OB associations and their origins

Nicholas J. Wright
Astrophysics Group, Keele University, Keele, ST5 5BG, UK
OB associations are unbound groups of young stars made prominent by their
bright OB members, and have long been thought to be the expanded remnants of
dense star clusters. They have been important in astrophysics for over a
century thanks to their luminous massive stars, though their low-mass members
have not been well studied until the last couple of decades. This has changed
thanks to data from X-ray observations, spectroscopic surveys and astrometry
from Gaia that allows their full stellar content to be identified and their
dynamics to be studied, which in turn is leading to changes in our
understanding of these systems and their origins, with the old picture of
Blaauw (1964a) now being superseded. It is clear now that OB associations have
considerably more substructure than once envisioned, both spatially,
kinematically and temporally. These changes have implications for the star
formation process, the formation and evolution of planetary systems, and the
build-up of stellar populations across galaxies.

Keywords: OB associations, star clusters, young stars, star formation, stellar
kinematics and dynamics.

##  1 Introduction

OB associations are gravitationally unbound groups of young stars, typically
containing many prominent OB stars as well as numerous low-mass stars. Their
low space densities ($<0.1$ M⊙ pc-3) make them dynamically unstable to
Galactic tidal forces and therefore over time they should disperse. The fact
that they exhibit some spatial and kinematic concentration (despite being
unbound) and contain short-lived OB stars implies that they must be young, a
fact first realised by Ambartsumian (1947), which provided the first evidence
that star formation was still ongoing in the Galaxy. OB associations continue
to be a valuable tracer of the distribution of young stars and the star
formation process. Their dimensions can range from a few to a few hundred
parsecs, while their total stellar masses cover a thousand to tens of
thousands of solar masses (systems less massive than this typically don’t
contain any OB stars and are thus known as T associations after the T-Tauri
stars that are often used to identify them).

OB associations are important objects to study because they represent a
transitional phase between the birth environment of stars (star forming
regions or star clusters) and the field star population in galaxies. They are
therefore useful for studying the star formation process (as they probe the
spatial and kinematic configuration of stars at birth) and for understanding
how and why young star clusters disperse and how the field star population is
built up. There are debates as to whether most stars form in dense clusters
that are then disrupted and briefly visible as dispersing OB associations
(Lada & Lada, 2003; Kroupa, 2011) or whether a large fraction of stars are
born in unbound, low-density groups like OB associations (Miller & Scalo,
1978; Kruijssen, 2012).

Historically OB associations were vital for identifying groups of OB stars and
calibrating their luminosity scale (e.g., Morgan et al., 1953; Humphreys,
1978). In addition OB associations provide large samples of unobscured young
stars that are useful for studies of the initial mass function, the frequency
and properties of multiple systems, protoplanetary disks and planetary systems
(e.g., Massey et al., 1995; Kouwenhoven et al., 2007; Kalas et al., 2015).
They also allow us to trace star formation over large areas and timescales and
thus study the propagation of star formation and the role of feedback in
triggering or halting star formation.

###  1.1 Historical Summary

The term “association” was coined by Ambartsumian (1947) to describe the low-
density groups of O- and B-type stars that had been known about for many years
(e.g., Kapteyn, 1914; Eddington, 1914; Pannekoek, 1929), placing them
alongside open and globular clusters as a new type of stellar grouping. The
term association would later be separated in OB- and T-associations based on
their most-prominent members (Ambartsumian, 1968). At the time it was thought
that only OB stars formed (and existed within) OB associations, with T-Tauri
stars restricted to T-associations and open clusters. This view was challenged
by Miller & Scalo (1978) who argued that unless OB associations also included
low-mass stars then there were insufficient numbers of birth sites for low-
mass stars to explain the observed field star population.

Concerning the stability of OB associations, Bok (1934) had already shown that
low-density systems (space densities $<0.1$ M⊙ pc-3) were unstable against
disruption by galactic tidal forces, which then led Ambartsumian (1947) to
realise that this meant such system must be young and in the process of
expanding. This implied that OB associations and the stars within them had
only recently formed, which provided one of the first pieces of evidence that
stars were still forming within the galaxy today, long before the discovery of
molecular clouds (Ambartsumian, 1947).

Since OB associations are gravitationally unbound, it was reasonable to assume
that they were in the process of expanding. Ambartsumian (1949) estimated a
typical expansion velocity of $\sim$5 km s-1 for OB associations, based on the
balance between their initial expansion and the Galactic tidal forces that
would be needed to produce the slightly flattened morphology observed in some
associations at the time. A similar view was suggested by Blaauw (1956) who
argued that the stars in associations “may have originated within a much
smaller volume than that occupied by the group at the present”. Observations
showing that the degree of proximity to the interstellar medium decreased with
the increasing size of associations, and that the more dispersed parts of
associations had older OB stars within them, solidified the expansion model
(Blaauw, 1964a).

