<figure><img src="image_1.png" /></figure>

# A Large Moving Group within the Lower Centaurus Crux Association

Bertrand Goldman¹,², Siegfried Röser³,⁴, Elena Schilbach³,⁴, Attila C. Moór⁵, and Thomas Henning¹

¹Max-Planck-Institut für Astronomie, Königstuhl 17, D-69117 Heidelberg, Germany; goldman@mpia.de

²Observatoire astronomique de Strasbourg, Université de Strasbourg—CNRS UMR 7550, 11 rue de l'Université, F-67000, Strasbourg, France

³Zentrum für Astronomie der Universität Heidelberg, Landessternwarte, Königstuhl 12, D-69117 Heidelberg, Germany

⁴Zentrum für Astronomie der Universität Heidelberg, Astronomisches Rechen-Institut, Mönchhofstraße 12-14, D-69120 Heidelberg, Germany

⁵Konkoly Observatory, Research Centre for Astronomy and Earth Sciences, Hungarian Academy of Sciences, H-1121 Budapest, Konkoly Thege Miklós út 15-17, Hungary

Received 2018 July 2; revised 2018 October 1; accepted 2018 October 1; published 2018 November 15

## Abstract

Scorpius–Centaurus is the nearest OB association, and its hundreds of members are divided into subgroups, including the Lower Centaurus Crux (LCC). Here we study the dynamics of the LCC area. We report the revelation of a large moving group containing more than 1800 intermediate- and low-mass young stellar objects and brown dwarfs that escaped identification until Gaia DR2 allowed a kinematic and photometric selection to be performed. We investigate the stellar and substellar content of this moving group using the Gaia DR2 astrometric and photometric measurements. The median distance of the members is 114.5 pc, and 80% lie between 102 and 135 pc from the Sun. Our new members cover a mass range of 0.02–5 $M_\odot$ and add up to a total mass of about 700 $M_\odot$. The present-day mass function follows a log-normal law with $m_c = 0.22 M_\odot$ and $\sigma = 0.64$. We find more than 200 brown dwarfs in our sample. The star formation rate had its maximum of $8 \times 10^{-5} M_\odot \text{ yr}^{-1}$ about 9 Myr ago. We grouped the new members into four denser subgroups, which have increasing age from 7 to 10 Myr, surrounded by “free-floating” young stars with mixed ages. Our isochronal ages, now based on accurate parallaxes, are compatible with several earlier studies of the region. The whole complex is presently expanding, and the expansion started between 8 and 10 Myr ago. Two hundred members show infrared excess compatible with circumstellar disks from full to debris disks. This discovery provides a large sample of nearby young stellar and substellar objects for disk and exoplanet studies.

**Key words:** brown dwarfs – open clusters and associations: individual (Lower Centaurus Crux) – protoplanetary disks – stars: formation – stars: luminosity function, mass function

**Supporting material:** tar.gz file

## 1. Introduction

Co-eval structures of the solar neighborhood are useful entities for various fields in astronomy. Either gravitationally bound (a.k.a. open clusters) or not (moving groups), they provide samples of stars that share the same composition, motion, and age. Possible deviations from this set of parameters also give insights into the formation of these structures: the velocity dispersion may allow the age and location of the birth place of the population to be constrained, and the composition and age dispersion may reveal a more complicated formation and multiple stages of star formation. Co-evolution, on the other hand, lets us constrain the stellar evolutionary models. To find such structures in the solar neighborhood allows spectroscopic follow-up of high spectral resolution and/or on small telescopes and easy detection of low-mass members, making them particularly valuable.

For almost a century, the Scorpius–Centaurus OB association (Sco–Cen) has been known as an ensemble of young, massive early-type stars that could hold the keys of star formation. Blaauw (1964) divided the huge nearby association into three subgroups, Upper Scorpius (US), Upper Centaurus Lupus (UCL), and Lower Centaurus Crux (LCC). Being on the southern hemisphere, the quality of the astrometric data historically was always poorer than that for associations on the northern hemisphere. Only with the advent of *Hipparcos* data did the situation change, but only for a subset of the brightest stars. In their fundamental paper, using the *Hipparcos* measurements, de Zeeuw et al. (1999) presented an overview of the OB associations in the wider neighborhood of the

Sun. For Scorpius–Centaurus, they obtained mean distances of 145 pc for the US, 140 pc for the UCL, and 118 pc for the LCC. Pecaut & Mamajek (2016) determined median ages of 11 (US), 16 (UCL), and 17 Myr (LCC), though each has a considerable spread of ages. Other authors obtained ages as low as 5 Myr for US (Preibisch & Zinnecker 1999; Preibisch & Mamajek 2008), and 16 to 18 Myr for the UCL and LCC, respectively (Mamajek et al. 2002; Sartori et al. 2003).

In their introduction, Wright & Mamajek (2018) gave an excellent overview of the importance of the Sco–Cen complex for questions related to the modes of star formation. Using the observations from Gaia DR1, they confirmed that all three subgroups, US, UCL, and LCC, are gravitationally unbound and found neither evidence for the expansion of the subgroups nor that they, or part of them, were formed by the disruption of star clusters.

Using the *Hipparcos* data and a convergent point method, de Zeeuw et al. (1999) found extended moving groups in the area of Sco–Cen, one in each subgroup. However, they and subsequent authors could only detect stars with masses typically above 0.7 $M_\odot$. The situation was best described by Preibisch & Mamajek (2008): “the surface of the census of low-mass Sco–Cen members has barely been scratched.” In this paper, we use the results of Gaia DR2 (Gaia Collaboration et al. 2018) to uncover a large population of young stellar objects located in the LCC area of the Sco–Cen region. About 70% of these stars belong to a large moving group and assemble in several subgroups of increasing age.
<figure><img src="image_2.png" /><figcaption>Figure 1. The Scorpius–Centaurus region. The axes are the Galactic longitude and Galactic latitude in degrees. The gray squares in the background show the stars from Hoogerwerf (2000) and outline the Sco–Cen association. The black dots are the seeds of our cluster search in TGAS (see the text). The small cyan dots outline the region studied in this paper and show the stars of our Basic Sample (BS; see text). The large cyan dot at (l, b) = (344°, +4°) is the young compact moving group around V1062 Sco (Röser et al. 2018). The large violet diamonds are the groups found by Oh et al. (2017), and the large red dots mark open clusters in MWSC, with distances less than 200 pc.</figcaption></figure>

The paper is structured as follows. In Section 2, we describe the identification of the moving group and its subgroups in the TGAS and Gaia DR2 catalogs; in Section 3, we determine and discuss the mass functions and age distributions of the groups. In Section 4, we discuss some interesting members, in particular those surrounded by circumstellar disks, and binaries. In Section 5, we study the internal dynamics and expansion of the groups. In Section 6, we summarize and discuss the major results of our paper.

## 2. Detection of the Crux Star-forming Region

### 2.1. Astrometric Detection

Parallel to the cluster search program described in Röser et al. (2016), we ran an alternative search process better adapted to the finding of nearby comoving groups. We give only a rough description of the procedure here as it is not crucial for this paper. The idea behind was to select a grid of convergent points homogeneously distributed over the sphere and to search for stars whose positions, proper motions, and parallaxes are consistent with a given convergent point. We then selected stars that have at least five other stars within a radius of 10 pc around them and whose tangential velocities are consistent with the same convergent point. These we regarded as seeds of possible comoving groups. Our search was restricted to stars having a parallax larger than 5.5 mas in TGAS. On the whole sky, we found 706 of these seeds. Here we concentrate onto the area of Sco–Cen between Galactic longitudes 280° and 360° and Galactic latitudes from −20° to +30°.

In Figure 1, we present the result of our search in the area of Sco–Cen. In the background we show the stars from Hoogerwerf (2000; US, UCL, LCC) with large gray squares. This serves as an indication of the extent of the Sco–Cen association. The seeds found by our search in TGAS are shown as black dots. Isolated seeds stand for small comoving groups

with not more than six stars in a 10 pc radius around the seed in TGAS. Spatial concentrations of seeds may suggest larger coherent groups. For instance, the concentration marked by the large cyan dot at (l, b) = (344°, +4°) is the young compact moving group around V1062 Sco (Röser et al. 2018) at a distance of 175 pc. For comparison, we also plot the centers of the two open clusters contained in MWSC (Milky Way Star Clusters; Kharchenko et al. 2013) up to 200 pc from the Sun (large red dots), as well as the groups found by Oh et al. (2017) with more than five neighbors (violet diamonds). Their and our findings overlap in parts. We retrieved the open cluster IC 2391 in MWSC at Galactic longitude $l = 289°$. Even the poorly populated cluster Feigelson 1 (also known as the $\epsilon$ Chameleontis group; Feigelson et al. 2013) at (l, b) = (300°, −15°) was found as one seed. Of the other seeds, the concentrations at (l, b) = (300°, −1°) and the large group at (l, b) = (297°, +8°) seemed to be the most interesting because of their proximity to the Sun; the mean of the TGAS parallaxes of these groups was 9.5 and 9.1 mas, respectively. We called the southern group Group A and the northern one Group B. The violet diamond in between Group A and B corresponds to the center of “Group 4” in Oh et al. (2017). Their Group 4 actually comprises our groups A and B, and in their notation, it has a size of 114 (number of stars in the group). It is worth mentioning that Chereul et al. (1997) already found a moving group of 33 nearby A-type dwarfs at a mean distance of 105 pc at (l, b) = (304°, +3°). For our Group A, we determined the phase-space coordinates on the basis of the TGAS positions, proper motions, and parallaxes, as well as with radial velocities retrieved from SIMBAD, and found mean values of

$$ 
\begin{aligned}
R_c &= (\xi_c, \eta_c, \zeta_c) = (-47.17, -5.23, -96.41) \text{ pc}, \\
V_c &= (v_{\xi,c}, v_{\eta,c}, v_{\zeta,c}) = (-2.97, 18.90, -14.76) \text{ km s}^{-1} \quad (1)
\end{aligned}
$$

in the barycentric equatorial coordinate system with the $\xi$-axis pointing to the vernal equinox, the $\eta$-axis to $\alpha = 90°$, and the
ζ-axis points to the equatorial north pole. The vector **R**<sub>c</sub> is measured in parsec and is the position vector of the center of the group, and **V**<sub>c</sub>, in km s<sup>-1</sup>, is the vector of the space velocity of the group; **V**<sub>c</sub>/||**V**<sub>c</sub>|| is the Cartesian equatorial coordinates of the convergent point.

