Gabriel A Pinto
email: gabriel.pinto@univ-lorraine.fr

Yves Marrocchi

Alessandro Morbidelli

Sébastien Charnoz

Maria Eugenia Varela

Kevin Soto

Rodrigo Martínez

Felipe Olivares

Constraints on Planetesimal Accretion Inferred from Particle-size Distribution in CO Chondrites

Keywords: Accretion, Asteroids, Carbonaceous chondrites, Protoplanetary disk, Planetesimals

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Introduction

Planetesimals are solid objects larger than 1 km in diameter that formed by the accumulation of orbiting bodies in the protoplanetary disk and whose internal strengths are dominated by self-gravity; they represent the main building blocks of the planets orbiting the Sun today. Chondrites are fragments of asteroids that were never sufficiently heated to melt their constituent silicates and thus preserve primitive grains of the materials from which they agglomerated, including refractory inclusions and chondrules, cemented together by a complex fine-grained matrix. Refractory inclusions are millimeter-to centimeter-sized particles that represent the oldest dated objects in the solar system (Connelly et al. 2016).

Chondrules are (sub-)millimeter-sized igneous spherules that formed by a still elusive hightemperature mechanism linked to either nebular or planetary processes (e.g., Johnson et al. 2015;Marrocchi et al. 2018Marrocchi et al. , 2019)). Although chondrules are the main constituents of chondrites and their accretion thus represents a key step in the formation of planetesimals, the mechanism by which they assembled into planetesimals remains poorly constrained.

Recent theoretical advances suggest that planetesimals formed from clumps of small particles (mostly chondrules in the case of chondrites) whose common gravitational attraction outweighed the dispersive action of turbulent diffusion (Klahr & Schreiber 2020, 2021). This process requires that clumps of particles be sufficiently dense and massive and would produce planetesimals of typically ~10 -100 km in diameter depending on the remaining gas mass in the solar nebula at the time of planetesimal formation. However, the formation of such particle clumps remains debated.

For instance, it was proposed that particles were concentrated into large vortices (Barge & Sommeria 1995) or in regions between small vortices that developed in the disk at the dissipation scale (Cuzzi et al. 2001(Cuzzi et al. , 2008)). More recently, it was proposed that particle clumps formed due to streaming instabilities, hydrodynamic instabilities due to the differential velocities of particles relative to the surrounding gas (Youdin & Goodman 2005;Johansen et al. 2009;Simon & Armitage 2014;[START_REF] Wahlberg Jansson | Proceeding Asteroids, Comets[END_REF], 2017;Johansen et al. 2015;Li et al. 2018Li et al. , 2019)).

For all these scenarios, the effectiveness of streaming instabilities in promoting clump formation depends on particle size, or rather the Stokes number, which here is the ratio between a particle's stopping time due to friction with the gas and the orbital period. For chondrule-sized particles, triggering the gravitational collapse of a pebble cloud in streaming instability (Gerbig et al., 2020) requires particles to be radially concentrated in an annulus (Drążkowska et al. 2016) to achieve a sufficiently large solid/gas ratio (Carrera et al. 2015;Yang et al. 2017). One of the strengths of the gravitational collapse scenarios is that it predicts the formation of binary planetesimals, which are observed in large numbers in the relatively pristine trans-Neptunian belt, and reproduces the observed statistics of their orbital orientations (Nesvorný et al. 2019).

Once planetesimals have formed, they can continue growing by accreting individual particles (if they exceed a critical size of about 1000 km diameter, Ormel & Klahr 2010) as they drift through the disk. This process is known as pebble accretion (Lambrechts & Johansen 2012;Johansen et al. 2015). The initial gravitational contraction of a clump of particles and later pebble accretion should produce characteristic variations of particle size with depth inside the resulting planetesimal. Here we report estimated particle-size distributions within different CO carbonaceous chondrites, chosen because they experienced limited alteration processes after their agglomeration, which could have affected the sizes of their constituent particles. We use our data to model the conditions of planetesimal accretion within the protoplanetary disk and the possible layered structure of the CO parent body(ies).

2-Material and method

We surveyed all particles in sections of three CO3 carbonaceous chondrites recovered in the Atacama Desert and provided by the Museo del Meteorito (San Pedro de Atacama, Chile): El Médano 216 (EM 216), El Médano 463 (EM 463) and Los Vientos 123 (LoV 123).

We also determined the particle-size distribution in Isna (thick section 3239 from the Muséum national d'Histoire naturelle, Paris, France). Backscattered electron (BSE) mosaics and X-ray compositional maps (Fe, Ni, Al, Mg, Ca, Si, S, Cr) were acquired using (i) a JEOL JSM-6510 scanning electron microscope (SEM) equipped with a Genesis EDX detector and operating with a 3 nA electron beam accelerated at 20 kV (CRPG, Nancy, France) and (ii) a JEOL 6400 SEM operating with a 1 nA electron beam accelerated at 15 kV (Naturhistorisches Museum, Vienna, Austria). The chemical compositions of ferroan olivine grains in FeO-rich chondrules of EM 216 and EM 463 were quantified using a Cameca SX100 electron microprobe at the Service Commun de Microscopies Electroniques et de Microanalyses X (SCMEM, Université de Lorraine, Nancy, France) using a 12 nA focused beam accelerated at 15 kV. LoV 123 was analyzed by wavelength dispersive spectroscopy with an ARL-SEMQ electron microprobe at ICATE (San Juan, Argentina) operating with a 15 nA electron beam accelerated at 15 kV. Natural and synthetic standards were used for both instrument calibrations.

Mosaics of all samples were prepared using the GNU image manipulation program.

Particle-size measurements were performed using the Fiji distribution of the ImageJ opensource image processing software (Schindelin et al. 2012). We analyzed all nebular components, including chondrules (types I and II being FeO-poor and -rich, respectively), calcium aluminum-rich inclusions (CAIs, types A and B), amoeboid olivine aggregates (AOAs), and isolated olivine grains (IOGs) over total surface areas of 96 mm 2 , 27 mm 2 , 32.4 mm 2 , and 30.6 mm 2 for EM 463, EM 216, LoV 123, and Isna, respectively. The sharpness and continuity of the borders in the X-ray compositional maps was improved in Fiji by first applying mean filter at 1 pixel and then enhancing the image contrast at 0.5%. Each particle was recorded in a mask layer by free-hand tracing (Fig. 1). We did not distinguish between AOAs or type-A or -B CAIs for refractory components in EM 216 and LoV 123. Each particle's diameter (d) was calculated assuming that its total area was circular in cross section (i.e., as d =

).

3-Results

The Cr 2 O 3 contents of subhedral FeO-rich olivine crystals were determined for 43, 49, and 62 type-II porphyritic chondrules in EM 216, EM 463, and LoV 123, respectively.

Chondrule olivine grains appear heterogeneous in both texture (Fig. 1) and composition , mean Fa 50.2 ± 10.7 ; Table S1). Ferroan olivines in type-II chondrules of EM 463 contain 0.04-0.57 wt.% Cr 2 O 3 (average 0.09 ± 0.10 wt.%, 1, Table S1); those of EM 216 and LoV 123 contain 0.09 ± 0.10 and 0.30 ± 0.13 wt.% Cr 2 O 3 , respectively (Table S1).

EM 216, EM 463, and LoV 123 show high modal abundances of FeO-poor type-I chondrules (41.59, 33.15, and 46.74%, respectively, Table 1) surrounded by fine-grained Ferich matrix (Fig. 1). Isna is dominated by type-II chondrules (33.26%) with occasional type-I chondrules (1.85%, Table 1). The modal abundances of refractory inclusion (CAIs + AOAs) range from 1.59% in Isna to 8.60% in LoV 123 (Table 1). We also observe a large variation in the modal abundances of type-II chondrules, from 5.22% in EM 216 to 33.26% in Isna (Table 1).

Our results show a significant difference between the mean spherical diameters of type-I and type-II chondrules: 92.78 and 162.52 μm, respectively (

Discussion

Correlation between chondrule-size distribution and degree of metamorphism

The degree of thermal alteration (i.e., petrographic grades) experienced by the carbonaceous chondrites can be estimated based on the mean and standard deviation of Cr 2 O 3 content of FeO-rich olivines in type-II chondrules (Grossman and Brearley 2005, Table S1).

We determined a petrographic grade of LoV 123, EM 216 and EM 462 to be 3.05, 3.2 and 3.2, respectively. The petrographic grade of Isna has been determined to be 3.75 ± 0.05 based on the Cr 2 O 3 content of Fe-rich olivines (Rubin & Li 2019) and detailled petrographic and mineralogical studies of AOAs (Chizmadia et al. 2002). Incorporating recent literature data for CO chondrites (Ebel et al. 2016), average chondrule diameters define a positive relationship with the petrographic grade of their parent chondrite (Fig. 2). Such a correlation was first noticed by Rubin (1989), although they reported a dissimilar relationship (Fig. 2A), likely due to their different estimation method and limited number of analyzed particles (< 1000). Because we used a similar method as Ebel et al. (2016), we here compare our results

to their dataset.

The deformation of chondrules (Fig. 1) could have either occurred during the evolution of the protoplanetary disk or within chondritic parent bodies. As lobate chondrules are commonly observed in carbonaceous chondrites that experienced minimal secondary deformation (Jacquet 2021) and CO carbonaceous chondrites show low impact-generated metamorphism transformation (Scott et al., 1992), we thus favor a nebular origin for the nonspherical chondrules. The observed correlation between chondrule diameter and petrographic grade in CO chondrites could result from (i) the conditions of planetesimal accretion (Scott & Jones 1990) and/or (ii) post-accretion thermal metamorphism processes resulting from impacts and/or 26 Al decay (Doyle et al. 2015;Vacher et al. 2018;Amsellem et al. 2020;Turner et al. 2021).

Although thermal metamorphism could result in mineral coarsening (Huss et al. 2006), this would mainly affect Fe-Ni metal beads and sulfides and would require temperatures >800 °C, significantly hotter than those estimated for CO chondrites (i.e., 300-600 °C; Jones & Rubie 1991;Keil 2000;Bonal et al. 2007). This is consistent with the fact that CV chondrules were only affected by Fe-Mg diffusion without any significant size increase, despite having undergone thermal metamorphism at temperatures higher than in CO chondrites (i.e., ≥600 °C; Ganino & Libourel 2017). Furthermore, type-I chondrule boundaries are well defined in CO chondrites, even in the most metamorphosed sample investigated here (i.e., CO3.8 Isna).

