LUX et al.

*Oliver Lux, Schillergäßchen 2, 07745 Jena.

# A search for runaway stars in twelve Galactic supernova remnants††thanks:
Based on service mode observations obtained from ESO-VLT in project
0100.D$-$0314 and NAOJ-Subaru in project S18B0195S.

Oliver Lux  Ralph Neuhäuser  Markus Mugrauer  Richard Bischoff  Astrophysical
Institute and University Observatory, Friedrich Schiller University Jena,
Thuringia, Germany<EMAIL_ADDRESS>
(09 November 2020; –; 07 December 2020)

###### Abstract

Runaway stars can result from core-collapse supernovae in multiple stellar
systems. If the supernova disrupts the system, the companion gets ejected with
its former orbital velocity. A clear identification of a runaway star can
yield the time and place of the explosion as well as orbital parameters of the
pre-supernova binary system. Previous searches have mostly considered O- and
B-type stars as runaway stars because they are always young in absolute terms
(not much older than the lifetime of the progenitor) and can be detected up to
larger distances. We present here a search for runaway stars of all spectral
types. For late-type stars, a young age can be inferred from the lithium test.
We used Gaia data to identify and characterise runaway star candidates in
nearby supernova remnants, obtained spectra of 39 stars with UVES at the VLT
and HDS at the Subaru telescope and found a significant amount of lithium in
the spectra of six dwarf stars. We present the spectral analysis, including
measurements of radial velocities, atmospheric parameters and lithium
abundances. Then we estimate the ages of our targets from the Hertzsprung-
Russell diagram and with the lithium test, present a selection of promising
runaway star candidates and draw constraints on the number of ejected runaway
stars compared to model expectations.

###### keywords:

supernova remnants – stars: fundamental parameters – stars: kinematics

††articletype: ORIGINAL ARTICLE

Deutsche Forschungsgemeinschaft (DFG), project NE 515/57-1.

##  1 Introduction

Runaway stars are characterised by having higher peculiar velocities than
typical field stars, $v_{\text{pec}}\gtrsim 20-30\,\text{km\,s}^{-1}$ (e.g.
Tetzlaff 2013; Renzo . 2019). There are two suggested ejection mechanisms: (i)
In the binary supernova ejection scenario (BES), the secondary component from
a binary or multiple system is ejected after the primary explodes in a
supernova (SN). A binary becomes unbound if more than half of the total system
mass gets suddenly ejected during the SN (Blaauw kick, Blaauw 1961). Systems
that do not fulfil this criterion can still get unbound if the newborn neutron
star gets a sufficiently high kick from an asymmetric SN (Renzo . 2019,
hereafter Rz19). The ejected companion flies away with its former orbital
velocity. The O9.5 star $\zeta$ Oph and PSR B1929+10 were the first pair of a
runaway star and a neutron star suggested to have been together as a massive
binary (Hoogerwerf ., 2000, 2001). However, later works (Chatterjee ., 2004;
Tetzlaff ., 2010; Zehe ., 2018) show that the association is very unlikely.
Recently, Neuhäuser . (2020) connected $\zeta$ Oph to PSR B1706$-$16\. Using
Monte Carlo simulations, they traced both objects back to a common origin in
Upper-Centaurus-Lupus, a subgroup of the Scorpius-Centaurus-Lupus OB
association, $1.78\pm 0.21$ megayears (Myr) ago. The moderate distance
$d=107\pm 4$ pc of the corresponding SN makes it likely that this event
contributed to the 60Fe content that was detected in the Earth crust and ocean
sediments, dated to an arrival time at the Earth $\sim$2 Myr ago (Knie . 1999;
Breitschwerdt . 2016; Wallner . 2016). (ii) In the dynamical ejection scenario
(DES), the runaway star is ejected by gravitational interaction of stars in
dense, multiple stellar systems, e.g. globular clusters or OB associations
(Poveda ., 1967). Examples are AE Aur and $\mu$ Col, which were ejected from
the Trapezium Cluster in the Orion Nebula, possibly by an encounter with the
massive binary $\iota$ Ori (Hoogerwerf ., 2001). This mechanism generally
produces higher velocities than the BES (Rz19), whereas the highest stellar
velocities can be explained by the Hills mechanism: The tidal encounter with
the supermassive black hole in the Galactic centre (Hills, 1988) can eject a
hypervelocity star with up to $v_{\text{pec}}\gtrsim 1000\,\text{km\,s}^{-1}$.

The limit $v_{\text{pec}}\gtrsim 20-30\,\text{km\,s}^{-1}$ is based on
statistical reasons. Tetzlaff (2013) fitted the velocity distribution of
nearby stars with two Maxwellians; one for the regular population and one for
the high-velocity population. Both intersect at $v_{\text{pec}}\approx
25\,\text{km\,s}^{-1}$. However, Rz19 performed an extensive population
synthesis of massive binaries and emphasised that up to 95 % of SN-ejected
main-sequence (MS) companions are expected to have $v_{\text{pec}}\leq
30\,\text{km\,s}^{-1}$ and should be called walkaway stars. In this work, we
will not divide between runaway and walkaway stars. All stars with a possible
BES origin will be classified as runaway star candidates, regardless of their
velocity.

While both ejection scenarios are verified by observations, it is still
unclear whether one of them is dominant. Therefore, it is important to observe
a large number of runaway stars to find whether they are produced by the BES
or the DES. One possibility to confirm a BES origin is to look for SN debris
in the stellar atmosphere, which can principally be done by high-resolution
spectroscopy (e.g. Przybilla . 2008). Another possibility is to identify
runaway stars inside SN remnants (SNRs) or in the vicinity of neutron stars
(e.g. Dinçel . 2015). SNRs are only visible for up to $\sim$0.1 Myr, which
constrains the flight time since the SN and therefore the search radius and
the maximum age of a runaway star. If the trajectory can be traced back to a
common location with the SNR centre and/or the neutron star position in the
past, the star should have been ejected during the SN. We then know the time
and place of the SN and can constrain SNR expansion models, using the
expansion time and the offset between the geometric centre and the actual
explosion site. If the age of the runaway star is known, e.g. from a birth
cluster, by subtracting the flight time we obtain the lifetime of the SN
progenitor and therefore its mass. From the space velocity of the runaway
star, we can determine the pre-SN orbital parameters and the SN birth kick
velocity imparted on the neutron star (e.g. Dinçel . 2015; Neuhäuser . 2020).

Previous searches for runaway stars have mainly focused on OB-type stars. They
are not only brighter, but also, due to their limited lifetime, they did not
travel too far away from their birth location (up to $\sim$820 pc for a B9.5
star travelling with $\sim$25 km s-1). OB runaway stars were searched for both
outside of SNRs (e.g. Hoogerwerf . 2000, 2001; Tetzlaff . 2010; Tetzlaff .
2011; Tetzlaff 2013; Tetzlaff . 2014), and inside SNRs (e.g. Dinçel . 2015;
Boubert . 2017, hereafter Bo17).

Dinçel . (2015) found the B0.5 V runaway star HD 37424 ($v_{\text{pec}}=74\pm
8\,\text{km\,s}^{-1}$) in the SNR S147 and PSR J0538+2817 to have been close
to the geometric centre of the SNR $30\pm 4$ kyr ago, confirming that the BES
does happen. Bo17 confirmed HD 37424 and suggested three further likely
candidates, located in the SNRs Cygnus Loop, HB21 and Monoceros Loop. They did
not exclude late-type runaways, but they considered early-type stars to be
more probable.

For our search in nearby SNRs, we know that an ejected companion has to be
young. A SN progenitor has a MS lifetime of up to 32 Myr (Ekström . 2012,
Table 3, in the case of an $8\,M_{\odot}$ star). The timespans of the
subsequent burning cycles and the lifetime of the SNR ($\sim$0.1 Myr) are
negligible compared to that. For dwarf stars of spectral types mid-F to M, we
can use lithium (Li I) as youth indicator, because it gets depleted in these
stars during the early phases of stellar evolution (e.g. Bodenheimer 1965;
D’Antona Mazzitelli 1984; Covino . 1997; Neuhäuser 1997); we use spectra of
these stars to search for the absorption line at 6707.8 Å. We can estimate the
age by measuring the effective temperature of the stars and the equivalent
width of the Li 6708 Å line and comparing it to stars in clusters with known
ages, e.g. the Pleiades (Soderblom . 1993; Soderblom 20101, 20102).

In this work, we present results of our search for runaway stars of all
spectral types in twelve nearby SNRs with data from the Gaia mission (Gaia
Collaboration, Prusti ., 2016). In Section 2 we explain how the SNRs and
corresponding runaway star candidates were selected. In Section 3 we describe
the observations. In Section 4 we show the spectral analysis and estimate the
ages of our candidates. In Section 5 we discuss the results and possible
implications for the frequency of runaway stars from the BES and, finally, in
Section 6 we conclude our work and give an outlook on future work.

##  2 Candidate selection

###  2.1 Selection of SNRs

Two main resources were used to find suitable SNRs, namely (i) the Green
catalogue of SNRs111http://www.mrao.cam.ac.uk/surveys/snrs/ (Green, 2014) for
positions and sizes and (ii) the catalogue from the University of
Manitoba222http://www.physics.umanitoba.ca/snr/SNRcat (Ferrand Safi-Harb 2012
and references therein), compiled mainly from high-energy observations, for
distances and ages.

The project was started in 2016 and the first target selection was done before
Gaia data release (DR) 2 (Gaia Collaboration ., 2018) was available.
Therefore, Gaia DR1 (Gaia Collaboration, Brown ., 2016) was used for these
targets. The much larger number of stars in Gaia DR2 had consequences for the
SNR selection, which are described in the following.

For observations of late-type stars, it was necessary to set a magnitude limit
and, hence, a distance limit for the SNRs. While the limiting magnitude of a
spectrograph at an 8 m-class telescope is $m_{V,\text{max}}\approx 19.5$ mag
(see UVES/VLT manual), for our selection we use $m_{G}\leq 17.0$ mag in order
to reach a sufficient signal-to-noise ratio, $S/N\gtrsim 30$, in the spectra
(see Section 3). Using this limit, mass-brightness relations for the main-
sequence from Henry McCarthy (1993), and choosing a limiting distance of
$d<500$ pc, where four SNRs are available, we get a mass limit of
$M=0.60\,M_{\odot}$, corresponding to spectral type K8. These limits were
applied for stars selected from Gaia DR2. For stars selected earlier from Gaia
DR1, Tycho-Gaia Astrometric Solution (TGAS), we chose $d<1600$ pc, including
up to 25 SNRs. With this limit we get down to $M=0.86\,M_{\odot}$,
corresponding to spectral type K0.5. The limits were chosen as a compromise
between investigating a high number of SNRs and getting down to the latest
possible spectral types. Due to the higher number of stars in Gaia DR2, we
limited the distance to $d<500$ pc in order to create feasible projects for
follow-up observations.

Analysing historical SNRs is particularly interesting, because their age is
accurately known, so we can directly study the SNR expansion and if we find a
runaway star, the determination of the distance and pre-SN binary properties
can be done much more accurately. Therefore, we also added four historical
SNRs to our sample, despite their larger distances. For the most distant one,
Cassiopeia A at $d\approx 3.5$ kpc, with the limits described above we could
detect stars down to spectral type $\sim$F9. The limiting spectral types of
the other SNRs, as well as further properties, are given in Table 1, where we
list our selection of the four SNRs with $d\leq 500$ pc, four further SNRs
with $500\text{\,pc}\leq d\leq 1600$ pc and the four historical SNRs
Cassiopeia A, 3C58 (SN 1181), Crab Nebula (SN 1054) and SNR G347.3$-$00.5 (SN
393).

Table 1: Properties of the SNRs studied in this work. The equatorial
coordinates RA/DEC and angular diameters $\Theta$ are taken from Green (2014),
the distances $d$ and ages $t$ from Ferrand Safi-Harb (2012) and references
therein, where the errors come from the range between the lowest and highest
values in the given literature. For 3C58, Crab Nebula and SNR G347.3$-$00.5
the ages are accurately known (assuming that the associations to the
historical SNe are correct). The last column gives the limiting stellar
spectral type SpT${}_{\text{lim}}$ that could be observed at the corresponding
nominal distance and a limiting magnitude of $m_{G}=17.0$ mag, as descibed
above.

SNR Name | Alternative Name | RA [h:m:s] | DEC [d:m] |  $\Theta$ [arcmin] |  $d$ [kpc] |  $t$ [kyr] | SpT${}_{\text{lim}}$
---|---|---|---|---|---|---|---
G074.0$-$08.5 | Cygnus Loop | 20:51:00 |  $+$30:40 | $230\times 160$ | $0.79\pm 0.21$ | $15\pm 5$ | K5
G160.9$+$02.6 | HB9 | 05:01:00 |  $+$46:40 | $140\times 120$ | $0.80\pm 0.40$ | $5.5\pm 1.5$ | K5
G180.0$-$01.7 | S147 | 05:39:00 |  $+$27:50 | $180\times 180$ | $1.30\pm 0.20$ | $30\pm 4$ | K1.5
G205.5$+$00.5 | Monoceros Loop | 06:39:00 |  $+$06:30 | $220\times 220$ | $1.44\pm 0.54$ | $90\pm 60$ | K1.5
G260.4$-$03.4 | Puppis A | 08:22:10 |  $-$43:00 | $60\times 50$ | $1.3\pm 0.3$ | $4.08\pm 0.38$ | K1.5
G263.9$-$03.3 | Vela | 08:34:00 |  $-$45:50 | $255\times 255$ | $0.275\pm 0.025$ | $18\pm 9$ | M0.5
G266.2$-$01.2 | Vela Jr. | 08:52:00 |  $-$46:20 | $120\times 120$ | $0.75\pm 0.25$ | $3.8\pm 1.4$ | K5
G330.0$+$15.0 | Lupus Loop | 15:10:00 |  $-$40:00 | $180\times 180$ | $0.33\pm 0.18$ | $23\pm 8$ | M0.5
G111.7$-$02.1∗ | Cassiopeia A | 23:23:26 |  $+$58:48 | $5\times 5$ | $3.50\pm 0.20$ | $0.334\pm 0.018$ | F9
G130.7$+$03.1∗ | 3C58 (SN 1181) | 02:05:41 |  $+$64:49 | $9\times 5$ | $2.60\pm 0.60$ | 0.839 | G2
G184.6$-$05.8∗ | Crab Nebula | 05:34:31 |  $+$22:01 | $7\times 5$ | $1.85\pm 0.35$ | 0.966 | G9
G347.3$-$00.5∗ | SN 393 | 17:13:50 |  $-$39:45 | $65\times 55$ | $1.3\pm 0.4$ | 1.627 | K1.5

  * •

∗ Historical SNR.

###  2.2 Selection of runaway star candidates

In order to define an area for the search for runaway star candidates, it is
important to know the position of the SN, which should be close to the
geometric centre (GC) of the SNR. In most cases, we chose the coordinates from
Green (2014), which are mainly based on radio observations, where the SNR
morphology usually is best visible. In the case of the Vela SNR, we combined
the centre from Green (2014) with two images from different wavelengths,
namely the 843 MHz radio image from Bock . (1998) and the X-ray image from
Sushch . (2011). The average coordinates of the three solutions were taken as
the actual GC (see Fig. 1)333Note that in the case of the Vela SNR we did not
trace back the stellar trajectories to the GC, but to the past position of the
Vela pulsar..

Figure 1: The central $$$\times$$$ of the Vela SNR. The crosses mark the three
geometric centres from X-rays (magenta, Sushch . 2011), radio (blue, Bock .
1998) and the Green catalogue (yellow, Green 2014). Their average and standard
deviation are shown by a red cross and ellipse, respectively. Background image
from ESO DSS-2-red.

Around the centre, we defined a search radius, depending on the maximum
possible velocity of runaway stars and the age of the SNR. Tauris (2015) gives
a maximum ejection velocity of $v_{\text{max}}=1050\,\text{km\,s}^{-1}$ for
stars with $0.9\,M_{\odot}$, corresponding to extreme cases which are very
rare (Rz19). We use $v_{\text{max}}=1000\,\text{km\,s}^{-1}$, which yields a
search radius of

$r_{\text{search}}=3.4934\times\frac{t\,\text{[yr]}}{d\,\text{[pc]}}\,\text{arcmin},$
(1)

with the age $t$ and the distance $d$ of the SNR. As SNR expansion velocities
are well above typical runaway star velocities during the Sedov-Taylor phase,
which is the current state of most of our SNRs, we expect potential runaway
stars to still be located inside the SNR.