**Gaia** DR2 now enables groups A and B to be studied in great detail. In a first step, we made a query of the *Gaia* DR2 catalog consisting of a cone search around (α, δ) = (186°.5, -60°.5) with a radius of 20°. The center of the cone was chosen midway between our groups A and B from TGAS, with radius large enough to contain both groups. Further, we applied the following restrictions: for the trigonometric parallax, ϖ ≥ 7 mas; for the proper motions, -43 mas yr<sup>-1</sup> ≤ μ<sub>α*</sub> ≤ -20 mas yr<sup>-1</sup> (μ<sub>α*</sub> ≡ μ<sub>α</sub> cos δ) and -35 ≤ μ<sub>δ</sub> ≤ +12 mas yr<sup>-1</sup>. The proper motion and parallax cuts were chosen to be consistent with the distribution of stars identified as members of these probable young moving groups we revealed in TGAS. This query yielded 20,138 objects, and we call this set of objects “Crux Cone.” For further processing, we followed the procedures described in Lindegren et al. (2018 Chapter 4.3 and Appendix C, Figures C.1 and C.2) to obtain a stellar sample cleaned of possible artifacts.

First, we required the “visibility periods used” to be equal to or larger than seven (Equation (11, ii) in Lindegren et al. 2018), which left us with 19,878 objects. The “unit weight error” cut (cf. Equation C.1 in Lindegren et al. 2018) removed a considerable portion of dubious entries, and 7,324 sources remained. Instead of the factor of 1.2 in this equation, we used a factor of 1.3, which seems more appropriate for our sample and keeps some 40 additional stars in the magnitude range 10.0 < *G* < 19.5 mag. Among these 40 stars, 11 (among 1844) will be selected in our final sample. It seems therefore likely that the astrometric parameters for those 11 stars found to be both young and comoving (with a narrow selection, as will be shown below) are correct and those stars should be kept in our analysis. The 1% increase in the statistics has of course no impact on the conclusion of this paper. As a next step, we applied the “flux excess ratio” cut (Equation C.2 in Lindegren et al. 2018), which reduced the sample to 3946 objects. Finally, we discarded three stars with relative errors of parallaxes σ*ρ*/ρ* larger than 10%. In summary, from the original 20,138 entries in our search cone, 3943 sources survived these filterings, forming an astrometrically clean sample. For the subsequent discussion, we will refer to this sample as “the Basic Sample (BS).” The stars of the BS are shown in Figure 1 as small cyan dots and outline the area of our search.

## 2.2. *Candidate Selection*

In Figure 2, we show the color–magnitude diagram (CMD) of the 3943 stars of the BS, where we plot the absolute magnitude M<sub>G</sub> (using the parallaxes from DR2) versus the color G − G<sub>RP</sub>. For colors redder than about 0.4 mag, two separate sequences show up, which we attribute to young stars (upper sequence) and old stars (lower sequence). We strengthen this in Figure 2, where we also plot the theoretical isochrones of CIFIST (Baraffe et al. 2015) and Parsec (Marigo et al. 2017). The red lines are the CIFIST isochrones for 10 and 400 Myr, while the black lines are the 10 and 400 Myr Parsec isochrones. From a comparison between the two sets of isochrones for different ages, we found that the brighter part of the CMD is appropriately described by the Parsec isochrones, whereas the CIFIST isochrones better fit the pre-main-sequence part for stars less massive than 1.4 M<sub>☉</sub>. Here we note that the CIFIST

<figure><img src="image_3.png" /><figcaption>**Figure 2.** Absolute magnitudes M<sub>G</sub> vs. G − G<sub>RP</sub> of the 3943 stars of the basic sample (BS). The black curves are the CIFIST isochrones from Baraffe et al. (2015) for 10 Myr (upper curve) and 400 Myr (lower curve). The red lines are the 10 Myr (upper curve) and 400 Myr (lower curve) isochrones from Parsec 3.0 (Marigo et al. 2017).</figcaption></figure>

<figure><img src="image_4.png" /><figcaption>**Figure 3.** Absolute magnitudes M<sub>G</sub> vs. G − G<sub>RP</sub> of the 3943 stars of the basic sample (BS). The 40 Myr isochrone of Baraffe et al. (2015; solid line, masses less than 1.4 M<sub>☉</sub>) and Parsec 3.0 (dashed line, masses higher than 1.4 M<sub>☉</sub>) separates the young population (cyan) from the old population (violet). An A<sub>G</sub> = 1 mag reddening vector is shown by an arrow.</figcaption></figure>

isochrones use the pre-launch filter profiles of Jordi et al. (2010), while the Parsec isochrones use the profiles of Evans et al. (2018), both in the Vega system.

The difference between the two observed sequences is too large to be explained by unresolved binaries of main-sequence stars in this field. Consequently, the upper sequences must be attributed to young pre-main-sequence stars. To separate the old and young populations, we made a cut along the 40 Myr CIFIST isochrone from the faintest stars up to the locus (G − G<sub>RP</sub>, M<sub>G</sub>) = (0.34, 3.22), corresponding to 1.4 M<sub>☉</sub>, and continuing on the Parsec 40 Myr isochrone for brighter stars. The actual separation is shown in Figure 3. We additionally discarded a small number of scattered stars, which we attribute to older red giants. The subsample of young stars (2659) in the BS is called “Young,” the complement (1,284 stars) “Old.” We note that beginning at the turnover at (G − G<sub>RP</sub>, M<sub>G</sub>) = (0.75, 6.0), there may be increasing contamination by old binaries on the sample “Young.” Moreover, a small number of main-sequence stars with G − G<sub>RP</sub> ≤ 0.34 may also be field stars, which can be identified and excluded by their different
kinematics. We discussed the CMDs without considering a possible extinction toward the sample of our stars from BS. One may expect that its influence is rather small since our stars are all nearby ($\bar{\omega} \ge 7$ mas). On the other hand, they are actually located within the Sco–Cen association and may not be far away from or behind gas and molecular clouds. However, from Figure 2, we note a good agreement between the 400 Myr isochrone and the main sequence of older BS stars of intermediate masses, and there is no reason to assume a significant extinction toward the stars in the BS. To verify this result, we also considered CMDs in other passbands including the near-infrared (J, Ks) from 2MASS (Skrutskie et al. 2006) and Gaia blue photometer (BP) integrated magnitudes. Nevertheless, any attempt to introduce a correction for de-reddening led to poorer fittings of isochrones to observations in the CMDs. Of course, we cannot exclude that some individual stars may be significantly reddened, especially if we consider the corresponding reddening vector in Figure 3. The extinction data published in Gaia DR2 do not allow this option to be checked as Andrae et al. (2018) mentioned in their paper. To summarize, we refrained from de-reddening the stars in the BS.

### 2.2.1. Kinematic Analysis

To reveal the members of the potential moving group characterized by Equation (1), we then studied the kinematics of the stars in the BS. Unfortunately, only about 12% of these stars have accurate radial velocities in Gaia DR2, with mean errors smaller than 2 km s⁻¹. Therefore, in order to identify comoving stars, we have to rely on criteria based on their tangential velocities only. The tangential velocity for each star i has the components $\kappa\mu_{\alpha*,i}/\bar{\omega}_i$ and $\kappa\mu_{\delta,i}/\bar{\omega}_i$, with $\kappa = 4.74047$ being the transformation factor from 1 mas yr⁻¹ at 1 kpc to 1 km s⁻¹. Stars with a common space motion should have tangential velocities directed toward a convergent point, in our case defined by Equation (1). We computed the components of the observed tangential velocity parallel and perpendicular to the local vector in the direction of this convergent point by a simple rotation around an angle $\psi$ which we determine below.

We followed the formalism of the convergent point method as described, e.g., in Röser et al. (2011) and transformed the Cartesian velocity vector of the group motion $V_c$ from Equation (1) by

$$ \begin{bmatrix} V_{c,\alpha,i} \\ V_{c,\delta,i} \end{bmatrix} = \begin{bmatrix} -\sin \alpha_i & \cos \alpha_i & 0 \\ -\cos \alpha_i \sin \delta_i & -\sin \alpha_i \sin \delta_i & \cos \delta_i \end{bmatrix} \cdot V_c. \quad (2) $$

With Equation (8) from van Leeuwen (2009), the angle $\psi$ is given as $\psi = \arctan(V_{c,\delta,i}/V_{c,\alpha,i})$, and the desired rotation reads

$$ \begin{bmatrix} v_{\parallel,i} \\ v_{\perp,i} \end{bmatrix} = \begin{bmatrix} \cos \psi & \sin \psi \\ -\sin \psi & \cos \psi \end{bmatrix} \cdot \begin{bmatrix} \kappa \ \mu_{\alpha*,i} / \bar{\omega}_i \\ \kappa \ \mu_{\delta,i} / \bar{\omega}_i \end{bmatrix}. \quad (3) $$

In an ideal case, where all the stars of a group would exactly move along the vector $V_c$, $v_\perp$ would become zero. This implies that we a priori know the vector $V_c$ with highest accuracy. But even in this case, the internal velocity dispersion and uncertainties of the measurements will cause scatter in $v_\perp$.

We also determined the covariance matrix for the velocities $v_\parallel$ and $v_\perp$ according to the error propagation from the covariance matrix of $\mu_{\alpha*}$, $\mu_\delta$, and $\bar{\omega}$. From the distribution of the variances, we determined the mean and the 90th

<figure><img src="image_5.png" /><figcaption>Figure 4. Distribution of the stars from the Old sample in the $v_\parallel$, $v_\perp$ plane with a binning of 0.5 × 0.5 (km s⁻¹)². $N_{\text{stars}}$ is the number of stars per bin.</figcaption></figure>

<figure><img src="image_6.png" /><figcaption>Figure 5. Distribution of the stars from the Young sample in the $v_\parallel$, $v_\perp$ plane with a binning of 0.5 × 0.5 (km s⁻¹)². $N_{\text{stars}}$ is the number of stars per bin.</figcaption></figure>

percentile of $\sigma_{v_\parallel}$ as 0.25 km s⁻¹ and 0.36 km s⁻¹, respectively. The corresponding values for $\sigma_{v_\perp}$ are 0.08 and 0.11 km s⁻¹.

