspathic sandstone with some beds of dolomite and green-gray shale, and the Lower Ordovician Oneota and Shakopee formations, which include beds that have been stained red owing to their relationship to the truncating, inter-regional unconformity at the base of directly overlying St. Peter Sandstone [6]. These two intervals produce the more easily traceable greenish and reddish clasts.

The calculated ‹N› = 43. The online "Earth impact effects program" calculator gives a 350m projectile diameter (d) for a final crater diameter of 5.8km. This results in 92m target water depth (H). This seems reasonable considering that resurge must have been able to overcome the elevated rim [cf. 1]. Likewise, certain benthic fossils in the Winneshiek Shale indicate a deposition within the photic zone (i.e., <200m) [cf. 13].

How to cite: Ormö, J., Sturkell, E., and King Jr., D. T.: Assessing target water depth and paleoenvironment at the Decorah impact structure, Iowa., Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-48, https://doi.org/10.5194/epsc2024-48, 2024.

On the 24th of December 2021, an impact produced a 150-m-wide crater on Mars (Fig. 1) and excavated subsurface water ice [1, 2]. The crater is located at 35.10°N, 189.82°E in Amazonis Planitia, and is the largest crater observed to have formed since MRO (Mars Reconnaissance Orbiter) began observations in 2006. The crater presents the lowest-latitude observation of subsurface ice exposed by an impact on Mars [12]. The ice is seen in the proximal ejecta, <700 m from the crater rim, with the highest concentration much closer to the crater rim [2].

The impact produced a magnitude-4 seismic event detected by the InSight (Interior Exploration using Seismic Investigations, Geodesy and Heat Transport) lander, 3500 km away [1]. Here we compare observations with numerical impact simulations, to constrain impact parameters and target structure, specifically the location of the pre-existing subsurface ice. Our results will provide important constraints for the climate history on Mars and understanding the properties of impact-generated seismic waves. 

Methods

We use iSALE3D shock physics code [3, 4] to simulate several scenarios that could have produced the “Christmas Eve crater”, informed by statistical analysis of impact parameters necessary to form a 150-m crater [1]. As asymmetry of the blast zone and the presence of an up-range ejecta exclusion zone suggest a highly oblique impact trajectory, we consider impact angles between 15–30°. We assume a spherical impactor of diameter 5–7 m striking the target at 12 km/s. We implement a two-layer target configuration. A fractured bedrock layer, modelled as 25% porous basalt (bulk density 2150 kg/m³), is overlain by a 50% porous basaltic regolith layer (bulk density 1430 kg/m³). Both layers are defined using the Tillotson equation of state for basalt [5], ϵ-α compaction model [6, 7] and the ROCK strength model [8]. We investigate three thicknesses of the upper layer: 10, 15 and 20 m, guided by observations of other craters in the area [2]. 

The state and position of ejecta are tracked using Lagrangian tracer particles throughout the simulation, and then projected to their final locations assuming ballistic trajectories [9]. We focus on the most proximal ejecta blanket, within 1 crater radius (75 m) of the crater rim, as it contains the highest concentration of visible ice. As we do not explicitly model ice as a separate material, we assume temperature and pressure thresholds of 0°C and 10 MPa, respectively, as the melting point and unconfined compressive strength of ice [10]. We identified ice patches visible around the crater using the quantitative multi-spectral method described in [11] to produce a map of icy ejecta for comparison with our simulations. We use this map to identify overlapping simulated ejecta particles and trace them back to their original positions, producing a possible pre-impact distribution of subsurface ice. 

Results and Discussion

Our simulations produce craters with rim diameters of 130–160m, consistent with observations. The best match to the observed crater morphology is achieved by impacts into the 15-m-thick regolith layer. The simulated craters in this target configuration are 17–19 m deep (depending on impact angle, measured from the pre-impact surface), consistent with the observed crater (17 m), though have steeper walls. This could be explained by later-stage crater modification, which is not simulated here. Fig. 2 shows an impact scenario of a 5.2-m-wide impactor at 30°, compared with the observed crater profile. 

Our results suggest that the proximal ejecta blanket originated from depths <12 m. The ejecta that have experienced temperatures and pressures below the thresholds described above (‘ice-compatible’) originated from even shallower depths, <10 m, and from 30–60 m radially away from the crater centre (Fig. 3). No ejecta with preserved ice originates from closer to the crater centre owing to higher temperatures and pressures closer to the impact point. 

