A Spectroscopic Survey of Primitive Main Belt Asteroid Populations
University of Central Florida STARS
Electronic Theses and Dissertations, 2020-
2021
A Spectroscopic Survey of Primitive Main Belt Asteroid Populations
Anicia Arredondo-Guerrero University of Central Florida
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STARS Citation Arredondo-Guerrero, Anicia, "A Spectroscopic Survey of Primitive Main Belt Asteroid Populations" (2021). Electronic Theses and Dissertations, 2020-. 469. https://stars.library.ucf.edu/etd2020/469 A SPECTROSCOPIC SURVEY OF PRIMITIVE MAIN BELT ASTEROID POPULATIONS
by
ANICIA ARREDONDO B.A. Astrophysics, Wellesley College, 2016
A dissertation submitted in partial fulfilment of the requirements for the degree of Doctor of Philosophy in Physics, Planetary Sciences Track in the Department of Physics in the College of Sciences at the University of Central Florida Orlando, Florida
Spring Term 2021
Major Professor: Humberto Campins © 2021 Anicia Arredondo
ii ABSTRACT
Primitive asteroids have remained mostly unprocessed since their formation, and the study of these populations has implications about the conditions of the early solar system and the evolution of the asteroid belt. This spectroscopic study of inner main-belt (IMB) primitive asteroids addresses three central objectives: 1) determine the origin and composition of objects in the near-Earth object population, particularly spacecraft targets; 2) test theories of how processes such as space weath- ering and aqueous alteration affect surface properties of small, low-albedo bodies; and 3) explore how primitive objects in the background population (i.e., asteroids not belonging to dynamical families) relate to each other and their implications for the evolution of the asteroid belt. In this work, I use the NASA Infrared Telescope Facility and the Telescopio Nazionale Galileo to obtain near-infrared (NIR; 0.7 to 2.5 microns) spectra of objects from three families and the background population. I compare the sample spectra with the published spectra of near-Earth objects and dy- namical studies to test arguments for origin. I compare the VNIR spectra with laboratory spectra of meteorites to constrain the asteroid compositions. I test for space-weathering effects by compar- ing the spectra of the younger families with the older, more-weathered families. I look for trends between the spectra of objects in the background family and their physical and orbital properties to uncover information about this primordial population at the time of formation and throughout its evolution. Chapter 3 describes the NIR characterization of the Klio family. Chapter 4 describes the NIR characterization of the Chaldaea family and its relationship to the Klio family. In Chapter 5, I characterize the Sulamitis family and compare with the Polana family. Finally, in Chapter 6 I characterize the primitive background population and compare the background objects with the families at similar locations.
iii Dedicated to every Latina astronomer who paved the way for me.
iv ACKNOWLEDGMENTS
I have many people to thank for their support and guidance throughout the time spent working on this dissertation. First and foremost, thank you to Dr. Humberto Campins for poaching me on my first visit to UCF and for advising me for the last four years. Your passion for science is infectious and I have enjoyed learning from you. Gracias a todo el grupo PRIMASS por todas las cosas que me han ensenado˜ y por perdonar el hecho de que no hablo espanol.˜
Rose and Rayna, my best pals, there are really no words to express how much I love and appreciate you so I will buy you each a pizza instead. Karley, thank you for being by my side throughout the writing process and for keeping the house at 65° so I can actually sleep. Mom, thank you for raising me in a way that made me confident that I can achieve anything that I put my mind to.
Thank you to everyone in the UCF Planetary group, especially my cohort: Amy, Michael, Keanna, and Isabel. Thank you to Jenny and Katie for taking me under your wing when I moved to Orlando and for sharing your cats with me. Thank you to my dissertation committee, Dr. Noemi Pinilla- Alonso, Dr. Yan Fernandez, Dr. Dan Britt, and Dr. Cristina Thomas for your time and counsel.
I’d like to extend an enormous thank you to Bobby Bus for having pity on a poor grad student and giving me telescope time when I really needed it. Also thank you to the wonderful TOs and the rest of the IRTF staff for all of your knowledge and assistance.
Lastly, to Dr. Stephen Slivan- You taught me everything I know about observing. There were many times during my PhD work that I would think back to something I learned while working with you. Without you, I wouldn’t be in grad school and I definitely wouldn’t be an astronomer. I am forever grateful to you.
v TABLE OF CONTENTS
LIST OF FIGURES ...... xii
LIST OF TABLES ...... xvi
CHAPTER 1: INTRODUCTION ...... 1
1.1 Primitive asteroids ...... 4
1.2 Aqueous alteration ...... 5
1.3 Near-Earth asteroids ...... 6
1.4 The Inner Main Belt ...... 8
1.5 PRIMitive Asteroid Spectroscopic Survey ...... 12
1.6 Overview of dissertation ...... 14
CHAPTER 2: METHODOLOGY ...... 15
2.1 Spectroscopy ...... 15
2.2 Observations ...... 17
2.2.1 Instruments ...... 19
2.3 Analysis of the reflectance spectra ...... 19
vi 2.3.1 Taxonomy ...... 20
2.3.2 Absorption features ...... 23
2.3.3 Spectral slope ...... 25
2.3.4 Curvature ...... 27
CHAPTER 3: NEAR-INFRARED SPECTROSCOPY OF THE KLIO PRIMITIVE INNER- BELT ASTEROID FAMILY ...... 29
3.1 Introduction ...... 30
3.2 Observations and data reduction ...... 33
3.2.1 IRTF ...... 36
3.2.2 TNG ...... 39
3.3 Results and analysis ...... 40
3.3.1 Characteristics of the sample ...... 40
3.3.2 Spectral slopes and curvature ...... 41
3.3.3 Spectral homogeneity of the Klio sample ...... 43
3.3.4 Comparison with the Polana family ...... 47
3.3.5 Interloper (12) Victoria ...... 48
3.3.6 Comparison with (101955) Bennu and (162173) Ryugu ...... 50
vii 3.4 Discussion ...... 52
3.4.1 The possibility of a third compositional group ...... 52
3.4.2 Spectral differences with Polana: consistent with space weathering? . . . . 54
3.5 Conclusion ...... 55
CHAPTER 4: NEAR-INFRARED SPECTROSCOPY OF THE CHALDAEA ASTEROID FAMILY: POSSIBLE LINK TO THE KLIO FAMILY ...... 58
4.1 Introduction ...... 59
4.2 Observations and data reduction ...... 64
4.2.1 IRTF ...... 67
4.2.2 TNG ...... 69
4.3 Results and analysis ...... 70
4.3.1 Characteristics of the sample ...... 70
4.3.2 Spectral slopes and curvature ...... 74
4.3.3 Spectral homogeneity of the Chaldaea sample ...... 78
4.3.4 Comparison with Klio ...... 80
4.4 Discussion ...... 83
4.4.1 Spectral similarities with Klio – one common parent body? ...... 83
viii 4.4.2 Spectral differences with Klio – inconsistent with space weathering . . . . 85
4.4.3 The existence of interlopers ...... 87
4.5 Summary and Conclusion ...... 88
CHAPTER 5: NEAR-INFRARED SPECTROSCOPY OF THE SULAMITIS ASTEROID FAMILY: SURPRISING SIMILARITIES IN THE INNER BELT PRIMITIVE ASTEROID POPULATION ...... 90
5.1 Introduction ...... 91
5.2 Observations and data reduction ...... 96
5.2.1 IRTF ...... 100
5.2.2 TNG ...... 102
5.3 Results and analysis ...... 102
5.3.1 Taxonomy determination ...... 103
5.3.2 Spectral slopes and curvature ...... 106
5.3.3 Reanalysis of Polana family data ...... 109
5.4 Discussion ...... 111
5.4.1 Spectral homogeneity in the NIR despite diversity in the visible ...... 111
5.4.2 Comparison with other IMB families ...... 112
5.4.3 Comparison with (101955) Bennu and (162173) Ryugu ...... 115
ix 5.5 Summary and conclusion ...... 117
CHAPTER 6: SPECTROSCOPY OF THE INNER BELT PRIMITIVE ASTEROID BACK- GROUND POPULATION ...... 119
6.1 Introduction ...... 120
6.2 Observations and data reduction ...... 124
6.2.1 Visible ...... 125
6.2.2 NIR ...... 130
6.2.2.1 IRTF ...... 136
6.2.2.2 TNG ...... 137
6.3 Results and analysis ...... 142
6.3.1 Characteristics of the sample ...... 142
6.3.2 Taxonomy determination ...... 143
6.3.3 Spectral slopes ...... 145
6.3.4 Curvature ...... 148
6.3.5 Aqueous alteration ...... 149
6.3.6 Comparison with the primitive IMB families ...... 153
6.3.7 Comparison with (101955) Bennu and (162173) Ryugu ...... 160
x 6.3.8 Comparison with meteorites ...... 164
6.3.9 Space weathering ...... 167
6.4 Discussion ...... 172
6.4.1 What the PBF tells us about the evolution of primitive asteroid families . . 172
6.4.2 What the PBF tells us about the origin of the background population and its relationship to the dynamical families ...... 173
6.4.3 What the PBF tells us about the evolution of the asteroid belt ...... 174
6.5 Conclusion ...... 175
CHAPTER 7: SUMMARY ...... 177
7.1 Primitive asteroid families ...... 177
7.2 The primordial background population ...... 179
7.3 Future work ...... 180
APPENDIX : COPYRIGHT PERMISSIONS ...... 184
LIST OF REFERENCES ...... 204
xi LIST OF FIGURES
1.1 The sequence of condensation from a gas of solar composition as it cools. . .2
1.2 The grand tack model...... 3
1.3 Asteroid (253) Mathilde...... 4
1.4 Size comparison of sample return mission targets...... 8
1.5 Primitive families in the inner main asteroid belt...... 9
1.6 Yarkovsky cones of IMB families...... 11
2.1 An example of a primitive and non-primitive VNIR spectrum...... 16
2.2 An example of the calibration spectra necessary to obtain an asteroid re- flectance spectrum ...... 18
2.3 Bus-DeMeo taxonomy key ...... 21
2.4 Concatenation of visible and NIR spectra ...... 22
2.5 An example of taxonomic classification of a VNIR spectrum ...... 23
2.6 An example of the CANA procedure for absorption band fitting ...... 24
2.7 An example of spectral slope determination ...... 26
2.8 An example of spectral curvature determination ...... 28
xii 3.1 Location of primitive asteroid families in the inner main belt ...... 32
3.2 Spectra of the 21 Klio family objects ...... 38
3.3 Example of concavity determination process ...... 43
3.4 Average spectrum of the Klio family ...... 44
3.5 Slope distribution of the Klio and Polana families ...... 45
3.6 Curvature distribution of the Klio family ...... 46
3.7 Yarkovsky cone for the Klio family ...... 49
3.8 Comparison of NEAs with the Klio family ...... 51
4.1 Location of primitive asteroid families in the inner main belt ...... 60
4.2 Yarkovsky cone for the combined Chaldaea-Klio group ...... 63
4.3 Spectra of the 15 Chaldaea family objects ...... 68
4.4 Composite VNIR spectra of 6 Chaldaea objects ...... 71
4.5 Slope distribution of the Chaldaea and Klio families ...... 75
4.6 Curvature distribution of the Chaldaea and Klio families ...... 77
4.7 Average spectrum of the Chaldaea family ...... 79
4.8 Average spectra of the Chaldaea and Klio families ...... 83
5.1 Location of primitive asteroid families in the inner main belt ...... 92
xiii 5.2 NIR spectrum of (302) Clarissa ...... 95
5.3 Spectra of the 19 Sulamitis family objects ...... 101
5.4 Composite VNIR spectra of 16 Sulamitis objects ...... 104
5.5 NIR taxonomy of the Sulamitis and Polana families ...... 105
5.6 Average spectrum of the Sulamitis family ...... 107
5.7 Curvature distribution of the Sulamitis family ...... 109
5.8 Slope distribution of the Sulamitis family ...... 112
5.9 Comparison of NEAs with the Sulamitis family ...... 116
6.1 Location of primitive asteroid families in the inner main belt ...... 121
6.2 Halos of background asteroids around dynamical families ...... 122
6.3 Visible spectra of the 21 PBF objects ...... 128
6.4 Visible slope distribution of objects in the IMB ...... 129
6.5 NIR spectra of the 55 PBF objects ...... 141
6.6 Taxonomic distribution of the PBF visible spectra compared to the other IMB families ...... 144
6.7 Taxonomy and location of PBF objects compared to the other IMB families . 145
6.8 All visible and NIR spectra of the PBF sample ...... 147
xiv 6.9 Slope and location of PBF objects compared to the other IMB families . . . . 148
6.10 Hydration amongst primitive IMB families ...... 150
6.11 Example of the output of the CANA program ...... 151
6.12 Hydration and location of PBF objects compared to the other IMB families . 152
6.13 Possible trend between spectral slope and band center ...... 153
6.14 Slope distributions of primitive IMB populations ...... 155
6.15 Spectral comparisons of the PBF and families at different inclination ranges . 158
6.16 Comparison of NEAs with the PBF ...... 161
6.17 Comparison of NEAs with the low inclination PBF ...... 162
6.18 Exploration of the possible trend between spectral slope and band center . . . 169
6.19 Relationship between diameter and NIR spectral slope ...... 171
7.1 3 µm spectroscopy of primitive asteroids...... 182
xv LIST OF TABLES
1.1 Primitive IMB families characterized by PRIMASS ...... 13
3.1 Physical and dynamical properties of the observed Klio family targets . . . . 34
3.2 Observational details of the Klio targets ...... 35
3.3 Coordinates and magnitudes of solar analog stars used for Klio targets . . . . 36
4.1 Summary of PRIMASS findings in the IMB ...... 61
4.2 Physical and dynamical properties of the observed Chaldaea family targets . . 65
4.3 Observational details of the Chaldaea targets ...... 66
4.4 Coordinates and magnitudes of solar analog stars used for Chaldaea targets . 67
4.5 Taxonomy of the Chaldaea family ...... 73
4.6 Comparison between NIR spectra of the Klio and Chaldaea families . . . . . 82
5.1 Summary of PRIMASS findings in the IMB ...... 93
5.2 Physical and dynamical properties of the observed Sulamitis family targets . . 97
5.3 Observational details of the Sulamitis targets ...... 99
5.4 Coordinates and magnitudes of solar analog stars used for Sulamitis targets . 100
5.5 Taxonomies and curvatures of Polana family objects ...... 110
xvi 5.6 Comparison between PRIMASS samples of the IMB families ...... 113
5.7 KS test to compare IMB families ...... 114
6.1 Physical and dynamical properties of the visible PBF spectra obtained. . . . . 126
6.2 Physical and dynamical properties of the NIR observed PBF targets . . . . . 131
6.3 Observational details of NIR observations...... 133
6.4 Coordinates and magnitudes of solar analog stars used for PBF targets . . . . 136
6.5 Summary of PRIMASS findings in the IMB ...... 154
6.6 KS test to compare IMB families ...... 156
6.7 Meteorite fits for PBF objects ...... 165
6.8 Comparison between PRIMASS samples of the IMB families in the NIR. . . 167
xvii CHAPTER 1: INTRODUCTION
Asteroids are leftover planetesimals from the formation of the solar system. Gravitational interac- tions prevented these planetesimals from growing into planets in the region between Jupiter and Mars. Instead, planetesimals remained relatively small and subsequently went though collisions which either fragmented them or, if the impact velocity was low enough, caused them to loosely stick together. Morbidelli et al. (2009) simulate coagulation caused by collisions in the asteroid belt to try identify the original size frequency distribution (SFD) of initial planetesimals. They find that the current SFD can only be matched if the original size of planetesimals was ∼100 km. This initial size means that most asteroids were not large enough to have differentiated into core, mantle, and crust layers.
The composition of an asteroid depends on where it formed in the solar disk and where it is now in relation to the Sun (e.g., Gradie & Tedesco 1982). Fig. 1.1 shows the types of solids that might condense at certain temperatures. In the temperature region where asteroids can form, Iron-Nickel metals, olivine, pyroxene, and feldspars condense closer to Mars while hydrated silicates, sulfates, and ices condense closer to Jupiter. This difference in initial composition forms the basis for the basic asteroid taxonomy proposed by Chapman et al. (1975), with carbonaceous C-types formed in the outer belt and stony S-types formed in the inner belt.
1 Figure 1.1: The sequence of condensation from a gas of solar composition as it cools. Shaded bars represent the temperatures where a material can condense and dashed lines indicate the lower temperatures that the materials can exist if condensation does not occur. This image was acquired from the web version of Christiansen & Hamblin (1995) (http://explanet.info/Chapter02.htm).
The asteroid belt is now well mixed, with C-types and S-types scattered throughout the belt. This mixing is likely remnant of a large scale event such as a giant planet migration that has been proposed by the Nice model (Tsiganis et al. 2005), the grand tack hypothesis (Walsh et al. 2011) and the jumping-Jupiter scenario (Brasser et al. 2009). An illustration of how planet migration can cause mixing in the asteroid belt is shown in Fig. 1.2. There have also been advances in our
2 classification schemes, with more recent taxonomic studies dividing asteroids into tens of different groups based on their surface properties (e.g., DeMeo & Carry 2014), rather than the 2 basic groups proposed by Chapman et al. (1975). Because of the wide diversity of objects within the asteroid belt, the study of asteroids addresses many different science objectives. Asteroids tell us about the building blocks of planets and the original composition of the solar nebula. Asteroids are also useful indicators of the distribution of volatiles in the solar system, which can inform us of how water arrived at Earth.
Figure 1.2: Example of how C-type asteroids (blue) that formed in the outer solar system and S- type asteroids (red) that formed in the inner system can be scattered by the inward-then-outward migration of giant planets (black) as explained by the grand tack model of Walsh et al. (2011). The dashed lines represent the approximate boundaries of the asteroid belt, and the asteroids inside the belt are now well mixed.
3 1.1 Primitive asteroids
Primitive asteroids are named as such because they have undergone minimal thermal and aqueous alteration, but not differentiation, since their formation. They are composed some of the most primitive material in the solar system. Carbonaceous chondrite meteorites, which contain high percentages of water and organic compounds, are believed to be fragments of primitive asteroids. We therefore believe that primitive asteroids consist of carbon-rich material, which makes their geometric albedos, or the amount of sunlight they reflect, less than 15% (e.g., Campins et al. 2018). The three primitive asteroids that have been imaged by spacecraft, (253) Mathilde, (101955) Bennu, and (162173) Ryugu, reflect 3%, 4%, and 4%, respectively, making all of them darker than a piece of coal (Fig. 1.3).
Figure 1.3: Asteroid (253) Mathilde imaged by the Near Earth Asteroid Rendezvous – Shoemaker (NEAR-Shoemaker) mission. This image has been artificially colorized to match its likely ap- pearance to a human observer by Daniel Macha´cek.ˇ This image was acquired from the Planetary Society (https://www.planetary.org/space-images/mathilde machacek).
4 Primitive asteroids likely formed in the giant planet region and are some of the C-type asteroids that were described in the previous section. Primitive asteroids can also be B-types, X-types, T- types, and D-types. The study of primitive asteroids also teaches us about the nature of mineral and organic material in the outer solar system and how prebiotic material may have been delivered to Earth.
1.2 Aqueous alteration
Because they likely formed past the snow line and then moved closer to the Sun, the ices that formed on primitive inner belt asteroids subsequently melted to form aqueously altered materials such as phyllosilicates, sulfates, oxides, carbonates, and hydroxides. Fornasier et al. (2014) find that the hydration on asteroids is correlated with size, with aqueous alteration processes dominant on asteroids between 50 and 240 km.
The two absorption features in primitive asteroid spectra that are associated with water and hydroxyl- bearing minerals are located at 0.7 µm and 3 µm. The 0.7 µm absorption is attributed to a charge transfer transition in oxidized iron found in phyllosilicates (Vilas & Gaffey 1989). Fornasier et al. (2014) analyze 600 primitive asteroids and characterize their 0.7 µm absorption band parameters. They find that the aqueous alteration sequence starts with P-type asteroids which are mostly un- altered, and increases in aqueous alteration through F-type, B-type, C-type, and G-type asteroids, in that order. They predict that 50% of all C-type asteroids have the 0.7 µm feature, and Vilas & Gaffey (1989) show that the depth of this absorption can be as deep as 5%. Bus et al. (2002) show that there are weaker bands that are also associated with iron oxide minerals centered at 0.43 µm, 0.60-0.65µm, and 0.80-0.90 µm.
5 The 3 µm absorption is formed by a combination of the OH-radical absorption feature and the first overtone of the 6 µm H2O fundamental. The depth of this feature varies based on the abundance, particle size, and temperature. Using 3 µm spectroscopy, we can observe water abundances less than a weight percent (Rivkin et al. 2002). Unfortunately, the 3 µm absorption is difficult to observe from ground based telescopes due to the H2O and CO2 absorption in Earth’s atmosphere, and due to the increased thermal background at these wavelengths. Rivkin et al. (2002) shows that in low- albedo asteroids, the 0.7 and 3 µm absorptions are correlated. The 3 µm feature is much more indicative of hydration, as objects can show no 0.7 µm feature but still have a feature at 3 µm, but not the other way around (Fornasier et al. 2014).
1.3 Near-Earth asteroids
Near-Earth asteroids (NEAs) are main belt asteroids who’s orbits have been perturbed so that they come close to Earth’s. NEAs are both interesting and dangerous. As the closest asteroids to Earth, NEAs are ideal spacecraft targets because of the low delta-v required to reach them. On the other hand, NEAs can come close enough to Earth to potentially collide, causing events like the Chelyabinsk meteor impact in 2013, and the Chixulub impact 65 million years ago. These potentially hazardous asteroids (PHAs) can currently be tracked, but not redirected.
6 There are two recent sample return missions with NEA targets: NASA’s OSIRIS-REx and JAXA’s Hayabusa2. OSIRIS-REx launched in 2016, orbited (101955) Bennu for 2 years, collected a sam- ple of the surface in 2020, and is currently on its way back to Earth. Hayabusa2 launched in 2014, sampled (162173) Ryugu in 2018, and the sample was returned to Earth in late 2020. In addition, (25143) Itokawa was sampled by JAXA’s Hayabusa mission in 2005, with the samples returned to Earth in 2010. Other NEAs that have been visited by spacecraft include the rendezvous with (433) Eros by NASA’s NEAR-Shoemaker in 2001, and the flyby of (4179) Toutatis by the Chinese Chang’e 2 probe in 2012.
These visits give us a glimpse into the diversity of NEAs, as seen in Fig. 1.4. Bennu and Ryugu are small (diameters of 525 and 870 m, respectively) and spheroidal, with extremely rocky terrain and equatorial bulges. Both Bennu and Ryugu are primitive NEAs (B-type and C-type, respectively). Itokawa, on the other hand, is a much larger S-type asteroid. It is elongated like a potato (largest axis 535 m), and has ponds of very smooth looking, fine-grained material. Eros is the largest of the five visited NEAs: it is elongated with a largest axis of 34.4 km. It is an S-type and shows smooth terrain with impact craters as large as 5.3 km. The diversity seen in the NEA sample teaches us about the representativeness and potential bias of our meteorite collection. Because the original source of NEAs is the main asteroid belt, constraining the main belt origins of these objects can offer context for the study of the NEAs themselves.
7 Figure 1.4: Size comparison of sample return mission targets. This image was acquired from NASA.
1.4 The Inner Main Belt
The main source region of NEAs is the inner main belt (IMB), located between the ν6 resonance and the 3:1 mean motion resonance. Collisions between the large primordial planetesimals created fragments of asteroids that follow the same orbital path and have roughly the same composition, called asteroid families. In the IMB, twenty families have been identified using the hierarchical clustering method (Nesvorny´ et al. 2015). Eight of those twenty are classified as primitive based on their low albedo: Polana-Eulalia (hereafter Polana), Clarissa, Erigone, Sulamitis, Klio, Chaldaea,
8 Chimaera, and Svea. A map of the IMB families is shown in Fig. 1.5. The families are named after the largest fragment, which is often assumed to be the largest remnant of the parent body. The study of families teaches us about the heterogeneity of primordial planetesimals, how meteorites and NEAs are delivered, and provide constraints on the dynamical and collisional evolution of the asteroid belt.
Figure 1.5: Semimajor axis versus inclination (left) and eccentricity (right) of primitive families in the inner main asteroid belt.
In the IMB 56% of asteroids do not belong to dynamical families. This population is called the background, and is represented by grey dots in Fig. 1.5. These asteroids can also be primitive in origin, and perhaps belonged to ancient dynamical families that have been erased over time by collisions, gravitational interactions, and Yarkovsky drift (see next paragraph). The background might also be the most important source of information about the IMB, because the background objects may be embryos of planetesimals that teach us about the diversity of material available for planetary accretion.
9 Over time, the Yarkovsky effect causes the small (10 cm to 10 km) IMB asteroids (both family and background) to drift in semimajor axis. This drift causes asteroids to spiral away from the Sun if they are prograde rotators or towards the Sun if they are retrograde rotators. The left plot of Fig. 1.6 shows the result of this drifting - a so called Yarkovsky cone. The orange diamond is the parent body and the black dots are family members. The black dots have drifted in semimajor axis so that there is an equal number on both sides, with smaller, fainter objects being more affected and moving farther from the center of the family. If the fragments drift far enough, they move into one of the two orbital resonances that border the IMB. An example of this is the right plot of Fig. 1.6 where there are significantly less objects farther out from the parent body than closer in. This is because the 3:1 resonance is located at 2.50 AU, which borders the Sulamitis family. When a fragment moves too close to these resonances, gravitational interactions increase the eccentricity of the asteroid, which has the potential to move them into NEA space. Bottke et al. (2015) use dynamical arguments to show that 60% of NEAs are delivered from either the ν6 or 3:1 resonance and Campins et al. (2010, 2013) use dynamics and spectroscopy to show that (101955) Bennu and
(162175) Ryugu were likely delivered through the ν6 resonance.
10 Figure 1.6: A representation of the Yarkovsky effect on members of the Erigone (left) and Sulamitis (right) families. The parent bodies are marked by orange and blue diamonds, respectively, and the black dots indicate family members. The grey lines indicate the Yarkovsky cones which are the boundaries of each family.
Another effect of the Yarkovsky drift is that ancient families may become blended with the back- ground population. Over a long enough period, this can cause families to be erased or very difficult to identify based on their orbits. Therefore, the background may be composed of original plan- etesimals, or it may be remnants of ancient families. The relationship between asteroid families and the background population has been studied with models (e.g., Dermott et al. 2018), but not
11 confirmed by observations. Furthermore, collisions between members of the original dynamical families may have created subfamilies (e.g., the Beagle family is a subfamily of Themis), meaning that the families we see today may be generations removed from the original primordial parent body. This and the fact that asteroids formed big imply that families may share a common origin and perhaps only a few parent bodies existed in the primordial solar system.
1.5 PRIMitive Asteroid Spectroscopic Survey
Because of the relevance of primitive asteroids to current spacecraft missions, the PRIMitive As- teroid Spectroscopic Survey (PRIMASS) was established to characterize possible source regions of primitive NEAs and offer a broader context for space missions. PRIMASS uses visible and near-infrared (NIR) spectroscopy to look for evidence of hydration and spectral diversity (or lack thereof) in different primitive asteroid populations. PRIMASS has already characterized all eight IMB families in the visible wavelength range (de Leon´ et al. 2016, Morate et al. 2016, 2018, 2019) and a summary of those results is given in Table 1.1.
12 Table 1.1: Primitive IMB families characterized by PRIMASS. Klio age is from Carruba & Nesvorny´ (2016) and all other ages are from Bottke et al. (2015). All other information is from the respective PRIMASS publications.
Family Inc (deg) Age (Myr) Visible Spectra 0.7 µm feature
Polana 2 1400 ± 150 Homogeneous Absent Clarissa 2 ∼60 Homogeneous Absent Erigone 5 130 ± 30 Diverse Present (58%) Sulamitis 5 200 ± 40 Diverse Present (60%) Klio 9 960 ± 250 Diverse Present (23%) Chaldaea 12 Homogeneous Present (79%) Chimaera 14 Diverse Present (20%) Svea 16 Homogeneous Absent
The lower inclination families (Polana, Clarissa, Erigone, and Sulamitis) fall into two distinct compositional groups: “Erigone-like”, which are hydrated (∼60% of the members) and spectrally diverse, and “Polana-like”, which have no 0.7 µm absorption and are spectrally homogeneous (Campins et al. 2018). The “Polana-like” group have inclinations near 2° while the “Erigone-like” group cluster tightly near 5°, suggesting common primordial parent bodies for each of the groups. The Polana family has also been characterized at NIR wavelengths by Pinilla-Alonso et al. (2016), and shows homogeneous, blue slopes for all asteroids in the sample.
The higher inclination families (Klio, Chaldaea, Chimaera and Svea), show a more complex picture than the two distinct groupings at lower inclinations. Despite overlapping somewhat in inclination, Klio and Chaldaea differ in both heterogeneity and the percentage of hydrated objects (23% for Klio and 79% for Chaldaea), and do not fit well into either of the “Polana-like” or “Erigone-like” categories. The Chimaera and Svea families show even greater differences in their visible spectra:
13 Chimaera is spectrally diverse and is 20% hydrated while Svea has homogeneous spectral slopes and no evidence of hydration. These sharp differences are suggestive of different primordial parent bodies for each family. The Svea family does fit the “Polana-like” group, but with a 14° higher inclination. The Chimaera family doesn’t fit into either of the lower inclination groups, but has a similar amount of hydration and spectral diversity as the Klio family.
1.6 Overview of dissertation
This dissertation extends the PRIMASS survey by characterizing the Klio, Chaldaea, and Sulamitis families at NIR wavelengths. The chosen families can be used to determine if the “Polana-like” and “Erigone-like” groups are also present at longer wavelengths and if there are any distinguishing features in the NIR that can differentiate between the families. In the next chapter I describe the methodology used to observe and characterize these populations. Chapter 3 shows the results for the Klio family. Chapter 4 shows the results for the Chaldaea family and its relationship to the neighboring Klio family. Chapter 5 shows the results for the Sulamitis family and explores the relationship between the low inclination families. In addition, this dissertation presents both visible and NIR spectra of the background population of asteroids. These background objects are compared with each other and also with the families. To date, this is the largest spectroscopic survey and most detailed study of the primitive IMB background. Chapter 6 details our findings, including that the background population is at least partially related to the primitive families.
14 CHAPTER 2: METHODOLOGY
2.1 Spectroscopy
Spectroscopy is a powerful tool for constraining asteroid surface composition via remote sensing measurements. The grains on an asteroids surface can absorb or reflect incident sunlight. The amount of light that is reflected is a property of the optical characteristics of its surface grains. Absorption features in reflectance spectra are due to electronic and vibrational transitions within specific mineral species (e.g., Gaffey et al. 1993). The wavelength of each absorption feature depends on the material present and can therefore be diagnostic of composition. For example, pyroxene absorbs near 1 and 2 µm, water-bearing minerals absorb near 3 µm, and spinel has an absorption feature at 1.5 µm. If these features are present in a spectrum, the asteroid may contain those materials. In practice, however, feature matching is sometimes not sufficient and more de- tailed analyses, including spectral modeling, are used to constrain composition and other surface characteristics.
The spectra presented in this dissertation are in the visible (0.4-0.9 µm) and near-infrared (NIR; 0.8-2.5 µm) wavelength ranges. At these wavelengths, asteroids in the main belt do not emit any radiation, but only reflect sunlight. The combination of VNIR spectra allows for a better constraint on composition and taxonomy, and is done whenever possible.
15 For silicate-rich asteroids such as S-types, the mineralogy derived from their spectra tends to be unambiguous. For primitive asteroids, however, the spectra lack diagnostic absorptions, and there are fewer features to fit. An example of the spectra of a primitive and non-primitive asteroid is given in Fig. 2.1. The non-primitive asteroid, (6) Hebe, shows two strong absorption features as- sociated with mafic minerals, more specifically olivine and pyroxene. The primitive asteroid, (54) Alexandra, does not have any strong absorption features. There is a shallow feature in Alexandra’s spectrum that is centered on 0.7 µm. This feature is associated with the charge transfer transition in oxidized iron found in phyllosilicates (Vilas & Gaffey 1989).