For each SNR, we first selected all Gaia stars in the given search radius with
$m_{G}=G\leq 17$ mag. We then traced back the projected trajectories of the
stars, using their proper motions and the age of the SNR to obtain their
coordinates at the time of the SN. We neglected the Galactic gravitational
potential because the lifetime of a SNR is so short that the potential will
not have a significant effect. Gaussian error propagation was used to
calculate the uncertainties of the past positions. We selected the stars that,
at the time of the SN, were located inside the error ellipse of the GC. For
stars from Gaia DR1, we used the standard deviations in right ascension and
declination from the determination of the Vela GC to define the error ellipse,
which was then also used for the other SNRs, scaled to the corresponding
diameter as given in Green (2014). For stars from Gaia DR2, we used the error
estimate given by Green (2009), which we translate here to

$\displaystyle\Delta\text{GC}$ $\displaystyle=$
$\displaystyle$$\times$$,\text{~{}if~{}}$$\leq\Theta\leq$$$
$\displaystyle\Delta\text{GC}$ $\displaystyle=$
$\displaystyle$$\times$$,\text{~{}if~{}}$$\leq\Theta\leq$$$ (2)
$\displaystyle\Delta\text{GC}$ $\displaystyle=$
$\displaystyle$$\times$$,\text{~{}if~{}}$$\leq\Theta$

where $\Theta$ is the angular diameter of the SNR. However, if a neutron star
is associated with the SNR, we selected only the stars that could be traced
back to a possible common origin with the neutron star. This was the case for
S147, Vela, 3C58 and the Crab Nebula. Questionable cases or neutron stars with
unknown proper motions were not considered.

Note that the strict limitation of Eqn. 2 means that we loose up to $\sim$33 %
of runaway star candidates that were located more than $1\,\sigma$ from the
nominal GC. We decided to make this limitation in order to create feasible
observing projects, concentrating on the most promising candidates.

Figure 2: HRD of the runaway star candidates with effective temperatures and
luminosities taken from Gaia DR2. In the key, for stars that were observed by
us, we give the names of the corresponding spectrographs, UVES and HDS. For
comparison, we show isochrones calculated from PARSEC models (Bressan .,
2012). The stars marked in red, dark-red and orange were considered to be
post-MS giants and ruled out from the investigation.

We also checked if the Gaia parallaxes of the stars are consistent with the
range of possible distances of the SNR. We compared these to distances derived
by Bailer-Jones . (2018) and found no significant deviations for the parallax
range considered in our work. All candidates are still within the correct
range when using their distances from Bailer-Jones . (2018). For stars from
Gaia DR2, we also checked the given values for radius, luminosity, effective
temperature and surface gravity, so that giants could be excluded beforehand.
The selection could be made by plotting the candidates into a Hertzsprung-
Russell diagram (HRD), shown in Fig. 2, where we compare our candidates to
PARSEC isochrones444PAdova and TRieste Stellar Evolution Code, see
http://stev.oapd.inaf.it/cgi-bin/cmd_3.3 (Bressan ., 2012). Reddish colours
indicate if a star was classified as a giant. If a star has an increased
luminosity compared to the zero-age main-sequence (ZAMS), it can be either a
pre-MS star or an evolved star (i.e. terminal-age MS or giant). As we are
looking for young stars, it might be better to leave a few evolved stars in
the sample than excluding too many pre-MS stars. The division was roughly made
at the isochrone for 1 Myr. For questionable cases and because luminosities
were not given for all possible targets in Gaia DR2, we also checked the
surface gravities $\log(g)$ of the targets in the StarHorse catalogue (Anders
., 2019). Four targets with $\log(g)<3$, where no luminosity was given, were
excluded.

Stars that fulfill all criteria were considered as good runaway star
candidates and selected as targets for the spectroscopic follow-up
observations. In Fig. 3 the selection is shown for the Monoceros Loop. Among
the 21 stars found here from Gaia DR1, we only show the remaining runaway star
candidates and two stars that show Li in their spectra. No neutron star is
known in the Monoceros Loop, so the GC was used as a reference position.

After the detection and characterisation of runaway star candidates, the goal
was to determine the radial velocity (RV) and the atmospheric parameters and
to search for the Li 6708 Å line as youth indicator in the spectra of late-
type stars.

Figure 3: The runaway star candidates in the Monoceros Loop. The image shows
the central $$$\times$$$ of the SNR. The geometric centre of the SNR is marked
by a red cross with corresponding $1\sigma$ error ellipse. The current
positions of runaway star candidates with and without detected lithium are
marked by yellow and blue crosses, respectively. They are connected with the
positions at the time of the SN (surrounded by a corresponding error ellipse)
by arrows in the direction of their proper motion. The star marked in black is
HD 261393, which was suggested as a runaway star by Bo17 and also observed by
us. The background image was taken by us in H$\alpha$ with the Schmidt-
Teleskop-Kamera (Mugrauer Berthold, 2010) at the University Observatory Jena.

##  3 Observations

###  3.1 VLT UVES Spectroscopy

For southern targets, we have used the UV-Visual Echelle Spectrograph (UVES),
which is mounted at the Nasmyth focus of Kueyen (UT2) at the Very Large
Telescope (VLT), run by the European Southern Observatory (ESO) and located on
the Cerro Paranal in Chile. We observed 33 runaway star candidates, selected
from Gaia DR1 TGAS, in five SNRs on the southern sky, namely (i) G180.0$-$01.7
(S147), (ii) G205.5+00.5 (Monoceros Loop), (iii) G260.4$-$03.4 (Puppis A),
(iv) G263.9$-$03.3 (Vela) and (v) G266.2$-$01.2 (Vela Jr.).

We used the Red arm (Ra), cross-disperser 3, with a central wavelength of 600
nm. In this setup, two charge-coupled device (CCD) detectors are used which
cover the spectral ranges $4986-5957$ Å and $6036-7003$ Å. Using $1\times 1$
binning, which is adequate for the relatively bright targets observed in P100,
a slit width of $None$ yields a resolving power of $R=32250$.

For the used magnitude limit of $G=17$ mag, corresponding to $V=17.13$
mag555For the calculation, $G=17.00$ mag was converted to $V=17.13$ mag by
using the polynomial relations by Evans . (2018) with $BP-RP=0.782$ given in
Pecaut Mamajek (2013, Table 5) for a G0 V star., with the given setup we reach
$S/N=33$ at 6711 Å in one hour integration time, derived with the UVES
exposure time calculator (ETC), version
P106.2666https://www.eso.org/observing/etc/bin/gen/form?INS.NAME$=$UVES$+$
INS.MODE$=$spectro, for a G0 V star at airmass 1.1, with a seeing of $None$
and with $2\times 2$ binning. A G0 V star was chosen for this example
calculation because it represents a typical target of our selection.

The VLT observations were executed between October and December 2017 in
service mode. In the appendix, Table 8, we list the 33 runaway star candidates
observed with UVES. The individual exposure times varied between 10 s and 500
s, depending on the brightness of the star. Two exposures were taken for each
star, where an average $S/N$ of 64.7 was achieved for the single exposures.
The total on-source integration time was 5 hours and 4 minutes.

Standard calibrations were used, i.e. Bias frames, Flatfield images of a
Halogen lamp and wavelength calibration images of a Thorium-Argon (ThAr) lamp
which are taken regularly for each standard setup as described in the UVES
calibration plan. Data reduction was done with the EsoReflex pipeline for
UVES. The individual reduction steps are bias subtraction, order detection,
flatfielding, wavelength calibration and spectrum extraction. We obtained four
spectra for each star; for each of the two exposures, we got one spectrum for
the lower and one for the upper wavelength regime.

The two exposures for each star were averaged to one spectrum with IRAF.
Normalisation was done with iSpec (Blanco-Cuaresma ., 2014).

###  3.2 Subaru HDS Spectroscopy

Six northern targets, all located in SNR G160.9+02.6 (HB9), were observed with
the High Dispersion Spectrograph (HDS), mounted at the Nasmyth focus of the
Subaru telescope, which is run by the National Astronomical Observatory of
Japan (NAOJ) and located on the Mauna Kea, Hawaii.

HDS uses two CCD detectors, where with the Ra setup the first one covers the
wavelength range $5062-6446$ Å and the second one $6509-7890$ Å. We used
$2\times 2$ binning and a slit width of $None$, yielding a resolving power of
$R=36000$.

For this setup, with the HDS ETC777https://www.naoj.org/cgi-bin/hds_etc.cgi,
we obtain $S/N\approx 24$ at 6697 Å for a G0 V star with $V=17.13$ mag at
airmass 1.1, with a seeing of $None$ and an exposure time of one hour.

At the bottom of Table 8 we list the six runaway star candidates that were
observed with Subaru in period S18B. All observations were done on 2018 Oct
26. The individual exposure times varied between 40 s and 300 s. Two exposures
were taken for each target and $S/N=105.7$ was reached on average for the
single exposures.

Standard calibrations were used and data reduction was done manually with
IRAF. We corrected for the bias level and bad pixels and applied order
tracing, flatfielding, spectrum extraction and wavelength calibration,
following the procedure described in Aoki Helminiak (2014). As described for
UVES/VLT, also here we obtained four spectra for each star. The last step of
the data reduction already included the normalisation of the spectra by
fitting a spline to the continuum, before the spectral orders were merged.

##  4 Spectral Analysis

Figure 4: Normalised, RV corrected and shifted UVES/VLT spectra between
$\sim$6700 Å and 6730 Å. The bottom six spectra show our Li-rich targets,
which have spectral types between F2 and G2, in particular (1) TYC159-251-1,
(2) TYC159-343-1, (3) TYC8150-2802-1, (4) TYC8150-3105-1, (5) TYC8152-1456-1
and (6) TYC8152-550-1. The yellow spectrum (0) shows TYC159-1896-1, a mid-F-
type star without significant Li absorption, for comparison. The varying
strengths of the Li 6707.8 Å line can be seen by comparison to the strong,
nearby lines: Fe I at 6704.5 Å, 6705.1 Å, 6713.1 Å, 6713.8 Å and 6726.7 Å, Ca
I at 6717.7 Å and Si I at 6722.0 Å. The individual spectral types and Li
equivalent widths are displayed in Table 4.

In the averaged spectra, we searched for the Li 6708 Å absorption line to
measure the equivalent width (see Section 4.4). Li was found in ten of our VLT
spectra. Four of these stars were found to be giants (see Fig. 2) and are
therefore too old to be runaway stars. Fig. 4 shows seven of our VLT spectra
around the Li line. While TYC159-1896-1 is only shown for comparison, six of
them show Li absorption in their spectra, namely (1) TYC159-251-1 and (2)
TYC159-343-1, located in the Monoceros Loop, (3) TYC8150-2802-1 and (4)
TYC8150-3105-1 (Vela) as well as (5) TYC8152-1456-1 and (6) TYC8152-550-1
(Vela Jr.). These six stars were analysed further with iSpec (Blanco-Cuaresma
., 2014) to determine their RVs and atmospheric parameters.

###  4.1 A new spectroscopic binary

The spectrum of TYC159-251-1 shows additional absorption lines, each
redshifted by $\sim$1.3 Å compared to the main absorption lines. Therefore,
the star can be identified as a double-lined spectroscopic binary (SB2). The
system is not yet listed in the 9th Catalogue of Spectroscopic Binary Orbits
(SB9, Pourbaix . 2009). In the following, we will use the designation
TYC159-251-1 A for the primary (brighter) component, TYC159-251-1 B for the
secondary component and TYC159-251-1 for the whole system. As we would get
misleading results from analysing the observed spectrum which contains light
from both stars, we disentangled the two components by subtracting a G0
template spectrum from the observed spectrum to obtain the first component and
a F6 template to obtain the second component (both templates taken from
Bagnulo . 2003)888The spectral types F6 and G0 were found to yield the best
representation of the observed spectrum after iterating with different
combinations of template spectra from Bagnulo . (2003) between F6 and G0
(F6+F9, F6+G0, F7+F9, F8+F8, F8+G0 and F9+F9. Other combinations were excluded
beforehand).. The continuum levels of the templates were normalised before to
have the flux level of the combined spectrum of TYC159-251-1, multiplied by
the factors 0.661 and 0.339, respectively, for the two components,
corresponding to the luminosities of the used spectral types according to
Pecaut Mamajek (2013, Table 5, hereafter PM13). Adding up these modified
templates yielded the best representation of the observed spectrum. Fig. 5
shows the observed spectrum (blue) as well as the individual components
(magenta, green) after subtracting the template spectra (black, yellow). The
residuals (light-blue) between the observed spectrum and the added templates
(red) only show a significant deviation from zero at the Li lines, which is
due to the manual removal of Li in the templates (to retain it in the
component spectra). In the following subsections, potential peculiarities in
the analysis of this SB2 will be described in additional paragraphs.

Figure 5: The spectrum of TYC159-251-1 (blue, 1) and its two components
TYC159-251-1 A (purple, 1A) and TYC159-251-1 B (green, 1B) between $\sim$6700
Å and 6730 Å. The first component was used for RV correction. For comparison
we also show the F6 template spectrum (black), scaled by the factor 0.661, and
the G0 template (yellow), scaled by the factor 0.339, which were subtracted
from the observed spectrum to obtain the two components. The red spectrum is
the addition of the two template spectra, while the residuals between the
observed spectrum and the added templates are shown in light-blue. Lithium was
manually removed from the template spectra.

###  4.2 Radial and space velocities

The barycentric correction was done with iSpec before comparing the spectral
lines to an atomic line mask (cross-correlation) to determine the RVs and
shift the spectra correspondingly. For the cross-correlation, a line list from
a Narval solar spectrum ($370-1048$ nm) was used. For the early-type stars,
the RV determination with iSpec was not possible due to the fewer and broader
lines. We determined their RVs manually with IRAF from the positions of the
following lines, if available: Fe II 5018.4 Å, Mg I 5172.7 Å, Mg I 5183.6 Å,
He I 5875.6 Å, Si II 6347.1 Å, Si II 6371.4 Å, Fe II 6456.4 Å, H$\alpha$
6562.8 Å. The central wavelength for each line was determined as the average
centre from three Gaussian fits, whereupon the $v_{r}$ for each line were
calculated. The $v_{r}$ in Table 2 correspond to average and standard
deviation of these measurements. By combining the RVs with proper motions and
parallaxes from Gaia DR2, we calculated the heliocentric space velocities of
the stars. For the calculation of the peculiar velocity with respect to the
local standard of rest, we first calculated the Galactic space velocity
components $(U,V,W)$, following the equations given in Johnson Soderblom
(1987). Then we corrected for the solar motion with respect to the local
standard of rest by adding
$(U_{\odot},V_{\odot},W_{\odot})=(8.5,13.38,6.49)\,\text{km\,s}^{-1}$
(Coskunoǧlu ., 2011). The peculiar velocity is the absolute value of the
resulting vector. The kinematic data of the six Li-rich stars and the best
early-type candidates are given in Table 2. TYC8150-2802-1, TYC8152-1456-1,
TYC8152-550-1, TYC3344-235-1, TYC159-2771-1 and HD 76060 have
$v_{\text{pec}}>25\,\text{km\,s}^{-1}$ and could therefore be classified as
runaway stars according to Tetzlaff (2013).

For the SB2 system TYC159-251-1, we determined the momentary RVs of the two
components with IRAF, measuring the positions of 38 absorption lines of Fe I,
Fe II, Ca I, Ni I and Si II. As line list we chose the Identification List of
Lines in Stellar Spectra (ILLSS, Coluzzi 1999). To obtain the systemic RV
$\gamma$, we need the mass ratio $q=M_{2}/M_{1}<1$. In Section 4.3 we will
find the components to have spectral types F6 $-$ F7 and F9 $-$ G0,
respectively, corresponding to $q=0.91^{+0.03}_{-0.05}$. The systemic RV can
be calculated as

$\gamma=v_{r_{1}}+\frac{q}{1+q}\times(v_{r_{2}}-v_{r_{1}})$ (3)

The following results were obtained:

$\begin{split}v_{r_{1}}&=-10.7\pm 0.6\,\text{km\,s}^{-1}\\\ v_{r_{2}}&=48.4\pm
0.7\,\text{km\,s}^{-1}\\\ \gamma&=17.5^{+0.9}_{-1.1}\,\text{km\,s}^{-1}\\\
\end{split}$ (4)

The systemic RV is then used to determine the space motion (see
$v_{\text{pec}}$ in Table 2). Extensive follow-up RV monitoring is required to
fit a RV curve and determine the orbital parameters. Only then the systemic
and space velocity of the system can be given reliably.