Figures 4 and 5 show the distribution of the stars in the Old and the Young samples in the $v_\parallel$, $v_\perp$ plane with a binning of 0.5 × 0.5 (km s⁻¹)². This bin size is large compared to the variances in $v_\parallel$ and $v_\perp$, as we just showed. So, the distributions in Figures 4 and 5 present the true physical velocity dispersions of the BS. The different behavior of the old and young populations is obvious. The old stars (left image) are homogeneously distributed on the plane with a typical density of fewer than three stars per 0.5 × 0.5 (km s⁻¹)². Such a flat distribution is fully consistent with the broad velocity dispersion of Galactic field stars. On the other hand, the velocity distribution of the young stars shows an impressive overdensity with a peak at (20.25, −0.75) km s⁻¹ with a maximum of more than 150 stars per 0.5 × 0.5 (km s⁻¹)². This overdensity is a strong signature for an autonomous coherent group consisting essentially of stars younger than about 40 Myr (Figure 3).
<figure><img src="image_7.png" /><figcaption>**Figure 6.** Density distribution of the stars from the CMG projected onto the Y, Z plane (Galactic coordinates). The bin size corresponds to 2.5 × 2.5 pc$^2$. $N_{\text{stars}}$ is the number of stars per bin. The cyan, blue, green, and black lines show the borders of the subgroups A0, A, B, and C, respectively.</figcaption></figure>

In Figure 5, we define this group by manually cutting the stars falling into the peak mentioned above. Indeed, they are populating the bins at a level of more than 30 $N_{\text{stars}}$. This cut is indicated by contour lines, and the selected sample contains 1844 objects. We call this sample the Crux Moving Group (CMG) henceforth.

In the following discussion, we use barycentric Galactic Cartesian coordinates X, Y, Z instead of the equatorial coordinates as before. The axes X, Y, Z are directed toward the Galactic center, the direction of Galactic rotation, and the Galactic north pole, respectively. The corresponding velocity coordinates are U, V, W. When we compare both samples, we find that, also in space, stars from the Old sample are distributed homogeneously, whereas the young stars, and especially the CMG stars, exhibit stronger local concentrations. In Figure 6, we present the distribution of the 1844 stars of the CMG sample in the YZ plane. Although these stars were selected in a narrow velocity range, they show a relatively broad distribution in space with several remarkable local overdensities. This led us to subdivide the CMG into four subgroups (CMG A0, CMG A, CMG B, and CMG C). In Figure 6, we draw the borders of the four subgroups with cyan, blue, green, and black lines respectively. The set of stars from the CMG not contained in groups A0 to C is called the CMG Intergroup (CMG Z). Let us mention here that this subdivision is somehow arbitrary and was guided by the appearance of these groups already in Figure 1.

The excellent quality (⟨$\bar{\sigma}/\sigma_{\bar{\sigma}}$⟩ ≥ 10) of the individual parallaxes from Gaia DR2 in the region of the CMG allows them to be inverted into individual distances. The four subgroups are distinguishable not only by their position on the sky, but also by their different distances from the Sun. Figure 7 shows the distribution of the individual distances of the stars in the four subgroups. There is a clear trend in distance, with A0 the closest to the Sun and the largest subgroup C as the farthest. Altogether, the whole CMG, including the Intergroup members, extends from 102 to 135 pc from the Sun (10th and 90th percentiles of the parallax distribution). The subgroups A0 to C are more compact. The

<figure><img src="image_8.png" /><figcaption>**Figure 7.** Distribution of the distances of the individual stars in the subgroups of the CMG derived from the trigonometric parallaxes from Gaia DR2.</figcaption></figure>

characteristic data for these groups are given in Table 1, while the data for individual members available in electronic format are described in Table 4. The spatial coordinates of the centers of the four subgroups are the means of the individual X, Y, Z of the $N_{\text{stars}}$ in each subgroup. The space velocities U, V, W of the groups are the averages of the individual velocities of those stars that have reliable radial velocity measurements in Gaia DR2. In Table 1, we give the standard deviations for the velocities, not the mean errors of the mean velocities. These are much more accurate and range from 0.06 to 0.50 km s$^{-1}$ with a median of 0.2 km s$^{-1}$.

There are no sharp, physically justified boundaries in space between CMG A0, A, B, and C and the stars from CMG Z. However, we will show in the following that, on average, they slightly differ not only in their location and kinematics, but also in their average age.

## 2.3. CMG Compared with Other LCC Samples

Before comparing our results with the work of other authors, we first summarize the selection process that led us to the large moving group CMG. The first selection from Gaia DR2 was a cone search around ($\alpha, \delta$) = (186°5, -60°5) with a radius of 20°, trigonometric parallax $\bar{\sigma} \ge 7$ mas, and proper motions -43 mas yr$^{-1} \le \mu_{\alpha*} \le -20$ mas yr$^{-1}$ ($\mu_{\alpha*} \equiv \mu_{\alpha} \cos \delta$) and -35 $\le \mu_{\delta} \le +12$ mas yr$^{-1}$ (see Section 2.1). This resulted in the Crux Cone with 20,138 objects. Then we applied quality filters as recommended by Lindegren et al. (2018, also described in Section 2.1) that led to the BS. These and the other relevant subsamples are summarized in Table 2.

The area studied in this paper covers most of the LCC and a small part of the UCL as shown in Figure 1. In the following, we compare the stellar content of our samples with the findings of other authors in the literature. A summary of these comparisons is given in Table 3. The first column lists the publications with which we compare our sample. The second column gives the number of stars in the respective data set. The third column is the size of the subset of stars from column 2, which are located in the cone around ($\alpha, \delta$) = (186°5, -60°5; radius 20°). In the fourth column, we list the number of stars from column 3 that have DR2 measurements and fall into the Crux Cone. The following columns list the cross-matches of the objects from column 4 with our BS (column 5), the “Old” sample (column 6), the “Young” sample (column 7), and the CMG (column 8).
<table><caption>**Table 1**

Characteristic Data of the Subgroups of the CMG</caption><thead><tr><th>Subgroup</th><th>N<sub>Stars</sub></th><th>X&#772; (pc)</th><th>Y&#772; (pc)</th><th>Z&#772; (pc)</th><th>U&#772; (km s<sup>-1</sup>)</th><th>V&#772; (km s<sup>-1</sup>)</th><th>W&#772; (km s<sup>-1</sup>)</th><th>D<sub>median</sub> (pc)</th><th>Age (yr)</th><th>Mass (M<sub>☉</sub>)</th></tr></thead><tbody><tr><td>A0</td><td>49</td><td>52 ± 3</td><td>-88 ± 5</td><td>-17 ± 3</td><td>-10.2 ± 0.5</td><td>-19.4 ± 1.0</td><td>-8.9 ± 0.3</td><td>103.0</td><td>7.0</td><td>13</td></tr><tr><td>A</td><td>149</td><td>53 ± 3</td><td>-92 ± 3</td><td>-3 ± 3</td><td>-9.5 ± 0.7</td><td>-20.1 ± 1.1</td><td>-7.9 ± 0.3</td><td>106.5</td><td>7.2</td><td>50</td></tr><tr><td>B</td><td>317</td><td>55 ± 8</td><td>-94 ± 4</td><td>+10 ± 3</td><td>-8.8 ± 1.1</td><td>-20.2 ± 1.8</td><td>-7.1 ± 0.3</td><td>109.5</td><td>9.3</td><td>98</td></tr><tr><td>C</td><td>487</td><td>63 ± 11</td><td>-97 ± 8</td><td>+23 ± 4</td><td>-8.2 ± 1.1</td><td>-20.9 ± 1.5</td><td>-6.1 ± 0.6</td><td>118.0</td><td>10.0</td><td>249</td></tr></tbody></table>

Note. The mean values are followed by their standard deviations.

<table><caption>**Table 2**

The Selection Window for the Crux Group of Young Stars (CMG)</caption><thead><tr><th>Data Set</th><th>N_stars</th></tr></thead><tbody><tr><td>Crux Cone</td><td>20,138</td></tr><tr><td>Basic sample (BS)</td><td>3943</td></tr><tr><td>"Old"</td><td>1284</td></tr><tr><td>"Young"</td><td>2659</td></tr><tr><td>CMG</td><td>1844</td></tr></tbody></table>

In their influential paper, de Zeeuw et al. (1999) applied a convergent point method to the *Hipparcos* measurements to find moving groups in 12 nearby OB associations; one of them was found in the LCC. Their LCC moving group contains 180 members. Table 3 shows that all but one star are in our area, but only 129 are in the Crux Cone. This means that 50 of the de Zeeuw stars have *Gaia* DR2 parallaxes smaller than 7 mas and/or proper motions outside our selection window. The other six stars did not survive the quality cuts (Lindegren et al. 2018; see also Section 2.1), and therefore, did not appear in the BS: they are probably too bright ($V < 6$ mag) to get accurate parameters in DR2. However, the majority of the de Zeeuw stars in the Crux Cone (85%) are classified as young by us, and two-thirds of the latter fulfill the strict kinematic conditions of being members of the CMG. About the same holds for the de Zeeuw UCL stars that are in our area. We note here that the faintest star from de Zeeuw in the CMG is HIP 108016 (F6/7V, $V = 9.78$ mag). This does not mean that the de Zeeuw sample is in any way complete down to $V \approx 10.0$ mag. The *Hipparcos* catalog is based on an input list with a restriction of 2.5 stars per square degree, a magnitude limit at about 12th magnitude in V, and has a completeness limit as bright as $V = 7.3$ mag.

Hoogerwerf (2000) made use of the “Astrographic Catalogue + Tycho” reference catalog (ACT) and the Tycho Reference Catalogue (TRC), which are complete to $V = 10.5$ mag. These catalogs only contain proper motions of a typical accuracy of 3 mas yr$^{-1}$, but no parallaxes. So, it is not surprising that among the 1190 Hoogerwerf stars in our area, only 289 are in the Crux Cone. About 80% of them are “Young,” and 136 stars fulfill the strong kinematic restrictions of the CMG.

