Free Access
Issue
A&A
Volume 647, March 2021
Article Number A37
Number of page(s) 13
Section Galactic structure, stellar clusters and populations
DOI https://doi.org/10.1051/0004-6361/202039390
Published online 04 March 2021

© ESO 2021

1. Introduction

With the second data release of the Gaia space mission (Gaia Collaboration 2018b), a catalogue comprising the full range of 6D phase space information has become available (Katz et al. 2019). This subset has enabled detailed studies of interconnected dynamical processes taking place in the Milky Way. For example, substructures in integrals-of-motion spaces of nearby disc stars have been related to the Galactic bar or spiral arms (e.g., Dehnen 2000; Monari et al. 2017, 2019; Khoperskov et al. 2019; Trick et al. 2019; Hunt et al. 2019), or to interactions with satellite galaxies (Antoja et al. 2018; Laporte et al. 2018; Laporte & Minchev 2019).

Furthermore, the assembly history of the Milky Way can be studied by identifying and characterising substructures in integrals-of-motion space, as first put forward by Helmi & de Zeeuw (2000). Structures comprising tens to hundreds of stars have been linked to accreted dwarf galaxies whose debris can still be traced in the solar neighbourhood (Helmi et al. 1999, 2017; Chiba & Beers 2000; Klement et al. 2008, 2009; Williams et al. 2011; Myeong et al. 2018; Koppelman et al. 2018, 2019b; Yuan et al. 2020; Naidu et al. 2020, see also the reviews of Newberg & Carlin 2016; Klement 2010). Besides these small groups, recently the relics of a massive dwarf galaxy (now known as Gaia-Enceladus) have been identified in the same kind of spaces by Helmi et al. (2018, see also Belokurov et al. 2018).

With the rise of machine learning, it has become popular to deploy automated classification algorithms to search for substructures (Necib et al. 2020; Borsato et al. 2020; Du et al. 2019; Koppelman et al. 2019b; Ratzenböck et al. 2020; Yuan et al. 2020). Substructures are not necessarily always accreted (e.g., Gómez & Helmi 2010; Gómez et al. 2013; Jean-Baptiste et al. 2017), and therefore the interpretation of their origin is not always straightforward. In this work, we investigate structures reported in Rmax − zmax space (e.g., Haywood et al. 2018), supplemented by information obtained from the orbital frequencies of the stars.

Substructures identified in the space characterised by Rmax − zmax may be amenable to intuitive interpretation. Schuster et al. (2012) noted that the low-[α/Fe] stars in the halo, identified in Nissen & Schuster (2010), have orbits that reach out to much larger Rmax and zmax than the high-[α/Fe] stars. A similar conclusion was reached by Haywood et al. (2018), who identified distinct ‘wedges’ in this space. These wedges are also present when using the chemically defined halo sample of Chiba & Beers (2000), with updated astrometry from Gaia DR2, ruling out that they are due to a kinematic bias. Haywood et al. (2018) relate the wedges to the merger with Gaia-Enceladus and link them to the impulsive heating of an ancient Milky Way disc. The authors suggest that the largest gap may hint at a phase transition in the assembly history of the Milky Way, from a significant to a quiescent accretion phase. A different view is presented by Amarante et al. (2020), who argue that these wedges are due to transitions between orbital families. Based on smoothly sampled halo orbits, these authors attribute the structures in this space to ‘a natural consequence of resonant effects’. This view is supported by Schuster & Allen (1997), who report on a gap in the distribution in zmax that is reminiscent of the wedges in zmax with Rmax space. These authors ascribe the gap in the zmax distribution to chaotic diffusion, where the orbits ‘stick’ to one set of quasi-periodic orbits or another.

In this work, we further explored the link between the Rmax − zmax space and orbital resonances, as mapped by frequency space. The space of orbital frequencies is particularly interesting for several reasons. Firstly, as shown by Gómez & Helmi (2010) and Gómez et al. (2010), individual streams associated with an accreted galaxy can be easily identified here. The time of accretion is quite precisely encoded in this space and can be determined from the separation between two adjacent streams in frequency space. Secondly, Valluri et al. (2012) argued that the strengths and locations of resonances for halo stars in frequency space are both dependent on the stellar distribution function, as well as on the global shape of the halo. In principle, the orientation of the halo with respect to the disc can be determined from the distribution of stars in frequency space. Moreover, Valluri et al. argued that the diffusion rates of the orbital frequencies can help in distinguishing between the true and an incorrect potential. Finally, the mapping of frequency space is a powerful tool to establish the different orbital families present in an ensemble. For separable and non-degenerate potentials the frequencies are directly related to the actions. Frequencies have the advantage that they can be easily and precisely calculated numerically. Also, in near-integrable potentials, for which the actions do not exist, the frequencies trace the different orbital families accurately (e.g., Laskar 1993).

We investigated the distribution of stars in the solar vicinity in spaces associated with different orbital parameters – such as pericentre, apocentre, eccentricity, zmax and orbital frequencies – and explored the kinematics of the substructures identified in those subspaces. The paper is organised as follows. In Sect. 2, we introduce the data and present the selection criteria applied. Then, in Sect. 3, we identify substructures in subspaces of different orbital parameters. In Sect. 4, we perform an orbital frequency analysis and show how the structures in the various subspaces, in particular Rmax − zmax, are linked to resonances in frequency space. We reflect on these results in Sect. 5 and present our conclusions in Sect. 6.