Taken together, these lines of evidence indicate that (i) the size characteristics of chondrules result from their formation processes during the evolution of the disk and (ii) aerodynamic sorting played a key role during the accretion of the CO parent body(ies). Rubin (1989) suggested that the relationship between chondrule size and the degree of metamorphism is related to monotonic changes in the aggregation of materials in the nebular CO reservoir. In this framework, larger chondrules would have been more centrally concentrated in the CO parent body(ies) than smaller chondrules, which could be explained by either simultaneous or sequential accretion of the two chondrule populations (Scott & Jones 1990). Larger chondrules located closer to the center of the CO parent body(ies) would have experienced more protracted thermal metamorphism than smaller chondrules closer to the surface, where heat generated by 26 Al decay was more readily evacuated. Aerodynamic sorting during accretion could thus have produce the co-variation of mean chondrule size with both depth and subsequently with metamorphic grade, as observed for CO chondrites.

As CO refractory inclusions are smaller than chondrules (Table 2), this process would also have generated positive relationships between CO metamorphic grade and their (i) bulk oxygen isotopic compositions and (ii) refractory inclusion/chondrule (RI/C) ratios.

Interestingly, the former has been reported in previous studies (Clayton & Mayeda 1999;Greenwood & Franchi 2004), and our data confirm that COs with lower metamorphic grades show higher RI/C ratios than more metamorphosed ones (Table 2). These results confirm that aerodynamic sorting during accretion concentrated larger chondrules toward the center of the CO parent body(ies), but smaller chondrules and CAIs were less concentrated at the core.

Based on this conclusion, in the following section we model and discuss which accretion process best matches these peculiar features of CO chondrites.

Size-sorting during planetesimal formation

In this section we attempt to provide a qualitative explanation of the interpretation discussed above, in which larger chondrules are more abundant than smaller chondrules at greater depths within the parent planetesimal, and vice versa. We first consider the case of planetesimal formation due to the self-gravitational contraction of a clump of particles.

Before they collapse into each other, the radial density distribution of particles in the group, ρ(r), is dictated by the equilibrium between gravity and turbulent diffusion of the gas within the group, resulting in (Klahr & Schreiber 2020a, 2020b):

ρ(r) = ρ(0) exp[-r 2 /(2l c 2 )], (1)
where

l c = 1/3 (δ/St) 1/2 H, ( 2
)
where St is the particle's Stokes number, r is the distance to the center, H is the pressurescaled height of the gas in the disk, and δ is the non-dimensional coefficient relating the turbulent diffusion coefficient D to the disk's scale height and orbital frequency Ω:

D = δH 2 Ω. ( 3
)
For particles smaller than the mean free path of gas molecules, St is proportional to particle size (i.e., the Epstein regime). In the Stokes regime, St is proportional to the square of particle size, although this case is rarely considered (nor would it change the considerations below). If there are particles of multiple sizes in the clump, even if the gravitational potential is set by one dominant particle size, each will follow a radial distribution like (1), with its own value of l c . Thus, combining Equations ( 1) and ( 2), smaller particles have a more distended radial distribution in the clump, whereas larger particles are more concentrated toward the center. This is because, for larger and smaller particle species 1 and 2, respectively, if St 1 > St 2 , then

l c1 < l c2 .
We now consider the settling of particles towards the center of the clump, forming a solid planetesimal. Particles are accelerated towards the center, but undergo more and more friction as their sedimentation rate increases; the two forces cancel when a particle attains its terminal velocity (v t ). Particles accelerated in a gas medium rapidly attain the so-called terminal velocity (v t ). The value of the terminal velocity increases approaching the center of gravitational attraction, but the slowest velocity defines the time that the pebble needs to reach the center, and that velocity is set by initial terminal velocity v_t. Thus, in the following we assume, without introducing much error and to simplify the final formulae, that v t is constant (but is particle-size dependent) during the whole contraction of the clump.

The planetesimal grows over time as more and more particles reach its surface during the sedimentation process. Thus, we can use the accretion timescale T as a proxy for planetesimal radius (not necessarily a linear relationship). In a time interval dT at time T, the planetesimal receives the particles that were originally in the clump at a distance between r and r + dr where r = v t T and dr = v t dT (here is where the assumption of constant v t is handy). At time T all particles from specie i (i=1 or 2) come from a spherical shell with radius r i = V ti T and thickness dr = v t1 dT, so the total mass accumulated in planetesimals between time T and T + dT

is dM i (T) = 4r i 2  i (r i )dr i = 4v ti 3 T 2  i (r i )dT.
So, the mass ratio of the particles with index 1 and 2 accumulated in the time interval T to T + dT is:

ρ 1 /ρ 2 (T) = (v t1 /v t2 ) 3 • [ρ 1 (v t1 T)/ρ 2 (v t2 T)], ( 4
)
where the first term (v t1 /v t2 ) 3 comes from the ratio of the volumes occupied at time 0 (i.e., at the beginning of the contraction of the clump) by the particles that sediment onto the planetesimal surface between T and T + dT. Let's now define x = v t1 T/(2 l c1 2 ). Recall that

v t1 /v t2 = St 1 /St 2 and l c1 /l c2 = (St 2 /St 1 ) 1/2 . Thus, (v t2 T) 2 /(2l c2 2 ) = (St 2 /St 1 ) 3 x 2 .
By applying the definition (1) into (4) and substituting for x, the mass ratio of the two population of particles landing on the planetesimal at time T = x(2l c1 2 ) 1/2 /v t1 is:

ρ 1 /ρ 2 (x) = ρ 1 (0)/ρ 2 (0) • (St 1 /St 2 ) 3 • exp[-x 2 ]/exp[-(St 2 /St 1 ) 3 x 2 ].
(5)

As an example, Figure 3 shows this function for St 2 = 0.6St 1 , which is appropriate for comparing chondrules with diameters of ~80-130 µm, i.e., in the Epstein regime. 5) normalized to the ratio ρ 1 (0)/ρ 2 (0) at the center of the particle clump from the beginning of a gravitational contraction leading to the formation of a planetesimal.

Our model shows that the ρ 1 /ρ 2 decreases monotonically as time progresses (i.e., when the timescale increases). This is in qualitative agreement with our interpretation that average chondrule size increases with depth in the parent body. Of note, chondrules are small compared to pebbles (i.e., 80-130 µm; Table 2) and would have Stokes number of 1 × 10 -4 at 2 AU for 100 µm particle-size (considering a minimum mass solar nebula). This is smaller than what is typically considered in gravitational collapse models of planetesimal formation (e.g. Yang et al., 2017). To have a more St =1e-3, as more traditionally considered, the gas should have been depleted by a factor of 10. This may well be possible, because chondrites formed at a late time (2-3 My) in the disk's chronology, well after the formation of Jupiter (Kruijer et al., 2017); cavity opening by Jupiter and photo-evaporation may well have reduced the density of gas by one order of magnitude.

Above, we considered the case of a gravitationally contracting planetesimal, but it is well accepted that after their contraction, planetesimals may have grown through pebble accretion, as discussed by Johansen et al. (2015). If a planetesimal accreted particles in the Bondi regime, the Stokes number of the preferentially accreted pebbles increases with the planetesimal's Bondi time, which is proportional to the planetesimal's mass (Lambrechts & Johansen 2012). This predicts that larger chondrules should have been predominant at shallower depths in their parent planetesimal, opposite to our observations. In the case that a larger protoplanet scattered a planetesimal into an orbit of greater inclination, the size of particles accreted by the planetesimal would have suddenly decreased because only small particles were available away from the disk's midplane. However, the mass accreted should have also declined, such that small particles would dominate only in a very thin surface layer (Johansen et al. 2015). We do not observe such a drastic change in chondrule size among the studied meteorites of different petrologic types, but rather a gradual trend. Thus, we conclude that our observations are consistent with the formation of the CO parent body(ies) via the contraction of a self-gravitating clump of chondrules of various sizes, and that the subsequent growth of the parent body(ies) due to pebble accretion was insignificant.

Conclusions

Our particle-size analysis of CO carbonaceous chondrites revealed that the mean spherical diameters of chondrules increase with increasing metamorphic degree. Combining our results with literature data, we show that this relationship was not established during postaccretion thermal metamorphism, but instead records aerodynamic size-sorting of particles during the accretion of the CO parent body(ies). By modeling the self-gravitating contraction of clumps of chondrules, we show that the accretion processes generated a gradual change in chondrule size, with larger chondrules being more centrally concentrated in the parent body(ies) than smaller ones. Our results also suggest that any subsequent planetesimal growth by pebble accretion should have been insignificant. We thus conclude that our observations are consistent with the formation of the CO parent body(ies) via the contraction of a selfgravitating clump of chondrules, supporting the theory that the formation of planetesimals occurred via gravitational collapse.

Appendices

Figure 1 .

 1 Figure 1. Representative examples of particle selection. (A) BSE image of LoV 123. (B) Xray compositional map of EM 463, with Mg, Al, Ca, and Fe shown in red, blue, green, and white. The colors of particle outlines indicate component type. In (A) type-I chondrules are outlined in red, type-II chondrules in blue, refractory components in green, IOGs in orange, and fine-grained rims in yellow. In (B), outline colors are as in (A), except that refractory components are distinguished between AOAs in green and CAIs in purple.

  Figure 2. (A) Mean chondrule diameter vs. petrologic type of CO chondrites (data from Rubin 1989, Ebel et al. 2016, and this study). (B) Probability density function of chondrule diameters in CO chondrites. Solid lines, this study (N, number of analyzed chondrules); dashed lines, Ebel et al. (2016).

Figure 3 .

 3 Figure 3. Equation (5) normalized to the ratio ρ 1 (0)/ρ 2 (0) at the center of the particle clump from the beginning of a gravitational contraction leading to the formation of a planetesimal.Our model shows that the ρ 1 /ρ 2 decreases monotonically as time progresses (i.e., when

Table 2

 2 , and 170.81 µm (54.90), respectively (Fig.2A, Table2). Similarly, type-II chondrules and refractory inclusions (CAIs + AOAs) show variable average sizes among the different COs (Table2): in LoV 123, EM 216, EM 463, and Isna, type-II chondrules have mean diameters of141.64 (112.92), 103.84 (95.26), 164.10 (100.38), and 184.85 µm (80.14), respectively, and refractory inclusions have mean diameters of62.63 (39.06), 75.32 (52.80), 109.76 (71.24), and 67.11 µm (47.81), respectively. The circularity and the mean aspect ratio of chondrules in our studied CO chondrites is 0.67 ± 0.15 and 1.60 ± 0.45, respectively.