Rizzuto et al. (2011) took the data from the *Hipparcos* catalog and combined these with radial velocities taken from the Second Catalog of Radial Velocities with Astrometric Data (Kharchenko et al. 2007). Using this data set and a Bayesian membership selection method, they found 436 members in the whole Sco–Cen complex. In the CMG, we find 76 stars out of their 436; 71 of them are common to de Zeeuw et al. (1999). Based upon kinematically and X-ray preselected candidates, Mamajek et al. (2002) carried out spectroscopic observations and characterized 110 pre-main-sequence stars as UCL and LCC members, with 41 of them in the Crux Cone. All but two stars are classified as young

by us also, and 30 of them are members of the CMG. Pecaut et al. (2012) investigated the 138 F-type members of the de Zeeuw et al. (1999) sample. Out of 48 of their stars in our BS, 42 are young and 28 are in the CMG.

Song et al. (2012) spectroscopically identified 104 G-, K-, and M-type members of the Sco–Cen complex based on X-ray and kinematically preselected samples. They confirmed the youth of these stars by comparing Li $\lambda$6708 absorption line strengths against those of stars in the TW Hydrae association and the $\beta$ Pictoris moving group. From their 57 members in the Crux Cone, 53 are in our BS, all of them are young, and 48 of them are members of the CMG.

Pecaut & Mamajek (2016) undertook a spectroscopic survey to search for new K- and M-type members of Sco–Cen. They obtained a sample of 493 solar-mass ($\approx 0.7–1.3$ M$_{☉}$) stars, out of which 156 were previously unknown. Of their 493 members, 120 are in the Crux Cone, 116 are young, and 95 are members in our CMG.

Finally, we mention the excellent review of Preibisch & Mamajek (2008) which, in great detail, describes the pre-*Gaia* DR2 situation with respect to the Sco–Cen OB association. From their paper, we have extracted the 121 stars in the LCC that they discussed and the 90 from the UCL. Ninety percent of their stars have masses larger than 0.8 M$_{☉}$. Out of the 211 stars, 124 are in our area, 95 fall in the Crux Cone, 90 are in the BS, 82 are young, and 62 are members in our CMG.

Among the stars in all these samples, which were found in the Crux Cone, the vast majority are also classified as young by us. The confirmation is especially important when we compare with such LCC samples where members were selected by not only using kinematic criteria (i.e., the last six lines of Table 3). Among these stars in the Crux Cone, 96% are in our BS and more than 90% are in the Young sample. Moreover, about 75% of these stars are members of the moving group CMG. The high-mass members found before the advent of *Gaia* DR2 form only the tip of the iceberg in the CMG moving group. On the other hand, we have another about 800 stars in our Young sample that are not as strictly comoving as the stars in the CMG. They may have a similar star formation history, but may have high relative velocities compared to the bulk of the CMG.

## 3. Stellar Masses and Ages

### 3.1. Individual Mass and Age Determination

The high photometric quality of *Gaia* DR2 enables a straightforward fit between the observed absolute magnitudes G, $G_{BP}$, and $G_{RP}$ and the ones determined from stellar evolution theory. This method has some limitations: it is of course model dependent, and the results will be as good as the assumptions made to derive the isochrones. The past accretion history may influence the results, although Baraffe et al. (2012)
<table><caption>**Table 3**
Comparison of the CMG with Other Data Sets of Members in the LCC</caption><thead><tr><th rowspan="2">Publication</th><th rowspan="2">Total Number</th><th colspan="6">Number of Stars in</th></tr><tr><th>Area</th><th>Crux Cone</th><th>BS</th><th>Old Sample</th><th>Young Sample</th><th>CMG</th></tr></thead><tbody><tr><td>de Zeeuw et al. (1999)-LCC</td><td>180</td><td>179</td><td>129</td><td>123</td><td>13</td><td>110</td><td>72</td></tr><tr><td>de Zeeuw et al. (1999)-UCL</td><td>221</td><td>19</td><td>15</td><td>15</td><td>1</td><td>14</td><td>10</td></tr><tr><td>Hoogerwerf (2000)-LCC</td><td>1036</td><td>1016</td><td>254</td><td>247</td><td>48</td><td>199</td><td>120</td></tr><tr><td>Hoogerwerf (2000)-UCL</td><td>1498</td><td>174</td><td>35</td><td>33</td><td>4</td><td>29</td><td>16</td></tr><tr><td>Rizzuto et al. (2011)</td><td>436</td><td>156</td><td>123</td><td>119</td><td>8</td><td>111</td><td>76</td></tr><tr><td>Mamajek et al. (2002)</td><td>110</td><td>50</td><td>41</td><td>41</td><td>2</td><td>39</td><td>30</td></tr><tr><td>Pecaut et al. (2012)</td><td>138</td><td>68</td><td>50</td><td>48</td><td>6</td><td>42</td><td>28</td></tr><tr><td>Song et al. (2012)</td><td>104</td><td>67</td><td>57</td><td>53</td><td>0</td><td>53</td><td>48</td></tr><tr><td>Pecaut & Mamajek (2016)</td><td>493</td><td>158</td><td>120</td><td>118</td><td>2</td><td>116</td><td>95</td></tr><tr><td>Preibisch & Mamajek (2008)-LCC</td><td>121</td><td>119</td><td>93</td><td>89</td><td>8</td><td>81</td><td>61</td></tr><tr><td>Preibisch & Mamajek (2008)-UCL</td><td>90</td><td>5</td><td>2</td><td>1</td><td>0</td><td>1</td><td>1</td></tr></tbody></table>

**Note.** For each publication, we report the total number of objects of the publication in column 2, and in the following columns, the number of those stars included in various subsamples, as defined in the text.

suggested that this may not affect clusters aged 10 Myr or older, except for the lowest-mass stars (Baraffe & Chabrier 2010). The impact of variability will be limited by the fact that Gaia is multi-epoch and we used the mean of several measurements. Finally, although a fraction of the binaries are resolved by Gaia (see Section 4.3), unresolved binaries may appear brighter and younger than if they were resolved. We do not attempt to quantify those effects.

We used two grids of models: the CIFIST 2011_2015⁶ models for intermediate-mass stars down to the brown dwarf regime (Baraffe et al. 2015) and the PADOVA tracks Parsec 3.0 for<fnref n="7" /> the higher-mass stars (Marigo et al. 2017). We adopted 0.7 $M_\odot$ as the limit between the two regimes, as the derived mass is identical for both models for that mass. We simultaneously determined the mass and age of each object by searching for the nearest model magnitudes to the data, after having linearly refined the model grid. We used the magnitudes in the filters $G$, $G_{BP}$, and $G_{RP}$ and only considered models up to 100 Myr.

We compared the model predictions in all three Gaia bands $G$, $G_{BP}$, and $G_{RP}$ and made a simple 3D match between observed and theoretical data with an inverse relative weight of 1.0, 15.0, and 2.5 respectively, roughly corresponding to the Gaia photometric uncertainties.

We discuss our results and compare with previous age determinations in Section 3.5.

## 3.2. Detection Efficiency

In order to derive the mass functions of the groups, we first needed to determine our detection efficiency, in particular what fraction of members are lost due the selection cuts we applied to obtain a kinematically clean sample, as a function of the $G$ magnitude and the sky location, which conjunctly for Gaia affects most the precision and accuracy of the measurements.

For this, we considered all entries in the Gaia catalog over the sky region of interest, with parallaxes corresponding to the expected values of the groups' members, i.e., $\bar{\omega} > 7$ mas, and any proper motion, in order to have good statistics. However, distant objects entering the nearby sample and hopefully, and

rightfully, removed by our selection cuts, should not lead us to underestimate the detection efficiency. To identify them, we used the Gaia CMD, looking for objects on the stellar main sequence to confirm the quality of their measured parallax and photometry. The main quality cuts that we applied (cf. Section 2.1) are

1. “visibility periods used” cut: in the magnitude range of interest, the cut just removes 1%–3% of the entries. Only 2% of the area has a detection efficiency lower than 70%, and we do not correct for this effect.

2. “unit weight error” cut: this cut removes 41% of the otherwise satisfactory entries over a wide range of apparent magnitudes. Among those stars, 17% fall on the main sequence (defined between the two 10 Myr and 1 Gyr isochrones; see red dots in Figure 8) and may therefore be bona fide nearby stars with a roughly correct Gaia parallax (but possibly an erroneous uncertainty). As we removed those stars from further analysis, we may have lowered our sensitivity.

3. “flux excess ratio” cut: 27% of the otherwise satisfactory stars fail this cut, mostly with $G > 19$ mag (see blue dots in Figure 8), causing a severe loss of sensitivity at the faint end. Among those stars, 11% are located on the main sequence.

To correct our loss of sensitivity, we considered that the stars removed by either the “unit weight error” cut or the “flux excess ratio” cut but falling on the main sequence should be added to the nearby sample. For the more complete sample, for each $G$ magnitude bin, we measured the fraction of stars passing all cuts and further analyzed. In the mass functions, we increased the weight of each member candidate by that fraction.

In addition to those cuts, we selected stars with $\bar{\omega} > 10 \times \sigma_{\bar{\omega}}$. We can readily determine which fraction of our $\bar{\omega} > 7$ mas sample satisfies this cut, with more than 90% completeness down to $G = 19$ mag, and we corrected for it. The luminosity functions of the various samples, i.e., passing or failing the quality cuts, are shown in Figure 9.