For this scenario, 96% of ice-containing pixels overlapped with a simulation tracer particle. When projected to their pre-impact positions, the overlapping ejecta imply a discontinuous ice distribution under the pre-impact surface (Fig.3b), consistent with the heterogeneous distribution of subsurface ice found at other ice-exposing impact sites at higher latitudes [12]. 

Conclusions

Our iSALE3D simulation results suggest the presence of a stronger bedrock layer 15 m beneath the surface, overlain by a porous regolith layer. We find that the ice blocks visible in orbital images most likely originated from shallow depths <10 m. Our results also suggest that the ice was laterally discontinuous across the pre-impact target. 

Figure 1: (a) CTX image of the Christmas Eve crater (image ID: K18_060561_2175_XI_37N170W). (b) HiRISE image of the crater and proximal ejecta (image ID: ESP_073077_2155_COLOR). Black circle marks the approximate crater rim. Dashed lines indicate the location of cross-sections in Fig. 2.

Figure 2: Depth profile of a simulation in (a) downrange and (b) cross-range direction, compared with the observed profiles (dashed lines) measured along the dashed lines in Fig. 1.

Figure 3: (a) Simulated ice-compatible ejecta at their final projected landing locations (light blue), the observed ice locations (dark blue squares) and observed ice pixels overlapping with simulated tracers (red). (b) Pre-impact locations of simulated ejecta overlapping with observed ice locations (red). Black circle marks the approximate crater rim. 

References: 

[1] Posiolova, L. V. et al. (2022) Science (New York, N.Y.) 378:412–417. 

[2] Dundas, C. M. et al. (2023) Geophysical Research Letters, 50. 

[3] Elbeshausen, D. et al. (2009) Icarus, 204:716–731.

[4] Elbeshausen, D. & Wünnemann, K. (2011) Proceedings, 11th Hypervelocity Impact Society Symposium.

[5] Tillotson, J. H. (1962) Report No. GA-3216, General Atomic, San Diego, CA,43. 

[6] Wünnemann, K. et al. (2006) Icarus, 180:514–527. 

[7] Collins, G. S. et al. (2011) International Journal ofImpact Engineering, 38:434–439. 

[8] Collins, G. S. et al. (2004) Meteoritics and Planetary Science, 39:217–231. 

[9] Raducan, S. D. et al. (2019) Icarus, 329:282–295. 

[10] Durham, W. B. et al. (1983) Journal of Geophysical Research: Solid Earth, 88:B377–B392. 

[11] Rangarajan, V. G. et al. (2023) Icarus,115849.

[12] Dundas, C. M. et al. (2021), Journal of Geophysical Research: Planets, 126(3), pp. 1–28.

How to cite: Wojcicka, N., Collins, G. S., Rangarajan, V. G., Dundas, C. M., and Daubar, I. J.: Oblique Impact Modelling of the Crater Formation and Icy Ejecta of the "Christmas Eve Crater" on Mars. , Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-771, https://doi.org/10.5194/epsc2024-771, 2024.

Introduction:

The surface and crustal structure of the terrestrial planets in the inner solar system have been influenced by large and energetic impact events. Impact basins have long been recognized and studied through satellite images, topographic data, and gravity data.

Peak-ring basins are characterized by a rim crest and an interior peak ring, while multi-ring basins are larger and defined by having additional concentric topographic rings (e.g., [1]). Peak-ring and multiring basins are widespread on the terrestrial planets and can be characterized by their gravity signature. GRAIL data showed that large lunar basins are characterized by a central gravitational anomaly. The size of this gravitational anomaly corresponds closely to the diameter of the inner peak-ring of a basin, while the main ring is approximatively twice the diameter of the peak ring [2]. This allowed to confirm the existence of previously proposed basins, to correctly identify which ring is the main crater rim, and to detect new basins that were not yet identified.

In this work we present an improved techniques based on the analysis of gravity and crustal thickness data to estimate the inner ring and rim crest diameters. This technique expands upon the work of [2] and allows us to better identify highly degraded basins. From this analysis, we quantify how lower resolution gravity and crustal thickness datasets (such as for Mars and Mercury) might bias the peak ring and main rim diameter estimates.