Table 2: Gaia DR2 $G$ magnitudes, parallaxes $\pi$ and kinematic parameters of the six Li-rich stars as well as the remaining early-type candidates in HB9, Monoceros Loop and Vela Jr. Parallaxes and proper motions are from Gaia DR2, the barycentric corrections ($BC$) were determined with iSpec. The radial velocities $v_{r}$ of the Li-rich stars were determined with iSpec, the $v_{r}$ of the early-type stars were determined with IRAF. For TYC159-251-1, we give the systemic RV $\gamma$ of the binary, approximated using a mass ratio of $q=0.91^{+0.03}_{-0.05}$. The individual RVs of the binary components were derived with IRAF. For the spectral types, see Table 4. Name | $G$ | $\pi$ | $\mu_{RA}$ | $\mu_{DEC}$ | $BC$ | $v_{r}$ | $v_{\text{pec}}$
---|---|---|---|---|---|---|---
| [mag] | [mas] | [mas yr-1] | [mas yr-1] | [km s-1] | [km s-1] | [km s-1]
Li-rich stars |  |  |  |  |  |  |
TYC159-251-1 | 11.51 | $1.774\pm 0.041$ | $+3.63\pm 0.07$ | $-9.44\pm 0.06$ | 5.98 | $+17.5^{+0.9}_{-1.1}$ | $19.3^{+1.1}_{-1.2}$
TYC159-343-1 | 10.87 | $1.383\pm 0.042$ | $+2.51\pm 0.09$ | $-2.76\pm 0.08$ | 6.03 | $+22.27\pm 0.41$ | $13.8\pm 0.7$
TYC8150-2802-1 | 12.03 | $3.226\pm 0.028$ | $-16.94\pm 0.05$ | $+29.17\pm 0.05$ | 15.14 | $+53.84\pm 0.07$ | $58.8\pm 0.6$
TYC8150-3105-1 | 11.94 | $3.351\pm 0.028$ | $-16.00\pm 0.05$ | $+10.05\pm 0.04$ | 15.01 | $+20.37\pm 0.29$ | $19.7\pm 0.5$
TYC8152-1456-1 | 12.47 | $1.661\pm 0.037$ | $+5.17\pm 0.06$ | $-11.56\pm 0.07$ | 15.13 | $-10.59\pm 0.08$ | $49.1\pm 1.0$
TYC8152-550-1 | 11.92 | $2.435\pm 0.026$ | $+4.95\pm 0.04$ | $-10.71\pm 0.05$ | 15.14 | $+36.12\pm 0.10$ | $37.1\pm 0.5$
HB9 |  |  |  |  |  |  |
TYC3344-235-1 | 11.07 | $0.91\pm 0.04$ | $-0.23\pm 0.09$ | $+1.37\pm 0.06$ | 19.15 | $-31.2\pm 2.8$ | $36.0\pm 2.9$
TYC3344-679-1 | 11.40 | $1.01\pm 0.06$ | $+0.74\pm 0.18$ | $-4.76\pm 0.11$ | 19.18 |  $+8.6\pm 7.3$ |  $8.2\pm 5.9$
TYC3344-683-1 | 12.33 | $1.21\pm 0.04$ | $-1.09\pm 0.07$ | $-7.02\pm 0.05$ | 19.19 |  $+9.3\pm 7.3$ | $15.3\pm 3.6$
TYC3344-553-1 | 10.87 | $1.53\pm 0.04$ | $-1.98\pm 0.08$ | $-7.22\pm 0.06$ | 19.18 | $-14.9\pm 7.3$ | $21.9\pm 5.7$
Monoceros Loop |  |  |  |  |  |  |
TYC159-2771-1 | 12.01 | $0.75\pm 0.04$ | $+1.46\pm 0.07$ | $-3.78\pm 0.07$ | 19.53 | $+31.8\pm 2.0$ | $25.3\pm 2.0$
HD261359 | 11.79 | $0.65\pm 0.05$ | $-0.48\pm 0.08$ | $-3.15\pm 0.07$ | 22.59 |  $+0.9\pm 5.1$ | $17.8\pm 4.5$
HD261393 | 10.05 | $0.79\pm 0.05$ | $-0.20\pm 0.09$ | $-0.97\pm 0.07$ | 24.55 | $+31.3\pm 5.8$ | $18.5\pm 5.3$
TYC159-2337-1 | 11.80 | $0.49\pm 0.05$ | $-1.70\pm 0.09$ | $-0.45\pm 0.08$ | 5.94 | $+22.1\pm 4.7$ | $18.0\pm 3.7$
TYC159-2671-1 | 12.25 | $0.69\pm 0.04$ | $+0.08\pm 0.07$ | $-0.78\pm 0.06$ | 19.57 | $+23.3\pm 7.3$ | $11.4\pm 6.3$
Vela Jr. |  |  |  |  |  |  |
HD 76060 | 7.85 | $2.06\pm 0.04$ | $-12.42\pm 0.08$ | $+10.31\pm 0.09$ | 9.38 | $+25.7\pm 7.3$ | $31.2\pm 2.7$

###  4.3 Atmospheric parameters

With the synthetic spectral fitting technique provided by iSpec, we determined
the atmospheric parameters of the six Li-rich stars. The routine computes
synthetic spectra and compares them to the observed spectra by nonlinear
least-squares fitting, minimising the $\chi^{2}$. We used the continuum-
normalised and RV-corrected spectra. The routine considers the following
parameters: (i) Effective temperature $T_{\text{eff}}$, (ii) Surface gravity
$\log(g)$, (iii) Metallicity [M/H], (iv) Microturbulence velocity
$v_{t}^{\text{mic}}$, (v) Macroturbulence velocity $v_{t}^{\text{mac}}$, (vi)
Rotation $v_{\text{rot}}\sin(i)$ and (vii) Limb darkening coefficient, taking
into account the resolving power $R=32250$.

To create the synthetic spectra, we used the radiative transfer code SPECTRUM
with the ATLAS9.Kurucz model atmosphere. Solar abundances were taken from
Kurucz and as line list we chose GESv5_atom_hfs_iso.420_92. For the fitting,
we chose the red part of the spectrum ($6036-7003$ Å), because it yields a
higher $S/N$. We chose a large number (> 100) of temperature-sensitive Fe I,
Fe II, Ca I, Ca II, Ni I, Si I and Si II lines as well as the wings of the
H$\alpha$ absorption line, which has a big effect on the results because it
spans a large wavelength range. After the automatic selection of the lines it
had to be checked if they were suitable for fitting, i.e. they should reach
back to the continuum within the segments that were drawn around each line. In
some cases the segments were modified in order to reach that. In Fig. 6, we
show examples of selected lines for TYC8152-550-1. The synthetic spectra are
created in the segments (grey) while the differences to the observed spectrum
are only computed within the line regions (yellow). If no continuum regions
were available around a line, or if strong deviations between spectrum and fit
were recognised within the line mask, the corresponding line was rejected.

Figure 6: Different example lines of TYC8152-550-1 (blue) that were used to
fit the atmospheric parameters (red). The yellow areas mark the line regions
used for the fit, the grey areas show segments of 0.5 Å on each side of each
line used for synthesising the spectra. In particular we show Fe I 6065, Fe II
6084, Ni I 6086, Fe I 6344, Si II 6347, Ca I 6472 and Si I 6527 Å.

As initial parameters for the fits, we used solar parameters as given in
Blanco-Cuaresma (2019) and the $T_{\text{eff}}$ of each star from Gaia DR2. In
each run, consisting of up to twelve iterations, it could be decided about
which parameters should be fixed or free. As all parameters were unknown, we
left all of them free, except either macroturbulence or
$v_{\text{rot}}\sin(i)$, which are strongly correlated with each other. By
running the fit several times with different combinations, we guaranteed that
the fixed parameter among $v_{t}^{\text{mac}}$ and $v_{\text{rot}}\sin(i)$ was
set to a reasonable value, while the other was left free. The limb darkening
coefficient was fixed to 0.6, because changing it did not show any effect.

The resulting atmospheric parameters of the six Li-rich stars are displayed in
Table 3. The given intrinsic errors are computed from the covariance matrix
which connects the errors of the individual parameters after the least-squares
fitting (Blanco-Cuaresma ., 2014). The flux errors for our spectra were
calculated from gain and read noise of UVES with the given setup. For
comparison, we also show $T_{\text{eff}}$ from Gaia DR2. We give the results
for both components of TYC159-251-1, where the fits were done individually
after disentangling the spectra. Note that the disentangled spectra have lower
$S/N$. The resulting $T_{\text{eff}}$ for TYC159-251-1 B corresponds to SpT
F9, but is also consistent with G0, which was used to disentangle the original
spectrum. The derived $T_{\text{eff}}$ for TYC159-251-1 A is fully consistent
with the assumption of an F6 template spectrum. $v_{t}^{\text{mac}}=0\,$km s-1
and $v_{\text{rot}}\sin(i)=0\,$km s-1 were obtained for both components in all
iterations, even when one or both of these parameters were left free.

Table 3: The atmospheric parameters of the six Li-rich stars determined with
iSpec. The effective temperatures from Gaia DR2 are also shown for comparison.
A missing error indicates that the parameter was kept constant for all
iterations. The results for TYC159-251-1 are given for both components
individually.

Name | $T_{\text{eff},Gaia}$ | $T_{\text{eff}}$ | $\log(g/\text{cm\,s}^{-2})$ | [M/H] | $v_{t}^{\text{mic}}$ | $v_{t}^{\text{mac}}$ | $v_{\text{rot}}\sin(i)$
---|---|---|---|---|---|---|---
| [K] | [K] |  |  | [km s-1] | [km s-1] | [km s-1]
TYC159-251-1 A |  $6257\pm 391^{*}$ | $6308\pm 59$ | $4.16\pm 0.22$ | $-0.19\pm 0.06$ | $1.49\pm 0.24$ | 0 | 0
TYC159-251-1 B |  $6257\pm 391^{*}$ | $6032\pm 73$ | $3.94\pm 0.27$ | $-0.16\pm 0.07$ | $1.86\pm 0.27$ | 0 | 0
TYC159-343-1 | $7079\pm 232$ | $6749\pm 69$ | $2.77\pm 0.35$ | $-0.57\pm 0.08$ | $4.9\pm 0.8$ | $16.4\pm 2.0$ | 1.5
TYC8150-2802-1 | $5853\pm 100$ | $5857\pm 75$ | $3.99\pm 0.20$ | $-0.08\pm 0.06$ | $1.28\pm 0.19$ | 0 | 0
TYC8150-3105-1 | $5614\pm 228$ | $5956\pm 73$ | $4.78\pm 0.21$ | $-0.19\pm 0.09$ | $1.6\pm 0.5$ | $4.8$ | $13.0\pm 0.9$
TYC8152-1456-1 | $5806\pm 77$ | $5955\pm 73$ | $4.14\pm 0.23$ | $-0.09\pm 0.07$ | $1.26\pm 0.26$ | $0.8\pm 1.9$ | 0
TYC8152-550-1 | $6061\pm 192$ | $6165\pm 63$ | $4.26\pm 0.24$ | $-0.27\pm 0.06$ | $1.50\pm 0.29$ | 0 | 0

  * •

∗ Gaia DR2 value for the unresolved system

###  4.4 Lithium equivalent widths and abundances

The Li equivalent widths $EW_{\text{Li}}$ of the six Li-rich stars were
measured with IRAF by integrating over the area of the Li 6708 Å line with
respect to the local continuum. The errors were calculated by adding the
uncertainty in fitting the continuum and the error related to read-noise
($2.8\,\text{e}^{-}$) and gain ($0.52\,\text{e}^{-}/\text{ADU}$) of the
instrument. The continuum error was derived by varying the flux level of the
local continuum for the integration.

If we do not disentangle the spectrum of a spectroscopic binary, the
equivalent widths will be underestimated due to the additional continuum flux
from the other component. Therefore, we measured the $EW_{\text{Li}}$ for both
components of TYC159-251-1 individually after disentangling. The absorption
line is even stronger for the secondary component. As the F6 and G0 template
spectra also show Li absorption, we removed the lines manually with IRAF
before subtracting the templates from the original spectrum. So it was
guaranteed that the Li lines of both components remained in the disentangled
spectra, where we could then measure the corresponding $EW_{\text{Li}}$.

The $EW_{\text{Li}}$ were converted to abundances by using the curves of
growth given by Soderblom . (1993, Table 2, hereafter So93). The abundance
scale is based on $\log(N_{\text{H}})=12$. The curves of growth are calculated
for $T_{\text{eff}}=$ 4000 $-$ 6500 K and local thermodynamic equilibrium. We
estimated the Li abundances $\log(N_{\text{Li}})$ by choosing the values from
So93 that best represent the given temperature and logarithmic equivalent
width (for $EW_{\text{Li}}$ in mÅ) of the Li-rich stars, taking into account
the errors. With $T_{\text{eff}}=6749\pm 69$ K, TYC159-343-1 is not covered by
the $T_{\text{eff}}$ range in So93. Therefore, here we extrapolated the given
abundances with a quadratic function and determined the Li abundance for
$T_{\text{eff}}=6750$ K and $\log(EW_{\text{Li}})=1.50$,
$\log(EW_{\text{Li}})=1.55$ and $\log(EW_{\text{Li}})=1.60$. The resulting
abundance error, $\Delta\log(N_{\text{Li}})=^{+0.07}_{-0.08}$, contains the
error from varying $EW_{\text{Li}}$ ($\Delta_{\text{var}}=^{+0.045}_{-0.048}$)
and the root mean square (RMS) of the fit ($\Delta_{\text{RMS}}=0.029$ in all
three cases). For the other Li-rich stars, the errors only come from the
variation of $EW_{\text{Li}}$ and/or $T_{\text{eff}}$, according to the
errors. The Li equivalent widths and abundances, together with temperatures,
spectral types and magnitudes are listed in Table 4. The abundances can be
compared to typical initial values, which are expected to be
$\log(N_{\text{Li}})=3.1-3.3$ for population I stars (Sestito Randich, 2005).

From D’Antona Mazzitelli (1984, Table 7) it can be seen that there is no
significant Li depletion in stars with $M\gtrsim 1.2\,M_{\odot}$ if no
additional mixing mechanisms occur. This corresponds to spectral type F7
(PM13). Therefore, the Li signal detected in TYC159-343-1 might not be
conclusive to estimate its age and also TYC159-251-1 and TYC8152-550-1 have to
be taken with care. Nevertheless, we will make an attempt to obtain rough age
estimates for the Li-rich stars in the next subsection.