Finally, one needed to know the completeness of Gaia before any cut is applied. Ideally, we would compare with a highly pure, deep optical catalog over the Crux sky region, but

<footnote n="6">⁶ phoenix.ens-lyon.fr/Grids/BT-Settl/CIFIST2011_2015/ISOCHRONES/</footnote>

<footnote n="7">⁷ stev.oapd.inaf.it/cmd</footnote>
<figure><img src="image_9.png" /><figcaption>Figure 8. Gaia color–magnitude diagram of all sources in a 20° radius area passing all of the quality cuts (2D histogram with grayscale), except for the unit weight error or flux excess ratio cuts (red and blue dots, respectively; 10% statistics for readability).</figcaption></figure>

<figure><img src="image_10.png" /><figcaption>Figure 9. Total number of Gaia entries in a 20° radius area with $\bar{\omega} > 7$ mas (red) and $\bar{\omega} > 10\sigma_{\bar{\omega}}$ and seven or more visibility periods (green). The pale blue histogram shows all the sources passing all of the quality cuts (2D histogram with color scale) except for the “unit weight error” (top-left panel) or “flux excess ratio” cut (top right), while the blue histogram shows, among those sources, those that are located on the main sequence (see Figure 8). The pink histogram shows the ratio of the complete nearby Gaia sources passing the quality cuts. Our method becomes too sensitive to the model isochones and low number statistics to be reliable for $G > 19$ mag.</figcaption></figure>

there are none. Comparing with the Pan-STARRS1 DR1 catalog (Chambers et al. 2016; Magnier et al. 2016), with a high-purity subset of objects with 20 or more detections, in two distinct fields with the same ecliptic latitude as the Crux field, $\epsilon = -45^\circ$. We found very different 90% completeness limits, from $G = 20.7$ mag at a longitude of 129°, to a brighter limit of $G = 19.2$ mag at 90°. Instead, we turned to the VVV near-infrared survey. Its second data release (Minniti et al. 2017) offers some overlap with the Crux region. We searched for Gaia counterparts to $5.10^5$ VVV stars with null pperrbits in the Y and J bands, within $1''$. In order to transform the VVV near-infrared photometry to the G magnitude, we derived the

color transformation$^8$:

$$G_{YJ} = Y_1 + 2.48(Y_1 - J_1) + 0.63, 0.4 < Y_1 - J_1 < 0.6 \text{ mag,}$$

We then determined that the Gaia 95% completeness limit is fainter than $G_{YJ} = 20$ mag.

Combining all of those effects, we estimated our 50% completeness magnitude of $G = 19$ mag, which corresponds to a mass of $0.02 M_{\odot}$ for a 10 Myr CIFIST isochrone. We found no fainter candidates and did not attempt to constrain the mass function for lower masses.

The kinematic selection may also lead to the loss of some CMG binaries, when their orbital motion is comparable to or larger than the selection box size in the tangential velocity plane (see Figure 5). For face-on, circular orbits, this will affect systems with a total mass $M_{\text{total}}/M_{\odot} > 10^{-2}a$, where a is the semimajor axis in astronomical units. Following Duchêne & Kraus (2013), this may affect a significant fraction of stars later than F, which have typical semimajor axis of tens of astronomical units. Early-type stars, with semimajor axis less than 1 au, may not be well fitted by the single-star astrometric model and have been removed from the Gaia catalog.

### 3.3. Contamination

In the paragraph above, we discussed the completeness of our sample using the observations of Gaia DR2. On the other hand, we also had to estimate a possible contamination of the CMG by “field” stars. Since we selected the 1844 stars of the CMG within a narrow range of tangential velocities (see Figure 5), the contamination should be rather low. However, we cannot exclude the possibility that there are, just by chance, a few young stars with the appropriate $v_\parallel$ and $v_\perp$ velocities but with an inadequate third component, the radial velocity. Provided that such stars have a roughly uniform distribution in Figure 5, we derived the number of contaminants by comparing the average number density of stars in Figure 5 inside and outside the area of the CMG. The 1844 stars of the CMG cover an area of $11 (\text{km s}^{-1})^2$. For comparison, we selected an area of $30 (\text{km s}^{-1})^2$ centered at (26.5, 0) km s$^{-1}$, where we find 47 stars that we rated as a typical background in Figure 5. This gives a possible contamination of 17 stars to our 1844, or about 1%. So, this has an insignificant impact on the mass function of the CMG.

### 3.4. Mass Functions

As stated in Section 3.1, we determined masses of individual members by using the isochrones of Marigo et al. (2017) for masses higher than $0.7 M_{\odot}$. For lower masses, we took the isochrones of Baraffe et al. (2015) as they model quite well the shape of the observed stellar and substellar sequences for our objects in the CMG. Making corrections for incompleteness (Section 3.2), we determined the mass functions for all four groups and for the intergroup members. They are shown in Figure 10, top panel. We caution that unresolved binaries are counted once, and we did not attempt to estimate the system total mass from the mass fitted in the combined photometry. We studied the case of resolved binaries in more detail in Section 4.3.

<footnote>$^8$ Because the VVV fields cover the Galactic plane with $|b_{\text{gal}}| < 2^\circ$, this relation depends strongly on extinction in the area and stellar sample considered. Here, $l_{\text{gal}} = 295^\circ-318^\circ$.</footnote>
<figure><img src="image_11.png" /><figcaption>**Figure 10.** Top: mass functions of all four groups (in color), the intergroup members (Z), and the whole cluster (black), together with a log-normal fit. Bottom: mass function of the whole cluster before (dashed histogram: stellar counts) and after (solid histogram) correction for incompleteness. We include various canonical mass functions, which we did not normalize for readability, as well as the power-law fit of the $\sigma$ Ori cluster (Peña Ramírez et al. 2012).</figcaption></figure>

We fitted a log-normal function (Scalo 1986) to our data and present the mass functions of Salpeter (1955), Chabrier (2003), and Kroupa (2002) for comparison (see Figure 10, bottom). We found that all the mass functions of our subgroups are very similar and well reproduced by a single log-normal function of mean mass $m_c = 0.22 M_\odot$ with a dispersion of 8%, and a very consistent standard deviation $\sigma = 0.64$, hence slightly wider than the canonical Chabrier (2003) mass function (for systems) with $\sigma = 0.55$. In particular, we found no difference between the groups' mass functions and that of the intergroup members. Hence, the latter are not only low-mass members ejected by dynamical interactions. There is an excess at about $1.4 M_\odot$, which may have been caused by an imperfect separation between young and old stars.

The added-up masses of the members detected in each group are listed in Table 1. The total mass of the stars and brown dwarfs in the CMG is 673 $M_\odot$.

We can compare our mass function to the mass functions reported for other young clusters or associations. Peña

Ramírez et al. (2012) also used the isochrones from the Lyon group, NextGen98 (Baraffe et al. 1998), AMES-Dusty (Chabrier et al. 2000), and AMES-Cond (Baraffe et al. 2003) to determine the low-mass part of the mass function of the $\sigma$ Orionis cluster with an age of 2–4 Myr. They slightly favor power laws over a log-normal form, with $\alpha = 0.6$ for masses less than $0.35 M_\odot$, where $\alpha$ is the index of the power law: $\Delta N/\Delta M = M^{-\alpha}$. Their log-normal fits in this mass range have $\sigma$ between 0.57 and 0.64, so are comparable to or slightly smaller than our value. They note, however, that this mass function considerably underestimates the number of planetary mass objects ($m < 0.012 M_\odot$), which is outside the mass range we are covering in this paper. While $\sigma$ Orionis is a very compact cluster with a characteristic radius of only 1 pc (Béjar et al. 2011), the CMG is a more extended group of young objects. So, the star formation mode of both may differ.

Lodieu (2013) studied the Upper Scorpius part of the Sco–Cen association, and found that for masses higher than $0.03 M_\odot$, it is well reproduced by the log-normal form of the Chabrier mass function, independent of the age one determines for Upper Scorpius. This is a steeper decline from $m_c = 0.22 M_\odot$ toward $0.03 M_\odot$ than what we found for the neighboring LCC part. Like Peña Ramírez et al. (2012) for $\sigma$ Orionis, Lodieu (2013) found an excess of objects with masses lower than $0.03 M_\odot$ depending on the adopted age. Again, this is beyond the scope of this paper.

Very recently, Mužić et al. (2017) studied the deeply embedded young (≈1 Myr), massive and very dense star cluster RCW 38, and found $\alpha = 0.71 \pm 0.11$ for masses between 0.02 and $0.5 M_\odot$, a decline toward low masses even shallower than in $\sigma$ Orionis.

Summarizing, we found that the decline of our mass function from the mean mass of $0.22 M_\odot$ down to $0.03 M_\odot$ is steeper than that found for the $\sigma$ Orionis cluster and RCW 38, but shallower than in the case of Upper Scorpius found by Lodieu (2013).

### 3.5. Age Distributions

As mentioned above, we simultaneously determined the model masses and ages of individual members. For the mass and age determination of low-mass stars and substellar objects, we relied on the CIFIST isochrones. First, because they model the observed loci of the object quite satisfactorily in all Gaia passbands, whereas in the case of the Parsec 3.0 isochrones the magnitudes in the $G_{RP}$ passband come out way too faint. This leads to the situation where ages from Parsec 3.0 are lower than the CIFIST ones when $G_{RP}$ is included and higher when it is not. The CIFIST isochrones for ages between 5 and 40 Myr are mostly parallel to each other in the Gaia CMD.

In Figure 11 (top), we show the age distribution for each subgroup. The bulk of the ages is between 5 and 20 Myr with few nominally older objects (note the logarithmic scale here). There are slight differences in the average ages of the subgroups of the CMG (shown in Table 1), increasing from as low as 7 Myr in A0 to 10 Myr for group C and the intergroups members.

Figure 11 (bottom) shows the star formation rate in the CMG. We do not distinguish here between the subgroups. Star formation essentially started 20 Myr ago and reached its maximum between 8 and 10 Myr before present with a rate of $8 \times 10^{-5} M_\odot \text{ yr}^{-1}$.

Those ages are similar to those reported by Song et al. (2012), derived from Li $\lambda$ 6708 absorption strength analysis. These authors determined a mean age of 10 Myr, without reporting individual ages. We note that we recovered 84% of their sample
<figure><img src="image_12.png" /><figcaption>**Figure 11.** Top: CIFIST age distribution for each subgroup and the intergroup sample (Z). We stress that the stellar counts, measured for bins of width 5 Myr, are not independent. Bottom: the star formation rate in the CMG as a whole as a function of time.</figcaption></figure>

over the area we studied (see Table 3). Although their results are model dependent, just like ours are, the two methods are independent and, if any, would have different biases. By comparing the lithium absorption strength, Song et al. (2012) derived a younger age for their LCC targets than for the β Pictoris moving group members, which they estimated at 12 Myr.

We can compare on an individual basis our ages with those of Pecaut & Mamajek (2016) for the sample of 95 stars in common. They agree well to within 15% on average for the low-mass stars (our mass $m < 0.8 M_\odot$), albeit with a large dispersion. However, the ages of Pecaut & Mamajek (2016) for higher-mass stars ($m > 1 M_\odot$) are 30% larger than ours. In fact, Pecaut & Mamajek (2016) detected a correlation between effective temperature and mean ages, while we did not.