Methods:

In our approach, we first quantify the regional value of the Bouguer gravity anomaly and crustal thickness, which is defined as the average value obtained from azimuthally averaged profiles in the radius range 1.5D to 2D, where D is the crater diameter. The diameter of the Bouguer gravity high, as well as the diameter of the crustal thickness anomaly, were then estimated as the radius where the profiles first intersect the background regional values. After the initial estimate of D was obtained, the procedure was iterated until there was no change in the obtained diameters.

We tested this method using Bouguer gravity data for certain lunar peak-ring and multi-ring basins (see table 1 in [2]), by considering the spherical harmonic degree range from 6 to 540 (which removes the effect of the hemispheric asymmetry and the South Pole–Aitken impact). We then filtered the data using the spherical harmonic degree range 6-49 in order to simulate the lower resolution of the Mars gravity models (e.g., GMM-3, [3],[4]). We then used the same approach using crustal thickness maps derived after GRAIL [5], both for the degree ranges 6-310 and 6-46, to simulate the loss of spatial resolution of Mars. Uncertainty estimates were obtained for the crustal thickness and the Bouguer anomaly diameter by considering the ±1σ values for the background values in the spatial range of 1.5D to 2D.

In Figure 1 we show an example for the Humboldtianum basin, which has a peak-ring diameter of 322 km (5.3° in angular radius). Our method gives a result of 323.1 ± 4.4 km from the Bouguer gravity data (using the degree range 6-540), and 316.3 km ± 4.9 km from crustal thickness data (using the degree range 6-310).

Conclusions and future work:

When considering the highest spatial resolution of the Bouguer gravity data and crustal thickness maps, our method properly detects peak-ring or inner ring sizes for lunar basins with main rim diameter greater than 250 km (i.e., for inner ring diameters greater than about 110 km). Nevertheless, when considering filtered versions of these datasets that correspond to the effective spatial resolution of the Mars gravity models, only basins with rim crest diameters greater than about 450 km can be detected with acceptable accuracy. Regardless, these results confirm a roughly one-to-one relationship between the Bouguer anomaly diameter and the inner peak-ring diameter of lunar basins, as well as between crustal thinning size and peak-ring size (Figure 2).

We first plan to apply this approach to the Moon in order to reassess the impact basins database of [2]. Following this, we will apply the same methodology to Mars to provide a consistent database of Martian basin sizes. Previous databases for Mars suffer from a difficulty of detection as a result of sedimentary and erosive processes, and also an imperfect understanding of their crustal thickness and gravity signatures that was only elucidated by the GRAIL mission. Future analyses will be applied to the planet Mercury. Results from these analyses will allow one to better constrain the impact rate during the early solar system.

Figure 1. Bouguer gravity anomaly (left) and crustal thickness (right) of the Humboldtianum impact basin.  Images of these datasets are shown in the top row, and azimuthally averaged profiles are shown in the bottom row. Black dashed lines represent the peak-ring or inner ring radius (km), while black solid lines denote the main rim radius (km). Regional values of the Bouguer gravity anomaly and crustal thickness are indicated by red dashed lines, while Bouguer anomaly and crustal thinning diameters are shown with red solid lines. 

Figure 2. Bouguer anomaly diameter (left) and crustal thinning diameter (right) versus peak-ring or inner-ring diameter (km) for certain lunar peak-ring and multi-ring basins. Basins include Schwarzschild, d’Alembert, Milne, Bailly, Planck, Schrödinger, Mendeleev, Birkhoff, Lorentz, Vaporum, Korolev, Moscoviense, Crüger-Sirsalis, Grimaldi, Apollo, Hertzsprung, Freundlich-Sharonov, Humboldtianum, Coulomb-Sarton, Humorum, Smythii, Nectaris, Orientale, Crisium, Imbrium. Red dashed lines indicate a 1:1 ratio.

References:

[1] Baker D. M. H., et al. (2011). Planet. Space Sci.

[2] Neumann G. A., et al. (2015). Sci. Adv.

[3] Genova A., et al. (2016). Icarus.

[4] Wieczorek M. A., et al. (2022). JGR: Planets.

[5] Wieczorek M. A., et al. (2013). Science.

Acknowledgements: We gratefully acknowledge funding from the Italian Space Agency (ASI) under ASI-INAF agreement 2017-47-H.0.

How to cite: Buoninfante, S., Wieczorek, M. A., Galluzzi, V., Ferranti, L., Milano, M., Fedi, M., and Palumbo, P.: Quantifying the size of impact basins on the Moon and Mars., Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-484, https://doi.org/10.5194/epsc2024-484, 2024.