Table 4: Temperatures, spectral types and magnitudes of the six Li-rich stars and the remaining observed early-type runaway star candidates. Effective Temperatures $T_{\text{eff}}$ of the Li stars and the stars in HB9 are from our parameter fits and from Bai . (2019), respectively. The spectral types of the stars in Monoceros and Vela Jr. were taken from the Skiff catalogue (Vo85, Houk 1978). Conversion between $T_{\text{eff}}$ and SpT was done with PM13. G, BP and RP magnitudes are from Gaia DR2, absolute $G$ magnitudes $M_{G}$ were calculated from Gaia G, extinction and parallax and $G$ band extinctions are from the StarHorse catalogue (Anders ., 2019). The last two columns show Li equivalent widths $EW_{\text{Li}}$ and abundances $\log(N_{\text{Li}})$ for the six Li-rich stars. The abundance errors correspond to the minimum and maximum possible abundance from So93, taking into account the uncertainties in $EW_{\text{Li}}$ and $T_{\text{eff}}$. For TYC159-251-1, we give the values for the unresolved system as well as for the individual components, assuming that TYC159-251-1 A contributes 66.1 % of the flux and TYC159-251-1 B 33.9 %. Name | $T_{\text{eff}}$ | SpT | $G$ | $BP$ | $RP$ | $M_{G}$ | $A_{G}$ | $EW_{\text{Li}}$ | $\log(N_{\text{Li}})$
---|---|---|---|---|---|---|---|---|---
| [K] |  | [mag] | [mag] | [mag] | [mag] | [mag] | [mÅ] |
Li-rich stars |  |  |  |  |  |  |  |  |
TYC159-251-1 | $6351\pm 78$ | F5–7 | 11.51 | 11.78 | 11.09 | $+2.61_{-0.19}^{+0.20}$ | $0.14_{-0.18}^{+0.20}$ | $17\pm 2$ | $2.15^{+0.30}_{-0.05}$
TYC159-251-1 A | $6308\pm 59$ | F6–7 | 11.96 | 12.23 | 11.54 | $+3.06_{-0.19}^{+0.20}$ | $0.14_{-0.18}^{+0.20}$ | $22\pm 3$ | $2.31^{+0.06}_{-0.11}$
TYC159-251-1 B | $6032\pm 73$ | F9–G0 | 12.68 | 12.96 | 12.26 | $+3.78_{-0.19}^{+0.20}$ | $0.14_{-0.18}^{+0.20}$ | $97\pm 7$ | $2.88^{+0}_{-0.08}$
TYC159-343-1 | $6749\pm 69$ | F2–4 | 10.87 | 11.08 | 10.54 | $+1.24_{-0.23}^{+0.29}$ | $0.33_{-0.22}^{+0.28}$ | $35\pm 5$ | $2.91^{+0.07}_{-0.08}$
TYC8150-2802-1 | $5857\pm 75$ | G0–G2 | 12.03 | 12.36 | 11.55 | $+4.53_{-0.13}^{+0.16}$ | $0.04_{-0.13}^{+0.15}$ | $50\pm 3$ | $2.25^{+0.03}_{-0.06}$
TYC8150-3105-1 | $5956\pm 73$ | F9.5–G1 | 11.94 | 12.26 | 11.42 | $+4.34_{-0.16}^{+0.11}$ | $0.22_{-0.16}^{+0.11}$ | $177\pm 7$ | $3.37\pm 0$
TYC8152-1456-1 | $5955\pm 73$ | F9.5–G1 | 12.47 | 12.81 | 11.97 | $+3.35_{-0.26}^{+0.21}$ | $0.22_{-0.25}^{+0.21}$ | $66\pm 4$ | $2.59^{+0.07}_{-0}$
TYC8152-550-1 | $6165\pm 63$ | F7–F9 | 11.92 | 12.21 | 11.48 | $+3.52_{-0.17}^{+0.19}$ | $0.33_{-0.17}^{+0.19}$ | $33\pm 4$ | $2.47^{+0.05}_{-0.26}$
HB9 |  |  |  |  |  |  |  |  |
TYC3344-235-1 | $7707\pm 377$ | A5–F0 | 11.07 | 11.38 | 10.61 | $-0.60_{-0.57}^{+0.21}$ | $1.45_{-0.57}^{+0.19}$ | – | –
TYC3344-679-1 | $7591\pm 435$ | A6–F0 | 11.40 | 11.72 | 10.92 | $-0.28_{-0.51}^{+0.16}$ | $1.69_{-0.49}^{+0.09}$ | – | –
TYC3344-683-1 | $7490\pm 360$ | A7–F0 | 12.33 | 12.58 | 11.94 | $+2.40_{-0.24}^{+0.30}$ | $0.35_{-0.23}^{+0.29}$ | – | –
TYC3344-553-1 | $7588\pm 249$ | A7–A9 | 10.87 | 11.08 | 10.53 | $+1.06_{-0.24}^{+0.46}$ | $0.73_{-0.24}^{+0.45}$ | – | –
Monoceros Loop |  |  |  |  |  |  |  |  |
TYC159-2771-1 | $10700_{-1250}^{+2550}$ | B8–A0 | 12.01 | 12.13 | 11.78 | $+1.11_{-0.25}^{+0.23}$ | $0.29_{-0.22}^{+0.20}$ | – | –
HD 261359 | $10700_{-1250}^{+2550}$ | B8–A0 | 11.79 | 11.87 | 11.62 | $+0.89_{-0.26}^{+0.26}$ | $-0.04_{-0.20}^{+0.20}$ | – | –
HD 261393 | $15700_{-1450}^{+1150}$ | B4–B6 | 10.05 | 10.07 | 10.02 | $-0.61_{-0.24}^{+0.17}$ | $0.16_{-0.20}^{+0.12}$ | – | –
TYC159-2337-1 | $10700_{-1250}^{+2550}$ | B8–A0 | 11.80 | 11.88 | 11.61 | $-0.04_{-0.27}^{+0.27}$ | $0.27_{-0.13}^{+0.12}$ | – | –
TYC159-2671-1 | $12500_{-1950}^{+1750}$ | B7–B9 | 12.25 | 12.31 | 12.08 | $+1.25_{-0.22}^{+0.19}$ | $0.18_{-0.19}^{+0.15}$ | – | –
Vela Jr. |  |  |  |  |  |  |  |  |
HD 76060 | $12500_{-950}^{+1750}$ | B7–B9 | 7.85 | 7.83 | 7.92 | $-0.63_{-0.11}^{+0.11}$ | $0.05_{-0.10}^{+0.10}$ | – | –

###  4.5 Age estimation

In the following, we will compare two methods to estimate the ages of our
targets. Firstly, we will use the positions of the stars in the HRD, where we
can compare them to isochrones (Figs. 7 and 8). Then we perform the Li test by
using the determined $T_{\text{eff}}$ and $EW_{\text{Li}}$ (or
$\log(N_{\text{Li}})$) to compare them to stars of clusters with known ages
(Figs. 9 and 10).

Figure 7: HRD with absolute G-band magnitudes of stars observed with UVES/VLT
and HDS/Subaru. For an age estimate, we show PARSEC isochrones (Bressan .,
2012) with metallicity $Z=0.152$, corresponding to solar metallicity. For the
early-type stars (blue), we used spectral types from the Skiff catalogue
(Skiff 2009 and references therein) and corresponding $T_{\text{eff}}$ ranges
from PM13. The remaining candidates in the Monoceros Loop (M) and Vela Jr.
(VJ) are marked by grey squares and labeled as (M1) TYC159-2771-1, (M2) HD
261359, (M3) HD 261393, (M4) TYC159-2337-1, (M5) TYC159-2671-1, (VJ1) HD
76060. For stars located in HB9 (cyan), labeled as (H1) TYC3344-235-1, (H2)
TYC3344-679-1, (H3) TYC3344-683-1 and (H4) TYC3344-553-1, we checked the
$T_{\text{eff}}$ from Bai . (2019), which are used here for the stars with
spectral type late-A. For stars with $T_{\text{eff}}\lesssim 7500$ K (SpT FG)
and no Li (purple), we use $T_{\text{eff}}$ from Gaia DR2. For the Li-rich
stars, marked with coloured squares, we show the $T_{\text{eff}}$ from our
parameter fits. They are labeled as (1) TYC159-251-1, (2) TYC159-343-1, (3)
TYC8150-2802-1, (4) TYC8150-3105-1, (5) TYC8152-1456-1 and (6) TYC8152-550-1.
TYC159-251-1 is a binary and the HRD positions of the individual components
correspond to the black triangles.

Figure 8: Left: Zoom-in of Fig. 7, centered on the Li-rich stars. Right: HRD
with absolute G-band magnitudes of TYC159-343-1. For an age estimate, we show
PARSEC isochrones (Bressan ., 2012) with metallicity [M/H] $=-0.57$. This
value and the $T_{\text{eff}}$ were determined from our parameter fits. The
axis scales are similar in both panels. For the key, see Fig. 7.

Figure 9: $T_{\text{eff}}$ versus $EW_{\text{Li}}$ for the six Li-rich stars,
shown with coloured error boxes and labeled as (1)${}_{\text{A}}$ TYC159-251-1
A, (1)${}_{\text{B}}$ TYC159-251-1 B, (2) TYC159-343-1, (3) TYC8150-2802-1,
(4) TYC8150-3105-1, (5) TYC8152-1456-1 and (6) TYC8152-550-1. For
TYC159-251-1, we only show the individual components. The curves represent
open clusters with different ages and were fitted by Eric Mamajek to the data
of the individual cluster stars. The original plot can be found under
http://www.pas.rochester.edu/$\sim$emamajek/images/li.jpg.

####  4.5.1 HRD isochrone placement

In the HRD (Fig. 7, for a zoom-in of the Li-rich stars see Fig. 8), we show
the stars observed with UVES/VLT and HDS/Subaru, using $T_{\text{eff}}$ either
from Gaia DR2, from their spectral type as given in the Skiff catalogue
(Skiff, 2009), from Bai . (2019) or from our parameter fits. Using Gaia DR2
parallaxes and StarHorse extinctions (Anders ., 2019), we converted the
apparent magnitudes $G$ to absolute magnitudes $M_{G}$. For the $G$-band
extinctions, no errors are given in the StarHorse catalogue. We converted the
$V$ band extinction errors to $G$ with the relation $\Delta A_{G}=0.77\Delta
A_{V}$ (Mugrauer, 2019). $\Delta M_{G}$ was then determined with Gaussian
error propagation.

We compare our stars to isochrones which were calculated from PARSEC models
(Bressan ., 2012), reflecting the whole stellar evolution until the first
thermal pulse. We used solar metallicity, $Z=0.152$, which corresponds to
$[\text{M/H}]=0.015$. This is close to $[\text{M/H}]=-0.20\pm 0.13$, the
average metallicity of our targets as given in the StarHorse catalogue. The
corresponding shift of isochrone positions is rather small compared to the
temperature and magnitude errors for early- and late-type stars, respectively.
However, TYC159-343-1 has a significantly lower metallicity of
$[\text{M/H}]=-0.57\pm 0.08$. Therefore, for this target we plotted seperate
isochrones to determine its age more precisely (see Fig. 8 and Table 5).

We directly see that the red giants are located on the giant branch of the
isochrones in Fig. 7, consistent with ages between 0.2 Gyr and $>$ 10 Gyr. The
Li-rich stars are mostly located relatively close to the ZAMS, only
TYC159-343-1 is located close to the isochrone for 5 Myr due to its high
brightness.

TYC159-251-1 was discovered by us to be a SB2. Photometrically unresolved
binarity changes the positions in the HRD. The brightness is then
overestimated, so the star is located significantly above the ZAMS. We correct
for this effect by assuming that the primary component contributes 66.1 % and
the secondary 33.9 % of the flux. This adds 0.45 mag and 1.17 mag,
respectively, to their $M_{G}$ magnitude, shifting them back towards the ZAMS.
Furthermore, the revised atmospheric parameter fits for the individual
components yield lower effective temperatures. The two individual components
differ by 276 K and 0.72 mag.

The positions of the early-type stars in the HRD have to be taken with care
because Gaia temperature estimates are trained only up to
$T_{\text{eff}}=10000$ K and tend to be underestimated when approaching this
limit. Therefore, for the early-type stars observed with UVES, we adopted the
spectral types from the Skiff catalogue (Skiff 2009). Only one reference is
given for each star (Houk 1978 for HD76060, Voroshilov . 1985, hereafter Vo85,
for the ten other early-type stars) and no error ranges are given there, so we
assume them to be $\pm 1$ subclass. The corresponding temperatures were
inferred from PM13. Although these stars were also observed by us, the spectra
were not suited for determining the spectral type. Our spectral range of
$5000-7000$ Å is quite narrow and misses most spectral lines that are usually
taken for the analysis of early-type stars. Additionally, the few lines we
found are often affected by strong rotational broadening. Nevertheless, we
checked the Skiff spectral types for consistency by comparing our spectra to
standard star spectra observed with the Fibre Linked ECHelle Astronomical
Spectrograph (FLECHAS, Mugrauer . 2014), operated at the 0.9 m telescope of
the University Observatory Jena. We mainly used the He I line at 5015.7 Å
(visible up to B9) and the Fe lines at $5018.0-5018.4$ Å (visible from A0 on).
Our results are largely consistent with the spectral types given by Vo85, with
deviations of up to four subclasses.

For the stars observed with HDS/Subaru, from comparison with the spectra we
found that the $T_{\text{eff}}$ in Gaia DR2 are underestimated. For these
targets, which are not listed in the Skiff catalogue, instead we used the
values given in Bai . (2019), which fit best to our spectra. Therefore, they
probably have spectral types A5 $-$ F0. However, their absolute magnitudes
differ a lot from each other, indicating different ages.

HD 261393, suggested by Bo17 as the best candidate in the Monoceros Loop, has
$T_{\text{eff}}=14250-16850$ K (Vo85) and $M_{G}\approx-0.6$ mag. Therefore,
it is probably an evolved but still young star of $\lesssim$ 105 Myr,
consistent with being the ejected companion of a SN-progenitor ($t\lesssim 32$
Myr).

Table 5: Estimated HRD ages (see Figs. 7 and 8) for the six Li-rich stars,
including both components of the SB2 TYC159-251-1, as well as for the
remaining observed runaway star candidates in HB9, Monoceros Loop and Vela Jr.
If a star is consistent with the zero-age main-sequence, only a lower limit is
given in column 2. Otherwise, the values in column 2 represent the case that
the star is on its pre-MS and the values in column 3 represent the case that
the star is on its terminal-age- or post-MS.

Name |  $t_{\text{HRD}}$ [Myr] |  $t_{\text{HRD}}$ [Myr]
---|---|---
| if (pre-)MS | if post-MS
Li-rich stars |  |
TYC159-251-1 A | 11 – 16 | 3200 – 4500
TYC159-251-1 B | 14 – 30 | 4800 – 7500
TYC159-343-1* | 2.1 – 3.8 | 1600 – 2300
TYC8150-2802-1 | $>30$ | –
TYC8150-3105-1 | $>25$ | –
TYC8152-1456-1 |  8 – 16 | 5000 – 8000
TYC8152-550-1 | 13 – 20 | 3800 – 6000
HB9 |  |
TYC3344-235-1 | $<1.6$ | 380 – 650
TYC3344-679-1 | $2.0$ | 490 – 800
TYC3344-683-1 | $>9$ | –
TYC3344-553-1 | 3.7 – 7 |  850 – 1300
Monoceros Loop |  |
TYC159-2771-1 | $>4.2$ | –
HD 261359 | $>3.5$ | –
HD 261393 | $<110$ | –
TYC159-2337-1 | 1.6 – 3.3 |  50 – 490
TYC159-2671-1 | $>5$ | –
Vela Jr. |  |
HD 76060 | 1.1 – 1.8 |  90 – 350

  * •

* HRD ages determined with isochrones for $[\text{M/H}]=-0.57$.

####  4.5.2 Lithium test

Fig. 9 shows $T_{\text{eff}}$ and $EW_{\text{Li}}$ for stars in several open
clusters of different
ages999http://www.pas.rochester.edu/$\sim$emamajek/images/li.jpg. Eric Mamajek
fitted polynomials to the data of cluster stars, allowing a rough localisation
of ages in the $T_{\text{eff}}-EW_{\text{Li}}$ space. By comparing the values
of the Li-rich stars, with $T_{\text{eff}}$ from our parameter fits, to these
curves, we obtained the age ranges displayed in Table 6. Further uncertainty
comes from the unknown initial Li abundances of our targets, depending on the
metallicity (Lambert Reddy, 2004). If the initial abundance was lower than in
the case of the cluster stars, less time would have been necessary to reach
the current abundance. Also, the data that were used to fit the polynomials
show a large scatter, e.g. due to different metallicities and rotational
velocities, as well as possible systematic effects like unresolved binarity or
starspots. Therefore, quantifying ages for individual stars is unreliable and
the values given in Table 6 should only be seen as a rough estimate.

Table 6: Ages of the six Li-rich stars, including both components of the SB2 TYC159-251-1, estimated from Fig. 9. Name |  $t_{\text{Li}}$ [Myr]
---|---
TYC159-251-1 A | $>625$
TYC159-251-1 B | 90 – 4000
TYC159-343-1 | 90 – 625
TYC8150-2802-1 | 625 – 4000
TYC8150-3105-1 | 5 – 50
TYC8152-1456-1 | 250 – 4000
TYC8152-550-1 | $>625$

####  4.5.3 Comparison

We summarise the ages inferred from the two diagrams in Tables 5 and 6.
TYC8150-2802-1 and TYC8150-3105-1 are consistent with the ZAMS in Fig. 7. For
TYC8150-2802-1, the low Li content indicates that it is an evolved star, at
least a few hundred Myr old. TYC8150-3105-1, however, shows very strong Li
absorption indicating that it is young. From combining both results, we obtain
an age of $25-50$ Myr, so we suggest that it is on its early MS. In the SB2
system TYC159-251-1, the primary is slightly above the ZAMS while the
secondary is consistent with it. Still, their ages are consistent with each
other. A young age of $\sim$15 Myr can not be excluded from the HRD, but we
consider it more likely that the binary is an evolved system having
$t_{\text{HRD}}\approx 4.5-4.8$ Gyr because their Li content is too small for
a young age.

TYC159-343-1 is a difficult case: From the parameter fits we obtained a very
low metallicity of [M/H] $=-0.57\pm 0.08$, so we inferred its HRD-age from
comparison to low-metallicity isochrones (see Fig. 8). Its high brightness of
$M_{G}=1.24\pm 0.19$ mag indicates an age of only $2.1-3.8$ Myr, if it is pre-
MS. However, it is more likely an evolved star with an age of $1.6-2.3$ Gyr.
The Li abundance, $\log(N_{\text{Li}})=2.91\pm 0.05$, is not a good age tracer
in this case, because of two caveats: (i) The low metallicity also means that
the star probably had a low initial Li abundance, e.g.
$\log(N_{\text{Li}})=2.64\pm 0.07$ according to Lambert Reddy (2004, Table 2),
which is even below the measured abundance. A lower initial abundance reduces
the time which is necessary to reach the current abundance and therefore the
age. (ii) TYC159-343-1 has spectral type F2 $-$ 4, meaning that no significant
Li depletion is expected from standard mixing mechanisms (D’Antona Mazzitelli,
1984). In Fig. 9, the star lies significantly below the 5 Myr curve and the
age can be given as $t_{\text{Li}}=90-625$ Myr. However, due to the caveats
stated above, we rather rely on the age ranges estimated from Fig. 8.