On the other hand, for the young population in the LCC, Mamajek et al. (2002) found mean ages of the pre-main-sequence stars from *Hipparcos* and TYC/ACT to range between 17 and 23 Myr. When using the Bertelli et al. (1994) tracks, they determined main-sequence turnoff ages for *Hipparcos* B-type members to be $16 \pm 1$ Myr. It is worth mentioning that much earlier, de Geus et al. (1989) and de Zeeuw & Brand (1985) obtained turnoff ages between 10 and 12 Myr, so much closer to the ages determined here. The difference between these earlier attempts to get an age for the LCC and ours lies in the availability of *Gaia* DR2. Indeed, we

detected a clear correlation ($R \ge 0.6$) between the relative offset of the kinematic parallaxes used in Preibisch & Mamajek (2008) and Pecaut & Mamajek (2016) to the *Gaia* DR2 parallaxes, and the relative offset of the ages derived in those articles, to ours. An overestimated parallax leads to an underestimated absolute luminosity and to an overestimated age. In addition, the majority of our members have masses way below those that could be used in the papers above, and the photometric quality is much superior.

# 4. Remarkable Members

## 4.1. Circumstellar Disks

With ages below 20 Myr, a fraction of our CMG members is likely to harbor a circumstellar disk at various stages of evolution.

We have used the Virtual Observatory SED analyzer (VOSA; Bayo et al. 2008) to analyze the spectral energy distribution of our members. In a nutshell, VOSA adjusts theoretical isochrones available in the virtual observatory (VO) to the photometric information provided by the user and/or collected on the VO by $\chi^2$ minimization. It detects deviations from the photospheric emission, such as the infrared excess of interest here, and ignores the corresponding data points when the fit is repeated. We used all available VO photometry and fitted the data with the CIFIST models with solar metallicity (Baraffe et al. 2015), refining once the infrared excess detection. For A- and B-type stars, which the CIFIST grid does not cover, we used the ATLAS9 atmosphere models (Castelli et al. 1997), again with solar metallicity and $E(B-V) = 0$. The best model parameters are listed in Table 4. We then selected for further inspection those stars with an excess of at least 0.25 mag and $5\sigma$, in at least one filter. We removed by eye a fraction of the candidates with clearly poor fits. Finally, we combined the official quality flags provided by the AllWISE catalog (Wright et al. 2010) and the results of a careful visual inspection of the relevant *WISE* images.

### 4.1.1. *Identification of Spurious WISE Detections*

Most of our excess candidates are located within $10^\circ$ of the galactic plane. The higher source density as well as the bright extended background emission results in enhanced contamination and spurious detections in this region (Koenig & Leisawitz 2014). We took special care to identify and discard candidate sources affected by such issues. The AllWISE catalog already provides useful information on the data quality. We removed all targets where detection in any of the four bands was qualified to be spurious either because of a diffraction spike (“D” flag), scattered light halo (“H”), optical ghost image (“O”) or latent image (“P”) from a nearby bright source based on the “confusion and contamination flag” parameter of the AllWISE catalog. Sources found to be extended in the W3 or W4 bands (i.e., the reduced $\chi^2$ of their profile-fit photometry, $w3rchi2$ or $w4rchi2$ is higher than 3) were also discarded. Furthermore, to ensure good quality data, only photometry with signal-to-noise ratio higher than 5 was kept in our analysis. According to Koenig & Leisawitz (2014), especially in regions with high sky background, a significant fraction of spurious detections may have remained unrecognized in the AllWISE catalog. To identify additional potentially problematic cases, we inspected the *WISE* images of our excess candidates visually. This process revealed several additional fake sources, and also objects whose photometry is
<table><caption>**Table 4**

Content of the Electronic Table of the Member Parameters</caption><thead><tr><th>Name</th><th>Description</th><th>Unit</th><th>Range</th></tr></thead><tbody><tr><td>designation</td><td>Gaia DR2 designation</td><td>...</td><td>...</td></tr><tr><td>Group</td><td>group membership</td><td>...</td><td>A0, A, B, C, Z</td></tr><tr><td>ra, dec, ...</td><td>Gaia DR2 parameters</td><td>...</td><td>...</td></tr><tr><td>vr_pred</td><td>predicted radial velocity</td><td>km s<sup>-1</sup></td><td>5.56 ~ 16.6</td></tr><tr><td>vpar</td><td>observed tangential velocity toward the convergent point</td><td>km s<sup>-1</sup></td><td>18.4 ~ 21.4</td></tr><tr><td>vperp</td><td>perpendicular to the convergent point direction</td><td>km s<sup>-1</sup></td><td>-2.3 ~ +2.0</td></tr><tr><td>e_vpar</td><td>uncertainty on vpar</td><td>km s<sup>-1</sup></td><td>0.053 ~ 1.28</td></tr><tr><td>e_vperp</td><td>uncertainty on vperp</td><td>km s<sup>-1</sup></td><td>0.012 ~ 0.41</td></tr><tr><td>vparvperpcor</td><td>(dimensionless) correlation between</td><td>...</td><td>-0.83 ~ 0.79</td></tr><tr><td>X_gal</td><td>X galactic coordinate</td><td>pc</td><td>+26.1 ~ +102.7</td></tr><tr><td>Y_gal</td><td>Y galactic coordinate</td><td>pc</td><td>-134.0 ~ -64.6</td></tr><tr><td>Z_gal</td><td>Z galactic coordinate</td><td>pc</td><td>-30.6 ~ +51.0</td></tr><tr><td>age_CIFIST</td><td>age according to the CIFIST isochrones</td><td>Gyr</td><td>10<sup>-3</sup> ~ 10<sup>-1</sup></td></tr><tr><td>mass_CIFIST</td><td>mass according to the CIFIST isochrones</td><td>M<sub>☉</sub></td><td>0.026 ~ 1.4</td></tr><tr><td>Teff_CIFIST</td><td>effective temperature according to the CIFIST isochrones</td><td>K</td><td>2632 ~ 6766</td></tr><tr><td>logL_CIFIST</td><td>log<sub>10</sub> of bolometric luminosity according to the CIFIST isochrones</td><td>CGS</td><td>-2.8 ~ 0.6</td></tr><tr><td>logg_CIFIST</td><td>log<sub>10</sub> of gravity according to the CIFIST isochrones</td><td>CGS</td><td>3.4 ~ 4.9</td></tr><tr><td>age_P</td><td>age according to the Parsec isochrones</td><td>years</td><td>5 10<sup>6</sup> ~ 10<sup>8</sup></td></tr><tr><td>mass_P</td><td>(initial) mass according to the Parsec isochrones</td><td>M<sub>☉</sub></td><td>0.1 ~ 5.3</td></tr><tr><td>logL_P</td><td>log<sub>10</sub> of bolometric luminosity according to the Parsec isochrones</td><td>L<sub>☉</sub></td><td>-2.7 ~ 3.0</td></tr><tr><td>logTe_P</td><td>log<sub>10</sub> of effective temperature according to the Parsec isochrones</td><td>K</td><td>3.4 ~ 4.2</td></tr><tr><td>logg_P</td><td>log<sub>10</sub> of gravity according to the Parsec isochrones</td><td>CGS</td><td>3.8 ~ 4.8</td></tr><tr><td>recovery_rate</td><td>fraction of recovered stars of same G magnitude</td><td>...</td><td>0.648 ~ 1.</td></tr><tr><td>age</td><td>best age</td><td>Gyr</td><td>10<sup>-3</sup> ~ 10<sup>-1</sup></td></tr><tr><td>mass</td><td>best mass</td><td>M<sub>☉</sub></td><td>0.026 ~ 5.3</td></tr><tr><td>Teff</td><td>best effective temperature</td><td>K</td><td>2632 ~ 16.10<sup>3</sup></td></tr><tr><td>LogL</td><td>best effective log<sub>10</sub> of bolometric luminosity</td><td>K</td><td>-2.8 ~ 3.0</td></tr><tr><td>AllWISE_ID</td><td>ID in the AllWISE catalog</td><td>...</td><td>...</td></tr><tr><td>logL_SED</td><td>log<sub>10</sub> bolometric luminosity according to the VOSA fit</td><td>Lbol<sub>☉</sub></td><td>-3.6 ~ 2.4</td></tr><tr><td>Teff_SED</td><td>effective temperature according to the VOSA fit</td><td>K</td><td>1200 ~ 7000</td></tr><tr><td>disk_type</td><td>Circumstellar disk classification</td><td>...</td><td>...</td></tr><tr><td>SIMBAD_Name</td><td>Primary SIMBAD name</td><td>...</td><td>...</td></tr><tr><td>Disk_ref</td><td>Reference for circumstellar disk</td><td>...</td><td>1 ~ 18</td></tr></tbody></table>

**Note.** Several Gaia DR2 columns, which are included for convenience without any change, are not described below. The content of the last seven columns is described in Section 4.1. In particular, the VOSA columns report the SED fit using either the CIFIST Grid, or for the A- and B-type stars, the ATLAS9 Grid. Reference for disk identification: 1—Carpenter et al. (2005), 2—Chen et al. (2005), 3—Chen et al. (2011), 4—Chen et al. (2012), 5—Clarke et al. (2005), 6—Cotten & Song (2016), 7—Cruz-Saenz de Miera et al. (2014), 8—Fujiwara et al. (2009), 9—Gregorio-Hetem et al. (1992), 10—Mannings & Barlow (1998), 11—Marton et al. (2016), 12—Melis et al. (2012), 13—Melis et al. (2013), 14—Oudmaijer et al. (1992), 15—Pecaut & Mamajek (2016), 16—Rizzuto et al. (2012), 17—Schneider et al. (2012), 18—this work.

seriously contaminated by nearby bright sources or background nebulosity. The applied quality criteria reduced the size of our sample to 170 systems. For 136 of them, both the W3 and W4 band data were found to have good quality; for 34 objects, the W4 band photometry was not used in the following analysis.