Introduction: The prevailing theory for the formation of Earth’s Moon is the Giant Impact hypothesis, in which a collision between the young Earth and a massive body ejects material into a circumplanetary disk from which the Moon then forms. In the leading version of this hypothesis, called the “canonical” Moon-forming impact, a Mars-sized impactor (Theia) collided with the Earth in an oblique impact at roughly the mutual escape velocity of the bodies. Prior simulations suggest that about one lunar mass of material is ejected into orbit, from which the Moon later accretes. A successful Moon-forming impact must be consistent with the following observations: 1) It must explain the angular momentum of the Earth-Moon system, 2) eject enough mass into orbit to allow the formation of the Moon, 3) the disk must be depleted in iron compared to the Earth and 4) the isotopic composition of the disk and Earth’s mantle must be similar. To date, no scenario has been found that satisfies all constraints without requiring either special initial conditions (isotopic similarity between target and impactor) or post-impact processes (removal of angular momentum, post-impact mixing). Studies investigating the formation of the Moon usually focus on meeting certain observational constraints and therefore a limited region of the parameter space. A systematic investigation of the parameter space of potential Moon-forming impacts has not yet been performed.

Methods: We present a systematic survey of the parameter space for Moon-forming giant impacts. This study consists of 7649 collision simulations which cover a wide range of the initial angular momentum, the impact velocity, and the impactor-to-target mass ratio. We also include pre-impact rotation for both the target (proto-Earth) and the impactor, where the rotational angular momentum can be either aligned or anti-aligned with the orbital angular momentum. The simulations are carried out using the 3D smoothed-particle hydrodynamics (SPH) and gravity code Gasoline, with numerical improvements for giant impact simulations. The colliding bodies are modelled as a two-layer sphere, using the ANEOS equation of state for both the iron core and the rocky mantle. Bodies with pre-impact rotation are created by evolving the particle representation in a co-rotating frame with an angular velocity that is slowly increasing until it reaches the desired value and then transferring it to the stationary frame. The collision outcomes are then analyzed using a novel disk finder algorithm to distinguish between the planet and the circumplanetary disk. From this we then calculate the bound angular momentum, the disk mass, the disk iron mass fraction, and the mixing between impactor and target material in the collision.

Results: The set of simulations presented here, which systematically samples the pre-impact parameter space, produces a very diverse set of outcomes (see Figure 1), including disk masses between 0-6.71 M_L (lunar masses). Notably, no simulation was identified that could satisfy all four of the known constraints.

Without pre-impact rotation, no collision produced a sufficiently massive disk (M_d > 2 M_L) below a bound angular momentum of 2 J_EM (see center frame of Figure 2). Collisions with pre-impact rotation can fill some parts of the post-impact parameter space that are not populated by collisions without pre-impact rotation. For example, massive disks can be created by collisions with very low impactor-to-target mass fractions (gamma < 0.1), and we find excellent mixing for all impactor-to-target mass fractions, but massive disks and good mixing are limited to impactor-to-target mass fractions >= 0.5 and the scenario with fast counter-rotating target and very low impactor-to-target mass fraction. By adding pre-impact rotation, we also find results that satisfy both the constraints on disk mass and angular momentum. However, in those cases the disk is predominantly composed of impactor material and therefore they do not satisfy the mixing constraint. Nonetheless, we find plenty of cases that satisfy any combination of three out of the four constraints.

Around 3% of the simulation results show a massive bound fragment (between 0.5 and 1.5 M_L) orbiting in the disk. These fragments either form directly by tidal deformation of impactor material and material excavated from the target or by fragmentation of spiral arms. Such direct formation of (proto) satellites can occur for a wide range of impact conditions. Their composition can be anything between target to impactor dominated and we also find cases with excellent mixing. The fate of such fragments is unclear as they could either be eroded by tidal forces or serve as a seed for the formation of the Moon. In the latter case, constraints on the disk mass and the mixing parameter may change, as this may increase the accretion efficiency. As the accreted material does not necessarily mix with the deeper layers of the Moon, this also allows for differences in the composition between the surface and the deep interior of the Moon.

