arXiv:2010.16079v1 [astro-ph.EP] 30 Oct 2020

LIS T O F C O N T RIB U T O RS

toshihiro kasuga
Public Relations Center
National Astronomical Observatory of Japan
2-21-1 Osawa Mitaka
Tokyo 181-8588
Japan
Department of Physics
Kyoto Sangyo University
Motoyama, Kamigamo Kita-ku
Kyoto 603-8555
Japan
david jewitt
Earth, Planetary and Space Sciences / Physics and Astronomy
University of California at Los Angeles
595 Charles Young Drive East Los Angeles
CA 90095-1567
USA

1

8
Asteroid–Meteoroid Complexes
TOSHIHIRO KASUGA AND DAVID JEWITT

8.1 Introduction

are comets is comparatively old, having been suggested by
Whipple (1939a,b, 1940) (see also Olivier, 1925; Hoﬀmeister, 1937). The Geminid meteoroid stream (GEM/4, from
Jopek and Jenniskens, 2011)1 and asteroid 3200 Phaethon
are probably the best-known examples (Whipple, 1983). In
such cases, it appears unlikely that ice sublimation drives
the expulsion of solid matter, raising the general questionof what produces the meteoroid streams? Suggested alternative triggers include thermal stress, rotational instability
and collisions (impacts) by secondary bodies (Jewitt, 2012;
Jewitt et al., 2015). Any of the above, if suﬃciently violent or prolonged, could lead to the production of a debris
trail that would, if it crossed Earth’s orbit, be classified asa meteoroid stream or an “Asteroid-Meteoroid Complex”,
comprising streams and several macroscopic, split fragments(Voloshchuk and Kashcheev, 1986; Jones, 1986; Ceplecha
et al., 1998).
The dynamics of stream members and their parent objects may diﬀer, and dynamical associations are not always
obvious. Direct searches for dynamical similarities employ a
distance parameter, DSH , which measures the separation inorbital element space by comparing q (perihelion distance),e (eccentricity), i (inclination), Ω (longitude of the ascending node), and ω (argument of perihelion) (Southworth and
Hawkins, 1963). A smaller DSH indicates a closer degree oforbital similarity between two bodies, with an empirical cutoﬀ for significance often set at DSH ! 0.10–0.20 (Williams
et al., 2019, Section 9.2.2). The statistical significance ofproposed parent-shower associations has been coupled with
DSH (Wiegert and Brown, 2004; Ye et al., 2016). Recent
models assess the long-term dynamical stability for highi and -e asteroids. Ohtsuka et al. (2006, 2008a) find the
Phaethon-Geminid Complex (PGC) and the Icarus complex
together using as criteria the C1 (Moiseev, 1945) and C2(Lidov, 1962) integrals. These are secular orbital variations
expressed by

Physical disintegration of asteroids and comets leads to the
production of orbit-hugging debris streams. In many cases,
the mechanisms underlying disintegration are uncharacterized, or even unknown. Therefore, considerable scientific interest lies in tracing the physical and dynamical propertiesof the asteroid-meteoroid complexes backwards in time, inorder to learn how they form.
Small solar system bodies oﬀer the opportunity to understand the origin and evolution of the planetary system. They
include the comets and asteroids, as well as the mostly unseen objects in the much more distant Kuiper belt and Oort
cloud reservoirs. Observationally, asteroids and comets are
distinguished principally by their optical morphologies, withasteroids appearing as point sources, and comets as diﬀuse
objects with unbound atmospheres (comae), at least whennear the Sun. The principal diﬀerence between the two isthought to be the volatile content, especially the abundanceof water ice. Sublimation of ice in comets drives a gas flux
into the adjacent vacuum while drag forces from the expanding gas are exerted on embedded dust and debris particles, expelling them into interplanetary space. Meteoroid
streams, consisting of large particles ejected at low speeds
and confined to move approximately in the orbit of the parent body, are one result.
When the orbit of the parent body intersects that of the
Earth, meteoroids strike the atmosphere and all but thelargest are burned up by frictional heating, creating the familiar meteors (Olmsted, 1834). Later, the phenomenon has
been realized as ablation by shock wave radiation heating(Bronshten, 1983). The first established comet-meteoroid
stream relationships were identified by G. Schiaparelli and
E. Weiss in 1866 (see Ceplecha et al., 1998). In the lasttwenty years, a cometary meteoroid stream theory has been
established enabling accurate shower activity prediction ofboth major (e.g. Leonids, Kondrat’eva et al., 1997; McNaught and Asher, 1999) and minor showers (2004 June
Boötids, Vaubaillon et al., 2005). This theory deals with theperturbed motion of streams encountering the Earth after
the ejection from relevant parent bodies. This improved theory has provided striking opportunities for meteor showerstudies of orbital trajectories, velocities and compositions,
resulting in a revolution in meteor science.
Some meteoroid streams seem to be made of debris released from asteroids. The notion that not all stream parents

C1=(1 − e2 ) cos2 (i)

(8.1)

C2

=

e2 (0.4 − sin2 (i) sin2 (ω)).

(8.2)

So-called time-lag theory is utilized to demonstrate longterm orbital evolution of complex members. When a streamcomplex is formed, the orbital energies (∝ a−1 , where a isthe semimajor axis) of ejected fragments are expected to beslightly diﬀerent from the energy of the precursor. The mo1

1

IAU Meteor Data Center, Nomenclature

2

Kasuga & Jewitt

tions of the released objects are either accelerated or decelerated relative to the precursor under gravitational perturbations (possibly including non-gravitational perturbations),
eﬀectively causing a time lag, ∆t, in the orbital evolution
to arise. Both C1 and C2 are approximately invariant during dynamical evolution, distinguishing the complex members. The PGC members (Phaethon, 2005 UD, 1999 YC),
for example, dynamically follow the Lidov-Kozai mechanism
based on secularly perturbed motion of the asteroids (Kozai,
1962).
Most known parent bodies are near-Earth Objects(NEOs), including both asteroids and comets. The comets
include various sub-types: Jupiter family comets (JFCs),
Halley type comets (HTCs) and Encke-type comets (Ye,2018) (Table 8.1). A classification of parents and theirassociated streams has been proposed based on their inferred evolutionary stages (Babadzhanov and Obrubov,
1987, 1991, 1992a,b). Some streams originating from asteroids (e.g. Phaethon), or from comets (e.g. 96P, 2P) are
the most evolved. For example, the Geminids, Quadrantids(QUA/10) and Taurid Complex (TAU/247) show stablesecular variation of the orbital elements under the mean
motion resonance with Jupiter and the Kozai resonance,
producing annual meteor showers. On the other hand, young
streams are usually from JFCs and HTCs. JFCs orbits, inparticular, are chaotically scattered by frequent close encounters with Jupiter, tending to produce irregular streams,
e.g. Phoenicids (PHO/254). HTCs orbit with widespread
inclinations, including retrograde orbits that are absent
in the JFCs, and may also generate regular showers, e.g.
Leonids (LEO/13) and Perseids (PER/7).
Non-gravitational force eﬀects can be important in theevolution of stream complexes but are diﬃcult to model,
since they depend on many unknowns such as the size, rotation, and thermal properties of the small bodies involved. In
the last decade, video- and radar- based surveys of meteors
have also been used to trace the trajectories back to potential parent NEOs, and so to find new stream complexes (e.g.
Brown et al., 2008a,b, 2010; Musci et al., 2012; Jenniskens,
2008a; Weryk and Brown, 2012; Rudawska et al., 2015; Jenniskens et al., 2016a,b; Jenniskens and Nénon, 2016; Ye
et al., 2016) (Reviewed in Jenniskens, 2017).
Meteor spectroscopy provides some additional constraints
on the composition of meteoroids (Millman and McKinley,
1963; Millman, 1980). Spectroscopy is typically capable ofobtaining useful data for meteors having optical absolutemagnitudes +3 to +4 or brighter, corresponding to meteoroid sizes " 1 mm (Lindblad, 1987; Ceplecha et al., 1998;
Borovička et al., 2010). Meteor spectra consist primarily of
atomic emission lines and some molecular bands in the visible to near-infrared wavelengths. The commonly identified
neutral atoms are Mg i, Fe i, Ca i, and Na i, while the singly
ionized atomic emissions of Ca ii and Mg ii also appear insome fast-moving meteors (e.g. Leonids with geocentric velocity Vg ∼72 km s−1 impact much more quickly than theGeminids, with Vg ∼ 35 km s−1 ). The abundances, the excitation temperatures and the electron densities can be deduced for each spectrum, assuming the Boltzmann distribution for electron energies, but the measurements are diﬃcult,and the resulting elemental abundances and/or intensity ratios (e.g. Nagasawa, 1978; Borovicka, 1993; Kasuga et al.,2005b; Borovička et al., 2005) are scattered (summarized in
Ceplecha et al., 1998; Kasuga et al., 2006). Generally, most
meteors are found to have solar abundance within factors of∼3–4. Some elements are noticeably underabundant, probably aﬀected by incomplete evaporation (Trigo-Rodrı́guez
et al., 2003; Borovicka et al., 1999; Kasuga et al., 2005b). In
particular, sodium (Na) is a relatively abundant and moderately volatile element that is easily volatilized. As a result,the abundance of Na in meteors is a good indicator of thermal evolution of meteoroids. Heating either during their residence in interplanetary space or within the parent bodies
themselves can lead to a sodium depletion.
In this chapter, we give a brief summary of observational
results on parents and their associated showers. We discuss
the properties of specific complexes, and tabulate their dynamical and physical properties (Tables 8.1, 8.2). We focus on the main meteoroid streams for which the properties and associations seem the most secure. Numerous
additional streams and their less certain associations are
discussed in the literature reviewed by (Jenniskens, 2006,
2008b; Borovička, 2007; Ye, 2018; Vaubaillon et al., 2019,
Section 7.5).