The HRD positions of TYC8152-1456-1 and TYC8152-550-1 also allow both
possibilities, i.e. they could be either pre- or post-ZAMS. Applying the Li
test to them excludes the possibility that they are very young. So, they are
probably evolved stars, unrelated to the birth association of the Vela Jr.
progenitor, whereas TYC8150-3105-1 could be from the same stellar group that
gave birth to the Vela progenitor.

####  4.5.4 The lithium depletion gap

The Hyades, at an age of $\sim$625 Myr, show an interesting feature that
contradicts standard mixing mechanisms: Stars of $T_{\text{eff}}\approx
6600\pm 300$ K show Li depletions which are enhanced by a factor $\sim$100
compared to neighbouring hotter and cooler stars (Boesgaard Tripicco, 1986;
Thorburn ., 1993; Deliyannis ., 2019). In a weaker extent, this feature can
also be seen in younger clusters like M35 ($\sim$175 Myr), where the gap is
just forming (Steinhauer, 2003; Steinhauer Deliyannis, 2004). The gap of
enhanced Li depletion is caused by non-standard mixing on the MS, acting
predominantly in the above stated $T_{\text{eff}}$ regime. Slow mixing, e.g.
by rotational and/or gravitational instabilities, and diffusion can play a
role in the formation of the gap (Steinhauer Deliyannis, 2004). Recent works
(Steinhauer Deliyannis, 2021) show the effect of rotational mixing for the
formation of the gap. The authors find that cluster stars with a high
$v_{\text{rot}}\sin(i)$ have a more strongly decreasing Li abundance with age.
Therefore, observations of stars within the gap and comparing their Li
abundances and $v_{\text{rot}}\sin(i)$ to clusters of known ages can give a
further estimate of their ages. Unfortunately, the method is not applicable
for us due to the insignificant measurements of $v_{\text{rot}}\sin(i)$ and
the unknown initial Li abundances. Additionally, the method only works for
stars within the Li depletion gap and only for ages of $\sim$100 $-$ 650 Myr.
However, by comparing the measured abundances to cluster data we can still
obtain a consistency check of the ages given in Tables 5 and 6. This
comparison also works for cooler stars, while it also suffers from the same
caveats as described for the Li test above, e.g. the large scatter of the
$\log(N_{\text{Li}})$.

In Fig. 10 we compare our Li targets to data of the clusters Pleiades (Butler
. 1987; Pilachowski . 1987; Boesgaard . 1988; So93) and Hyades (Thorburn .,
1993; Boesgaard Tripicco, 1986). The Li depletion gap of the Hyades is
indicated by the brown arrow and the attached horizontal bar at
$T_{\text{eff}}=6300-6900$ K. Note that the hot edge of the gap is steeper
than the cold edge, where a larger scatter is observed. The $T_{\text{eff}}$
of TYC159-251-1 A and TYC159-343-1 are consistent with the Li gap.
TYC159-251-1 A lies at the cold edge of the gap and its Li abundance is lower
than of Hyades stars with the same $T_{\text{eff}}$. Therefore, the Hyades age
could be seen as a lower limit for the SB2 TYC159-251-1. TYC159-343-1 has one
of the highest Li abundances among our sample, despite its low metallicity.
Its Li abundance is comparable to Pleiades stars in this
$T_{\text{eff}}$-region, so the Pleiades age ($\sim$125 Myr) can be seen as an
upper limit. Possibly the star is younger than inferred from Fig. 9 ($90-625$
Myr), where we did not consider the metallicity.

The other four Li-rich targets have lower $T_{\text{eff}}$ and do not fall in
the Li depletion gap. The abundances of TYC8150-2802-1, TYC8152-1456-1 and
TYC8152-550-1 are consistent with the Hyades or older, whereas TYC8150-3105-1
has more Li than comparable Pleiades stars. Therefore, it is very young,
consistent with what we found from Fig. 9.

Figure 10: Lithium abundances of Pleiades and Hyades stars. The data for the
Pleiades (green) were taken from Soderblom . (1993, Table 1), including data
from Boesgaard . (1988); Butler . (1987) and Pilachowski . (1987). The data
for the Hyades (brown) were taken from Thorburn . (1993) and Boesgaard
Tripicco (1986). Inverted triangles show upper limits in the case of non-
detections. Our Li targets, shown with coloured squares representing the error
boxes, are labeled as in Fig. 9. The stars (1)${}_{\text{A}}$ TYC159-251-1 A
and (2) TYC159-343-1 are consistent with the Li depletion gap, which is marked
by the brown arrow and the attached horizontal bar.

##  5 Discussion

###  5.1 Overview of the candidates

Table 7: Numbers of runaway star candidates with $G<17$ mag in the SNRs
covered in this work. Columns 2 and 3 give the numbers of candidates
identified in Gaia DR1 TGAS and DR2, respectively. Column 4 gives the number
of observed stars and if they were observed with UVES/VLT (U) or with
HDS/Subaru (H). These observations are all based on the DR1 selection,
whereupon Gaia DR2 was used for the further characterisation of the targets.
Column 5 gives the names of associated neutron stars and column 6 the numbers
of remaining candidates that could not be excluded during the analysis.

SNR Name | DR1 | DR2 | Obs. | PSR | Rem. cands. | Additional info
---|---|---|---|---|---|---
G074.0$-$08.5 | 16 | 19 | 0 | – | 19 |
G160.9+02.6 | 6 | 1 | 6 (H) | SGR 0501+4516(?) | 5 | SGR & PSR prob. unrelated
G180.0$-$01.7 | 3 | – | 3 (U) | J0538+2817 | 1 | HD 37424 confirmed as runaway1
G205.5+00.5 | 21 | 116∗∗ | 21 (U) | – | 5 | incl. HD 2613932
G260.4$-$03.4 | 2 | 9 | 2 (U) | – | 9 | PSR J0821$-$4300 prob. unrelated
G263.9$-$03.3 | 2 | 1 | 2 (U) | J0835$-$4510 | 1 |
G266.2$-$01.2 | 5 | 12 | 5 (U) | – | 13 | PSR J0855$-$4644 prob. unrelated
G330.0+15.0 | 6 | 3 | 0 | – | 3 |
G111.7$-$02.1∗ | 0 | 0 | 0 | CCO J232327.9+584842 | 0 |
G130.7+03.1∗ | 0 | 0 | 0 | J0205+6449 | 0 | No cand. consistent with PSR
G184.6$-$05.8∗ | 0 | 0 | 0 | J0534+2200 | 0 |
G347.3$-$00.5∗ | 2 | 18 | 0 | CCO(?) | 18 |

  * •

∗ Historical SNR; ∗∗ not suggested for observations, more precise constraints
required; 1 Dinçel . (2015); 2 Boubert . (2017)

We list in Table 7 the SNRs with their corresponding numbers of runaway star
candidates from the different selection steps and if they were already
observed, as well as additional information (e.g. the best candidates). A list
of the individual candidate stars can be found in Table 9.

Note that for the candidates selected from Gaia DR2 the allowed locations
around the SNR centres at the time of the SN were reduced from an error
estimate based on the determination of the Vela SNR centre to a smaller error
based on the description in Green (2009) (see Section 2.2). This was necessary
to limit the high number of stars from DR2 more strictly, in order to create
feasible projects for follow-up observations. In both cases the errors scale
with the size of the SNR. Trajectories of associated pulsars were considered
for the DR2 selection.

In S147, we can confirm the B0.5 V star HD 37424 based on DR2 data. Its
position at the time of the SN is well within the 1 $\sigma$ error ellipse of
the GC used for the DR1 selection and only slightly outside (by 8.5 %) of the
more strict error ellipse used for the DR2 selection. We also checked for the
possibility that more than one star might be ejected from a SN in a multiple
system. But the three other runaway star candidates besides HD 37424, which
projected trajectories could originate from the GC and which were observed
with UVES, are neither consistent with PSR J0538+2817 nor with HD 37424.

Two of the DR2 candidates in the Lupus Loop are particularly promising. Their
positions at the time of the SN were $$$\pm$$$ and $$$\pm$$$ off the GC,
respectively (see Table 9) and both have very high proper motions clearly
directing away from the centre. Spectroscopic follow-up observations are
highly suggested for them.

The historical SNRs generally have very small diameters due to their low age.
This reduces the number of stars in the cone search. Just as Fraser Boubert
(2019), for Cas A and the Crab Nebula we did not find any candidates. For SNR
G130.7+03.1 (from SN 1181) we found one star with $G<17.0$ mag to be
consistent with the GC, but it was ruled out due to the missing kinematical
consistency with PSR J0205+6449. SNR G347.3$-$00.5 probably originates from SN
393 and is much larger than the other three historical SNRs
($\Theta=$$\times$$$). Here we found 18 candidates to be consistent with the
GC. The other SNRs are described in more detail in the following subsection.

###  5.2 Individual SNRs

####  5.2.1 HB9

HB9 is a $$$\times$$$ SNR on the northern hemisphere, where we found six
runaway star candidates around the GC from Gaia DR1, which were observed with
HDS/Subaru. The adopted SNR distance and age are $d=0.8\pm 0.4$ kpc and
$t=5.5\pm 1.5$ kyr, respectively (Leahy Tian, 2007). Among the observed
candidates, we found two giants. The parameters of the other four are given in
Tables 2 and 4. They have a large range of effective temperatures in the
literature, where the Gaia DR2 values are probably underestimated, as can be
seen from comparison with the spectra. Unfortunately, the chosen spectral
range was not convenient for spectral type determination of early-type stars.
Therefore, we rely on the $T_{\text{eff}}$ given by Bai . (2019), which
correspond to spectral type late-A. The RVs were calculated from the positions
of individual absorption lines (see Section 4.2), where the uncertainties are
large due to the low number of available lines. The most precise value is
given for TYC3344-235-1, which also shows the most promising kinematics.
Although the proper motion is moderate, its high radial velocity (pointing
towards us) gives it the highest peculiar velocity among all observed early-
type stars ($v_{\text{pec}}=36.0\pm 2.9$ km s-1). Its high distance of
$d=1103^{+48}_{-45}$ pc would then require that the SNR lies at the upper edge
of its distance range. TYC3344-235-1 and TYC3344-679-1 have high luminosities
of $34\,L_{\odot}$ and $20\,L_{\odot}$, respectively. From their spectra it is
clear that they are not giants. They could be very young A-type stars but it
is more likely that they are evolved, with ages between $\sim$0.2 Gyr and 2
Gyr, especially as TYC3344-235-1 is even located above the isochrone for 1 Myr
in Fig. 7.

There are two neutron stars in the vicinity of HB9. For PSR J0502+4654, its
high characteristic age of $\tau_{c}=1.81$ Myr (Manchester ., 2005) and its
proper motion indicate that it is unrelated. The magnetar SGR 0501+4516 is
located $\sim$$None$ to the south of the centre. If the location of the GC is
close to the explosion site, an association would mean that SGR 0501+4516
travels with more than two times the largest 2D velocity of any other known
pulsar (Hobbs ., 2005), which makes the association very unlikely.

From Gaia DR2, we found one fainter, yet unobserved object consistent with the
GC. As both neutron stars are probably unrelated, we suggest that this is the
most probable candidate. It has $G=16.27$ mag and was located $$$\pm$$$ from
the GC at the time of the SN. Its spectral type is K5.5 $-$ K7 according to
its Gaia DR2 $T_{\text{eff}}=4172_{-67}^{+184}$ K. From its position in Fig. 2
($L=0.199\pm 0.021\,L_{\odot}$), it should have a young age of $\sim$15 $-$ 50
Myr, so it could indeed be the ejected companion of the HB9 progenitor.

####  5.2.2 Monoceros Loop

The Monoceros Loop is one of the largest ($\Theta=$$$) and probably the oldest
SNR of our sample, although the age determination is very uncertain ($30-150$
kyr, Welsh . 2001). From Ferrand Safi-Harb (2012), we adopted $d=1.44\pm 0.54$
kpc, as recent works (Yu ., 2019; Zhao ., 2018) place it at
$d=0.941_{-0.094}^{+0.096}$ kpc or $d=1.257_{-0.101}^{+0.092}$ kpc and
$d=1.98$ kpc, respectively. The discrepancy connects to the question if the
Monoceros Loop interacts with the Rosette Nebula (3C 163), an HII region
bordering the SNR at the south-western edge.

We observed 21 stars with UVES/VLT in the Monoceros Loop, where no associated
neutron star is known. Eight of those were ruled out because they were found
to be giants. Among the others, the most interesting cases are (i) HD 261393,
a B5 V star according to Vo85, which was proposed as the most likely runaway
companion by Bo17, (ii) TYC159-251-1, which was newly discovered by us to be a
double-lined spectroscopic binary and (iii) TYC159-343-1, an F2–F4 star.
TYC159-343-1 and both components of TYC159-251-1 show some Li absorption.
However, the ages inferred in Section 4.5 are higher than what we expect for a
BES runaway star. Also, from the updated distance to the Monoceros Loop and
from using proper motions and parallaxes from Gaia DR2, TYC159-343-1
($d=0.723_{-0.021}^{+0.023}$ kpc) and TYC159-251-1
($d=0.564_{-0.014}^{+0.013}$ kpc) turn out to be foreground stars, even when
the lower limit of the SNR distance given by Yu . (2019) is extended to
$2\,\sigma$ ($0.941\,\text{kpc}-0.188\,\text{kpc}=0.753$ kpc).

Bo17 did not have Gaia DR2, so we can still ask the question why HD 261393
gained a higher probability in their work than TYC159-343-1. There are two
reasons: (1) It is/was located closer to the GC. Bo17 used a prior to describe
the location of the centre as a 2D-Gaussian with a full width at half maximum
(FWHM) of $\tilde{\Theta}=\text{max}($$,0.05\,\Theta)$. This corresponds to
$\tilde{\Theta}=$$$ in the case of the Monoceros Loop. The angular distance
between TYC159-343-1 and the Green centre at the time of the SN was $$$\pm$$$,
almost twice the FWHM. This compares to $$$\pm$$$ for HD 261393. Note that the
errors also take into account the uncertainty of the GC. (2) It is more
massive. Multiplicity studies (e.g. Pinsonneault Stanek 2006) have shown that
massive stars tend to form binaries with similar masses. Therefore, Bo17
assign higher runaway probabilities to more massive, early-type stars.
However, in wide systems the companions can form more independently and the
distribution of mass ratios $q$ becomes almost consistent with random pairings
from the initial mass function (IMF) (Moe Di Stefano, 2017). Sana Evans (2011)
and Sana . (2012) present relatively flat mass ratio distributions down to
$q=0.2$. Unfortunately, the regime of the most extreme mass ratios, $q<0.1$,
is usually not covered by multiplicity studies, because the extreme flux
ratios as well as decreasing RV precision in the case of early-type stars make
low-mass companions hard to detect. However, there is no reason to assume that
the mass ratio distributions should change for $q<0.1$.

Among the other observed stars in the Monoceros Loop, there is TYC159-1896-1
which has spectral type F3 $-$ 9 (according to its $T_{\text{eff}}$ as listed
in Gaia DR2) and which does not show Li, hence is too old, and nine stars of
spectral types mid-B to mid-A. Five of those turn out to be spatially
inconsistent with the GC after reevaluating them with DR2 data.

So in total, we have five good runaway star candidates in the Monoceros Loop,
including HD 261393 and four stars of spectral types B7 $-$ A1, which
trajectories are displayed in Fig. 3 and which kinematic data are noted in
Table 2. Their RVs were determined from the positions of individual absorption
lines (see Section 4.2). For TYC159-2671-1, H$\alpha$ was the only visible
absorption line, so here the error was adopted from the highest error among
the other stars. Their spectral types were adopted from Vo85, where the given
range comes from our error estimate of $\pm 1$ subclass.

A search for further targets in the DR2 catalogue yields more than 100
candidates with $G<17$ mag, due to the large uncertainties of distance and age
of the SNR. We abstain from listing them all in Table 9 because it is not
realistic to observe them all in the near future. It appears more fruitful to
first concentrate on the brighter candidates observed with UVES/VLT.

Considering the large discrepancy of possible distances and ages, the
identification of a runaway star would be particularly important for the
Monoceros Loop to constrain these parameters significantly.

####  5.2.3 Puppis A

The Vela region, located on the southern sky in the Vela constellation, is a
very extended and complex region with several stellar clusters, star-forming
regions and three SNRs.