Conclusions: In our systematic survey of the Moon-forming giant impact parameter space, we find no single case that can satisfy all observational constraints without requiring very specific pre-impact conditions or post-impact processes. We also find that the canonical Moon-forming impact cannot produce sufficiently massive disks without significant pre-impact rotation. A modified canonical impact with a counter-rotating target can produce massive enough disks but these disks are then strongly impactor dominated. In general, collisions with pre-impact rotation can produce massive disks at much lower angular momentum, allowing to reconcile the angular momentum and disk mass constraints. The presence of gravitationally bound fragments found in a small fraction of cases could significantly alter the disk mass and mixing constraints; as a result, the community may be forced to reconsider the currently accepted set of constraints.

Figure 1: Post-impact variables of all 6247 simulations that result in a merger.

Figure 2: Top panel of Figure 1 split into the 9 different spin configurations.

How to cite: Meier, T., Reinhardt, C., Timpe, M., Stadel, J., and Moore, B.: A Systematic Survey of Moon-Forming Giant Impacts, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-74, https://doi.org/10.5194/epsc2024-74, 2024.

Introduction: Planetary formation models suggest that Earth experienced multiple high-energy impacts. They can produce substantial melt in the proto-Earth’s silicate mantle, possibly forming a global magma ocean. Mixing of the impactor’s metallic core into the molten Earth's silicate mantle controls the chemical equilibration between metal and silicates, which defines the respective compositions of Earth's core and mantle. Previous studies explore mixing upon large impacts either with numerical modelling or with analog laboratory experiments. Numerical simulations are efficient in reproducing the shock physics of hypervelocity impacts. However, their spatial resolution is limited and does not allow for reproducing the turbulent features that are responsible for metal-silicate mixing in a magma ocean. On the other hand, liquid impact experiments that do produce small-scale mixing and turbulence are subsonic: they neglect compressibility effects.  Here, we investigate the degree of mixing upon impact by coupling various results from fluid impact experiments [1] and numerical modelling extending the crater depth from experiments to supersonic conditions [2].

Methods: The laboratory experiments used to explore mixing and extend to the supersonic regimes consist in a volume of fluid impacting into a tank of water. The denser impactor fluid is dyed, enabling for the optical estimate of the volume of the evolving sinking plume right after the impact. The volume of the plume gives, once corrected from the volume of the impactor material, an analogue of the volume of entrained silicates upon a given collision. We have extended laboratory results on the crater growth from subsonic to supersonic conditions using the grid-based Eulerian shock physics code iSALE [3,4,5,6] to simulate fluid impacts. The scaling-law that is produced [2] to extend laboratory experiments results to supersonic, hypervelocity conditions, is used here to be further applied to the mixing. The experiments from [1] suggest that the impact between an impactor of density ρ_i and a lighter target of density ρ_t  involves two main stages: the opening of the crater at early times and the fall of central jet that had been formed from the collapse of the crater. That latter event controls the release of the impactor material into the target. At later times, buoyancy forces become important, which controls the sinking of the impactor material into a turbulent thermal descending into the target. The competition between total buoyancy and the momentum of the collapsing jet controls the dynamics of the descending thermal [1], hence the mixing between impactor material and surrounding target silicates. It has been showed that the jet height scales as the maximum crater depth  and that the maximum jet volume scales as the maximum crater volume [7]. We use these scalings, coupled to those regarding mixing from [1] and those on the extension of maximum crater depth from experiments to supersonic conditions from [2] to extend the mixing results from laboratory experiments to hypervelocity regimes accounting for both the effects of the Froude number (measure of the importance of the impactor kinetic energy to its gravitational energy at impact) and the Mach number (impact velocity to sound speed ratio).

Results: Figure 1 shows the mass of silicates that is entrained in the thermal prior to its descent to the magma. It provides a direct measure of the so-called mixing between metal and silicates during a large impact into a magma ocean. We find that the Mach number decreases the estimates of metal-silicate mixing upon impacts derived from laboratory experiments. The extent of its effect however depends on the other impact parameters such as the impactor size. The larger the impactor compared to the target size, the larger the effect of the Mach number. For M>3, the mixing estimates can be underestimated by more than a factor 3 if not accounting for the shock.