8.2 General Properties
To date, twelve objects in the six asteroid-meteoroid complexes have been studied in detail. Figure 8.1 represents
their distributions in the semimajor axis versus eccentricity
plane, while Figure 8.2 shows semimajor axis versus inclination (Table 8.1 summarizes the orbital properties).
Traditionally, the Tisserand parameter with respect toJupiter, TJ , is used to characterize the dynamics of small
bodies (Kresak, 1982; Kosai, 1992). It is defined by!

aa
TJ = J + 2 (1 − e2 )
a
aJ

"1/2

cos(i)(8.3)

where a, e and i are the semimajor axis, eccentricity and inclination of the orbit and aJ = 5.2 AU is the semimajor axisof the orbit of Jupiter. This parameter, which is conserved in
the circular, restricted 3-body problem, provides a measureof the close-approach speed to Jupiter. Jupiter itself has TJ= 3. Main belt asteroids have a < aJ and TJ > 3 while
dynamical comets from the Kuiper belt have 2 ≤ TJ < 3
and comets from the Oort cloud have TJ < 2. In principle,
the asteroids and comets can also be distinguished compositionally. The main belt asteroids are generally rocky, non-icy
objects which probably formed inside snow-line in the protoplanetary disk, while comets contain a larger ice fraction
and formed beyond it. In practice, it is diﬃcult or impossible to measure the compositions of most small bodies inthe solar system, so that composition is not often a usefuldiagnostic.
The use of TJ as a discriminant breaks down near TJ ≃ 3,
since the definition assumes that Jupiter’s orbit is a circle,

Asteroid–Meteoroid Complexes

3

Figure 8.1 Distribution of the parent bodies (black circles) in
the semimajor axis vs. orbital eccentricity plane (cf. Table 8.1).
The distributions of asteroids (brown dots) and comets(light-blue dots) are shown for reference. The lines for TJ = 3.08
with i=0◦ (solid curve) and with i=9◦ (dotted curve) broadly
separate asteroids and comets. Perihelion distances q = 0.25, 0.5
and 1 AU are shown as red curves.

Figure 8.2 Same as Figure 8.1 but semimajor axis vs.
inclination.

4Kasuga & Jewitt

Table 8.1. Orbital Properties
Complex
Geminids

Quadrantids
Capricornids

Taurids
Taurids-Perseids
Phoenicids

ObjectPhaethon2005 UD
1999 YC
2003 EH1
96P/Machholz 1
169P/NEAT
P/2003 T122017 MB1
2P/Encke
1566 Icarus
2007 MK6
2003 WY25

aa
1.271
1.275
1.422
3.123
3.018
2.604
2.568
2.372
2.215
1.078
1.081
3.046

eb
0.890
0.872
0.831
0.619
0.959
0.767
0.776
0.753
0.848
0.827
0.819
0.685

ic
22.25328.68238.22670.83859.97511.304
11.4758.508
11.78122.85225.1385.9000

qd
0.140
0.163
0.241
1.191
0.125
0.607
0.575
0.586
0.336
0.187
0.196
0.961

ωe
322.174
207.573
156.395
171.361
14.622217.977217.669264.628
186.54231.29725.4669.839Ωf
265.23119.74664.791282.979
94.548176.219176.465
126.974
334.569
88.08292.88768.931

Qg
2.402
2.387
2.603
5.056
5.911
4.602
4.561
4.158
4.094
1.969
1.966
5.132

Porb h
1.43
1.44
1.70
5.52
5.24
4.20
4.12
3.65
3.30
1.12
1.12
5.32

C1 i
0.178
0.184
0.191
0.066
0.020
0.396
0.382
0.424
0.269
0.268
0.270
0.525

C2 i
0.274
0.267
0.234
0.146
0.324
0.227
0.232
0.215
0.287
0.246
0.246
0.188

TJ j
4.509
4.504
4.114
2.065
1.939
2.887
2.894
3.071
3.025
5.296
5.284
2.816

Notes: Orbital data are obtained from NASA JPL HORIZONS (https://ssd.jpl.nasa.gov/horizons.cgi)aSemimajor axis (AU)b
EccentricitycInclination (degrees)dPerihelion distance (AU)e
Argument of perihelion (degrees)fLongitude of ascending node (degrees)gAphelion distance (AU)h
Orbital period (yr)iDynamical invariants: C1 = (1 – e2 ) cos2 (i), C2 = e2 (0.4 – sin2 (i) sin2 (ω))
jTisserand parameter with respect to Jupiter (asteroids with TJ > 3.08, comets with TJ < 3.08; a < aJ (5.2 AU))
the gravity of other planets is neglected, and so on. Accordingly, a functional definition for the boundary employed here
is TJ ≃ 3.08 (Jewitt et al., 2015), shown in Figure 8.1, where
the solid curve represents TJ computed assuming i = 0◦ ,while the dotted curve is the same but with i = 9◦ (equal
to the average inclination of 516,633 numbered objects in
the JPL Horizons database). They are very similar, indicating the definition broadly works independently of i. Thisavoids chaotic cases caused by deviation of the real solar
system from the circular, restricted 3-body approximation.
Also, comet 2P/Encke with TJ ∼ 3.03 and the quasi-Hilda
comets with TJ ∼ 2.9 – 3.04 are appropriately classified withthis criterion.
As seen in Figure 8.1, all of the parent bodies have e "0.6. Seven objects fall on the right side of the boundary with
TJ < 3.08, corresponding to the region of comets. Massloss activity has been directly detected in most of them,although 2003 EH1 and 2017 MB1 have yet to show evidence
for current activity. Objects on the left side of the diagram
are classified in the region of the asteroids, with TJ > 3.08.
In these objects, a range of physical processes appear todrive the mass loss. They have small q ! 0.25 AU and are
categorized as near-Sun objects (see Figure 8.1). Recurrent
activity of Phaethon at perihelion, including the formationof a tail, has been reported.
In Figure 8.2, 2003 EH1 and 96P show remarkably high-i" 60◦ . Five objects with a = 1.0 ∼ 1.4 AU have moderate iof 20◦ ∼ 40◦ . Another five objects at a = 2.2 ∼3.1 AU have
low-i, compatible with those of most main belt asteroids.

8.3 Known Asteroid-Meteoroid Complexes
The physical properties of the main asteroid-meteoroid complexes are listed in Table 8.2, and we review them focusing
on the observational evidence.
8.3.1 Geminids – (3200) Phaethon
The Geminid meteor shower is one of the most active annual showers (Whipple, 1939b; Spurný, 1993). The shower
currently has zenithal hourly rate (ZHR), the number of
meteors visible per hour under a clear-dark sky (! limiting
magnitude +6.5), of ∼ 120 and is expected to continue toincrease to a peak ZHR ∼ 190 in 2050 as the Earth moves
deeper into the stream core (Jones and Hawkes, 1986; Jenniskens, 2006). The Geminids are dynamically associatedwith the near-Earth asteroid (3200) Phaethon (1983 TB)(Whipple, 1983) (Figure 8.3). A notable orbital feature ofboth is the small perihelion distance, q ∼ 0.14 AU, raising
the possibility of strong thermal processing of the surface.
Indeed, the peak temperature at perihelion is ∼1000 K (Ohtsuka et al., 2009).
Spectroscopy of Geminid meteors shows an extreme diversity in Na content, from strong depletion of Na abundancein some to sun-like values in others. Kasuga et al. (2005a)
found a huge depletion in the Na/Mg abundance ratio of aGeminid meteor, with a value an order of magnitude smaller
than the solar abundance ratio (Anders and Grevesse, 1989;
Lodders, 2003). Line intensity ratios of Na i, Mg i, and Fe i
emissions of the Geminids also show a wide range of Na line
strengths, from undetectable to intense (Borovička et al.,Asteroid–Meteoroid Complexes

5

Table 8.2. Physical Properties
Complex
Geminids

Quadrantids
Capricornids

Taurids
Taurids-Perseids
Phoenicids

ObjectPhaethon1
2005 UD2
1999 YC3
2003 EH1 4
96P/Machholz 15
169P/NEAT6
P/2003 T122017 MB1 7
2P/Encke8
1566 Icarus92007 MK6
2003 WY25 10

De a
5–6
1.3±0.1
1.7±0.2g
4.0±0.3
6.4
4.6±0.60.52
4.8±0.4
1.0–1.3
0.18i≤ 0.32

pv b
0.09–0.13
0.11f
0.09±0.03g
0.04f
0.04f
0.03±0.01
0.05±0.02
0.30–0.50h
0.40f
0.04f

Prot c
3.604
5.249
4.495
12.650
6.38
8.410
6.69
11
2.273
-

B−Vd
0.59±0.01
0.66±0.03
0.71±0.04
0.69±0.01
0.73±0.02
0.73±0.06
0.76±0.02-

V − Rd
0.35±0.01
0.35±0.02
0.36±0.03
0.39±0.01
0.40±0.03
0.43±0.02
0.39±0.06
0.41±0.02-

R − Id
0.32±0.01
0.33±0.02
0.38±0.01
0.44±0.04
0.28±0.02
-

a/be
∼1.45
1.45±0.06
1.89±0.09
1.50±0.01
1.4 ≤1.31±0.03∼1.2
1.44±0.06
1.2–1.4-

Notes:
aEﬀective diameter (km)bGeometric albedocRotational period (hr)dColor index. Solar colors are B − V =0.64±0.02, V − R=0.35±0.01 and R − I=0.33±0.01 (Holmberg et al., 2006).
e
Axis ratiofAssumed valueg
De =1.4±0.1 with the assumed pv =0.11 (Kasuga and Jewitt, 2008).
hExtremely high pv > 0.7 is given De < 0.8 km (Mahapatra et al., 1999; Harris and Lagerros, 2002).
i
Estimated from the absolute magnitude H=20.3 (MPO386777) with the assumed pv =0.40.
1Green et al. (1985); Tedesco et al. (2004); Dundon (2005); Ansdell et al. (2014); Hanuš et al. (2016); Taylor et al. (2018)2
Jewitt and Hsieh (2006); Kinoshita et al. (2007) 3 Kasuga and Jewitt (2008); Mainzer et al. (2011); Warner (2017)4
Kasuga and Jewitt (2015) 5 Licandro et al. (2000); Meech et al. (2004); Lamy et al. (2004) 6 Kasuga et al. (2010);
Weissman et al. (2004); A’Hearn et al. (2005); DeMeo and Binzel (2008) 7 Warner (2018) 8 Kokotanekova et al. (2017);
Lamy et al. (2004); Campins and Fernández (2002); Lowry and Weissman (2007b); Fernández et al. (2000) 9 Gehrels et al.
(1970); Veeder et al. (1989); Chapman et al. (1994); Nugent et al. (2015); Miner and Young (1969); Harris (1998); Angelis(1995); Dundon (2005) 10 Jewitt (2006)

Figure 8.3 Composite R-band image of Phaethon taken at the
University of Hawai’i 2.2-m telescope (UH2.2) on UT 2003
December 19. The full image width is 1’ and this is a 3600 s
equivalent integration, with Phaethon at heliocentric distance R= 1.60 AU, geocentric distance ∆ = 1.39 AU, and phase angle α= 0.3◦ . The image shows no evidence of mass-loss activity. From
Hsieh and Jewitt (2005).