Figure 2.1: An example of a primitive and non-primitive VNIR spectrum. The spectrum of (6) Hebe shows two strong absorption features near 1 and 2 µm, indicative of iron bearing silicates. The spectrum of (54) Alexandra is less diagnostic. These two VNIR spectra were obtained from the Catalog of Asteroid Spectra by the MITHNEOS MIT-Hawaii Near-Earth Object Spectroscopic Survey.
16 2.2 Observations
The NIR spectra presented in this thesis were obtained using two different telescopes. Most of the spectra were obtained using the NASA Infrared Telescope Facility (IRTF) in Hawaii, and some were observed using the Telescopio Nazionale Galileo located on the island of La Palma in Spain.
We selected targets for observation based on the lists of asteroid families made by Nesvorny´ et al. (2015) and Delbo´ et al. (2017). We prioritized objects in these lists that had visible spectra available from SMASS, S3OS2, or PRIMASS (Bus & Binzel 2002, Lazzaro et al. 2004, Morate et al. 2018, 2019). We typically observed asteroids with visual magnitude V<18 and at an airmass X<1.3. We varied the exposure time based on the magnitude of the target and sky conditions, and we observed in an ABBA scheme where A is the position of the object in the slit during the first exposure and B is a position shifted along the slit. The ABBA scheme was repeated to increase the signal to noise based on the magnitude of the object.
For calibration of the final reflectance of the asteroid, we observed a local standard star for each object as well as at least one solar analog for each night of observations. To subtract the telluric absorption from the Earth’s atmosphere, we chose one local standard star for each asteroid with spectral type G2-5V and solar-like B–V and V–K colors and observed it at an airmass similar to that of the asteroid. To remove the solar contribution(i.e., the portion of the reflectance that is from the Sun) from the asteroid spectrum, we observed solar analogs, which are stars with spectra very similar to that of the Sun. Each asteroid spectrum and solar analog spectrum are divided by the local standard star spectrum to remove the telluric contribution from both. The modified asteroid spectrum is then divided by the modified solar analog spectrum to obtain the final reflectance. This reflectance spectrum is normalized to unity at 1.0 µm. An example of the spectra of an asteroid, local standard star, and solar analog are given in Fig. 2.2 along with an example of a final reflectance spectrum of an asteroid.
17 Figure 2.2: An example of the calibration spectra necessary to obtain an asteroid reflectance spec- trum. The top left panel shows a raw asteroid spectrum. The top middle and top right panels show raw spectra of a local standard star and a solar analog star, respectively. The bottom panel is an example of a final reflectance spectrum after the stars have been divided out from the asteroid spectrum.
18 2.2.1 Instruments
The IRTF is a 3.2 meter telescope funded by NASA and operated and managed by the University of Hawaii Institute for Astronomy. The telescope is sited on Mauna Kea in Hawaii and is opti- mized for infrared observations. As of 2018, 95% of IRTF science is conducted with wide-band spectrographs. We used the SpeX instrument (Rayner et al. 2003) which is a medium-resolution spectrograph that covers the wavelength range 0.7-5.3 µm. Specifically, we used the PRISM mode which covers 0.7-2.52 µm and has a resolving power of R∼200. We also used the MORIS Gulbis et al. (2011) CCD camera attached to the side of SpeX for guiding on our faint targets. The IRTF can be controlled remotely, and all of the observations presented in this dissertation were done from Orlando, Florida.
The Telescopio Nazionale Galileo (TNG) is a 3.58 meter telescope funded by the Italian Istituto Nazionale di Astrofisica (INAF) and managed by The Fundacion´ Galileo Galilei. The telescope is sited at the top of the Roque de Los Muchachos Observatory in the Canary Islands of Spain. The TNG is capable of observing at both visual and infrared wavelengths, however we only used it for infrared. We used the NICS instrument (Baffa et al. 2001) which covers a wavelength range of 0.9-2.5 µm with a resolving power of R∼35 (Oliva 2000). Observations using the TNG were obtained in person by my colleague and co-author, Vania Lorenzi.
2.3 Analysis of the reflectance spectra
The methods for analysis of the final reflectance spectra include taxonomic classification, mea- surement of absorption bands, and calculation of spectral slope and spectral curvature.
19 2.3.1 Taxonomy
We analyze each of our final reflectance spectra to categorize them into taxons based on the Bus- DeMeo taxonomy (DeMeo et al. 2009) which extends the well-known Bus taxonomy scheme to the NIR. This taxonomic system is split into three taxonomic groups (Fig. 2.3). The S-complex in- cludes Tholen S-type asteroids that are composed mainly of olivine and pyroxene. The C-complex includes Tholen B-type, C-type, F-type, and G-type asteroids and are closely related to carbona- ceous chondrite meteorites. The X-complex includes Tholen E-type, M-type, and P-type asteroids which are distinguished by their albedos. Finally, there is an end member group of spectrally distinct types, including T-types, D-types, V-types, and A-types. As mentioned in Section 2.1, primitive asteroids are mostly in the C-complex and X-complex and do not exhibit as many spec- tral features as those in the S-complex.
20 Figure 2.3: Bus-DeMeo taxonomy key from DeMeo et al. (2015). Spectra presented in this work are mostly from the C-complex and X-complex.
The taxonomic classification is done using an online tool created by Popescu et al. (2012) called Modeling for Asteroids (M4AST). M4AST consists of an asteroid spectral database of about 2700 spectra from various sources as well as a set of applications for analyzing asteroid spectra. We upload our visible and NIR spectra to the website, and then concatenate the spectra if both visible and NIR spectra are available. An example of this concatenation is shown in Fig. 2.4.
21 Figure 2.4: Concatenation of visible (blue) and NIR (red) spectra in M4AST.
We then use the concatenated VNIR spectrum for taxonomic classification. M4AST first fits the spectrum with a polynomial curve and then compares that curve to the standard spectrum of each class. The standard spectrum of the class with the smallest χ2 residual is returned along with the 2 next most close matches. An example of this classification is shown in Fig. 2.5. The three best matches for the spectrum are Ch-type, Xc-type, and Cgh-type spectra. The right image shows the comparison of the spectrum with the best match standard Ch-type spectrum, including error bars.
22 Figure 2.5: An example of taxonomic classification of a VNIR spectrum. The left plot shows that the three best matches for the spectrum are Ch-type, Xc-type, and Cgh-type spectra. The right plot shows a closer inspection of the Ch-type match, with both spectra normalized to unity at 1.0 µm.
2.3.2 Absorption features
Each visible spectrum was searched for an absorption band at 0.7 µm , as explained in Section 1.2 to be associated with hydrated minerals. To compute the center and depths of absorption bands, we used a program called Codes for ANalysis of Asteroids (CANA; De Pra et al. 2018). The program first computes a fourth-order polynomial fit to the asteroid spectrum and then fits a straight line tangent to the polynomial fit at the limits of where the absorption band should be. This straight line is assumed to be the continuum of the spectrum, and is divided from the spectrum so that what remains is the band itself (Fig. 2.6). To find the central wavelength of the band, CANA finds the local minimum of the continuum removed spectrum. The band depth is the percent difference between the reflectance at the central wavelength and the reflectance of the continuum (1 in most
23 cases). To make the final band center and depth measurements more robust, CANA repeats the above procedure 1000 times. The mean of the 1000 band centers and 1000 band depths is the final center and depth. The standard deviation of the mean is the error in the final center and depth. To be considered hydrated in this study, an asteroid must present a band centered near 0.7 µm and have a band depth of >1%.
Figure 2.6: An example of the CANA procedure for absorption band fitting. The black line is the asteroid spectrum after the continuum is removed. The blue line is the fourth order polynomial fit to the spectrum. The blue dot indicates the local minimum that is the band center. The calculated values for band minimum and band depth are given in the plot.
24 2.3.3 Spectral slope
One way of characterizing primitive asteroid spectra is by measuring spectral slope. The spectral slope of a primitive asteroid can be indicative of composition and degree of alteration (de Leon´ et al. 2012), and can suggest compositional diversity within a population (Campins et al. 2018). For the visible and NIR wavelength ranges, we calculate spectral slope using the equation
dS/dλ S0 = SλU
where dS/dλ is the rate of change of the reflectance and SλU is the reflectance at 0.55 µm for visible and 1.0 µm for NIR. The units of S0 are %/1000A˚ and is visualized by the pink line in Fig. 2.7. S0 of each asteroid is calculated in the range from 0.55 to 0.90 µm for visible spectra and 0.95 to 2.30 microns for NIR spectra.
25 Figure 2.7: An example of spectral slope determination for an asteroid. The black line is the normalized reflectance spectrum, and the pink line is the measured spectral slope between 0.95 to 2.30 µm.
The uncertainty in the measurement of S0 is dominated by the systematic uncertainty caused by the solar analog stars. If each solar analog was identical, then dividing them by any other spectrum would result in a continuum of 1. This is not the case, as no two stars are exactly alike spectrally. The deviation from unity is systematic error associated with using different solar analog stars, and dominates all other sources of uncertainty.
26 2.3.4 Curvature
Ziffer et al. (2011) showed that spectral concavity was a distinguishing feature to distinguish be- tween objects in the Themis and Veritas primitive outer belt families. In their study of NIR spectra, the Themis objects exhibit positive curvature (convex) while the Veritas objects have a negative cur- vature (concave). We quantify the curvature of each of our NIR spectra following the procedure in Ziffer et al. (2011). First, we fit a fourth order polynomial to the spectrum, omitting the data points within the telluric absorption bands. We then take the second derivative of the polynomial and take the average of that to obtain a value for curvature, c. We show an example of this process in Fig. 2.8. The units for c are %/µm2. The uncertainty in the quantification of curvature is found by repeating the process on the spectrum plus and minus the systematic error of the slope and taking the difference between the maximum and minimum values.
27 Figure 2.8: An example of spectral curvature determination for an asteroid. Top: A fourth order polynomial (red line) was fit to the spectra in the 0.9–2.4 µm region (green), excluding the telluric bands (in blue). Bottom: The second derivative of the fit (blue) is summed in the region between 1.0 and 2.3 µm (green) to quantify the curvature.
28 CHAPTER 3: NEAR-INFRARED SPECTROSCOPY OF THE KLIO
PRIMITIVE INNER-BELT ASTEROID FAMILY
Anicia Arredondo1, Vania Lorenzi2,3, Noem´ı Pinilla-Alonso4, Humberto Campins1, Andrew Malfavon1, Julia de Leon´ 3,5, David Morate6
1 Physics Department, University of Central Florida, P.O. Box 162385, Orlando, FL 32816, USA 2 Fundacion´ Galileo Galilei, INAF, La Palma, Tenerife, Spain 3 Instituto de Astrof´ısica de Canarias, C/V´ıa Lactea´ s/n, 38205 La Laguna, Tenerife, Spain 4 Florida Space Institute, University of Central Florida, Orlando, FL 32816, USA https://www.overleaf.com/project/5fa95ff3d83b04199c6e6875 5 Departamenta de Astrof´ısica, Uni- versidad de La Laguna, 38205 La Laguna, Tenerife, Spain 6 Observatorio´ Nacional, Coordenac¸ao˜ de Astronomia e Astrof´ısica, Rio de Janeiro 20921, Brazil
This chapter was published in Icarus in Jan. 2020.
29 3.1 Introduction
Primitive asteroids have been minimally altered since their formation and therefore present a glimpse of what the solar nebula looked like at the time of solar system formation. Campins et al. (2018) define primitive asteroids as those with low reflectivity (geometric albedo pV ≤ 0.15) and mostly featureless spectra. The slopes of these objects’ spectra can range from blue (S’<0) to red (S’>0), however the differences between the spectra of C, B, G, and F class objects are not very pronounced (Cellino et al. 2001). Primitive asteroids can be aqueously altered, generating hydrated materials like phyllosilicates, sulfates, oxides, carbonates, and hydroxides. Evidence of this alteration is mostly seen in the visible (with an absorption feature at 0.7 µm ) and 3 µm regions. Understanding the organic compounds and water abundances on asteroids can help explain how the solar system evolved and how life may have formed on Earth.
The inner main belt (IMB), located between the ν6 secular resonance and the 3:1 mean motion resonance with Jupiter (2.15–2.5 AU), is the main source of near-Earth asteroids (NEAs) (Bottke et al. 2015). Asteroid impacts create fragments with similar orbital elements, called families. The Yarkovsky/YORP effect alters the orbits of these collisional fragments, delivering them into the nearby orbital resonances which can then deliver them to NEA space. NEAs and their sources have become of particular interest recently due to the two sample return missions with NEAs as targets: NASA’s OSIRIS-REx and JAXA’s Hayabusa2 are completing missions to asteroids (101955) Bennu and (162173) Ryugu, respectively. Dynamical models and spectral comparisons show that both objects were likely delivered from primitive, low-inclination, inner-asteroid belt populations (Bottke et al. 2015, Campins et al. 2010, 2013); however, we cannot discard high- inclination IMB families as the source of either object due to their close proximity to the ν6 and 3:1 resonances. Because one cannot send a sample return mission to every NEA, we must rely on
30 remote methods such as spectroscopy to determine their composition and provide a framework for interpreting results from these missions by integrating smaller-scale details with the bigger picture of solar system origin and evolution. Knowing the origin of an NEA will also offer more context for hazard mitigation and provide insight into whether an object has valuable resources at or below its surface.
At least eight primitive asteroid families have been identified in the inner belt (Fig. 3.1): Polana- Eulalia (hereafter Polana), Erigone, Clarissa, Chaldaea, Klio, Sulamitis, Svea, and Chimaera (Nesvorny´ et al. 2015, Walsh et al. 2013). The PRIMitive Asteroid Spectroscopic Survey (PRI- MASS) is a spectroscopic survey in both the visible and near-infrared (NIR) that aims to charac- terize asteroids in these families and other primitive populations in the main belt. We combine our NIR and visible spectra with spectroscopy at other wavelengths and with geometric albedos to: [1] determine the presence and distribution of hydrated minerals in the IMB; [2] explore dif- ferences between and within families; [3] predict the properties of primitive NEAs; [4] test for space weathering effects by comparing the spectra of the younger families with the older, more- weathered families; and [5] constrain asteroid compositions by comparing with laboratory spectra of minerals and carbonaceous chondrites.
31 Figure 3.1: Semimajor axis vs. inclination (left) and eccentricity (right) of the primitive inner belt asteroid families. Objects in this study are highlighted with red diamonds.
PRIMASS has observed the Polana-Eulalia family in the visible and NIR (de Leon´ et al. 2016, Pinilla-Alonso et al. 2016) and the other seven families have all been observed in the visible (Morate et al. 2016, 2018, 2019). The results of these surveys show that primitive inner belt asteroid families fall into at least two compositional groups based on the dispersion of spectral slopes and the presence of the 0.7 µm absorption feature: The Polana and Clarissa families are “Polana-like” (anhydrous and spectrally homogeneous) and the Erigone and Sulamitis families are “Erigone-like” (hydrated and spectrally diverse). Morate et al. (2019) characterize Klio in the vis-
32 ible and find that Klio spectra are heterogeneous and 23% of their sample was hydrated, which is a smaller fraction of hydrated objects compared to Erigone (58%). They use a two-sample Kolmogorov-Smirnov test comparing the visible slope distributions of Klio objects with the slope distributions of Polana-like and Erigone-like populations to show that in the visible, the Klio family is not compatible with either group.
The goal of this paper is to use NIR spectroscopy to characterize the Klio family and compare it with the two previously determined compositional groups. Details of the observations and the data reduction process are given in Section 3.2. The analysis of the computed reflectance spectra is shown in Section 3.3. We discuss the results of our survey and the possibility of a “Klio-like” compositional group in Section 3.4. We compare our spectra with the older Polana family as well as NEAs (101955) Bennu and (162173) Ryugu and discuss implications for space weathering in Section 3.4. Conclusions and future work are given in Section 3.5.
3.2 Observations and data reduction
This work focuses on the Klio family as defined by Nesvorny´ et al. (2015) which contains 330 total objects. Our targets were chosen if they were bright enough (V<18) to achieve an acceptable S/N of ∼50 in ∼1 h of observing time in PRISM mode on the IRTF. The same criterion was used to choose objects for the TNG. Because of the low number of objects with a magnitude above our threshold, we observed one object with V>18. Objects were observed at an airmass of X<1.30 whenever possible. Priority was given to objects that had already been observed in the visible (e.g. Morate et al. 2019). We observed 21 of the 330 total Klio objects (Table 3.1), including the two
33 brightest members of the family, (12) Victoria and (84) Klio, the latter being the most likely parent body of the family and the former a potential interloper (see Section 3.3.5). Observational details are shown in Table 3.2 and include asteroid number, date and UT start of observation, telescope used, airmass, phase angle (α) and apparent visual magnitude (V) at the time of observations, total exposure time and solar analog star used.
Table 3.1: Physical and dynamical properties of the observed targets. We include our computed values for the spectral slope (S’) and curvature (c) and their associated errors. See Section 3.3.2 for more details. Also included are the Tholen-Bus-SMASS taxonomies for each object found in the literature. A * indicates that a visible spectrum is published by SMASS (Bus & Binzel 2002) and a † indicates that a visible spectrum is published by Morate et al. (2019).
2 2 ˚ 0 ˚ Number a [AU] e i [deg] H D [km] errD [km] pV errpV Tax T-B-S S’ [%/1000A] errS [%/1000A] c [%µm ] errc[%µm ] (12) Victoria * 2.334 0.220 8.373 7.2 133.32 36.012 0.149 0.068 S - - 2.059 0.012 0.897 0.045 (84) Klio *† 2.362 0.236 9.315 9.3 82.33 23.038 0.037 0.021 G Ch - 1.118 0.011 -0.325 0.008 (3146) Dato † 2.433 0.200 8.366 13.5 7.68 1.707 0.121 0.061 - - C 1.167 0.019 -0.406 0.021 (3627) Sayers * 2.349 0.147 9.712 13.7 10.137 3.959 0.054 0.043 - B - 0.371 0.007 0.641 0.011 (5327) 1989 EX1 † 2.342 0.162 9.755 13.5 11.131 3.275 0.057 0.03 - - C 1.489 0.013 -0.337 0.044 (6857) Castelli † 2.313 0.196 9.200 14.0 8.873 2.243 0.056 0.069 - - C 2.023 0.014 0.294 0.028 (7274) Washioyama † 2.326 0.141 10.743 14.0 8.951 2.859 0.055 0.064 - - C 0.202 0.096 1.277 0.037 (8091) 1992 BG † 2.330 0.152 9.900 13.8 10.268 3.22 0.046 0.047 - - C 1.143 0.009 -0.027 0.284 (15561) 2000 GU36 † 2.399 0.244 10.469 14.2 6.863 1.555 0.047 0.044 - - C 1.272 0.01 0.05 0.005 (20080) Maeharatorakichi 2.340 0.143 10.144 14.3 7.983 2.008 0.058 0.035 - - - 0.889 0.01 0.808 0.014 (20168) 1996 VY4 2.403 0.174 8.723 15.0 4.969 1.091 0.072 0.036 - - - 0.96 0.008 0.611 0.067 (20432) 1999 BD12 † 2.342 0.227 11.333 14.0 6.166 1.508 0.117 0.09 - - C 0.469 0.029 −3.742 0.008 (20604) Vrishikpatil 2.366 0.162 8.482 14.3 5.172 1.59 0.117 0.097 - - - 0.427 0.006 0.228 0 (37917) 1998 FJ103 2.442 0.204 7.905 14.5 4.823 1.104 0.12 0.088 - - C 0.484 0.006 1.031 0.006 (42347) 2002 AV155 2.328 0.222 10.556 15.1 6.584 2.165 0.037 0.045 - - CX 1.357 0.014 0.374 0.025 (42802) 1999 GE15 2.423 0.164 9.099 14.6 6.568 2.273 0.059 0.053 - - - 2.234 0.031 2.312 0.194 (43346) 2000 RT103 † 2.391 0.191 8.510 14.3 8.456 2.426 0.047 0.033 - - - 0.903 0.01 0.774 0.012 (65540) 7628 P-L 2.398 0.201 10.708 14.5 6.22 1.114 0.066 0.041 - - - 1.578 0.048 0.342 0.107 (70158) 1999 NZ37 2.336 0.254 9.867 15.2 5.415 2.144 0.055 0.034 - - - 1.822 0.013 0.786 0.04 (70394) 1999 RP237 2.440 0.224 8.933 14.9 5.764 1.604 0.058 0.058 - - - 0.891 0.005 0.774 0.006 (106797) 2000 XX27 2.317 0.141 10.006 15.7 3.744 0.936 0.06 0.055 - - - 0.927 0.009 0.21 0.019
34 Table 3.2: Observational details. A † indicates that the number corresponds to star in Table 3.3
Number Date UT Start UT Telescope Airmass Alpha [°] V Texp [s] SA† 12 20170203 5:09 TNG 1.56 23.2 11.8 64 4,6 84 20170203 6:50 TNG 1.73 21.5 14.6 640 4,6 3146 20170202 2:10 TNG 1.15 10.8 17.3 3060 4,5 3627 20170823 9:48 IRTF 1.59 10.4 17.0 1920 1 5327 20180628 9:40 TNG 1.14 27.8 16.8 2160 6 6857 20180622 11:55 IRTF 1.36 21.7 16.7 1440 2,3 7274 20170202 9:49 TNG 1.02 17.7 17.7 1800 4,5 8091 20170130 11:56 TNG 1.14 3.0 16.2 3600 1,2,6 15561 20190330 14:41 IRTF 1.20 14.7 18.4 2160 5,6 20080 20190330 6:48 IRTF 1.11 16.9 16.7 1440 5,6 20168 20181004 5:56 IRTF 1.50 7.3 17.7 2400 2 20432 20180622 13:20 IRTF 1.37 30.1 16.9 1680 2,3 20604 20170823 12:33 IRTF 1.09 9.1 17.7 1200 1 37917 20180901 14:15 IRTF 1.22 13.7 17.4 2400 1,2 42347 20170823 15:54 IRTF 1.84 9.4 17.2 1920 1 42802 20180628 12:28 TNG 1.56 18.6 17.6 2160 1,2,6 43346 20180622 10:48 IRTF 1.36 10.9 16.3 1440 2,3 65540 20170202 11:01 TNG 1.12 26.4 17.7 3060 4,5 70158 20170925 6:39 IRTF 1.18 10.4 16.5 2640 1 70394 20180901 11:48 IRTF 1.16 10.1 17.4 2400 1,2 106797 20190330 7:50 IRTF 1.08 11.1 17.8 2400 5,6
35 Table 3.3: Coordinates and magnitudes of solar analog stars used.
ID Star RA Dec V
1 SA 115-271 23 42 41.8 +00 45 14 9.7 2 SA 112-133 20 43 11.8 +00 26 15 10.0 3 SA 107-684 15 37 18.1 -00 09 50 8.4 4 SA 98-978 06 51 34.0 -00 11 28 10.5 5 SA 102-1081 10 57 04.0 +00 13 10 9.9 6 SA 107-998 15 38 16.0 +00 15 24 10.5
3.2.1 IRTF
We observed 13 asteroids over 6 nights between August 2017 and March 2019 using the NASA Infrared Telescope Facility (IRTF) at the Mauna Kea Observatory (Hawaii). We used the SpeX spectrograph (Rayner et al. 2003) in the low-resolution PRISM mode with a wavelength coverage of 0.7–2.52 µm. The 0.8×1500slit was oriented at the parallactic angle for all nights. The asteroids were observed using a beam switching pattern of AB separated by 7.500. All objects were observed near the meridian if possible.
To correct for telluric absorptions and to remove the solar spectrum from the asteroid spectra, we observed at least one local standard star (with a spectral type G2-5 V and B–V and V–K colors similar to those of the Sun) for each object at the same airmass that we observed the asteroid. We also observed at least one well-studied solar analog star at several different airmasses throughout the night (Table 3.3). We acquired flat field frames and argon arc calibration frames at least once
36 per night. The program SpeXtool (Cushing et al. 2004) was used to apply flat fielding, combine, and extract spectra of each AB pair of asteroid, local standard star, and solar analog. The telluric contribution of each spectrum was modeled and removed using the IRTF supplied program xtellcor (Vacca et al. 2003). The final spectra are shown in Fig. 3.2. All spectra are normalized to unity at 1.0 µm.
37 Figure 3.2: Spectra of the 21 objects in our sample. A blue number indicates the object was observed by TNG and green indicates IRTF. The spectral slope S’ between 0.95 and 2.3 µm is shown by a pink line. All spectra have been normalized to unity at 1 µm.
38 3.2.2 TNG
We observed 8 asteroids over 4 nights between January 2017 and June 2018 using the Telescopio Nazionale Galileo (TNG), a 3.58 m telescope sited on the El Roque de Los Muchachos Obser- vatory, La Palma Island, Canary Islands (Spain). The observations were performed with the near infrared instrument NICS, using the AMICI prism and the 1.500slit. This setup provides spectra in the range 0.8–2.4 µm, with a resolving power of 35 quasi-constant along the spetrum (Oliva 2000). We observed at least two solar analog stars each night in order to correct for telluric absorptions and to obtain relative reflectance. We observed each object in four different positions along the slit, following an ABBA scheme (positions A and B separated by 1000) in order to subtract the sky contribution. The ABBA scheme was repeated several times to increase the signal to noise. The exposure time and the number of ABBA repetitions was chosen depending on the sky conditions and the magnitude of the object. All observations were performed in parallactic angle and, for asteroids, the tracking of telescope was set to the asteroid proper motion.
Flat fielding, removal of sky and background, and extraction of a 1D spectrum were done using standard IRAF packages. Emission lines from Argon and Xenon lamps are not resolvable using the AMICI prism’s low resolution, and so a look-up table based on theoretical dispersion predicted by ray tracing was used for wavelength calibration. To calibrate in wavelengths and to divide the asteroid spectrum by the solar analog spectrum, we used a Python script developed by this group. Details of the procedure followed are given in Pinilla-Alonso et al. (2016). The final reflectance spectrum of each object, normalized to unity at 1.0 µm, is shown in Fig. 3.2.
39 3.3 Results and analysis
3.3.1 Characteristics of the sample
Our final sample contains 21 different objects from the Klio family. Table 3.1 shows the physical and dynamical properties of each target. Orbital elements (a, e, i) and absolute magnitude (H) have been extracted from the Minor Planet Center Orbit Database file (MPCORB.DAT4). Taxonomical classification (Tholen, Bus, and SDSS) has been taken from the Small Bodies Node of the NASA
1 Planetary Data System . Diameter (D), visible geometric albedo (pV ) and their respective associ- ated errors are from NEOWISE (Mainzer et al. 2019). The spectral slope (S’) of each object was calculated using the method described in Morate et al. (2018) in the range from 0.95 to 2.3 µm. The curvature (c) of each spectrum was calculated as using the method in Ziffer et al. (2011) in the region from 0.9 to 2.4 µm.
Of the 330 total family members, 74 have SMASS taxonomies determined. Of those 74, 86% are C-type, 10% X-type, and the rest are non-primitive types (S and Q); however, in the sample of Morate et al. (2019), they find that the Klio family is 56.7% C-types, 26.7% X-types, 13.3% B-types, and 3.3% T/D-types. All 330 objects have geometric albedos from NEOWISE (Mainzer et al. 2019), and 92% have pV <0.1. The mean albedo of our sample is 0.071±0.013 which is similar to the mean albedo for the entire family (0.065±0.002). Objects in the sample range from 4 to 133 km with an average of 17 km.
1http://pdssbn.astro.umd.edu/
40 Asteroids (12) Victoria and (84) Klio have published visible and NIR spectra in the literature (Bus & Binzel 2002, Chapman et al. 1993, Fornasier et al. 2014, Lazzaro et al. 2004, Reddy & Sanchez 2016, Xu et al. 1995), and (3627) Sayers has one published visible spectrum (Bus & Binzel 2002). Both (84) Klio and (3627) Sayers show the 0.7 µm feature associated with hydration, but (12) Victoria does not. 9 of our objects overlap with those observed by Morate et al. (2019) in the visible (denoted by † in Table 3.1).
3.3.2 Spectral slopes and curvature
Because primitive asteroid spectra are featureless in the NIR, we calculate the spectral slope (S0) of each asteroid in the range from 0.95 to 2.3 µm following the procedure outlined in Morate et al. (2016) using the expression dS/dλ S0 = S1.0
0 where dS/dλ is the rate of change of the reflectance and S1.0 is the reflectance at 1.0 µm. S is measured in units of %/1000A˚ . The spectral slope can be indicative of composition and degree of alteration (e.g. de Leon´ et al. 2012). To estimate the uncertainty in calculated slopes, we follow the procedure in Pinilla-Alonso et al. (2016) and compare the spectra of solar analog stars each night. Each solar analog spectrum was divided by a reference solar analog spectrum. The result of the division is 1 if the solar analogs are identical. The range of deviation from unity is taken to be the systematic error of computed spectral slopes which dominates all other sources of uncertainty. The error in the slope for each object is given for each night in Table 3.1. In Fig. 3.2 we show each individual spectrum in our sample with the linear fit for the spectral slope S0 plotted over it (pink line). The S0 fit is always linear, even if the spectrum has curvature.
41 The concavity (c<0) or convexity (c>0) of the spectrum can be used to describe differences be- tween families. We quantify the curvature, c, of each spectrum by averaging the second derivative of a polynomial fit (e.g. Ziffer et al. 2011). We fit a fourth order polynomial to the region between 0.9 and 2.4 µm. Because the signal in the telluric bands (1.3–1.45 µm and 1.8–1.95 µm) is essen- tially zero, these points are omitted. The average concavity is calculated by summing the second derivative of the polynomial fit at each point between 1.0 and 2.3 µm and dividing by the number of points. We exclude 0.1 µm from each end from the summation to avoid introducing error caused by the divergence of the fit if we had used a higher order. Fig. 3.3 shows this process for a concave and a convex object. The error in the curvature measurement is found by doing the same fitting process on the spectrum plus the associated error and again on the data minus the associated error (e.g., red dashed lines in Fig. 3.4). The total error in the curvature is the difference between the average c value of the spectrum and the average c value of the spectrum with errors.
42 Figure 3.3: Concavity determination of a concave (left) and convex (right) spectrum. Top: A fourth order polynomial (red line) was fit to the spectra in the 0.9–2.4 µm region (green), excluding the telluric bands (in blue). Bottom: The second derivative of the fit (blue) is summed in the region between 1.0 and 2.3 µm (green) to quantify the concavity.
3.3.3 Spectral homogeneity of the Klio sample
Fig. 3.4 shows the spectrum for each asteroid in our sample compared with the average spectra of the Klio and Polana families. The spectral slopes of the sample range between flat (0.202±0.440%/1000A˚ ) to red (2.234±0.540%/1000A˚ ) and the slope of the mean spectrum is 1.052±0.425%/1000A˚ . Our sample is spectrally homogeneous: the distribution of spectral slopes is shown in Fig. 3.5. The average Klio spectrum is convex with a c value of 0.511±0.113%/µm2and there are 4 concave, 2 flat, and 15 convex objects (Fig. 3.6). We did not find any correlation between slope and curva- ture. We also did not find any significant trends between either slope or curvature with physical attributes, observational conditions, or orbital parameters.
43 Figure 3.4: Each individual spectrum from our sample (grey) compared with the average spectra of the Klio (red) and Polana (purple) families. The red and purple dashed lines are ±1σ of the mean of the average Klio and Polana spectra. Both average spectra fit within 1σ of the other. The Klio family spectra are mostly convex, red and homogeneous. All spectra have been normalized to unity at 1 µm. Polana data are from Pinilla-Alonso et al. (2016).
44 Figure 3.5: Comparison of the NIR spectral slope distributions of the Klio and Polana primitive inner belt families. The range of slopes is similar for both families regardless of different orbital elements (Fig. 3.1). The bin size is 0.33%/1000A˚ . Polana data are from Pinilla-Alonso et al. (2016).
45 Figure 3.6: Distribution of the curvature of Klio objects. There is a mixed distribution between concave (c<0) and convex (c>0) shapes. The average curvature of 0.511%/µm2 is plotted as a pink dashed line. The bin size is 0.5%/µm2.