	). The mean

Table 1 .

 1 Modal abundances (the ratio of component pixel area relative to the total pixel area of the chondrite, in %) of refractory components and chondrules in the analyzed CO3 chondrites.

	Component	LoV 123	EM 216	EM 463	Isna
	Type I chondrule	46.74	41.59	33.15	1.85
	Type II chondrule	8.64	5.22	7.75	33.26
	All chondrules	55.37	46.81	40.90	35.11
	CAI A	-	-	0.58	0.66
	CAI B	-	-	0.55	0.93
	AOA	-	-	2.37	-
	All refractory inclusions (RI)	8.60	6.82	3.49	1.59
	RI/chondrules ratio	0.16	0.15	0.09	0.05

Table 2 .

 2 The number as well as mean and median diameters (µm) of chondrules and refractory inclusions in each surveyed section.

	Component	LoV 123	EM 216	EM 463	Isna	All CO
	Type I chondrule
	n	2210	1140	2539	25	5914
	Mean diameter	71.60	90.91	111.37	170.81	92.78
	1σ	56.51	65.61	89.44	54.90	74.92
	Median diameter	55.68	71.73	88.56	158.10	72.77
	Type II chondrule
	n	109	91	253	318	771
	Mean diameter	141.64	103.84	164.10	184.85	162.52
	1σ	112.92	95.26	100.38	80.14	97.33
	Median diameter	108.01	73.56	138.00	164.99	138.34
	All chondrules
	Mean diameter	74.89	91.86	116.15	183.82	100.47
	1σ	62.11	68.28	89.48	71.62	80.27
	Median diameter	57.16	71.80	92.60	164.03	79.00
	CAIs+AOAs
	n	221	278	252	92	843
	Mean diameter	62.63	75.32	109.76	67.11	81.35
	1σ	39.06	52.80	71.24	47.81	43.08
	Median diameter	53.11	61.47	88.63	49.03	64.54

Table S1 .

 S1 Chemical compositions of ferroan olivine grains in type-II chondrules of three CO carbonaceous chondrites(EM 463, EM 216, LoV 123).

			37.12 35.51 37.27	34.90 41.70 27.96	0.04 0 -	0.16 0.15 -	0.45 0.31 0.39	29.55 24.15 34.91	0.08 0.07 0.21	0.01 0 0.07	102.3 101.9 100.8
			38.95 36.93 36.62	23.89 36.28 28.41	0.01 0.01 -	0.03 0.13 -	0.16 0.27 0.44	38.35 28.88 34.30	0.07 0 0.22	0 0.02 0	101.4 102.5 100.0
			36.14 36.77 37.66	37.74 35.47 23.47	0.04 0.06 -	0.24 0.08 -	0.43 0.29 0.37	26.92 29.26 38.07	0.08 0.07 0.38	0 0 0.05	101.6 102.0 100.0
			38.19 36.56 35.49	27.40 35.44 36.77	0.03 0.06 -	0.14 0.24 -	0.32 0.14 0.32	35.56 28.97 27.34	0.15 0.06 0.21	0 0 0.04	101.8 101.5 100.2
	CO3	#	SiO2 36.79 34.61 35.86	FeO 35.39 41.14 33.14	Al2O3 0.01 0.21 -	CaO 0.22 0.25 -	MnO 0.39 0.36 0.25	MgO 28.60 23.82 30.56	Cr2O3 0.08 0.40 0.51	TiO2 0 0 0	Total 101.5 100.8 100.3
	EM 463	49	38.34 38.26 37.98 36.80	27.85 28.99 29.72 25.71	0 0.01 0 -	0.09 0.06 0.10 -	0.31 0.42 0.22 0.25	34.58 34.50 33.24 36.44	0 0.06 0 0.55	0 0 0 0.01	101.2 102.3 101.3 99.7
			36.05 36.28 37.46 38.40	39.18 35.30 31.18 17.55	0.03 0.02 0.06 -	0.11 0.24 0.10 -	0.28 0.36 0.28 0.18	24.89 28.41 33.48 43.11	0.15 0.06 0.18 0.23	0 0.01 0.01 0.04	100.7 100.7 102.7 99.5
		Mean	37.41 37.04 36.98 36.00	30.69 29.66 34.42 30.58	0.05 0.03 0.04 -	0.09 0.18 0.12 -	0.25 0.35 0.33 0.36	31.50 33.41 29.80 32.64	0.05 0.06 0.09 0.22	0 0.02 0.01 0	100.0 100.7 99.8
		Mean STD	36.48 37.71 1.09 37.61	33.55 29.33 5.44 24.04	0.16 0.05 0.05 -	0.08 0.12 0.08 -	0.28 0.31 0.10 0.22	30.73 33.56 4.51 37.69	0.09 0.09 0.10 0.45	0 0.01 0.01 0	101.4 100.0
	LoV 123	STD 62	38.51 1.13 35.67 37.93	23.92 5.57 31.03 24.38	0.03 0.05 --	0.01 0.08 --	0.31 0.10 0.31 0.30	38.32 4.57 32.75 36.94	0.01 0.10 0.24 0.34	0 0.01 0 0	101.1 100.0 99.9
	EM 216	43	38.27 37.04 35.54 36.08	26.49 35.98 33.42 31.88	0.10 0.04 --	0.02 0.08 --	0.30 0.41 0.31 0.34	36.47 28.47 30.54 31.55	0.01 0.02 0.23 0.16	0.01 0 0 0.03	101.6 102.0 100.0 100.0
			37.83 37.47 35.62 33.24	28.32 33.21 32.24 49.31	0.10 0.03 --	0.14 0.14 --	0.34 0.42 0.31 0.53	35.69 30.91 32.20 16.72	0.14 0.11 0.27 0.18	0.01 0.04 0.04 0.02	102.6 102.3 100.7 100.0
			38.53 36.13 35.69 37.14	24.64 38.93 33.58 24.45	0.04 0.09 --	0.11 0.28 --	0.35 0.39 0.20 0.26	37.80 26.55 31.47 37.49	0 0.10 0.27 0.23	0 0.01 0 0.01	101.5 102.5 101.2 99.6
			37.00 37.49 33.56 40.74	31.29 28.17 42.69 6.99	0.05 0.10 --	0.14 0.22 --	0.40 0.33 0.43 0.11	31.79 35.20 23.09 52.25	0.06 0.08 0.27 0.25	0.01 0 0 0.02	100.7 101.6 100.0 100.3
			39.12 38.52 37.40 36.96	23.16 26.03 21.97 28.77	0.04 0.01 --	0.13 0.06 --	0.26 0.30 0.25 0.37	39.32 36.89 40.04 33.99	0.06 0.11 0.31 0.19	0 0 0 0.01	102.1 101.9 100.0 100.3
			35.92 37.41 33.61 37.19	35.52 31.21 41.06 25.44	0 0.03 --	0.11 0.01 --	0.36 0.39 0.33 0.29	28.72 31.99 25.03 37.01	0.03 0.08 0.32 0.29	0 0 0.03 0	100.7 101.1 100.4 100.2
			37.92 35.96 34.18 35.33	26.29 39.55 36.98 38.88	0.01 0.04 --	0.15 0.25 --	0.38 0.45 0.26 0.30	36.24 24.97 28.02 25.30	0.08 0 0.40 0.26	0 0 0 0.01	101.1 101.2 99.8 100.1
			37.91 38.79 36.65 38.64	30.51 24.12 27.88 15.59	0.06 0.03 --	0.09 0.02 --	0.17 0.27 0.22 0.35	32.96 37.98 34.58 45.25	0.08 0.41 0.13 0.66	0.01 0 0.04 0.03	101.8 101.6 99.5 100.5
			36.94 37.53 35.73 38.82	34.02 33.13 27.90 17.52	0.09 0 --	0.02 0.15 --	0.38 0.27 0.15 0.15	29.42 32.12 36.32 43.18	0.18 0.09 0.27 0.19	0 0 0.01 0	101.0 103.3 100.4 99.9
			39.29 37.46 35.92 39.06	23.36 30.82 31.78 15.93	0 0.27 --	0.18 0.05 --	0.25 0.24 0.32 0.26	38.18 31.97 33.20 44.54	0.19 0.40 0.22 0.39	0.04 0 0 0	101.5 101.2 101.4 100.2
			40.21 36.11 33.28 38.61	17.47 39.77 38.93 20.66	0.10 0 --	0.10 0.14 --	0.22 0.22 0.59 0.15	42.78 25.86 25.75 40.72	0.11 0.03 0.13 0.23	0 0 0.01 0	101.0 102.1 98.7 100.4
			37.97 38.41 33.17 35.17	24.12 28.06 42.60 35.55	0.23 0 --	0.30 0.08 --	0.26 0.29 0.34 0.32	36.13 34.43 23.95 28.56	0.57 0.13 0.15 0.18	0.01 0 0 0	99.6 101.4 100.2 99.8
			38.61 37.11 33.48 34.18	26.09 34.77 42.07 42.09	0.02 0 --	0.12 0.08 --	0.37 0.47 0.52 0.36	35.62 29.28 24.00 22.67	0.08 0.05 0.13 0.32	0.01 0 0.03 0.01	100.9 101.8 100.2 99.6
			38.40 38.43 33.39 37.36	24.24 25.50 41.49 21.05	0.09 0.03 --	0.09 0.10 --	0.16 0.25 0.45 0.19	36.61 36.57 24.46 40.44	0.08 0 0.36 0.49	0 0.02 0.01 0.04	99.7 100.9 100.2 99.6
			38.40 35.91 32.97 36.10	26.45 38.17 42.15 29.09	0.03 0.08 --	0 0.19 --	0.22 0.48 0.46 0.23	35.04 25.92 24.82 34.65	0.05 0.19 0.25 0.26	0 0 0 0.05	100.2 100.9 100.6 100.4
		Mean	37.89 37.73 37.30 35.88	30.34 32.52 26.38 31.84	0.04 0.02 -	0.09 0.02 -	0.38 0.22 0.30 0.34	33.23 31.88 35.92 31.62	0.09 0.02 0.22 0.30	0 0 0.05 0.02	102.1 102.4 100.2
		STD	38.34 35.51 40.68 1.92	25.83 41.54 20.79 9.27	0 0.06 -	0.04 0.34 -	0.20 0.42 0.23 0.12	35.77 23.33 31.63 7.50	0 0.12 0.33 0.13	0.02 0.03 0.12 0.02	100.2 101.3 93.8
			39.23 36.01 34.73	21.67 37.16 39.52	0.01 0.03 -	0.03 0.15 -	0.27 0.46 0.43	39.77 27.34 24.90	0.12 0.06 0.27	0 0.02 0.02	101.1 101.2 99.8
			38.30 38.80 34.63	28.82 25.71 39.66	0.08 0.02 -	0.13 0.11 -	0.43 0.47 0.44	33.43 37.31 24.74	0.16 0.04 0.32	0 0 0	101.4 102.4 99.8
			37.41 36.89 34.68	31.61 37.57 39.59	0 0.02 -	0.18 0.06 -	0.31 0.37 0.43	31.75 26.87 24.82	0.04 0.11 0.29	0.03 0 0.01	101.3 101.9 99.8
			38.65 38.56 33.82	23.64 26.92 43.32	0.03 0 -	0.02 0.00 -	0.12 0.13 0.49	38.09 35.91 22.21	0.12 0.07 0.50	0 0 0.02	100.6 101.6 100.3
			38.63 36.18 35.19	24.94 38.87 38.42	0.10 0 -	0.06 0.15 -	0.16 0.35 0.46	37.24 26.31 25.68	0.19 0.04 0.30	0.01 0 0	101.3 101.9 100.1
			38.24 36.91 35.78	27.19 31.87 29.81	0.08 0.08 -	0 0.17 -	0.30 0.24 0.26	35.15 31.28 33.91	0.38 0.16 0.32	0 0 0.03	101.3 100.7 100.1
			38.28 35.93 35.00	27.97 40.76 34.25	0.10 0.02 -	0.01 0.22 -	0.26 0.55 0.32	33.84 24.52 29.95	0.05 0.05 0.21	0 0.04 0.01	100.5 102.1 99.7
			40.41 37.44 38.25	17.10 32.88 18.43	0.03 0.06 -	0.07 0.05 -	0.22 0.47 0.15	44.05 31.49 42.55	0.03 0.02 0.75	0.03 0.00 0	101.9 102.4 100.1
			37.81 35.80 34.12	29.89 40.93 38.47	0.06 0.01 -	0.17 0.28 -	0.28 0.38 0.52	32.57 24.63 26.47	0.10 0.08 0.36	0 0 0.01	100.9 102.1 100.0
			36.27 36.68 34.13	37.11 37.19 37.41	0.02 0.02 -	0.10 0.11 -	0.48 0.46 0.41	27.79 26.80 27.88	0.07 0.09 0.32	0 0.03 0.01	101.8 101.3 100.1
			36.66 37.16 34.97	33.88 34.65 40.21	0.07 0.03 -	0.14 0.07 -	0.24 0.32 0.53	30.13 30.04 24.48	0.06 0.03 0.27	0 0 0.04	101.2 102.3 100.5
			37.74 35.96 34.17	21.68 39.53 42.01	0.10 0 -	0.06 0.07 -	0.20 0.19 0.62	40.56 25.62 22.93	0.10 0.06 0.10	0 0 0	100.4 101.4 99.8
			37.55 36.82 32.64	30.09 36.62 48.74	0.01 0 -	0.11 0.12 -	0.22 0.32 0.53	33.38 28.42 18.20	0.04 0.06 0.32	0 0 0.02	101.4 102.4 100.4
			37.26 38.34 32.33	32.46 28.48 52.22	0.08 0.04 -	0.21 0.02 -	0.44 0.32 0.61	31.28 35.09 14.18	0 0.09 0.25	0.05 0.02 0.03	101.8 102.4 99.6
			36.82 36.15 35.88	34.74 37.60 28.68	0 0.03 -	0.08 0.07 -	0.53 0.32 0.25	29.35 27.12 35.17	0.05 0 0.28	0.04 0.02 0.03	101.6 101.3 100.3
			35.60 37.88 35.78	37.13 30.32 30.47	0.03 0.01 -	0.37 0.06 -	0.19 0.22 0.39	27.65 32.70 33.34	0.02 0.05 0.31	0 0.02 0	100.9 101.2 100.3
			35.32 38.28 36.19	40.69 27.78 30.39	0.02 0 -	0.27 0.05 -	0.46 0.28 0.35	23.44 35.41 32.38	0.07 0.17 0.59	0 0 0	100.3 102.0 99.9
			36.00 34.40 36.64	38.28 48.35 30.31	0.01 0.01 -	0.18 0.14 -	0.40 0.45 0.21	26.07 18.21 33.17	0.07 0.07 0.24	0 0 0.02	101.0 101.6 100.6
			37.63 37.04 37.65	31.64 34.55 22.20	0.05 0.04 -	0.14 0.09 -	0.43 0.32 0.25	31.66 30.23 39.53	0.13 0.01 0.28	0 0 0	101.7 102.3 99.9