2. Data

We used the subset of Gaia DR2 with line-of-sight velocities (Gaia Collaboration 2018d). This subset – also referred to as the Gaia radial-velocity spectrometer (RVS) subset – of more than 7 million stars thus contains accurate measurements of the 3D positions and 3D velocities of the stars. Gaia DR2 parallaxes are known to suffer from systematic parallax errors that vary with position, G-band magnitude, and stellar type. Therefore, we used Bayesian distances derived by McMillan (2018). The author derived distances by assuming an overall parallax offset of −29 mas with an RMS error of 43 mas for the Gaia RVS subset as found for the full Gaia DR2 release (Lindegren et al. 2018). We focus here on a local sample with d < 2.5 kpc of high-quality parallaxes (parallax_over_error > 5). These quality criteria leave us with a sample of 5 015 006 stars.

For these stars, we computed positions and velocities in a galactocentric cylindrical coordinate system (R, z, ϕ, vR, vz, vϕ). We placed the solar position at R = 8.2 kpc (McMillan 2017)1 and z = 0.014 kpc (Binney et al. 1997). The Local Standard of Rest (LSR) velocity (i.e. vc(R) = 232.8 km s−1) is, for consistency, also taken from McMillan (2017). We further set the motion of the Sun relative to the LSR to (U, V, W) = (11.1, 12.24, 7.25) km s−1 (Schönrich et al. 2010). Here, U expresses a radially inward motion, V, into the direction of Galactic rotation, and W, perpendicular to the Galactic plane and towards the Galactic north pole. We further define vϕ such that it is positive in the sense of Galactic rotation.

In this work, we focused on nearby thick-disc and halo samples. Stars belonging to these samples are ‘isolated’ based on their kinematics (see also Venn et al. 2004; Bensby et al. 2014; Posti et al. 2018). We selected stars that satisfy vmin < |v(x)−vc(R)| < vmax. For the thick-disc sample, we chose km s−1 and km s−1, while for the halo we only set a minimum speed of ( km s−1). These selections are illustrated in Fig. 1. The thick-disc sample comprises 216 672 stars, and the halo sample 17 704 stars. For the traditional definition of the halo, this consists of two main components, namely a hot thick-disc and a proper halo (Gaia Collaboration 2018a; Koppelman et al. 2018; Haywood et al. 2018; Di Matteo et al. 2019; Gallart et al. 2019). Therefore, we split the halo sample further (as shown in Fig. 1) into a ‘thick-disc tail’ (TDtail, blue dots) and a ‘pure-halo’ sample (pureHalo, black dots). We stress that this split is based purely on the kinematics of the stars, and both samples likely contain both accreted stars and those formed in-situ. We set the boundary at km s−1, such that the pureHalo sample contains 10 370 stars and the TDtail sample 7334 stars.

thumbnail Fig. 1.

Velocity distribution of the stars in our sample, highlighting the various kinematic selections. A logarithmic map of star counts is shown in the background (the smooth blue map). Stars within the cyan semi-circles are considered to belong to the thick-disc. The stars outside the outermost cyan curve are considered halo stars, they are split into a pure halo and a TDtail sample.

3. Results: Phase-space and integrals of motion

For most of the analysis presented in this paper, we adopt the Milky Way potential determined by McMillan (2017), unless stated otherwise. This axisymmetric potential consists of a stellar thin and thick-disc, HI gas disc, molecular gas disc, a flattened bulge, and a spherical halo component (NFW). This system has a virial mass of 1.3 × 1012 M, with 5.43 × 1010 M in stars.

We integrated the orbits of halo stars in this potential for 80 Gyr, and 40 Gyr for the thick-disc stars. These long timescales were chosen because of our interest in computing the orbital frequencies (see Sect. 4). Numerically estimating the frequencies requires densely sampled orbits with many periods, typically more than 20. Halo stars are integrated for a longer time than thick-disc stars because the latter, in general, have much shorter orbital periods. The orbits were integrated using AGAMA (Vasiliev 2019), which uses an adaptive time step (it uses the eighth-order Runge-Kutta DOP853 integrator). We checked that the energies of the stars do not deviate more than 0.01% of their initial values. The data is output at fixed time intervals of roughly 1.25 Myr and 2.5 Myr for the halo and thick-disc samples, respectively.

3.1. Velocity space

The local stellar distribution is enriched by several moving groups that are apparent as clumps in velocity space (vR, vz, vϕ). These structures might be related to stars born together or may be of dynamical nature (e.g., see Antoja et al. 2018, on the outer Lindblad resonance (OLR) mode due to the Galactic bar).

Figure 2 shows the velocity distribution of the three subsets that we analysed, as indicated by the insets. As is clearly apparent from this figure, these distributions are affected by the kinematic selection criteria used. We also highlight how the top (thick-disc) and middle rows (TDtail) look similar, whereas the bottom row (pureHalo) shows a different distribution, mostly in the middle and right columns corresponding to vϕ vs. vR and vϕ vs. vz, respectively. The TDtail component (middle row) consists of stars with ‘hotter’ thick-disc-like kinematics, and it is likely the result of the proto-disc being puffed up during the merger with Gaia-Enceladus (Helmi et al. 2018; Di Matteo et al. 2019).

thumbnail Fig. 2.