46 3.3.4 Comparison with the Polana family
Pinilla-Alonso et al. (2016) observed the Polana family in the NIR. The Klio and Polana fami- lies are very distinct from each other in orbital space (Fig. 3.1) and differ in the visible in both spectral homogeneity and hydration. In the NIR, the differences between the two families are more subtle. Fig. 3.4 shows that the average spectra of the Klio and Polana families lie within 1σ of each other. Both families have similar average slopes, 1.052±0.425%/1000A˚ for Klio and 0.68±0.68%/1000A˚ for Polana, i.e., indistinguishable within the uncertainties (Fig. 3.5). A qual- itative comparison of the average NIR spectra of these two families (Fig. 3.4) shows the Polana spectrum is more convex in the region between 0.8 and 1.5 µm.
While the average spectra of the families are similar, the slope distributions are not. The range of NIR slopes is small for both families (2.032%/1000A˚ for the Klio family and 1.58%/1000A˚ for the Polana family). We run a two-sample Kolmogorov-Smirnov test comparing the NIR slope distri- butions shown in Fig. 3.5 to quantify the differences between the two families. The Kolmogorov-
Smirnov statistic, Dm,n, is given by the equation
Dm,n = sup |F (x)1,m − F (x)2,n| x
where F (x)1,m and F (x)2,n are the cumulative distribution functions of the two compared distri- butions. A critical value, Dcrit,m,n is defined as
rm + n D = c(α) crit,m,n mn
47 where c(α) is 1.36 for α=0.05 and m and n are the sizes of the two samples. If Dm,n>Dcrit,m,n, then both samples originated in different distributions. We note that the size of the Klio family might not be large enough for these tests to be conclusive, however the test can be informative and a guide for future studies of this family. For the Klio and Polana families, D is 0.402 and Dcrit is 0.364, indicating that the two populations are distinct. It is interesting that the average spectra for the families are so similar while the individual spectra are not.
3.3.5 Interloper (12) Victoria
Fig. 3.7 shows the Klio family in (a, H) space along with the Yarkovsky cone which represent the farthest distance a family member is likely to drift from the center of the family due only to the Yarkovsky effect. Six of the objects in our study fall outside the Yarkovsky cone for the Klio family; however, the albedos and spectra of five of those objects are consistent with other Klio family members. These five objects are also not very far from the Yarkovsky cone given their sizes and the likely ejection velocities that would have put them outside this cone. Asteroid (12) Victoria is a different case, it was defined as a Klio member by Nesvorny´ et al. (2015) based on its orbital elements; however, there are many arguments against Victoria’s membership in the Klio family. First, Victoria does not lie within the Yarkovsky cone of the family and given its size it’s hard to explain this with a large ejection velocity during the family forming collision. In addition, Victoria is not a primitive asteroid, its geometric albedo is 0.149 and was classified by Tholen (1984) as an S-type asteroid, which is consistent with Victoria’s spectra in the visible (Bus & Binzel 2002, Chapman et al. 1993, Lazzaro et al. 2004) and NIR (Bell et al. 1984, Reddy & Sanchez 2016). Our own spectrum (Fig. 3.2) looks similar to those published by Bell et al. (1984) and Reddy & Sanchez (2016); the shape and depth of absorptions near 1 and 2 µm confirm Victoria’s S-type classification. For these reasons, we do not use asteroid Victoria in our analysis.
48 Figure 3.7: Absolute magnitude (H) of the asteroids in the Klio family (red dots) as a function of their semimajor axis. The solid red lines represent the Yarkovsky cone for the family. Dashed lines mark the ν6 and 3:1 resonances. The objects in this study are marked with larger red circles, and the red diamond is the parent body (84) Klio.
49 3.3.6 Comparison with (101955) Bennu and (162173) Ryugu
NASA’s OSIRIS-REx and JAXA’s Hayabusa2 are currently completing missions to asteroids (101955) Bennu and (162173) Ryugu, respectively. Both NEAs are primitive: Bennu is a B-type and Ryugu is a C-type. Dynamical models and spectral comparisons show that both objects were likely deliv- ered from primitive, low-inclination, inner-asteroid belt populations (Bottke et al. 2015, Campins et al. 2010, 2013); however, due to their proximity to the ν6 and 3:1 resonances and high percentage of B and C-type objects, we cannot discard high-inclination IMB families as the source of either object (i.e., Klio).
We compare the NIR spectra of Bennu and Ryugu with the spectra of our sample following the same comparative approach as in Pinilla-Alonso et al. (2016) to show that both objects’ spectra could be compatible with an origin in the Klio family (Fig. 3.8). Binzel et al. (2015) present multi- ple spectra of Bennu obtained between 1990 and 2013 which show a range of spectral slopes. We plot the reddest and bluest spectra of Bennu in blue and cyan, respectively, as well as the mean slope of the Klio family (red). While the Binzel et al. spectrum fits within 1σ of the mean Klio spectrum, the Clark et al. spectrum looks visibly different than the Klio average. The spectral vari- ation of Bennu from ground-based spectra is not well understood, but in-situ measurements from OSIRIS-REx will hopefully help solve this puzzle. The most recent spectrum from the OSIRIS- REx spacecraft (Hamilton et al. 2019) shows that Bennu has a blue spectrum similar to the spec- trum from Clark et al.. While we do not favor the Klio family for the origin of Bennu because of the mismatch between the Klio spectrum and that of Clark et al. and Hamilton et al., we cannot rule it out based on our spectra alone. The spectrum of Ryugu from Abe et al. (2008) fits well within the spectra of the Klio family. The most recent spectrum of Ryugu from Kitazato et al. (2019) shows a similar red slope. Based off of spectral comparisons, we conclude that the Klio family cannot be discounted as a possible origin for Ryugu or Bennu.
50 Figure 3.8: Comparison of the spectra of the Klio family from this work to spectra of (101955) Bennu and (162173) Ryugu. The spectra from Binzel et al. (2015) and Abe et al. (2008) are compatible with the spectra in the Klio family. The Bennu spectrum from Clark et al. (2011) is bluer than the average Klio spectrum.
51 3.4 Discussion
3.4.1 The possibility of a third compositional group
The Klio family is spectrally homogeneous in the NIR; however, this spectral homogeneity does not extend to the visible (Morate et al. 2019). The parent body (84) Klio shows evidence of hydrated minerals (Fornasier et al. 2014) and so do at least some of the other family members (23% of the sample in Morate et al. 2019), making the Klio family similar to the Erigone family in the visible, though with a smaller percentage of hydrated objects. Morate et al. (2019) perform a two sample Kolmogorov-Smirnov test comparing the visible slope distributions of Klio objects with the slope distributions of the Polana-like and Erigone-like populations to show that the Klio family is not compatible with either group.
The only families we have characterized in the NIR are Polana and now Klio, so our results on the NIR behavior of inner-belt families are limited. The similarity between the average spectra of the Klio and Polana families in the NIR is consistent with two possibilities: 1) There is a third com- positional group that is “Klio-like”, i.e., somewhat hydrated and spectrally diverse in the visible but homogeneous in the NIR, or 2) NIR spectra of inner-belt families cannot be used to distin- guish between primitive families. We have observed and are beginning to analyze NIR spectra of the Erigone, Sulamitis, and Chaldaea families in order to determine if this second possibility is correct.
We note that NIR spectroscopy of other primitive families has proven to be a diagnostic tool show- ing significant differences between and even within families. Fornasier et al. (2016) compared the spectra of the Themis family and the Beagle subfamily to show that the Themis objects showed a larger range of spectral slopes than the younger Beagle subfamily, and that the Beagle objects look spectrally bluer on average than the Themis ones. Ziffer et al. (2011) show that there are
52 striking differences between the spectra of the Themis and Veritas families: Themis objects are convex and red sloped while Veritas objects are concave and blue sloped. de Leon´ et al. (2010) show homogeneity within objects in the Pallas family, and that the average Pallas spectrum has a very blue slope. In these three cases, the spectra of each family are homogeneous in the NIR, as the Klio objects are; however, it is clearly possible to detect differences between families using NIR slope, curvature, and slope distribution. Therefore, the result that the Klio and Polana fami- lies look so similar in the NIR could be suggestive of a similar origin; for example, the primitive IMB population may have originated from just a few primordial parent bodies. This observational result supports the conclusions of Dermott et al. (2018) that asteroids in the IMB originate from the disruption of only a few large bodies.
In the visible, Polana and Clarissa are “Polana-like” and Sulamitis and Erigone are “Erigone-like”, suggesting that composition is similar at similar inclinations (Fig. 3.1). If there is indeed a third “Klio-like” composition, this trend with inclination will be strengthened. This result is being used to constrain models of giant planet formation (e.g. Lowry 2018). It is logical to suspect that the Chaldaea family, with an inclination close to Klio, is also “Klio-like”. Morate et al. (2019) report these two families present a similar slope distribution with similar mean values in the visible; however, they also report a difference in the percentage of hydrated objects in the Chaldaea family (79%) compared to the Klio family (23%). Morate et al. (2019) suggest that the difference in the hydration fraction of the Klio and Chaldaea families could be due to a common parent body with varying degrees of hydration with depth.
53 3.4.2 Spectral differences with Polana: consistent with space weathering?
Carruba & Nesvorny´ (2016) estimate the age of the Klio family to be 960±250 Myr. The Polana family is 1400±150 Myr old (Bottke et al. 2015), which is marginally older than the Klio family given the published uncertainties. Lantz et al. (2017) show that the effects of space weathering can be seen after timescales as short as 103 years, so the Polana family should show more signs of space weathering than Klio. Next, we consider if space weathering could explain small differences between these families. Our comparison assumes these two families started out with similar com- positions, although this is not clear, their NIR spectral similarities suggest that this assumption is not completely unjustified.
The effects of space weathering on low albedo objects are not well known. For example, Lantz et al. (2017) show that ion irradiating low albedo meteorites can cause their spectra to become bluer, brighter, and more convex in the visible to near-infrared range. However, the laboratory results of Thompson et al. (2019) indicate that space weathering effects on primitive meteorites can be more complex than expected. The calculated S0 for the mean Polana family spectrum is 0.68±0.68%/1000A˚ , which is marginally bluer than the S0 for the Klio family (1.052±0.425%/1000A˚ ). We have determined the albedo of the Polana family from a subset of objects that are not contam- inated by the high albedo Nysa family. The average albedo of these Polana family objects with albedos determined from NEOWISE is 0.084±0.003, which is brighter than the value for 330 Klio objects (0.065±0.002). Polana appears to be more convex than the Klio family as shown in the comparison of both their average spectra. The comparison of these two families with marginally different ages is consistent (Polana is older, bluer, brighter, and more convex than Klio) with the results predicted by Lantz et al. (2017).
54 This result should be regarded with healthy suspicion for several reasons, including that the initial compositions of the two populations might not be the same (i.e., many Klio objects show the 0.7 µm absorption but Polana objects do not). However; just because the Polana family member spectra do not show absorption features at 0.7 µm does not necessarily mean that they are not hydrated. It has been suggested that hydrated minerals could be present, and that the paucity of hydration features may be due to attenuation of the bands by space weathering (Campins et al. 2018). This is the case for asteroid (101955) Bennu, which has not shown any absorption at 0.7 µm (so far), however there is a strong absorption at 2.7 µm that is due primarily to hydrous clay minerals (Hamilton et al. 2019). Both Bennu and Ryugu most likely originated in the Polana family (Bottke et al. 2015, Campins et al. 2010, 2013) so the evidence of hydration despite the lack of 0.7 µm absorption on the spacecraft targets could mean that the Polana family is also hydrated. We searched the literature but did not find any published spectra of asteroid (142) Polana in the 3 µm region where the 2.7 µm band may be present. Keeping this in mind, we believe that the agreement between the some of the predictions of Lantz et al. (2017) and the differences between the Klio and Polana spectra are consistent with greater space weathering of Polana asteroids. As mentioned, this is a tenuous result but one that can be pursued further with additional in-situ observations of Bennu and Ryugu, as well as ground-based studies.
3.5 Conclusion
In the NIR, the Klio family is spectrally homogeneous, red sloped, and objects show a range of positive and negative curvature. In comparison with the Polana family, the Klio family has a similar mean slope and average spectrum, though Polana is bluer and slightly more convex. This similarity could be indicative of two possibilities: 1) There is a third compositional group of asteroids in the inner main belt that are somewhat hydrated, spectrally heterogeneous in the visible, and spectrally
55 homogeneous in the NIR, or 2) the spectra of inner-belt primitive families in the NIR region are not very diverse (as we discussed in Section 3.4.1, this would be in contrast with outer-belt primitive families, e.g., Campins et al. 2018). In principle this is consistent with the findings of Dermott et al. (2018) that IMB families originate from a small number of primordial planetesimals, however we do point out that the Klio and Polana families look significantly different in the visible. These possibilities will be explored more by doing this same characterization on other primitive families in the IMB (i.e., Erigone, Sulamitis, Chaldaea, Svea, Chimaera). We have already observed a sufficient fraction of the Erigone family in the NIR and have a partial sample of Sulamitis and Chaldaea family members. The Svea and Chimaera families are our next priority for spectroscopy in the NIR.
If there is indeed a “Klio-like” compositional group, the fact that the three groups occur at three different inclinations may be used as a constraint on the early history of the inner belt and the migration of the giant planets. As mentioned, the Polana and Klio families are not very different in the NIR; however, NIR studies of the Themis, Beagle, Pallas, and Veritas families show consid- erable NIR differences. The Klio and Polana families are similar but not identical, and the small differences in slope and geometric albedo for the Klio and the older Polana families appear to be consistent with space weathering effects described by Lantz et al. (2017) and suggested by the vis- ible differences between Clarissa and Polana (Morate et al. 2018, Campins et al. 2018). Spectral comparisons show that the Klio family cannot be ruled out as a source for both asteroids (101955) Bennu and (162173) Ryugu. Future work to support these claims includes:
• Complete our NIR survey of the Erigone and Sulamitis families to test if inner-belt primitive families are homogeneous in the NIR, regardless of diversity seen in the visible.
• A NIR survey of the Chaldaea family to compare with the similar inclination Klio family and test the apparent correlation between composition and inclination.
56 • A survey of primitive background objects that don’t belong to any of these families to test the correlation between composition and inclination and to determine if IMB families are the main source of the background, or if the background has different spectral characteristics and possibly a different origin. A group of inner belt primordial background objects has been identified by Delbo´ et al. (2017), and a compositional survey of Delbo objects in the NIR is being worked on currently by us.
57 CHAPTER 4: NEAR-INFRARED SPECTROSCOPY OF THE
CHALDAEA ASTEROID FAMILY: POSSIBLE LINK TO THE KLIO
FAMILY
Anicia Arredondo1, Humberto Campins1, Noem´ı Pinilla-Alonso2, Julia de Leon´ 3,4, Vania Lorenzi5,3, David Morate6
1 Physics Department, University of Central Florida, P.O. Box 162385, Orlando, FL 32816, USA 2 Florida Space Institute, University of Central Florida, Orlando, FL 32816, USA 3 Instituto de Astrof´ısica de Canarias, C/V´ıa Lactea´ s/n, 38205 La Laguna, Tenerife, Spain 4 Departamenta de Astrof´ısica, Universidad de La Laguna, 38205 La Laguna, Tenerife, Spain 5 Fundacion´ Galileo Galilei, INAF, La Palma, Tenerife, Spain 6 Observatorio´ Nacional, Coordenac¸ao˜ de Astronomia e Astrof´ısica, Rio de Janeiro 20921, Brazil
This chapter was published in Icarus in Jan. 2021.
58 4.1 Introduction
A primitive asteroid is one that has been minimally processed and therefore maintains a composi- tion similar to that with which it formed; hence the study of primitive asteroids can tell us about the composition of the early Solar System. Campins et al. (2018) define primitive asteroids as those with low reflectivity (geometric albedo pV ≤0.15) and mostly featureless spectra. Signatures of aqueous alteration can be seen in the reflectance spectra of primitive asteroids through absorption features near 0.7 and 3 µm.
Evidence suggests that asteroids initially formed big through a fast accretion mechanism (Der- mott et al. 2018, Delbo´ et al. 2019, Morbidelli et al. 2009, e.g.). When these primordial asteroids suffered catastrophic collisions, they broke up into smaller asteroids that share similar orbital ele- ments, called asteroid families. Depending on the composition of the parent body (i.e., differenti- ated or not), the resulting family members may share similar compositions or may provide us with a sample of the interior of the parent body. Over time, the orbits of family members will be altered by effects such as the Yarkovsky/YORP effects, the hypothesized giant planet migration, and sub- sequent, sometimes catastrophic, collisions. Evidence of further catastrophic collisions can be seen throughout the current asteroid belt. As an example, within the Themis family there is a subfamily called Beagle that is much younger than Themis (Nesvorny´ et al. 2008). Therefore, the current identified families may be generations removed from the original primordial planetesimals, and it is possible that multiple families could have originated in a single primordial progenitor body.
The PRIMitive Asteroid Spectroscopic Survey (PRIMASS) is a spectroscopic survey in both the visible and near infrared (NIR) that aims to characterize primitive asteroids in the main belt. By combining our spectra with geometric albedos and laboratory spectra of carbonaceous chondrites we explore the differences between and within asteroid families, test for space weathering effects through comparison of older and younger families, and obtain constraints on asteroid composition.
59 Most of our work has focused on primitive asteroids located in the inner main belt (IMB) between 2.15 and 2.5 AU. At least eight primitive asteroid families have been identified in the IMB (Fig. 4.1): Polana-Eulalia (hereafter Polana), Clarissa, Sulamitis, Erigone, Klio, Chaldaea, Chimaera, and Svea (Nesvorny´ et al. 2015, Walsh et al. 2013). PRIMASS has observed and characterized the Polana and Klio families in the visible and NIR (Arredondo et al. 2020, de Leon´ et al. 2016, Morate et al. 2019, Pinilla-Alonso et al. 2016) and the other six families in the visible only (Morate et al. 2016, 2018, 2019).
Figure 4.1: Semimajor axis vs. inclination (left) and eccentricity (right) of the primitive inner belt asteroid families defined by Nesvorny´ et al. (2015). Objects in this study are highlighted with red diamonds.
60 For primitive asteroids, it was expected that family members should look spectrally similar; how- ever, we have found that this is not always the case. A summary of previous PRIMASS results is given in Table 4.1. We find that primitive IMB asteroid families differ based on the dispersion of spectral slopes (homogeneous vs diverse) and the percentage of objects that show the 0.7 µm absorption feature associated with hydration. This work focuses on the Chaldaea family, which in the visible has homogeneous spectral slopes and 79% of the sample from Morate et al. (2019) showed evidence of hydration at 0.7 µm. The Chaldaea family has 132 identified members, 28 of which have SMASS taxonomies (96% C-type and 4% X-type). The mean albedo for the entire family is 0.07±0.03 and 85% of family members have pV <0.1(Mainzeret al. 2019).
Table 4.1: Summary of PRIMASS findings in the IMB.
Family Inc [°] Number of objects 0.7 µm feature Visible Spectra NIR Spectra
Polana 3 1924 Absent Homogeneous Homogeneous Clarissa 3 179 Absent Homogeneous Erigone 5 1776 Present (58%) Diverse Sulamitis 5 303 Present (60%) Diverse Klio 9 330 Present (23%) Diverse Homogeneous Chaldaea 12 132 Present (79%) Homogeneous Homogeneous (This work) Chimaera 14 108 Present (20%) Diverse Svea 16 48 Absent Homogeneous
Morate et al. (2019) present an interesting hypothesis about the relationship between the Chaldaea family and the nearby Klio family. They posit that because of their similar location and comple- mentary compositions (Chaldaea is hydrated but homogeneous, Klio is diverse but significantly less hydrated), Chaldaea and Klio may have formed from one large primordial object, with the more hydrated Chaldaea objects near to the surface and the less hydrated Klio objects in the inner
61 layers. It is also possible that the contrary is true, that Klio are the objects from the outside layers and the more hydrated Chaldaea objects are the inside layers. Young et al. (1999) use oxygen- isotope compositional data to infer that CI, CM, and CV chondrites could have been derived from different zones in a single body, with the more altered CI objects nearer to the center. In this case, the Chaldaea object would be the CI analog that would have made up the inner layers of the body, while the less hydrated Klio objects would be located closer to the surface.
Morate et al. (2019) present a Yarkovsky cone for the combined Chaldaea-Klio family, which shows the maximum distance that family members will drift over the age of the family due to the Yarkovsky effect (Fig. 4.2). This cone enveloped most objects from both families including the two parent bodies of each family. We note that there is not a clear explanation for why the Chaldaea objects seem to cluster so tightly to one side of the parent body while the Klio objects are spread evenly on both sides. It could be that they are majority prograde rotators and therefore
only drift outward from the Sun or that the higher inclination Chaldaea objects are closer to the ν6 resonance and therefore are depleted at lower semimajor axes. The common origin hypothesis of Morate et al. (2019) is supported by the small difference in mean visual slope of the two families (Chaldaea mean: 0.88±1.25%/1000A˚ , Klio mean: 1.16±2.04%/1000A˚ ). Because the outer layer objects would be more exposed to space weathering, they may be more altered in slope.
62 Figure 4.2: Adapted from Morate et al. (2019). Absolute magnitude of Chaldaea (green) and Klio (red) family asteroids as a function of their proper semi-major axes. The two large markers represent the parent bodies of each family. The black solid line represents the probable Yarkovsky envelope for the combination of both families.
The goal of this paper is to use our NIR spectroscopy to characterize the Chaldaea family and to test the link to the Klio family proposed by Morate et al. (2019). Details of the observations and the data reduction process are given in Section 4.2. The analysis of the computed reflectance spectra is shown in Section 4.3. We compare our spectra with the Klio family and discuss implications in Section 4.4. Conclusions and future work are given in Section 4.5.
63 4.2 Observations and data reduction
The Chaldaea family has 132 objects identified by Nesvorny´ et al. (2015). We did not impose any observing constraints (i.e., low magnitude or low albedo) on our targets because this is a small and faint family. We did, however, prioritize observing objects with visible spectra published in Morate et al. (2019). Our sample includes 15 Chaldaea objects (11% of the total members) including asteroid (313) Chaldaea, the likely parent body of the family. Two of the objects were observed twice and one of the objects was identified as an interloper and is not used in the analysis, bringing our total spectra to 16. We used essentially the same observing approach for the two different telescopes we observed with. All observations were made as close to the meridian as possible, ideally with an airmass of X<1.30, and the slits were aligned with the parallactic angle. The total exposure time for each object was determined depending on sky conditions and magnitude of the object, and we used an AB beam switching pattern to subtract the sky contribution.
The physical and dynamical properties of each target are given in Table 4.2. Orbital elements (a, e, i) and absolute magnitude (H) have been extracted from the Minor Planet Center Orbit Database file
(MPCORB.DAT4). Diameter (D) and visible geometric albedo (pV ) are from NEOWISE (Mainzer et al. 2019). The spectral slope (S0) of each object was calculated using the method described in Pinilla-Alonso et al. (2016) in the region from 0.95–2.3 µm. The curvature (c) of each spectrum was calculated as using the method in Ziffer et al. (2011) in the region from 0.9–2.3 µm. Observational details are shown in Table 4.3 and include asteroid number, date and UT start of observation, telescope used, airmass, phase angle (α), apparent visual magnitude (V) at the time of observations, total exposure time and solar analog star used.
64 Table 4.2: Physical and dynamical properties of the observed targets. We include our computed values for the spectral slope (S0) and curvature (c) and their associated errors. See Section 4.3.2 for more details. a Indicates that a visible spectrum is published by Morate et al. (2019).
0 Number Date a [AU] e i [°] HV D [km] errD [km] pV errpV Slope [%/1000A˚ ] errS Curvature errC (313) Chaldaea a 20170130 2.376 0.181 11.652 8.9 96.00 7.81 0.05 0.01 1.57 0.40 -0.07 0.00 (5333) Kanaya a 20180901 2.345 0.168 10.973 13.1 13.59 0.04 0.04 0.00 0.60 0.19 -0.12 0.03 (7030) Colombini a 20190322 2.440 0.238 9.335 13.8 6.53 0.48 0.07 0.01 1.55 0.56 non-primitive (8106) Carpino a 20190716 2.412 0.218 9.703 13.7 8.07 1.93 0.10 0.04 0.57 0.49 -0.27 0.01 (19145) 1989 YC 20200203 2.401 0.271 10.408 14.1 7.94 0.23 0.08 0.02 0.61 0.13 -0.37 0.00 ” 20200129 ” ” ” ” ” ” ” ” 0.70 0.26 -0.33 0.02 (28424) 1999 XA 20181004 2.419 0.246 11.931 14.2 6.98 0.41 0.08 0.02 1.13 0.20 too noisy (33913) 2000 LK14a 20190716 2.366 0.223 11.230 14 7.76 0.17 0.10 0.02 0.79 0.49 -0.56 0.02 (35135) 1992 RO1 20180622 2.406 0.258 12.790 14.9 6.23 0.08 0.07 0.00 3.15 1.05 too noisy (46566) 1991 RW21 20181228 2.423 0.257 9.857 14.4 7.31 1.87 0.05 0.02 1.33 0.09 -0.57 0.01 (53178) 1999 CT35 20180622 2.397 0.207 12.084 14.5 6.27 0.03 0.06 0.00 0.46 1.05 -0.21 0.00 (66432) 1999 NL46 20180901 2.428 0.273 8.742 14.8 6.84 0.07 0.03 0.01 1.56 1.05 -0.51 0.01 ” 20180628 ” ” ” ” ” ” ” ” 1.097 0.54 -0.473 0.00 (74832) 1999 TH26 20190118 2.385 0.197 10.074 14.8 6.30 2.97 0.06 0.11 3.76 0.40 -0.30 0.04 (95018) 2002 AZ9 a 20170203 2.418 0.266 8.879 14.9 5.31 0.31 0.08 0.02 0.58 0.16 -0.37 0.03 (98818) 2000 YH125 20200203 2.423 0.184 11.252 15.5 3.84 0.64 0.08 0.02 1.03 0.13 -0.09 0.02 (134740) 2000 AX187 20190322 2.403 0.187 12.277 15 n/a n/a n/a n/a -0.10 0.56 -0.06 0.02
65 Table 4.3: Observational details. A blank box indicates that no solar analog was observed for that night. a Indicates that the number corresponds to star in Table 4.4.
a Number Date UT start Telescope Airmass alpha V Texp [s] SA 313 20170130 22:12 TNG 1.566 27.4 13.2 420 5 5333 20170204 3:56 TNG 1.38 23.7 16 2160 4,6 7030 20190322 13:05 IRTF 1.461 20.6 17.4 2400 5 8106 20190716 9:46 IRTF 1.486 12.5 16.6 1440 1 19145 20200203 9:49 IRTF 1.04 3.9 17.4 2400 5 19145 20200129 22:15 TNG 1.23 6.4 17.4 1440 4,7 28424 20181004 7:40 IRTF 1.949 19.9 16 480 2 33913 20190716 11:04 IRTF 1.393 24.7 17.8 2640 1 35135 20180622 9:32 IRTF 1.25 13.6 16.8 2160 2,3 46566 20181228 15:08 IRTF 1.064 19.9 17.5 2400 53178 20180622 7:58 IRTF 1.216 13.2 16.9 1920 2,3 66432 20180628 4:04 TNG 1.24 18.8 17 2880 1,2,6 66432 20180901 5:07 IRTF 1.336 21.3 16.8 1680 1,2 74832 20190118 12:22 IRTF 1.124 19.6 17.3 1800 95018 20170203 0:14 TNG 1.34 1.1 17.9 2880 4,6 98818 20200203 12:32 IRTF 1.01 9.9 17.8 2640 5 134740 20190322 8:12 IRTF 1.001 17.4 17.7 1200 5
66 Table 4.4: Coordinates and magnitudes of solar analog stars used.
ID Star RA Dec V
1 SA 115-271 23 42 41.8 +00 45 14 9.7 2 SA 112-133 20 43 11.8 +00 26 15 10.0 3 SA 107-684 15 37 18.1 -00 09 50 8.4 4 SA 98-978 06 51 34.0 -00 11 28 10.5 5 SA 102-1081 10 57 04.0 +00 13 10 9.9 6 SA 107-998 15 38 16.0 +00 15 24 10.5
4.2.1 IRTF
For this study we primarily used the SpeX spectrograph (Rayner et al. 2003) at the 3.0 m NASA Infrared Telescope Facility (IRTF) at the Mauna Kea Observatory in Hawaii. We used the low- resolution PRISM mode with a wavelength coverage of 0.7–2.52 µm, a slit size of 0.8×1500, and an AB beam switching pattern separated by 7.500. We observed 12 asteroids over 9 nights between June 2018 and February 2020.
We chose a local standard star for each asteroid with spectral type G2-5V and solar-like B–V and V–K colors. These local standard stars were observed at similar airmasses as the asteroid and then used to correct for telluric absorptions and to remove the solar spectrum from the asteroid spectrum. We also observed at least one well-defined solar analog star (Table 4.4) at multiple airmasses throughout each night. Standard IRAF packages were used to apply flat fielding, combine, and extract spectra. The telluric contribution of each spectrum was modeled and removed using the ATRAN atmospheric model (Lord 1992). The final spectra are shown in Fig. 4.3, normalized to unity at 1.0 µm.
67 Figure 4.3: Spectra of the 15 objects in our sample. A blue spectrum indicates the object was observed by TNG and green indicates IRTF. The spectral slope S0 between 0.95 and 2.3 µm is shown by a pink line. Grey regions correspond to telluric absorptions. All spectra have been normalized to unity at 1 µm.
68 4.2.2 TNG
Supplementary observations were made with the NICS instrument on the 3.6 m Telescopio Nazionale Galileo (TNG) at the El Roque de Los Muchachos Observatory in Canary Islands (Spain). The ob- servations were performed using the AMICI prism and the 1.500slit, resulting in a wavelength range of 0.8–2.4 µm and a resolving power of ∼35 quasiconstant along the spectrum (Oliva 2000). We used an ABBA beam pattern separated by 1000. We observed 5 asteroids over 5 nights between January 2017 and January 2020.
We did not use local standard stars but did observe at least two solar analog stars each night for telluric absorption correction and to obtain relative reflectance. Standard IRAF packages were used for flat fielding, removal of sky and background, and spectrum extraction. The AMICI prism’s low resolution does not allow for resolution of emission lines from Argon and Xenon lamps, so wavelength calibration was done using a look-up table based on theoretical dispersion predicted by ray tracing. We used a Python script developed by our group to calibrate wavelengths (details in Pinilla-Alonso et al. 2016). The final reflectance spectrum of each object, normalized to unity at 1.0 µm, is shown in Fig. 4.3.
69 4.3 Results and analysis
4.3.1 Characteristics of the sample
The mean albedo of our sample is 0.08±0.04 which is slightly higher than the average for the whole family and might be due to our observing bias for brighter objects. The diameters of objects in our sample range from 4 to 96 km with a median diameter of 7 km. We identified one interloper object based on the shape of its spectrum and its taxonomic type, (7030) Colombini, which we discuss more in Section 4.4.3. We exclude this non-primitive asteroid from the rest of our analysis.
The probable parent body of the Chaldaea family is asteroid (313) Chaldaea. The parent asteroid has been observed in the visible (Chapman et al. 1993, Morate et al. 2019) and the 3 µm region (Feierberg et al. 1985, Lebofsky et al. 1995). The 3 µm spectrum shows an absorption indicative of hydrated material, and Morate et al. (2019) show that 79% of the members of the Chaldaea family show evidence of hydration in the 0.7 µm band. This makes the Chaldaea family the IMB family with the highest fraction of hydrated objects (Table 4.1). Six of the objects in our study overlap with those observed by Morate et al. (2019) in the visible (denoted by a in Table 4.2). The composite spectra for these six objects are presented in Fig. 4.4.
70 Figure 4.4: Composite spectra of the six objects with both visible (from Morate et al. 2019) and NIR spectra.
71 Taxonomical classification (Tholen, Bus, and SDSS) for all of our objects has been taken from the Small Bodies Node of the NASA Planetary Data System1 and is shown in Table 4.5. To determine the taxonomy of each asteroid from our NIR spectra, we classified each NIR (and VNIR, if available) spectrum using the M4AST online tool.2 This tool uses the taxons defined in DeMeo et al. (2009) as a reference. It takes the uploaded spectrum, fits a curve and then compares this curve to each of the taxons at the corresponding wavelengths, using a χ2 fitting procedure. The results of our analysis is given in Table 4.5. Our sample contains 55% C-types, 27% X-types, 9% S types. The analysis of four of the objects was inconclusive due to noise.