2005). Studies of Geminid meteor spectra, reported sincethe 1950s (e.g. Millman, 1955; Russell et al., 1956; Harvey,

1973; Borovička, 2001; Trigo-Rodrı́guez et al., 2003, 2004),
mostly fit this pattern.
As a summary for Na in the Geminids in the last decade,Kasuga et al. (2006) investigated perihelion dependent thermal eﬀects on meteoroid streams. The eﬀect is supposed toalter the metal abundances from their intrinsic values intheir parents, especially for temperature-sensitive elements:
a good example is Na in alkaline silicate. As a result, meteoroid streams with q ! 0.1 AU should be depleted in
Na by thermal desorption, because the corresponding meteoroid temperature (characterized as blackbody) exceeds
the sublimation temperature of alkaline silicates (∼ 900 K
for sodalite: Na4 (AlSiO4 )3 Cl). For this reason, the Na loss
in Geminids (q ∼ 0.14 AU) is most likely to be caused bythermal processes on Phaethon itself (see section 8.6.1).
The parent body Phaethon (diameter 5 – 6 km, from
Tedesco et al., 2004; Taylor et al., 2018) has an optically
blue (so-called “B-type”) reflection spectrum that distinguishes it from most other asteroids and from the nuclei ofcomets. Specifically, only ∼ 1 in 23 asteroids is of B-spectral
type and most cometary nuclei are slightly reddish, like theC-type asteroids or the (redder) D-type Jovian Trojans.
A dynamical pathway to another B-type, the large mainbelt asteroid (2) Pallas, has been reported with ∼ 2% prob-

6

Kasuga & Jewitt

Figure 8.4 Dynamical evolution of (3200) Phaethon and2005 UD for the six orbital elements, representing similar
behavior with a time-shift of ∼4600 yr. The abscissa shows timein Julian Terrestrial Date (JTD). From Ohtsuka et al. (2006)

Figure 8.5 R-band image of 2005 UD in 1800 s exposure
recorded using the UH2.2 on UT 2005 November 21. The region
shown is ∼ 150” in width and the distances and phase angle ofthe object were R = 1.59 AU, ∆ = 0.96 AU and α = 35.8◦ ,respectively. From Jewitt and Hsieh (2006).

ability (de León et al., 2010, see also Todorović, 2018). However, the colors of Phaethon and Pallas are not strictly identical, as would be expected if one were an unprocessed chipfrom the other, although both are blue. Color diﬀerences
might result from preferential heating and modification of
Phaethon, with its much smaller perihelion distance (0.14
AU vs. ∼2.1 AU for Pallas). Recently, Ohtsuka et al. (2006,
2008a) suggested the existence of a “Phaethon-Geminid
Complex (PGC)”–consisting of a group of dynamically associated split fragments, and identified the 1 km-sized asteroids 2005 UD and 1999 YC as having a common origin
with Phaethon (cf. Figure 8.4). Photometry of 2005 UD and1999 YC revealed the optical colors (Figures 8.5, 8.6). The
former is another rare blue object (Jewitt and Hsieh, 2006;
Kinoshita et al., 2007) while the latter is spectrally neutral(Kasuga and Jewitt, 2008) (see section 8.5.1).
The key question is how the Geminid meteoroid stream
was produced from Phaethon. The extreme possibilities arethat the Geminids are the products of a catastrophic event(for example an energetic collision, or a rotational disrup-

Figure 8.6 Keck telescope R-band image of 1999 YC on UT2007 October 12 showing a point source with F W HM ∼ 0.65”
in 400 s integration, centered within a frame 40” wide.
Heliocentric, geocentric distances and phase angle were R =2.60 AU, ∆ = 1.82 AU and α = 16.4◦ , respectively. From
Kasuga and Jewitt (2008).

tion) or that they are produced in steady-state by continuing
mass-loss from Phaethon.
In the steady-state case, the entire stream mass, Ms ∼1012 – 1013 kg (Hughes and McBride, 1989; Jenniskens,
1994; Blaauw, 2017)(cf. 1–2 orders larger in Ryabova, 2017),
must be released over the last τ ∼ 103 years (the dynamical
lifetime of the stream, cf. (Jones, 1978; Jones and Hawkes,
1986; Gustafson, 1989; Williams and Wu, 1993; Ryabova,
1999; Beech, 2002; Jakubı́k and Neslušan, 2015; Vaubaillon
et al., 2019, Section 7.5.4)). This gives dMs /dt ∼ Ms /τ =30 – 300 kg s−1 , comparable to the mass loss rates exhibitedby active Jupiter family comets. However, while the Jupiterfamily comets are notable for their distinctive comae ofejected dust, Phaethon generally appears as a point source,
devoid of coma or other evidence for on-going mass loss
(Chamberlin et al., 1996; Hsieh and Jewitt, 2005; Wiegert
et al., 2008) (cf. Figure 8.3).
Recently, this general picture has changed with the detection of mass loss using near-perihelion observations taken

Asteroid–Meteoroid Complexes

Figure 8.7 STEREO-A image taken on 2016 August 20 of (a)(3200) Phaethon and tail aligned on the position angles of theanti-solar direction and (b) the whole image sequence on thenegative heliocentric velocity vector projected on the sky plane.
Geometric parameters are R ∼ 0.14 AU, Phaethon-STEREO
distance ∼ 0.9 AU, and the angle ∼ 120◦ . The projected
Sun-Phaethon line (the narrow white line across the left panel)is drawn to better illustrate that the direction of the tail wasanti-solar, while the white dotted line across the right panel isthe orientation. From Hui and Li (2017).

7

Figure 8.8 2003 EH1 in a 360 s R-band image taken at Keckon UT 2013 October 2. No coma or tail is visible on the object
having an FWHM of 0”.86 in the frame of 40” × 25”. R=2.1 AU,∆=2.0 AU and α=27.6◦ . From Kasuga and Jewitt (2015).

8.3.2 Quadrantids – 2003 EH1with the Solar Terrestrial Relations Observatory (STEREO)
spacecraft in 2009, 2012 and 2016 (Jewitt and Li, 2010; Jewitt et al., 2013; Li and Jewitt, 2013; Hui and Li, 2017). In
addition to factor-of-two brightening at perihelion relative
to the expected phase-darkened inverse-square law brightness, a diﬀuse, linear tail has been resolved, as shown in
Figure 8.7. Because of its close association with the hightemperatures experienced at perihelion, the activity is likely
to result from thermal fracture or the desiccation of hydrated
minerals. The observed particles from Phaethon are micronsized, and are highly susceptible to solar radiation pressure
sweeping. They are rapidly accelerated by radiation pressure and so cannot be retained in the Geminid stream (Jewitt and Li, 2010). In addition, their combined mass, ∼ 3
× 105 kg per perihelion, is at least 107 × smaller than thestream mass, Ms (Jewitt et al., 2013). It is possible thatmuch more mass is contained in larger particles which, however, present a small fraction of the scattering cross-section
and which, therefore, are unsampled in the optical data fromSTEREO.
In this regard, larger particles (sizes >10 µm) were recently reported in thermal emission at 25 µm (Arendt,
2014), but evidently contribute little to the optical scattering cross-section. The particles in the stream are estimatedto be near mm-scale or larger (e.g. Blaauw, 2017), up to ∼1
to 10 cm, as measured in lunar impacts (Yanagisawa et al.,2008; Suggs et al., 2014; Ortiz et al., 2015; Szalay et al.,2018). The limit to optical depth of a Phaethon’s trail is≤ 3×10−9 (Jewitt et al., 2018), consistent with those ofcometary dust trails (10−9 –10−8 , from Sykes and Walker,
1992; Ishiguro et al., 2009) (see also section 8.6.3). Whilecontinued long wavelength observations of Phaethon to detect large particles will be helpful (Jewitt et al., 2015), thetrue nature of Phaethon and the PGC complex objects may
await spacecraft missions resembling NASA’s “Deep Impact” (A’Hearn et al., 2005; Kasuga et al., 2006, 2007a) and
JAXA’s “DESTINY+” (Sarli et al., 2018).