Different projections of velocity space for the various subsets considered in this chapter: thick-disc, TDtail, and pure halo from top to bottom. The lower density of stars present in these top panels for small values of vR and vz is due to the exclusion of thin-disc stars in our selection. We note that even our pureHalo subset contains some stars with ‘hot’ thick-disc-like kinematics (as evidenced by the arch-like structure in the bottom central panel). Red dashed curves in the bottom panel outline the structures that are mentioned in the text.

The strong, radially elongated feature at vϕ = 0 in the pureHalo sample (middle panel) is indicative of the debris of Gaia-Enceladus (cf. ‘the sausage’, Belokurov et al. 2018). The top and bottom parts of the crescent shape at vϕ ≈ 100 km s−1 imply that this sample still contains a small fraction of heated thick-disc stars. It can be seen as an extension of the distribution shown in the middle panel of the TDtail (middle row). The elongated structure in vR and the crescent shape have been outlined with thin, red dashed lines.

In the bottom-left panel, we also see a clear overdensity of stars in an arch near vz ∼ ±250 km s−1 for a range of values of vR. These stars define the two clumps in the bottom-right panel seen at (vϕ ∼ 120 km s−1, vz ∼ ±250 km s−1), and they correspond to the Helmi streams (Helmi et al. 1999). The Helmi streams are marked with thin, red dashed lines.

3.2. Characteristic orbital parameters

In a static axisymmetric potential, such as the one considered here, the integrals of motion energy (E) and angular momentum in the z-direction (Lz) are conserved. However, stars that share a common origin phase-mix as they evolve in the Galactic potential (Helmi & White 1999), which results in complex configurations in phase space (Tremaine 1999), their distributions in, for example, E − Lz space remain clumped.

We show the distribution of the stars in the three samples in E − Lz space in Fig. 3 (see the left column). The middle and right columns show a combination of other orbital parameters that are sometimes used to identify substructures. The pureHalo sample (bottom row) shows by far the largest spread in the orbital parameters (we note, in particular, the larger dynamical range with regard to the top row). Its distribution is dominated by stars near Lz ∼ 0 kpc km s−1, which is the debris identified as Gaia-Enceladus (Koppelman et al. 2018; Helmi et al. 2018). Besides this prominent structure, we see contamination from the hot thick-disc and several small overdensities such as that tentatively linked to an accretion event labelled Sequoia (Lz ∼ −2000 kpc km s−1, E ∼ −120 000 km2 s−2), and (Lz ∼ −1200 kpc km s−1, E ∼ −170 000 km2 s−2) for Thamnos (see Myeong et al. 2019 and Koppelman et al. 2019b, for more details).

thumbnail Fig. 3.

Orbital parameters spaces for the different subsets defined in Sect. 2. From top to bottom: thick-disc, TDtail, and pureHalo. The left panels correspond to the most often used space to identify merger debris (i.e. Lz,  E). The substructure seen near Lz = 250 kpc km s−1 for highly bound energies in the middle-left panel is the globular cluster M4 (also known as NGC 6221), whose extent is largely the result of limitations in the Gaia DR2 astrometry for very dense stellar fields. The subspaces shown in the middle and right panels depict a large amount of structure, which appears to be independent of the subset or population considered. We note that the top row shows a smaller dynamical range than the other two rows.

In the middle panels of Fig. 3 we plot Rmax vs. zmax, while in the right panels we show the eccentricity vs. zmax/Rmax (i.e. a proxy for orbital inclination). The eccentricity is defined as e = [ra − rp]/[ra + rp], where rp and ra denote the orbital pericentre and apocentre (measured as the minimum and maximum galactocentric distances over the whole integration interval). Particularly in these orbital parameters spaces, we see a large amount of structure. These structures look similar to the ‘wedges’ reported by Haywood et al. (2018) and computed for a different Milky Way potential than used here. Figure 3 shows that the structures are present in all the star samples explored, although with a different prominence. The fact that they are present independently of the population suggests that the phenomenon that causes them must be of a global nature and is unlikely to be due to individual (ancient) accretion events.

The shape and exact location of the structures shift slightly when changing the Galactic potential, but they appear to be generic. For example, we also integrated our samples using the potentials provided by Piffl et al. (2014) and the MWPotential2014 (Bovy 2015) and found similar features (see Appendix A and also Haywood et al. 2018).

To gain further insight, we also decided to explore a less generic but integrable potential, namely a Milky Way-like potential of Stäckel form (derived from Batsleer & Dejonghe 1994; Famaey & Dejonghe 2003). Potentials of this form are separable and contain only regular orbits. Figure 4 shows the distributions of the subspace Rmax vs. zmax for the stars in our various subsets now determined using a two-component Stäckel Milky Way potential of Famaey & Dejonghe (2003). This potential has a total mass of Mtot = 4 × 1011M, a disc mass with a fraction of k = 0.11 of Mtot, scale length ra = 7 kpc, and a flattening of the disc and halo of ϵd = 75 and ϵh = 1.02, and it provides a good representation of, for example, the Galactic rotation curve. Figure 4 reveals rather smooth and featureless distributions, especially in comparison to the previous figure. This implies that the structures seen in Fig. 3 must be induced by the non-integrability of the generic Galactic potentials used. The structures could thus be due to resonances and the result of chaotic diffusion, a hypothesis that we explore further in Sect. 4.

thumbnail Fig. 4.