### 4.1.2. Classification of Disk Candidates

The observed infrared excesses, outlined by the WISE photometry, indicate the presence of circumstellar disks around our 170 selected stars. To assess the evolutionary stage of the revealed disks, we applied the classification scheme proposed by Luhman & Mamajek (2012) and Pecaut & Mamajek (2016). This scheme allows four different classes to be distinguished: full disks, transitional disks, evolved disks, and debris disks (for details of the different classes, see Espaillat et al. 2012) based on the $E(K_s - W3)$ and $E(K_s - W4)$ color excesses of the objects. To compute these color excesses, the predicted photospheric $K_s - W3$ and $K_s - W4$ colors were taken from the VOSA models described above. Where no good quality W4 photometry

was available, the full, transitional, and evolved disk categories could not be distinguished unambiguously based on the applied criteria. For these cases, we could discern three different categories based on the $E(K_s - W3)$ excesses: full or evolved disks, transitional or evolved disks, and debris disks. In three special cases, the resulting classification was reconsidered based on literature data. Using the above scheme, HD 110058, HD 121191, and HD 121617 were classified as primordial disks. However, though these objects exhibit strong mid-infrared excess and harbor detectable amounts of CO gas (Lieman-Sifry et al. 2016; Moór et al. 2017), they were generally categorized as debris disks based on their spectral energy distributions and dust masses (Melis et al. 2013; Lieman-Sifry et al. 2016; Moór et al. 2017). Considering these more detailed studies, the three targets finally have been classified as debris disks. Color excesses as well as the obtained classifications are displayed in Figure 12. As a result of our analysis, we identified 115 primordial (full disks, transitional disks, or evolved disks) and 55 debris disks in our sample. The objects with disks are listed in the
<figure><img src="image_13.png" /><figcaption>**Figure 12.** Color excesses for disk candidates having good quality photometry in both W3 and W4 bands. Classification of the disks is based on criteria proposed by Luhman & Mamajek (2012) and Pecaut & Mamajek (2016; for details, see Section 4.1.2).</figcaption></figure>

machine-readable table described in Table 4, along with their main Simbad name and the discovery reference.

We note that 15 debris disks of the Sco–Cen association listed by Chen et al. (2011, 2012) are not recovered by our search. Nine are not in our CMG sample, generally because of poor Gaia data, an isochrone age larger than 40 Myr, or because their parallax is smaller than 7 mas.

## 4.2. Brown Dwarfs

Given the small distance of the LCC and its young age, Gaia is able to detect even low-mass brown dwarfs in this region. Among the 1844 objects in the CMG sample, there are 214 with masses smaller than $0.073 M_\odot$, i.e., brown dwarfs. Only for one of them, [FLG2003] eps Cha 16, did we find an entry in SIMBAD. This young stellar object (YSO in SIMBAD) has spectral type M5.75 (we determined $M_G = 11.51$, mass = $0.05 M_\odot$, age = 7 Myr), is located at the farthest edge of the CMG sample, and is a member of our A0 subgroup. Among those 214 brown dwarfs, 34 harbor a circumstellar disk. While we are not complete with respect to transitional and debris disks that are too faint to be detected with our WISE selection, we note than $(14 \pm 3)$% of the brown dwarfs bear full disks or full/evolved disks. This is eight times higher (a $5\sigma$ effect) than for stars with a fraction of $(1.6 \pm 0.5)$%, possibly meaning that brown dwarfs keep their disks optically thick at infrared wavelengths longer than stars.

## 4.3. Resolved Binaries

The Gaia DR2 catalog has a nominal spatial resolution of 0.4 arcsec, i.e., 46 au at the mean member distance. According to Duchêne & Kraus (2013), about half the stellar binaries made of components with masses between 0.7 and $5 M_\odot$ have semimajor axis larger than that separation. These binaries, through their additional orbital motion, perturb the study of the internal dynamics of the cluster and may lead to incompleteness (see Section 3.2). In addition, both components of resolved binaries would be counted separately in luminosity and mass functions, such as those that we analyzed in Section 3.4, while unresolved binaries are counted only once, so it is important to identify resolved binaries in order to derive coherent luminosity and mass functions. Finally, the orbits of high-mass systems with semimajor

axis of a few tens of astronomical units will ultimately be determined by Gaia, and therefore the dynamical masses of their components will be measured.

We searched our member catalog for pairs sharing the same proper motions and parallaxes. To estimate the contamination by random pairing, we used the standard procedure of creating a mock catalog of stars derived from the observed sample, shifting one coordinate by some angle, so that any pairing between the observed and mock catalogs can only be a spurious match (see Lépine & Bongiorno 2007 for an illustration of the method). To estimate the quality of the match, we calculated the sum $\Sigma$ of the differences between the proper motions and parallaxes, expressed in $\sigma$. We include in the error budget the orbital motion, which will introduce a difference in the proper motions, calculated for a face-on, circular orbit.

We selected observed pairs up to 10 arcmin and $\Sigma < 9$. For each pair, we determined the number of spurious pairs with parameters (sky separation, $\Sigma$) smaller than the observed pair, i.e., better matches. To improve the statistics, we shifted the observed sample, either the R.A. or decl., by 0.5, 1, and $1^\circ$5. We identified eight pairs, with separation up to $3'$, for which the number of better spurious matches is smaller than 2.2. They are listed in Table 5. The next best candidate, with a separation of 400'' and $\Sigma = 7.2$, have 20 better spurious matches.

Although the statistics is limited, we find that the mass ratios increase from 0.1 to 0.2 for F and G stars to almost unity at the end of the main sequence, as has been found in previous studies (for a review, see Duchêne & Kraus 2013). The lowest-mass pair we identified has a total mass of $0.4 M_\odot$, and the closest separation is 3 arcsec.

# 5. The Motion of the CMG

Table 1 gives the full 6D average phase-space coordinates for the four subgroups of GCYS in the Galactic Cartesian coordinate system. These coordinates are determined exclusively from the observations in Gaia DR2, i.e., only for those members that have measured radial velocities in DR2. In Table 1, we find correlations between the spatial coordinates and the respective velocities of the subgroups. This suggests an expansion of the whole CMG. The good accuracy of the mean space velocity components (median $0.2 \text{ km s}^{-1}$) allows the groups to be traced back in time to find the minimum distance of the groups from each other. Given the ages of the stars in the CMG, we trace them back for not more than about 20 Myr. In the Galactic Cartesian coordinate system, we used a simple 3D linear approach, $X(t) = X(0) + t \times V$, with the components of the space vector $X$ in pc and the components of the velocity vector $V$ in pc Myr$^{-1}$ ($1 \text{ km s}^{-1} = 1.0227 \text{ pc Myr}^{-1}$). We traced back the centers of the subgroups from Table 1 in steps of 0.1 Myr using the data given in this table. To characterize the size of the CMG, we used the quantity $R = \sqrt{\sum_{i=0}^{N} \sum_{j=i+1}^{N} r_{ij}^2}$, where $r_{ij}$ is the 3D distance between subgroup i and subgroup j. We found that R had a minimum at $-10 \text{ Myr}$. That means that the centers of the four subgroups had its closest distance from each other 10 Myr before present and have been drifting away from each other since. This finding is compatible with the ages estimated for the individual stars in the CMG.

In this connection, the question arose of whether the subgroups we see now were also separated from the beginning, or whether the CMG is expanding as a whole. To test this hypothesis, it is necessary to trace back the individual stars of
<table><caption>**Table 5**
Binary Members</caption><thead><tr><th>Gaia ID 1</th><th>Group 1</th><th>Mass 1 (M<sub>☉</sub>)</th><th>Gaia ID 2</th><th>Group 2</th><th>Mass 2 (M<sub>☉</sub>)</th><th>Separation (arcsec)</th><th>Contaminants</th></tr></thead><tbody><tr><td>6075310478057303936</td><td>C</td><td>0.75</td><td>6075310478050174592</td><td>C</td><td>0.24</td><td>3.232</td><td>0</td></tr><tr><td>6065138792911302784</td><td>Z</td><td>1.62</td><td>6065138792911300096</td><td>Z</td><td>0.13</td><td>19.413</td><td>0</td></tr><tr><td>5861825550891696000</td><td>A</td><td>0.20</td><td>5861825349107118976</td><td>Z</td><td>0.19</td><td>31.073</td><td>0</td></tr><tr><td>6065047911406764288</td><td>Z</td><td>1.37</td><td>6065047636528853888</td><td>Z</td><td>0.26</td><td>51.768</td><td>0</td></tr><tr><td>5894194318558985984</td><td>Z</td><td>1.17</td><td>5894194348576228992</td><td>Z</td><td>0.66</td><td>59.699</td><td>0.2</td></tr><tr><td>6087959014999185152</td><td>Z</td><td>0.92</td><td>6087958675701081600</td><td>Z</td><td>0.10</td><td>113.306</td><td>0.7</td></tr><tr><td>6090729754245639552</td><td>Z</td><td>0.39</td><td>6090730132201936256</td><td>Z</td><td>0.18</td><td>149.843</td><td>1.2</td></tr><tr><td>5854812388999707008</td><td>Z</td><td>1.60</td><td>5854813312408616320</td><td>A0</td><td>0.30</td><td>166.388</td><td>3.4</td></tr></tbody></table>

the CMG sample, which means we need accurate radial velocity measurements for a representative subset of stars in the CMG. In Gaia DR2 (Katz et al. 2018), we find observed radial velocities for 247 stars of the CMG. The formal precision of the radial velocity measurements varies between 0.2 and more than $10 \text{ km s}^{-1}$ with a median of $2 \text{ km s}^{-1}$, which is at least one order of magnitude inferior to the tangential velocities. To reduce the influence of the poorly measured radial velocities, we considered only stars (123 in number) having $\sigma_{V_{\text{obs}}}$ better than $2 \text{ km s}^{-1}$ (median of $1 \text{ km s}^{-1}$) for the back-tracing. The lower accuracy of the radial velocities influences primarily the Y coordinate, because it is closest to the line of sight. To a lesser extent, this holds for the X coordinate, while the Z coordinate is little affected.

As in the case of the subgroups, we also find correlations between the spatial coordinates and the respective velocity components of the 123 stars. In Figure 13, we show the velocities U, V, W versus the coordinates X, Y, Z. A linear fit between coordinates and velocities delivers slopes of $0.087 \pm 0.008$ in X, $0.064 \pm 0.016$ in Y, and $0.067 \pm 0.004 \text{ km s}^{-1} \text{ pc}^{-1}$ in Z. All of the slopes are significantly positive, which indicates expansion in all three spatial directions. As expected, the expansion rate is worst determined ($4\sigma$) in the Y-direction, reflecting the lesser accuracy of the radial velocities.