The Quadrantid meteor shower was first reported in 1835
(Quetelet, 1839) and appears annually in early January.
The shower consists of two diﬀerent components, the socalled young and old meteoroid streams, which represent
a very short duration of core activity (lasting ∼0.5 day)and a broader, longer-lived (∼4 days) background activity(Wiegert and Brown, 2005; Brown et al., 2010, and references therein). The width of a meteoroid stream depends
on its age, as a result of broadening by accumulated planetary perturbations. The small width of the Quadrantid corestream indicates ejection ages of only ∼200–500 years (Jenniskens, 2004; Williams et al., 2004; Abedin et al., 2015),
and there is even some suggestion that the first reports ofmeteoroid stream activity coincide with the formation of thestream. On the other hand, the broader background stream
implies larger ages of perhaps ∼3,500 years or more (Ohtsuka et al., 1995, 2008b; Kańuchová and Neslušan, 2007).
Two parent bodies of the Quadrantid complex have been
proposed. The 4 km diameter Near-Earth Object (196256)
2003 EH1 (hereafter 2003 EH1 ), discovered on UT 2003
March 6 by the Lowell Observatory Near-Earth-Object
Search (LONEOS) (Skiﬀ, 2003), may be responsible for the
young core stream (Jenniskens, 2004; Williams et al., 2004;
Wiegert and Brown, 2005; Babadzhanov et al., 2008; Jopek,
2011; Abedin et al., 2015). The orbit of 2003 EH1 has a =3.123 AU, e = 0.619, i = 70◦ .838 and q = 1.191 AU (Table 8.1). The TJ (= 2.07) identifies it as a likely Jupiterfamily comet, albeit one in which on-going activity has yet
to be detected (Koten et al., 2006; Babadzhanov et al., 2008;
Borovička et al., 2009; Tancredi, 2014). The steady-state
production rates ! 10−2 kg s−1 estimated from 2003 EH1
at R = 2.1 AU are at least five orders of magnitude too small
to supply the core Quadrantid stream mass Ms ∼ 1013 kg(Kasuga and Jewitt, 2015) (see Figure 8.8). Even at q=0.7–
0.9 AU a few hundred years ago, sublimation-driven activity
from the entire body takes ∼ 10s of years in the whole orbit,
being hard to reconcile. In order to form the core Quadrantid

8Kasuga & Jewitt

stream, we consider episodic replenishment by an unknown
process to be more likely.
Comet 96P/Machholz 1 has been suggested as the sourceof the older, broader part of the Quadrantid complex(Vaubaillon et al., 2019, Section 7.5.2), with meteoroids released 2,000–5,000 years ago (McIntosh, 1990; Babadzhanov
and Obrubov, 1991; Gonczi et al., 1992; Jones and Jones,
1993; Wiegert and Brown, 2005; Abedin et al., 2017, 2018).
Comet 96P currently has a small perihelion orbit (a =3.018 AU, e = 0.959, i = 59◦ .975 and q = 0.125 AU from
Table 8.1) substantially diﬀerent from that of 2003 EH1 . Despite this, calculations show rapid dynamical evolution thatallows the possibility that 2003 EH1 is a fragment of 96P, orthat both were released from a precursor body (together
defining the Machholz complex: Sekanina and Chodas,
2005). One or both of these bodies can be the parents of
the Quadrantid meteoroids (Kańuchová and Neslušan, 2007;
Babadzhanov et al., 2008; Neslušan et al., 2013a,b, 2014;
Vaubaillon et al., 2019, Section 7.5.2).
A notable dynamical feature of 2003 EH1 is the strong
evolution of the perihelion distance (Wiegert and Brown,
2005; Neslušan et al., 2013b; Fernández et al., 2014). Numerical integrations indicate that the minimum perihelion
distance q ∼ 0.12 AU (e ∼ 0.96) occurred ∼1500 yr ago(Neslušan et al., 2013b; Fernández et al., 2014), and the perihelion has increased approximately linearly with time from
0.2 AU 1000 years ago to the present-day value of 1.2 AU. At
its recently very small (Phaethon-like) perihelion distance, it
is reasonable to expect that the surface layers should havebeen heated to the point of fracture and desiccation (seesection 8.5.4).
As described above, the Phaethon-produced Geminid meteoroids (q ∼0.14 AU) show extreme diversity in their Naabundance, from strong depletion to near sun-like Na content (Kasuga et al., 2005a; Borovička et al., 2005). Curiously, the Quadrantid meteoroids from the core stream areless depleted in Na than the majority of Geminid meteoroids (Koten et al., 2006; Borovička et al., 2009). The interpretation of this observation is unclear (see section 8.6.2).
The optical colors of 2003 EH1 are similar to, but slightly
redder than, those of the Sun. They are most taxonomicallycompatible with the colors of C-type asteroids (Kasuga andJewitt, 2015) (see section 8.5.1).

8.3.3 Capricornids – 169P/NEATThe α-Capricornids (CAP/1) are active from late July toearly August, usually showing slow (∼22 km s−1 ) and brightmeteors. The shower, with an ascending nodal intersection of
ω=270◦ with the Earth, is expected to be a twin stream also
producing a daytime shower (Jenniskens, 2006). Because of
the low entry velocities, the meteor plasma excitation temperature is Tex ! 3600K and no trace of high temperaturegas (i.e. hot component of Tex ∼ 104 K) is found (Borovicka
and Weber, 1996). The metal contents of the α-Capricornids
are unremarkable, being within a factor of a few of the Solar abundance (Borovicka and Weber, 1996; Madiedo et al.,2014).

Figure 8.9 Comet 169P/NEAT in a 600 sec, r’-band image
taken at the Dominion Astrophysical Observatory 1.8 m′′
telescope on UT 2010 February 17. The frame size is 200 ×′′
150 . No coma or tail is visible on the object having an FWHM
of 2.8”. R=1.43 AU, ∆=0.47 AU and α=16.1◦ . From Kasuga
et al. (2010).

Recently Brown et al. (2010) suggested that the Daytime Capricornids-Sagittariids (DCS/115) are closely dynamically related to the α-Capricornids. One of the parent
body candidates, comet 169P/NEAT, has been identified as
the parent body of the α-Capricornid meteoroid stream bynumerical simulations (Jenniskens and Vaubaillon, 2010).
The object was discovered as asteroid 2002 EX12 by the
NEAT survey in 2002 (cf. Warner and Fitzsimmons, 2005;
Green, 2005) and was re-designated as 169P/NEAT in 2005
after revealing a cometary appearance (Green, 2005). The
orbital properties (Levison, 1996, TJ = 2.89) and opticalobservations reveal that 169P/NEAT is a ∼ 4 km diameter,
nearly dormant Jupiter family comet with tiny mass-loss
rate ∼ 10−2 kg s−1 (Kasuga et al., 2010) (Figure 8.9).
In the steady state, the stream mass Ms ∼1013 –1015 kg
and the age τ ∼5,000 yr (Jenniskens and Vaubaillon, 2010)
together require a mass-loss rate four orders of magnitudelarger than measured in 2010. In the case of 169P, cometaryactivity (a dust tail) was confirmed in 2005 (almost 1 orbital
period before) and episodic mass-loss should be expected.
This is a diﬀerent case from other asteroidal parents of complexes. Kasuga et al. (2010) used the fractional change in thespin angular velocity to estimate the mass loss from 169P as∼ 109 –1010 kg per orbit. With the Ms and τ , the conclusion
is that the origin of the α-Capricornids meteoroid stream
could be formed by the steady disintegration of 169P.
Other parent body candidates continue to be proposed
for the Capricornids. P/2003 T12 (SOHO) was suggested
to share a common parent with 169P, following a breakup∼2900 yr ago (Sosa and Fernández, 2015). Comet 169P isa large, almost inactive body (Kasuga et al., 2010), while
P/2003 T12 seems to be a very small comet, with a sub-km
radius nucleus (Sosa and Fernández, 2015) accompanied by
dust-tails in near-Sun STEREO-B observations (Hui, 2013).
The orbit of 2017 MB1 was suggested to resemble that ofAsteroid–Meteoroid Complexes

9the α-Capricornids meteor shower (Wiegert et al., 2017).
2017 MB1 has not been reported to show any sign of massloss activity.
8.3.4 Taurid Complex – 2P/Encke
The Taurid meteor shower includes the Northern, the Southern and other small branches (Vaubaillon et al., 2019, Section 7.5.5), possibly originating from more than one parent body. The Taurids show protracted, low-level activity
with many fireballs from September to December, peaking
in early November every year. The Taurid meteoroid complex has been suggested to be formed by a disrupted giant
comet (40 km-sized) 104 years ago (Clube and Napier, 1984,
1987; Asher et al., 1993), although the very recent breakup of such a large (i.e. rare) body is statistically unlikely.
Comet 2P/Encke has, for a long time, been considered as
the most probable parent of the shower (Whipple, 1940).
The Taurid complex has a dispersed structure with low
inclination and perihelia between 0.2 and 0.5 AU. The low
inclination of the stream orbit enhances the eﬀect of planetary perturbations from the terrestrial planets (Levison
et al., 2006), resulting in the observed, diﬀuse structure of
the complex (Matlovič et al., 2017). Furthermore, 2P has
a relatively small heliocentric distance (aphelion is Q =4.1AU) allowing it to stay mildly active around its orbit,
and producing a larger spread in q ∼0.34 AU, than could
be explained from ejection at perihelion alone (Gehrz et al.,2006; Kokotanekova et al., 2017).
An unfortunate artifact of the low inclination of the
Taurid complex (i ∼ 12◦ , Table 8.1) is that many nearEarth asteroids are plausible parent bodies based on orbital dynamical calculations (e.g. Asher et al., 1993; Steeland Asher, 1996; Babadzhanov, 2001; Porubčan et al., 2004,
2006; Babadzhanov et al., 2008). As a result, many of the
proposed associations are likely coincidental (Spurný et al.,2017; Matlovič et al., 2017). Actually no spectroscopic linkage between 2P/Encke and the 10 potential Taurid-complex
NEOs has been confirmed (the latter are classified variously
as X, S, Q, C, V, O, and K-types) (Popescu et al., 2014;
Tubiana et al., 2015), this being totally diﬀerent from thecase of the Phaethon-Geminid Complex. Here, we focus on
the physical properties of the Taurid meteor shower and themost strongly associated parent, 2P/Encke.
Spectroscopic studies of some Taurid meteors find a carbonaceous feature (Borovička, 2007; Matlovič et al., 2017).
The heterogeneity (large dispersion of Fe content) and low
strength (0.02–0.10 MPa) of the Taurids (Borovička et al.,2019, Section 2.3.4) suggest a cometary origin, consistent
with but not proving 2P as a parent (Borovička, 2007; Matlovič et al., 2017). Note that the current perihelion distance(q ∼0.34 AU, where T ∼ 480 K) is too large for strong thermal metamorphism to be expected (Borovička, 2007).
Comet 2P/Encke is one of the best characterized shortperiod comets, with published determinations of its rotation period, color, albedo and phase function (see reviews;
Lamy et al., 2004; Kokotanekova et al., 2017). For example, the eﬀective radius is 2.4 km, the average color indices,
B − V = 0.73±0.06 and V − R = 0.39±0.06 (e.g. Lowry