Rmax vs. zmax for the thick-disc (left), TDtail (middle), and pureHalo (right) populations, computed for an integrable two-component Milky Way Stäckel potential. In all subsets, the amount of structure has drastically decreased. The clear overdensity of stars in the TDtail sample and at Rmax ∼ 6.5 kpc and zmax ∼ 0.5 kpc corresponds to the globular cluster M4 and is thus of no specific interest.

3.3. Comparison to known substructures

To further strengthen this preliminary conclusion, we investigated the distribution of stars associated with known accreted substructures in the orbital parameter spaces just explored. Many candidate substructures have been identified since Gaia DR2, such as Gaia-Enceladus (Helmi et al. 2018), Sequoia (Myeong et al. 2019), and Thamnos (Koppelman et al. 2019b), and some have also been confirmed, such as the Helmi streams (Helmi et al. 1999; Koppelman et al. 2019a).

We took the stars that belong to these substructures according to Koppelman et al. (2019b) and noted their intersection with our pureHalo sample. In this way, we find 1186 Gaia-Enceladus stars, 54 members of the Helmi streams, 92 in Sequoia, and 1098 Thamnos stars.

The velocity distribution of the stars in these various substructures is shown in the top panels of Fig. 5. The bottom-left panel directly reflects the selection criteria used to identify them (where a lower limit in the energy E has been somewhat artificially imposed on the Gaia-Enceladus stars to reduce overlap with the low-energy halo and the tail of the thick-disc (e.g., Feuillet & Feltzing 2020). This constraint depletes the number of stars with small vR and eccentricities e < 0.8. On the other hand, the middle and right-hand-side panels of the bottom row show that these accreted substructures are also themselves split into substructures or wedges, just like the stars in the various subsets considered in the previous section. For example, the nearly horizontally aligned gaps in (Rmax, zmax)-space continue to be very prominent. Given the (tentative) range of accretion times for the various objects (∼6 to 10 Gyr ago), and that their debris seems to be influenced in the same way, this analysis further supports the idea that the origin of the features seen in Fig. 3 must lie in a global phenomenon affecting the whole Galaxy.

thumbnail Fig. 5.

Distribution of accreted substructures in the pureHalo subset following the identification by Koppelman et al. (2019b): in velocity space (top) and in orbital parameters spaces (bottom).

4. Analysis: Orbital frequencies

4.1. Methods

In this section, we describe a frequency analysis we performed to explore whether resonances might play a role in the substructures seen in the orbital parameter spaces. We estimated the frequencies numerically using the SuperFreq package2 (Price-Whelan 2015, see also Price-Whelan et al. 2016), which is a Python implementation very similar to the Numerical Analysis of Fundamental Frequencies (or NAFF) code written by Valluri & Merritt (1998) and Valluri et al. (2010, 2012), which in itself is an implementation of the NAFF technique of Laskar (1990; 1993, see also Wang et al. 2016, who provide a comparison between the NAFF algorithm and another commonly used orbital frequency and classification code by Carpintero & Aguilar 1998). The software finds the fundamental frequencies by computing the Fourier spectra for the phase space coordinates (more precisely: of a complex time series) used to describe the orbit. Since the orbits are computed in an axisymmetric potential, we followed Valluri et al. (2012) in using a slightly different form of the cylindrical polar coordinates, namely the Poincaré symplectic polar variables. In this case, the complex time series were chosen as follows: fR = R + ivR, fz = z + ivz, . The frequencies in R and z are defined to be positive, while the sign of the ϕ-frequency corresponds to the direction of rotation (i.e. positive when aligned with the Galactic rotation)3.

In SuperFreq, the first fundamental frequency is the non-zero frequency with the highest amplitude. The algorithm then continues by moving down along the remaining frequency lines in the Fourier spectrum corresponding to the other coordinates, in order of decreasing amplitude. The next fundamental frequency must be different from the first one found, and so on. These fundamental frequencies Ω = (ΩR, Ωz, Ωϕ) are not necessarily equal to the dominant frequencies ω in the Fourier spectra of each coordinate. The fundamental frequencies found are also not necessarily equal to the frequencies in which the angle variables in action-angle coordinates vary with time. For regular orbits, the recovered fundamental frequencies will, however, be a linear combination of these ‘more fundamental’ (action-angle) frequencies (Valluri et al. 2012).

Resonant orbits are those for which the fundamental frequencies are commensurable. That is, a resonance is defined to exist if n ⋅ Ω = 0, for any vector n = (nz, nR, nϕ) comprising integer numbers (where at most one element is zero). In theory, there are an infinite number of resonances for combinations of arbitrarily large vector elements ni; here, we only consider those with |ni| ≤ 5. In a spherical potential, the true ‘fundamental’ frequencies corresponding to the nature of the potential satisfy Ωr/2 ≤ |Ωϕ| ≤ Ωr (Binney & Tremaine 2008), where the lower limit corresponds to any orbit in a homogeneous sphere (harmonic oscillator), and the upper limit corresponds to orbits in a Kepler potential (point mass). In the epicyclic approximation, for a flat rotation curve [vcirc(r) = const.]. In an axisymmetric disc potential, vertical oscillations for most disc stars are expected to have the shortest periods (Ωz is the largest in Ω).