Acknowledgments

We thank Johan Villeneuve, Nicolas Schnuriger, and Laurette Piani for helpful discussion. G.A.P and F.O.E. acknowledge support from FONDECYT project 1201223. G.A.P. was supported by an Eiffel excellence scholarship (grant 968045D). We thank the anonymous review for helpful comments and Frederic Rasio for editorial handling. This is CRPG contribution #2777.

Research data

Original data from this study are available on the Ordar database; http://doi.org/10.24396/ORDAR-68


Yves Marrocchi
email: yvesm@crpg.cnrs-nancy.fr

Marco Delbo

Matthieu Gounelle

The astrophysical context of collision processes in meteorites

Keywords: chondrites, clasts, brecciation, asteroids, impacts

Chondrites are leftover solids from the early evolution of the solar protoplanetary disk that never experienced melting since their formation. They comprise unequilibrated assemblages of low-and high-temperature components, including volatile-rich, fine-grained matrices, Fe-Ni metal, sulfides, refractory inclusions, and chondrules. Consequently, chondrites are commonly described as pristine, primitive, or primordial rocks of the solar system. However, impact-generated secondary features are abundant in chondrites, suggesting that collisions among early-formed planetesimals and their fragmentation and reassembly have been effective throughout the evolution of the solar system. In this report, we review evidence of the major role of impacts in generating the current mineralogical and petrographic characteristics of chondrites. We provide perspective to these meteoritic features by discussing recent analyses of large-scale structures of the main asteroid belt and remotesensing observations of asteroids. Observations at various spatial scales all attest that the 'primitive' materials formed during the evolution of the solar system have largely been reprocessed, confirming previous studies that primitivity is relative, not absolute. This implies that (i) chondrites (and some differentiated meteorites) should systematically be envisioned as reprocessed and heterogeneous materials and (ii) brecciated meteorites should be considered the norm and unbrecciated meteorites the exception.

1-Introduction

Meteorites are asteroidal (or rarely cometary; [START_REF] Gounelle | The orbit and atmospheric trajectory of the Orgueil meteorite from historical records[END_REF]) fragments recovered at the Earth's surface either promptly after their falls (e.g., [START_REF] Hiroi | The Tagish Lake Meteorite: A possible Sample from a D-Type Asteroid[END_REF][START_REF] Jenniskens | Radar-Enabled Recovery of the Sutter's Mill Meteorite, a Carbonaceous Chondrite Regolith Breccia[END_REF][START_REF] Marrocchi | The Piancaldoli meteorite: A forgotten primitive LL3.10 ordinary chondrite[END_REF] or during field trips to cold/hot deserts (e.g., Bischoff and Geiger 1995;[START_REF] Evatt | The spatial flux of Earth's meteorite falls found via Antarctic data[END_REF][START_REF] Gattacceca | The densest meteorite collection area in hot deserts: The San Juan meteorite field (Atacama Desert, Chile)[END_REF][START_REF] Harvey | The Origin and Significance of Antarctic Meteorites[END_REF][START_REF] Schlüter | The Dar al Gani meteorite field (Libyan Sahara): Geological setting, pairing of meteorites, and recovery density[END_REF]. Meteorites' parent bodies are leftovers from the planetary formation processes that operated during the evolution of the solar system. Differentiated parent bodies recorded processes such as metal-silicate differentiation, early volcanism, granite-like magmatism and magma ocean episodes (e.g., [START_REF] Bischoff | Trachyandesitic volcanism in the early Solar System[END_REF][START_REF] Terada | Asteroidal granite-like magmatism 4.53 Gyr ago[END_REF]. In contrast, chondritic parent bodies did not experience melting and comprise high-temperature components formed in the solar protoplanetary disk (SPD) before planetesimal accretion, namely refractory inclusions (Ca-Al-rich inclusions, CAIs, and amoeboid olivine aggregates, AOAs; e.g., Marrocchi et al. 2019a) and chondrules (e.g., [START_REF] Connolly | Chondrules: The canonical and noncanonical views[END_REF].

Chondrites thus represent invaluable witnesses of the SPD's evolution because they are made of components that formed at different times and locations in the SPD [START_REF] Bollard | Combined U-corrected Pb-Pb dating and 26Al-26Mg systematics of individual chondrules -Evidence for a reduced initial abundance of 26Al amongst inner Solar System chondrules[END_REF][START_REF] Connelly | The absolute chronology and thermal processing of solids in the solar protoplanetary disk[END_REF][START_REF] Kawasaki | Variations in initial 26Al/27Al ratios among fluffy Type A Ca-Al-rich inclusions from reduced CV chondrites[END_REF][START_REF] Macpherson | Wellresolved variations in the formation ages for Ca-Al-rich inclusions in the early Solar System[END_REF][START_REF] Pape | Time and duration of chondrule formation: Constraints from 26 Al-26 Mg ages of individual chondrules[END_REF][START_REF] Petitat | 53 Mn -53 Cr ages of Kaidun carbonates[END_REF][START_REF] Villeneuve | Homogeneous Distribution of 26 Al in the Solar System from the Mg Isotopic Composition of Chondrules[END_REF]).

However, most chondrites experienced subsequent thermal and hydrothermal alteration during asteroidal evolution, which blurred their initial textures, mineralogies, and chemical and isotopic compositions [START_REF] Amsellem | Timing of thermal metamorphism in CM chondrites: Implications for Ryugu and Bennu future sample return[END_REF][START_REF] Brearley | The action of water[END_REF][START_REF] Fujiya | Comprehensive study of carbon and oxygen isotopic compositions, trace element abundances, and cathodoluminescence intensities of calcite in the Murchison CM chondrite[END_REF][START_REF] Huss | Thermal Metamorphism in Chondrites[END_REF]Marrocchi et al. 2018a[START_REF] Marrocchi | The Tarda Meteorite: A Window into the Formation of D-type Asteroids[END_REF][START_REF] Turner | Carbonaceous chondrite meteorites experienced fluid flow within the past million years[END_REF][START_REF] Vacher | Petrographic and C & O isotopic characteristics of the earliest stages of aqueous alteration of CM chondrites[END_REF]Vacher et al. , 2019aVacher et al. , 2019b[START_REF] Vacher | Hydrogen in chondrites: Influence of parent body alteration and atmospheric contamination on primordial components[END_REF]. Based on these modifications, chondrites are assigned petrologic types on a scale from 1 to 6, reflecting the progressive roles of low-temperature aqueous alteration (from type 3 to 1) and thermal metamorphism (from type 3 to 6; van [START_REF] Van Schmus | A chemical-petrologic classification for the chondritic meteorites[END_REF].