The origin of this expansion remained unclear. Opik (1953) and Oort (1954)
suggested the idea that an unbound group of stars must form from an unbound
cloud of gas and thus be born in a state of expansion. An alternative
formation model was first put forward by Tutukov (1978) in which a
gravitationally bound cluster of stars, embedded within a molecular cloud,
becomes unbound and begins to expand when the gas is dispersed by feedback
processes. This idea was consistent with the growing observational evidence
that many giant molecular clouds were gravitationally bound (e.g., Duerr et
al., 1982) and that O-type stars can efficiently disperse the residual gas in
molecular clouds (e.g., Whitworth, 1979). Later studies, including numerous
N-body simulations of the reaction of the stellar system to gas expulsion,
reinforced this idea (Lada, 1987; Goodwin, 1997; Kroupa et al., 2001;
Baumgardt & Kroupa, 2007).

The next major advance in the study of OB associations was brought about by
data from the astrometric satellite Hipparcos (Perryman et al., 1997), thanks
in part to the work of the SPECTER consortium to get many thousands of
candidate members of known OB associations included in the Hipparcos input
catalogue. These data lead to a landmark census of the nearby OB associations
presented by de Zeeuw et al. (1999), which included significant improvements
in the definition, membership, and characterisation of these associations.
This work, and others driven by Hipparcos data, has since lead to many
advances in the study of OB associations, including the first studies of the
3D structure of OB associations (de Bruijne, 1999b), the identification of
runaway OB stars ejected from associations (Hoogerwerf et al., 2000), the
discovery of new OB associations (de Zeeuw et al., 1999) and studies of
Galactic structure derived from OB association kinematics (Mel’Nik & Dambis,
2009).

One of the major limitations of Hipparcos for the study of OB associations was
its depth, confining samples to only the brightest (and most massive) OBA-type
members of even the closest associations. Follow-up surveys to identify the
less-massive (and more numerous) members of these associations has been
possible, but has, in general, been hindered by the large areas over which OB
associations span on the sky and the difficulty of separating young, low-mass
association members from field stars (Brown et al., 1999; Briceño et al.,
2007). In this regard, astrometry from Gaia, the successor to Hipparcos has
quickly become invaluable for the identification and study of huge numbers of
low-, intermediate- and high-mass stars in OB associations. It is against this
backdrop that this review is presented, with the study of OB associations at
the cusp of a revolution in our understanding of these important systems.

###  1.2 Scope of this review

The purpose of this review is to summarise our current knowledge of the
properties and origin of OB associations, focussing mainly on observational
results, but considering recent ideas from theory and simulations where
relevant. In particular I focus on work from the past two decades since the
last major reviews on OB associations (Blaauw, 1991; Brown et al., 1999). This
period includes work that has built up on the shoulders of the Hipparcos
studies and leads into early work with Gaia.

This review is limited to OB associations in the Milky Way, because while
similar systems have been identified and studied in other galaxies (e.g., van
den Bergh, 1964; Efremov et al., 1987), with NGC 604 worthy of note as the
prototype extragalactic OB association (Maíz-Apellániz et al., 2004), the
tools used to study these systems and the results derived from their study are
necessarily different. For a recent review of these systems we refer the
reader to Gouliermis (2018).

Section 2 outlines the various methods used to identify OB association
members, which are necessarily different and more complex to those used to
identify the members of compact star clusters. Section 3 summarises the
properties of some of the most prominent and well-studied OB associations in
the Solar vicinity, while Section 4 includes a discussion of these properties
in a broader sense, taking into account large-scale studies of multiple
associations. In Section 5 these results are discussed and their implications
for our understanding of the origin, formation and evolution of OB
associations considered. Finally, a summary is provided in Section 6 with some
concluding remarks.

##  2 Observing OB associations and identifying members

One of the obstacles in studying OB associations is the difficulty of reliably
identifying their members. While high-mass stars can be identified just from
photometry, young low-mass stars can be difficult to distinguish from the
older (and more numerous) field stars. When studying star clusters the density
of members on the sky is significantly higher than the background density that
members can be reasonably effectively identified simply from photometry, but
the stellar density in OB associations is sufficiently low that this is not
possible. The fact that the stars in OB associations are typically older than
stars in star forming regions means that the use of H$\alpha$ excess emission
or infrared excesses to identify members is inefficient and could certainly
lead to an age bias in the resulting sample.

The key methods used for the last few decades have included X-ray observations
or multi-object spectroscopy, though recently it has been possible to identify
members based on the enhanced luminosity of pre-MS stars using Gaia
parallaxes. Even when young stars are identified and confirmed, assigning them
into different OB associations or groups is still a challenging task requiring
spatial and kinematic information. This section summarises some of the
different techniques used to identify young stars, confirm their youth, and
assign them into different groups or associations.

###  2.1 Identifying young stars

The first task in identifying members of OB associations is to efficiently
select candidate young stars. For high-mass stars this can usually be done
simply with photometry, but for less massive stars that can photometrically
appear very similar to their older counterparts it requires some indicator of
youth such as elevated luminosity, enhanced X-ray emission, or rapid rotation.

####  2.1.1 Photometric identification of O and B-type members

Figure 1: Photometric selection of OB stars. The left panel shows the original
$(U-B,B-V)$ colour-colour diagram from Johnson & Morgan (1953) with O- and
B-type stars shown as open squares and filled circles respectively. An
unreddened main sequence can be traced out following the B-type stars while
reddening vectors are shown extending to redder colours. The right panel shows
a more modern representation of this method in the $(u-g,g-r)$ colour-colour
diagram from Mohr-Smith et al. (2015). The unreddened main sequence can be
seen extending to later spectral types, with reddening curves (for
$R_{V}=3.8$) extending from spectral types B1V, B3V, G0V, and an ideal
Rayleigh-Jeans tail spectrum. Grey points show all stars in the field of view,
blue crosses show candidate OB stars selected from this method, and red
triangles show known OB stars in the field of view.