The NAFF methodology provides reliable frequencies for regular orbits that have been traced for 20 − 30 periods, so we limited our analysis to stars with a minimum of 20 orbital periods. All thick-disc stars and almost all halo stars (99.76%) are integrated long enough to satisfy this criterion. Most halo stars are integrated for 100 − 600 orbital periods, while most thick-disc stars for 150 − 300 orbital periods.

4.2. Results

We commence the frequency analysis by showing the distribution of fundamental orbital frequencies for each of the populations or subsets considered in Fig. 6. The top row shows a density map, and because of the size of the sample, the colours indicate the logarithm of the number of stars, with bright orange being the densest and blue being the least dense. Resonances are clearly discernible in this space as well-populated thin straight lines. Most of the resonances are present in the ΩR vs. Ωz space (left column). The regions around these resonances are depleted of stars, which is common for non-integrable potentials and indicates the presence of chaotic zones (e.g., see Fig. 1 of Price-Whelan et al. 2016, for an illustration of this phenomenon). There are no very obvious resonances in Ωϕ vs. ΩR (middle panels, although these might be expected for thin-disc stars, particularly in relation to the Galactic bar, see Dehnen 2000).

thumbnail Fig. 6.

Frequency maps for the different populations. From top to bottom: thick-disc, TDtail, and pureHalo population. From left to right: cycling through the combinations of frequencies. The panels, especially those in the left column, show a clear presence of resonances in the form of straight lines and of depleted regions around them.

We find that roughly 10% of the stars in the thick-disc sample are on resonant orbits in ΩzR. For the TDtail and for the pureHalo samples, these fractions are much higher: namely 25% and 28%, respectively. Stars are said to belong to a certain resonance if the orbital frequency ratio for the pair of coordinates deviates at most 0.001 from the rational number of the resonance under consideration. To establish if the resonances are related to the structures we identified in the previous section, we now discuss the most dominant resonances in each of the subsets.

A selection of the most dominant resonances in the thick-disc sample is shown in Fig. 7 (top row). These are the Ωz : ΩR= 1:1 (blue), 2:1 (yellow), 1:2 (green), 3:2 (red), 3:4 (purple), and 5:4 (brown) resonances. The percentages of stars in each of these resonances are listed in Table 1. The middle row of Fig. 7 shows that the stars on these resonances occupy specific regions in the spaces of orbital parameters. There is a strong connection between the frequencies and features seen in these spaces. The overdensities first shown in Fig. 3 correspond directly to some of the highlighted resonances. On the other hand, depleted regions and gaps in the orbital parameter spaces indicate regions of unstable orbits. Such orbits are likely trapped by the resonances explored here, leaving empty regions around the corresponding frequencies. This is, for example, clearly seen for the Ωz : ΩR= 1:2 (in green) and Ωz : ΩR= 3:2 (in red) resonances. The bottom panels of Fig. 7 show that stars on a given resonance have a broad range of velocities. However, some resonances (e.g., the 1:2) occupy only a small region in certain projections of velocity space.

thumbnail Fig. 7.

Selected resonances in the thick-disc subset and their mapping onto other informative subspaces. Top row: frequency maps highlighting the main resonant families. Middle row: characteristic orbital parameters colour-coded according to the different resonances. Bottom row: their distribution in velocity space.

Table 1.

Percentage of stars in the various sub-samples associated with the listed Ωz : ΩR resonances.

Figure 7 shows that most of the resonances can be associated with a contiguous region in the orbital parameters spaces. However, this figure also reveals a higher complexity for the Ωz : ΩR= 1:1 and the Ωz : ΩR= 2:1 resonances. The stars of these families are found in two regions that surround the most prominent gap of Fig. 3, located near Rmax ∼ 8 kpc and zmax ∼ 2 kpc. This potentially means that these resonances are not solely responsible for this depleted region. We return to this point in the next section.

Following the same procedure for the halo samples, we show the TDtail and the pureHalo samples separately in Figs. 8 and 9, respectively. In addition to the Ωz : ΩR resonances identified in the thick-disc sample, we also find the 3:4 and 5:2 being populated. Just as we saw for the thick-disc sample, the stars associated with the resonances also occupy specific regions in the orbital parameter spaces.

thumbnail Fig. 8.

Similar to Fig. 7, but for the TDtail sample.

thumbnail Fig. 9.

Similar to Figs. 7 and 8, but for the pureHalo sample.

4.3. Commonalities between sub-samples

In all three sub-samples, we find similar families of resonances to be dominant. The Ωz : ΩR= 1:1 is the most prominent resonance in all samples. In the thick-disc sample, the other dominant resonance is 3:2. For the TDtail sample, this resonance is also the second most dominant frequency, but the 2:1 and 3:4 resonances are also prominent. In the pureHalo sample, the second most dominant resonance is 2:1, followed by 3:4 and 3:2.

We now return to the most prominent gap seen in Fig. 3 and present in all sub-samples, which is the gap that goes through (Rmax, zmax)∼(8, 2) kpc. For practical purposes, we focused on the thick-disc subset. We have previously noted from the middle row panels of Fig. 7 that the gap also divides the stars associated with the resonances Ωz : ΩR= 1:1, 2:1, and 1:2 in two separate zones. This implies that these resonances on their own can not fully explain the presence of a gap in the Rmax vs. zmax space.