Figure 8.10 The 24 µm infrared image of comet 2P/Encke
obtained in 2004 June by the Spitzer Space Telescope. The
image field of view is about 6’ and centered on the nucleus. The
near horizontal emission is produced by recent cometaryactivity, and the diagonal emission across the image is themeteoroid stream along the orbit (Kelley et al.,2006; see also
Reach et al., 2007). Courtesy NASA/JPL-Caltech/M. Kelley(Univ. of Minnesota).

and Weissman, 2007b), the rotational period is about 11 hr
or 22 hr (but changing with time in response to outgassing
torques (Belton et al., 2005; Kokotanekova et al., 2017, andreferences therein); with minimum axis ratio 1.4 (Fernández
et al., 2000; Lowry and Weissman, 2007a; see also, Lamyet al., 2004).
Optically, 2P appears dust-poor because optically brightmicron and sub-micron sized dust particles are underabundant in its coma (Jewitt, 2004a). However, thedust / gas ratio determined from thermal emission is an
extraordinary µ ∼ 30, suggesting a dust-rich body (Reachet al., 2000; Lisse et al., 2004) compared to, for example,
Jupiter family comet 67P/Churyumov-Gerasimenko, whereµ = 4±2 (Rotundi et al., 2015). The total mass loss is 2–6
× 1010 kg per orbit, mostly in the form of large particles
that spread around the orbit and give rise to 2P’s thermal
dust trail as the source of Taurid meteor showers (Asher
and Clube, 1997; Reach et al., 2000) (see Figure 8.10). The
entire structured stream mass, Ms ∼ 1014 kg (Asher et al.,1994), must have been released over the last τ ∼ 5,000 –
20,000 years (the dynamical lifetime of the stream, cf. Whipple, 1940; Babadzhanov and Obrubov, 1992a; Jenniskens,
2006). This gives dMs /dt ∼ Ms /τ = 200 – 600 kg s−1 , a
few times larger than the mass loss rates typically reported
for active Jupiter family comets. Various small near-Earthobjects and some meteorite falls have been linked with theorbit of the stream as potentially hazardous (Brown et al.,2013; Olech et al., 2017; Spurný et al., 2017).
8.3.5 Sekanina’s (1973) Taurids-Perseids – Icarus
The Icarus asteroid family was reported as the first family found in the near-Earth region, which dynamically relates asteroids 1566 Icarus, 2007 MK6 and Sekanina’s (1973)
Taurid-Perseid meteor shower (Ohtsuka et al., 2007) (seeFigure 8.11).

10

Kasuga & Jewitt

Figure 8.12 2003 WY25 (=289P/Blanpain) imaged in R-band,
500 sec integration at R=1.6 AU, ∆=0.7 AU and α=20.7◦ using
UH2.2 on UT 2004 March 20. A faint coma is apparent,
extending to the southeast. Arrows show the directions of thenegative heliocentric velocity vector (marked ”-V”) and theanti-solar direction (”AS”). The estimated radius ∼ 160 m is thesmallest active cometary parent ever observed. From Jewitt(2006).

8.4 Possible Complexes
Figure 8.11 Dynamical evolution process of 1566 Icarus and2007 MK6 in JTD. The orbital elements (q, a, e, ω, Ω and i)and the C1 and C2 integrals are plotted (cf. Table 8.1).
Time-shifting is only ∼1,000 yr. From Ohtsuka et al. (2007).

Here, we describe two examples of less well-characterized
complexes suspected to include stream branches and one ormore parent bodies.

8.4.1 Phoenicids – Comet D/1819 W1 (Blanpain)
Near-Earth Apollo asteroid 1566 Icarus (= 1949 MA) was
discovered in 1949, having distinctive small q=0.19 AU andhigh i=23◦ (Baade et al., 1950). The object has diameterDe ∼ 1 km (e.g. Table 1, from Chapman et al., 1994), a moderately high albedo of 0.30 – 0.50 (cf. Gehrels et al., 1970;
Veeder et al., 1989; Chapman et al., 1994; Nugent et al.,2015) and a short rotational period, ∼2.273 hr (e.g. Miner
and Young, 1969) (see also, Harris and Lagerros, 2002). A
reflection spectrum close to Q- or V-type asteroids is found(Gehrels et al., 1970; Hicks et al., 1998; Dundon, 2005) (seesection 8.5.1).
On the other hand, near-Earth asteroid 2007 MK6 (=2006 KT67 ) was discovered in 2007 (Hill et al., 2007). Assuming an albedo like that of Icarus, then 2007 MK6 is∼180 m in diameter as computed from the absolute magnitude H=20.3 (MPO386777) (see Table 8.2). The breakup
hypothesis from Icarus, if true, could be due to near a critical rotation period and thermal stress induced at small
q∼0.19 AU (subsolar temperature ∼ 900 K), which might berelated to the production of the meteoroid stream (see section 8.5.2). The Taurid-Perseid meteoroids can be dynamically related with the Icarus asteroid family (DSH ∼ 0.08),
speculated to cross the Earth’s orbit (Sekanina, 1973). The
rare detection of the Taurid-Perseid meteor shower may result from the intermittent stream (swarm) due to very limited dust supply phase from the parent body.

The Phoenicid meteor shower was first reported more than
50 years ago, on December 5 in 1956 (Huruhata andNakamura, 1957). Promptly, the lost Jupiter family comet
D/1819 W1 (289P/Blanpain) was proposed as the potential
source (Ridley, 1957 reviewed in Ridley, 1963). In 2003, theplanet-crossing asteroid 2003 WY25 was discovered (Ticha
et al., 2003), with orbital elements resembling those of
D/Blanpain (Foglia et al., 2005; Micheli, 2005). The related
Phoenicids’ activity in 1956 and 2014 (Watanabe et al.,2005; Jenniskens and Lyytinen, 2005; Sato et al., 2017;
Tsuchiya et al., 2017), raised the possibility that 2003 WY25
might be either the dead nucleus of D/Blanpain itself ora remnant of the nucleus surviving from an earlier, unseen disintegration. Jewitt (2006) optically observed asteroid 2003 WY25 , finding the radius of 160 m (an order ofmagnitude smaller than typical cometary nuclei), and revealing a weak coma consistent with mass-loss rates of
10−2 kg s−1 (Figure 8.12). The latter is too small to supply
the estimated 1011 kg stream mass on reasonable timescales(≤10,000 yrs, Jenniskens and Lyytinen, 2005). Indeed, themass of 2003 WY25 (assuming density 1000 kg m−3 and aspherical shape) is ∼2×1010 kg, smaller than the stream
mass. Either the stream was produced impulsively by the
final stages of the break-up of a once much larger precursor
to 2003 WY25 , or another parent body may await discovery(Jewitt, 2006).

Asteroid–Meteoroid Complexes
8.4.2 Andromedids – Comet 3D/Biela

1.1

A

1.0

R

0.9

B-V

The Andromedid meteor shower (AND/18) was firstly reported in 1798 (Hawkins et al., 1959). The dynamics werelinked with Jupiter family comet 3D/Biela (Kronk, 1988,
1999). The shower is proposed to result from continuous
disintegration of 3D/Biela from 1842 until its sudden disappearance in 1852, resulting in irregular meteor shower appearances (e.g. Olivier, 1925; Cook, 1973; Jenniskens andVaubaillon, 2007). The estimated stream mass is 1010 kg(Jenniskens and Lyytinen, 2005), however, the absence ofparent candidates means that little can be determined about
the production of the meteoroids. Nonetheless, the Andromedid meteor shower was actually detected by radar in2011 and is numerically predicted to appear in the coming
decades (Wiegert et al., 2013).

11

S

VQ

0.8

T

G
C

0.7

D

Phaethon2005 UD
1999 YC
2003 EH1
169P
1566

X

B
F
0.6

0.5
0.30

0.35

0.40

0.45

0.50

0.55

0.60

V-R

Figure 8.13 Color plots of V − R vs. B − V showing parent
bodies (filled symbols) from Table 8.2, and Tholen taxonomic
classifications (Tholen, 1984), as tabulated by Dandy et al.
(2003). The color of the Sun (red circle) is also plotted.

8.5 Parent Bodies
0.60

A

0.55

0.50

R
V-R

In this section we discuss group physical properties ofthe parent bodies (cf. Table 8.2). Most objects (e.g.
3200 Phaethon, 2005 UD, 1999 YC, 2003 EH1 and 169P)
show point-like images (Figures 8.3, 8.5, 8.6, 8.8, 8.9) fromwhich we can be confident that the measured properties refer to the bare objects (or nuclei) alone. However, 2P/Encke,
2003 WY25 and some other comets may sometimes be active (e.g. Figures 8.10, 8.12) leading to potential confusion
between the properties of the nucleus and the near-nucleus
comae.

S
T

0.45

0.40

V

Q

X
C
BF G

0.35

Phaethon2005 UD
2003 EH1
169P
1566

0.30

8.5.1 Colors
Figures 8.13 and 8.14 show distributions of the colors of theparent bodies from Table 8.2. In addition, Tholen taxonomy
classes are plotted from photometry of NEOs from Dandy
et al. (2003). Here, 2P is not included because of the coma
contamination suggesting mild activity during the whole orbit.
The asteroids of the PGC (3200 and 2005 UD, 1999 YC)
show colors from nearly neutral to blue. Asteroids 3200
Phaethon and 2005 UD are classified as B-type asteroids(cf. Dundon, 2005; Jewitt and Hsieh, 2006; Kinoshita et al.,2007; Licandro et al., 2007; Kasuga and Jewitt, 2008; Jewitt,
2013; Ansdell et al., 2014), while 1999 YC is a C-type asteroid (Kasuga and Jewitt, 2008). Heterogeneity on the surfaces of Phaethon and 2005 UD may be due to intrinsicallyinhomogeneous composition, perhaps aﬀected by hydration
processes (Licandro et al., 2007), and by thermal alteration(Kinoshita et al., 2007). The rotational color variation of2005 UD shows B-type for 75% of the rotational phase but
C-type for the remainder (Kinoshita et al., 2007). The colors of the PGC objects are broadly consistent with being
neutral-blue.
Optical colors of 2003 EH1 are taxonomically compatible
with those of C-type asteroids (Kasuga and Jewitt, 2015)(Figures 8.13 and 8.14). The V-R color (0.39±0.01) is similar to that of 96P (V − R = 0.40±0.03, from Licandro et al.,

0.25
0.20

0.25

0.30

D

0.35

0.40

0.45

0.50

0.55

R-I

Figure 8.14 The same as Figure 8.13 but in the R − I vs.
V − R color plane. The color of the Sun is exactly coincident
with that of 2005 UD.