The process for photometrically identifying OB stars was developed and
pioneered by Johnson & Morgan (1953) in their landmark paper on the $UBV$
photometric system. They showed that field OB stars could be identified (and
their reddening determined) using the $Q$ method in the $U-B$ versus $B-V$
colour-colour diagram (see Figure 1). This method is effective because in this
diagram early-type stars do not redden onto themselves (or onto late-type
stars) and thus trace out a unique region of the colour-colour diagram (in
part due to the response of the $U-B$ colour to the strength of the Balmer
break in hot stars). The method is quite reliable and contamination rates are
low (e.g., Mohr-Smith et al. 2017 show that when using modern photometry the
method is reliable 97% of the time) and as such validation of photometrically-
identified candidates is useful, but not necessary.

This work, combined with the absolute-magnitude calibration of B-type stars
(Johnson & Morgan, 1953), allowed the identification of OB stars in open
clusters and associations and the determination of distances to these systems.
This allowed numerous authors to identify groups of OB stars across the sky
and thus the first detailed cataloguing of OB association high-mass members
(e.g., Johnson & Morgan, 1954; Harris, 1956; Blaauw, 1956, 1964a; Ruprecht,
1966). These membership lists were instrumental in determining the fundamental
properties of these stars (e.g., Humphreys, 1978) and estimating turn-off ages
(e.g., de Geus et al., 1989) \- for more details see Section 4.5.

While photometric identification of field OB stars is most effective in the
blue part of the optical spectrum, within the Galactic plane high extinction
can make this approach difficult. In regions of high extinction identifying OB
stars in the near-IR is sometimes necessary. For example, Comerón et al.
(2002) used a combination of near-IR colour-magnitude and colour-colour
diagrams to identify 46 new candidate OB stars in Cygnus OB2, significantly
increasing the census of such objects in the association.

The availability of wide-field, deep and multi-band optical photometry across
the Galactic plane (Drew et al., 2005, 2014) allowed Mohr-Smith et al. (2015)
to extend the $Q$ method of Johnson & Morgan (1953) to the large datasets of
the modern survey era. They adapted the standard $Q$ method to the Sloan
filter system, selecting OB stars in the $u-g$ versus $g-r$ colour-colour
diagram (see Figure 1) and used follow-up spectroscopy to confirm their method
was successful 97% of the time (Mohr-Smith et al., 2017). In addition to
identifying large numbers of early-type stars across Carina, these authors
have also shown how this method can be effectively used to identify new OB
associations (Mohr-Smith et al., 2017; Drew et al., 2018).

####  2.1.2 X-ray identification of FGKM members

Young stars are highly luminous X-ray sources, typically hundreds to thousands
of times more luminous than older, main-sequence stars (Preibisch & Feigelson,
2005). For solar and late-type stars these X-rays are emitted from a
magnetically-confined plasma at temperatures of millions of Kelvin known as a
corona. The corona is heated by the stellar dynamo, which itself is driven by
stellar rotation, as evidenced by the strong relationship between rotation and
X-ray activity (Wright et al., 2011). Since stars spin down as they age they
also gradually become less X-ray luminous, with X-ray to bolometric luminosity
ratios dropping from $\sim$10-3 for pre-main sequence stars to
$10^{-8}$–$10^{-4}$ for older main-sequence stars (e.g., Feigelson et al.,
2004; Wright et al., 2010b). This significant X-ray luminosity difference
between young stars and older field stars makes X-rays a powerful tool for
identifying young stars.

Observations from the Einstein Observatory and the ROSAT All-Sky Survey
uncovered many thousands of X-ray sources associated with young stars in star
forming regions and associations (e.g., Walter et al., 1988; Neuhaeuser et
al., 1995; Preibisch et al., 1998; Webb et al., 1999; Naylor & Fabian, 1999;
Pozzo et al., 2000), though the limited sensitivity and low spatial resolution
prevented it from being used to identify large samples of low-mass association
members. These limitations were overcome following the launch of the Chandra
and XMM-Newton X-ray observatories, but the small fields of view and high
demand for observing time with these facilities limited their application to
the wide areas needed for fully observing most OB associations (though compact
observations of parts of OB associations were still possible, e.g., Argiroffi
et al., 2006; Getman et al., 2006; Wright & Drake, 2009). For more distant and
compact OB associations, such as Cyg OB2, tiled surveys do facilitate coverage
of the entire association (e.g., Wright et al., 2014a). Upcoming all-sky data
from eROSITA (Merloni et al., 2020) will allow members of nearby associations
to be identified, though this will mostly be limited to the brighter solar-
type stars in the closest ($d<500$ pc) associations.