To gain further insight, we selected stars around the gap for the Ωz : ΩR= 1:1 resonance, as shown in the leftmost panel of Fig. 10. The rightmost panel of this figure shows that the stars in this region cluster around Ωϕ : Ωz= 2:3, but that those above the gap have Ωϕz < 2/3, while those below have values greater than 2/3. Similar results are found for the other resonant families, with the gap seen in the characteristic orbital parameters of stars associated with the Ωz : ΩR= 2:1 resonance being related to Ωϕ : Ωz= 1:3. Overlapping chaotic zones, which are related to resonances, enhance chaotic diffusion (e.g., Laskar 1993). We surmise that the most prominent gap in (Rmax, zmax) shown here is due to the 3D (Ωϕ : ΩR : Ωz) = (2:2:3) resonance.

thumbnail Fig. 10.

Selection of stars (left panel) around the gap in Rmax vs. zmax for the Ωz : ΩR= 1:1 resonance for stars in the thick-disc sample. Their mapping to eccentricity vs. zmax/Rmax and the distributions of their frequency ratios Ωϕz are, respectively, shown in the middle and right panels. The last panel shows that the sample is bifurcated at Ωϕz= 2:3.

Interestingly, a similar gap in zmax space was already discussed by Schuster & Allen (1997). Our analysis agrees well with the assessment of these authors, who ascribed the gap to the diffusion of chaotic orbits between two quasi-periodic orbit families. That is, the chaotic orbits ‘stick’ to the quasi-periodic regions –one on each side of the gap– and quickly diffuse in the regions between them. Schuster & Allen suggested that such a structure appears because > 40% of the halo stars are on chaotic orbits in their Milky Way model (which is different to the one used in this work).

Figures 79 show that the velocity distributions of stars associated with resonances reveal the presence of clumps, which could potentially be confused with merger debris (e.g., the Arcturus stream, see Navarro et al. 2004 or Kushniruk & Bensby 2019). A more recent example is the reported prograde stellar stream, Nyx, with azimuthal velocities around 140 km s−1 (Necib et al. 2020), which could also be due to the Galactic bar acting on thick-disc-like stars, but which we did not consider here (e.g., see Monari et al. 2013 or Antoja et al. 2015). We note that, in most cases, the clumps associated with resonances do not correspond to overdensities in velocity space for the full data set, as can be seen from a comparison of Figs. 79 to Fig. 2.

4.4. Known accreted substructures in frequency space

We now explore the regions of frequency space occupied by the accreted substructures introduced in Sect. 3.3. Their kinematics and orbital parameters are shown in Fig. 5, while Fig. 11 shows their distribution in frequency space.

thumbnail Fig. 11.

Frequency maps of the known accreted substructures for the stars present in our pureHalo sample. Some of the clearest resonances in the panel on the left are marked with light grey lines.

Figure 11 reveals that Gaia-Enceladus and Thamnos are seen to span rather large regions in frequency space. They are affected by resonances just like the whole halo subset. The Sequoia stars occupy an small region close to and around the Ωϕ : ΩR= −2:3 resonance. On the other hand, the Helmi streams (in red) occupy two relatively narrow regions in Ωϕ vs. ΩR, and they also appear to be split in ΩR vs. Ωz, with one of the two groups being on the Ωz : ΩR= 1:2 resonance. There are stars both from the positive and negative z-velocity clumps on the 1:2 resonance, with the associated stars being slightly colder kinematically. The splitting into clumps in frequency space could potentially be due to different wraps of the debris streams (Gómez & Helmi 2010), but clearly a much higher number of stars would be needed to confirm this interpretation.

5. Discussion

The analysis carried out in this work shows that substructures present in spaces associated with the orbital parameters of stars can be driven by properties of the gravitational potential in which the stars move. In particular, we see that a realistic (but non-integrable) Milky Way potential – with a disc and halo components – leads to the presence of well-populated orbital resonances. The trapping of orbits at resonances and chaotic zones leads to the depletion of stars in the regions around these resonances.

These results do not strongly depend on the choice for the gravitational potential assumed for the Milky Way, provided this is generic enough (and not integrable, e.g., of Stäckel form). We also find that in two other commonly used Galactic potentials (Piffl et al. 2014; Bovy 2015), very similar substructures can be observed in Rmax vs. zmax, as is shown in Appendix A, and that they are associated with resonant families, and hence with the presence of non-integrability in the system.

Our analysis (as also concluded by Amarante et al. 2020) therefore directly shows that not all substructures are related to accretion, as is often considered in the literature, nor to the settling of the gravitational potential after major merger activity as had been suggested by Haywood et al. (2018). Nonetheless, the characteristics of substructures from merger events may sometimes be related to the presence of resonances. For example, we see that some of the stars in the Helmi streams appear to be close to a resonance, Ωz : ΩR∼ 1:2. This perhaps explains why its stars are distributed asymmetrically in velocity space (the stream with vz < 0 has more stars than that with vz > 0). This asymmetry was used to constrain its time of accretion to approximately 5-8 Gyr ago, which is a puzzlingly low value given that these stars are on relatively bound orbits. Since stars near a resonance take longer to spread out in space (Vogelsberger & White 2008), this can lead to an underestimation of their time of accretion. This means that chaotic dynamics may have to be considered when modelling the evolution of tidal debris, as already hinted by Price-Whelan et al. (2015) in their modelling of cold thin streams farther away in the halo.