The most distant SNR in this region is Puppis A. Here, two runaway star
candidates were observed with UVES, namely TYC7669-1336-1 and TYC7669-1414-1,
which could both originate from the GC in projection. According to their
$T_{\text{eff}}$ from Gaia DR2, their spectral types are F3 $-$ F6 and F1 $-$
F2, respectively. No Li was found in the spectra. After reevaluating them with
the DR2 parallax, it was found that they are too close to be consistent with
the distance given for Puppis A ($d=1.3\pm 0.3$ kpc, Reynoso . 2017).

A search in Gaia DR2 yields nine candidates with $G<17$ mag, where eight of
them have proper motions pointing to north-west, similar to most other stars
in this region of the sky. Only one candidate, a K1.5 $-$ K3 star with
$G=13.2$ mag, has a higher proper motion and a slightly different direction
compared to the other stars, pointing to west-north-west. It might, therefore,
be the best candidate. It was located $$$\pm$$$ from the GC at the time of the
SN. With $T_{\text{eff}}=5042_{-180}^{+115}$ K and $L=4.64\pm
0.21\,L_{\odot}$, from Fig. 2 it could be consistent with a young age of
$\sim$1 $-$ 2 Myr, although it is more likely an evolved star with $4.5-11$
Gyr.

####  5.2.4 Vela

With $d=0.275\pm 0.025$ kpc, Vela is the closest SNR, therefore having a large
diameter of $\Theta=$$$. Its age is only poorly constrained to $t=18\pm 9$ kyr
(Ferrand Safi-Harb, 2012). However, the characteristic age of the Vela pulsar,
$\tau_{c}=11.3$ kyr, indicates that the SNR age could be close to the lower
limit.

Two stars were observed, TYC8150-2802-1 and TYC8150-3105-1, both showing Li in
their spectra. Their spectral types are G0 $-$ G2 and F9.5 $-$ G1,
respectively. From the Li content and the position in the HRD, we can conclude
that TYC8150-3105-1 has a relatively young age of $25-50$ Myr, while
TYC8150-2802-1 is probably an evolved MS star. Anyway, they cannot be
associated runaway stars, because their trajectories are not consistent with
the past position of the Vela pulsar (J0835$-$4510). TYC8150-3105-1 could be a
member of the young association to which also the Vela progenitor belonged,
whereas TYC8150-2802-1 is probably an interloper.

Fig. 11 shows the motion of the Li-rich stars in yellow and the Vela pulsar in
magenta. Two further candidates are marked, found in Gaia DR2, which are
consistent with the Vela pulsar: The star marked in blue was identified by us
and could be observed in the near future. The star marked in black was
identified by Fraser Boubert (2019) (hereafter FB19), therein denoted as Star
A. It is closer to the nominal past position of the Vela pulsar, but with
$G=20.1$ mag it is much too faint to obtain a spectrum with sufficient $S/N$.
FB19 note that it is unlikely that the ejected companion of the Vela SN
progenitor would be such faint (absolute magnitude $M_{G}=12.7$ mag), so
probably Star A is an unrelated background star. FB19 used the distribution of
runaway star velocities given by Rz19 to constrain their search radius, which
limits their selection more strictly than the
$v_{\text{max}}=1000\,\text{km\,s}^{-1}$ used by us. Therefore, the star
marked in blue in Fig. 11 was not found by them.

Figure 11: The central $$$\times$$$ of the Vela SNR. The red labels mark the
geometric centre and its error ellipse, the yellow labels the Li-rich stars
observed in ESO P100, namely (3) TYC8150-2802-1 and (4) TYC8150-3105-1. The
crosses mark the current positions, connected to the positions at the time of
the SN (with error ellipse) by an arrow in proper motion direction. The motion
of the Vela PSR (J0835$-$4510) is shown in magenta, our runaway candidate from
Gaia DR2 in blue and the candidate identified by FB19 in black. Background
image from ESO DSS-2-red.

####  5.2.5 Vela Junior

Vela Jr. is located about $None$ east of Vela. With an age between 2.4 kyr and
5.1 kyr it is younger than Vela and its distance is between 0.5 kpc and 1.0
kpc (Allen ., 2015).

Five stars were observed with UVES: TYC8152-120-1 is a mid-F star, where no Li
was discovered. TYC8152-104-1 is a K giant, and TYC8152-1456-1 (F9.5 $-$ G1)
and TYC8152-550-1 (F7 $-$ F9) show Li absorption in their spectra. However,
the measured abundances indicate that they are much too old to be runaway
stars associated with the Vela Jr. progenitor (see Figs. 9, 10, Table 6). HD
76060 is the brightest star of our observed sample with $G=7.85$ mag. Its
spectral type is B8 IV–V, according to Houk (1978). From its Gaia DR2
parallax, we infer a distance of $d=0.486_{-0.010}^{+0.011}$ kpc, which could
still be consistent with the lower distance limit of Vela Jr. Its position at
the time of the SN is consistent with the GC according to Eqn. 2 and its
peculiar velocity of $31.2\pm 2.7$ km s-1 (see Table 2 for further parameters)
is relatively high. Its age, according to Fig. 7, is either $1.1-1.8$ Myr (if
pre-MS) or $90-350$ Myr (if post-MS). Although the post-MS age, which would be
too high, is more likely, the pre-MS cannot be excluded. Furthermore, with
$2\,\sigma$ error bars it could also be young enough if it is post-MS.
Therefore, we still consider HD 76060 as a promising candidate.

As the association to PSR J0855$-$4644 is unlikely, we made a search for DR2
candidates which trace back to a position around the GC according to Eqn. 2
and found twelve objects with $G<17$ mag which are considered good runaway
candidates together with HD 76060 and are suggested for follow-up
observations.

###  5.3 Constraints on runaway star production

Here, we want to give an estimate how the findings of this work fit model
populations of runaway stars in the literature (Bo17, Rz19). We investigated
twelve SNRs and can exclude runaway stars with $G<17$ mag in three of them.
For SNR S147, we can confirm HD 37424 with Gaia DR2 data. Based on this
association, we find no further runaway stars within S147. In each of the
remaining eight SNRs, we find one or more candidates, but none of those can be
confirmed yet.

So the minimum number of BES runaway stars within the twelve investigated SNRs
is one, while the maximum number is nine, corresponding to a fraction of
ejected runaway stars per SNR of $8-75$ %. The upper limit could be even
higher if we also count the cases where more than one runaway star could have
been ejected. In three of the SNRs (Cygnus Loop, Vela Jr., Lupus Loop), two of
the remaining candidates could have a common origin. Although many more error
ellipses of the runaway star candidates are overlapping, it was considered
that for a certain age of the SNR the error regimes would shrink
correspondingly.

This compares to theoretical values for runaway stars ejected from core-
collapse SNe between 32.5 % (Bo17) and 68 % (Rz19). The latter value comes
from $78_{-22}^{+9}$ % of binary systems that do not merge before the first
SN, multiplied by $86_{-22}^{+10}$ % that get disrupted during the SN, which
gives $67_{-26}^{+17}$ % of high-mass binary systems that eject a runaway
star. Furthermore, Bo17 state that $\sim$1.22 core-collapse SNe happen in each
high-mass binary system, so we get $55_{-21}^{+13}$ % of core-collapse SNe
that eject a runaway star.

Both values are consistent with our findings. We need follow-up observations
of our best candidates to search for SN debris as well as more precise
estimates for the ages, distances and explosion sites of the SNRs to further
constrain the numbers of ejected runaway stars in SNRs. From the theoretical
values, we expect to verify $3-7$ of our runaway star candidates within the
eleven SNRs besides S147 in the future.

##  6 Conclusions

Our search in the Gaia DR2 catalogue for runaway stars in twelve Galactic SNRs
yielded no further certain runaway stars besides HD 37424, but 73 new
promising candidates, which are listed together with HD 37424 in Table 9.
Among these, five stars in the Monoceros Loop and one in Vela Jr. were
observed with UVES/VLT and four stars in HB9 with Subaru/HDS.

In total, we observed 39 stars, where 29 were ruled out, e.g. after revisiting
their motion with Gaia DR2 data or they turned out to be giants. We found six
dwarf stars with lithium; two in Monoceros, two in Vela and two in Vela Jr.
TYC159-343-1, located in the Monoceros Loop in projection, might be a young
star. We obtained $90-625$ Myr from Fig 9 but its Li signal is hard to
interpret due to the early spectral type (F$3\pm 1$) (D’Antona Mazzitelli,
1984), and its low metallicity ([M/H] $=-0.57\pm 0.08$) (Lambert Reddy, 2004).
Even $2.1-3.8$ Myr, the pre-MS solution from Fig. 7, is possible. However,
using DR2 data and an updated distance of the Monoceros Loop, TYC159-343-1 had
to be ruled out because it is too close. The same applies to TYC159-251-1,
which was discovered by us to be a double-lined spectroscopic binary.
TYC8150-3105-1 in Vela shows the strongest Li absorption among all observed
stars, so it was found to be younger than $\sim$50 Myr. The Li-rich stars in
Vela (TYC8150-3105-1 and TYC8150-2802-1) were ruled out because they are not
consistent with the motion of the PSR. They could be interlopers near the
volume of the SNR while TYC8150-3105-1 could also be a member of the same OB
association from which the SNR formed. The Li-rich stars in Vela Jr. were
ruled out because they are too old, which can be seen from their relatively
small Li abundances.

The B-type star HD 261393 is the most likely candidate in the Monoceros Loop.
In Vela, we point out two stars which are consistent with the motion of the
Vela PSR. The star marked in blue in Fig. 11 is suggested for observations,
while Star A (FB19) is too faint to be investigated spectroscopically.

In total, we found 74 good runaway star candidates with $G<17$ mag, mostly
with spectral type K. We suggest to take high-resolution spectra to search for
the Li 6708 Å line. Among the twelve investigated SNRs, nine SNe could have
ejected a runaway star, while only HD 37424 in S147 is confirmed. Three of
these SNe could possibly have ejected two runaway stars.

Note that we might miss some stars with $G>17$ mag which would mainly be
spectral type M. Due to missing observations of such late-type stars around
high-mass primaries, we do not know how many M-type runaway stars we could
expect. Furthermore, we might miss some Gaia DR2 candidates due to our strict
selection limit from the allowed angular distance to the geometric centre (GC)
at the time of the SN, which was necessary to create feasible observing
projects.

In order to finally proof a BES origin, we need to observe the best candidates
with very high resolution and $S/N$ to be able to detect SN debris, e.g.
heavy- and $\alpha$-elements, in the stellar atmospheres. We also emphasise
that it is important to precisely know the explosion site. The case of HD
37424 in SNR S147 is rather exceptional, because both the pulsar and the
runaway star trace back to the Green centre and meet it within just a few
arcminutes. We did not find other cases where a pair of a pulsar and a runaway
star candidate was located so close to the GC. In general, the explosion site
will not coincide with the GC, due to the motion of the local standard of rest
and due to asymmetries of the SNR expansion (e.g. Meyer . 2015, 2020). A
careful analysis of the time-dependent SNR morphology is needed to locate the
explosion site.

## Acknowledgements

We want to thank our present and former colleagues for many hints and fruitful
discussions, in particular Baha Dinçel, Anna Pannicke and Christian Adam. We
acknowledge Sergi Blanco-Cuaresma for many useful hints for the usage of iSpec
and the analysis of stellar spectra in general; Eric Mamajek for providing the
data and fits for $T_{\text{eff}}$ versus $EW_{\text{Li}}$ for clusters of
different ages; and Aaron Steinhauer for providing the data that relate
$\log(N_{\text{Li}})$, $v_{\text{rot}}\sin(i)$ and age. This work made use of
ESO VLT data from run 0100.D-0314 and of NAOJ Subaru data from run S18B0195S.
Therefore, we would like to thank the ESO and NAOJ staff, in particular the
local support astronomers as well as John Pritchard (ESO), Chie Yoshida (NAOJ)
and Akito Tajitsu (NAOJ), for the observations and many useful hints and
comments about the data reduction. The NAOJ Subaru telescope is located on the
Mauna Kea and we would like to acknowledge the important cultural and
spiritual role, which the mountain has for the indigenous Hawaiian community.
This work made extensive use of data from the ESA-mission
Gaia101010https://www.cosmos.esa.int/gaia, which were processed and provided
by the Gaia Data Processing and Analysis Consortium (DPAC). We further used
VizieR, Simbad and Aladin, provided by the Centre de Donnée astronomiques de
Strasbourg (CDS), and the Australia Telescope National Facility (ATNF) Pulsar
Catalogue (Manchester ., 2005).

Funding was provided by the Deutsche Forschungsgemeinschaft, project NE
515/57-1.