While X-ray observations are very efficient for identifying young stars, there
are also many other types of object that are luminous in X-rays and that can
act as contaminants in studies of young clusters or associations. These can
include foreground main-sequence stars, background giants, X-ray binaries and
cataclysmic variables, as well as extragalactic sources such as quasars. The
frequency of these different contaminants can depend on the sightline towards
the region targeted as well as the background galactic population and any
obscuring material that might limit the level of extragalactic contamination.
Getman et al. (2011) show that many of these contaminants can be identified
either from their significantly harder X-ray spectra or from the lack of an
optical counterpart, particularly for extragalactic contaminants. However
spectroscopic follow-up is still useful for identifying stellar contaminants
such as foreground or background sources.

####  2.1.3 Luminosity identification of association members

Figure 2: Luminosity identification of young stars in the $M_{G}$ vs.
$G_{\mathrm{BP}}-G_{\mathrm{RP}}$ absolute colour-magnitude diagram. An
unreddened main sequence is visible across the diagram with a pre-main
sequence clearly visible above it. Figure adapted from Zari et al. (2019).

Pre-main sequence stars are more luminous than main-sequence stars and
therefore in a colour – absolute magnitude diagram they stand out above both
the single-star and binary star sequences (see Figure 2), allowing them to be
identified and separated. Until recently this method was not easy to utilise
due to the dearth of parallaxes for pre-MS stars (though it could be used for
individual star clusters where the distance was known, e.g., Hernández et al.,
2014), but following Gaia’s second data release in April 2018 this approach
can be used to identify most young stars ($<$10–20 Myrs, dependent on colour)
with reasonably high-precision parallaxes (roughly within 500 pc for Gaia DR2
data).

Zari et al. (2018) used this method with the Gaia $M_{G}$ vs. $G_{BP}-G_{RP}$
diagram (Figure 2) to identify stars younger than 20 Myrs within 500 pc of the
Sun, mapping the 3D distribution of the known YSO populations in the solar
neighbourhood. Damiani et al. (2019) combined this technique with a proper
motion selection to map out the stellar population of the Scorpius-Centaurus
OB association. This approach can reduce the level of contamination,
particularly for groups of stars with distinct proper motions relative to the
field star population (such as Sco-Cen), though it could also introduce a
kinematic bias that should be considered if the sample is to be used for any
structural or kinematic study. Recently Zari et al. (2019) used this method to
select members of the Orion OB1 association and map out its 3D structure.

This method does require a careful correction for interstellar extinction to
produce an accurate colour – absolute magnitude diagram (e.g., Zari et al.,
2018), and can introduce biases to the sample, such as an age–mass bias (since
low-mass pre-MS stars are more over-luminous at a given age than more massive
stars, and therefore easier to identify) and both distance and magnitude
biases (since a cut is often placed on the relative parallax error,
$\sigma_{\varpi}/\varpi$, which typically increases with increasing distance
and apparent magnitude). There is also the risk of contamination from binaries
or from field stars with uncertain parallaxes, though such contamination is
expected to be minor and spectroscopic verification may only be necessary for
pre-MS intermediate-mass stars that are less offset from the main sequence
than their lower-mass siblings. Despite this, the all-sky availability of Gaia
data and the improved precision expected from future data releases suggests
this method is likely to be the most effective for stars within (at least) 1
kpc in the future.

####  2.1.4 Other methods

In addition to the methods discussed above, there are a number of other
techniques that can or have been used to identify young stars in OB
associations. Historically one of the most common methods has been the use of
objective prism surveys or H$\alpha$ photometry to identify young stars from
the excess H$\alpha$ emission due to accretion or chromospheric activity
(e.g., Liu et al., 1981; Wiramihardja et al., 1989; Mikami & Ogura, 2001). The
recent H$\alpha$ Galactic plane surveys IPHAS and VPHAS+ (Drew et al., 2005,
2014) have facilitated deep, large-area surveys for young stars with the
H$\alpha$ excess method, with recent applications in Cyg OB2 and Cep OB2 (Vink
et al., 2008; Barentsen et al., 2011). However, this method has a strong bias
towards younger Classical T-Tauri stars, relative to older Weak-lined T-Tauri
stars, and therefore carries with it an age bias. There is also the risk of
contamination from foreground dMe stars (Briceño et al., 2001) and therefore
spectroscopic verification can be useful.

Infrared surveys are also effective for identifying young stars. Spitzer Space
Telescope photometry has been used to identify many thousands of young stars
in star forming regions and clusters (e.g., Povich et al., 2013), though the
limited time available hindered its application to wide-field studies of OB
associations outside of the Galactic plane or major star-forming complexes
(e.g., Guarcello et al., 2013). The availability of all-sky infrared images
from the Wide-Field Infrared Survey Explorer (WISE) has facilitated large-area
surveys for young stars in OB associations (e.g., Azimlu et al., 2015; Fischer
et al., 2016). Contamination from star-forming galaxies, asymptotic giant
branch stars and shock emission regions is possible however, though cuts in
magnitude or colour space can be used to reduce the contamination rates to
$\sim$10-20% (e.g., Fischer et al., 2016). Spectroscopic verification is
therefore useful.