Probably the best way to tell whether substructures are related to accretion events is via a chemical tagging analysis. Stars that originate from a satellite will follow characteristic tracks in chemical abundance space. On the other hand, stars that group together because of dynamical resonances have no reason to be chemically distinct from other stars.

How stars populate different resonant families depends on their distribution function and the gravitational potential in which they move. In this work, we only considered that of the Milky Way, but a recent analysis of nearby thin-disc stars has shown that the Sagittarius dwarf also plays a role in their dynamics (see Antoja et al. 2018). Interestingly, if we integrate the orbit of Sagittarius in the McMillan (2017) potential for 40 Gyr with the initial conditions derived by the Gaia Collaboration (2018c) and apply the SuperFreq code to obtain its frequencies (i.e. we repeat the analysis done for the stars in our sample), we find that Sagittarius falls on the Ωz : ΩR= 1:2 and Ωz : Ωϕ= 1:1 resonance. The stars in our sample on the Ωz : ΩR= 1:2 resonance (green points in Figs. 79) define a rather distinct branch in frequency space. They also have rather different velocities, with high vz values. This intriguing coincidence probably deserves further investigations, as this is also the resonance where some of the Helmi stream stars lie.

External influences such as Sagittarius or the Large Magellanic Cloud are known to have an effect on the dynamics of the Galactic disc and of streams, both of which are so cold kinematically, easily revealing their presence (Garavito-Camargo et al. 2019; Petersen & Peñarrubia 2020b,a; Cunningham & Garavito-Camargo 2020; Erkal et al. 2020; Vasiliev et al. 2021). The long orbital periods of these satellites, in comparison to those of halo stars in the solar vicinity, would suggest that the pertubations will act adiabatically on their orbits. More precise conclusions would need orbital integrations including those systems as well as the ‘live’ response of the Galaxy. Of course, all of this does not take away the main conclusions of this work, no matter the exact equilibrium state of the Milky Way halo: structures that arise in the space of orbital parameters can be purely due to the gravitational potential model considered and are not necessarily related to merger events.

6. Conclusions

We studied the dynamical properties of nearby stars in the thick-disc and stellar halo using data from the second data release of the Gaia mission. We explored spaces of characteristic orbital parameters for a generic Galactic potential and found that these host a large amount of structures. Further analysis using orbital frequencies shows that these structures are due to the presence of resonant families (as claimed also by Amarante et al. 2020, see also Schuster & Allen 1997), and to the depletion of orbits around them due to non-integrability. Therefore, these structures reflect intrinsic properties of generic gravitational potentials of the Milky Way.

These findings are interesting on their own and highlight that a large number of stars in these components (nearly 30% for the halo sample and specifically for the McMillan (2017) potential) are on resonances. This gives us hope to use them to more precisely pinpoint the gravitational potential of the Milky Way. Valluri et al. (2012) argued that most stars are launched on regular tori, and that therefore chaotic diffusion should occur. The expectation would thus be that the fraction of stars on such irregular orbits would be the lowest for the true potential (as this would indicate self-consistency between the distribution function and the gravitational potential).

Frequency analysis is also interesting, because it can reveal the individual streams originating in an accretion event (Gómez & Helmi 2010). Our analysis shows hints of such imprints, but the number of stars is too low to reach meaningful conclusions. Furthermore, it may be particularly challenging to identify the individual streams or wraps, as these may be disguised and trapped by the resonances present in the phase space of halo stars.


1

This value is consistent with the Gravity Collaboration (2019), who recently determined R = 8178 ± 13stat. ± 22sys. pc.

3

By definition, the fundamental frequencies are non-zero. However, a zero frequency line in the spectrum of a coordinate can indicate that the orbit’s centre is not on the origin. For example, a zero-frequency in z indicates an asymmetric orbit such as a ‘banana’ orbit.

Acknowledgments

We gratefully acknowledge financial support from a VICI grant and a Spinoza Prize from the Netherlands Organisation for Scientific Research (NWO) and HHK is grateful for the support from the Martin A. and Helen Chooljian Membership at the Institute for Advanced Study. This work has made use of data from the European Space Agency (ESA) mission Gaia (http://www.cosmos.esa.int/gaia), processed by the Gaia Data Processing and Analysis Consortium (DPAC, http://www.cosmos.esa.int/web/gaia/dpac/consortium). Funding for the DPAC has been provided by national institutions, in particular the institutions participating in the Gaia Multilateral Agreement. For the analysis, the following software packages have been used: vaex (Breddels & Veljanoski 2018), numpy (Van Der Walt 2011), matplotlib (Hunter 2007), jupyter notebooks (Kluyver et al. 2016).

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Appendix A: Variations between Milky Way models

We explore here the level of the model dependency of the structures in Rmax − zmax space analysed in the main text by comparing two key figures for a set of often-used Milky Way models. These are the MWPotential2014 model (Bovy 2015) and the Piffl et al. (2014) model. For simplicity, we only show our findings for the pureHalo sample, which naturally has the largest range in Rmax − zmax space and therefore populates the full range of structures seen in this space. Figures A.1 and A.2 summarise our findings, and they clearly show that the variations between the models are small.

thumbnail Fig. A.1.