Type-3 chondrites thus correspond to materials that underwent minimal modifications after their agglomeration. Despite having suffered hydrothermal alteration of varying intensity, high-temperature components in type-2 chondrites are thought to preserve evidence of the conditions of dust formation in the SPD (e.g., [START_REF] Chaumard | Oxygen isotope systematics of chondrules in the Murchison CM2 chondrite and implications for the CO-CM relationship[END_REF][START_REF] Kööp | A link between oxygen, calcium and titanium isotopes in 26Al-depleted hibonite-rich CAIs from Murchison and implications for the heterogeneity of dust reservoirs in the solar nebula[END_REF]Marrocchi et al. 2018bMarrocchi et al. , 2019b;;[START_REF] Piani | Origin of hydrogen isotopic variations in chondritic water and organics[END_REF]. Consequently, these peculiar chondrites are commonly described as pristine, primitive, or primordial rocks of the solar system [START_REF] Gounelle | The Asteroid-Comet Continuum: In Search of Lost Primitivity[END_REF][START_REF] Libourel | Search for primitive matter in the Solar System[END_REF]).

In addition to secondary alteration processes, hypervelocity impacts are of fundamental importance and played a key role in the early solar system and during its evolution [START_REF] Bischoff | The old, unique C1 chondrite Flensburg -Insight into the first processes of aqueous alteration, brecciation, and the diversity of water-bearing parent bodies and lithologies[END_REF][START_REF] Raymond | Planet formation: key mechanisms and global models[END_REF]), yet have often been overlooked or treated in unrelated contexts (e.g., [START_REF] Lauretta | Meteorites and the early solar system II[END_REF]. The classical model of planetary formation involves three main steps, all entailing collisional processes. In the first, dust particles experience low-velocity collisions while migrating toward the disk's midplane, thus forming fractal aggregates that subsequently form pebbles through collisioninduced compaction [START_REF] Blum | The Growth Mechanisms of Macroscopic Bodies in Protoplanetary Disks[END_REF][START_REF] Johansen | Rapid planetesimal formation in turbulent circumstellar disks[END_REF]. The second step is characterized by the collision-induced coagulation of planetesimals, which forms massive planetary embryos and cores (Morbidelli et al. 2016). In the third step, gravitational instabilities induce collisions between planetary embryos, forming terrestrial planets on a time scale of 10 7 -10 8 years [START_REF] Raymond | Planet formation: key mechanisms and global models[END_REF][START_REF] Raymond | Building the terrestrial planets: Constrained accretion in the inner solar system[END_REF]. Furthermore, the Moon, whose presence next to the Earth has shaped its evolution, is the result of a giant impact within the first 100 million years of the formation of CAIs [START_REF] Canup | Forming a Moon with an Earth-like Composition via a Giant Impact[END_REF][START_REF] Canup | Origin of the Moon in a giant impact near the end of the Earth's formation[END_REF]. Fundamentally, all steps of planetary formation involve collisional processes, implying that they played a key role in shaping the present architecture of the solar system and that they are not an epiphenomenon that can be treated independently of other processes.

Collisions have occurred continually through the present and are now the main factor driving the evolution of the asteroid belt, planetary and moon surfaces, and the formation and sustainability of Saturn's rings [START_REF] Crida | Formation of Regular Satellites from Ancient Massive Rings in the Solar System[END_REF]. This is exemplified by the complex structures of boulders on near-Earth asteroids (101955) Bennu and (162173) Ryugu, as recently revealed by high-spatial-resolution images obtained by NASA's and JAXA's sample return missions OSIRIS-REx and Hayabusa2, respectively (DellaGiustina et al. 2019;[START_REF] Jaumann | Images from the surface of asteroid Ryugu show rocks similar to carbonaceous chondrite meteorites[END_REF]Lauretta et al. 2019aLauretta et al. , 2019b;;[START_REF] Sugita | The geomorphology, color, and thermal properties of Ryugu: Implications for parent-body processes[END_REF]. In particular, the discovery of craters in the surfaces of Bennu's boulders revealed the effects of meteoroid impacts on this asteroid, which occurred only in the last 1.75 Myr [START_REF] Ballouz | Bennu's near-Earth lifetime of 1.75 million years inferred from craters on its boulders[END_REF].

Furthermore, recent analyses of the main asteroid belt have revealed that only a very limited number of the asteroids that we catalogue today have survived intact from primordial times [START_REF] Delbo | Identification of a primordial asteroid family constrains the original planetesimal population[END_REF], whereas the large majority of asteroids are fragments [START_REF] Dermott | The common origin of family and non-family asteroids[END_REF]) that experienced disruption of their parent bodies and re-accumulation [START_REF] Michel | Collisions and Gravitational Reaccumulation: Forming Asteroid Families and Satellites[END_REF][START_REF] Michel | Collision and gravitational reaccumulation: Possible formation mechanism of the asteroid Itokawa[END_REF][START_REF] Michel | Collisional formation of top-shaped asteroids and implications for the origins of Ryugu and Bennu[END_REF]. In this report, we thus aim to (i) detail the astrophysical context of collisional processes in meteorites and (ii) put recent astrophysical and planetological observations in a chondritic perspective.

2-Meteoritic evidence of widespread collision processes in the asteroid belt

Asteroidal collisions represent a major evolutionary process for solid bodies in the solar system and have played a key role in establishing chondritic features, even in the putative 'primitive' type-2 and -3 chondrites. This is attested by the ubiquitous occurrence of brecciated textures in meteorites, with ~30% of ordinary chondrites (OCs) and almost all CM carbonaceous chondrites showing brecciation features (Fig. 1A; Bischoff et al. 1998;[START_REF] Brearley | The action of water[END_REF][START_REF] Gattacceca | Young asteroid mixing revealed in ordinary chondrites: The case of NWA 5764, a polymict LL breccia with L clasts[END_REF][START_REF] Lentfort | Classification of CM chondrite brecciasimplications for the evaluation of samples from the OSIRIS-REx and Hayabusa 2 missions[END_REF][START_REF] Suttle | The aqueous alteration of CM chondrites, a review[END_REF]. Such characteristics are fundamental because (i) OCs are by far the most common meteorites falling on Earth (~80%) and (ii) CMs represent 25% of all carbonaceous chondrite falls [START_REF] Gounelle | Hydrogen isotopic composition of water from fossil micrometeorites in howardites[END_REF]. Even rarer chondrite types like CI carbonaceous chondrites [START_REF] Alfing | Modal abundances of coarse-grained (>5 μm) components within CI-chondrites and their individual clasts -Mixing of various lithologies on the CI parent body(ies)[END_REF]), Rumuruti chondrites [START_REF] Bischoff | The Rumuruti chondrite group[END_REF]), C1 chondrites [START_REF] Bischoff | The old, unique C1 chondrite Flensburg -Insight into the first processes of aqueous alteration, brecciation, and the diversity of water-bearing parent bodies and lithologies[END_REF], CB/CH chondrites [START_REF] Briani | Pristine extraterrestrial material with unprecedented nitrogen isotopic variation[END_REF][START_REF] Bonal | Chondritic lithic clasts in the CB/CH-like meteorite Isheyevo: Fragments of previously unsampled parent bodies[END_REF], and enstatite chondrites (Fig. 1B; [START_REF] Piani | Magmatic sulfides in the porphyritic chondrules of EH enstatite chondrites[END_REF][START_REF] Piani | Earth's water may have been inherited from material similar to enstatite chondrite meteorites[END_REF] have brecciated textures, thus implying that all chondrite types have been reprocessed after the accretion of their respective parent bodies.

Different types of brecciation are commonly observed among meteorites, revealing that collision and re-agglomeration processes occurred at varying scales and intensities [START_REF] Bischoff | Nature and origins of meteoritic breccias[END_REF][START_REF] Rubin | The Adhi Kot breccia and implications for the origin of chondrules and silica-rich clasts in enstatite chondrites[END_REF]. Monomict breccias are the most common and correspond to meteorites formed by the in situ progressive fracturing and compaction of a single lithology at a scale much smaller than their asteroidal parent bodies. Impact-melt breccias (Fig. 1B) are rocks in which the matrix cementing the brecciated fragments crystallized from impactgenerated melts. Genomict breccias define meteorites containing fragments of the same compositional group, but with different metamorphic or hydrothermal alteration histories (Fig. 2; [START_REF] Gattacceca | Young asteroid mixing revealed in ordinary chondrites: The case of NWA 5764, a polymict LL breccia with L clasts[END_REF][START_REF] Vacher | Collisional and alteration history of the CM parent body[END_REF][START_REF] Verdier-Paoletti | Testing the genetic relationship between fluid alteration and brecciation in CMchondrites[END_REF][START_REF] Stöffler | Shock metamorphism of planetary silicate rocks and sediments: Proposal for an updated classification system[END_REF]; these textures reflect large-scale brecciation processes occurring within a given parent body. In contrast, polymict breccias define meteorites composed of fragments originating from different parent bodies (Fig. 3; [START_REF] Bischoff | Nature and origins of meteoritic breccias[END_REF][START_REF] Zolensky | The Kaidun Microbreccia Meteorite: A Harvest from the Inner and Outer Asteroid Belt[END_REF]; these are of fundamental importance as they (i) record asteroidal mixing processes and (ii) contain unique clasts not observed as individual samples [START_REF] Patzek | Mineralogy of volatile-rich clasts in brecciated meteorites[END_REF][START_REF] Patzek | Hydrogen isotopic composition of CI-and CM-like clasts from meteorite breccias -Sampling unknown sources of carbonaceous chondrite materials[END_REF]. The violent evolution of the solar system is also evidenced by commonly observed petrofabrics and/or flattened chondrules and refractory inclusions (Fig. 4; [START_REF] Gattacceca | An impact origin for the foliation of chondrites[END_REF]Greashake 2014;[START_REF] Hanna | Impact-induced brittle deformation, porosity loss, and aqueous alteration in the Murchison CM chondrite[END_REF][START_REF] Rubin | Collisional facilitation of aqueous alteration of CM and CV carbonaceous chondrites[END_REF][START_REF] Vacher | Collisional and alteration history of the CM parent body[END_REF], which are generally interpreted to result from impact-induced preferential orientations. Furthermore, all classes of meteorites contain xenolith clasts (Fig. 5)-characterized by different lithologies than their respective host rocks (e.g., [START_REF] Bischoff | Nature and origins of meteoritic breccias[END_REF][START_REF] Bonal | Chondritic lithic clasts in the CB/CH-like meteorite Isheyevo: Fragments of previously unsampled parent bodies[END_REF][START_REF] Briani | Pristine extraterrestrial material with unprecedented nitrogen isotopic variation[END_REF][START_REF] Briani | Xenoliths and microxenoliths in H chondrites: Sampling the zodiacal cloud in the asteroid Main Belt[END_REF][START_REF] Liu | NanoSIMS isotopic investigation of xenolithic carbonaceous clasts from the kapoeta howardite[END_REF][START_REF] Patzek | Hydrogen isotopic composition of CI-and CM-like clasts from meteorite breccias -Sampling unknown sources of carbonaceous chondrite materials[END_REF][START_REF] Zolensky | The Kaidun meteorite: Mineralogy of an unusual CM1 lithology[END_REF]resulting from the low-velocity incorporation (i.e., <1 km s -1 ; [START_REF] Briani | Evidence for an asteroidcomet continuum from simulations of carbonaceous microxenolith dynamical evolution: Carbonaceous microxenoliths and the asteroid-comet continuum[END_REF]) of micrometeoroids by asteroidal parent bodies [START_REF] Briani | Evidence for an asteroidcomet continuum from simulations of carbonaceous microxenolith dynamical evolution: Carbonaceous microxenoliths and the asteroid-comet continuum[END_REF][START_REF] Rubin | On the origin of shocked and unshocked CM clasts in Hchondrite regolith breccias[END_REF].