Another method is to identify young stars based on their rapid rotation, since
stars spin-up as they form and then gradually spin-down over many Gyrs. Young
($\lesssim$50–100 Myrs) stars can therefore be identified from their rapid
($<$3–5 day, Patten & Simon, 1996) rotation periods. This has become feasible
in recent years due to the availability of time-domain, high-precision
photometric surveys designed to detect transiting extrasolar planets, e.g.,
Kepler, SuperWASP, and very recently, TESS (e.g., Binks et al. 2015 used this
method to identify several thousand rapidly-rotating young FGK stars from
SuperWASP data). Contamination from tidally-locked binaries can be significant
however, and so spectroscopic follow-up to confirm the youth of the stars is
necessary.

Multi-epoch photometry can also be used to identify young stars from their
inherent variability (e.g., Carpenter et al., 2001), though Briceño (2009)
find this method needs spectroscopic verification of the members and has about
a 50% success rate. Certain types of stars can be identified from single-epoch
photometry, for example Drew et al. (2008) used H$\alpha$ photometry to
identify young A-type stars and Damiani (2018) used optical and near-IR
photometry to identify pre-main sequence M-type stars. Despite this, single
epoch photometry is generally used to measure the surface density of young
stars above the galactic background level, rather than actually identifying
individual young stars (e.g., Alves & Bouy, 2012; Zari et al., 2017; Armstrong
et al., 2018).

###  2.2 Confirming the youth of stars

Once candidate young stars have been identified it is often necessary to
confirm their youth using other means. None of the methods mentioned above are
perfect and can lead to contamination from other objects, either stellar or
extra-galactic. Candidate young stars can be verified by obtaining multiple
indicators of youth from those listed above, particularly if the contaminants
of each method are mutually exclusive (e.g., most contaminants to infrared
surveys for young stars are likely to be in the background and so parallax
measurements can be effective to confirm such sources). However the most
effective approach is to obtain spectroscopy and use one or more spectroscopic
indicators of youth.

####  2.2.1 Lithium in young stars

Figure 3: Equivalent width of the Li 6708 Å feature versus $V-I$ colour for
stars towards the Gamma Velorum cluster in Vela OB2 from Jeffries et al.
(2014). The solid red lines show predictions from the models of Baraffe et al.
(1998) at 10 (upper) and 20 (lower) Myr, while the dashed blue line shows the
upper envelope for EW(Li) for IC 2391 and IC 2602 at $\simeq$50 Myr. Objects
whose lithium content is consistent with being a young star are shown in blue
and red, with all other stars in the field of view shown as black circles.

Lithium is probably the most effective indicator of youth in K and early
M-type stars because it is depleted during their pre-main sequence phase
(e.g., D’Antona & Mazzitelli, 1994; Soderblom, 2010), making it a commonly-
used spectroscopic indicator of youth (e.g., Preibisch et al., 1998; Jeffries
et al., 2014; Pecaut & Mamajek, 2016). Lithium is burnt at approximately 3
million K, a temperature that is reached in the core and at the bottom of the
convective zone during pre-main sequence evolution. For a short period of
time, low-mass stars deplete lithium in their cores whilst mixing Li-depleted
material to the surface of the star. This continues until the expansion of the
radiative core causes the temperature at the base of the convection zone to
drop below 3 million K, bringing lithium depletion to a halt. This point is
reached at a later time in lower-mass stars due to their slower pre-MS
evolution, and in stars with $M<0.6$ M⊙ total lithium depletion will
eventually occur (Baraffe et al., 1998).

The main Li feature in stellar spectra is the (unresolved) doublet at 6708 Å,
which is relatively easy to detect and measure, having an equivalent width of
$>$100 mÅ for young stars. The strength of this absorption line decreases
during the early life of a star, with a timescale for significant depletion of
$\sim$10–20 Myrs for mid-M stars, $\sim$100 Myrs for K-type stars, and up to
$\sim$1 Gyr for G-type stars. This makes the presence of a strong lithium
absorption line in stellar spectra a clear indicator of youth for K- and
M-type stars, but less effective for G-type stars. However, since older stars
up to the age of the Pleiades can exhibit Li absorption it is necessary to
define a threshold equivalent width as a function of colour or mass and use
this to select stars (though rotation and other factors may also determine the
lithium abundance in stars, e.g., Soderblom et al., 1993). Fortunately,
lithium equivalent widths have been measured in many young clusters, providing
a valuable calibration for this technique (e.g., Randich et al., 1997). Figure
3 shows lithium equivalent width as a function of colour for young stars in
Vela OB2 from Jeffries et al. (2014) showing how this can be used to identify
young stars.

There are very few types of contaminant for a sample of lithium-rich stars. A
small number of field stars may be young enough to have photospheric lithium
(1% for K-type stars and 10% for G-type stars, based on the depletion
timescales), but these objects are rare. Even rarer are lithium-rich field
giants, of which $\sim$1% of G/K giants are expected to have sufficient
lithium to reach most detection thresholds for young stars (Brown et al.,
1989).

####  2.2.2 Surface gravity indicators

Another valuable spectroscopic method for confirming the youth of stars is the
use of spectral lines that are dependent on surface gravity. Pre-main sequence
stars have surface gravities intermediate between those of main-sequence stars
and giants. Notable examples of such gravity-sensitive lines include the Na i
doublets at 5889, 5896 Å and 8183–8195 Å or the Ca i lines at 6102, 6122 and
6162 Å. Mohanty et al. (2004) present a selection of spectroscopic features in
M-type stars that can be used to infer the surface gravity to within 0.25 dex.
Some of these features are degenerate with effective temperature and therefore
a selection of features are needed to measure surface gravity, possible with
sufficiently high-resolution spectroscopy.