Comparison of the Rmax − zmax space for three often-used Milky Way models. The left panel is equivalent to the bottom middle panel of Fig. 3 in the main text and is repeated here for convenience.

Figure A.1 shows the same space as the one shown in the middle panel in the bottom row of Fig. 3. For convenience, we also show the version using the McMillan17 potential here, this panel is the same as the one in the main text. The wedges in Rmax − zmax space are present in all models, although they are clearly not equivalent. The large-scale structures (i.e. the three large wedges) are in approximately the same location. However, the fine substructure on top of these large wedges appears to be more sensitive to the model. We do note that some of this is due to the visualisation method and the sample size. For a larger sample that populates a larger fraction of the Rmax − zmax space, we expect to find more and finer substructures.

Similarly, Fig. A.2 shows the variation of the bottom-left panel of Fig. 6 as a function of the three Milky Way models (the McMillan17 panel is also equivalent to the one in the main text and is shown to facilitate comparison). Also in this space, the similarities between the three models are apparent.

thumbnail Fig. A.2.

Comparison of the ΩR − Ωz space for three often-used Milky Way models. The left panel is equivalent to the bottom-left panel of Fig. 6 in the main text and is repeated here for convenience.

All Tables

Table 1.

Percentage of stars in the various sub-samples associated with the listed Ωz : ΩR resonances.

All Figures

thumbnail Fig. 1.

Velocity distribution of the stars in our sample, highlighting the various kinematic selections. A logarithmic map of star counts is shown in the background (the smooth blue map). Stars within the cyan semi-circles are considered to belong to the thick-disc. The stars outside the outermost cyan curve are considered halo stars, they are split into a pure halo and a TDtail sample.

In the text
thumbnail Fig. 2.

Different projections of velocity space for the various subsets considered in this chapter: thick-disc, TDtail, and pure halo from top to bottom. The lower density of stars present in these top panels for small values of vR and vz is due to the exclusion of thin-disc stars in our selection. We note that even our pureHalo subset contains some stars with ‘hot’ thick-disc-like kinematics (as evidenced by the arch-like structure in the bottom central panel). Red dashed curves in the bottom panel outline the structures that are mentioned in the text.

In the text
thumbnail Fig. 3.

Orbital parameters spaces for the different subsets defined in Sect. 2. From top to bottom: thick-disc, TDtail, and pureHalo. The left panels correspond to the most often used space to identify merger debris (i.e. Lz,  E). The substructure seen near Lz = 250 kpc km s−1 for highly bound energies in the middle-left panel is the globular cluster M4 (also known as NGC 6221), whose extent is largely the result of limitations in the Gaia DR2 astrometry for very dense stellar fields. The subspaces shown in the middle and right panels depict a large amount of structure, which appears to be independent of the subset or population considered. We note that the top row shows a smaller dynamical range than the other two rows.

In the text
thumbnail Fig. 4.

Rmax vs. zmax for the thick-disc (left), TDtail (middle), and pureHalo (right) populations, computed for an integrable two-component Milky Way Stäckel potential. In all subsets, the amount of structure has drastically decreased. The clear overdensity of stars in the TDtail sample and at Rmax ∼ 6.5 kpc and zmax ∼ 0.5 kpc corresponds to the globular cluster M4 and is thus of no specific interest.

In the text
thumbnail Fig. 5.

Distribution of accreted substructures in the pureHalo subset following the identification by Koppelman et al. (2019b): in velocity space (top) and in orbital parameters spaces (bottom).

In the text
thumbnail Fig. 6.

Frequency maps for the different populations. From top to bottom: thick-disc, TDtail, and pureHalo population. From left to right: cycling through the combinations of frequencies. The panels, especially those in the left column, show a clear presence of resonances in the form of straight lines and of depleted regions around them.

In the text
thumbnail Fig. 7.

Selected resonances in the thick-disc subset and their mapping onto other informative subspaces. Top row: frequency maps highlighting the main resonant families. Middle row: characteristic orbital parameters colour-coded according to the different resonances. Bottom row: their distribution in velocity space.

In the text
thumbnail Fig. 8.

Similar to Fig. 7, but for the TDtail sample.

In the text
thumbnail Fig. 9.

Similar to Figs. 7 and 8, but for the pureHalo sample.

In the text
thumbnail Fig. 10.

Selection of stars (left panel) around the gap in Rmax vs. zmax for the Ωz : ΩR= 1:1 resonance for stars in the thick-disc sample. Their mapping to eccentricity vs. zmax/Rmax and the distributions of their frequency ratios Ωϕz are, respectively, shown in the middle and right panels. The last panel shows that the sample is bifurcated at Ωϕz= 2:3.

In the text
thumbnail Fig. 11.

Frequency maps of the known accreted substructures for the stars present in our pureHalo sample. Some of the clearest resonances in the panel on the left are marked with light grey lines.

In the text
thumbnail Fig. A.1.

Comparison of the Rmax − zmax space for three often-used Milky Way models. The left panel is equivalent to the bottom middle panel of Fig. 3 in the main text and is repeated here for convenience.

In the text
thumbnail Fig. A.2.

Comparison of the ΩR − Ωz space for three often-used Milky Way models. The left panel is equivalent to the bottom-left panel of Fig. 6 in the main text and is repeated here for convenience.

In the text

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