1http://pdssbn.astro.umd.edu/ 2http://m4ast.imcce.fr/
72 Table 4.5: Taxonomies of Chaldaea objects from the literature and our sample. a Indicates that a visible spectrum is published by Morate et al. (2019).
Number Tholen Bus SDSS Morate et al. (2019) This work
(313) Chaldaea a C Cgh Inconclusive (5333) Kanaya a Ch C Ch Ch (7030) Colombini a A Sv (8106) Carpino a C Cb Inconclusive (19145) 1989 YC C Cb (28424) 1999 XA Xc (33913) 2000 LK14 a Ch Inconclusive (35135) 1992 RO1 Inconclusive (46566) 1991 RW21 Cgh (53178) 1999 CT35 X Xe (66432) 1999 NL46 Cgh (74832) 1999 TH26 D (95018) 2002 AZ9 a Ch Cgh (98818) 2000 YH125 C Cgh (134740) 2000 AX187 C Xc
Asteroid (74832) was classified as a D-type, which are primitive asteroids that are most abundant in the trojan population. Asteroid (35135) also has a very red slope typical of D-type objects, though we were unable to constrain its taxonomy. It could be possible that these objects are interlopers. Though D-types are not common in the IMB, the IMB does have 1% of the total D-types (Gartrelle et al. 2020). Furthermore, we have seen D-types in both the Klio and Chimaera families. Therefore, we keep these two objects in our analysis.
73 4.3.2 Spectral slopes and curvature
The spectral slope of primitive asteroid spectra can be indicative of composition and degree of alteration (e.g. de Leon´ et al. 2012). We follow the procedure from Pinilla-Alonso et al. (2016) to calculate spectral slope (S0) for each asteroid in the range from 0.95 to 2.3 µm with the expression
dS/dλ S0 = S1.0
where dS/dλ is the rate of change of the reflectance and S1.0 is the reflectance at 1.0 µm. The units of S0 are %/1000A˚ . We follow the procedure in Pinilla-Alonso et al. (2016) to estimate the uncertainty in S0 by comparing the spectra of the solar analog star each night. The calculated S0 and uncertainty are given in Table 4.2. The slope can be visualized by the pink line plotted over the individual spectra in Fig. 4.3. We note that the spectra in Pinilla-Alonso et al. (2016) have a better signal to noise ratio than the spectra presented in this work and therefore their measurement of slope is more diagnostic than ours. Still, we present our results for the slope and curvature of our spectra as a baseline for future work to build upon. We hope that future studies will allow us to add more spectra to the sample, which would increase the signal to noise ratio and therefore confidence in our results.
In our sample, the bluest slope is -0.10±0.56%/1000A˚ and the reddest slope is 3.76±0.40%/1000A˚ . We computed the average of the spectra of the 14 members of the Chaldaea family (we did not include asteroid 7030). The slope of the average spectrum between 0.95 and 2.3 µm is moder- ately red (0.85±0.42%/1000A˚ ). The uncertainty in the slope of the average spectrum is derived by quadratically combining the systematic uncertainties of each asteroid with the uncertainties of each asteroid spectrum and averaging the combinations. The distribution of spectral slopes of the
74 members of the family, as well as the slope of the average spectrum is shown in Fig. 4.5. We note that a recent work derived an uncertainty of σS0 = 0.42%/1000A˚ on slope measurements over 0.8 to 2.4 µm using the SpeX instrument at the IRTF (Marsset et al. 2020). This value is coincidentally the same value for uncertainty in our average spectrum. We note the work of Marsset et al. (2020) for completeness, but do not use their value in our analysis.
Figure 4.5: Comparison of the NIR spectral slope distributions of the Chaldaea and Klio primitive inner belt families. In comparison with Klio, the Chaldaea objects have a wider range of slopes. The slope of the average spectrum of Chaldaea family members (green vertical line) is slightly less red than the slope of the average spectrum of Klio family members (red vertical line). The bin size is 0.33%/1000A˚ . Klio data are from Arredondo et al. (2020).
75 The concavity (c<0) or convexity (c>0) of spectra is another way to quantify differences between families. We follow the procedure in Ziffer et al. (2011) to quantify the curvature (c) of each spectrum by averaging the second derivative of a fourth order polynomial fit in the region between 0.9 and 2.3 µm. We omit the points in the telluric bands because the signal there is essentially zero and we also exclude 0.1 µm from each end to avoid introducing error caused by the divergence of the fit if we had used a higher order. The average value of the second derivative of the polynomial fit to the remaining points is taken to be the value for c. The error in the curvature measurement is found by doing the same fitting process on the spectrum plus and minus the error in the spectrum and taking the difference between the two. Two of the spectra were too noisy for a good fit (Table 4.2), and of the remaining 14 there are 11 concave, 3 flat, and 0 convex objects in our sample (Fig. 4.6). The curvature of the average spectrum of the Chaldaea family members is concave with c = -0.270±0.213%/µm2.
76 Figure 4.6: Distribution of the curvature of Chaldaea and Klio objects. The Chaldaea objects are mostly concave (c<0). The average of the Chaldaea family (green vertical line) is concave whereas the average of the Klio family (red vertical line) is convex. The bin size is 0.5%/µm2. Klio data are from Arredondo et al. (2020).
For more details on spectral slope and curvature calculation see Arredondo et al. (2020). We did not find any correlation between spectral slope and amount of curvature, nor did we find any sig- nificant trends between either parameter and physical attributes (diameter, albedo), observational conditions (airmass, phase angle), or orbital parameters (semimajor axis, eccentricity, inclination).
77 4.3.3 Spectral homogeneity of the Chaldaea sample
The overwhelming majority of Chaldaea objects are C-types. We did find one blue NIR spectrum, which could indicate that the Chaldaea family is diverse in the NIR, however with the present sam- ple the statistical significance of diversity is marginal. This diversity, if real, would be surprising, because the Chaldaea family is homogeneous in the visible (Morate et al. 2019). In Fig. 4.7, we show the spectrum for each individual Chaldaea object (grey lines) as well as the average spectrum of the Chaldaea family (green line). The dashed lines indicate 1σ from the mean. It is possible that the one blue object we observe, asteroid (134740) 2000 AX187, is an interloper or perhaps it has just been more processed than the rest of the family. This object is the only one that does not
have a pV from NEOWISE and it is one of the few that shows a nonconcave shape. Furthermore, the uncertainty in the calculated slope of this object (Table 4.2) is large. The rest of the Chaldaea sample is red sloped, so we conclude that the Chaldaea family is homogeneous in the NIR.
78 Figure 4.7: Each individual spectrum from our sample (grey) compared with the average spectrum of the Chaldaea family (green). The dashed lines represent ±1σ of the mean. All spectra have been normalized to unity at 1 µm.
79 In our previous NIR spectroscopic studies of primitive IMB families (Arredondo et al. 2020, Pinilla-Alonso et al. 2016), we found that the Klio and Polana families are both homogeneous and look similar to each other in the NIR despite different compositions and orbital location. How- ever, NIR spectroscopy has proven to be diagnostic for primitive families in other parts of the main belt (e.g., Fornasier et al. 2016, Ziffer et al. 2011), so it is possible that our conclusions so far for the IMB are limited by a small sample size. Our NIR study of the Chaldaea family offers one more sample. It seems there is no evidence of diversity within the Chaldaea family, which agrees from our previous results. We see the strongest evidence that the Chaldaea family looks different from the Klio and Polana families in its average curvature, which we discuss more in Section 4.3.4.
4.3.4 Comparison with Klio
We explore the relationship between the Chaldaea and Klio families to test the hypothesis of Morate et al. (2019) that they potentially originated in the same primordial parent body. We do not compare Chaldaea with Polana because the Klio and Polana families look almost identical in the NIR and also because we do not expect Polana and Chaldaea to originate from the same parent body. The Chaldaea and Klio families are close in semimajor axis, eccentricity, and inclination (Fig. 4.1), closer than the other IMB families. In fact, there is significant overlap in all three or- bital elements. Both of these families are relatively small compared to the larger IMB families like Polana and Erigone (Table 4.1). Nesvorny´ et al. (2015) calculate that the diameter of a sphere with the volume equivalent of all fragments of the Klio family would be 33 km. This does not include the largest member or suspected interlopers, and is much less that the ∼100 km minimum diameter of primordial planetesimals proposed by Morbidelli et al. (2009). There is no equivalent calcula- tion for the Chaldaea family, but based on the fact that it has less than half the number of fragments
80 of the Klio family, we can assume that the equivalent diameter is significantly less than 33 km. If one were to combine both families, the diameter of the primordial parent body would increase but not reach the ∼100 km minimum just mentioned. The remaining pieces of the primordial parent body could be part of the background population or could have been ejected from the inner belt through the ν6 resonance.
Both families show signs of aqueous alteration through the absorption feature at 0.7 µm, though the fraction of hydration of Chaldaea objects is much higher than the fraction of hydration of Klio objects (Table 4.1). In the visible wavelength range, the Chaldaea family are homogenous while the Klio family present diverse spectra. This outcome was unexpected because we previously believed that diversity and hydration were correlated (i.e., de Leon´ et al. 2016, Morate et al. 2016, 2018). However, if the Chaldaea and Klio families are treated as one population, the combined Chaldaea-Klio group would be hydrated and spectrally diverse, which agrees with our previous observations of other IMB families at lower inclinations.
Arredondo et al. (2020) characterized the Klio family in the NIR and found that Klio objects are red, convex, and homogeneous. A comparison between those results and this work is given in Table 4.6. Fig. 4.5 shows the spectral slope distributions for the Chaldaea and Klio families. The Chaldaea objects span almost double the range of slopes than the Klio objects but the average slopes of each family are similar within the standard deviations in the slopes (Table 4.6). We ran a two-sample Kolmogorov-Smirnov (KS) test comparing the slope distributions of the to families to quantify how different they are. The KS statistic, Dm,n, is given by the equation
Dm,n = sup |F (x)1,m − F (x)2,n| x
81 where F (x)1,m and F (x)2,n are the cumulative distribution functions of the two compared distri-
butions. A critical value, Dcrit,m,n is defined as
rm + n D = c(α) crit,m,n mn
where c(α) is 1.36 for α=0.05 and m and n are the sizes of the two samples. If Dm,n>Dcrit,m,n, then
both samples originated in different distributions. We compute that DC,K = 0.250 and Dcrit,C,K = 0.456, or in other words, the distributions are not statistically different.
Table 4.6: Comparison between NIR spectra of the Klio and Chaldaea families.
Klio (Arredondo et al. 2020) Chaldaea (this work)
Average albedo of the sample 0.07±0.01 0.08±0.04 Average slope [%/1000A˚ ] 1.05±0.43 0.85±0.42 Average curvature [%/µm2] 0.51±0.11 (convex) -0.27±0.21 (concave) Slope range [%/1000A˚ ] 2.03 3.86
Therefore, the only difference between the two families is the curvature of the spectra. In Fig. 4.6 we show the distribution of curvature for objects in both families. Unlike Klio, the Chaldaea average is negative, or concave, whereas the average Klio spectrum is convex. A KS test comparing
the c distribution of both families (i.e., Fig. 4.6) gives DC,K = 0.750 and Dcrit,C,K = 0.474, or in other words, the distributions are statistically different. This difference in curvature is prominent when comparing the average spectra of both families (Fig. 4.8), however, the average spectrum of each family is within 1σ of the other (i.e., the solid lines in Fig. 4.8 are within the dashed lines of the opposite colour).
82 Figure 4.8: Comparison of the average spectra for the Chaldaea (green) and Klio (red) families, showing how concave the Chaldaea family is. The dashed lines represent ±1σ of the mean. Klio data are from Arredondo et al. (2020). Spectra have been normalized to unity at 1 µm.
4.4 Discussion
4.4.1 Spectral similarities with Klio – one common parent body?
The results of PRIMASS studies find a clear trend with composition and inclination in the lower- inclination families. The Polana and Clarissa families are homogeneous, anhydrous and cluster tightly around 3°, and the Erigone and Sulamitis families are heterogeneous, hydrated and clus- ter around 5°(Table 4.1). This led us to consider that the families at similar inclinations could be linked. Our working hypothesis (Lowry 2018) is that the parent bodies of families with similar
83 spectra and inclination, but different eccentricities, were fragments of a common progenitor at the same inclination. During the hypothesized giant planet instability, the parent bodies of each spec- trally similar family acquired different eccentricities but preserved their original common inclina- tion; finally, they fragmented generating the families that we see today (e.g., Brasil et al. 2016). It is possible that Polana/Clarissa and Erigone/Sulamitis originated in 2 distinct, non-differentiated parent bodies and that Chaldaea/Klio originated in a third non-differentiated parent body.
The Chaldaea and Klio families share similar inclinations, though they are more spread in incli- nation than the Polana/Clarissa and Erigone/Sulamitis groups. However, as stated before, Chal- daea/Klio overlap in eccentricity while Polana/Clarissa and Erigone/Sulamitis do not. In Section 4.3.4 we compared the Chaldaea and Klio families in the NIR and found only one difference be- tween the two families. Our comparisons show that the Klio and Chaldaea samples are similar in average slope, albedo, and slope dispersion but differ greatly in curvature. Our results in the NIR therefore support the hypothesis of Morate et al. (2019), that the difference in the hydration fraction of the Klio and Chaldaea families could be due to a common parent body with varying degrees of hydration with depth. The reason for the difference in curvature is unknown but could potentially be an effect of space weathering or depth within the parent body.
84 4.4.2 Spectral differences with Klio – inconsistent with space weathering
The age of an asteroid family is defined as the time since the collision that broke the parent body apart. We note that while there is a published age for the Klio family (960±250 Myr from Carruba & Nesvorny´ 2016), we could not find an age for the Chaldaea family in the literature. The age could be inferred from the slope of the Yarkovsky cone of the family (e.g., Fig. 4.2), however, for the purposes of this section we are assuming the Chaldaea/Klio same progenitor hypothesis is true and therefore the Chaldaea objects, which would have been on the outside of the parent body, should show more signs of space weathering (SpWe) than the Klio objects. This comparison also assumes that these two families started out with similar compositions. If they originated in the same primitive parent body, this assumption is not baseless.
The effects of SpWe are well known for high albedo S-type asteroids (redder slopes, darker albedo, attenuated absorption bands), however, laboratory experiments have shown that these trends are not necessarily the same for low albedo objects (e.g., Lantz et al. 2017, 2018, Nakamura et al. 2019, Thompson et al. 2019), and that the results of experiments depend significantly on sample prepa- ration method (chip vs. powder vs. pellet), type of SpWe simulated (solar wind vs micrometeorite bombardment), and energy level. These conflicting results are also mirrored in the ground-based observations of low albedo asteroids (e.g., Campins et al. 2018, Fornasier et al. 2016, Lantz et al. 2013). In Arredondo et al. (2020), we showed that the differences between the older Polana and younger Klio families were consistent with the SpWe predictions made by Lantz et al. (2017): Polana is bluer, brighter, and more convex than Klio in the NIR. In the visible, Morate et al. (2019) showed that the Chaldaea family is bluer than the Klio family, which also agrees with Lantz et al. (2017), however they do not calculate curvature for either family. Our NIR spectra show that Chal-
85 daea is slightly bluer, slightly brighter, but much more concave than the Klio family (Table 4.6), which does not completely agree with Lantz et al. (2017). Our results imply that the Chaldaea family may be older than the Klio family, though these differences are not large. It could be that the families are so close in age that the effects of SpWe aren’t that prevalent, but it is more likely that the effects of SpWe on low albedo objects are more complex than we currently understand.
Another difference between the Chaldaea and Klio families that Morate et al. (2019) suggested could be explained by SpWe is the marked difference between the fraction of hydrated objects in the two populations. We note that primitive asteroids tend to be relatively featureless besides the 0.7 µm absorption band, so the effect of band attenuation is less noticeable in primitive spectra. Campins et al. (2018) suggest that the paucity of the 0.7 µm hydration feature in the Polana family (the oldest family in the IMB) may be due to attenuation by SpWe. We therefore would expect that the outer layer Chaldaea objects should have shallower absorption bands than the inner Klio objects. From the values reported in Morate et al. (2019), the average band depth of Chaldaea objects is 3.10±0.84% and the average band depth of Klio objects is 2.42±0.76%. It seems the Chaldaea bands are deeper than the Klio bands, however they are equal within the uncertainties. This could be a sign that our knowledge about SpWe on primitive objects is incomplete, or it could mean that the Chaldaea family is younger than the Klio family. Without a defined age for the Chaldaea family determined by dynamics, this will remain an open question.
86 4.4.3 The existence of interlopers
In the past, it was assumed that primitive families would look spectrally homogeneous because they are composed of the same material from the original parent body. However, we have recently shown that not all families are homogeneous (Table 4.1). This has interesting implications about how asteroids break up and reaccumulate during catastrophic collisions. Furthermore, we have discovered interlopers within the families based on spectra and albedo that we believe to be mem- bers of the nearby non-primitive Vesta and Flora families. The existence of interlopers such as the S-type (7030) Colombini may help explain the recent identification of exogenous rocks on the surfaces of NEAs. More specifically, Dellagiustina et al. (2019) showed the discovery of Vesta-like material on (101955) Bennu and Tatsumi et al. (2019) showed the presence of ordinary chondrite material on (162173) Ryugu. Similar silicate-rich interlopers have been found in all other prim- itive IMB families that PRIMASS has studied. Possible explanations for this material could be that the interloper objects were perturbed at some point and just happen to have the same orbital elements as the families, or perhaps they are a remainder of one of the bodies in the catastrophic collision (e.g., a primitive and non-primitive body collide, leaving mostly primitive material behind but some remains of the non-primitive object as well). While we do not have the spatial resolution necessary to observe IMB asteroids at the same level of detail that we have observed Bennu and Ryugu, the existence of non-primitive interlopers in the families suggests that IMB asteroids might also present some exogeneous material on their surfaces.
87 4.5 Summary and Conclusion
In the NIR, the Chaldaea family is spectrally homogeneous, red sloped, and concave shaped. Our taxonomy determination using NIR and VNIR spectra shows that the Chaldaea family is dominated by C-type asteroids and that B-type asteroids are not common. In comparison with the similar inclination Klio family, the Chaldaea family has essentially the same average slope, albedo, and slope dispersion. The major difference between the Klio and Chaldaea families is that the Chaldaea family presents very concave spectra, whereas Klio presented a mix of concave and convex spectra. Our NIR comparison is consistent with the hypothesis that the Chaldaea and Klio families have a similar parent body, as suggested by Morate et al. (2019) if the difference in curvature can be explained by a mechanism such as space weathering. Our spectra do not completely agree with the predictions of Lantz et al. (2017) that older primitive objects are bluer, brighter, and more convex. We show that IMB families can look different from each other in the NIR through differences in concavity, which we didn’t see when characterizing the Klio and Polana primitive families. Future work to support these claims include:
• Obtain a better sample of Chaldaea family members with greater signal to noise. One sugges- tion to address this is to compare the significance of the data derived from the observations to that of a simulated data set of similar statistical properties.
• Obtain NIR spectra of the Erigone and Sulamitis families to test if primitive IMB families that are spectrally diverse at visible wavelengths are always homogeneous in the NIR.
• Obtain NIR spectra of the Svea and Chimaera families to test the hypothesis that families with similar inclinations have similar composition.
88 • Obtain visible and NIR spectra of objects from the primitive background population (i.e., asteroids not belonging to any families) to test the correlation between inclination and com- position. These spectra can also be used to test the “asteroids formed big” hypothesis by examining if the background objects near the families are similar in composition to family members.
89 CHAPTER 5: NEAR-INFRARED SPECTROSCOPY OF THE
SULAMITIS ASTEROID FAMILY: SURPRISING SIMILARITIES IN
THE INNER BELT PRIMITIVE ASTEROID POPULATION
Anicia Arredondo1, Humberto Campins1, Noem´ı Pinilla-Alonso2, Julia de Leon´ 3,4, Vania Lorenzi5,3, David Morate3,6
1 Physics Department, University of Central Florida, P.O. Box 162385, Orlando, FL 32816, USA 2 Florida Space Institute, University of Central Florida, Orlando, FL 32816, USA 3 Instituto de Astrof´ısica de Canarias, C/V´ıa Lactea´ s/n, 38205 La Laguna, Tenerife, Spain 4 Departamenta de Astrof´ısica, Universidad de La Laguna, 38205 La Laguna, Tenerife, Spain 5 Fundacion´ Galileo Galilei, INAF, La Palma, Tenerife, Spain 6 Observatorio´ Nacional, Coordenac¸ao˜ de Astronomia e Astrof´ısica, Rio de Janeiro 20921, Brazil
This chapter was published in Icarus in Apr. 2021.
90 5.1 Introduction
The PRIMitive Asteroid Spectroscopic Survey (PRIMASS) aims to characterize primitive colli- sional families across the main belt and some groups in the outer belt such as Cybeles and Hildas. One of PRIMASS’ goals is to use spectral information to determine or constrain the origin of near-Earth asteroids (NEAs). The ν6 resonance is the inner edge of the main asteroid belt, and is responsible for delivering 60% of NEAs to near-Earth space (Bottke et al. 2002) by modifying the object’s eccentricity and semi-major axis while preserving the object’s inclination. At least eight primitive asteroid families have been identified in the inner main belt (IMB) and are a likely source for NEAs: Polana-Eulalia (hereafter Polana), Clarissa, Sulamitis, Erigone, Klio, Chaldaea, Chimaera, and Svea (Nesvorny´ et al. 2015, Walsh et al. 2013, Fig. 5.1). PRIMASS has observed and characterized the Polana, Klio, and Chaldaea families in the visible and NIR (Arredondo et al. 2020, 2021, de Leon´ et al. 2016, Morate et al. 2019, Pinilla-Alonso et al. 2016) and the other five families in the visible only (Morate et al. 2016, 2018, 2019). A summary of previous PRIMASS results is given in Table 5.1. We characterize the families using visible spectra based on the diver- sity of spectral slopes (homogeneous versus diverse), the presence or absence of hydrated minerals (indicated by the presence of the 0.7 µm absorption feature), and the percentage of the observed asteroids showing the 0.7 µm band. In total, PRIMASS has spectroscopically observed over ∼700 asteroids.
91 Figure 5.1: Semimajor axis vs. inclination (left) and eccentricity (right) of the primitive inner belt asteroid families as defined by Nesvorny´ et al. (2015). Objects in this study are highlighted with red diamonds.
92 Table 5.1: Summary of PRIMASS findings in the IMB.
Number observed Family Inc [°] Total number 0.7 µm Visible Spec- NIR Spectra References
of family (Vis|NIR) feature tra members
Polana 3 1924 65 45 Absent Homogeneous Homogeneous de Leon´ et al. (2016), Pinilla- Alonso et al. (2016) Clarissa 3 179 33 - Absent Homogeneous Too faint Morate et al. (2018) Erigone 5 1776 101 - Present Diverse In progress Morate et al. (2016) (58%) Sulamitis 5 303 64 19 Present Diverse Homogeneous Morate et al. (2018) (60%) (this work) Klio 9 330 30 21 Present Diverse Homogeneous Arredondo et al. (2020), (23%) Morate et al. (2019) Chaldaea 12 132 15 15 Present Homogeneous Homogeneous Arredondo et al. (2021), (79%) Morate et al. (2019) Chimaera 14 108 20 - Present Diverse Morate et al. (2019) (20%) Svea 16 48 8 - Absent Homogeneous Morate et al. (2019)
We emphasize the importance of the low inclination, primitive IMB population because the two targets of sample return missions, (101955) Bennu and (162173) Ryugu, are low inclination (6°and 5.9°, respectively), primitive (B-type and C-type, respectively) NEAs. These two objects are dy- namically likely to originate from this primitive IMB population (Campins et al. 2010, 2013, Bot- tke et al. 2015), and understanding the source region of these sample return targets will enhance the mission by offering context to connect the small-scale details seen by the spacecraft with the large-scale characteristics of IMB families.
93 In the lower inclination families (Polana, Clarissa, Erigone, and Sulamitis), we see a clear trend with inclination and visible spectral properties. The Polana and Clarissa families are clustered at 3°inclination, have homogeneous visible spectra, and do not present the absorption band at 0.7 µm indicative of hydration. In contrast, the Erigone and Sulamitis families cluster at 5°inclination, present diverse spectra (i.e., a wide range of taxonomic types and spectral slopes), and both fam- ilies show ∼60% of their members being hydrated. This clear dichotomy in lower inclination families is suggestive of two different parent bodies, one for Polana/Clarissa and the other for Erigone/Sulamitis. Our working hypothesis (Lowry 2018) is that during the postulated giant planet instability (e.g., Brasil et al. 2016), the parent bodies of each spectrally similar family acquired different eccentricities but preserved their original common inclination. These parent bodies then fractured, leaving the families that we see today in Fig. 5.1. If so, we would expect the NIR spectra of Polana/Clarissa and Erigone/Sulamitis to show the same sharp contrasts.
Our observations of the Polana family (Pinilla-Alonso et al. 2016) show that it is homogeneous in the NIR with a range of slopes varying from -0.12%/1000A˚ to 1.78%/1000A˚ , with an average slope of 0.68±0.68%/1000A˚ . The Clarissa family is too faint for observations in the NIR; how- ever, we obtained NIR observations of the largest object and likely parent asteroid (302) Clarissa. The spectrum of (302) Clarissa shows the same general slope and shape that the Polana family do (Fig. 5.2). From this simple comparison, we argue that the Clarissa and Polana families are similar in both the visible and NIR, in agreement with previous PRIMASS findings.
94 Figure 5.2: Spectral comparison of asteroid (302) Clarissa (yellow) and the average spectrum of the Polana family (purple) from Pinilla-Alonso et al. (2016). The shaded purple region indicates ±1σ from the Polana mean spectrum. The spectra are similar, implying that the Clarissa and Polana families are similar in the NIR as they are in the visible. Both spectra are normalized to unity at 1 µm.
95 We have not yet published our results on the characterization of the Erigone family in the NIR; however, observations have been completed and analysis is underway. For this reason, we cannot compare to see if the Sulamitis and Erigone families are similar to each other like the Polana and Clarissa families are. In this work, we characterize the Sulamitis family and compare it to the Polana family in order to test if the difference seen between these families in the visible is also present in the NIR. Our observations are described in Section 5.2. The computed reflectance spectra are shown in Section 5.3. We compare our spectra with other IMB families (especially Polana) and with Bennu and Ryugu in Section 5.4 and we present our conclusions in Section 5.5.
5.2 Observations and data reduction
We selected the asteroids in this study from the list of family members compiled by Nesvorny´ et al. (2015). We observed 19 of the 303 total family members, which is the practical limit for this faint family (average H = 16.2) based on the size of the telescopes used in this study. We prioritized observing objects that PRIMASS has published visible spectra for in Morate et al. (2018). Table 5.2 shows the properties of each asteroid observed in this study, including osculating orbital elements (a,e,i) from the Minor Planet Center Orbit Database file (MPCORB.DAT4), absolute magnitude
(H), diameter (D), geometric albedo (pV ) from NEOWISE (Mainzer et al. 2019), Bus-DeMeo taxonomy calculated in this work, and values for slope (S0) and curvature (c) calculated in this work. The average albedo for our sample is 0.07±0.02, which is consistent with the Sulamitis family being a primitive family. Diameters of objects in our sample range from 4 to 61 km, with a median size of 8.6 km.
96 Table 5.2: Physical and dynamical properties of the observed targets. We include our computed values for the spectral slope (S0) and curvature (c) and their associated errors. See Section 5.3.2 for more details. We include all observed objects in this table even if they were non-primitive and are not used in our analysis (Section 5.3.1). Also included are the parameters for asteroid (302) Clarissa presented in Fig. 5.2. * Indicates that a visible spectrum is published by Morate et al. (2018).
0 2 ˚ 0 Asteroid a [AU] e i [°] H D [km] errD [km] pV errpV Tax S [%/1000A] errS c [%/µm ] errc (752) Sulamitis * 2.463 0.074 5.962 10.3 60.85 1.59 0.03 0.00 X 1.84 0.40 -1.29 0.10 (1923) Osiris * 2.436 0.063 4.958 13.6 13.46 0.21 0.03 0.01 X 0.08 0.19 -0.24 0.02 (6806) Kaufmann * 2.454 0.093 4.670 14.1 8.99 3.56 0.06 0.05 T 1.62 0.08 -0.46 0.05 (9476) 1998 QQ36 * 2.453 0.109 4.488 14.0 8.61 2.34 0.07 0.07 Cg 0.54 0.10 0.67 0.00 (12421) Zhenya * 2.450 0.122 5.556 14.2 6.90 2.26 0.09 0.08 Cg 1.31 0.08 0.10 0.09 (13509) Guayaquil * 2.475 0.097 4.118 14.4 10.23 0.94 0.04 0.27 C 1.44 0.13 0.95 0.07 (24726) 1991 VY * 2.454 0.093 5.856 14.1 9.65 3.48 0.05 0.11 X 2.40 0.52 -1.94 0.60 (24907) Alfredhaar * 2.454 0.069 4.659 15.5 N.A. N.A. N.A. N.A. X 0.09 0.13 -0.22 0.10 (25036) Elizabethof * 2.476 0.084 4.175 14.5 7.12 0.12 0.06 0.01 Cg 1.02 0.08 0.53 0.01 (26807) 1982 RK1 * 2.473 0.104 4.217 14.0 N.A. N.A. N.A. N.A. X 0.16 0.19 0.10 0.04 (28894) Ryanchung * 2.477 0.115 5.034 14.3 4.91 1.41 0.07 0.03 Ch 0.06 0.44 0.38 0.10 (34890) 2001 VS62 2.467 0.083 4.149 14.8 N.A. N.A. N.A. N.A. Sq 3.01 0.49 0.54 0.10 (42411) 3249 T-1 2.467 0.107 4.870 14.6 8.15 0.05 0.05 0.01 Cb 1.27 0.49 -0.06 0.01 (50239) 2000 BW3 * 2.475 0.106 5.965 14.0 9.22 0.33 0.06 0.01 T 0.75 0.10 0.18 0.18 (53537) 2000 AZ239 * 2.449 0.079 4.837 14.4 N.A. N.A. N.A. N.A. Sq 3.82 0.44 1.13 0.09 (59065) 1998 UB43 * 2.446 0.122 5.358 15.1 6.58 2.01 0.05 0.03 Xc 0.76 0.19 0.27 0.10 (79143) 1992 BQ2 * 2.489 0.065 4.354 14.4 3.90 0.25 0.24 0.06 V 2.58 0.13 -1.03 0.06 (119526) 2001 UF175 * 2.479 0.079 5.574 15.6 N.A. N.A. N.A. N.A. Xk 1.44 0.44 0.62 0.07 (244905) 2003 WJ120 2.482 0.118 4.460 16.0 N.A. N.A. N.A. N.A. V 3.37 0.52 -1.19 0.51
(302) Clarissa 2.407 0.111 3.412 11.1 38.53 3.1 0.05 0.01 Not calculated as part of this study
Our observations were obtained using two telescopes, the NASA InfraRed Telescope Facility (IRTF) in Hawaii and the Telescopio Nazionale Galileo (TNG) in Spain. Table 5.3 shows the observing conditions for each target, including the UT date and time, telescope used, airmass at
the start of observations (X), phase angle (α), visual magnitude (V), total exposure time (Texp), and solar analog star used (SA). For all observations we tried to observe as close to the meridian as
97 possible and with the slit aligned with the parallactic angle. We utilized an ABBA beam switching pattern to subtract the sky contribution and we observed well known solar analog stars (Table 5.4) each night to subtract the solar contribution. The exposure time of each object was determined based on sky conditions and visual magnitude. Further details of observations and reduction for each telescope are described below.
98 Table 5.3: Observational details. We include all observed objects in this table even if they were non-primitive and were not used in our analysis (Section 5.3.1). Also included are the parameters for asteroid (302) Clarissa presented in Fig. 5.2. * Indicates that a visible spectrum is published by Morate et al. (2018). † Indicates that numbers in this column correspond to stars in Table 5.4.