2000; Meech et al., 2004). We note that the optical colors of2003 EH1 are significantly less red than the average colorsof cometary nuclei (Jewitt, 2002; Lamy et al., 2004). This
could be a result of past thermal processing when the object had a perihelion far inside Earth’s orbit. Indeed, theweighted mean color of 8 near-Sun asteroids having perihelion distances ! 0.25 AU (subsolar temperatures " 800 K)
is V-R = 0.36±0.01 (Jewitt, 2013), consistent with the colorof EH1 (cf. section 8.5.4).
The optical colors measured for 169P/NEAT are less red
than D-type objects, as found in normal cometary nuclei
and Trojans, but similar to those of T- and X- type asteroids (Figures 8.13 and 8.14). The near-infrared spectrum
measurement (0.8–2.5 µm) classified 169P as a T-type asteroid based on the Bus taxonomy with pv =0.03±0.01 (DeMeoand Binzel, 2008). The T-type asteroids represent slightly
redder-sloped visible wavelength spectra than those of Ctype. Perhaps a refractory rubble mantle has formed on the

12Kasuga & Jewitt

Figure 8.15 Diameter vs. rotational period for asteroids (dots)
and the parent bodies of meteoroid streams (large circles). The
asteroid data are taken from Chang et al. (2015) while theparent body parameters are listed in Table 8.2. The horizontal
dotted line shows the spin barrier period of ∼2.2 hr (e.g. Warner
et al., 2009) and asteroid 1566 Icarus rotates nearby.

consider rotational instability as a possible cause of theirpast separation.
In principle, rotation rates of asteroids can be accelerated to critical limits by torques exerted from solar radiation through the YORP eﬀect (Vokrouhlický et al., 2015).
The YORP e-folding timescale of the spin, τY , is estimatedfrom the ratio of the rotational angular momentum, L, to
the torque, T . The relation may be simply expressed as2
τY ∼ K De2 Rh(Jewitt et al., 2015), where K is a constant,
De is the asteroid diameter (km) and Rh is the heliocentric
distance (AU). The value of constant K is sensitive to many
unknown parameters (the body shape, surface texture, thermal properties and spin vector of the asteroid and so on),but can be experimentally estimated from published measurements of YORP acceleration in seven well-characterized
asteroids (Table 2 from Rozitis and Green, 2013). Scaling K
to the bulk density of 1566 Icarus ρ = 3400 kg m−3 (V-type
or ordinary chondrite from Wilkison and Robinson, 2000;
Britt et al., 2002) and its rotation period Prot =2.273 hr, wefind K ∼ 7×1013 s km−2 AU−2 . The approximation is represented as (cf. Equation (3) of Jewitt et al., 2015),
τY (Myr) ≈ 2

169P surface, driven by volatile sublimation, and red matterhas been lost (Jewitt, 2002).
Asteroid 1566 Icarus is taxonomically classified as a Q- or
V-type (Figures 8.13 and 8.14). These types suggest thermal
evolution (perhaps at the level of the ordinary chondrites)
relative to the more primitive carbonaceous chondrites. The
formation process of the associated complex is unknown, but
we speculate that processes other than comet-like sublimation of ice are responsible.

8.5.2 Dust Production Mechanism: Example of
1566 Icarus and 2007 MK6Here we consider possible dust production mechanisms fromasteroidal parents. A diameter – rotation plot compiled fromthe data in Table 8.2 is shown in Figure 8.15. The rotational
period of 1566 Icarus (Prot =2.273 hr) is near the spin barrier
period of ∼2.2 hr (Warner et al., 2009; Chang et al., 2015).
Asteroids rotating near or faster than this barrier are presumed to have been destroyed when centrifugal forces haveovercome the gravitational and cohesive forces binding them
together (Pravec et al., 2008).
The aftermaths of recent and on-going asteroid break-up
have been identified observationally (e.g. P/2013 R3, Jewitt
et al., 2014, 2017) and studied theoretically (Hirabayashi
et al., 2014). Additionally, diﬀerent mechanisms can operate
together. Rotational instability in P/2013 R3, for instance,might have been induced by YORP torques, or by outgassing
torques from sublimated ice, or by a combination of the two.
Thermal disintegration, electrostatic ejection and radiationpressure sweeping may all occur together on near-Sun object
3200 Phaethon (Jewitt and Li, 2010). Figure 8.15 shows that
1566 rotates near the ∼2.2 hr spin barrier period, implying
a rotational breakup in the past. Both 1566 and its possible
fragment 2007 MK6 have small perihelia (q ∼0.19 AU). We

!

De
1 km

"2 !

Rh
1 AU

"2

.

(8.4)

For 1566 with De =1 km orbiting at Rh ∼1.08 AU, Equation (8.4) gives τY ≈ 2 Myr. This is two orders of magnitude
smaller than the collisional lifetime of 1-km near-Earth asteroids (Bottke et al., 1994), suggesting that YORP torquespin-up is plausible.
Asteroids rotating faster than the spin-barrier cannot beheld together by self-gravitation only, but require cohesivestrength (e.g. Scheeres et al., 2010). The cohesive strength
at rotational breakup of a body can be estimated by thedispersed fragmental sizes, initial separation speed, and thebulk density using Equation (5) of Jewitt et al. (2015),
S∼ρ

#

′

De
De

$

(∆v)2 .

(8.5)

Both fragmental asteroids (1566 and 2007 MK6 ) are as′
sumed to have the same bulk density ρ, De and De are diameters of 1566 and 2007 MK6 respectively (see Table 8.2),
and ∆v is the excess velocity of escaping fragments, assumed
comparable to the escape velocity from 1566. Adopting the′same value for ρ (see above) and substituting (De /De ) =0.18 (the diameter ratio between MK6 and 1566), and ∆v= 0.69 m s−1 , we find S ∼ 290 N m−2 .
This small value is comparable to strengths ∼10–
100 N m−2 modeled by a rubble-pile asteroid bounded
by weak van der Waals forces (reviewed in Scheeres andSánchez, 2018), but more than five orders of magnitude
smaller than the values of typical competent rocks (107 –108 N m−2 ). A rotational break-up origin of 1566 and MK6
is possible provided they have a weak, rubble-pile structure, as is thought likely for a majority of kilometer-sizedasteroids as a result of past, non-destructive impacts.
Several processes could eject dust from the surface of 1566.
Firstly, thermal disintegration can be induced by thermal
expansion forces that make cracks on the surfaces of aster-Asteroid–Meteoroid Complexes
oids and produce dust particles. The characteristic speeds ofdust particles produced by disintegration can be derived byconversion from thermal strain energy into kinetic energy ofejected dust particles. The necessary conversion eﬃciency,
η, is given by (cf. Equation (3) of Jewitt and Li, 2010)
η∼

% v

e

α δT

&2 % ρ &
Y

,

(8.6)

where, ve = 0.69 m s−1 is the escape velocity from 1566, α ∼10−5 K−1 is the characteristic thermal expansivity of rock
(Lauriello, 1974; Richter and Simmons, 1974), δT ∼450 K is
the temperature variation between the q and Q, and Y =(1–10) × 1010 N m−2 are typical Young’s moduli for rock(Pariseau, 2006, p.474). With ρ as above we find η " 0.1–
1 % is needed for the velocities of ejected dust particles tosurpass the escape velocity. This very small value of conversion eﬃciency is suﬃcient for most dust particles producedby thermal disintegration to be launched into interplanetary
space.
Secondly, electrostatic forces caused by photoionization
by solar UV can eject small particles. The critical size for a
1 km-diameter asteroid is ae ! 4 µm (Equation (12) Jewitt
et al., 2015). Millimeter-sized particles cannot be electrostatically launched and this process may contribute little ornothing to meteoroid stream formation.
Finally, radiation pressure sweeping can remove small particles from an asteroid once they are detached from thesurface by another process (i.e. once the surface contact
forces are temporarily broken). Radiation pressure sweeping
is most eﬀective at small heliocentric distances. The criticalsize to be swept away, arad (µm), is estimated by equating
the net surface acceleration (gravitational and centripetal)
with the acceleration due to radiation pressure, given by
Equation (6) of Jewitt and Li (2010)arad ∼

3 g⊙
2
2 π RAU f 1/2 De

!