Such clasts differ from polymict breccias because they could have been incorporated into asteroids throughout the evolution of the solar system. The mixing of the host and xenoliths probably at low-velocity is also attested by the lack of thermal metamorphism between the clasts and their host meteorites (Fig. 6). Interestingly, the very few occurrences of CC clasts in OC chondrites could result from the fragile nature of carbonaceous materials relative to their OC counterparts. Taken together, brecciated textures, xenolith clasts, and petrofabrics thus all attest to the violent evolution of the solar system, both during the first ~5-10 Myr of the SPD's evolution and continuing today.

Finally, the presence of high-pressure minerals in most meteorite groups is strong evidence of pervasive impacts on their parent bodies. Indeed, most of these minerals have been found within or adjacent to dark shock veins. Following the discovery of ringwooditethe spinel-structured polymorph of olivine-in Tenham, a large variety of shock-induced minerals have been discovered, including majorite, akimotoite (the ilmenite-structured polymorph of MgSiO 3 ), and plagioclase with the hollandite structure [START_REF] Sharp | Shock Effects in Meteorites[END_REF]. At present, high-pressure minerals are used to constrain the pressures, temperatures, and timescales of collision-induced shocks.

3-The chaotic origin of the main asteroid belt

The first step of asteroid formation occurred ~4.56 Ga when dust grains within the protoplanetary disk surrounding the Sun accreted to form rocky/icy planetesimals ~100 km in diameter. These were the first solid objects in the solar system whose internal strengths were dominated by self-gravity [START_REF] Johansen | Rapid planetesimal formation in turbulent circumstellar disks[END_REF][START_REF] Klahr | Turbulence Sets the Length Scale for Planetesimal Formation: Local 2D Simulations of Streaming Instability and Planetesimal Formation[END_REF][START_REF] Marcus | Formation of the planets by accretion of planetesimals: Some statistical problems[END_REF][START_REF] Morbidelli | Challenges in planet formation: Challenges in planet formation[END_REF][START_REF] Safronov | Relative sizes of the largest bodies during the accumulation of planets[END_REF]. Although these objects represent the main building blocks of terrestrial planets and the cores of giant planets, their formation conditions and initial size distribution remain a matter of debate [START_REF] Klahr | Turbulence Sets the Length Scale for Planetesimal Formation: Local 2D Simulations of Streaming Instability and Planetesimal Formation[END_REF]. Classical models are based on the bottom-up progressive coagulation of small solids into planets via merging collisions. However, during the disk's evolution, such bottom-up planetesimal accretion is fundamentally limited by the 'meter-sized barrier': metersized planetesimals will experience (i) rapid drift toward the central star and associated evaporation and (ii) destruction during high-speed collisions [START_REF] Wetherill | Formation of the Earth[END_REF][START_REF] Wurm | Growth of planetesimals by impacts at ∼25 m/s[END_REF]. Recent models thus suggest that planetesimals were initially large (100-1,000 km) due to high concentrations of small solid particles in turbulent structures of the SPD [START_REF] Cuzzi | Toward planetesimals: Dense chondrule clumps in the protoplanetary nebula[END_REF][START_REF] Johansen | Rapid planetesimal formation in turbulent circumstellar disks[END_REF][START_REF] Morbidelli | Asteroids Were Born Big[END_REF]). Such models are in agreement with studies reconstructing the size distribution of planetesimals from analyses of the current asteroid population [START_REF] Bottke | The fossilized size distribution of the main asteroid belt[END_REF][START_REF] Delbo | Identification of a primordial asteroid family constrains the original planetesimal population[END_REF][START_REF] Delbo | Ancient and primordial collisional families as the main sources of X-type asteroids of the inner main belt[END_REF][START_REF] Morbidelli | Asteroids Were Born Big[END_REF][START_REF] Tsirvoulis | Reconstructing the size distribution of the primordial Main Belt[END_REF]). However, most known asteroids are much smaller than planetesimals [START_REF] Masiero | Main-belt asteroids with wise /neowise: near-infrared albedos[END_REF], indicating that they are remnants produced by the collisional fragmentation of larger parent bodies, including planetesimals, throughout the history of the solar system. An investigation of the dynamical behavior of asteroids reached the same conclusion [START_REF] Dermott | The common origin of family and non-family asteroids[END_REF]). These results imply that (i) the current asteroid population is the result of a dramatic collisional evolution that must also be reflected in the physical and textural properties of meteorites and (ii) similar to meteorites, the current architecture of asteroid belt cannot be considered as a direct witness of the early solar system.

Most of planetesimal collisions have occurred in the main asteroid belt, between the orbits of Mars and Jupiter (heliocentric distances of ~2-3.2 AU), where most known asteroids are currently located. Contrary to common perceptions, the main belt is mostly empty, containing only 5 × 10 -4 Earth masses (i.e., ~4 % of the mass of the Moon), with asteroids spread over an annular area more than three times as wide as the distance between Mercury and Mars. In addition, about half of the main-belt mass is contained in the four largest asteroids: Ceres, Vesta, Pallas, and Hygiea (DeMeo and Carry 2014). The main belt is dynamically excited, with orbital eccentricities (e) up to 0.3 and inclinations (i) of 0 to >20°, and is characterized by the presence of spectroscopically distinct asteroids that, in general, correspond to distinct compositions (Fig. 7; Bus andBinzel 2002a, 2002b; DeMeo and Carry 2014): C-type asteroids are water-rich and genetically linked to carbonaceous chondrites, whereas S-type asteroids are dry and related to OCs. The compositions of asteroids vary as a function of their orbital semimajor axis (a): S-type asteroids are more abundant in the inner and central main belt (2.1 < a < 2.7 AU) whereas C-type asteroids are more abundant in the outer regions (2.7 < a < 3.2 AU).

Since the discovery of the first asteroids in the early nineteenth century, the main belt was long considered to be a 'living relic' of planetary formation. Canonical planetary formation theories postulated that (i) asteroids remained at the heliocentric distances at which they formed (as for planetesimals; [START_REF] Gradie | Compositional Structure of the Asteroid Belt[END_REF] and (ii) the compositional gradient of the main belt reflects that of the SPD from the inner border of the protoplanetary disk (a few tenths of an AU from the Sun) out to Jupiter [START_REF] Mandell | Exotic Earths: Forming Habitable Worlds with Giant Planet Migration[END_REF]). However, these views began to shift in the mid-2000s. First, it was proposed that iron-meteorite parent bodies could have formed in the region of terrestrial planets before being scattered to the main belt by gravitational interactions with the growing terrestrial protoplanets [START_REF] Bottke | Iron meteorites as remnants of planetesimals formed in the terrestrial planet region[END_REF].

Second, dynamic models suggested that the D-type asteroids observed today in the main belt are not indigenous, but instead correspond to trans-Neptunian objects that were scattered into the main belt by the giant planets during their instability phase [START_REF] Fujiya | Migration of D-type asteroids from the outer Solar System inferred from carbonate in meteorites[END_REF]Levison et al. 2009;[START_REF] Marrocchi | The tumultuous childhood of the Solar System[END_REF][START_REF] Morbidelli | Evidence from the asteroid belt for a violent past evolution of jupiter's orbit[END_REF][START_REF] Tsiganis | Origin of the orbital architecture of the giant planets of the Solar System[END_REF]. Following these pioneering works, main-belt asteroid implantation models were advanced using new dynamic models involving giant-planet formation and migration early in the solar system.

These models were developed to address two principle stumbling blocks of the classical 'static' models: they systematically produce (i) Mars analogs as massive as the Earth, rather than about a tenth of the Earth's mass, and (ii) eccentricity and inclination distributions of asteroids that do not match those of the current belt. Although important uncertainties remain, the following three recent models (Fig. 7) have successfully reproduced these peculiar characteristics and suggest that the entire main belt population of dynamically excited and radially mixed S-and C-type asteroids are not indigenous, but were implanted in their current locations between the terrestrial-and giant-planet regions.

(i) The "grand tack" model proposed that the inward and subsequent outward migrations of Jupiter and Saturn in the protoplanetary disk truncated the disk at 1 AU, limiting the material available for the formation of Mars (Fig. 7; [START_REF] Walsh | A low mass for Mars from Jupiter's early gas-driven migration[END_REF]. In this scenario, an initial population of planetesimals was expelled from the main belt by the inward migration of Jupiter, which plunged through the main belt. The main belt was subsequently repopulated by planetesimals during the outward migration of Jupiter [START_REF] Walsh | A low mass for Mars from Jupiter's early gas-driven migration[END_REF]. In this scenario, planetesimals suffered mutual collisions at velocities > 10 km s -1 [START_REF] Johnson | Timing of the formation and migration of giant planets as constrained by CB chondrites[END_REF]).