Knowing that GCYS is expanding, we performed a full 3D back-tracing for the 123 stars in the Galactic Cartesian coordinate system to determine the time of minimum extension. We evaluated the quantity $R$, which we rewrite as

$$R^2 = \sum_{i=0}^{N} \sum_{j=i+1}^{N} X_{i,j}^2 + \sum_{i=0}^{N} \sum_{j=i+1}^{N} Y_{i,j}^2 + \sum_{i=0}^{N} \sum_{j=i+1}^{N} Z_{i,j}^2$$

with $X_{i,j}^2 = (X_i - X_j)^2$ and so on. This allowed the influences on $R$ by the motions along the X, Y, Z axes to be studied. At $t = 0$, the extension of the CMG is largest in the z-direction, and hence the partial sum over $Z_{i,j}^2$ is largest. Going back in time, the partial sums are steadily decreasing. Already, at $-5$ Myr, the partial sum in Y has reached its minimum, strongly increasing afterwards and dominating the value of $R$. We rate this as a kind of “virtual expansion” caused by the poor knowledge of radial velocities (see also the scatter in the middle panel of Figure 13).

If we ignore the Y-component and take only the 2D differences, the minimum for the 2D $R$ is reached at $-7.2$ Myr for all 123 stars. We show the back-tracing results in the X, Z plane in Figure 14. From left to right, we show the distributions of the 123 stars at present, 9 Myr ago, and 18 Myr ago. The positions of the stars 18 Myr ago are shifted by 50 pc in the positive Z-direction to avoid overlap with the situation at

<figure><img src="image_14.png" /><figcaption>**Figure 13.** Galactic Cartesian velocities U, V, W vs. coordinates X, Y, Z (from top to bottom) for 123 stars in the CMG with radial velocities having $\sigma_{V_{\text{obs}}}$ better than $2 \text{ km s}^{-1}$. The regression fits with slopes of $0.087 \pm 0.008$, $0.064 \pm 0.016$, and $0.067 \pm 0.004 \text{ km s}^{-1} \text{ pc}^{-1}$ are shown in red.</figcaption></figure>

$-9$ Myr. Already, at $-9$ Myr, and even more at $-18$ Myr, we notice 11 outliers. If we exclude the latter from the calculation of the radius $R$, we find the minimum of the 3D $R$ at $-7.0$ Myr and of the 2D $R$ at $-9.2$ Myr.

We conclude that the CMG as a whole is expanding since at least $9-10$ Myr (group centers), and the subgroups A0 to C, which we observe today, were parts of a single structure at that time. The clumps are probably the results of slightly different initial space velocities. We cannot exclude that expansion already started more than 10 Myr ago as back-tracing is hampered by the poor knowledge of radial velocities ($2 \text{ km s}^{-1}$ at 10 Myr make an offset of 20 pc). Although 70% of the stars in the CMG were estimated to be younger than 10 Myr, the onset of star formation could have started earlier, and a later event triggered expansion.

Wright & Mamajek (2018) explicitly stated that they found no indication of expansion in their careful analysis of the kinematic history of the LCC on the basis of Gaia DR1, except that they got evidence for expansion along the Y-axis with a significance of almost $2\sigma$ for their LCC sample. Also, de Geus (1992) did not
<figure><img src="image_15.png" /><figcaption>**Figure 14.** Density distribution of 123 CMG stars having accurately measured radial velocities (see text). The leftmost distribution shows the configuration at present, the central one from 9 Myr ago, and the right one from 18 Myr ago. The latter is shifted by 50 pc in the positive Z-direction to avoid overlap with the configuration at −9 Myr.</figcaption></figure>

detect an expansion in the region of Crux, but from a completely different approach. He examined the distribution of neutral hydrogen gas (H1) connected to the Sco–Cen association. Whereas the column density of H1 is low over most of the area of CGRS, he found an H1 loop (shell) at (l, b) = (285°, +18°) and determined its center to be at (l, b) = (300°, +8°) from the curvature of the loop. He estimated the mass of the molecular cloud in the loop to be (1.0 ± 0.5) × 10⁵ M☉. Remarkably, the center of the loop is situated in the middle between our groups B and C, and the loop itself marks the northern border of group C at this Galactic longitude. From his H1 21 cm observations, he found no evidence for this loop to be part of an expanding shell.

Contrary to de Geus (1992), Ortega et al. (2009), who also analyzed the motion of the H1 shell in Crux, found that the shell had been expanding. Between 13 and 9 Myr before present, the shell increased its radius from 15 to almost 30 pc, but then met together with the UCL shell. In this case, this collision could have triggered the star formation of the CMG, which had its maximum about 9 Myr ago.

## 6. Summary and Discussion

The data from the second *Gaia* data release are an invaluable tool to study the kinematics of stellar ensembles in the Milky Way. Moreover, due to its excellent photometric data, *Gaia* enables accurate determinations of stellar masses and ages from color–absolute-magnitude diagram. However, it becomes apparent that the wealth of new observations of low-mass stars in the *Gaia* passbands requires new effort to construct adequate theoretical isochrones. We studied in this paper the large complex in the LCC where a few million years ago, star formation took place. We found some 2800 stars in a cone around (α, δ) = (186°:5, −60°:5) with a radius of 20°, which we could rate as younger than 40 Myr based on the location in the (M<sub>G</sub>, G − G<sub>RP</sub>) CMD. Among these, 1844 stars share tangential motion compatible with a common 3D

space motion. Their distances from the Sun range from 102 to 135 pc. The total mass of the stellar and substellar objects is about 700 M☉. On top of this common motion the whole group is expanding, and expansion started about 9 Myr ago. Also 9 Myr ago the peak in the star formation rate occurred with 8 × 10⁻⁵ M☉ yr⁻¹.

So, the event that was responsible for the beginning of expansion could also have triggered a higher star formation rate. Or, did a collision between the UCL and LCC molecular shells occur 9 Myr ago, as Ortega et al. (2009) claim, and did such a collision trigger star formation?

The ages we determined for the stars in the CMG are consistent with the age of the LCC according to Song et al. (2012). Moreover, although Pecaut & Mamajek (2016) determined an age of 17 Myr for their LCC sample, their average age of the subset in common with the CMG is identical with ours within 20%, proving that the mean age of the CMG is quite reliable. Similarly, we find that the LCC association is much younger than the 18 Myr found by Mamajek et al. (2002). This may have the following reason. Their age determination was based upon main-sequence turnoff ages for *Hipparcos* B-type members. The vast majority of stars in our sample have masses less than 0.3 M☉, and the ages of these pre-main-sequence stars and brown dwarfs are determined by comparing the three *Gaia* DR2 photometric bands (absolute magnitudes) with the theoretical isochrones of CIFIST, and we rely on these age marks. As the ages from Mamajek et al. (2002) are based upon high-mass stars, there may be some inconsistency on the models underlying both methods.

Given the youth of the group, the present-day mass function should not differ too much from the initial mass function of the complex. We find that a log-normal mass function with mean mass m<sub>c</sub> = 0.22 M☉ and a standard deviation σ = 0.64 fits the observations quite well between 0.05 and 1.0 M☉. We observe, however, a steeper decline in the brown dwarf regime at masses lower than 0.05, which cannot be attributed to the incompleteness of our sample. In total, the CMG contains 214 brown dwarfs. There is also an excess at about 1.4 M☉ which may have been caused by an imperfect separation between young and old stars near this mass. Our log-normal mass function differs from the canonical Chabrier mass function (Chabrier 2003) essentially because of its broader standard deviation (0.64 versus 0.55). A Chabrier-type mass function was also found by Lodieu (2013) when he studied the star formation in the Upper Scorpius part of Sco–Cen. While the slope of the Kroupa mass function (Kroupa 2002) describes the high-mass end quite well, it fails at low-mass stars and brown dwarfs. First, the observed maximum of our mass function is at higher masses (0.2 observed versus 0.1 M☉ in Kroupa), and our slope in the brown dwarf regime is much steeper. The latter is also valid for a comparison of our mass function in the LCC with the mass functions in the σ Orionis cluster (Béjar et al. 2011; Peña Ramírez et al. 2012) and the very young embedded open cluster RCW 38 (Mužić et al. 2017).

With members' ages ranging from 5 to 20 Myr, we detected no less than 170 circumstellar disks showing a wide variety of evolution, from full disks to debris disks. Among these, 34 are associated with brown dwarfs. The brown dwarf disks are in its majority full/evolved, and 14% of the brown dwarfs in our sample carry these.

Summing it all up, the LCC is a rich association very close to the Sun, and this makes it an ideal laboratory for
studying very young stars and very young brown dwarfs and their circumstellar disks. The total stellar mass of our sample in the LCC is about 700 $M_\odot$, and 80% of its stars are located between 102 and 135 pc from the Sun, with a median distance of 114.5 pc. According to the ages of our objects determined from the CIFIST isochrones, star formation started about 20 Myr before present and reached its maximum 9 Myr ago. This is equal to the expansion age of our CMG. Hence, the event that was responsible for this expansion may also have been an additional trigger for star formation.

Our findings and our comparison with previous studies of young stars in the western part of Sco–Cen show how the much deeper and more accurate astrometry of Gaia will shed light on moving groups, and thus will require a similarly more structured classification of the known moving groups.

We thank the anonymous referee for comments that helped improve the manuscript. We are grateful to Ulrich Bastian from ZAH for very helpful discussions on the astrometric quality of Gaia DR2 and on the procedures needed to obtain and astrometrically clean the sample of stars.

This study was supported by Sonderforschungsbereich SFB 881 “The Milky Way System” (subprojects B5 and B7) of the German Research Foundation (DFG). This research has made use of the SIMBAD database and of the VizieR catalog access tool, operated at CDS, Strasbourg, France.

This work has made use of data from the European Space Agency (ESA) mission Gaia (https://www.cosmos.esa.int/gaia), processed by the Gaia Data Processing and Analysis Consortium (DPAC, https://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.

This publication makes use of VOSA, developed under the Spanish Virtual Observatory project supported from the Spanish MICINN through grant AyA2011-24052.

This publication makes use of data products from the Wide-field Infrared Survey Explorer, which is a joint project of the University of California, Los Angeles, and the Jet Propulsion Laboratory/California Institute of Technology, funded by the National Aeronautics and Space Administration.

Facilities: Gaia, CDS, WISE.

Software: Astropy (Astropy Collaboration et al. 2013).

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