Asteroid Date UT start Telescope X α VTexp [s] SA† (752) Sulamitis * 20170131 03:14 TNG 1.11 17.7 14.0 720 4 (1923) Osiris * 20170929 08:02 IRTF 1.20 1.4 16.5 1440 2 (6806) Kaufmann * 20200202 13:22 IRTF 1.23 21.1 18.6 4320 5 (9476) 1998 QQ36 * 20170823 11:14 IRTF 1.70 13.6 17.0 1920 1 (12421) Zhenya * 20200202 12:05 IRTF 1.02 4.5 18.0 2880 5 (13509) Guayaquil * 20200203 13:38 IRTF 1.11 16.3 18.6 4080 5 (24726) 1991 VY * 20180828 05:46 TNG 1.05 26.6 18.5 3560 1 (24907) Alfredhaar * 20200203 10:48 IRTF 1.00 2.9 18.2 3120 5 (25036) Elizabethof * 20200202 10:27 IRTF 1.05 5.5 18.0 2880 5 (26807) 1982 RK1 * 20170929 07:15 IRTF 1.24 3.6 16.8 1440 2 (28894) Ryanchung * 20190424 08:08 IRTF 1.22 3.9 18.1 4560 5 (34890) 2001 VS62 20190716 06:46 IRTF 1.32 19.8 18.8 2880 1 (42411) 3249 T-1 20190716 12:45 IRTF 1.28 24.2 18.8 4560 1 (50239) 2000 BW3 * 20170823 14:01 IRTF 1.46 8.2 16.9 1920 1 (53537) 2000 AZ239 * 20190424 05:57 IRTF 1.00 23.3 18.7 3120 5 (59065) 1998 UB43 * 20170929 09:00 IRTF 1.10 3.1 17.4 3360 2 (79143) 1992 BQ2 * 20200203 15:18 IRTF 1.21 22.7 18.5 2640 5 (119526) 2001 UF175 * 20190422 09:11 IRTF 1.10 5.1 18.4 3600 5 (244905) 2003 WJ120 20180828 04:11 TNG 1.70 2.6 18.4 3170 1
(302) Clarissa 20171004 04:39 TNG 1.25 9.1 13.9 640 3
99 Table 5.4: Coordinates and magnitudes of solar analog stars used. The number in the ID column corresponds to the number in the last column of Table 5.3.
ID Star RA Dec V
1 SA 115-271 23 42 41.8 +00 45 14 9.7 2 SA 112-133 20 43 11.8 +00 26 15 10.0 3 SA 93-101 01 53 18.0 +00 22 25 9.7 4 SA 98-978 06 51 34.0 -00 11 28 10.5 5 SA 102-1081 10 57 04.0 +00 13 10 9.9
5.2.1 IRTF
The IRTF is a 3.0 m telescope sited at the Mauna Kea Observatory in Hawaii. We used the SpeX spectrograph (Rayner et al. 2003) in its low-resolution PRISM mode which yields a wavelength coverage of 0.70–2.52 µm. We used the 0.8×1500slit for all observations. In addition to the solar analog stars in Table 5.4, we also chose one local standard star for each asteroid with spectral type G2-5V and solar-like B–V and V–K colors and observed it at an airmass similar to that of the asteroid to correct for telluric absorptions. In total we observed 16 asteroids between August 2017 and February 2020.
To reduce the data, we used the IRTF supplied program Spextool (Cushing et al. 2004), which applies flat fielding, combines, and extracts the spectrum from each frame. We modelled and removed the telluric contribution from each spectrum using the IRTF supplied program xtellcor (Vacca et al. 2003). Fig. 5.3 shows the final spectra, normalized to unity at 1.0 µm.
100 Figure 5.3: Spectra of the 19 objects in our sample. A blue spectrum indicates the object was observed by TNG and green indicates IRTF. The spectral slope S0 between 0.95 and 2.30 µm is shown by a pink line. Grey regions correspond to telluric absorptions. We include all observed objects in this figure even if they were non-primitive and are not used in our analysis (Section 5.3.1). All spectra have been normalized to unity at 1 µm. 101 5.2.2 TNG
The TNG is a 3.58 m telescope sited on the El Roque de Los Muchachos Observatory in La Palma, Canary Islands, Spain. We used the NICS instrument with its AMICI prism and the 1.500slit. This setup provides spectra in the range 0.8–2.4 µm, with a resolving power of ∼35 quasi-constant along the spectrum (Oliva 2000). In total we observed 3 asteroids between January 2017 and August 2018.
We used standard IRAF packages for flat fielding, removal of sky and background, and extraction of the 1D spectrum. The low resolution of the AMICI prism does not allow for emission lines from Argon and Xenon lamps to be resolved, and so wavelength calibration was done using a look-up table based on theoretical dispersion predicted by ray tracing. We used a Python script developed by our group to calibrate wavelength and to divide the asteroid spectrum by the solar analog spectrum and remove the solar contribution. Fig. 5.3 shows the final spectra, normalized to unity at 1.0 µm.
5.3 Results and analysis
Analysis of our observed sample includes taxonomy determination, spectral slope and curvature calculations, classification as a homogeneous family, comparison with spectra of other primitive IMB families, and comparison with primitive NEAs Bennu and Ryugu.
102 5.3.1 Taxonomy determination
Sixteen of the objects in our sample were also observed at visible wavelengths by Morate et al. (2018). We used the M4AST1 online tool (Popescu et al. 2012) to combine the spectra of these objects to create VNIR spectra (Fig. 5.4). The ability to combine visible and NIR wavelengths allows a more accurate taxonomy determination for each object. For example, de Leon´ et al. (2012) show that while B-types always have a blue slope in the visible, their spectra can vary from blue sloped to red sloped in the NIR.
1http://m4ast.imcce.fr/
103 Figure 5.4: Composite spectra of 16 objects with visible (from Morate et al. 2018) and NIR spectra. All spectra are normalized to 0.55 µm.
104 We also used M4AST for taxonomic classification. The tool compares the submitted spectra with the Bus-DeMeo taxonomies (DeMeo et al. 2009) and returns the taxonomic class with the lowest χ2 residual. We performed this on the composite VNIR spectrum (if available) or the NIR spectrum alone of each object. Our computed taxonomic class for each object is given in Table 5.2. Our sample contains six C-Types, seven X-types, two T-types, and two non-primitive types (Fig. 5.5).
Figure 5.5: Taxonomic classes from NIR spectra of the Polana family (Pinilla-Alonso et al. 2016) and the Sulamitis family (this work).
We note that 4 of the objects returned non-primitive taxonomic types (34890, 53537, 79143, and 244905). All four of these objects also have either high or unreported albedos (Table 5.2). Al- though these objects were listed as family members by Nesvorny´ et al. (2015), we conclude they are interlopers (possibly from the nearby Flora or Vesta families) and we exclude them from the rest of our analysis because our focus is the primitive population. Another method for determining
105 interlopers is looking at the C parameter for each object given by Nesvorny´ et al. (2015). Bodies with |C|>1 are suspected interlopers. In our sample, 2 objects have |C|>1: 79143, which we have already found to be an interloper, and 1923. Because 1923 has albedo and spectral properties con- sistent with the rest of the Sulamitis family, we do not exclude it from the rest of the analysis. This brings our total number of objects for the rest of the analysis to 15.
5.3.2 Spectral slopes and curvature
We calculated the spectral slope (S0) of each spectrum between 0.95 and 2.30 µm following the procedure outlined in Pinilla-Alonso et al. (2016). We used the equation
dS/dλ S0 = S1.0
where dS/dλ is the rate of change of the reflectance and S1.0 is the reflectance at 1.0 µm. Spectral slope can be blue (S0<0), flat (S0≈0), or red (S0>0). We estimate the systematic uncertainty in S0 by dividing each solar analog spectrum by a reference spectrum, and then measuring the spread in the slopes of those spectra. We also note that a recent work derived an uncertainty of σS0 = 0.42%/1000A˚ on slope measurements over 0.8 to 2.4 µm using the SpeX instrument at the IRTF (Marsset et al. 2020). The computed S0 for each object is visualized by the pink line in Fig. 5.3 and given in Table 5.2.
The spectral slopes for our 15 primitive objects range from 0.06±0.44%/1000A˚ to 2.40±0.52%/1000A˚ . We computed the average spectrum of the primitive objects in our sample, and the slope of the av- erage spectrum is 0.89±0.26%/1000A˚ . Fig. 5.6 shows the average spectrum (blue) compared with the spectra of the individual objects (grey) along with ±1σ of the mean (shaded blue).
106 Figure 5.6: Each individual spectrum from our sample of 15 primitive objects (grey) compared with the average spectrum of the Sulamitis family (blue). The blue shaded region represents ±1σ of the mean. All spectra have been normalized to unity at 1 µm.
107 The curvature (c) of a spectrum is the amount of concavity (c<0) or convexity (c>0), or the deviation from being flat (c≈0). The average curvature of a family is a useful tool for distinguishing between families using their NIR spectra (e.g., Arredondo et al. 2021). We quantify the curvature of each spectrum using the method detailed in Ziffer et al. (2011). First, we omit points in the telluric bands and 0.1 µm from each edge, to ensure a reliable polynomial fit. We then fit a fourth order polynomial to the region between 0.9 and 2.3 µm and take the average of the second derivative of that fit. The uncertainty in this measurement is calculated by doing the same process on the spectrum plus and minus the associated error (i.e., shaded portion in Fig. 5.6), and taking the difference between them.
We apply this method to the 15 primitive objects in our sample and find that there are 5 concave, 4 flat, and 6 convex objects (Fig. 5.7). This is a fairly even distribution with no preference towards a particular curvature. The curvature of the average spectrum is semi-flat (0.242±0.072%/µm2), which reflects the fact that the curvatures are roughly an even mix of concave and convex.
108 Figure 5.7: Distribution of the curvature in the NIR of IMB family objects. The curvature of the average spectrum for each family is presented as a vertical line of the same color. Klio data are from Arredondo et al. (2020) and Chaldaea data are from Arredondo et al. (2021). Curvatures for the Polana family were not calculated by Pinilla-Alonso et al. (2016) but we ran our analysis on their data to compute curvatures. The bin size is 0.4%/µm2.
5.3.3 Reanalysis of Polana family data
In the next section we compare the Sulamitis family spectra analyzed in this work with the spectra of other families observed by PRIMASS. To do this effectively, we need to add on to the work done by Pinilla-Alonso et al. (2016). We computed the NIR (or VNIR if available) taxonomies of the Polana family objects and calculate the curvature of each object between 0.9 and 2.3 µm following the same methods used on the Sulamitis objects. Our results are given in Table 5.5 and are visualized in Fig. 5.5 and Fig. 5.7. We classify 91% of the Polana as C-types, 7% as X-types, and 2% other. The curvature of the average Polana NIR spectrum is 0.113±0.024%/µm2.
109 Table 5.5: Results of our reanalysis of Polana family NIR spectra obtained in Pinilla-Alonso et al. (2016). * Indicates that the taxonomy was determined from VNIR spectra (visible from de Leon´ et al. 2016).
2 Number Tax c [%/µm ] errc 2 Number Tax c [%/µm ] errc 112 Iphigenia X * -0.11 0.01 3566 Levitan C * 0.56 0.02 142 Polana C * -0.11 0.06 3999 Aristarchus C * 0.02 0.21 495 Eulalia C * 0.11 0.01 4173 Thicksten C * 1.05 0.05 557 Violetta C 0.02 0.00 4219 Nakamura C * 0.68 0.13 750 Oskar C * 0.05 0.21 6118 1986 QX3 C 0.18 0.01 1012 Sarema C * 0.08 0.04 6142 1999 FP C * 0.23 0.01 1183 Jutta C * -0.17 0.01 6278 Ametkhan C -0.12 0.05 1190 Pelagia C 0.22 0.01 6471 Collins C 0.03 0.15 1267 Geertrudia C -0.02 0.06 6661 Ikemura C * 0.77 0.02 1493 Sigrid C * 0.20 0.62 6712 Hornstein C 0.10 0.03 1650 Heckmann C * 0.13 0.01 6769 Brokoff C * -0.10 0.04 1740 Paavo Nurmi C -0.08 0.08 8212 Naoshigetani S N/A N/A 2007 McCuskey C 0.07 0.24 8927 Ryojiro C 0.50 0.03 2066 Palala C -0.15 0.02 11118 Modra C -0.20 0.00 2081 Savaza C * 0.84 0.13 11338 Schiele C 0.37 0.33 2139 Makharadze C 0.12 0.07 13997 1993 FB32 C 0.02 0.02 2276 Warck C 0.39 0.12 14215 1999 TV6 C 0.55 0.03 2441 Hibbs X 0.10 0.02 16132 Angelakim C 0.36 0.02 2662 Kandinsky X 0.48 0.01 23893 Lauman C 0.19 0.01 3064 Zimmer C 0.12 0.08 24656 1987 QT7 C 0.03 0.02 3130 Hillary C * 0.35 0.03 25024 Calebmcgraw C 0.35 0.04 3228 Pire C 0.07 0.15 66336 1999 JB62 C 0.14 0.03 3247 Di Martino C 0.18 0.02 110 5.4 Discussion
5.4.1 Spectral homogeneity in the NIR despite diversity in the visible
In Fig. 5.8 we show the distribution of the NIR spectral slopes in our sample (blue histogram) and the slope of the average spectrum (blue dashed vertical line). The Sulamitis objects span a 2.34%/1000A˚ range of slopes. This, along with the distribution of taxonomies shown in Fig. 5.5 (i.e., mostly X-types and C-types and no B-types), allow us to classify Sulamitis as a spectrally homogeneous family in the NIR. This is not the first time that PRIMASS has found clear homo- geneity in a NIR sample despite heterogeneity at visible wavelengths (Table 5.1). Arredondo et al. (2020) posit that all primitive IMB families are homogeneous in the NIR, which is supported by our observations of the Sulamitis family.
111 Figure 5.8: Comparison of the NIR spectral slope distributions of the Sulamitis family (blue) with the Polana (purple), Klio (red) and Chaldaea (green) families. The slope of the average spectrum for the Sulamitis family is presented by a dashed vertical line. The slopes of the average spectrum for the other families are presented as solid vertical lines. The range of Sulamitis slopes overlap almost perfectly with the Polana and Klio families. The Sulamitis average is almost indistinguish- able from the average of the Chaldaea family. Klio data are from Arredondo et al. (2020), Chaldaea data are from Arredondo et al. (2021), and Polana data are from Pinilla-Alonso et al. (2016). The bin size is 0.5%/1000A˚ .
5.4.2 Comparison with other IMB families
The three other families PRIMASS has observed in the NIR are Polana, Klio, and Chaldaea. We mainly focus on the comparison with the Polana family to see if the distinct differences between Sulamitis and Polana in the visible extend to the NIR; however, we compare with Klio and Chal- daea for completeness. The average albedo, average slope, average curvature, and range of slopes for our sample of each family are given in Table 5.6.
112 Table 5.6: Comparison between PRIMASS samples of the IMB families in the NIR.
Polana Klio Chaldaea Sulamitis (Pinilla-Alonso et al. 2016) (Arredondo et al. 2020) (Arredondo et al. 2021) (this work)
Average pV of 0.05 ± 0.01 0.07 ± 0.01 0.08 ± 0.04 0.07 ± 0.02 the sample Average curvature 0.11 ± 0.02 0.51 ± 0.11 -0.30 ± 0.21 0.24 ± 0.07 [%/µm2] Average slope 0.68 ± 0.68 1.05 ± 0.43 0.85 ± 0.42 0.89 ± 0.26 [%/1000A˚ ] Slope range [-0.12,1.78] [0.20, 2.23] [-0.10, 1.57] [0.06,2.40] [%/1000A˚ ]
The Sulamitis family is similar in albedo to the three other IMB families. The average curvature of the Sulamitis family is similar to that of the Polana family (flat) but is less than the Klio curvature and significantly larger than the Chaldaea curvature. Fig. 5.7 shows a comparison of the distribu- tion of curvatures of each family. Like Klio, there is an even mix of concave and convex Sulamitis objects, and like Polana, the average curvature is essentially flat. Arredondo et al. (2021) show that the curvature distribution between the Klio family and Chaldaea family are statistically different and using the same analysis we find that Sulamitis and Chaldaea are significantly different as well.
The average slope of the Sulamitis family is comparable to that of the other families, and all are within 1σ of the mean of each other. Fig. 5.8 shows a comparison of the slope distributions of each family. All four families have a narrow range in slope (∼2%/1000A˚ ), though the center of this distribution (i.e., average slope) differs slightly. We use a Kolmogorov-Smirnov test (KS-test) to quantify if the slope distributions are statistically similar or not. We use the expression:
Dm,n = sup |F (x)1,m − F (x)2,n| x
113 where F (x)1,m and F (x)2,n are the cumulative distribution functions of the two compared distri-
butions. A critical value, Dcrit,m,n is defined as
rm + n D = c(α) crit,m,n mn
where c(α) is 1.36 for α=0.05 and m and n are the sizes of the two samples. If Dm,n>Dcrit,m,n, then the null hypothesis is rejected, meaning both samples originated in different distributions.
Table 5.7 shows the results of the test. For all 3 cases, D Table 5.7: Results of the KS-test comparing NIR spectral slopes of the IMB families. D Dm,n Dcrit,m,n Polana vs Sulamitis 0.29 0.40 Klio vs Sulamitis 0.28 0.49 Chaldaea vs Sulamitis 0.33 0.47 Fig. 5.8 shows that the slope distributions of the Sulamitis and Polana families are essentially identical, and the KS test confirms that the populations cannot be differentiated using spectral slope in the NIR. The slope distribution of the Sulamitis family differs more from the Klio and Chaldaea families, but the KS test shows that this difference is not statistically significant. Arredondo et al. (2020) concluded that the Klio and Polana families essentially identical in the NIR, and this work shows that the Sulamitis family is essentially identical as well. 114 We emphasize that the result of the Sulamitis family looking similar to the Polana family is coun- terintuitive based on our past the PRIMASS findings in the visible and the difference in NIR taxonomies shown in Fig. 5.5. The dichotomy of composition seen in the visible spectra of the Polana family vs the Sulamitis family is not present in the NIR sample. This conclusion can be further tested by comparing the VNIR spectra of all four low inclination families (Polana, Clarissa, Erigone, and Sulamitis) however that is outside of the scope of this paper. Arredondo et al. (2020) posit that NIR spectroscopy is not as diagnostic as visible spectroscopy for primitive families in the IMB, and our results support this. However, we point out that NIR has proven to be quite diagnostic for primitive families elsewhere in the asteroid belt (e.g., de Leon´ et al. 2010, Fornasier et al. 2016, Ziffer et al. 2011). 5.4.3 Comparison with (101955) Bennu and (162173) Ryugu Dynamical models and spectral comparisons show that the large majority of NEAs were likely delivered from low-inclination (i<8°), inner-asteroid belt (2.15 AU cluding the Sulamitis family, through the ν6 resonance and the 3:1 mean motion resonance with Jupiter (Bottke et al. 2002, Granvik et al. 2018). Fig. 5.1 shows that the Sulamitis family is closer to the 3:1 resonance than the ν6 resonance. Interactions with these resonances will perturb the object’s semi-major axis and eccentricity while preserving its inclination. Both spacecraft target NEAs, (101955) Bennu and (162173) Ryugu, are located at ∼6°inclination so we must consider the 5°inclination Sulamitis family as a possible source region for them. 115 We compare our spectra with the ground-based (Clark et al. 2011, Abe et al. 2008) and space-based (Hamilton et al. 2019, Kitazato et al. 2019) spectra of both objects in order to test the hypothesis that either object could have originated in the Sulamitis family. A plot of the spectra of the NEAs compared with the spectra of the Sulamitis family is shown in Fig. 5.9. Bennu is a B-type asteroid which is seen in its blue-sloped NIR spectrum (indicated by red and dark red lines). There are no B-types in our Sulamitis sample, and both of the Bennu spectra lie outside of 1σ of the average Sulamitis spectrum (indicated by the blue dashed lines). Based on our spectra, we can rule out the Sulamitis family as a source of Bennu. Pinilla-Alonso et al. (2016) and de Leon´ et al. (2018) show that a more likely source for Bennu is the Polana family. Figure 5.9: Comparison of the average spectrum of the Sulamitis family to spectra of spacecraft targets (101955) Bennu and (162173) Ryugu. The average Sulamitis spectrum (blue) appears to be compatible with Ryugu (green; right), but not Bennu (red; left). Blue shaded region is ±1σ of the mean Sulamitis spectrum. The spectra are normalized at 2.2 µm for a better comparison. 116 The Ryugu spectrum from Abe et al. (2008) (light green line) is almost an exact match to the average Sulamitis spectrum and fits well within 1σ of the average Sulamitis spectrum. The Ryugu spectrum from Kitazato et al. (2019) (dark green line) covers a shorter wavelength range than our data (1.8–2.5 µm), but the wavelengths that overlap with our spectra agree very well. de Leon´ et al. (2018) show that the most likely source for Ryugu is the Polana family. From our NIR spectra, we conclude that the Sulamitis family cannot be discounted as a possible source for asteroid Ryugu. 5.5 Summary and conclusion In the NIR, the Sulamitis family is spectrally homogeneous, red sloped, and there is an even mix of concave and convex curvatures. We do not see diversity within the Sulamitis family in the NIR, which is consistent with our observations of homogeneity in the other IMB families in the NIR. The Sulamitis family looks similar to the other IMB families in albedo, average spectral slope, slope dispersion, and average curvature (except Chaldaea), despite the sharp differences in visible slope distribution and percentage of hydrated objects. We do not see evidence of significant differences between low inclination families (i.e., Polana and Sulamitis) like we see in the visible, which is consistent with our hypothesis that the NIR is not diagnostic for primitive IMB families. Our taxonomy determination shows that the majority of the Sulamitis family are X-types and C-types, which combined with spectral comparisons shows that the NEA spacecraft target Ryugu could potentially have originated in the Sulamitis family. Future work to support these claims include: • Obtain NIR spectra of the Erigone family to test if the Erigone family is similar to Su- lamitis but different from Polana and Clarissa. This would further test the links between Polana/Clarissa and Erigone/Sulamitis that PRIMASS has observed in the visible. We have completed observations of the Erigone family, and the data are currently in analysis. 117 • Obtain visible and NIR spectra of objects from the Svea and Chimaera families and from the primitive background population (i.e., asteroids not belonging to any families) to test the correlation between inclination and composition in the IMB. 118 CHAPTER 6: SPECTROSCOPY OF THE INNER BELT PRIMITIVE ASTEROID BACKGROUND POPULATION Anicia Arredondo1, Humberto Campins1, Noem´ı Pinilla-Alonso2, Julia de Leon´ 3,4, Vania Lorenzi3,5, David Morate6, Juan Luis Rizos3,4 Mario´ De Pra´2, 1 Physics Department, University of Central Florida, P.O. Box 162385, Orlando, FL 32816, USA 2 Florida Space Institute, University of Central Florida, Orlando, FL 32816, USA 3 Instituto de Astrof´ısica de Canarias, C/V´ıa Lactea´ s/n, 38205 La Laguna, Tenerife, Spain 4 Departamenta de Astrof´ısica, Universidad de La Laguna, 38205 La Laguna, Tenerife, Spain 5 Fundacion´ Galileo Galilei, INAF, La Palma, Tenerife, Spain 6 Observatorio´ Nacional, Coordenac¸ao˜ de Astronomia e Astrof´ısica, Rio de Janeiro 20921, Brazil The work in this chapter was previously submitted to Icarus. 119 6.1 Introduction The PRIMitive Asteroid Spectroscopic Survey (PRIMASS) was created to study the origin of near- Earth asteroids (NEAs) with the goal of providing context for the OSIRIS-REx and Hayabusa2 missions by determining the origins and likely compositions of mission targets (101955) Bennu and (162173) Ryugu. We focus on the primitive (pV <0.15) asteroid population because both Bennu and Ryugu are primitive asteroids. The compositional study of NEAs is also important for choosing future spacecraft mission targets, in-situ resource utilization (ISRU), hazard mitigation, and exploring the origins of the solar system. Bottke et al. (2002, 2015) show that the main source region of NEAs is the inner main belt (IMB) located between the ν6 secular resonance and the 3:1 mean motion resonance with Jupiter (2.15 - 2.5 AU). These resonances perturb the orbits of main belt asteroids enough to potentially send them into Earth orbit while preserving their inclinations. At least eight primitive asteroid families have been identified in the inner belt based on dynamics and albedo: Polana-Eulalia (hereafter Polana), Erigone, Clarissa, Chaldaea, Klio, Sulamitis, Svea, and Chimaera (Nesvorny´ et al. 2015, Walsh et al. 2013) (Fig. 6.1). For more details on how families are identified, we refer the reader to Nesvorny´ et al. (2015). 120 Figure 6.1: Semimajor axis vs. inclination (left) and eccentricity (right) of the primitive inner belt asteroid families. Red circles indicate objects in the primordial background family (PBF) identified by Delbo´ et al. (2017). Objects in this study are highlighted with red diamonds for objects with NIR spectra, blue diamonds for objects with visible spectra, and purple diamonds for objects with both visible and NIR spectra. A map of the IMB asteroid population with the family members removed leaves a population of background objects that seem to cluster into halos around the location of the families (Fig. 6.2). This suggests that the families are the source of the background population. Dermott et al. (2018) show that the sizes of background asteroids are correlated with eccentricity and anti-correlated with inclination, which suggests that both family and non-family asteroids originated from a small number of large primordial planetesimals. If this is true, then the background objects should be spectrally similar to the family objects in the same region. This hypothesis is supported by studies of near-infrared colors using data within the MOVIS catalog by Morate et al. (2018b). 121 Figure 6.2: The inner belt before and after all families (pink) have been removed, leaving only the background population behind (black). Asteroids cluster around the family locations, forming ha- los. The dynamical families extend beyond the specific dynamical boundaries chosen by Nesvorny´ et al. (2015) and the families are potentially the main source of the background population. There are also regions of the IMB far from the families and family halos, such as the region bor- dering the ν6 resonance. Background objects near this resonance are far enough from the families that they are probably not members but are also closer to the ν6 and therefore more likely to be perturbed into NEA space. Campins et al. (2013, 2018) propose that the background population may be as important as the families as sources of NEAs. It is possible that the background objects may have belonged to a larger primordial family that has been erased over time (e.g., Fieber-Beyer & Gaffey 2020). In a plot of semi-major axis vs 1/diame- ter, a dynamical family will plot in a V shape. Smaller objects are more affected by the Yarkovsky effect than larger objects, which creates the V shape, or Yarkovsky cone. The Yarkovsky cone represents the farthest distance a family member can drift in semimajor axis due to the Yarkovsky 122 effect. The Yarkovsky force occurs at a rate of da/dt, and so the slope of the V shape is indicative of the age of the family, or the time when the parent body initial fragmented. It is difficult to iden- tify ancient families older than 2.5 Gyr using standard clustering methods (e.g., Nesvorny´ et al. 2015) because objects have drifted in semimajor axis until the Yarkovsky cone eventually becomes blended with the background population. Furthermore, Brasil et al. (2016) show that planetary perturbations can also disperse primordial families beyond recognition by drastically affecting the inclination of family members. Delbo´ et al. (2017) searched for V-shaped groups in (semi-major axis vs 1/diameter) space of low-albedo background asteroids and identified a primordial family of dark asteroids in the IMB. Deienno et al. (2021) use multiple V-shape searching techniques to confirm the presence of this ancient family. This proposed primordial background family (PBF) is ∼4 billion years old which is significantly older than the oldest IMB family and spans an inclination range of 12°. This inclination range indicates that the PBF may have formed before the giant planet migration (Delbo´ et al. 2017), so trends in the spectra of these PBF objects with their orbital and physical properties will enhance our knowledge of the asteroid belt before and during the mixing event caused by the motion of the giant planets. Delbo´ et al. (2017) propose that the majority of what were known as primitive background asteroids are actually part of the PBF. We therefore can assume that the PBF is indicative of the entire population of primitive background asteroids, though there are still some IMB asteroids that are not connected with the PBF or any other family. From here on, when we refer to “background” asteroids, we mean the population of primitive IMB asteroids that are not affiliated with any of the dynamical families defined by Nesvorny´ et al. (2015). We use the PBF to study two important processes with one group. We can compare the PBF to itself to learn about the evolution of an ancient asteroid family and we can compare the individual objects in the PBF to the primitive families to learn about the relationship between the families and the primitive background. 123 The goal of this study is to spectrally characterize the primordial background population in the inner belt and to compare with the population of primitive IMB families. This will allow us to explore the probability of NEAs coming from the background, provide evidence for the hypothesis that asteroids were born big (i.e., Morbidelli et al. 2009), and map the distribution of hydration throughout the IMB. Details of our observations and data reduction are given in Section 6.2. In Section 6.3 we analyze each spectrum for taxonomy, slope, curvature, and the presence of the 0.7 µm band, and then we compare the entire sample with other IMB families, Bennu and Ryugu, me- teorites and space weathering experiments. We discuss the implications of our findings in Section 6.4 and present our conclusions in Section 6.5. 6.2 Observations and data reduction This work focuses on objects identified in the PBF by Delbo´ et al. (2017). We only used the asteroids marked as “certain” members in the Data S1 text file supplementary material. These 108 asteroids all have D<50 km and are in the section of space between the inward border of the primitive family V-shape and that of Polana. In addition to the 108 certain members, Delbo´ et al. (2017) identify 8 objects within the same region of interest with D≥50 km which they posit could be family members or small planetesimals. They also identify 5,121 objects that are outside of the V-shape of the family and do not belong to any families. For simplicity, we only looked at objects that were marked by Delbo´ et al. as certain members. We present new near-infrared (NIR; 0.8-2.5 µm) spectra of 55 objects and visible spectra of 21 objects from other surveys. There are 17 objects that overlap in both samples. 