Gρ
3π
− 2
f2
Prot

"−1

,

(8.7)

where, g⊙ is the gravitational acceleration to the Sun at
1 AU, RAU is the heliocentric distance expressed in AU, f is
the limit to the axis ratio (=a/b), G is the gravitational constant. We substitute g⊙ = 0.006 m s−2 , RAU = 0.187 (Table 8.1), f =1.2 (Table 8.2), G = 6.67 × 10−11 m3 kg−1 s−2and adopt the same values of De , ρ and Prot (see above) into
Equation (8.7), then obtain arad ∼ 4,500 µm ≈ 5 mm. Thissize is large enough to contribute meteoroid-sized particles
to a stream (see 8.6.3, Meteoroid Streams).
In summary, asteroid 1566 Icarus is a possible product ofrotational breakup and, given its small perihelion distance,
potentially experiences a mass loss process similar to thoseinferred on Phaethon. Near-perihelion observations of 1566
and/or 2007 MK6 may indeed show Phaethon-like mass-loss.
8.5.3 End State
Active objects on comet-like orbits with TJ < 3.08 (Table 8.1
and Figure 8.1) are presumed to be potential ice sublimators.
The timescales for the loss of ice from a mantled body τdv ,for the heat propagation into the interior of a body τc , anddynamical lifetime of short-period comets τsp ∼ 105 –106 yr

13(Duncan et al., 2004) can be compared to predict an object’s
end state (Jewitt, 2004b).
The τdv is calculated using ρn De /2f (dm/dt), where ρn =600 kg m−3 is the cometary bulk density (Weissman et al.,2004) and f = 0.01 is the mantle fraction (A’Hearn et al.,1995). The orbit with averaged ā ∼ 2.7 AU and ē ∼ 0.8 fromthe seven objects has a specific mass loss rate of water ice
dm/dt ! 10−4 kg m−2 s−1 (Figure 6, in Jewitt, 2004b). Thenwe find τdv " 104 De in yr, where De is an eﬀective diameterin km (Table 8.2). The τc ∼ 8.0×104 De2 in yr is derived fromthe equation of heat conduction given by De2 /4κ, where κ =10−7 m2 s−1 is the assumed thermal diﬀusivity of a porous
object. The critical size of object to form an inactive, devolatilized surface is constrained by the relation τdv ! τc ,which gives De " 0.13 km. Likewise, the size capable of containing ice in the interior of a body for the dynamical lifetime
is given by the relation τc " τsp , which gives De " 1.1 km.
Most comet-like objects are expected to be dormant, withice depleted from the surface region, but potentially still
packed deep inside. We note that 2003 WY25 of the Phoenicids is on its way to the dead state due to its small size
(De = 0.32 km). The τc ∼ 8,000 yr suggests that solar heatcan reach into the body core approximately 1 or 2 orders ofmagnitude sooner than the end of the dynamical lifetime.
The core temperature around the orbit (Jewitt and Hsieh,2006), Tcore ∼ 180 K, exceeds the sublimation temperatureof water ice 150 K (Yamamoto, 1985). The extremely weak
activity of 2003 WY25 may portend its imminent demise (cf.
section 8.4.1).
8.5.4 Lidov-Kozai Mechanism2
The Lidov-Kozai mechanism works on the secular dynamics of small solar system objects. Large-amplitude periodic
oscillations of the e and i (in antiphase) are produced,
whereas the a is approximately conserved, while the ω librates around π/2 or 3π/2 if C2 < 0 or circulates if C2 > 0(Kozai, 1962; Lidov, 1962).
Perihelia can be deflected into the vicinity of the Sunby this mechanism (on timescale ∼1000s of years), perhaps
causing physical alteration (or even breakup) due to enormous solar heating (Emel’yanenko, 2017). We find all parent
bodies stay in the circulation region (C2 > 0, see Table 8.1).
The minimum perihelion distance qmin can be computed
using maximum eccentricity, emax , given by (Equations (5)and (28) of Antognini (2015))
emax ='1−(1
ζ − ζ 2 − 60 C1 ,
6
%

&

(8.8)

where ζ = 3+5(C1 +C2 ) (Equation (31) of Antognini, 2015),
C1 and C2 are from Table 8.1. With the obtained qmin ,we find the Geminids (PGC), Quadrantids and Sekanina’s(1973) Taurid-Perseids complexes are near-Sun objects (Table 8.3). Amongst them, 2003 EH1 turns itself into a nearSun object with qmin ∼ 0.12 AU (cf. section 8.3.2), albeit
2 In Chapter 7 (Vaubaillon et al., 2019) the Lidov-Kozai mechanism focuses on secular changes in e, i and ω to find how e andω relate to whether an orbit intersects Earth’s orbit to produce a
meteor shower.

14

Kasuga & Jewitt

Table 8.3. Closest Approach to the Sun by the Lidov-Kozai Mechanism
Complex
Geminids

Quadrantids
Taurids-Perseids

ObjectPhaethon2005 UD
1999 YC
2003 EH1
96P/Machholz 1
1566 Icarus
2007 MK6

emax a
0.90
0.90
0.89
0.96
0.99
0.85
0.85

qmin b
0.13
0.13
0.16
0.12
0.03
0.16
0.16

TpeakcRef.d
1100∼0.1311100 0.13–0.141
1000
1100
∼0.1222300 0.03–0.053
1000∼0.1741000∼0.174

Notes: Near-Sun objects have perihelia ! 0.25 AU.
Maximum eccentricity (Equation (8.8))bMinimum perihelion distance (AU) estimated by ≃ a(1 − emax ), where a is from Table 8.1.
c
Peak temperature at qmin (K)dReferred minimum perihelion distance (AU) from numerical integrations and analytical methods
1
Ohtsuka et al. (1997, 2006) (cf. Figure 8.4) 2 Neslušan et al. (2013b); Fernández et al. (2014) (see section 8.3.2) 3 Bailey
et al. (1992); Sekanina and Chodas (2005); Abedin et al. (2018) 4 Ohtsuka et al. (2007) (cf. Figure 8.11)
aq ∼ 1.2 AU at present (see Figure 8.1). The peak temperature " 1000 K is similar to that experienced by Phaethon (cf.
1566 Icarus), and could likewise cause strong thermal anddesiccation stresses, cracking and alteration in EH1 , withthe release of dust (Jewitt and Li, 2010; Molaro et al., 2015;
Springmann et al., 2018). This example reminds us that,even in objects with q presently far from the Sun, we cannot exclude the action of extreme thermal processes in pastorbits.

8.6 Meteors and Streams
8.6.1 Na Loss: Thermal
Sodium loss in Geminid meteors results from the actionof a thermal process in or on the parent body, Phaethon(section 8.3.1). Čapek and Borovička (2009) calculated thetimescale for thermal depletion of Na from an assumed initial solar value down to 10% of solar abundance in Geminid meteoroids during their orbital motion in interplanetary
space (using assumed albite (NaAlSi3 O8 ) and orthoclase(KAlSi3 O8 ) compositions). With particle diameters ≥ mmscale, they found depletion timescales " 104 –105 yr (Figure 8.16), some 1 to 2 orders of magnitude longer than thestream age of ! 103 yr. On the other hand, the dynamical
lifetime of Phaethon, while very uncertain, is estimated tobe ∼ 3 × 107 yr (de León et al., 2010). This is surely longcompared to the age of the Geminid stream and long enoughfor Na to be thermally depleted from Phaethon.
In principle, other processes might aﬀect the sodium
abundance. Sputtering by the solar wind, photon stimulated
Na desorption (on Mercury and the Moon) (McGrath et al.,
1986; Potter et al., 2000; Killen et al., 2004; Yakshinskiy and
Madey, 2004) and cosmic ray bombardment (Sasaki et al.,2001) have been well studied. These processes act only on
the surface, and are ineﬃcient in removing Na from deeper
layers (Čapek and Borovička, 2009).

Figure 8.16 Timescale for escape of 90% of initial Na content
from Geminid meteoroids in the stream, as a function of themeteoroid size. The low diﬀusion data for Na in orthoclase wasused as a limit for the slowest loss, while the faster diﬀusion
data for Na in albite (10× higher than for orthoclase) is
considered as a more realistic value for the Geminids. The timeinterval in the shaded region corresponds to the estimated age ofthe Geminid meteoroid stream (1000–4000 yr). From Čapek andBorovička (2009).

8.6.2 Abundances vs. Intensity Ratios
In the interpretation of meteor spectra, two types of evaluation methods are employed in the literature. Some investigators calculate elemental abundances while others use simple
line intensity ratios. Here we describe advantages and disadvantages hidden in both methods.
Elemental abundances are quantitatively utilized for comparing with those of meteor showers and the solar abundances (Anders and Grevesse, 1989; Lodders, 2003). The
derivation of abundances is influenced by the complicatedphysics of the ionized gas at the head of the meteor(Borovicka, 1993; Kasuga et al., 2005b). For example, theSaha equation is needed to calculate the neutral vs. ion balance, but this depends on the assumption of a plausible excitation temperature, Tex , for the emitting region. This is par-

Asteroid–Meteoroid Complexes
ticularly important for Na, which has a smaller first ionization energy (≈ 5.1 eV) compared to other species (e.g. Mg: ≈7.6 eV, Fe: ≈ 7.9 eV). The selection of the appropriate Tex is
problematic. Fireball-like spectra have been suggested to be
the combination of some thermal components with typicalTex ∼ 5,000, 8,000 and 10,000 K, respectively (Borovicka,
1993; Kasuga et al., 2005b, 2007b). Borovicka (1993) considered Ca ii lines in the hot components (Tex ∼ 10,000 K)
and estimated electron density from the radiating volume of
the meteor in the direction of its flight. This implies that
Ca ii may not reflect the actual electron density from theirspectral emission profile. In Kasuga et al. (2005b), on theother hand, Ca ii is taken from the main component (Tex
∼ 5,000 K) instead of the hot one, and they derive electron
density reflected upon the measured spectral profile.
The definition of hot component theory has to satisfy theequality of total metal abundances (Ca/Mg) and pressureof the radiant gas between the main and the hot components (Borovicka, 1993). However, the relation does not fit
in most spectral data (e.g. Leonids). The Saha function isused to verify the definition but mostly finds negative values of electron density, which is clearly unrealistic (Kasuga
et al., 2005b). This proves that the original hot component
theory may go against its own definition (Borovicka, 1993).
Kasuga et al. (2005b) suggest that the Ca ii lines do not always belong to the hot component, but instead to the main
component. Plausibility is also found in their low excitationenergies (Ca ii ≈ 3.1 eV), which are compatible with thoseof other neutral metals (e.g. Na i ≈ 2.1 eV) identified in themain component (Kasuga et al., 2007c). On the other hand,
the hot component primarily consists of species with highexcitation energies " 10 eV. Accordingly the Ca ii lines aremost likely to belong to the main component.
Given the diﬃculties in calculating absolute abundances,many researchers have employed simple line intensity ratios(e.g. Borovička et al., 2005). The method analyzes neutral
atomic emission lines of Na i, Mg i, and (weak) Fe i only. Notethat the intensity ratios do not directly reflect the elemental
abundances due to no consideration for excitation propertiesof the elements and ions, including electron densities (e.g.
Borovička, 2001; Borovička et al., 2005; Koten et al., 2006).
Laboratory spectroscopic experiment proves that intensity
ratios are not informative of the abundance and cannot be
used to determine the meteorite analogue (Drouard et al.,2018). Line ratios can be used, at best, to study the trend ofelemental content in meteors which happen to possess similar physical conditions (including similar entry velocities,
similar strengths and similar excitation temperatures). Lineintensity ratios can suggest the trend of elemental content
in meteor showers but cannot provide the abundances.
8.6.3 Meteoroid Streams
Physical properties of meteoroids in meteoroid streams ordebris in dust trails whose orbits do not intersect that of
the Earth can be revealed by thermal and optical observations. The streams (or trails) and meteoroids mostly consist of mm to cm-scale compact aggregates, as estimatedfrom the ratio of solar radiation pressure to solar gravity of15