(ii) The low-mass asteroid belt model preserves the concept of implanting planetesimals into the main belt from the terrestrial-(S-types) and giant-planet regions (C-types; Raymond and Izidoro 2017a, 2017b). However, contrary to other models such as the grand tack model, these authors showed that the main belt could have initially had a low mass (or could have even been empty) and that large-scale migrations of the giant planets are not required to populate it with the current mixture of asteroid types (Fig. 7). In this model, C-and S-type asteroids were implanted not only from different sources but also at different times: C-type planetesimals that formed in the giant-planet region were scattered inward during the gas phase of the protoplanetary disk (i.e., within a few million years after the formation of CAIs; Raymond andIzidoro 2017a, 2017b), whereas S-type asteroids were implanted from the terrestrial-planet region into the main belt after the gas in the protoplanetary disk had dissipated (i.e., later than ~10 Myr after the formation of CAIs). The implantation of S-type asteroids was slower and more prolonged than that of C-type asteroids, possibly lasting 100 Myr (Raymond and Izidoro 2017b) because, in this model, S-type planetesimals must have been scattered via multiple close encounters with planetary embryos. In such models, Mars' small mass results from a primordial mass deficit between the present-day orbits of Earth and Jupiter [START_REF] Hansen | Formation of the terrestrial planets from a narrow annulus[END_REF]).

(iii) The early instability model recently provided an alternative view in which the giant planets did not undergo large migrations during the gas phase (Fig. 7; [START_REF] Clement | Mars' growth stunted by an early giant planet instability[END_REF]Clement et al. , 2019aClement et al. , 2019b)). In this model, as in the canonical scenario, a sea of planetesimals spanning the inner rim of the protoplanetary disk (out to Jupiter) was perturbed by an early orbital instability of the giant planets around 10 Myr after gas dispersal, which triggered the formation of the terrestrial planets. This instability also removed significant mass from the main-belt region and broadened the ranges of the eccentricities and inclinations of main-belt planetesimals to the current values, thereby starting the violent collisional evolution of these bodies.

Although the current orbital structure (semi-major axis, eccentricity, inclination) of the asteroid belt is consistent with an initially uniform distribution of asteroids [START_REF] Minton | A record of planet migration in the main asteroid belt[END_REF], its self-excitation requires several orders of magnitude more mass than is present today (e.g., O'Brien and Sykes 2011). This implies that the belt experienced violent processes that considerably modified its initial structure. Whatever the process that dominantly shaped the asteroid belt, collisional processes were important and common during the early evolution of the solar system. Depending on the scenario invoked, the main belt reached its current dynamic configuration after tens to hundreds of millions of years (e.g., after giant-planet instability). Since then, collisions have continued to occur with typical impact velocities between asteroids of ~5.3 km s -1 [START_REF] Bottke | The fossilized size distribution of the main asteroid belt[END_REF] and references therein), although the impact velocities define a broad distribution and collisions at higher-thanaverage speeds are not uncommon, particularly for objects with high i and/or e values. One such case is the third-largest asteroid (2) Pallas (i = 34.8°, e = 0.23), for which impact velocities are estimated to be ~12 km s -1 on average [START_REF] Marsset | The violent collisional history of aqueously evolved (2) Pallas[END_REF]. Collisions at such speeds can yield temperatures high enough to induce the dehydration and thermal metamorphism of asteroidal materials, which should be observable in the meteorites that we may receive from these bodies (Amsellem et al. 2020 and references therein). This collisional evolution is responsible for the formation of the families of asteroid fragments (Nesvorny et al. 2015) that we observe in the main belt. The effects of these collisions, which can lead to the mixing of asteroidal materials of different compositions, is also visible in some physical properties of asteroids, as it is detailed in the next section.

4-Evidence for collisional processes from in-situ exploration of asteroids

In situ exploration by space missions and telescopic observations of asteroids provide growing evidence of the effects of impacts and the resulting material mixing produced by the infall of exogenic impacting objects.

The first evidence of the violent collisional evolution of asteroids is the presence of craters on their surfaces. This has been clear since NASA's Galileo mission, the first spacecraft to fly by an asteroid, which imaged (951) Gaspra in 1991 [START_REF] Veverka | Galileo's Encounter with 951 Gaspra: Overview[END_REF]) and showed that its surface is covered by impact craters [START_REF] Stooke | The Surface of Asteroid 951 Gaspra[END_REF]. Since then, impact craters appear to be ubiquitous structures as they have been observed on all imaged asteroids. These craters range in size from the ~500-km-diameter Rheasilvia basin near the south pole of ( 4) Vesta [START_REF] Schenk | The Geologically Recent Giant Impact Basins at Vesta's South Pole[END_REF][START_REF] Thomas | Impact Excavation on Asteroid 4 Vesta: Hubble Space Telescope Results[END_REF]) to decimeter-scale craters in boulders on the surface of the 500-m-diameter near-Earth asteroid (101955) Bennu [START_REF] Ballouz | Bennu's near-Earth lifetime of 1.75 million years inferred from craters on its boulders[END_REF]). On Vesta, the deposition of the so-called 'dark material' that is often associated with craters has been attributed to impacts by carbonaceous asteroids [START_REF] Mccord | Dark material on Vesta from the infall of carbonaceous volatilerich material[END_REF]. This dark material is then (i) diffused throughout the regolith of the surface of Vesta over time [START_REF] Mccord | Dark material on Vesta from the infall of carbonaceous volatilerich material[END_REF]) and (ii) compressed by subsequent impacts on the surface and cemented together with original Vestan material (brighter in color) into breccia-like rocks. This process is in good agreement with the common presence of clasts of carbonaceous chondritic materials within Howardites (e.g., Fig. 5; [START_REF] Buchanan | Carbonaceous chondrite clasts in the howardites Bholghati and EET87513[END_REF][START_REF] Liu | NanoSIMS isotopic investigation of xenolithic carbonaceous clasts from the kapoeta howardite[END_REF][START_REF] Mcsween | HED Meteorites and Their Relationship to the Geology of Vesta and the Dawn Mission[END_REF][START_REF] Van Drongelen | Petrology and oxygen isotopic compositions of clasts in HED polymict breccia NWA 5232[END_REF], regolith breccia meteorites belonging to the HED group and originating from the surface of Vesta.

Because Vesta is a differentiated asteroid, it has been proposed that impacts could excavate its basaltic crust, revealing its underlying olivine mantle. However, very little olivine has been detected at Vesta's surface, even in and around the giant impact basins Rheasilvia and Venenia [START_REF] Ammannito | Olivine in an unexpected location on Vesta's surface[END_REF][START_REF] Sanctis | Spectroscopic Characterization of Mineralogy and Its Diversity Across Vesta[END_REF]. Therefore, the rare occurrences of olivine signatures on Vesta's surface may rather correspond to material implanted by olivine-rich impactors [START_REF] Turrini | Olivine on Vesta as exogenous contaminants brought by impacts: Constraints from modeling Vesta's collisional history and from impact simulations[END_REF]. Indeed, laboratory impact experiments performed at velocities similar to the average impact velocity on Vesta demonstrated that impactor material could be preserved at Vesta's surface [START_REF] Daly | Delivering a projectile component to the vestan regolith[END_REF]. Although the survivability and implantation of impactor material is a function of impact velocity, laboratory hypervelocity impact experiments on porous and non-porous targets showed that impact-generated material implantation is a general process occurring throughout the solar system [START_REF] Avdellidou | Survival of the impactor during hypervelocity collisions -I. An analogue for low porosity targets[END_REF][START_REF] Avdellidou | Survival of the impactor during hypervelocity collisions -II. An analogue for high-porosity targets[END_REF]. This implies that the mixing of asteroidal materials was possible not only during the early gas phase of the SPD when impact velocities were lower, but throughout the history of the solar system. Interestingly, high-speed collisions (i.e., 12 km s -1 ; [START_REF] Marsset | The violent collisional history of aqueously evolved (2) Pallas[END_REF]) can also yield temperatures high enough to induce the dehydration and thermal metamorphism of asteroidal materials (Amsellem et al. 2020 and references therein).

Very recent observations at cm-scale spatial resolution of the surfaces of the near-Earth asteroids (101955) Bennu and (162173) Ryugu by NASA's OSIRIS-REx (Lauretta et al. 2019a(Lauretta et al. , 2019b) and JAXA's Hayabusa2 [START_REF] Okada | Highly porous nature of a primitive asteroid revealed by thermal imaging[END_REF][START_REF] Tatsumi | Collisional history of Ryugu's parent body from bright surface boulders[END_REF]) sample return missions, respectively, provide a wealth of indications about the mixing of asteroidal materials and the presence of brecciated rocks on these asteroids. OSIRIS-REx spectroscopic and spectrophotometric observations revealed the presence of bright, pyroxene-rich boulders of Vesta-like composition on the dark, carbonaceous surface of Bennu (DellaGiustina et al. 2020). This suggests that some asteroids of the Vesta collisional family impacted Bennu, or more likely its parent body, mixing the two materials [START_REF] Dellagiustina | Variations in color and reflectance on the surface of asteroid (101955) Bennu[END_REF]. This is also attested by the peculiar textures of some of Bennu's pyroxene-rich, hummocky boulders, which show potential layering or brecciation [START_REF] Dellagiustina | Variations in color and reflectance on the surface of asteroid (101955) Bennu[END_REF]. Furthermore, two bright, pyroxene-bearing clasts appear embedded within a large, partially buried boulder, the latter having an albedo similar to the average of Bennu (DellaGiustina et al. 2020).

Interestingly, the Hayabusa2 orbiter also detected bright material on the dark surface of Ryugu [START_REF] Tatsumi | Collisional history of Ryugu's parent body from bright surface boulders[END_REF], thus suggesting that brecciation is a common process during asteroidal evolution.

In contrast to Bennu, the spectral signatures of Ryugu indicate two distinct units of bright material. One appears to be more compatible with an S-type asteroidal composition (rather than a pyroxene-rich, Vesta-like nature; [START_REF] Tatsumi | Collisional history of Ryugu's parent body from bright surface boulders[END_REF]) that could have been delivered by an impact (or series of impacts) on Ryugu's parent body. The second unit has spectral features similar to the bulk of Ryugu, but with a less pronounced upturn at ultraviolet frequencies compared to the Ryugu average [START_REF] Tatsumi | Collisional history of Ryugu's parent body from bright surface boulders[END_REF]. This unit may represent a more thermally metamorphosed material compared to that covering most of the surface of the asteroid [START_REF] Tatsumi | Collisional history of Ryugu's parent body from bright surface boulders[END_REF]. It thus appears that Ryugu's surface comprises materials originating from various depths within its parent body. Such diverse materials may also be present on other similar asteroids, compacted together by small-scale impacts to form breccias.