## References

  * Allen . (2015) Allen2015Allen, GE., Chow, K., DeLaney, T., Filipović, MD., Houck, JC., Pannuti, TG. Stage, MD.  2015Jan,  ApJ798282.
  * Anders . (2019) Anders2019Anders, F., Khalatyan, A., Chiappini, C. et al.  2019Aug,  A&A628A94.
  * Aoki Helminiak (2014) Aoki2014Aoki, W. Helminiak, K.  2014,  Data reduction of echelle spectra with IRAF Data reduction of echelle spectra with IRAF [].  Tokyo, Japan.
  * Bagnulo . (2003) Bagnulo2003Bagnulo, S., Jehin, E., Ledoux, C., Cabanac, R., Melo, C., Gilmozzi, R. ESO Paranal Science Operations Team.  2003Dec,  The Messenger11410-14.
  * Bai . (2019) Bai2019Bai, Y., Liu, J., Bai, Z., Wang, S. Fan, D.  2019Aug,  AJ158293.
  * Bailer-Jones . (2018) Bailer-Jones2018Bailer-Jones, CAL., Rybizki, J., Fouesneau, M., Mantelet, G. Andrae, R.  2018Aug,  AJ156258.
  * Blaauw (1961) Blaauw1961Blaauw, A.  1961May,  Bull. Astron. Inst. Netherlands15265.
  * Blanco-Cuaresma (2019) Blanco-Cuaresma2019Blanco-Cuaresma, S.  2019Jun,  MNRAS48622075-2101.
  * Blanco-Cuaresma . (2014) Blanco-Cuaresma2014Blanco-Cuaresma, S., Soubiran, C., Heiter, U. Jofré, P.  2014Sep,  A&A569A111.
  * Bock . (1998) Bock1998Bock, DCJ., Turtle, AJ. Green, AJ.  1998Oct,  AJ11641886-1896.
  * Bodenheimer (1965) Bodenheimer1965Bodenheimer, P.  1965Aug,  ApJ142451.
  * Boesgaard . (1988) Boesgaard1988Boesgaard, AM., Budge, KG. Ramsay, ME.  1988Apr,  ApJ327389.
  * Boesgaard Tripicco (1986) Boesgaard1986Boesgaard, AM. Tripicco, MJ.  1986Mar,  ApJ302L49.
  * Boubert . (2017) Boubert2017Boubert, D., Fraser, M., Evans, NW., Green, DA. Izzard, RG.  2017Sep,  A&A606A14.
  * Breitschwerdt . (2016) Breitschwerdt2016Breitschwerdt, D., Feige, J., Schulreich, MM., Avillez, MAD., Dettbarn, C. Fuchs, B.  2016Apr,  Nature532759773-76.
  * Bressan . (2012) Bressan2012Bressan, A., Marigo, P., Girardi, L., Salasnich, B., Dal Cero, C., Rubele, S. Nanni, A.  2012Nov,  MNRAS4271127-145.
  * Butler . (1987) Butler1987Butler, RP., Cohen, RD., Duncan, DK. Marcy, GW.  1987Aug,  ApJ319L19.
  * Chatterjee . (2004) Chatterjee2004Chatterjee, S., Cordes, JM., Vlemmings, WHT., Arzoumanian, Z., Goss, WM. Lazio, TJW.  2004Mar,  ApJ6041339-345.
  * Coluzzi (1999) Coluzzi1999Coluzzi, R.  1999Sep,  VizieR Online Data CatalogVI/71A.
  * Coskunoǧlu . (2011) Coskunoglu2011Coskunoǧlu, B., Ak, S., Bilir, S. et al.  2011Apr,  MNRAS41221237-1245.
  * Covino . (1997) Covino1997Covino, E., Alcala, JM., Allain, S., Bouvier, J., Terranegra, L. Krautter, J.  1997Dec,  A&A328187-202.
  * D’Antona Mazzitelli (1984) Dantona1984D’Antona, F. Mazzitelli, I.  1984Sep,  A&A138431-442.
  * Deliyannis . (2019) Deliyannis2019Deliyannis, CP., Anthony-Twarog, BJ., Lee-Brown, DB. Twarog, BA.  2019Oct,  AJ1584163.
  * Dinçel . (2015) Dincel2015Dinçel, B., Neuhäuser, R., Yerli, SK., Ankay, A., Tetzlaff, N., Torres, G. Mugrauer, M.  2015Apr,  MNRAS44843196-3205.
  * Eastman . (2010) Eastman2010Eastman, J., Siverd, R. Gaudi, BS.  2010Aug,  PASP122894935.
  * Ekström . (2012) Ekstroem2012Ekström, S., Georgy, C., Eggenberger, P. et al.  2012Jan,  A&A537A146.
  * Evans . (2018) Evans2018Evans, DW., Riello, M., De Angeli, F. et al.  2018Aug,  A&A616A4.
  * Ferrand Safi-Harb (2012) Ferrand2012Ferrand, G. Safi-Harb, S.  2012May,  Advances in Space Research4991313-1319.
  * Fraser Boubert (2019) Fraser2019Fraser, M. Boubert, D.  2019Jan,  ApJ871192.
  * Gaia Collaboration . (2018) Gaia2018Gaia Collaboration, Brown, AGA., Vallenari, A. et al.  2018Aug,  A&A616A1.
  * Gaia Collaboration, Brown . (2016) Gaia2016aGaia Collaboration, Brown, AGA., Vallenari, A. et al.  2016Nov,  A&A595A2.
  * Gaia Collaboration, Prusti . (2016) Gaia2016bGaia Collaboration, Prusti, T., de Bruijne, JHJ. et al.  2016Nov,  A&A595A1.
  * Green (2009) Green2009Green, DA.  2009Mar,  Bulletin of the Astronomical Society of India37145-61.
  * Green (2014) Green2014Green, DA.  2014Jun,  Bulletin of the Astronomical Society of India42247-58.
  * Henry McCarthy (1993) Henry1993Henry, TJ. McCarthy, J., Donald W.  1993Aug,  AJ106773.
  * Hills (1988) Hills1988Hills, JG.  1988Feb,  Nature3316158687-689.
  * Hobbs . (2005) Hobbs2005Hobbs, G., Lorimer, DR., Lyne, AG. Kramer, M.  200507,  MNRAS3603974-992.
  * Hoogerwerf . (2000) Hoogerwerf2000Hoogerwerf, R., de Bruijne, JHJ. de Zeeuw, PT.  2000Dec,  ApJ5442L133-L136.
  * Hoogerwerf . (2001) Hoogerwerf2001Hoogerwerf, R., de Bruijne, JHJ. de Zeeuw, PT.  2001Jan,  A&A36549-77.
  * Houk (1978) Houk1978Houk, N.  1978,  Michigan catalogue of two-dimensional spectral types for the HD stars Michigan catalogue of two-dimensional spectral types for the HD stars.
  * Johnson Soderblom (1987) Johnson1987Johnson, DRH. Soderblom, DR.  1987Apr,  AJ93864.
  * Knie . (1999) Knie1999Knie, K., Korschinek, G., Faestermann, T., Wallner, C., Scholten, J. Hillebrandt, W.  1999Jul,  Phys. Rev. Lett.83118-21.
  * Lambert Reddy (2004) Lambert2004Lambert, DL. Reddy, BE.  2004Apr,  MNRAS3492757-767.
  * Leahy Tian (2007) Leahy2007Leahy, DA. Tian, WW.  2007Jan,  A&A46131013-1018.
  * Manchester . (2005) Manchester2005Manchester, RN., Hobbs, GB., Teoh, A. Hobbs, M.  2005Apr,  AJ12941993-2006.
  * Meyer . (2015) Meyer2015Meyer, DMA., Langer, N., Mackey, J., Velázquez, PF. Gusdorf, A.  2015Jul,  MNRAS45033080-3100.
  * Meyer . (2020) Meyer2020Meyer, DMA., Petrov, M. Pohl, M.  2020Apr,  MNRAS49333548-3564.
  * Moe Di Stefano (2017) Moe2017Moe, M. Di Stefano, R.  2017Jun,  ApJS230215.
  * Mugrauer (2019) Mugrauer2019Mugrauer, M.  2019Dec,  MNRAS49045088-5102.
  * Mugrauer . (2014) Mugrauer2014Mugrauer, M., Avila, G. Guirao, C.  2014Jan,  Astronomische Nachrichten3354417.
  * Mugrauer Berthold (2010) Mugrauer2010Mugrauer, M. Berthold, T.  2010Apr,  Astronomische Nachrichten3314449.
  * Neuhäuser (1997) Neuhaeuser1997Neuhäuser, R.  1997Jan,  Science2761363-1370.
  * Neuhäuser . (2020) Neuhaeuser2019Neuhäuser, R., Gießler, F. Hambaryan, VV.  2020Oct,  MNRAS4981899-917.  10.1093/mnras/stz2629
  * Pecaut Mamajek (2013) Pecaut2013Pecaut, MJ. Mamajek, EE.  2013Sep,  ApJS20819.
  * Pilachowski . (1987) Pilachowski1987Pilachowski, CA., Booth, J. Hobbs, LM.  1987Dec,  PASP991288-1291.
  * Pinsonneault Stanek (2006) Pinsonneault2006Pinsonneault, MH. Stanek, KZ.  2006Mar,  ApJ6392L67-L70.
  * Pourbaix . (2009) Pourbaix2009Pourbaix, D., Tokovinin, AA., Batten, AH. et al.  2009Aug,  VizieR Online Data CatalogB/sb9.
  * Poveda . (1967) Poveda1967Poveda, A., Ruiz, J. Allen, C.  1967Apr,  Boletin de los Observatorios Tonantzintla y Tacubaya486-90.
  * Przybilla . (2008) Przybilla2008Przybilla, N., Fernanda Nieva, M., Heber, U. Butler, K.  2008Sep,  ApJ6842L103.
  * Renzo . (2019) Renzo2019Renzo, M., Zapartas, E., de Mink, SE. et al.  2019Apr,  A&A624A66.
  * Reynoso . (2017) Reynoso2017Reynoso, EM., Cichowolski, S. Walsh, AJ.  2017Jan,  MNRAS46433029-3039.
  * Sana . (2012) Sana2012Sana, H., de Mink, SE., de Koter, A. et al.  2012Jul,  Science3376093444.
  * Sana Evans (2011) Sana2011Sana, H. Evans, CJ.  2011Jul,  The multiplicity of massive stars The multiplicity of massive stars.  C. Neiner, G. Wade, G. Meynet . (), Active OB Stars: Structure, Evolution, Mass Loss, and Critical Limits. Proc. International Astronomical Union Active OB Stars: Structure, Evolution, Mass Loss, and Critical Limits. Proc. International Astronomical Union 272, 474-485.
  * Sestito Randich (2005) Sestito2005Sestito, P. Randich, S.  2005Nov,  AAP4422615-627.
  * Skiff (2009) Skiff2009Skiff, BA.  2009Feb,  VizieR Online Data CatalogB/mk.
  * Soderblom (20101) Soderblom2010aSoderblom, DR.  20101Sep,  ARA&A48581-629.
  * Soderblom (20102) Soderblom2010bSoderblom, DR.  20102Apr,  Using lithium to estimate ages for solar-type stars Using lithium to estimate ages for solar-type stars.  C. Charbonnel, M. Tosi, F. Primas . (), Light Elements in the Universe. Proc. International Astronomical Union Light Elements in the Universe. Proc. International Astronomical Union 268, 359-360.
  * Soderblom . (1993) Soderblom1993Soderblom, DR., Jones, BF., Balachand ran, S., Stauffer, JR., Duncan, DK., Fedele, SB. Hudon, JD.  1993Sep,  AJ1061059.
  * Steinhauer (2003) Steinhauer2003Steinhauer, A.  2003\.  Formation and evolution of the open cluster lithium gap Formation and evolution of the open cluster lithium gap .  Indiana University.
  * Steinhauer Deliyannis (2004) Steinhauer2004Steinhauer, A. Deliyannis, CP.  2004Oct,  ApJ6141L65-L68.
  * Steinhauer Deliyannis (2021) Steinhauer2021Steinhauer, A. Deliyannis, CP.  2021,  submitted.
  * Sushch . (2011) Sushch2011Sushch, I., Hnatyk, B. Neronov, A.  2011Jan,  A&A525A154.
  * Tauris (2015) Tauris2015Tauris, TM.  2015Mar,  MNRAS448L6-L10.
  * Tetzlaff (2013) Tetzlaff2013Tetzlaff, N.  2013\.  Identifying birth places of young neutron stars to determine their kinematic age Identifying birth places of young neutron stars to determine their kinematic age doctoral thesis.  Friedrich-Schiller-Universität Jena (FSU).
  * Tetzlaff . (2014) Tetzlaff2014Tetzlaff, N., Dinçel, B., Neuhäuser, R. Kovtyukh, VV.  2014Mar,  MNRAS43843587-3593.
  * Tetzlaff . (2011) Tetzlaff2011Tetzlaff, N., Eisenbeiss, T., Neuhäuser, R. Hohle, MM.  2011Oct,  MNRAS4171617-626.
  * Tetzlaff . (2010) Tetzlaff2010Tetzlaff, N., Neuhäuser, R., Hohle, MM. Maciejewski, G.  2010Mar,  MNRAS40242369-2387.
  * Thorburn . (1993) Thorburn1993Thorburn, JA., Hobbs, LM., Deliyannis, CP. Pinsonneault, MH.  1993Sep,  ApJ415150.
  * Voroshilov . (1985) Voroshilov1985Voroshilov, VI., Guseva, NG., Kalandadze, NB., Kolesnik, LN., Kuznetsov, VI., Metreveli, MD. Shapovalov, AN.  1985,  Catalogue of BV magnitudes and spectral classes for 6000 stars. Ukrainian Acad. Nauk, Kiev, 1-140. Catalogue of BV magnitudes and spectral classes for 6000 stars. Ukrainian Acad. Nauk, Kiev, 1-140.
  * Wallner . (2016) Wallner2016Wallner, A., Feige, J., Kinoshita, N. et al.  2016Apr,  Nature532759769-72.
  * Welsh . (2001) Welsh2001Welsh, BY., Sfeir, DM., Sallmen, S. Lallement, R.  2001Jun,  A&A372516-526.
  * Yu . (2019) Yu2019Yu, B., Chen, BQ., Jiang, BW. Zijlstra, A.  2019Sep,  MNRAS48833129-3142.
  * Zehe . (2018) Zehe2018Zehe, T., Mugrauer, M., Neuhäuser, R. et al.  2018Jan,  Astronomische Nachrichten339146-52.
  * Zhao . (2018) Zhao2018Zhao, H., Jiang, B., Gao, S., Li, J. Sun, M.  2018Mar,  ApJ855112.

##  Appendix A Overview of the runaway candidates

Table 8 shows the individual stars that were observed with UVES/VLT in P100
and with HDS/Subaru in S18B, respectively. All observations were executed in
service mode. For the HDS spectra, the $S/N$ was measured in continuum regions
between 6620 Å and 6760 Å in the combined spectra with IRAF, while for the
UVES spectra we obtained the $S/N$ from the ESO Archive Science
Portal111111http://archive.eso.org/scienceportal/home. Julian date and
barycentric Julian date were calculated with the online tools from
AAVSO121212https://www.aavso.org/jd-calculator and Eastman .
(2010)131313http://astroutils.astronomy.ohio-state.edu/time/utc2bjd.html,
respectively.

Table 9 gives an overview of the 74 possible runaway stars that were
identified during this analysis. The spectral types SpT were estimated from
the Gaia DR2 $T_{\text{eff}}$ if available, with the following exceptions: HD
37424 was observed, described and confirmed as a runaway star by Dinçel .
(2015). The stars in SNR G205.5+00.5 as well as HD76060 in G266.2$-$01.2 were
observed by us with UVES and their spectral types were adopted from the Skiff
catalogue (Skiff 2009; Houk 1978; Vo85). The stars in G160.9+02.6 (except the
faintest one) were observed by us with HDS and their spectral types were
adopted from the $T_{\text{eff}}$ given in Bai . (2019).

Table 8: List of observed stars. The columns give target name, instrument (U for UVES/VLT; H for HDS/Subaru), equatorial coordinates (J2000) calculated with VizieR from Gaia DR2 data, the SNR where the target is located, Gaia DR2 G magnitude, barycentric Julian date BJD, total on source integration time $t_{\text{exp}}$ and the $S/N$ reached in the fully reduced and averaged spectra. Name | Instr. | RA | DEC | SNR | $G$ | BJD$-2450000$ | $t_{\text{exp}}$ | $S/N$
---|---|---|---|---|---|---|---|---
|  | [h:m:s] | [d:m:s] |  | [mag] | [d] | [s] |
TYC1869-1435-1 | U | 05:38:15.15 |  $+$27:49:18.48 | G180.0$-$01.7 | 12.20 | 8056.84606 | 900 | 112
TYC1869-1596-1 | U | 05:38:23.71 |  $+$27:41:33.07 | G180.0$-$01.7 | 12.16 | 8060.80811 | 800 | 106
TYC1869-1670-1 | U | 05:39:03.05 |  $+$27:32:35.26 | G180.0$-$01.7 | 12.57 | 8058.82311 | 900 | 87
TYC159-2540-1 | U | 06:37:52.95 |  $+$06:39:53.48 | G205.5$+$00.5 | 12.73 | 8060.84497 | 300 | 48
HD 260990 | U | 06:37:57.18 |  $+$06:28:47.47 | G205.5$+$00.5 | 11.63 | 8075.79833 | 300 | 76
HD 261117 | U | 06:38:19.74 |  $+$06:23:40.70 | G205.5$+$00.5 | 9.74 | 8036.85181 | 60 | 81
TYC159-241-1 | U | 06:38:25.23 |  $+$06:11:31.56 | G205.5$+$00.5 | 10.47 | 8076.79745 | 240 | 104
TYC159-2006-1 | U | 06:38:32.24 |  $+$06:39:02.91 | G205.5$+$00.5 | 9.28 | 8076.80388 | 60 | 95
TYC159-2408-1 | U | 06:38:32.28 |  $+$06:30:22.52 | G205.5$+$00.5 | 9.94 | 8076.80850 | 140 | 103
TYC159-2771-1 | U | 06:38:50.77 |  $+$06:35:26.03 | G205.5$+$00.5 | 12.01 | 8076.81354 | 480 | 69
TYC159-1896-1 | U | 06:38:51.46 |  $+$06:39:15.88 | G205.5$+$00.5 | 11.48 | 8076.82462 | 360 | 85
TYC159-2082-1 | U | 06:38:59.13 |  $+$06:32:37.00 | G205.5$+$00.5 | 11.08 | 8076.83232 | 300 | 94
HD 261359 | U | 06:39:07.47 |  $+$06:27:37.99 | G205.5$+$00.5 | 11.79 | 8067.76600 | 720 | 114
HD 261393 | U | 06:39:13.09 |  $+$06:37:53.98 | G205.5$+$00.5 | 10.05 | 8060.81954 | 120 | 110
TYC159-2337-1 | U | 06:39:20.12 |  $+$06:21:00.47 | G205.5$+$00.5 | 11.80 | 8107.61880 | 800 | 109
TYC159-2671-1 | U | 06:39:24.25 |  $+$06:33:08.53 | G205.5$+$00.6 | 12.25 | 8076.84023 | 480 | 64
TYC159-892-1 | U | 06:39:24.38 |  $+$06:37:25.98 | G205.5$+$00.5 | 9.44 | 8107.65051 | 100 | 109
HD 261527 | U | 06:39:33.69 |  $+$06:06:58.87 | G205.5$+$00.5 | 8.41 | 8107.65524 | 20 | 83
TYC159-2962-1 | U | 06:39:36.43 |  $+$06:16:36.10 | G205.5$+$00.5 | 9.30 | 8107.65808 | 100 | 106
HD 261589 | U | 06:39:44.89 |  $+$06:27:35.57 | G205.5$+$00.5 | 11.06 | 8107.64301 | 360 | 99
TYC159-251-1 | U | 06:39:47.09 |  $+$06:11:16.54 | G205.5$+$00.5 | 11.51 | 8107.66221 | 360 | 83
TYC159-2564-1 | U | 06:40:00.11 |  $+$06:17:36.17 | G205.5$+$00.5 | 11.97 | 8107.63233 | 600 | 80
HD 261715 | U | 06:40:13.41 |  $+$06:16:45.02 | G205.5$+$00.5 | 11.25 | 8107.66940 | 480 | 98
TYC159-343-1 | U | 06:40:15.57 |  $+$06:11:36.61 | G205.5$+$00.5 | 10.87 | 8107.67796 | 180 | 75
TYC7669-1336-1 | U | 08:22:05.66 |  $-$42:58:08.41 | G260.4$-$03.4 | 12.27 | 8078.79809 | 1000 | 100
TYC7669-1414-1 | U | 08:22:10.86 |  $-$43:04:17.62 | G260.4$-$03.4 | 11.97 | 8078.81243 | 800 | 94
TYC8150-2802-1 | U | 08:34:02.97 |  $-$45:32:59.68 | G263.9$-$03.3 | 12.03 | 8078.82413 | 1000 | 107
TYC8150-3105-1 | U | 08:34:35.19 |  $-$45:11:37.87 | G263.9$-$03.3 | 11.94 | 8078.83783 | 600 | 90
TYC8152-120-1 | U | 08:52:02.14 |  $-$46:23:22.99 | G266.2$-$01.2 | 11.58 | 8107.68437 | 600 | 80
HD 76060 | U | 08:52:02.45 |  $-$46:17:19.84 | G266.2$-$01.2 | 7.85 | 8035.85245 | 20 | 116
TYC8152-1456-1 | U | 08:52:05.71 |  $-$46:13:37.61 | G266.2$-$01.2 | 12.47 | 8077.80900 | 720 | 72
TYC8152-104-1 | U | 08:52:08.05 |  $-$46:12:09.38 | G266.2$-$01.2 | 11.53 | 8077.82179 | 480 | 91
TYC8152-550-1 | U | 08:52:14.68 |  $-$46:14:50.00 | G266.2$-$01.2 | 11.92 | 8077.83090 | 480 | 80
TYC3344-235-1 | H | 05:00:31.29 |  $+$46:33:27.25 | G160.9$+$02.6 | 11.07 | 8418.99274 | 600 | 145
TYC3344-679-1 | H | 05:00:53.98 |  $+$46:30:00.61 | G160.9$+$02.6 | 11.40 | 8419.00315 | 600 | 151
TYC3344-771-1 | H | 05:01:10.74 |  $+$46:42:11.04 | G160.9$+$02.6 | 10.60 | 8419.01078 | 360 | 185
TYC3344-683-1 | H | 05:01:10.87 |  $+$46:27:38.11 | G160.9$+$02.6 | 12.33 | 8419.02190 | 1200 | 142
TYC3344-124-1 | H | 05:01:11.84 |  $+$46:47:57.30 | G160.9$+$02.6 | 10.45 | 8419.03231 | 240 | 136
TYC3344-553-1 | H | 05:01:16.00 |  $+$46:33:21.68 | G160.9$+$02.6 | 10.87 | 8419.03787 | 360 | 139
Table 9: List of identified runaway star candidates. Equatorial coordinates
(J2000) were calculated with VizieR from Gaia DR2 data. The fifth column shows
the angular distance $\rho$ to the geometric centre at the time of the SN,
except for the candidate in Vela (G263.9$-$03.3), where the distance to the
location of the Vela pulsar at the time of the SN is given. The sixth column
shows the $G$ magnitudes, taken from Gaia DR2, and the last column shows the
spectral types SpT.