10−5 ! β ! 10−3 (e.g. Sykes and Walker, 1992; Reach et al.,2000, 2007; Ishiguro et al., 2002; Sarugaku et al., 2015).
Meteor ablation models in the Earth atmosphere are alsoavailable to estimate the size of meteoroids (Bronshten,
1983; Ceplecha et al., 1998). The classical models find typicalmeteors are 10µm – 10 cm in size, however, with uncertainties caused by various parameters (e.g. luminous eﬃciency,
ablation coeﬃcient, and fragmentation) which are sensitive
to the meteoroid entry speed, tensile strength and brightness(Ceplecha et al., 1998; Babadzhanov, 2002). Faint meteors
are estimated to be 10µm ∼1 mm, but model improvements
are needed to better represent fragmentation (CampbellBrown and Koschny, 2004; Borovička et al., 2007).
8.6.4 Zodiacal Cloud
The zodiacal cloud is a circumsolar disk consisting of small
dust particles supplied by comets and asteroids. The total
mass is ∼ 4 × 1016 kg, most (∼90%?) of which is suppliedby JFC disruptions and the rest by Oort cloud comets (!
10%) and asteroids (! 5%) (Nesvorný et al., 2010, 2011;
Jenniskens, 2015). The supply rate needed to maintain thezodiacal cloud in steady-state is 103 to 104 kg s−1 (Nesvorný
et al., 2011).
The fate of dust particles released from comets into the zodiacal cloud is traceable (cf. Grun et al., 1985). Sub-micronparticles, with β > 0.5, are immediately blown out of thesolar system on hyperbolic orbits by radiation pressure andare referred-to as β-meteoroids (Zook and Berg, 1975; Grunet al., 1985). The JFCs frequently disintegrate when nearperihelion and form dust trails or meteoroid streams, mainly
consisting of mm – cm sized dust particles (see section 8.6.3).
The collisional lifetime of mm-scale particles at 1 AU is estimated to be τcol " 105 yr, modeled with the orbital distribution of sporadic meteors measured by radar (Nesvorný
et al., 2011) (cf. τcol ∼ 104 or 105 yr, Grun et al., 1985; Soja
et al., 2016; Yang and Ishiguro, 2018). On the other hand,
Poynting-Robertson (P-R) and solar-wind drag cause dust
particles to spiral down to the Sun. The P-R drag timescale,τPR , to drift down from ā ∼ 2.7 AU with ē ∼ 0.8 (objects
with TJ < 3.08, see Table 8.1 and Figure 8.1) to 1 AU aroundthe Earth orbit (e ∼ 0.017) is calculated using the equationin Wyatt and Whipple (1950) (cf. Dermott et al., 2002),
τPR ≃

730
yr,
β(1 + sw)

(8.9)

where β ! 10−3 (see 8.6.3) for dust particles of radius "1 mm, with bulk density of 600 kg m−3 (Weissman et al.,2004), and sw = 0.3 is eﬃciency of solar-wind drag on aparticle normalized to the P-R drag eﬀect (Gustafson, 1994).
The estimated P-R drag lifetime is τPR " 6×105 yr. This being somewhat longer than τcol , the mm-scale dust particles
are subject to collisional disruption while spiralling down to
the Sun by P-R drag. As a result of competition between
these two eﬀects (loss to Poynting-Robertson at small sizes,
loss to collisional shattering at large sizes), the 100–200µm
dust particles are the most abundant in the zodiacal cloud(Love and Brownlee, 1993; Grun et al., 1985; Ceplecha et al.,
1998; Nesvorný et al., 2010).

16

Kasuga & Jewitt

Nesvorný et al. (2011) noted that the collisional lifetime
for mm-scale particles is long compared to the plausible lifetimes of most meteoroid streams (! 104 yr). They speculatethat cm-scale particles are sources of smaller dust grains.
Centimeter-scale particles are also released from JFCs. The
sequence may result in a population of mm-sized or smaller
particles which could be more resistant to collisions. Recent
meteor observations suggest a relative lack of large particles (∼7 mm) (Jenniskens et al., 2016b), and also suggestthat some of these larger particles disappear on timescales∼104 yr, not from collisions, but from other processes. Moorhead et al. (2017) finds a two-population sporadic meteoroid bulk density distribution suggesting that the physical character of freshly ejected dust particles could be altered over time. As another example, Rosetta dust collectors sampled both very pristine fluﬀy aggregates and compact particles (" 4 cm in diameter) with a possible rangeof the dust bulk density from 400 to 3000 kg m−3 (Rotundi
et al., 2015). This variety could have resulted from aggregate fragmentation into the denser collected grains as thespacecraft approached, while the packing eﬀect is proposed
as a plausible mechanism theory for fluﬀy dust particles released from comets (Mukai and Fechtig, 1983). The particle
size could be reduced on a timescale of 104 –105 yr, comparable with the meteoroid stream lifetime. The eﬀect makes
the bulk density increase from 600 kg m−3 to 3000 kg m−3 ,corresponding to shrinking the particle size approximately
down to half. This could be a potential explanation for disappearing larger-scale dust particles in the meteoroid streams.

8.7 Summary and Future WorkIn the last decade, a growing understanding of parent bodies and meteoroid streams has been achieved by combining new physical observations and dynamical investigations.
Still, even where the associations are relatively clear, mostcomplexes have multiple potential parent bodies and it remains unclear how the streams were formed.
Observationally, a major challenge is posed by the difficulty of measuring the physical characteristics of parent
bodies, most of which are faint by virtue of their small size
(typically ! a few km). They are also frequently observationally inaccessible because of their eccentric orbits, whichcause them to spend most of the time far away near aphelion.
Long-term surveillance of NEOs around their entire orbits
might better reveal how and when parent bodies disintegrate
and produce debris.
Dynamically, there are at least two challenging problems.
One concerns the identification of parent bodies through thecomparison of the orbital elements of meteoroids and potential parents by a D-criterion. Such methods work best forparents of young streams, where the eﬀects of diﬀerential
dynamical evolution are limited. However, in older systems,
the dynamical elements have evolved enough to seriously undercut the use of the D-criteria. For this reason, for example,
numerous Taurid parent bodies continue to be proposed. A
key objective is to find a way to more reliably associateolder meteoroid streams with their parent bodies. A second problem is the use of long-term dynamical simulations
in which the initial conditions and/or potentially important
non-gravitational eﬀects are partly or wholly neglected.
We list key questions to be answered in the next decade.
1. Geminids: What process can act on ∼1000 yr timescales
to produce the Geminid meteoroid stream? Phaethon appears dynamically associated with at least two kilometersized asteroids (2005 UD and 1999 YC) suggesting a past
breakup or other catastrophe. But the likely timescale forsuch an event is ≫1000 yr. What caused the breakup
and is it related to the Geminids? How many other
PGC-related objects await discovery? Are Geminids represented in the meteorite collections and, if so, how can
we identify them?
2. Quadrantids: Presumed parent 2003 EH1 is currently inactive, but was recently as close to the Sun as is Phaethon
at perihelion. Can residual mass loss in EH1 be detected?
Is the Quadrantid sodium abundance depleted as a resultof the previously smaller perihelion? What physical difference is to be found in 96P which has a near-Sun orbit
even now?
3. Capricornids: Several parent bodies have been proposed
including both active comets and inactive asteroids. Did
they originate from a common precursor? Is asteroid
2017 MB1 related?
4. Taurids: The prime parent body is 2P/Encke but numerous additional parents with diverse properties continue
to be proposed (mostly based on the D-criterion). How
can we establish the relevance of these other objects to
the Taurid stream? Can activity be detected? Is the Dcriterion appropriate to judge?
5. Sekanina’s (1973) Taurid-Perseids: Do 1566 Icarus and2007 MK6 share common physical properties? Is thesodium abundance in Taurid-Perseids depleted by solar
heating due to the small perihelion?
6. Phoenicids, Andromedids and other minor complexes:
Fragmentation is expected to produce a wide range ofobject sizes, with many bodies being too small to havebeen detected so far. What role can be played in the
search for stream-related bodies by upcoming deep sky
surveys, like the Large Synoptic Survey Telescope?
7. Which is the better index, the D-criterion (e.g. DSH ) or
the dynamical invariants (C1 , C2 )?
8. How many near-Sun objects, driven by the Lidov-Kozai
mechanism, exist?
9. What more can we learn from meteor spectroscopy, particularly of faint meteors?
10. Sporadic meteoroid populations tend to lose large dust
particles (sizes " 7mm) on timescales of 104 yr. Why?

Acknowledgments
We are grateful to David Asher for his enthusiastic guidance to improve this chapter. TK thanks Takaya Okamoto
for wholehearted assistance with this study and also Junichi Watanabe, Hideyo Kawakita, Mikiya Sato and Chie

Asteroid–Meteoroid Complexes
Tsuchiya for support. We appreciate David Čapek, Jiřı́
Borovička, Man-To Hui, Jing Li and Michael S. P. Kelley
for figure contributions. We acknowledge reviews by Tadeusz
Jopek and an anonymous reviewer. Lastly, we thank Galina
Ryabova, again David Asher and Margaret Campbell-Brown
for organizing this Meteoroids-book project.

17

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