124 6.2.1 Visible We searched for published visible spectra of family objects using the Small Bodies Data Ferret sup- ported by the NASA Planetary System1. We found 15 spectra from the Small Main-belt Asteroid Spectroscopic Survey (SMASS) and 5 spectra from the Small Solar System Objects Spectroscopic Survey (S3OS2) (Bus & Binzel 2002, Lazzaro et al. 2004). We obtained one more spectrum from a study by Matlovicˇ et al. (2020) for a total of 21 spectra. The objects are listed in Table 6.1. Orbital elements (a,e,i) and absolute magnitude (H) were extracted from the Minor Planet Center Orbit Database file (MPCORB.DAT4). Diameter and albedo were obtained from NEOWISE (Mainzer et al. 2019). Also included are our calculated values for spectral slope, depth and center of the 0.7 µm absorption band, and Bus taxonomic classification (discussed more in Section 6.3). Four of the objects (1806, 2575, 4422, and 5524) were determined to be non-primitive types based on their albedo and spectral shape. Delbo´ et al. (2017) identifies these as being members of the Flora family, and so we excluded them from the rest of the analysis, even though one of them, (1806), is dark enough to be a primitive asteroid. There are other members of the PBF that are also members of other families, indicated by an a or b in Table 6.1, but we did not exclude these because they had spectra consistent with primitive asteroids. Fig. 6.3 shows the individual spectrum for each object, normalized to unity at 0.55 µm. 1http://sbn.psi.edu/ferret/ 125 Table 6.1: Details of the visible spectra obtained. A + indicates that the NIR spectrum for the object is also presented in this paper. A * indicates that the object is not primitive and is therefore excluded from the rest of the analysis. An a indicates that the object was identified as being a member of the Flora family, and b indicates that the object was identified as being a member of the Chaldaea family. $ At the time of publication of Delbo´ et al. (2017) all PBF members had pV <0.12 however this work uses a newer source for albedo which is why some are higher. $ 0 0 Number Survey a [AU] e i [°] H D [km] pV S [%/1000A˚ ] ErrS hydrated depth [%] center [µm] Visible Taxonomy 220 + S3OS2 2.349 0.257 7.59 11.2 30.0 0.06 1.543 0.002 yes 1.223 0.692 Cb 249 + S3OS2 2.377 0.218 9.62 11.3 36.6 0.04 2.066 0.005 yes 1.552 0.673 Cgh 282 + SMASS 2.340 0.080 9.03 11.0 37.9 0.04 -2.251 0.011 B 783 SMASS 2.342 0.230 9.34 11.1 41.7 0.03 0.873 0.026 Cgh 853 + SMASS 2.312 0.106 9.22 11.6 22.0 0.06 -1.024 0.014 yes 1.704 0.716 Ch 917 + S3OS2 2.382 0.200 5.13 11.6 26.6 0.05 1.591 0.003 Cb 1244 + S3OS2 2.343 0.099 8.70 11.5 30.4 0.04 3.273 0.001 X 1705 + SMASS 2.299 0.245 7.71 13.3 9.3 0.10 -1.671 0.017 B 1806 +*a S3OS2 2.236 0.107 3.84 12.0 8.0 0.04 N/A N/A Sv 2259a SMASS 2.294 0.186 4.68 12.6 19.9 0.05 1.375 0.011 yes 2.677 0.663 Cgh 2328 + SMASS 2.342 0.146 10.01 13.0 12.7 0.07 -0.373 0.021 Cg 2503 SMASS 2.192 0.213 7.11 14.4 5.6 0.12 -1.291 0.023 B 2575 +*a SMASS 2.240 0.123 4.67 12.7 6.4 0.39 N/A N/A Sa 2772 + SMASS 2.314 0.205 9.79 14.0 9.8 0.05 -1.299 0.018 B 2778 + SMASS 2.281 0.122 4.61 13.2 9.2 0.12 -0.843 0.039 B 3684 SMASS 2.287 0.154 6.81 13.9 7.6 0.09 -0.500 0.023 Cg 4422 +*a SMASS 2.238 0.179 4.80 12.6 6.3 0.12 N/A N/A S 4750 +a SMASS 2.184 0.091 4.91 13.9 10.7 0.04 1.616 0.025 Xk 5081 + SMASS 2.319 0.114 13.20 12.6 16.5 0.06 -1.838 0.032 yes 2.982 0.715 B 5333 +b SMASS 2.345 0.168 10.97 13.1 11.5 0.08 -1.249 0.030 yes 3.697 0.665 B 5524 +* Matlovicˇ et al. (2020) 2.367 0.028 7.48 13.0 19.9 0.03 N/A N/A V 126 127 Figure 6.3: Individual visible spectra of PBF members. Objects with blue labels are from SMASS and those with green labels are from S3OS2. The object with an orange label is from Matlovicˇ et al. (2020). Pink lines show the computed spectral slope between 0.55-0.9 µm. All spectra are normalized to unity at 0.55 µm. To further explore the PBF in the visible wavelength range, Fig.6.4 shows the visible slope dis- tribution of the IMB objects with pV <0.12 from SDSS (Carvano et al. 2010, Hasselmann et al. 2011). For details of how the slopes are calculated, see De Pra´ et al. (2020). The first column includes all objects, the second is primitive families only, the third is the PBF, and the fourth is background objects only. In the plot of the entire IMB, the blue objects cluster at low inclinations while the red objects are more prevalent at the higher inclination. The plot of the IMB families shows that this is mostly due to the overwhelming amount of B-types in the Polana family. When 128 we remove the families and plot the background only, the colors are much more well mixed, though the end points (sin(i)>0.2 and sini(i)<0.05) are majority red and blue, respectively. The slopes of the PBF objects are also well mixed; the B-types don’t cluster near each other and are spread out in inclination. The visualization of these colors shows that the PBF is an accurate sample of the entire background, and doesn’t appear to be influenced by the families. Figure 6.4: SDSS gri slope distribution of objects in the IMB with pV <0.12. The columns from left to right are (1) all IMB asteroids, (2) primitive families only, (3) PBF only, and (4) background asteroids only. Objects with blue slopes are colored with blue dots and objects with red slopes are colored with red dots. 129 6.2.2 NIR For our NIR observations we primarily used the NASA Infrared Telescope Facility (IRTF) at the Mauna Kea Observatory. We also used the Telescopio Nazionale Galileo (TNG) at the El Roque de Los Muchachos Observatory in the island of La Palma (Canary Islands, Spain). We prioritized observing objects with visible spectra published by SMASS or S3OS2 (Section 6.2.1). In total we observed 55 objects. Table 6.2 lists the orbital properties, absolute magnitude, diameter and albedo as well as our calculated values for Bus-DeMeo taxonomy, spectral slope, and curvature (discussed more in Section 6.3). Observing details and reduction method vary based on telescope used and are described below. Table 6.3 gives observational details for each object including number, date of observation, telescope, phase angle, apparent V magnitude, exposure time, and solar analog star. Delbo´ et al. (2017) identify some members of this sample as being members of other IMB families, indicated by a letter subscript in Table 6.2. We include the objects with spectra consistent with primitive asteroids in our sample, despite them being listed as members of other families. 130 Table 6.2: Details of the NIR spectra obtained. A + indicates that the visible spectrum for the object is also presented in this paper. A * indicates that the object is not primitive and is therefore excluded from the rest of the analysis. Also included are our calculated values for slope, curvature, and Bus-DeMeo Taxonomy. Details of these calculations are given in Section 6.3. An a, b, c, d, or e superscript indicates that the object was identified as being a member of the Flora, Chaldaea, Erigone, Nysa, or Vesta family, respectively. $ At the time of publication of Delbo´ et al. (2017) all PBF members had pV <0.12 however this work uses a newer source for albedo which is why some are higher. $ 0 0 2 Number a [AU] e i [°] H D [km] errD pV errpV NIR Tax S [%/1000A˚ ] errS c [%/µm ] errc (220) Stephania + 2.349 0.257 7.59 11.2 30.00 10.67 0.06 0.08 B -3.555 0.720 0.332 0.176 (249) Ilse + 2.377 0.218 9.62 11.3 36.64 13.61 0.04 0.03 Xk 1.400 0.800 -0.491 0.269 (282) Clorinde + 2.340 0.080 9.03 10.9 37.87 12.07 0.04 0.03 C 1.174 0.500 0.320 0.071 (370) Modestia 2.324 0.091 7.86 10.7 32.45 7.64 0.11 0.05 Xk 1.392 0.315 0.206 0.003 (689) Zita 2.316 0.229 5.74 12.2 9.10 1.68 0.27 0.10 Cg 1.251 1.050 0.218 0.119 (853) Nansenia + 2.313 0.106 9.22 11.7 22.04 5.85 0.06 0.09 B -0.187 0.156 -0.024 0.021 (916) America 2.365 0.237 11.09 11.5 34.20 12.07 0.04 0.05 Cb 1.441 0.120 -0.263 0.216 (917) Lyka + 2.381 0.200 5.13 11.6 26.62 8.14 0.05 0.03 Cb 1.341 0.010 -0.485 0.651 (933) Susi c 2.368 0.165 5.54 12.4 21.93 6.34 0.04 0.02 Xe 1.559 1.050 0.460 0.268 B -0.636 0.092 0.261 0.021 (1159) Granada 2.379 0.059 13.03 11.6 28.69 9.87 0.03 0.04 Xk 1.148 0.260 -0.098 0.012 (1216) Askania 2.232 0.179 7.60 13.5 9.62 2.53 0.16 0.09 Sq 2.148 0.447 0.309 0.074 (1244) Deira + 2.343 0.099 8.70 11.5 30.43 9.14 0.04 0.03 Xk 1.404 0.305 0.746 0.006 (1544) Vinterhansenia 2.373 0.105 3.33 11.9 22.42 6.88 0.06 0.03 N.A. 7.290 0.447 -1.394 0.025 (1700) Zvezdara 2.361 0.226 4.51 12.5 18.46 4.90 0.06 0.04 Cg 0.895 0.447 0.560 0.198 (1705) Tapio + 2.299 0.245 7.71 13.2 9.28 2.49 0.10 0.11 Cg 0.927 0.312 0.344 0.013 (1806) Derice +* a 2.236 0.107 3.84 12.0 7.98 0.76 0.04 0.15 Sv 2.196 0.493 0.697 0.280 (1924) Horus d 2.340 0.131 2.73 13.5 9.06 1.98 0.11 0.04 C 1.634 0.312 0.202 0.007 (2171) Kiev *a 2.255 0.166 7.52 12.7 6.45 0.31 0.39 0.05 Sq 2.176 0.557 0.998 0.146 (2259) Sofievka a 2.294 0.186 4.68 12.6 19.89 6.12 0.05 0.04 B -0.420 0.460 -0.119 0.016 (2322) Kitt 2.293 0.041 2.40 13.1 11.42 3.31 0.10 0.05 Cgh 0.894 0.493 0.070 0.092 (2328) Robeson + 2.342 0.146 10.01 13.0 12.66 3.32 0.07 0.04 Ch 0.910 0.213 -0.358 0.039 (2536) Kozyrev * 2.307 0.228 4.87 12.5 7.97 0.33 0.31 0.04 S 1.445 0.447 0.584 0.051 Sa 2.610 0.312 0.386 0.006 (2575) Bulgaria +*a 2.240 0.123 4.67 12.7 6.41 0.29 0.39 0.06 Sq 2.926 0.170 0.573 0.036 (2772) Dugan + 2.314 0.205 9.79 14.0 9.76 3.13 0.05 0.03 C 1.012 0.213 1.349 0.010 (2773) Brooks 2.328 0.143 3.67 13.1 10.81 2.57 0.07 0.03 C 1.360 0.557 0.309 0.040 (2778) Tangshan + 2.281 0.122 4.61 13.1 9.16 2.14 0.12 0.07 D 3.711 0.557 0.764 0.219 (2839) Annette *a 2.216 0.150 4.81 12.9 5.42 0.73 0.45 0.21 Sq 2.002 0.220 0.478 0.136 (3633) Mira a 2.312 0.102 3.31 13.6 9.14 2.97 0.05 0.04 Cg 0.797 0.447 0.604 0.125 131 $ 0 0 2 Number a [AU] e i [°] H D [km] errD pV errpV NIR Tax S [%/1000A˚ ] errS c [%/µm ] errc (3723) Voznesenskij 2.255 0.111 1.38 14.0 9.12 2.54 0.05 0.03 Ch 0.868 0.081 0.948 0.172 (4231) Fireman 2.260 0.066 8.59 13.3 N.A. N.A. N.A. N.A. Ch 0.938 0.493 -0.487 0.113 (4422) Jarre +*a 2.238 0.180 4.80 12.6 6.34 0.81 0.12 0.11 S 0.493 0.156 -0.040 0.052 D 4.016 0.447 -0.107 0.100 (4524) Barklajdetolli e 2.320 0.133 7.29 12.9 9.91 2.84 0.18 0.09 D 4.199 0.213 -0.560 0.203 (4750) Mukai a 2.184 0.090 4.90 13.9 10.65 2.89 0.04 0.03 X 1.882 0.180 -0.869 0.632 (5081) Sanguin + 2.319 0.114 13.20 12.6 16.45 4.09 0.06 0.03 B -0.104 1.050 0.086 0.410 (5333) Kanaya +b 2.345 0.168 10.97 13.1 11.53 3.01 0.08 0.09 Ch 0.597 0.161 -0.285 0.009 (5524) Lecacheux + * 2.366 0.027 7.49 12.9 19.90 12.77 0.03 0.10 V 0.652 0.447 -1.169 0.014 (5924) Teruo d 2.346 0.109 4.09 13.3 11.37 3.27 0.07 0.03 C 0.766 0.600 0.129 0.155 (6125) Singto * 2.223 0.136 1.31 13.8 4.42 0.93 0.29 0.15 Sa 4.532 0.557 1.301 0.285 C 1.321 0.188 1.151 0.003 (6542) Jacquescousteau 2.302 0.120 3.74 14.0 7.70 2.57 0.05 0.04 C 1.858 0.180 0.126 0.572 (6647) Josse 2.206 0.196 1.14 14.9 5.66 1.44 0.03 0.02 Ch 0.449 1.299 -0.355 0.000 (7132) Casulli 2.309 0.209 5.49 13.7 9.02 0.06 0.09 0.00 Cg 0.838 0.120 0.170 0.263 (9723) Binyang * 2.148 0.118 4.98 14.3 3.69 0.10 0.12 0.03 V 0.891 0.444 -0.116 0.007 (10446) Siegbahn 2.223 0.115 6.72 14.4 7.40 1.96 0.07 0.07 D 3.874 0.444 1.010 0.129 (10542) Ruckers 2.220 0.085 3.30 14.4 6.48 1.72 0.08 0.08 C 1.676 0.557 0.534 0.155 (14179) Skinner 2.288 0.213 8.30 14.2 7.74 1.83 0.06 0.03 T 2.903 0.220 -0.326 0.078 (15415) Rika * 2.202 0.228 7.48 14.2 2.83 0.49 0.61 0.15 V 1.038 0.305 -0.742 0.059 (15998) 1999 AG2 2.207 0.044 4.26 14.2 6.52 1.79 0.15 0.06 N.A. 10.367 0.444 0.496 0.247 (20771) 2000 QY150 2.307 0.047 11.59 13.9 8.75 2.38 0.03 0.02 Sa 4.791 0.444 0.670 0.718 (28620) 2000 FE26 2.276 0.110 4.63 14.4 7.35 2.32 0.05 0.04 Ch 0.949 0.220 0.620 0.188 (28736) 2000 GE133 2.281 0.068 9.74 13.9 7.27 2.54 0.07 0.07 Ch 0.716 0.188 0.185 0.106 (30514) Chiomento 2.223 0.184 3.20 15.5 5.98 1.75 0.03 0.03 Ch 1.135 0.188 0.056 0.097 (31487) Parthchopra a 2.231 0.107 6.34 14.9 5.08 1.44 0.07 0.08 D 3.486 0.557 1.156 0.158 (32898) 1994 PS1 a 2.211 0.223 5.47 15.5 4.29 1.04 0.07 0.04 Cg 1.132 0.194 -0.221 0.203 (48876) 1998 HE103 a 2.202 0.182 6.64 15.0 4.77 0.11 0.10 0.03 Cg 0.929 0.188 0.412 0.138 (49863) 1999 XK104 2.177 0.169 2.05 16.0 2.69 0.76 0.10 0.07 Cg 0.948 0.170 0.614 0.605 132 Table 6.3: Observational details of NIR observations. The number in the SA column corresponds to the solar analog star used, given in Table 6.4. Number Date UT start Telescope Airmass alpha V Texp [s] SA 220 20180623 22:16 TNG 2.21 2.2 12.9 60 2,6 249 20181214 4:42 TNG 1.22 22.2 16.1 540 4,5 282 20171225 7:02 TNG 1.12 25.5 15.3 720 4 370 20190424 5:13 IRTF 1.08 23.4 15.7 720 5 689 20180622 10:21 IRTF 1.28 12.4 14.7 720 2,3 853 20180504 10:27 IRTF 1.12 9.7 14.1 960 2 916 20181214 5:40 TNG 1.22 22.2 16.8 720 4,5 917 20180816 14:10 IRTF 1.02 26.2 14.5 960 1 917 20180628 4:18 TNG 1.52 32.2 15.4 720 1,2 933 20180622 5:20 IRTF 1.27 26.1 16.9 960 2,3 1159 20181031 5:35 TNG 1.04 23.1 16.3 2160 4 1216 20190330 9:55 IRTF 1.06 6.6 16 720 5,6 1244 20191015 9:00 IRTF 1.04 11.3 15.4 1200 1 1544 20190330 6:07 IRTF 1.24 26 16.4 1200 5,6 1700 20190330 12:25 IRTF 1.34 13.2 17.2 2160 5,6 1705 20181215 1:49 TNG 1.1 6 15.9 2880 4,5 1806 20190716 15:15 IRTF 1.08 26.5 16.7 720 1 1924 20181215 6:02 TNG 1.11 18.1 16.5 2160 4,5 2171 20190322 6:56 IRTF 1.17 22.3 17.7 2400 5 2259 20171225 22:03 TNG 1.11 6.2 16.3 2520 4,7 2322 20190716 5:44 IRTF 1.38 18.2 16.8 2160 1 2328 20190615 9:35 IRTF 1.55 13.8 16.7 2160 6 133 Number Date UT start Telescope Airmass alpha V Texp [s] SA 2536 20190330 13:39 IRTF 1.43 17.7 17.2 2160 5,6 2575 20181101 3:43 TNG 1.16 22.4 17.3 2160 1,2,7 2772 20190615 5:53 IRTF 1.05 24.5 18.8 4320 6 2773 20190322 5:15 IRTF 1.02 26.3 16.4 2400 5 2778 20190322 12:00 IRTF 1.24 16.1 17.2 2160 5 2839 20180812 13:33 IRTF 1.16 16.8 15.7 720 1 3633 20190330 11:34 IRTF 1.28 13.9 16.7 1440 5,6 3723 20200202 15:17 IRTF 1.12 23.6 17.9 2640 5 4231 20190716 8:10 IRTF 1.34 18.4 16.6 2160 1 4422 20180504 9:41 IRTF 1.17 10.6 16.4 1200 2 4524 20190615 8:36 IRTF 1.4 23.3 17.1 960 6 4750 20180628 3:10 TNG 1.15 6.5 16.9 2160 1,2 5081 20180622 14:28 IRTF 1.36 26.3 16.7 1440 2,3 5333 20170204 3:56 TNG 1.38 23.7 16 2160 4,6 5333 20180901 6:15 IRTF 1.24 9.7 17.2 2160 1,2 5524 20190330 9:13 IRTF 1 10.4 15.9 720 5,6 5924 20130709 13:34 IRTF 1.22 15.1 17.1 3360 2 6542 20180901 12:50 IRTF 1.12 8.7 17.5 2160 1,2 6542 20181031 20:11 TNG 1.34 17.9 18 3240 1,2,7 6647 20180816 13:34 IRTF 1.49 5.3 16.2 1200 1 134 Number Date UT start Telescope Airmass alpha V Texp [s] SA 7132 20181101 2:08 TNG 1.08 3 17.4 2700 1,2,7 9723 20190424 9:23 IRTF 1.32 4.3 16.9 1200 5 10446 20190422 10:56 IRTF 1.19 1.4 17.4 3360 5 10542 20190322 14:46 IRTF 1.33 23.9 17.7 2400 5 14179 20180812 12:07 IRTF 1.07 11.1 16.6 1440 1 15415 20191015 9:41 IRTF 1.18 19.7 15.9 1200 1 15998 20190422 8:02 IRTF 1.03 20.2 17.6 2400 5 20771 20190422 5:22 IRTF 1.05 26.9 18.4 3120 5 28620 20180812 11:39 IRTF 1.59 5.9 16.5 1200 1 28736 20180901 7:58 IRTF 1.4 6.7 17.1 2160 1,2 30514 20180901 9:01 IRTF 1.21 2.3 17 2160 1,2 31487 20190322 9:30 IRTF 1.02 14.5 17.3 2160 5 32898 20170929 5:46 IRTF 1.15 26.5 17.5 2160 2 48876 20180901 10:15 IRTF 1.28 6.3 17 2160 1,2 49863 20181101 0:37 TNG 1.11 11.8 17.9 2880 1,2,7 135 Table 6.4: Solar analog stars used. ID Star RA Dec V 1 SA 115-271 23 42 41.8 +00 45 14 9.7 2 SA 112-133 20 43 11.8 +00 26 15 10.0 3 SA 107-684 15 37 18.1 -00 09 50 8.4 4 SA 98-978 06 51 34.0 -00 11 28 10.5 5 SA 102-1081 10 57 04.0 +00 13 10 9.9 6 SA 107-998 15 38 16.0 +00 15 24 10.5 7 SA 93-101 01 53 19.0 +00 22 22 9.7 6.2.2.1 IRTF With IRTF we observed 41 asteroids over 15 nights between July 2013 and October 2019. We used the SpeX spectrograph (Rayner et al. 2003) in PRISM mode which covers a wavelength range of 0.7-2.52 µm. We used the 0.8×1500slit oriented at the parallactic angle for all nights except 20180816 where we used the 0.8×6000slit. The observations made use of a beam switching pattern of AB separated by 7.500(3000for 20180816). The exposure time and the number of exposures varies based on magnitude. For each object we observed at least one local standard star at a similar airmass to correct for the solar spectrum and for telluric absorptions. In some cases, we observed multiple standard stars in which case we use all of them and average the different results to find a final relative reflectance. For each night we also observed a solar analog star at several airmasses (Table 6.4). We reduced the data using SpeXtool (Cushing et al. 2004) and removed the telluric contribution using the ATRAN atmospheric model (Lord 1992). The spectra are shown in Fig. 6.5 normalized to unity at 1.0 µm. 136 6.2.2.2 TNG With TNG we observed 15 asteroids over 9 nights between February 2017 and December 2018. We used the AMICI prism and the 1.500slit on the NIR instrument, NICS, which covers a wavelength range of 0.8-2.4 µm with a resolving power of ∼35 quasi-constant along the spectrum (Oliva 2000). We followed an ABBA beam pattern separated by 1000. The exposure time and number of exposures vary based on magnitude and sky conditions (Table 6.3). All observations were performed in parallactic angle and the tracking of telescope was set to the asteroid proper motion. To obtain relative reflectance and to correct for telluric absorption, we observed at least one solar analog star per night. Reduction of the data was done using standard IRAF packages. Emission lines from Argon and Xenon lamps are not resolvable using the AMICI prism’s low resolution, and so a look-up table based on theoretical dispersion predicted by ray tracing was used for wavelength calibration. The Python script used to calibrate wavelengths and to divide by the star spectrum was developed by our group and available upon request (details in Pinilla-Alonso et al. 2016). The final spectra are shown in Fig. 6.5 normalized to unity at 1.0 µm. 137 138 139 140 Figure 6.5: Individual NIR spectra for each object. A blue spectrum indicates the object was observed by TNG and green indicates IRTF. The spectral slope S0 between 0.95-2.3 µm is shown by a pink line. All spectra have been normalized to unity at 1 µm. 141 6.3 Results and analysis 6.3.1 Characteristics of the sample We present 21 visible and 55 NIR spectra of objects from the PBF proposed by Delbo´ et al. (2017). Seventeen objects in our sample have both visible and NIR spectra presented in this paper and are denoted by a + in Tables 6.1 and 6.2. Diameters of members range from 3 to 41 km, which is expected as all of the members of the PBF have D<50 km. The average diameter is 13 km which is similar to that of the families in the IMB (Polana average 16 km, Klio average 8 km) however, the combined size of the 10 largest objects in the entire PBF are larger than the combined size of the 10 largest objects from either the Polana or Erigone families. The mean albedo of the sample is 0.08 ± 0.08. There is a range of taxonomies from B-types to D-types including 12 non-primitive interlopers, which were identified by their spectral shape and/or high albedo. These interlopers are the result of the PBF being significantly blended with the background population and are to be expected. We denote each by a * in the Table 6.1 and Table 6.2 and we exclude them from the rest of the analysis. In the past, it was assumed that families would look spectrally homogeneous because it was be- lieved that the parent body was not differentiated. However, we have recently shown that not all families are homogeneous (e.g., Morate et al. 2019) which has interesting implications about how asteroids break up and reaccumulate during collisions. Furthermore, the ancient age of this family means that the objects have had a long time to be altered by effects such as space weathering, which can cause changes to the asteroid surfaces over time. Because all members of the family have been exposed to space weathering for the same amount of time, any differences within the family will most likely be due to differences in initial composition, which will tell us about the heterogeneity of the primordial parent body. 142 Two of our NIR sample objects were observed by both the IRTF and the TNG telescopes on different nights, and one of those (1159) shows variations in spectral slope over time. The two spectra of 1159 were observed 2 months apart. The rotation period of 1159 is 31 hours (Binzel 1987), however no error in this number is published so we are not able to calculate the phase angle of the two observations. The slope of the first spectrum is blue (-0.636 ± 0.092 %/1000A˚ ), and the slope of the second spectrum is red (1.148 ± 0.260 %/1000A˚ ). The diameter of 1159 is 34 km which is large enough to have surface heterogeneity. The rotational inhomogeneity of this object is supported by the recent findings of exogenous material on Bennu and Ryugu (Dellagiustina et al. 2019, Tatsumi et al. 2019). For instance, Binzel et al. (2015) show spectral slope variation in ground-based spectra of Bennu similar to those we present here. The asteroids in our sample are much larger than Bennu and Ryugu, so it is logical that they are even more likely to be composed of more than one material. 6.3.2 Taxonomy determination We classified each visible and NIR spectrum using the M4AST online tool2 (Popescu et al. 2012). We used the option for the Bus-DeMeo taxonomy scheme (DeMeo et al. 2009) and looked for the match with the lowest χ2 residual. When it was not clear which spectral type was the best fit, we used our best judgement. The results are shown in Table 6.1 and Table 6.2. Our analysis revealed 8 C-types, 7 B-types, 2 X-types, 3 S-types, and 1 V-type in the visible sample and 26 C-types, 2http://m4ast.imcce.fr/ 143 5 B-type, 6 X-types, 5 T/D types and 12 non-primitive (S,V) types in our NIR sample. Two of the objects, 1544 and 15998, did not have any taxonomy that matched well, but are still primitive based on their albedos. As stated before, the non-primitive types are excluded from the rest of the analysis. Of the 12 non-primitive types, 5 are classified into the Flora family by Nesvorny´ et al. (2015). The remaining seven objects (3 V-types and 4 S-types) are true interlopers. Figure 6.6 shows the taxonomies of the objects observed in the visible in the PBF compared to the taxonomies of the visible spectra of other IMB families (de Leon´ et al. 2016, Morate et al. 2016, 2018a, 2019). It is interesting to note that for each of the families there is a clear majority of taxonomic types (i.e., mostly B-types, mostly C-types) however the background has roughly the same number of B-types and C-types. In Figure 6.7, we show that the PBF B-types are not near the location of the predominately B-type families, ergo they are different populations. Figure 6.6: Taxonomic distribution of the PBF visible spectra compared to the other IMB families. The PBF is mostly B-types and C-types and has one of the highest abundances of B-types. Other families’ visible spectra taxonomic distribution are from de Leon´ et al. (2016), Morate et al. (2016, 2018a, 2019). 144 Figure 6.7: Similar to Figure 6.1, but the IMB families are now colored green if they are predom- inately C-types and blue if they are predominately B-types. The location of different taxonomic types in the PBF are shown by colored diamonds. The PBF B types do not cluster around the B-type families and the taxonomic types look well mixed (i.e., no pattern or clustering). 6.3.3 Spectral slopes Primitive spectra are generally featureless, but may display some amount of curvature as well as a gradient in slope. We calculate both of these parameters and their associated errors to compare with the spectra of the families. Details of the calculations for spectral slope and curvature are given in Arredondo et al. (2020) but are summarized here. The spectral slope (S0) of each asteroid is calculated in the range from 0.55 to 0.9 µm for visible and 0.95 to 2.3 µm for NIR using the expression dS/dλ S0 = SN 145 where dS/dλ is the rate of change of the reflectance and SN is the reflectance at 0.55 µm for visible and 1.0 µm for NIR. S0 is measured in units of %/1000A˚ . The uncertainty in calculated NIR slopes is dominated by the systematic error caused by using different solar analogs. The systematic error for each night is given in Table 6.2. For the visible spectra, the error in the slope is the standard deviation of the calculation. The visible slope distribution appears to be bimodal with 7 red objects and 10 blue objects, however the gap between the 2 groups is just 0.5 %/1000A˚ which is the same order of magnitude of most of our slope uncertainties. We do not see any hint of bimodality in the sample of NIR spectra, possibly because NIR spectra are less altered by space weathering effects (Section 6.3.9). The NIR spectra are overwhelmingly red sloped, and the objects that are blue in the visible are not always also blue in the NIR, in agreement with what was found in de Leon´ et al. (2012). We did not find any correlation with slope and orbital parameters or properties of the 0.7 µm absorption band. As expected based on the range of spectral types, the PBF is diverse; containing objects with red slopes and objects with blue slopes (Fig. 6.8). The visible spectral slopes of the sample range between -2.251 ± 0.011 %/1000A˚ to 3.273 ± 0.001 %/1000A˚ with a mean slope of 0.00 ± 1.60 %/1000A˚ . The NIR spectral slopes show even more of a variation ranging between blue (-3.555 ± 0.720 %/1000A˚ ) to red (4.199 ± 0.213 %/1000A˚ ) and the mean slope is 1.155 ± 1.265 %/1000A˚ . As shown in Fig. 6.9, the distribution of blue and red sloped objects is well mixed in eccentricity and orbital space. It could be said that the blue objects are closer to lower semimajor axis and vice- versa, however this could be due to small number statistics. We note that until now, PRIMASS has not detected diversity in the NIR; 95% of the Sulamitis and Polana samples had slopes between 0 and 2 %/1000A˚ despite being very different in the visible (Arredondo et al. 2021b). In contrast, only about 75% of our PBF sample show the same range of slopes, so the diversity seen in the NIR from the background objects is a unique feature of this population. 146 Figure 6.8: All visible (left) and NIR (right) spectra of the PBF sample (grey lines). The green, blue, orange and red lines are the average spectra of the C-types, B-types, X-types, and T/D-types in our sample, respectively. In both wavelength ranges we see diversity, i.e., a wide spread of both blue slopes and red slopes. 147 Figure 6.9: Similar to Figure 6.1 but showing the distribution of blue and red visible sloped PBF objects (marked by blue and red diamonds, respectively) in semimajor axis vs inclination (left) and eccentricity (right) space. Objects with similar slope do not seem to cluster together, suggesting that there were multiple progenitor bodies or a major mixing event. 6.3.4 Curvature The curvature, c, of each spectrum is calculated following the approach of Ziffer et al. (2011) by averaging the second derivative of a polynomial fit to the spectrum in the region from 0.9-2.4 µm for the NIR spectra. The curvature of the spectra is a useful tool for distinguishing between families; e.g., the Chaldaea family’s concave curvature distinguishes it from all other IMB families in the NIR (Arredondo et al. 2021a). We did not calculate the curvature for the visible sample because none of the other IMB families have curvature calculated and so we have nothing to compare our values with. Objects with a c value less than zero are concave, while objects with a c value greater than zero are convex. The uncertainty in this process is calculated by doing the same analysis on the highest and lowest points of the error bars and then taking an average. 148 The calculated c values are given in Table 6.2. There are 11 concave, 21 convex, and 10 flat objects. The average curvature for the sample is 0.2 ± 0.5 %/µm2. The other families we have computed curvature for are the Klio (0.5 ± 0.1 %/µm2), Chaldaea (-0.3 ± 0.2 %/µm2), Polana (0.11 ± 0.02 %/µm2), and Sulamitis (0.24 ± 0.07%/µm2) families, which show the same general distribution of concave and convex shapes (except for Chaldaea which is overwhelmingly concave). The even distribution of concave and convex shapes in the PBF spectra is expected because the background objects are more diverse than the objects in the families. 6.3.5 Aqueous alteration Primitive asteroids are the transitional objects between rocky and icy bodies and are expected to contain hydrated silicate minerals (e.g., De Pra´ et al. 2018). While the reflectance spectra of prim- itive asteroids are generally featureless, a considerable number of IMB primitive asteroids show evidence of aqueous alteration through the presence of an absorption feature at 0.7 µm attributed to charge transfer transitions in oxidized iron (e.g., De Pra´ et al. 2018, and reference therein). The PBF spans the entire inner belt (Figure 6.1) and the presence of hydrated material in background objects can be used to map the distribution of hydrated material in the IMB. Understanding the distribution and abundance of hydration in the IMB can help explain how water may have been delivered to Earth. The 0.7 µm absorption band is present in a large fraction of the spectra of primitive IMB families (Morate et al. 2016, 2018a, 2019, ; Fig. 6.10). Visual inspection of our sample in Figure 6.3 shows that this feature is present in the spectra of the background population as well. We used the CANA toolkit (Codes for ANalysis of Asteroids; De Pra et al. 2018) to compute the center and depth of absorption bands in the visible spectra. CANA computes a fourth-order polynomial to the data and finds the continuum by fitting a straight line tangent to the polynomial fit at the 149 limits of the absorption band. This is divided out so that only the band is left (Fig. 6.11). The central wavelength of the absorption band is the local minimum of the absorption band fit. The band depth is the percent difference between the reflectance of the linear fit (1 after continuum removal) and the reflectance at the central wavelength. This process is repeated 1000 times for robustness and the mean of all the runs is the final center and depth. The error in center and depth is the corresponding 1 sigma standard deviations of the mean. An asteroid is considered hydrated if there is an absorption band near 0.7 µm with a band depth of more than 1%. The results of the calculations are seen in Table 6.1. Figure 6.10: Hydration amongst primitive IMB families adapted from Morate et al. (2016). Each bar shows the total number of objects observed by PRIMASS, and the blue represents the fraction of the whole sample (blue + green) that showed the 0.7 µm absorption band indicative of hydration. 150 Figure 6.11: Example of the output of the CANA program showing the band center and band depth calculation for an object. the continuum is removed by fitting a line tangent to the band ends and then dividing it out. The blue line is a polynomial fit to the data and the blue dot is the furthest point from the linear fit. The percent difference is the band depth. Of our 17 primitive objects with visible spectra, 6 show the 0.7 µm band (30% of the sample). The average band center is 0.69 ± 0.02 µm and the average band depth is 2.3 ± 0.9%. Fig. 6.12 shows the physical location of hydrated and non-hydrated objects. The hydrated objects do not cluster around themselves or around the hydrated families, indicating again that the sample is well-mixed. 