OSIRIS-REx images of the surface of Bennu show diverse rocks including (i) boulders with albedo and/or color heterogeneities and (ii) non-fractured boulders containing meter-sized clasts (Lauretta et al. 2019a). Such assemblages have been interpreted as impact breccias (e.g., see Fig. 1b of [START_REF] Walsh | Craters, boulders and regolith of (101955) Bennu indicative of an old and dynamic surface[END_REF]. However, the collisional processes capable of producing breccia-like boulders several meters in size and containing meter-scale clasts are probably too energetic to be withstood by Bennu itself [START_REF] Bischoff | Nature and origins of meteoritic breccias[END_REF]). Hence, these boulders probably formed on Bennu's parent body and may record (i) the parent body's accretion in the SPD [START_REF] Walsh | Craters, boulders and regolith of (101955) Bennu indicative of an old and dynamic surface[END_REF]), (ii) compaction of regolith blocks on the parent body, or (iii) collisional events that fragmented the parent body to form Bennu. The surfaces of some boulders larger than ten meters in size also appear to be "littered" with brighter centimetric rocks that, in some cases, appear to be cemented together [START_REF] Bennett | A high-resolution global basemap of (101955) Bennu[END_REF].

Although the origin of this material is unclear, it may result from (i) impacts on the boulders themselves [START_REF] Bennett | A high-resolution global basemap of (101955) Bennu[END_REF] or (ii) particles falling back onto the asteroid after particleejection events, such as were observed on Bennu (Lauretta et al., 2019b).

To summarize, the presence of clasts and exogenic components in Bennu and Ryugu imply that material-mixing via impacts is a common process on these bodies, and likely throughout the solar system. These observations also give context to what was once a breakthrough in asteroid and meteoritic science, namely the discovery of distinct meteorite classes within the strewn field of the Almahata-Sitta meteorites, which were produced by the atmospheric breakup of asteroid 2008 TC 3 [START_REF] Bischoff | Asteroid 2008 TC3-Almahata Sitta: A spectacular breccia containing many different ureilitic and chondritic lithologies: Almahata Sitta-A spectacular breccia[END_REF][START_REF] Goodrich | Origin and history of ureilitic material in the solar system: The view from asteroid 2008 TC 3 and the Almahata Sitta meteorite[END_REF][START_REF] Goodrich | The first samples from Almahata Sitta showing contacts between ureilitic and chondritic lithologies: Implications for the structure and composition of asteroid 2008 TC 3[END_REF][START_REF] Horstmann | The Almahata Sitta polymict breccia and the late accretion of Asteroid 2008 TC3 -Invited Review[END_REF][START_REF] Jenniskens | The impact and recovery of asteroid 2008 TC3[END_REF]. Before the Almahata-Sitta recovery, the canonical concept was that one class of meteorite originates from one class of asteroid.

These recent spacecraft observations have thus demonstrated that 2008 TC 3 is perhaps not an exception and that individual asteroids can host materials corresponding to diverse meteorite classes.

5-Implications and perspectives

Brecciated textures, xenolith clasts, and petrofabrics in meteorites attest to the chaotic evolution of the solar system, both during the first ~5-10 Myr of the evolution of the SPD and continuing until today. Given our knowledge of the structure of the asteroid belt and our theoretical understanding of the evolution of the solar system, this testifies that so-called primitive materials formed during the evolution of the SPD have been largely reprocessed, implying that primitivity is relative, not absolute [START_REF] Bischoff | Shock stage distribution of 2280 ordinary chondrites-Can bulk chondrites with a shock stage of S6 exist as individual rocks?[END_REF][START_REF] Gounelle | The Asteroid-Comet Continuum: In Search of Lost Primitivity[END_REF][START_REF] Libourel | Search for primitive matter in the Solar System[END_REF]. This is not greatly surprising as the spectra of ~20% of nearby solartype stars of various ages show infrared peaks related to the presence of debris disks, lowmass (i.e., less than the mass of the Earth), low-density, dust-dominated disks that are continuously replenished by collisions between large planetesimals [START_REF] Hughes | Debris Disks: Structure, Composition, and Variability[END_REF].

Debris disks are analogous to our solar system's Kuiper Belt and zodiacal light and are fundamentally different from protoplanetary disks, which are dominated by agglomeration processes whereas debris disks are destruction-dominated. This implies that early-formed planetesimals (i) were largely reprocessed during the evolution of the solar system and (ii) experienced fragmentation and subsequent reassembly as asteroids observed today. Indeed, the astronomical survey of the asteroid belt suggests that only a few hundreds of asteroids have remained unchanged since their formation [START_REF] Delbo | Identification of a primordial asteroid family constrains the original planetesimal population[END_REF], clearly indicating that solar system materials have been largely reprocessed via impacts. In addition, planetesimals that accreted during the two million years following CAI formation likely experienced melting due to 26 Al decay [START_REF] Bizzarro | Rapid Timescales for Accretion and Melting of Differentiated Planetesimals Inferred from 26Al-26Mg Chronometry[END_REF][START_REF] Hevey | A model for planetesimal meltdown by 26 Al and its implications for meteorite parent bodies[END_REF][START_REF] Sahijpal | Numerical simulations of the differentiation of accreting planetesimals with 26 Al and 60 Fe as the heat sources[END_REF], which, at least partially (i.e., primitive achondrites), erased their primordial features.

The complex evolution of asteroids is an ongoing process; recent reports suggest that impact-generated fluid circulations, probably triggered by hypervelocity impacts, have occurred within the past few hundreds of thousands of years [START_REF] Amsellem | Timing of thermal metamorphism in CM chondrites: Implications for Ryugu and Bennu future sample return[END_REF][START_REF] Turner | Carbonaceous chondrite meteorites experienced fluid flow within the past million years[END_REF]. This is also highlighted by the ubiquitous presence of xenolith clasts in chondrites and achondrites [START_REF] Bischoff | Nature and origins of meteoritic breccias[END_REF][START_REF] Bonal | Chondritic lithic clasts in the CB/CH-like meteorite Isheyevo: Fragments of previously unsampled parent bodies[END_REF][START_REF] Briani | Pristine extraterrestrial material with unprecedented nitrogen isotopic variation[END_REF][START_REF] Briani | Xenoliths and microxenoliths in H chondrites: Sampling the zodiacal cloud in the asteroid Main Belt[END_REF][START_REF] Cohen | Feldspathic clast populations in polymict ureilites: Stalking the missing basalts from the ureilite parent body[END_REF][START_REF] Liu | NanoSIMS isotopic investigation of xenolithic carbonaceous clasts from the kapoeta howardite[END_REF][START_REF] Patzek | Mineralogy of volatile-rich clasts in brecciated meteorites[END_REF][START_REF] Patzek | Hydrogen isotopic composition of CI-and CM-like clasts from meteorite breccias -Sampling unknown sources of carbonaceous chondrite materials[END_REF][START_REF] Zolensky | The Kaidun meteorite: Mineralogy of an unusual CM1 lithology[END_REF], which may have been incorporated into their current bodies by implantation processes and/or regolith gardening throughout the evolution of the solar system. Therefore, 'primitive' samples are probably extremely rare, and the concept of primitivity should only be used to relatively describe two or several objects, and always in terms of processes.

A direct consequence is the limitation of the so-called bulk composition of meteorites, as -depending on the scale sampledthe average composition could have been modified by various processes operating all along the solar system evolution. For example, although CI chondrites are commonly used as a reference for chemical normalization [START_REF] Lodders | Solar system abundances and condensation temperatures of the elements[END_REF], CIs show brecciated structures and are significantly heterogeneous at a scale of just ~100 mg [START_REF] Morlok | Brecciation and chemical heterogeneities of CI chondrites[END_REF] or up to 1 g for elements sensitive to fluid circulation [START_REF] Barrat | Geochemistry of CI chondrites: Major and trace elements, and Cu and Zn Isotopes[END_REF].

Therefore, virtually all chondrites (and some differentiated meteorites) should be viewed as reprocessed, heterogeneous materials, and analyses of their bulk chemical and isotopic compositions can only provide usable results when specifying the sampling scale. Taken together, all these observations suggest the need for a new perspective on meteorites, considering (i) highly brecciated meteorites (i.e., Kaidun and Almahata Sitta, typically viewed as anomalies) as the norm and (ii) monomict breccias as the exception. In addition, the numerous outstanding questions regarding the origin, distribution, and evolution of xenolith clasts require (i) extensive mineralogical, petrographic, and isotopic characterizations of xenolith clasts (especially in OCs, for which only a limited number have been described), (ii) an understanding of the dynamics and chronology of their incorporation into their host bodies, and (iii) further development of the physics of small-scale implantation processes. [START_REF] Bischoff | Nature and origins of meteoritic breccias[END_REF][START_REF] Rubin | Impact features of enstatite-rich meteorites[END_REF]. Courtesy of Laurette Piani.

Figure captions

Fig. 1 :

 1 Fig. 1: A) Reflected-light photograph of a piece of the Boriskino CM chondrite, courtesy of

Fig. 2 :

 2 Fig. 2: Back-scattered electron image of the brecciated Cold Bokkeveld CM chondrite. In-

Fig. 3 :

 3 Fig. 3: Energy-dispersive X-ray elemental map of the Mg, Al, Si, Ca, and Fe distributions in

Fig. 4 :

 4 Fig. 4: Back-scattered electron image of CV3 Efremovka, showing a high degree of

Fig. 5 :

 5 Fig. 5: Back-scattered electron images of xenolith clasts (central in each image) observed in

Fig. 6 :

 6 Fig. 6: A) Back scattered electron image of the Paris chondrites showing the presence of

Fig. 7 :

 7 Fig. 7: Schematic comparison of the three global models currently proposed for explaining

  Fig. 1

Acknowledgments

We thank Laurette Piani and Maximilien J. Verdier-Paoletti for providing chondrite pictures. This work is the result of intense scientific discussions during the February 2020 workshop "First solids and planetesimals: formation conditions and evolution" hosted in les Houches (France) and organized by the Programme National de Planétologie (PNP). Nan Liu and Markus Patzek are warmly thanked for positive and constructive reviews and we thank Tim Jull for handling the article. This is CRPG contribution #2751.