Figure 1. Mass of silicates mixed with metal by the impact stage, prior to the descent into the magma, Δ, as a function of the impactor to target radius, R/R_t , varying the Mach number, M, for two cases: a) impact velocity is escape velocity, U = Ue  and b) impact velocity is twice larger than the escape velocity, U = 2 Ue. Circle: impactor of 100 km in radius onto an Earth-sized target. Diamond: canonical Moon-forming scenario with a Mars-sized impactor and U = Ue [8]. Square: Moon-forming scenario with an impactor mass 20 times smaller than the target onto a fast-spinning
Earth at U = 2 Ue [9].

Discussion: We have studied the statistics on how often in typical classical scenarios of accretion such collisions would occur for Earth analogs through its growth history. We use collision files from N-body simulations in the Grand Tack scenario [10] and estimate that, depending on the sound speed of the impacted material, 24% to 74% of the total amount of collisions endured by Earth analogs may occur at M>3. If considering only giant impacts, as these are those when the effect of the Mach number is the most significant, this drops to 4% to 28%. The effect of the Mach number on the mixing, however not extreme, may need to be accounted for when following the metal silicates reequilibration through an entire stage of planetary formation.  

Acknowledgments: We gratefully acknowledge the developers of iSALE-2D, including Gareth Collins, Kai Wünnemann, Dirk Elbeshausen, Tom Davison, Boris Ivanov and Jay Melosh. This work was funded by the Deutsche Forschungsgemeinschaft (SFB-TRR170, subproject C2 and C4).

References: [1] Landeau M. et al. (2021) Earth Plan. Sci. Lett. 564, 116888. [2] Allibert L. et al. (2023) JGR: Planets 128 (8), e2023JE007823. [3] Collins. G. S. et al. (2004) Meteoritics & Planetary Science 39:217-231. [4] Wünnemann K. et al. (2006) Icarus 180:514-527. [5] Elbeshausen, D. et al. (2009). Icarus, 204(2), 716-731. [6] Elbeshausen et al. (2011). Proceedings of 11th Hypervelocity Impact Symposium (HVIS), Fraunhofer Verlag. [7] Ghabache É. et al. (2014) . Fluid Mech. 761, 206–219. [8] Canup R. M. (2004) Icarus 168 (2), 433–456. [9] Ćuk, M. and Stewart S. (2012) Science 338 (6110), 1047–1052. [10] Jacobson S. A. and Morbidelli A. (2014) Phil. Trans. Roy. Soc. A 372, 20130174.

How to cite: Allibert, L., Landeau, M., Röhlen, R., Maller, A., Nakajima, M., and Wünnemann, K.: Metal-silicate mixing upon Giant Impacts into magma oceans, Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-1126, https://doi.org/10.5194/epsc2024-1126, 2024.

1. Introduction

The launch of the Hera mission is scheduled for October 2024 [1]. The target of the mission is the binary Didymos system, especially the secondary Dimorphos, which was impacted by the DART mission on 27^th September 2022 at 01:15 am (CEST) [2]. The mission shall provide a detailed characterisation of the two bodies and discover what happened to Dimorphos after the impact.  To better understand the cratering process, numerical impact simulations have been conducted to analyse the effects of various target properties, (e.g. as strength and porosity), interior structures and exterior boulders, on the outcome of the cratering event [e.g. 3-5]. Nevertheless, these shock physics codes need to be validated against laboratory experiments [e.g. 6-8]. In this study, we explore the effects of curvature of the asteroid strength (cohesion, crush resistance), porosity, and density of boulders on the cratering process in laboratory experiments at the EPIC facility in CAB, and we use these experiments to validate numerical codes.   

2. Method

Impact experiments are performed using the EPIC accelerator, which is a 20 mm calibre compressed N2 (300 bar) cannon that launches projectiles at velocities up to ≈420 m/s. The experiments, half- or quarter-space, can be recorded with two high-speed cameras. In the experiments presented here we have kept the impact velocity constant at approximately 380 m/s and used 20 mm spherical delrin projectiles.  We used motion tracking sensors with a 100 Hz sampling rate to measure the seismic signal that is generated by the impact. For the first experiment, approximating the surface curvature of an asteroid, we used a target of quartz sand piled in the shape of a half-frustum of ~17 cm height (quarter-space). For the target properties-studies we used two unconsolidated targets of projectile-sized boulders of different strength (cohesion, crush resistance) in quarter-space configuration. The weaker boulder material is made from a blend of sand and gypsum with a lower crush-strength than the porous ceramic boulders in the other experiment. The material of the sand/gypsum boulders was also tested in a homogeneous cylinder impact setup (half-space).