151 Figure 6.12: Similar to Figure 6.1 but showing the location of hydrated PBF objects. Blue dia- monds are PBF members that showed the 0.7 µm feature and green diamonds are objects that did not. There does not appear to be a trend with hydration and location or vicinity to families. Fig. 6.13 shows the relationship between the center of the absorption band and the spectral slope. From our sample, we show that redder objects correspond to a shorter central wavelength. This could be due to a small sample size, but it could also be due to a process such as space weathering which we discuss in Section 6.3.9. 152 Figure 6.13: Possible trend between spectral slope and band center. Redder objects have a shorter central wavelength. 6.3.6 Comparison with the primitive IMB families PRIMASS has observed all 8 primitive families in the visible (de Leon´ et al. 2016, Morate et al. 2016, 2018a, 2019) and the Polana, Klio, Chaldaea, and Sulamitis families in the NIR (Arredondo et al. 2020, 2021a,b, Pinilla-Alonso et al. 2016). A summary of our previous results is shown in Table 6.5. In the lower inclination families, there is a trend of visible spectra with inclination (Polana and Clarissa are located at 3°and have homogeneous spectra and no evidence of aqueous alteration. Erigone and Sulamitis are located at 5°and have diverse spectra and evidence of aqueous alteration). Klio, Chaldaea, Chimaera and Svea do not follow this trend at all in the visible. In the NIR, we have found that all families look almost identical, except for the significant curvature of the Chaldaea family. We compare the spectra of the background objects to the families at similar inclinations to test for similar origins between the families and the background. 153 Table 6.5: Summary of PRIMASS findings in the IMB. Klio age is from Carruba & Nesvorny´ (2016), PBF age is from Delbo´ et al. (2017), and all other ages are from Bottke et al. (2015). All other information is from the respective PRIMASS publications. Family Inc [°] Age [Myr] 0.7 µm feature Visible Spectra NIR Spectra Polana 3 1400 ± 150 Absent Homogeneous Homogeneous Eulalia 3 830 ± 370/100 Absent Homogeneous Homogeneous Clarissa 3 ∼60 Absent Homogeneous Too faint for NIR Erigone 5 130 ± 30 Present (58%) Diverse ? Sulamitis 5 200 ± 40 Present (60%) Diverse Homogeneous Klio 9 960 ± 250 Present (23%) Diverse Homogeneous Chaldaea 12 ? Present (79%) Homogeneous Homogeneous Chimaera 14 ? Present (20%) Diverse ? Svea 16 ? Absent Homogeneous ? Present (30%) Diverse Diverse PBF 1-13 ∼4 billion (This work) (This work) (This work) Fig. 6.14 shows a box and whisker plot comparing the slope distributions of IMB families com- pared to the PBF. In the NIR, the average slopes are identical within the uncertainties for all fami- lies. The PBF has the largest range of spectral slopes, highlighting its diversity. In the visible, the average spectral slope of each family is more varied. The range of the PBF is smaller than most of the families’. 154 Figure 6.14: The visible (left) and NIR (right) slope distribution of each IMB family and the PBF. Each box represents the average and standard deviation of the slope distribution, while the solid orange line represents the median. The colored points represent individual slope measurements for each family member. The tails of the whiskers show the minimum and maximum values. We run a two-sample Kolmogorov-Smirnov test comparing the slope distributions to quantify the differences between families. The Kolmogorov-Smirnov statistic, Dm,n, is given by the equation Dm,n = sup |F (x)1,m − F (x)2,n| x 155 where F (x)1,m and F (x)2,n are the cumulative distribution functions of the two compared distri- butions. A critical value, Dcrit,m,n is defined as rm + n D = c(α) crit,m,n mn where c(α) is 1.36 for α=0.05 and m and n are the sizes of the two samples. If Dm,n>Dcrit,m,n, then the null hypothesis is rejected, meaning both samples originated in different distributions. We compare the slopes of IMB family objects with the PBF objects that are at similar inclinations (± 1°of the average in Table 6.5). There are no visible PBF spectra at Polana/Clarissa inclination and no NIR family spectra published from the Clarissa or Erigone families and so we do not compare those. We are also unable to compare any PBF spectra with Chimaera and Svea because the PBF does not extend to high enough inclinations. The number of objects in our sample at each inclination along with the results of our KS tests are given in Table 6.6. In all cases except Polana we cannot reject the null hypothesis that the PBF and the family members come from similar distributions. We note that our sample size may be too small for these tests to be conclusive, however the results are still enlightening. Table 6.6: Results of the KS tests. The PBF spectral slopes are compatible with the spectral slopes of all families except for Polana. Family mF amily nPBF DP BF,vis DP BF,vis,crit Compatible? mF amily nPBF DP BF,NIR DP BF,NIR,crit Compatible? Polana - - - - 25 9 0.569 0.529 no Erigone 90 4 0.422 0.695 yes - - - - Sulamitis 59 4 0.441 0.703 yes 15 11 0.371 0.557 yes Klio 30 7 0.448 0.571 yes 20 8 0.222 0.546 yes Chaldaea 14 1 0.929 1.408 yes 16 3 0.278 0.848 yes 156 We split the PBF NIR spectra into three different inclination regions and compare with the family spectra at those inclinations in Fig. 6.15. The average spectrum for each family is plotted in a colored line with ±1σ shaded around it. In most cases, the PBF spectra overlap with the family averages; however, there are some outliers. The range of spectral slopes ∆S0 is 1.068 %/1000A˚ for the low inclination group, 4.128 %/1000A˚ for the middle inclination group, and 3.537 %/1000A˚ for the highest inclination group. The spread of slopes in the mid and high inclination groups is more than double that of the low inclination group. We now discuss each inclination region in detail. 157 Figure 6.15: A comparison of the NIR spectra of PBF members (grey) and family averages (col- ors) at three different inclination ranges (3°, 5°, and 11°). The PBF spectra match well with the family spectra at each inclination, except for some outliers. Polana average is from Pinilla-Alonso et al. (2016), Klio, Chaldaea, and Sulamitis averages are from Arredondo et al. (2020, 2021a,b), respectively. 158 The lowest inclination range is near ∼3°and includes the Polana and Clarissa families. We were unable to compare visible slopes because none of our visible PBF sample were located around 3°(see Fig. 6.1). In our NIR sample, the PBF objects at Polana-like inclinations have similar spectra to the average spectrum of the Polana family (Fig. 6.15) and both populations are homo- geneous, but the results of the KS test show that they are not similar distributions (Table 6.6). The Polana family has the highest abundance of B-type asteroids of all the IMB families (Figure 6.6) however the B-types of the PBF are nowhere near the Polana family in orbital space (Figure 6.7). de Leon´ et al. (2012) show that B-type asteroids can have a wide range of slopes in the NIR de- spite all being blue in the visible, so it is possible that some of our NIR objects were mislabeled as non-B-types because we do not have visible information for them. None of the Polana objects and only 3 of the Clarissa objects show evidence of hydration, as opposed to the 30% of the PBF. These results support the idea that the Polana and Clarissa families and the background near them do not come from the same primordial body, though the homogeneity in the two samples suggests that perhaps they originate from two primordial bodies located at similar inclination. The second inclination group around ∼5°includes the Erigone and Sulamitis families. Both fam- ilies have similar ages, which supports the idea that they originated from the same source. They also present roughly the same amount of hydrated objects and very similar distribution of taxo- nomic types. The PBF is less hydrated (30% compared to 60%) and has a larger fraction of B-type objects, however all three populations show spectral diversity in the visible. In the NIR, Arredondo et al. (2021b) show that the Sulamitis family is homogeneous and looks almost identical to all the other families. We compare the available visible and NIR slopes of the Sulamitis and Erigone fam- ilies and the PBF to show that they are compatible with each other (Table 6.6). These similarities support the idea that the Sulamitis and Erigone families and background objects originated from the same primordial object. 159 The third inclination group is centered near ∼11°and includes the Klio and Chaldaea families. Un- like the low inclination IMB, the high inclination families do not follow a trend with composition and inclination in their visible spectra. The Klio family is diverse and hydrated in the visible, but homogeneous in the NIR (Arredondo et al. 2020, Morate et al. 2019). Contrastingly, the Chaldaea family is homogeneous and hydrated in the visible, but diverse in the NIR (Arredondo et al. 2021a, Morate et al. 2019). Also, the Chaldaea objects present a much higher percent of hydration (80%) than the nearby Klio objects (23%). Morate et al. (2019) suggest that this disparity might be due to a common parent body with varying degrees of hydration with depth. If so, the background objects around the families should also follow this trend. Fig. 6.12 shows that out of the objects that overlap with the Klio and Chaldaea families, 2 hydrated objects overlap the Chaldaea family, while 4 unhydrated and 3 hydrated objects overlap with the Klio family. Table 6.6 shows that in the visible and NIR, the spectral slopes of our sample are compatible with the Klio and Chaldaea families, i.e., they might come from the same distribution. This result supports both the “asteroids formed big” hypothesis as well as the hypothesis of Morate et al. (2019) that Klio and Chaldaea come from one differentiated body. 6.3.7 Comparison with (101955) Bennu and (162173) Ryugu The targets of the NASA OSIRIS-REx and JAXA Hayabusa2 missions are the primitive asteroids (101955) Bennu and (162173) Ryugu, respectively. Campins et al. (2010, 2013) and Bottke et al. (2015) use dynamical models and spectral comparisons to show that both objects likely originate from low inclination IMB populations (i.e., the Polana, Clarissa, Erigone, and Sulamitis families and the low albedo background). We obtained the latest spectra of Bennu and Ryugu taken with the spacecraft (Hamilton et al. 2019) as well as ground-based spectra from before the encounters (Abe et al. 2008, Moskovitz et al. 2013) and we compare these spectra with the objects in the PBF. Both NEAs likely were delivered by the ν6 resonance which does not have a large effect on inclination 160 according to Campins et al. (2010, 2013); However, Granvik et al. (2018) modeled the inner belt to show that asteroids escaping through the ν6 and 3:1 resonance are systematically shifted to higher inclinations. Both Bennu and Ryugu have inclinations of ∼6°, so we also compare their spectra with the low inclination (i <8°) PBF objects only. Fig. 6.16 shows the comparison of the Bennu and Ryugu spectra with our entire sample and Fig. 6.17 shows the comparison of the Bennu and Ryugu spectra with the low inclination sample only. Figure 6.16: Comparison of the spectra of Bennu and Ryugu from Hamilton et al. (2019), Abe et al. (2008) and Moskovitz et al. (2013) with the spectra of our sample. The Ryugu spectrum is very similar to the rest of the background objects in both wavelength regions. The Bennu spectrum is similar to the few B-types in the sample. 161 Figure 6.17: Same as Fig. 6.16 but only including spectra from our sample with inclinations i <8°. There is only one blue NIR spectra like Bennu at these inclinations. The Ryugu spectra still match well with our sample. The spectrum of Bennu (navy line in Figs. 6.16 and 6.17) is blue sloped and featureless (i.e., the 0.7 µm band is not present). We note that not all ground-based spectra of Bennu show such a blue slope (see discussion in Arredondo et al. 2020), however, we take the spacecraft spectrum to be the most accurate. While the entire PBF sample includes 33% B-type objects (Section 6.3.2), it appears that the low inclination sample contains only one B-type like Bennu. Based off of this, we do not think that Bennu originated from the low albedo background population. The spectrum for Ryugu (green line in Figs. 6.16 and 6.17) is neutral to slightly red sloped. We note that there are many visible spectra of Ryugu published (Binzel et al. 2001, Lazzaro et al. 2013, Moskovitz et al. 2013, Perna et al. 2017, Sugita et al. 2013, Vilas 2008) and they do not all agree with each other (see discussion in de Leon´ et al. 2018). We choose to compare with the spectrum published by Moskovitz et al. (2013) because it has the highest signal to noise. The 162 Ryugu spectrum does not have the 0.7 µm feature, though 4 of our 8 C-type objects do. The Ryugu visible spectrum seems to lie right between the red PBF objects and the blue PBF objects but does not belong to either group. The NIR spectrum fits well within the spectra of the entire family and the low inclination population. Because of these spectral similarities, we conclude that the background is a plausible source region of Ryugu. Campins et al. (2010, 2013) identify the Polana family as the most likely source for Bennu and the background population as the most likely source of Ryugu. Our spectral comparisons support both of these claims. Besides the Polana family, the IMB family with the most amount of B- types is the Svea family (Fig. 6.6), however the inclination of Svea is too high to be a plausible source for Bennu and Ryugu. Of the low inclination populations, the background has the second largest abundance of B-types, however none of those B-types are at the inclination of Bennu. We therefore think that the more likely source for Bennu is the Polana family, though it is also possible that our sample is just too small to be conclusive. We do believe that the background is a very plausible source for Ryugu. In the low inclination IMB, most families are majority C-types, including the background population. Moreover, the location of the background is much closer to the ν6 resonance than any of the low inclination families (Fig. 6.1) and appears to deliver several times more objects to this resonance than any of the families (Campins et al. 2013). Our VNIR spectral sample of the background is very similar to Ryugu. The differences in origin of Bennu and Ryugu can help explain the differences we see on the surfaces (Lauretta et al. 2019). 163 6.3.8 Comparison with meteorites For a first order constraint on composition, we compare our spectra to the spectra of meteorites in the RELAB database (Pieters & Hiroi 2004). Comparison with meteorites is not the most accurate way to determine composition because the meteorites have been altered by the Earth’s atmosphere and do not accurately represent what asteroids look like in space, however it is still an educational exercise. For all of our NIR spectra, we use M4AST to join each spectrum with its visible component (if available), and then compare the full spectrum with the spectrum of every meteorite available in the RELAB database. M4AST returns the ten meteorite samples with the lowest χ2 residual from our asteroid spectra. If there was not an obvious best match (i.e., the top two matches had a difference of ∼0.001), we used our best judgement to decide which meteorite was the best fit based on visual inspection of the plots. Because primitive asteroid spectra are generally featureless, we can usually only find meteorite spectra that are the same general shape as our asteroids, as there are no absorption bands or features to fit to. We did this analysis for the primitive and non-primitive objects in our sample. Three of the objects were too noisy for a good fit and are excluded from the rest of the analysis (220, 2772, and 4750). Carbonaceous chondrites (CC) are primitive, undifferentiated meteorites, and so we expect CCs to be the closest match to our primitive asteroids. Of the asteroids with NIR spectra available, the majority of matches are indeed CCs (Table 6.7). CI and CM meteorites show evidence of aqueous alteration so we would expect our hydrated objects (indicated by a + in Table 6.7) to match best with those, and they do. There are also matches with ordinary chondrites (OC), achondrites, irons, and silicates. The non-primitive asteroids (marked with a * in Table 6.7) match with non- primitive meteorites, however some have CC matches. Some of the primitive asteroids also have non-primitive meteorite matches. These discrepancies highlight our point that this is not a perfect comparison and that fitting a featureless spectrum leaves room for error. 164 Table 6.7: The best meteorite fit for our samples. The most common match is with carbonaceous chondrites. A * indicates that the asteroid is a non-primitive type. A + indicates that the object is hydrated. Asteroid Sample name Sample ID General type Type Subtype Resid. 220 + no good match 249 + Murray MR-MJG-110 Rock Carbonaceous Chondrite CM2 0.00085 282 Y-82162,79 <125 µm MB-CMP-019-A Rock Carbonaceous Chondrite CI Unusual 0.00083 370 Y-82054,106 MP-TXH-058 Rock Carbonaceous Chondrite CM2 0.00147 689 A-881632,95 CO3 <125 µm MP-TXH-104 Rock Carbonaceous Chondrite CO3 0.00137 853 + Kaidun >250 µm MA-LXM-076 Rock Carbonaceous Chondrite CM or CR 0.00131 916 Bells #1226B(1) MB-TXH-053 Rock Carbonaceous Chondrite CM2 0.00119 917 Esquel MB-TXH-043 Rock Stony Iron Pallasite 0.00151 933 LEW87009,16 LM-LAM-011 Rock Carbonaceous Chondrite CK6 0.07080 1159 A-881594,62 MP-TXH-057 Rock Carbonaceous Chondrite CM2 0.00304 1159 Canon Diablo MB-CMP-016 Rock Iron IA 0.00103 1216 Almahata Sitta #44 125-500 µm MT-PMJ-108-B Rock Achondrite Ureilite Anomalous Polymict 0.00064 1244 Ozona MH-FPF-051-E Rock Ordinary Chondrite H6 0.00138 1544 QUE93005,9 CM2 thin section LM-LAM-032 Rock Carbonaceous Chondrite CM2 0.00623 1700 Divnoe MB-CMP-015-L Rock Primitive Achondrite Anomalous 0.00146 1705 Y-82162,79 <125 µm MB-CMP-019-A Rock Carbonaceous Chondrite CI Unusual 0.00086 1806 * Hamlet (LL4) chip irradiated OC-TXH-002-A60 Rock Ordinary Chondrite LL4 0.00486 1924 ALH-77307,95 CO3 <125 µm MP-TXH-102 Rock Carbonaceous Chondrite CO3 0.00102 2171 Chicora (LL6) <125 µm OC-TXH-014-C Rock Ordinary Chondrite LL6 0.00231 2259 + Murchison IOM 55-63 µm OG-CMA-004 Biological Organics CM2 0.00130 2322 Tsarev 1a/c RS-CMP-060-P4 Rock Ordinary Chondrite L5 0.01019 2328 Tsarev 1a/c RS-CMP-060-P4 Rock Ordinary Chondrite L5 0.00237 2536 Benares (a) MT-HYM-083 Rock Ordinary Chondrite LL4 0.00156 2575 * Chateau Renard (L6) <125 µm pellet irradiated OC-TXH-011-D35 Rock Ordinary Chondrite L6 0.00184 2575 EETA79001 lithology A whole rock DD-MDD-019 Rock Igneous Shergottite 0.00225 2772 Too noisy for a good fit 2773 PCA91008,15 MP-TXH-049 Rock Carbonaceous Chondrite C2 0.00263 2778 ALHA77005,193 olivine (not pure) DD-MDD-032 Mineral Silicate (Neso) Olivine 0.00342 2839 NWA1948 (LL6) <250 µm OC-SXS-023-D Rock Ordinary Chondrite LL6 0.00227 3633 MAC87300,46 MP-TXH-045 Rock Carbonaceous Chondrite C2 0.00284 3723 ALH85002,25 MB-TXH-081-1 Rock Carbonaceous Chondrite CK4 0.00582 4231 Tsarev 1a/c RS-CMP-060-P4 Rock Ordinary Chondrite L5 0.00402 4422 * META78008,34 MP-LAM-006-C2 Rock Achondrite Ureilite 0.01192 4524 Mundrabilla MB-TXH-034 Rock Iron Iranom 0.00167 4750 + no good match 5081 + EET87526,17 LM-LAM-010 Rock Carbonaceous Chondrite CK5 0.03699 5333 + Boriskino RS-CMP-046 Rock Carbonaceous Chondrite C2 0.00095 5524 Bereba MR-MJG-085 Rock Achondrite Eucrite (AEUC) 0.00445 5924 Y-82162,79 <125 µm MB-CMP-019-A Rock Carbonaceous Chondrite CI Unusual 0.02205 165 Asteroid Sample name Sample ID General type Type Subtype Resid. 6125 Y-984028,74 exterior whole rock 22 mg DD-MDD-118 Rock Igneous Lherzolitic Shergottite 0.00360 6542 MAC88100,30 pressed powder for laser irradiation MP-TXH-022-L0 Rock Carbonaceous Chondrite CM2 0.00671 6647 ALH78113,76 <45 µm AR-ASR-003-C Rock Achondrite Aubrite Enstatite Achondrite 0.02166 7132 Hvittis MR-MJG-024 Rock Enstatite Chondrite E6 0.00307 9723 Millbillillie No.4 MS-JTW-052-4 Rock Achondrite Laser-Irradiated Eucrite 0.00316 10446 Meteorites with steel MI-CMP-001 Rock - - 0.01070 10542 Esquel MB-TXH-043 Rock Stony Iron Pallasite 0.01758 14179 Alais MR-MJG-106 Rock Carbonaceous Chondrite CI 0.03095 15415 EET87503,97 45-75 µm MB-TXH-068-C Rock Achondrite Basaltic HED Howardite 0.01157 15998 .2 µm Grit Smo MI-CMP-012 Rock Iron IA 0.08014 20771 A-881955,70 MP-TXH-037 Rock Carbonaceous Chondrite CM2 0.05444 28620 ALH-77307,95 CO3 <125 µm MP-TXH-102 Rock Carbonaceous Chondrite CO3 0.04733 28736 Y-793592,101 Aubrite <125 µm MP-TXH-105 Rock Enstatite Achondrite Aubrite 0.00573 30514 GRO95577,6 MP-TXH-061 Rock Carbonaceous Chondrite C2 0.01089 31487 Tagish Lake ET01-B MT-MEZ-011 Rock Carbonaceous Chondrite Unusual C 0.00474 32898 Y-793592,101 Aubrite <125 µm MP-TXH-105 Rock Enstatite Achondrite Aubrite 0.01788 48876 Ivuna MB-TXH-060 Rock Carbonaceous Chondrite CI 0.01072 49863 PCA91084,8 MP-TXH-050 Rock Carbonaceous Chondrite C2 0.00853 We do note that even if the best match for an asteroid was a non-CC, the second or third best match was a CC. The point of this analysis is not necessarily to find the closest meteorite analog for each asteroid, but to search for an overall trend of composition for the primitive background. Because we see a range of spectral diversity in both our visible and NIR samples, it makes sense that we see some diversity in composition, however, the overall trend is that these primitive asteroids match best with the primitive material in carbonaceous chondrites. 166 6.3.9 Space weathering For silicate material, it is well known that space weathering causes reddening, darkening, and absorption band attenuation (Clark et al. 2002). It is unclear how space weathering affects low albedo asteroids and laboratory experiments seem to produce conflicting results (e.g., Thompson et al. 2019). Lantz et al. (2017, 2018) use ion irradiation of different carbonaceous chondrite samples to show that for low albedo material, the VNIR spectra of carbonaceous chondrites get bluer slopes, brighter albedos, and more convex spectra as a result of space weathering. This comparison assumes that the families all started with similar compositions, which we showed in Section 6.3.6 might not be true. Regardless, the PBF is older than any of the other families (Table 6.6) and should therefore show effects of space weathering compared to the other families. At each inclination region, we test whether the PBF objects are bluer, brighter, and more convex than the family objects at the same region. A summary of the results of our comparison is given in Table 6.8. Table 6.8: Comparison between PRIMASS samples of the IMB families in the NIR. Polana PBF Sulamitis PBF Klio Chaldaea PBF (Pinilla-Alonso et al. 2016) at Polana inclinations (Arredondo et al. 2021b) at Sulamitis inclinations (Arredondo et al. 2020) (Arredondo et al. 2021a) at Klio/Chaldaea inclinations Average 1.25 ± 0.02 1.03 ± 0.02 0.77 ± 0.01 NIR slope 0.68 ± 0.68 0.89 ± 0.26 1.05 ± 0.43 0.85 ± 0.42 (redder) (redder) (bluer) [%/1000A˚ ] 0.06 ± 0.01 0.05 ± 0.02 0.04 ± 0.03 Average pV 0.05 ± 0.01 0.07 ± 0.02 0.07 ± 0.01 0.08 ± 0.04 (brighter) (darker) (darker) Average 1.46±1.17 0.02 ± 0.13 -0.05 ± 0.03 NIR curvature 0.11 ± 0.02 0.24 ± 0.07 0.51 ± 0.11 -0.30 ± 0.21 (more convex) (flatter) (flatter) [%/µm2] 167 The NIR slope of the average spectrum of the low inclination (3°) PBF group is 1.25 ± 0.015 %/1000A˚ which is redder than the average slope of the Polana family, though they are similar within the uncertainties. The average albedo of the low inclination sample is 0.06 ± 0.01 which is brighter than the average albedo of the Polana family. The curvature of the average spectrum shows that the PBF sample (1.46±1.17 %/µm2) is more convex than the average Polana spectrum. The comparison of the low inclination group agrees with the predictions made by Lantz et al. (2017). The NIR slope of the average spectrum of the mid inclination (5°) PBF group is 1.03 ± 0.02 %/1000A˚ which is redder than the average slope of the Sulamitis family. The average albedo of the mid inclination sample is 0.05 ± 0.02 which is darker than the average albedo of the Erigone family (0.06 ± 0.01) and the Sulamitis family. The curvature of the average spectrum of the mid inclination PBF group is 0.02 ± 0.13%/µm2 which is flatter, and therefore more concave, than that of the Sulamitis family. The comparison of the mid inclination group does not agree with the predictions made by Lantz et al. (2017). The NIR slope of the average spectrum of the high inclination (11°) PBF group is 0.77 ± 0.01 %/1000A˚ which is bluer than the average slope of both the Klio and Chaldaea families. The average albedo of the high inclination sample is 0.04 ± 0.03 which is darker than the average albedo of both the Klio and Chaldaea families. The average curvature of our PBF sample is -0.05 ± 0.03 %/µm2 which is more concave than the Klio family and more convex than the Chaldaea family. The comparison of the high inclination group does not agree with the predictions made by Lantz et al. (2017). Much like the lab experiments in Thompson et al. (2019), our observational results do not seem to follow a singular trend. There must be other factors at play such as grain size, composition, process (micrometeorite bombardment vs solar wind erosion), etc. 168 Space weathering can also have an effect on absorption bands by attenuating the depth of the band. In our sample, we did not find any evidence for band attenuation and there was no trend with band depth and spectral slope. As shown in Fig. 6.13, there may be a trend with band center and spectral slope, with the redder objects having a shorter central wavelength. If this is the case, then another effect of space weathering on primitive bodies is a shift to shorter wavelengths. We look to see if this trend is also present in our sample of primitive families (Fig. 6.18). The Klio and Chaldaea families (black and green dots, respectively) follow this trend, however the difference between band centers is less pronounced. The Erigone, Sulamitis, and Chimaera families do not show this trend and are essentially a scatter plot. Figure 6.18: Same as Fig. 6.13 but included are the slope/band center relationships for the other IMB primitive families. The families in the left plot follow the trend of redder spectral slopes having shorter band centers that the PBF objects do. The families on the right do not. 169 The Pearson correlation coefficient is a measurement of how related two variables are, with 0 being not related and ±1 being linearly related. We calculated this coefficient, r, to show that visible slope and band center are correlated for Klio, Chaldaea, and the PBF (rK=-0.58, rC=-0.44, rPBF = -0.74) but not for Erigone, Sulamitis, and Chimaera (rE = 0.08, rS = -0.07, rChi = -0.20). At each inclination, the PBF objects have shallower average band depths and shorter average band centers than the families at those inclinations, suggesting that bands are attenuated and wavelengths are shifted shorter with age. We see another trend when comparing spectral slope with diameter (Fig. 6.19). All 6 of our blue sloped objects are larger than the average diameter of 16 km. Furthermore, the five reddest sloped objects are all smaller than 16 km. This could mean that there is a size limit to the objects that are affected by space weathering. For example, larger objects could be more susceptible to space weathering and that is why they are bluer. A similar idea of size-correlated slopes was proposed in De Pra´ et al. (2018) and Morate et al. (2019). Five of the six blue sloped objects in our sample are much darker than the average (pV <0.04). This would mean that space weathering causes objects to get darker, contrary to what was shown by Lantz et al. (2017, 2018), but in agreement with our comparison of the mid inclination and high inclination groups. 170 Figure 6.19: Relationship between diameter and spectral slope of the NIR sample. The average size of object in our sample is 16 km. The reddest (S0>2 %/1000A˚ ) objects are all smaller than the average and the bluest (S0<0 %/1000A˚ ) objects are all larger than the average. 171 6.4 Discussion 6.4.1 What the PBF tells us about the evolution of primitive asteroid families The PBF proposed by Delbo´ et al. (2017) does not look spectrally uniform; i.e., there is hetero- geneity in both the visible and NIR spectra. This suggests that there are processes at work that cause families to diversify over time. Our sample shows a diversity in taxonomic types, spectral slope, curvature, and composition suggested by meteorite matches. In Arredondo et al. (2021b), the surprising similarities in the NIR spectra of the Polana and Sulamitis families led us to suspect that NIR spectroscopy may not be very diagnostic when trying to distinguish between primitive families in the IMB. However, this is not so for our PBF sample, and NIR diversity seems to be a unique feature of the background population, as the other IMB families we have studied are all homogeneous in the NIR. There are a few possible explanations for the diversity seen in both the visible and NIR samples. The first is that the background could come from multiple different progenitor bodies with different compositions or perhaps one progenitor body that was differentiated. If the latter were true, we would expect similar objects to cluster together in orbital space. As shown in Figs. 6.7, 6.9, and 6.12, the distribution of similar objects is well mixed in orbital space. This mixing would have required a mechanism more powerful than the orbital resonances, such as the migration of the giant planets. This would mean that the formation of the PBF predates the mixing mechanism which puts a constraint on the timescale of the giant planet migration. 172 A second explanation for this diversity is that one of the groups could have been more altered by processes such as space weathering than the other group. The ancient age of the PBF has left more than enough time for all of the objects to be space weathered. If this were the case, we should be able to see similarities related to initial composition between the two populations such as the trend we saw with all the bluer objects being large (Fig. 6.19). The PBF can therefore serve as an example of what the asteroids in the primitive IMB families will look like in the future. This hypothesis can be tested by laboratory experiments of space weathering on primitive meteorites of different sizes, however that is outside the scope of this paper. 