Damiani et al. (2014) introduce a gravity-sensitive index $\gamma$ as part of
the analysis of spectra from the Gaia-ESO Survey, based on various sets of
gravity-sensitive lines around 6490–6500 and 6760–6775 Å. With high-resolution
spectroscopy the index allows log $g$ to be inferred to a precision of 0.2
dex, sufficient to allow young stars to be separated from giants and main-
sequence stars at temperatures $\lesssim$ 5000 K.

####  2.2.3 Other spectroscopic indicators of youth

Other spectroscopic indicators of youth include emission lines due to strong
accretion in very young stars, such as He i 5876, 6678 Å or the Ca ii triplet
at 8500, 8544, and 8665 Å (e.g., Briceño et al., 2019) or Balmer lines in
emission from active late-type stars with strong chromospheric activity
(though this is also seen in field stars up to $\sim$1 Gyr old, Stauffer &
Hartmann, 1986).

###  2.3 Assigning young stars into groups

Once a sample of young stars have been identified and confirmed it is often
necessary to divide them into groups or identify contaminating young stars.
Even within a compact area of the sky there may be young stars that are not
part of the group under study, either young field stars or stars from other
nearby groups. Stars in clusters or associations have very similar space
motions and therefore members of a given group can be identified as a coherent
structure in velocity space, for which there are various different methods
that have been used over the last few decades that we will discuss here. It is
worth noting that the use of proper motions (or parallaxes for nearby systems)
to assign stars into groups can introduce a spatial or kinematic bias to the
sample, which should be considered when studying it.

####  2.3.1 Classical methods not requiring radial velocities

Figure 4: Division of young stars in Orion OB1 into subgroups from Kounkel et
al. (2018) using hierarchical clustering. The clustering algorithm was applied
to 10,248 stars and resulted in the identification of 190 groups across the
association, which were then manually combined into five larger structures
shown here: Orion A (green; known group), Orion B (orange; known group), Orion
C (cyan; new group), Orion D (red; new group) and $\lambda$ Ori (blue; known
group). In the left panel, black circles show the position of the major bright
stars in Orion with Orion’s Belt across the centre of the image.

The classical method for identifying members of a nearby moving group or
association is the convergent point method, which is based on the fact that
stars with a common space motion will have proper motions that appear to
converge from or to a point on the sky (depending on whether the group is
moving towards or away from us, e.g., Jones, 1971; de Bruijne, 1999a). The
convergent point for a group of stars is found by minimising the proper motion
perpendicular to the direction towards the convergent point (accounting for
measurement uncertainty and an intrinsic velocity dispersion). New members of
the group may be added to the group if their motion perpendicular to the
convergent point is below some threshold. This method has been valuable
because it only requires proper motions and sky positions, and not parallaxes
or radial velocities that are generally harder to obtain. This method was
extensively used by de Zeeuw et al. (1999) to perform a census of nearby OB
associations using Hipparcos data. They also used the “Spaghetti method” of
Hoogerwerf & Aguilar (1999), which uses parallaxes as well as positions and
proper motions, and identifies velocity cylinders or “spaghetti” that must
overlap in velocity space for all stars that are members of a group.

####  2.3.2 Modern methods using clustering algorithms

The availability of radial velocities for many Hipparcos stars has allowed the
membership of OB associations to be determined in 6-dimensional position and
velocity space. This can be achieved using a model for the spatial and
kinematic distribution of the group, such as a 6-dimensional Gaussian or a
mixture of Gaussians if one or more subgroups are known. Rizzuto et al. (2011)
applied such a method to re-establish the membership of Sco-Cen using linear
models in various dimensions of 6-dimensional space and calculating membership
probabilities for individual sources.

The availability of Gaia data for billions of stars provided 5-dimensional
(position, parallax and proper motions) data suitable for the application of
more advanced clustering algorithms. These include techniques such as
$k$-means clustering (where the user assumes a number of groups, but not their
positions, and the data is then iteratively divided into these groups),
mixture models (where the data is modelled as a collection of distributions –
usually multi-dimension Gaussians – and each source is assigned a probability
of belonging to a group based on this distribution), and hierarchical
clustering (where sources are divided into groups that have source-to-source
separations less than some threshold value). Many of these techniques are
unsupervised, meaning that they do not require pre-existing identification of
the clusters and involve a minimum of human intervention. Each of these
methods has its own advantages and disadvantages in identifying groups of
stars and rejecting non-clustered field stars, including whether or not the
number of groups must be assumed a-priori, whether the groups have to be
modelled in a certain way (e.g., as a collection of Gaussians), how to
determine the dividing line between different groups, and how incomplete data
is accounted for (such as the lack of radial velocities for many stars).