Gaia DR2 | RA [h:m:s] | DEC [d:m:s] | SNR |  $\rho$ [arcmin] |  $G$ [mag] | SpT
---|---|---|---|---|---|---
1859462217726265728 | 20:50:22.21 |  $+$30:30:17.37 | G074.0$-$08.5 | $2.3\pm 4.7$ | 16.23 | K2–K3.5
1859461771049581952 | 20:50:47.80 |  $+$30:30:43.62 | G074.0$-$08.5 | $3.4\pm 5.1$ | 16.88 | K5–K6
1859469532044083968 | 20:50:54.62 |  $+$30:39:06.92 | G074.0$-$08.5 | $3.5\pm 3.7$ | 15.20 | K2–K3
1859469364551815040 | 20:50:54.66 |  $+$30:37:29.63 | G074.0$-$08.5 | $3.4\pm 3.8$ | 15.65 | K2.5–K3.5
1859469467631489920 | 20:50:55.57 |  $+$30:38:13.63 | G074.0$-$08.5 | $1.1\pm 3.5$ | 15.47 | K2–K3.5
1858709945619185664 | 20:50:56.65 |  $+$30:26:27.61 | G074.0$-$08.5 | $4.9\pm 5.9$ | 15.06 | G9–K2
1859471357406309120 | 20:50:57.42 |  $+$30:42:38.52 | G074.0$-$08.5 | $2.6\pm 4.3$ | 14.44 | –
1859471254338021120 | 20:50:58.70 |  $+$30:40:51.90 | G074.0$-$08.5 | $0.6\pm 3.4$ | 16.67 | K1.5–K3.5
1859471288697195136 | 20:50:59.09 |  $+$30:41:00.61 | G074.0$-$08.5 | $1.1\pm 4.3$ | 14.22 | G9
1859471082528259840 | 20:51:00.80 |  $+$30:40:47.78 | G074.0$-$08.5 | $1.7\pm 4.2$ | 15.94 | G9–K3.5
1859470975153131904 | 20:51:01.85 |  $+$30:39:17.77 | G074.0$-$08.5 | $2.0\pm 4.2$ | 14.04 | G0–G3
1859474823444598144 | 20:51:02.39 |  $+$30:44:36.83 | G074.0$-$08.5 | $3.7\pm 4.1$ | 15.67 | K2–K3.5
1859471185617972864 | 20:51:03.14 |  $+$30:41:42.71 | G074.0$-$08.5 | $3.5\pm 4.3$ | 16.27 | K2.5–K3.5
1859471872812767104 | 20:51:04.07 |  $+$30:43:47.86 | G074.0$-$08.5 | $1.9\pm 4.1$ | 15.72 | K1.5–K3
1859471219977693056 | 20:51:09.28 |  $+$30:42:10.58 | G074.0$-$08.5 | $1.7\pm 4.2$ | 16.88 | K5–M0
1858715683695452800 | 20:51:10.06 |  $+$30:26:53.64 | G074.0$-$08.5 | $1.7\pm 6.8$ | 13.01 | G2–K0
1858719291463059840 | 20:51:17.37 |  $+$30:36:47.21 | G074.0$-$08.5 | $1.7\pm 3.5$ | 16.95 | F2–K6
1858719085309596672 | 20:51:30.74 |  $+$30:36:47.96 | G074.0$-$08.5 | $4.0\pm 3.6$ | 13.84 | G9–K3.5
1859572714338261504 | 20:51:35.72 |  $+$30:58:15.42 | G074.0$-$08.5 | $6.8\pm 6.5$ | 15.02 | K3–K6
2069667129926577921 | 05:00:31.29 |  $+$46:33:27.25 | G160.9$+$02.6 |  $8\pm 14$ | 11.07 | A5–F0
2069187811576350722 | 05:00:53.98 |  $+$46:30:00.61 | G160.9$+$02.6 | $10\pm 17$ | 11.40 | A6–F0
206994059046621824 | 05:00:55.07 |  $+$46:40:37.90 | G160.9$+$02.6 | $1.7\pm 3.0$ | 16.27 | K5.5–K7
2069382631292887043 | 05:01:10.87 |  $+$46:27:38.11 | G160.9$+$02.6 | $12\pm 17$ | 12.33 | A7–F0
2069425194375804164 | 05:01:16.00 |  $+$46:33:21.68 | G160.9$+$02.6 |  $7\pm 16$ | 10.87 | A7–A9
34424902642611745285 | 05:39:44.40 |  $+$27:46:51.19 | G180.0$-$01.7 | $4.8\pm 7.0$ | 8.91 | B0.5
31334623164327712006 | 06:38:50.77 |  $+$06:35:26.03 | G205.5$+$00.5 | $12\pm 26$ | 12.01 | B8–A0
31334120180761625607 | 06:39:07.47 |  $+$06:27:37.99 | G205.5$+$00.5 |  $3\pm 22$ | 11.79 | B8–A0
31334860933717186568 | 06:39:13.09 |  $+$06:37:53.98 | G205.5$+$00.5 | $10\pm 25$ | 10.05 | B4–B6
31334071046336005129 | 06:39:20.12 |  $+$06:21:00.47 | G205.5$+$00.5 | $11\pm 22$ | 11.80 | B8–A0
313341449197726681610 | 06:39:24.25 |  $+$06:33:08.53 | G205.5$+$00.5 |  $7\pm 20$ | 12.25 | B7–B9
5526325078219949440 | 08:22:00.32 |  $-$42:59:34.28 | G260.4$-$03.4 | $1.6\pm 1.8$ | 13.58 | G2–G9
5526328136236654208 | 08:22:00.50 |  $-$42:57:26.84 | G260.4$-$03.4 | $2.1\pm 2.7$ | 13.17 | K1–K3
5526325112579690496 | 08:22:05.08 |  $-$42:58:43.34 | G260.4$-$03.4 | $0.9\pm 2.5$ | 16.86 | K7–M0
5526324803342055168 | 08:22:05.62 |  $-$43:01:23.19 | G260.4$-$03.4 | $1.8\pm 2.8$ | 16.92 | M1–M2
5526324906421263744 | 08:22:06.33 |  $-$42:59:23.89 | G260.4$-$03.4 | $0.5\pm 2.1$ | 15.47 | K3.5–K5.5
5526325284378390784 | 08:22:10.14 |  $-$42:59:29.27 | G260.4$-$03.4 | $0.3\pm 2.5$ | 14.87 | K5.5–K7
5526324833397363840 | 08:22:10.47 |  $-$43:00:49.73 | G260.4$-$03.4 | $1.2\pm 2.8$ | 15.99 | K5.5–M0
5526325250018655616 | 08:22:15.57 |  $-$42:58:31.52 | G260.4$-$03.4 | $1.7\pm 2.4$ | 13.56 | G9–K3
5526325353097874176 | 08:22:17.60 |  $-$42:58:40.09 | G260.4$-$03.4 | $1.9\pm 2.2$ | 13.00 | F2–F9
5521967782363042432 | 08:35:48.51 |  $-$45:18:46.12 | G263.9$-$03.3 | $5.1\pm 5.4$ | 15.85 | K3.5–M1.5
5329655463417782528 | 08:51:46.03 |  $-$46:21:01.11 | G266.2$-$01.2 | $2.1\pm 2.9$ | 16.38 | K4.5–K7
5329655257259343232 | 08:51:47.10 |  $-$46:21:17.64 | G266.2$-$01.2 | $2.6\pm 3.8$ | 16.87 | K6–K9

  * •

1TYC3344-235-1; 2TYC3344-679-1; 3TYC3344-683-1; 4TYC3344-553-1; 5HD 37424;
6TYC159-2771-1; 7HD 261359; 8HD 261393; 9TYC159-2337-1; 10TYC159-2671-1

Table 2: Continued

Gaia DR2 | RA [h:m:s] | DEC [d:m:s] | SNR |  $\rho$ [arcmin] |  $G$ [mag] | SpT
---|---|---|---|---|---|---
5329655845686016256 | 08:51:50.00 |  $-$46:20:20.25 | G266.2$-$01.2 | $1.7\pm 3.4$ | 15.13 | K4–K5.5
5329655983110009856 | 08:51:53.68 |  $-$46:18:19.05 | G266.2$-$01.2 | $1.7\pm 3.9$ | 16.57 | K6–K9
5329655776966532864 | 08:51:54.73 |  $-$46:19:37.86 | G266.2$-$01.2 | $0.7\pm 2.8$ | 15.24 | K2.5–K3.5
5329657563672287872 | 08:51:55.40 |  $-$46:16:42.01 | G266.2$-$01.2 | $2.8\pm 4.3$ | 16.26 | K5–K8
5329656017469763456 | 08:51:59.13 |  $-$46:18:09.17 | G266.2$-$01.2 | $1.0\pm 4.1$ | 15.30 | K2–K4
5329655811325638272 | 08:51:59.71 |  $-$46:19:01.93 | G266.2$-$01.2 | $0.7\pm 4.1$ | 14.91 | G8–K2
5329657288780088704 | 08:52:00.36 |  $-$46:17:57.93 | G266.2$-$01.2 | $1.8\pm 4.3$ | 16.62 | K2.5–K6.5
5329654849253600896 | 08:52:00.56 |  $-$46:21:49.15 | G266.2$-$01.2 | $2.9\pm 4.3$ | 14.97 | K1.5–K5
5329654849241704320 | 08:52:00.64 |  $-$46:21:48.85 | G266.2$-$01.2 | $2.5\pm 4.3$ | 16.82 | –
532965749063837427211 | 08:52:02.45 |  $-$46:17:19.84 | G266.2$-$01.2 | $2.2\pm 4.1$ | 7.85 | B8 IV,V
5329654883601436032 | 08:52:05.09 |  $-$46:22:19.13 | G266.2$-$01.2 | $2.5\pm 4.0$ | 14.91 | K0.5–K2.5
6005336183671977344 | 15:09:15.49 |  $-$40:00:55.10 | G330.0$+$15.0 | $1.1\pm 4.4$ | 16.31 | K7–M0
6005322405417134848 | 15:09:22.17 |  $-$40:17:50.46 | G330.0$+$15.0 | $4.3\pm 6.9$ | 15.95 | K3.5–M1
6005309932830763648 | 15:09:27.22 |  $-$40:24:26.21 | G330.0$+$15.0 | $1.7\pm 4.1$ | 15.70 | K9–M2.5
5972266305587974656 | 17:13:45.20 |  $-$39:46:11.77 | G347.3$-$00.5 | $1.3\pm 2.5$ | 14.71 | K1.5–K3.5
5972266443017191680 | 17:13:47.23 |  $-$39:44:36.86 | G347.3$-$00.5 | $0.7\pm 2.4$ | 15.97 | K5.5–M0.5
5972266683540067328 | 17:13:47.48 |  $-$39:43:14.68 | G347.3$-$00.5 | $1.9\pm 2.8$ | 16.84 | K6–M0.5
5972266443017192960 | 17:13:47.51 |  $-$39:44:29.27 | G347.3$-$00.5 | $0.7\pm 2.6$ | 16.65 | K6–M1.5
5972266447337050496 | 17:13:49.35 |  $-$39:44:01.38 | G347.3$-$00.5 | $1.1\pm 2.9$ | 16.98 | K6–M1
5972266447374127872 | 17:13:49.64 |  $-$39:44:03.04 | G347.3$-$00.5 | $1.0\pm 2.9$ | 15.01 | K4–K6
5972266378654585856 | 17:13:50.45 |  $-$39:45:35.03 | G347.3$-$00.5 | $0.5\pm 2.8$ | 16.59 | K6.5–K9
5972266477376942848 | 17:13:51.06 |  $-$39:43:28.12 | G347.3$-$00.5 | $1.8\pm 2.9$ | 14.81 | K1–K4.5
5972219340115696512 | 17:13:51.18 |  $-$39:46:39.54 | G347.3$-$00.5 | $1.7\pm 2.9$ | 16.54 | –
5972266378617336832 | 17:13:51.38 |  $-$39:45:25.30 | G347.3$-$00.5 | $0.5\pm 2.6$ | 16.40 | K3.5–K6.5
5972266481696795008 | 17:13:52.37 |  $-$39:43:28.21 | G347.3$-$00.5 | $1.6\pm 2.8$ | 16.92 | K3.5–K9
5972266408657455872 | 17:13:53.26 |  $-$39:44:22.17 | G347.3$-$00.5 | $0.9\pm 2.5$ | 16.88 | K5–M1
5972266413014335616 | 17:13:53.46 |  $-$39:44:41.28 | G347.3$-$00.5 | $0.7\pm 2.1$ | 16.59 | K7–M2.5
5972219374522872576 | 17:13:54.47 |  $-$39:46:45.27 | G347.3$-$00.5 | $1.9\pm 2.7$ | 15.04 | K4.5–K6
5972219477602131840 | 17:13:55.22 |  $-$39:45:21.84 | G347.3$-$00.5 | $1.0\pm 2.0$ | 13.47 | G8–K1.5
5972266413014326528 | 17:13:55.45 |  $-$39:44:36.52 | G347.3$-$00.5 | $1.3\pm 2.2$ | 15.62 | K5–K9
5972267886141988864 | 17:13:56.10 |  $-$39:44:03.01 | G347.3$-$00.5 | $1.6\pm 2.4$ | 13.94 | M3.5
5972267890445877120 | 17:13:56.96 |  $-$39:43:57.12 | G347.3$-$00.5 | $1.8\pm 2.4$ | 16.91 | K4.5–K9

  * •

11HD 76060

## Author Biography

Oliver Lux. The author obtained a Bachelor of Science in Physics and a Master
of Science in Astrophysics at the Universities of Bochum and Bonn,
respectively, gaining knowledge about neutron stars. In his current work as a
PhD student at the University of Jena, he works on runaway stars and supernova
remnants. Further research interests are variable stars and exoplanets.