6.4.2 What the PBF tells us about the origin of the background population and its relationship to the dynamical families As stated before, there is a trend with composition and inclination in the lower IMB families. The Polana and Clarissa families both show homogeneous spectra and no evidence of aqueous alteration. The Erigone and Sulamitis families both show heterogeneous spectra and evidence of aqueous alteration. This trend does not extend into the high inclination families – Klio, Chal- daea, Chimaera and Svea present different amounts of hydration (Fig. 6.10) and vary in spectral diversity. When we split our PBF sample into inclination regions that match with the families, we see similar trends. In the low inclination group, the PBF spectra are homogeneous and in the two higher inclination groups, the PBF spectra are diverse. Based on the spectral comparisons and KS tests, the PBF follows the same compositional trend with inclination that the families do, with the exception of the Polana family. This suggests that the families are potentially the source of the background and might have originated from one or a few similar progenitor bodies. This 173 supports the conclusions of Dermott et al. (2018) and Morbidelli et al. (2009) that asteroids formed big. So while the background is distinct from itself, it is similar to the families at similar inclina- tions, except for Polana. Brasil et al. (2016) show that the giant planet instability could lead to a strong dispersion of asteroid eccentricities but would preserve inclination. This agrees with our hypothesis that PBF predates the giant planet migration. 6.4.3 What the PBF tells us about the evolution of the asteroid belt Dermott et al. (2018) showed that there is an anticorrelation with size and inclination in the IMB background objects, meaning that the higher inclination objects are larger. This makes sense log- ically as there are more collisions at lower inclinations because there are more asteroids nearer to the plane of the solar system. We should then see a trend with size and diversity, and we do note that all the blue objects in our NIR sample are in the top third of size distribution (Fig. 6.19). Therefore, the larger objects are more diverse and smaller objects are not as blue. The reddest objects all have D<13 km. This trend with size and slope could be related to space weathering. The presence of aqueous material based on the 0.7 µm absorption is distributed evenly around the IMB; we did not find any clusters of hydrated objects in the background or the families. The PBF shows a moderate amount of hydration (30%), not as hydrated as the most hydrated families but still much more hydrated than the “dry” or anhydrous families. While it may seem from Fig. 6.12 that there is no inclination trend with hydration, we note that our PBF sample does not extend down to Polana/Clarissa inclinations, and therefore is inconclusive. We also did not find correlations between band depth/center and any orbital, physical or dynamical properties. 174 We note that just because an object does not present the 0.7 µm band does not mean it is not hydrated. For example, the ground-based spectrum of Bennu from Clark et al. (2011) did not show the 0.7 µm absorption feature, however spectra from closer to the object (Hamilton et al. 2019) showed an absorption feature near 2.7 µm and thermal infrared spectral features that are indicative of a hydrated surface. In our own sample, five of the seven B-types are similar to Bennu in that they do not show the 0.7 µm feature. We searched the literature for 3-micron spectroscopy of any of our objects but did not find any. We decided not to apply for telescope time to do the observations ourselves because the SNR would have been very low for objects this faint. So, while only 30% of our sample shows the 0.7 µm feature, it is possible that more of these objects are hydrated. 6.5 Conclusion Based on our spectral comparison in the visible and NIR range, we show that the PBF is diverse and well mixed in terms of slope, size, curvature, composition, and taxonomy. The lack of homogeneity within this ancient family is meaningful. Possible explanations for this are: • The “asteroids formed big” hypothesis (Morbidelli et al. 2009) is incorrect • Asteroids formed big but those big objects were well differentiated • There was a large scale mixing event, such as the hypothesized giant planet migration • Space weathering has mainly affected larger objects The lack of trends we observed within the PBF indicates that this is the fate of the currently well- defined families in the IMB. 175 We compare our spectra of the background population with families at similar inclinations to show that the Sulamitis, Klio, and Chaldaea families and the background population near those families originated from the same or very similar sources. We show that the Polana family is statistically different from the PBF in that area, and probably did not come from the same source. There may be bias in our result because we only observed objects that were part of a primordial background family proposed by Delbo´ et al. (2017), however based on the distribution of these objects in orbital space (Fig. 6.1), we do believe that our sample is a good representation of the entire IMB background population. Based on our spectral comparisons, the primordial background cannot be ruled out as a possible source for Ryugu and is a more likely source of Ryugu than the Polana family. Future work to strengthen these conclusions includes obtaining spectra of background objects that are not in the PBF (i.e., non-family asteroids in Dermott et al. 2018) and obtaining NIR spectra of the families we have not characterized yet (Table 6.6). We have already observed and are working on analyzing the NIR spectra for the Erigone family and have begun to observe the Svea and Chimaera families. 176 CHAPTER 7: SUMMARY In this dissertation, we used visible and NIR spectroscopy to characterize several primitive IMB asteroid families and the background near them. In this chapter, I discuss these results and the future work of PRIMASS. 7.1 Primitive asteroid families Our analysis of objects in the Klio, Chaldaea, and Sulamitis families in Chapters 3, 4, and 5 led to three major conclusions. First, we learned that NIR spectroscopy is not very diagnostic to differen- tiate between families. With the exception of the significant curvature difference of the Chaldaea family, the NIR spectra of primitive IMB families are all similar within their uncertainties. Sec- ond, we learned that the extremely different groups in the low inclination families (“Polana-like” and “Erigone-like” groups in Campins et al. 2018) are not noticeable with NIR spectra. Lastly, we learned that some laboratory simulations of space weathering match well to observations of asteroids with similar compositions but very different ages. The 20 Klio family spectra are spectrally homogeneous, have red slopes, and show a fairly even distribution of positive and negative curvature. The 14 Chaldaea family spectra are spectrally homogeneous, have red slopes, and have overwhelmingly concave curvature. The 15 Sulamitis family spectra are spectrally homogeneous, have red slopes, and show a fairly even distribution of positive and negative curvature. In addition to the work in this dissertation, Pinilla-Alonso et al. (2016) characterized 45 NIR spectra of objects in the Polana family. They found that the Polana family is spectrally homogeneous, has red slopes, and we showed in Arredondo et al. (2021) that the Polana family spectra have mostly convex curvature. All four families are similar 177 within the uncertainties in terms of albedo, average spectral slope, and slope dispersion. The one distinguishing feature in these NIR spectra is the significant concavity of the Chaldaea family spectra which we showed was significantly different using a KS test. The other three families’ curvatures are similar to each other within the uncertainties. We conclude that NIR spectroscopy is not very diagnostic for primitive IMB families, though we acknowledge that it has been diagnostic for primitive outer belt families (e.g., Fornasier et al. 2016, Ziffer et al. 2011). Our analysis of the Sulamitis family showed that it is very similar to the Polana family in the NIR, and that the diversity seem in the visible spectrum of the Sulamitis family does not extend into the NIR. This is in contrast with the two groupings based on the visible spectra, “Polana- like” and “Erigone-like”, discussed by Morate et al. (2018) and Campins et al. (2018); however, it is consistent with our first conclusion that NIR spectroscopy is not diagnostic for primitive IMB families. Though the NIR spectra of these families are very similar, they are not identical. The small differences between the spectra of families could be consistent with space weathering effects. Our comparison of the Polana family and the younger Klio family showed that the spectra are consistent with the bluer, brighter, and more convex spectra theorized in laboratory studies by Lantz et al. (2017). The curvature difference of the Chaldaea family could also be explained by space weathering, if the hypothesis of Morate et al. (2019) that the Klio and Chaldaea families originated in one body is true. This same effect is theorized between the Polana and Clarissa families in Morate et al. (2018) and Campins et al. (2018). 178 7.2 The primordial background population Our analysis of objects in the primordial background family in Chapter 6 showed that the PBF is diverse and well mixed in terms of slope, size, curvature, composition, and taxonomy. This diversity leads to three main points. First, the heterogeneity of the PBF shows the fate of the current dynamical families, and this diversity could be due to a difference in how space weathering affects different objects. Second, we compared the background spectra with family spectra to show that the background is similar to the families near it. This suggests that the families are the source of the background and that the families and the background come from the same source. Lastly, the background is the most likely source of asteroid Ryugu based on its abundance of C-type asteroids and its proximity to the ν6 resonance. The 17 visible and 41 NIR spectra of PBF objects do not look spectrally uniform, suggesting that there are processes that change the surfaces of asteroids over time. The most realistic explanation for this is that space weathering affects different objects in different ways. Weather this alteration is more prominent based on size, albedo, mineralogy, grain size, or hydration is unknown. The idea of space weathering correlated with size was proposed by De Pra´ et al. (2018) and Morate et al. (2019) and supported by the observations in Section 6.3.9; however, space weathering on airless bodies is likely a more complex process which will require much laboratory studies and observational surveys to fully understand. We grouped our PBF spectra by inclination and compared with the spectra of the primitive IMB families at similar inclination to show that with the exception of the low-inclination Polana family, the families and the background have similar composition. This result supports the “asteroids formed big” hypothesis of Morbidelli et al. (2009) because all asteroids at similar inclinations (family and background) have similar compositions. This would mean that the objects in current 179 dynamical families are actually the grandchildren of a primordial body that broke up to form the family parent bodies and the background. While the PBF is diverse within itself, it follows the same general spectral trends with inclination as the primitive families (e.g., Morate 2018). This could also be an indication of several primitive progenitors at several inclinations (e.g., Lowry 2018). The NEA (162173) Ryugu is a C-type asteroid with an inclination of 5.88°that was delivered to NEA space from the ν6 resonance Campins et al. (2013). Ground based spectra of Ryugu do not show a absorption feature at 0.7 µm , which is similar to half of the C-type asteroids in our sample. The visible and NIR spectra of Ryugu match well with the spectra of PBF objects with inclinations <8°. While other low inclination primitive IMB families have C-type objects without the 0.7 µm feature, the background population is significantly closer to the ν6 resonance and therefore much more likely to have delivered Ryugu to NEA space. 7.3 Future work The characterization and comparison of primitive IMB families provides insights to the relation- ships between families, NEAs, and the background as well as to the potential links between fam- ilies. The work presented in this thesis is a continuation of the work started by PRIMASS to spectrally characterize the primitive asteroid population. The next steps in this work are to • Finish observing the IMB families in the NIR to test links between families and the spectral trends with inclination • Observe parent bodies using 3 µm spectroscopy to gain a better understanding of the distri- bution of hydrated minerals in the IMB 180 • Characterize the background in the other regions of the asteroid belt to test the hypothesis that the families are the source of the background population and that asteroids formed big As Table 6.5 shows, our knowledge of the primitive IMB families is incomplete. There are four more families (Erigone, Clarissa, Svea, and Chimaera) that need to be observed and characterized with NIR spectroscopy. The Erigone and Clarissa families will provide a test similar to the one in Section 5.4.2 between the Polana and Sulamitis families. The Svea and Chimaera families will help to test the hypothesis that families at similar inclinations have similar compositions. Recent findings of the 3 µm absorption feature on in-situ measurements from asteroid Bennu de- spite ground and spaced based spectra not showing the 0.7 µm absorption lead to a natural con- clusion that more 3 µm spectra are necessary. Campins et al. (2010) used dynamics and spectral comparisons to show that Bennu most likely originated in the Polana family. An spectrum from Rivkin (Unpublished) (Fig. 7.1), shows that the parent body of the Polana family, (142) Polana, does not show any absorption feature near 3 µm, thus making the argument of Campins et al. (2010) more complicated. Bennu could still originate in Polana family if for instance, there is some re- lationship between the strength of the 3 µm feature and size. If there were a regolith process that works on smaller objects like Bennu but does not affect larger objects like Polana, it could be hiding the 3 µm feature in the Polana spectrum. If Polana is completely anhydrous, then we must look elsewhere for the origin of Bennu. In this case, a plausible origin family are the Sulamitis and Erigone families. While these two families have a much smaller fraction of B-type asteroids than the Polana family, they are both hydrated and closer in inclination to Bennu. Fig. 7.1 shows that the parent body of the Sulamitis family (752) Sulamitis, does present an absorption feature in the 3 µm region. 181 Figure 7.1: 3 µm spectroscopy of the parent bodies of three primitive families and the asteroid Bennu. The grey regions represent the atmospheric absorption bands, where spectral data is not useful. The parent bodies of the two families with majority B-types like Bennu (Polana and Svea) do not present the 3 µm feature associated with hydration. The Sulamitis family, however, does have this feature, making it a plausible source for Bennu. 182 While 3 µm spectroscopy is difficult due to absorption by Earth’s atmosphere and other factors, objects that are bright can be observed from Earth and fainter objects can be observed from tele- scopes such as the Stratospheric Observatory for Infrared Astronomy (SOFIA) or the future James Webb Space Telescope (JWST) which is scheduled for launch in late 2021. The parent body of each family should be extensively studied, because the composition of that body is most likely the same as the composition of all the other fragments and because it is often the largest family member. Asteroids (142) Polana, (752) Sulamitis and (329) Svea already have 3 µm data available (Morate 2018, Rivkin Unpublished, Fig. 7.1); however, the parent bodies of the other five families do not have any 3 µm data available in the literature and should be the targets of future study. As shown in Chapter 6, the population of background asteroids is fascinating and not well studied. Future work to fully understand the formation and evolution of the asteroid belt must include detailed study of background asteroids similar to the characterization done in this work. In the inner belt, we should obtain spectra of background objects that are not included in the PBF, especially those at higher and lower inclinations and those closest to the ν6 resonance. In the middle and outer asteroid belt, the background should be studied and compared with families at similar inclinations to test if they are compositionally similar. New developments in technology such as JWST and the ESA Gaia mission’s data release 3 will add to our understanding of asteroid formation and the evolution of the main belt. 183 APPENDIX : COPYRIGHT PERMISSIONS 184 185 186 187 188 189 190 191 192 193 194 LIST OF REFERENCES Abe, M., Kawakami, K., Hasegawa, S., Kuroda, D., Yoshikawa, M., Kasuga, T., Kitazato, K., Sarugaku, Y., Kinoshita, D., Miyasaka, S., Urakawa, S., Okumura, S., Takagi, Y., Takato, N., Fujiyoshi, T., Terada, H., Wada, T., Ita, Y., Vilas, F., & Weissman, R. P. 2008, in 39th Lunar and Planetary Science Conference, 1594 Arredondo, A., Campins, H., Pinilla-Alonso, N., de Leon,´ J., Lorenzi, V., & Morate, D. 2021a, Icarus, 354, 114028 —. 2021b, Icarus, 358, 114210 Arredondo, A., Lorenzi, V., Pinilla-Alonso, N., Campins, H., Malfavon, A., de Leon,´ J., & Morate, D. 2020, Icarus, 335, 113427 Baffa, C., Comoretto, G., Gennari, S., Lisi, F., Oliva, E., Biliotti, V., Checcucci, A., Gavrioussev, V., Giani, E., Ghinassi, F., Hunt, L. K., Maiolino, R., Mannucci, F., Marcucci, G., Sozzi, M., Stefanini, P., & Testi, L. 2001, A&A, 378, 722 Bell, J. F., Hawke, B. R., Gaffey, M. J., & Owensby, P. D. 1984, Bulletin of the American Astro- nomical Society, 16, 692 Binzel, R. P. 1987, Icarus, 72, 135 Binzel, R. P., DeMeo, F. E., Burt, B. J., Cloutis, E. A., Rozitis, B., Burbine, T. H., Campins, H., Clark, B. E., Emery, J. P., Hergenrother, C. W., Howell, E. S., Lauretta, D. S., Nolan, M. C., Mansfield, M., Pietrasz, V., Polishook, D., & Scheeres, D. J. 2015, Icarus, 256, 22 Binzel, R. P., Harris, A. W., Bus, S. J., & Burbine, T. H. 2001, Icarus, 151, 139 Bottke, W. F., Morbidelli, A., Jedicke, R., Petit, J.-M., Levison, H. F., Michel, P., & Metcalfe, T. S. 2002, Icarus, 156, 399 195 Bottke, W. F., Vokrouhlicky,´ D., Walsh, K. J., Delbo,´ M., Michel, P., Lauretta, D. S., Campins, H., Connolly, H. C., Scheeres, D. J., & Chelsey, S. R. 2015, Icarus, 247, 191 Brasil, P. I. O., Roig, F., Nesvorny,´ D., Carruba, V., Aljbaae, S., & Huaman, M. E. 2016, Icarus, 266, 142 Brasser, R., Morbidelli, A., Gomes, R., Tsiganis, K., & Levison, H. F. 2009, A&A, 507, 1053 Bus, S. J. & Binzel, R. P. 2002, Icarus, 158, 106 Bus, S. J., Vilas, F., & Barucci, M. A. 2002, in Asteroids III (University of Arizona Press), 169–182 Campins, H., de Leon,´ J., Morbidelli, A., Licandro, J., Gayon-Markt, J., Delbo,´ M., & Michel, P. 2013, AJ, 146, 26 Campins, H., Leon,´ J. d., Licandro, J., Hendrix, A., Sanchez,´ J. A., & Ali-Lagoa, V. 2018, in Primitive Meteorites and Asteroids, ed. N. Abreu (Elsevier), 345–369 Campins, H., Morbidelli, A., Tsiganis, K., de Leon,´ J., Licandro, J., & Lauretta, D. 2010, ApJ, 721, L53 Carruba, V. & Nesvorny,´ D. 2016, MNRAS, 457, 1332 Carvano, J. M., Hasselmann, P. H., Lazzaro, D., & Mothe-Diniz,´ T. 2010, A&A, 510, A43 Cellino, A., Zappala,` V., Doressoundiram, A., Di Martino, M., Bendjoya, P., Dotto, E., & Miglior- ini, F. 2001, Icarus, 152, 225 Chapman, C. R., Gaffey, M., & McFadden, L. 1993, NASA Planetary Data System, EAR-A-DBP- 3-RDR-24COLOR-V2.1 Chapman, C. R., Morrison, D., & Zellner, B. 1975, Icarus, 25, 104 Christiansen, E. H. & Hamblin, K. W. 1995, Exploring the Planets, 2nd edn. (Prentice Hall) 196 Clark, B. E., Binzel, R. P., Howell, E. S., Cloutis, E. A., Ockert-Bell, M., Christensen, P., Barucci, M. A., DeMeo, F., Lauretta, D. S., Connolly, H., Soderberg, A., Hergenrother, C., Lim, L., Emery, J., & Mueller, M. 2011, Icarus, 216, 462 Clark, B. E., Hapke, B., Pieters, C., & Britt, D. 2002, in Asteroids III (University of Arizona Press), 585–599 Cushing, M., Vacca, W., & Rayner, J. 2004, PASP, 116, 362 de Leon,´ J., Campins, H., Morate, D., De Pra,´ M., Al´ı-Lagoa, V., Licandro, J., Rizos, J. L., Pinilla- Alonso, N., DellaGiustina, D. N., Lauretta, D. S., Popescu, M., & Lorenzi, V. 2018, Icarus, 313, 25 de Leon,´ J., Campins, H., Tsiganis, K., Morbidelli, A., & Licandro, J. 2010, A&A, 513, A26 de Leon,´ J., Pinilla-Alonso, N., Campins, H., Licandro, J., & Marzo, G. A. 2012, Icarus, 218, 196 de Leon,´ J., Pinilla-Alonso, N., Delbo,´ M., Campins, H., Cabrera-Lavers, A., Tanga, P., Cellino, A., Bendjoya, P., Gayon-Markt, J., Licandro, J., Lorenzi, V., Morate, D., Walsh, K. J., DeMeo, F., Landsman, Z., & Al´ı-Lagoa, V. 2016, Icarus, 266, 57 De Pra, M. N., Carvano, J., Morate, D., Pinilla-Alonso, N., & Licandro, J. 2018, in AAS/Division for Planetary Sciences Meeting Abstracts #50, 315.02 De Pra,´ M. N., Pinilla-Alonso, N., Carvano, J., Licandro, J., Morate, D., Lorenzi, V., Leon,´ J. d., Campins, H., & Mothe-Diniz,´ T. 2020, A&A, 643, A102 De Pra,´ M. N., Pinilla-Alonso, N., Carvano, J. M., Licandro, J., Campins, H., Mothe-Diniz,´ T., De Leon,´ J., & Al´ı-Lagoa, V. 2018, Icarus, 311, 35 Deienno, R., Walsh, K. J., & Delbo,´ M. 2021, Icarus, 357, 114218 197 Delbo,´ M., Walsh, K., Avdellidou, C., Fornasier, S., Deienno, R., Van Belle, G., & Morbidelli, A. 2019, in EPSC-DPS Joint Meeting, EPSC–DPS2019–877 Delbo,´ M., Walsh, K., Bolin, B., Avdellidou, C., & Morbidelli, A. 2017, Science, 357, 1026 Dellagiustina, D. N., Emery, J. P., Golish, D. R., Rozitis, B., Bennett, C. A., Burke, K. N., Ballouz, R.-L., Becker, K. J., Christensen, P. R., Drouet D’Aubigny, C. Y., Hamilton, V. E., Reuter, D. C., Rizk, B., Simon, A. A., Asphaug, E., Bandfield, J. L., Barnouin, O. S., Barucci, M. A., Bierhaus, E. B., Binzel, R. P., Bottke, W. F., Bowles, N. E., Campins, H., Clark, B. C., Clark, B. E., Connolly, H. C., Daly, M. G., Delbo,´ M., Deshapriya, J. D. P., Elder, C. M., Fornasier, S., Hergenrother, C. W., Howell, E. S., Jawin, E. R., Kaplan, H. H., Kareta, T. R., Le Corre, L., Li, J.-Y., Licandro, J., Lim, L. F., Michel, P., Molaro, J., Nolan, M. C., Pajola, M., Popescu, M., Garcia, J. L. R., Ryan, A., Schwartz, S. R., Shultz, N., Siegler, M. A., Smith, P. H., Tatsumi, E., Thomas, C. A., Walsh, K. J., Wolner, C. W. V., Zou, X.-D., De Leon,´ J. D., Lauretta, D. S., & Team, O.-R. 2019, Nat. Astron., 3, 341 DeMeo, F. E., Alexander, C. M. O., Walsh, K. J., Chapman, C. R., & Binzel, R. P. 2015, in Asteroids IV (University of Arizona Press), 13–41 DeMeo, F. E., Binzel, R. P., Slivan, S. M., & Bus, S. J. 2009, Icarus, 202, 160 DeMeo, F. E. & Carry, B. 2014, Nat., 505, 629 Dermott, S. F., Christou, A. A., Li, D., Kehoe, T. J. J., & Robinson, J. M. 2018, Nat. Astron., 2, 549 Feierberg, M. A., Lebofsky, L. A., & Tholen, D. J. 1985, Icarus, 63, 183 Fieber-Beyer, S. K. & Gaffey, M. J. 2020, Planet. Sci. J., 1, 68 Fornasier, S., Lantz, C., Barucci, M. A., & Lazzarin, M. 2014, Icarus, 233, 163 198 Fornasier, S., Lantz, C., Perna, D., Campins, H., Barucci, M. A., & Nesvorny, D. 2016, Icarus, 269, 1 Gaffey, M. J., Burbine, T. H., & Binzel, R. P. 1993, Meteorit. Planet. Sci., 28, 161 Gartrelle, G. M., Hardersen, P. S., Izawa, M. R. M., & Nowinski, M. C. 2020, in 51st Lunar and Planetary Science Conference, 1872 Gradie, J. & Tedesco, E. 1982, Science, 216, 1405 Granvik, M., Morbidelli, A., Jedicke, R., Bolin, B., Bottke, W. F., Beshore, E., Vokrouhlicky,´ D., Nesvorny,´ D., & Michel, P. 2018, Icarus, 312, 181 Gulbis, A. a. S., Bus, S. J., Elliot, J. L., Rayner, J. T., Stahlberger, W. E., Rojas, F. E., Adams, E. R., Person, M. J., Chung, R., Tokunaga, A. T., & Zuluaga, C. A. 2011, PASP, 123, 461 Hamilton, V. E., Simon, A. A., Christensen, P. R., Reuter, D. C., Clark, B. E., Barucci, M. A., Bowles, N. E., Boynton, W. V., Brucato, J. R., Cloutis, E. A., Connolly, H. C., Hanna, K. L. D., Emery, J. P., Enos, H. L., Fornasier, S., Haberle, C. W., Hanna, R. D., Howell, E. S., Kaplan, H. H., Keller, L. P., Lantz, C., Li, J.-Y., Lim, L. F., McCoy, T. J., Merlin, F., Nolan, M. C., Praet, A., Rozitis, B., Sandford, S. A., Schrader, D. L., Thomas, C. A., Zou, X.-D., & Lauretta, D. S. 2019, Nat. Astron., 3, 332 Hasselmann, P. H., Carvano, J. M., & Lazzaro, D. 2011, NASA Planetary Data System, EAR-A- I0035-5-SDSSTAX-V1.0 Kitazato, K., Milliken, R. E., Iwata, T., Abe, M., Ohtake, M., Matsuura, S., Arai, T., Nakauchi, Y., Nakamura, T., Matsuoka, M., Senshu, H., Hirata, N., Hiroi, T., Pilorget, C., Brunetto, R., Poulet, F., Riu, L., Bibring, J.-P., Takir, D., Domingue, D. L., Vilas, F., Barucci, M. A., Perna, D., Palomba, E., Galiano, A., Tsumura, K., Osawa, T., Komatsu, M., Nakato, A., Arai, T., Takato, N., Matsunaga, T., Takagi, Y., Matsumoto, K., Kouyama, T., Yokota, Y., Tatsumi, E., Sakatani, 199 N., Yamamoto, Y., Okada, T., Sugita, S., Honda, R., Morota, T., Kameda, S., Sawada, H., Honda, C., Yamada, M., Suzuki, H., Yoshioka, K., Hayakawa, M., Ogawa, K., Cho, Y., Shirai, K., Shimaki, Y., Hirata, N., Yamaguchi, A., Ogawa, N., Terui, F., Yamaguchi, T., Takei, Y., Saiki, T., Nakazawa, S., Tanaka, S., Yoshikawa, M., Watanabe, S., & Tsuda, Y. 2019, Science, 364, 272 Lantz, C., Binzel, R. P., & DeMeo, F. E. 2018, Icarus, 302, 10 Lantz, C., Brunetto, R., Barucci, M. A., Fornasier, S., Baklouti, D., Bourc¸ois, J., & Godard, M. 2017, Icarus, 285, 43 Lantz, C., Clark, B. E., Barucci, M. A., & Lauretta, D. S. 2013, A&A, 554, A138 Lauretta, D. S., Hergenrother, C. W., Chesley, S. R., Leonard, J. M., Pelgrift, J. Y., Adam, C. D., Al Asad, M., Antreasian, P. G., Ballouz, R.-L., Becker, K. J., Bennett, C. A., Bos, B. J., Bottke, W. F., Brozovic,´ M., Campins, H., Connolly, H. C., Daly, M. G., Davis, A. B., de Leon,´ J., Del- laGiustina, D. N., Drouet d’Aubigny, C. Y., Dworkin, J. P., Emery, J. P., Farnocchia, D., Glavin, D. P., Golish, D. R., Hartzell, C. M., Jacobson, R. A., Jawin, E. R., Jenniskens, P., Kidd, J. N., Lessac-Chenen, E. J., Li, J.-Y., Libourel, G., Licandro, J., Liounis, A. J., Maleszewski, C. K., Manzoni, C., May, B., McCarthy, L. K., McMahon, J. W., Michel, P., Molaro, J. L., Moreau, M. C., Nelson, D. S., Owen, W. M., Rizk, B., Roper, H. L., Rozitis, B., Sahr, E. M., Scheeres, D. J., Seabrook, J. A., Selznick, S. H., Takahashi, Y., Thuillet, F., Tricarico, P., Vokrouhlicky,´ D., & Wolner, C. W. V. 2019, Science, 366 Lazzaro, D., Angeli, C. A., Carvano, J. M., Mothe-Diniz,´ T., Duffard, R., & Florczak, M. 2004, Icarus, 172, 179 Lazzaro, D., Barucci, M. A., Perna, D., Jasmim, F. L., Yoshikawa, M., & Carvano, J. M. F. 2013, A&A, 549, L2 200 Lebofsky, L. A., Feierberg, M. A., & Jones, T. D. 1995, NASA Planetary Data System, EAR-A-3- RDR-THREEMICRON-V1.2 Lord, S. D. 1992, A New Software Tool for Computing Earth’s Atmospheric Transmission of Near- and Far-infrared Radiation (Ames Research Center) Lowry, V. 2018, in American Astronomical Society, DPS meeting #50, 414.11 Mainzer, A., Bauer, J., Cutri, R., Grav, T., Kramer, E., Masiero, J., Sonnett, S., & Wright, E. 2019, NASA Planetary Data System Marsset, M., DeMeo, F. E., Binzel, R. P., Bus, S. J., Burbine, T. H., Burt, B., Moskovitz, N., Polishook, D., Rivkin, A. S., Slivan, S. M., & Thomas, C. 2020, ApJ, 247, 73 Matlovic,ˇ P., De Leon,´ J., Medeiros, H., Popescu, M., Rizos, J. L., & Mansour, J.-A. 2020, A&A, 643, A107 Morate, D. 2018, PhD thesis, Universidad de La Laguna, Tenerife, Spain Morate, D., de Leon,´ J., De Pra,´ M., Licandro, J., Cabrera-Lavers, A., Campins, H., & Pinilla- Alonso, N. 2018a, A&A, 610, A25 Morate, D., de Leon,´ J., De Pra,´ M., Licandro, J., Cabrera-Lavers, A., Campins, H., Pinilla-Alonso, N., & Al´ı-Lagoa, V. 2016, A&A, 586, A129 Morate, D., de Leon,´ J., De Pra,´ M., Licandro, J., Pinilla-Alonso, N., Campins, H., Arredondo, A., Carvano, J. M., Lazzaro, D., & Cabrera-Lavers, A. 2019, A&A, 630, A141 Morate, D., Licandro, J., Popescu, M., & de Leon,´ J. 2018b, A&A, 617, A72 Morbidelli, A., Bottke, W. F., Nesvorny,´ D., & Levison, H. F. 2009, Icarus, 204, 558 201 Moskovitz, N. A., Abe, S., Pan, K.-S., Osip, D. J., Pefkou, D., Melita, M. D., Elias, M., Kitazato, K., Bus, S. J., DeMeo, F. E., Binzel, R. P., & Abell, P. A. 2013, Icarus, 224, 24 Nakamura, T., Lantz, C., Kobayashi, S., Nakauchi, Y., Amano, K., Brurnetto, R., Matsumoto, M., Takahashi, M., Matsuoka, M., Noguchi, T., Matsumoto, T., Miyake, A., Tsuchiyama, A., & Zolensky, M. E. 2019, in 82nd Annual Meeting of The Meteoritical Society, 6211 Nesvorny,´ D., Bottke, W. F., Vokrouhlicky,´ D., Sykes, M., Lien, D. J., & Stansberry, J. 2008, ApJ, 679, L143 Nesvorny,´ D., Broz,ˇ M., & Carruba, V. 2015, in Asteroids IV (University of Arizona Press), 297– 321 Oliva, E. 2000, Mem. Soc. Astron. Ital., 71, 861 Perna, D., Barucci, M. A., Ishiguro, M., Alvarez-Candal, A., Kuroda, D., Yoshikawa, M., Kim, M.-J., Fornasier, S., Hasegawa, S., Roh, D.-G., Muller,¨ T. G., & Kim, Y. 2017, A&A, 599, L1 Pieters, C. M. & Hiroi, T. 2004, in 35th Lunar and Planetary Science Conference, 1720 Pinilla-Alonso, N., de Leon,´ J., Walsh, K. J., Campins, H., Lorenzi, V., Delbo, M., DeMeo, F., Licandro, J., Landsman, Z., Lucas, M. P., Al´ı-Lagoa, V., & Burt, B. 2016, Icarus, 274, 231 Popescu, M., Birlan, M., & Nedelcu, D. A. 2012, A&A, 544, A130 Rayner, J., Toomey, D., Onaka, P., Denault, A., Stahlberger, W., Vacca, W., Cushing, M., & Wang, S. 2003, PASP, 115, 362 Reddy, V. & Sanchez, J. A. 2016, NASA Planetary Data System, EAR-A-I0046-3- REDDYMBSPEC-V1.0 Rivkin, A. S. Unpublished, (142) Polana spectrum, (Personal Communication) 202 Rivkin, A. S., Howell, E. S., Vilas, F., & Lebofsky, L. A. 2002, in Asteroids III (University of Arizona Press), 235–253 Sugita, S., Kuroda, D., Kameda, S., Hasegawa, S., Kamata, S., Hiroi, T., Abe, M., Ishiguro, M., Takato, N., & Yoshikawa, M. 2013, in 44th Lunar and Planetary Science Conference, 2591 Tatsumi, E., Sugita, S., Kameda, S., Honda, R., Kouyama, T., Yokota, Y., Sakatani, N., Honda, C., Michikami, T., Tomokatsu, M., Yamada, M., Suzuki, H., Cho, Y., Matsuoka, M., Hayakawa, M., Yoshioka, K., Ogawa, K., Sawada, H., Vilas, F., Domingue, D., Le Corre, L., Sasaki, S., Nakamura, T., & Hiroi, T. 2019, in 50th Lunar and Planetary Science Conference, 1753 Tholen, D. J. 1984, PhD thesis, The University of Arizona, Tucson, Arizona Thompson, M. S., Loeffler, M. J., Morris, R. V., Keller, L. P., & Christoffersen, R. 2019, Icarus, 319, 499 Tsiganis, K., Gomes, R., Morbidelli, A., & Levison, H. F. 2005, Nat., 435, 459 Vacca, W., Cushing, M., & Rayner, J. 2003, PASP, 115, 389 Vilas, F. 2008, AJ, 135, 1101 Vilas, F. & Gaffey, M. J. 1989, Science, 246, 790 Walsh, K. J., Delbo,´ M., Bottke, W. F., Vokrouhlicky,´ D., & Lauretta, D. S. 2013, Icarus, 225, 283 Walsh, K. J., Morbidelli, A., Raymond, S. N., O’Brien, D. P., & Mandell, A. M. 2011, Nat., 475, 206 Xu, S., Binzel, R. P., Burbine, T. H., & Bus, S. J. 1995, Icarus, 115, 1 Young, E. D., Ash, R. D., England, P., & Rumble, D. 1999, Science, 286, 1331 203 Ziffer, J., Campins, H., Licandro, J., Walker, M. E., Fernandez, Y., Clark, B. E., Mothe-Diniz, T., Howell, E., & Deshpande, R. 2011, Icarus, 213, 538 204