A review of the various methods can be found in Feigelson & Babu (2012),
though we highlight some of them and their recent applications here. Wilkinson
et al. (2018) apply the Density-Based Spatial Clustering of Applications with
Noise algorithm (DBSCAN, Simoudis et al., 1996, one of the most commonly-used
clustering algorithms in the literature) to Gaia DR1 data in Upper Sco,
increasing the known members by $\sim$50%. Cantat-Gaudin et al. (2019b) used
the classification scheme UPMASK (Unsupervised Photometric Membership
Assignment in Clusters, Krone-Martins & Moitinho, 2014), which uses principal
component analysis and $k$-means clustering to identify groups of stars,
finding 7 groups across Vela and Puppis with distinct ages and kinematics.

Table 1: List of known OB associations compiled from Ruprecht (1966), de Zeeuw et al. (1999) and Brown et al. (1999). The list is divided into well-studied associations, those that are clearly identified but have not been extensively studied, and those for which very little is known (including some which may not be real associations and are labelled with ”?”). Distances and ages are gathered from the literature. Positions are taken from Ruprecht (1966) and updated from de Zeeuw et al. (1999) where possible. Distances are taken from the source with the smallest uncertainty, preferring Gaia measurements or single-association studies where possible. For Sco-Cen and Orion OB1 the subgroups are listed instead of the entire associations due to their prominence (the Sco-Cen subgroups are also known as Sco OB2_2, OB2_3 and OB2_4). Name | l | b | Distance | Ref. | Age | Ref. | Clusters | Nebulae
---|---|---|---|---|---|---|---|---
|  $[$deg.$]$ |  $[$deg.$]$ |  $[$pc$]$ |  |  $[$Myr$]$ |  |  |
_Well-studied OB associations:_ |
Sco-Cen: US | 351.5 | +20.0 | $143\pm 6$ | W18 | $10\pm 7$ | P16 |  |
Sco-Cen: UCL | 331.0 | +12.5 | $136\pm 5$ | W18 | $16\pm 7$ | P16 |  |
Sco-Cen: LCC | 298.5 | +5.5 | $115\pm 4$ | W18 | $15\pm 6$ | P16 |  |
Ori OB1a | 201.0 | -17.3 |  $\sim$360 | K18 | 8–12 | B94 | 25 Ori |
Ori OB1b | 205.0 | -18.0 | 360–420 | K18 | 2–8 | B08 |  $\epsilon$ Ori |
Ori OB1c | 211.3 | -19.5 |  $\sim$385 | K18 | 2–6 | B08 | NGC 1980, NGC 1981 |
Ori OB1d | 209.0 | -19.5 |  $\sim$380 | Z17 | 1–2 | Z17 | Orion Nebula Cluster |
Vela OB2 | 262.8 | -7.7 | $411\pm 12$ | DZ 99 | 10–30 | CG19 | Gamma Vel, P Puppis |
Trumpler 10 | 262.8 | +0.7 | $372\pm 23$ | DZ99 | 45–50 | CG19 |  |
Cyg OB2 | 80.2 | +0.8 | 1350–1750 | B19 | 1–7 | W15 |  |
_High-confidence OB associations that have not been extensively studied:_ |
Ara OB1 | 338.0 | 0.0 |  $\sim$1100 | M73 |  $\sim$2 | A87 | NGC 6193 |
Aur OB1 | 173.1 | -1.6 |  $\sim$1060 | M17 | 11-22 | T10 | NGC 1912, NGC 1960 | Sh2-227
Aur OB2 | 173.0 | +0.1 |  $\sim$2420 | M17 |  $\sim$5.5 | T10 | NGC 1893, Stock 8 | IC 410, IC 417
Cam OB1 | 142.5 | +2.0 |  $\sim$800 | M17 | 7–14 | S85 | NGC 1502 | S202
Car OB1 | 286.5 | -0.5 | $2300\pm 50$ | S06 | 1–10 | D01 | NGC 3293, Tr 14–16 | NGC 3372
Car OB2 | 290.4 | +0.1 |  $\sim$1830 | M20 |  $\sim$4 | G94 |  |
Cas OB6 | 135.9 | +1.3 |  $\sim$1750 | M17 |  $\sim$4 | T10 | IC 1805, IC 1848 | W3, W4, W5
Cas-Tau | 169.0 | -16.5 | 125–300 | DZ99 |  $\sim$50 | DZ99 |  |
Cen OB1 | 304.2 | +1.4 |  $\sim$1920 | M17 | 6–12 | K94 | Stock 16 | RCW 75