To simulate these experiments, we use the iSALE shock physics code [9-11], applying different strength models and the ε-α-porosity compaction model: For sand parameters, we refer to [6].

3. Results

The crater in the frustum target (Figure 1, top) reaches a transient stage after ~85 ms. The crater is larger than a crater in a flat sand target (Figure 1, bottom).  First models of the frustum experiment yield a similar crater diameter (+16%).

Figure 1: Frustum. Top: Images of the original target (left half) and the final crater (right half). Colours denote: Blue – 26.7 cm crater width at pre-impact surface level when the maximum depth (8.2 cm, red) is reached (49 ms). Green – 27.3 cm crater width of final crater. Bottom: Simulation of frustum (left) and flat sand target (right) after 70.5 ms. The dashed line indicates the shape of the frustum.

The crater in a cohesive homogeneous target of low strength (Figure 2) resembles typical craters in competent rock with fractures and spall pieces.

Figure 2: Homogeneous weak (but cohesive) material buried in sand. The cylindric material is 20 cm in diameter and 20 cm in depth. The crater is ~7.1 ± 0.5 cm in diameter.

The experiments with targets made of boulders of two different strengths produce craters of different characteristics (Figure 3). While the diameters are similar, the depth of the crater with weaker boulders is much deeper than in the case of stronger boulders, yielding a depth-diameter ratio of ~0.56 instead of ~0.19. In both cases, impact generated dust is injected into the pore space in the target. In the weak target, more boulders are crushed, and fragments are ejected. In the stronger case, ejection is mostly suppressed.

Figure 3: Craters in unconsolidated targets of strong (top) and weak (bottom) cohesive, projectile-size boulders. The diameters (green) are 20 cm & 17.7 cm, respectively, and the depth (orange) is 3.8 cm & 9.9 cm.

4. Conclusion and Discussion

Our experiments serve as validation scenarios for shock physics codes. First simulations can reproduce the resulting crater shape and diameter (Figure 1), but further modelling is ongoing.

The experiments show that surface curvature plays an important role for the cratering event if a large fraction of the crater growths beyond the flat surface. In our case, the flat top diameter of the frustum is about the size of the crater radius in a flat target, i.e. about half of the cratering is ongoing beyond the flat surface. In contrast to a flat target, the release of the pressure and the resulting pressure gradients are not vertical (i.e. there is a horizontal component influencing material movement), and a larger lateral growth of the crater occurs.

The experiments with targets of homogeneous projectile-sized particles (“boulders”) show some  transition of the cratering regime from typical crater shapes to deep craters and funnels [cf. 12, 13]. Besides impact velocity or the projectile-target density ratio, the transition seems to depend on the strength and crushing behaviour of the target.

Acknowledgements

We gratefully acknowledge the developers of iSALE (www.isale-code.de). The work by J.O. and S.D.R. was supported by grant PID2021-125883NB-C22 by the Spanish Ministry of Science and Innovation/State Agency of Research MCIN/AEI/ 10.13039/501100011033 and by “ERDF A way of making Europe”. R.L., J.O., I.H. S.D.R., M.J., and K.W. were supported by the Spanish Research Council (CSIC) support for international cooperation: I-LINK project ILINK22061.

References

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[4] Raducan S. D. et al., Nature Astronomy,8,445-455,2024.

[5] DeCoster M. E. et al., PlanSciJour,5,21,2024.

[6] Ormö J. et al., M&PS, 50(12),2067-2086,2015.

[7] Luther R. et al., PlanSciJour,3,227,2022.

[8] Ormö J. et al., EPSL,594,117713,2022.

[9] Amsden A. et al., LANL,LA-8095,101,1980.

[10] Wünnemann, K. et al.,Icarus,180,514-527, 2006.

[11] Collins, G. S. et al., Int.Journ.Imp.Eng.38,434-439,2011.

[12] Kadono T., PSS,47,305-318,1999.

[13] Luther R. et al., M&PS,58(12),1832-1847, 2023.

How to cite: Luther, R., Ormö, J., Herreros, I., Benavidez, P., Raducan, S. D., Jutzi, M., Baldauf, S., and Wünnemann, K.: Impact Experiments and Model Validation in the frame of the Hera mission , Europlanet Science Congress 2024, Berlin, Germany, 8–13 Sep 2024, EPSC2024-194, https://doi.org/10.5194/epsc2024-194, 2024.