Cep OB1 | 104.2 | -1.0 |  $\sim$2780 | M17 | 1–5 | C11 | NGC 7380 |
Cep OB2 | 102.1 | +4.6 |  $\sim$730 | M20 | 5 | DZ99 | Tr 37, NGC 7160 | IC 1396
Cep OB3 | 110.4 | +3.0 |  $\sim$700 | M17 | 5–8 | J96 | Cep OB3b | S155
Cep OB4 | 118.3 | +5.3 |  $\sim$660 | M17 | 1–6 | M68 | B59 | S171
Cep OB6 | 105.1 | +0.1 | $270\pm 12$ | DZ99 |  $\sim$50 | DZ99 |  |
CMa OB1 | 224.0 | -1.3 |  $\sim$1200 | Z20 | 1–10 | S18 | NGC 2353, NGC 2327 | IC 2177, Sh2-296
Collinder 121 | 235.7 | -10.0 | $543\pm 23$ | DZ99 | 5 | DZ99 | Collinder 121 | Sh2-306
Cyg OB1 | 75.5 | +1.1 |  $\sim$1460 | M17 | 6–8 | R08 | NGC 6913, IC 4996 |
Cyg OB3 | 72.9 | +1.9 |  $\sim$1830 | M17 | 2–12 | R08 | NGC 6871, NGC 6883 |
Cyg OB4 | 82.8 | -7.6 |  $\sim$800 | M17 |  $\sim$8.3 | U01 |  |
Cyg OB5 | 67.1 | +2.1 |  $\sim$1610 | R66 |  |  |  |
Cyg OB6 | 86.0 | +1.0 |  $\sim$1700 | R66 |  |  |  |
Cyg OB7 | 89.0 | 0.0 |  $\sim$630 | M17 | 1–13 | U01, W13 |  | Northern Coalsack
Cyg OB8 | 77.9 | +3.4 |  $\sim$1830 | M17 | 4–6 | M15 |  |
Cyg OB9 | 77.7 | +1.9 |  $\sim$960 | M17 | 2–4 | M15 |  | Sh2-108
Gem OB1 | 189.1 | +1.0 |  $\sim$1210 | M20 |  $\sim$9 | T10 | NGC 2175 | IC 443
Lac OB1 | 96.7 | -17.6 | $368\pm 17$ | DZ99 | 2–25 | C08 |  | S126
Mon OB1 | 202.1 | +1.0 |  $\sim$580 | M17 | 1–10 | F99 | NGC 2264 | NGC 2264, Mon. Ring
Mon OB2 | 206.3 | -2.1 |  $\sim$1210 | M17 | 2–15 | T76 | NGC 2244 | Rosette Neb.
Per OB2 | 159.2 | -17.1 | $296\pm 17$ | DZ99 | 1–10 | A15 | IC 348, NGC 1333 |
Per OB3 | 147.0 | -5.5 | $175\pm 3$ | DZ99 | 50 | DZ99 |  $\alpha$ Per |
Pup OB1 | 243.5 | +0.3 |  $\sim$2010 | M17 |  $\sim$4 | H72 | NGC 2467 | S311
Sco OB1 | 343.3 | +1.2 | $1560\pm 35$ | Y20 | 1–10 | D18 | NGC 6231, Tr 24 | G345.45+1.50,
|  |  |  |  |  |  |  | Gum 55, IC 4628
Sct OB2 | 23.1 | -0.4 |  $\sim$1600 | M09 |  $\sim$6 | S85 | NGC 6705 |
Ser OB1 | 16.7 | 0.0 |  $\sim$1530 | M17 | 8–13 | T10 | NGC 6611 | M16, M17
Ser OB2 | 18.2 | +1.7 |  $\sim$1600 | M17 |  $\sim$4.5 | T10 | NGC 6604 | Sh2-54
Sgr OB1 | 7.6 | -0.9 |  $\sim$1260 | M17 | 5–8 | S85 | NGC 6530, Col. 367 | M20
_OB associations for which very little is known:_ |
Aql OB1 (?) | 37.3 | -0.6 |  $\sim$2750 | R66 |  |  |  |
Cam OB3 | 147.0 | +3.0 |  $\sim$2650 | M17 |  $\sim$11 | T10 | Alicante 1 |
Cas OB1 | 124.1 | -1.4 |  $\sim$2010 | R66 |  $\sim$10 | L86 |  |
Cas OB2 | 112.0 | 0.0 |  $\sim$2100 | M17 |  $\sim$10 | L86 |  | NGC 7538, Sh1-57
Cas OB4 | 120.1 | -0.3 |  $\sim$2300 | M17 |  $\sim$8 | L86 |  |
Cas OB5 | 116.2 | -0.5 |  $\sim$2010 | M17 | 6–8 | S85 |  |
Cas OB7 | 122.8 | +1.2 |  $\sim$2010 | M17 |  $\sim$8 | L86 |  | Sh2-180
Cas OB8 | 129.2 | -1.1 |  $\sim$2300 | M17 |  $\sim$20 | T10 | NGC 581, NGC 663 |
Cas OB9 (?) | 113.5 | -2.5 |  $\sim$800 | R66 |  |  |  |
Cas OB10 (?) | 130.8 | -6.3 |  $\sim$3800 | R66 |  |  |  |
Cas OB14 | 120.4 | +0.7 |  $\sim$880 | M17 |  $<$10 | T10 |  |
Cen OB2 | 294.3 | -1.0 |  $\sim$2100 | R66 | 3–10 | B14 | IC 2944 | IC 2948
Cep OB5 | 108.5 | -2.8 |  $\sim$2090 | H78 |  $\sim$10 | L86 |  |
Cir OB1 | 315.5 | -2.8 |  $\sim$2010 | M17 |  |  | Pismis 20 |
Cor. Aust. (?) | 0.0 | -18.0 |  $\sim$130 | DZ99 |  |  | Coronet cluster | NGC 6726, NGC 6729
Cru OB1 (?) | 294.9 | -1.1 |  $\sim$2010 | M17 | 5–7 | K94 |  | IC 2944
Mon OB3 | 217.6 | -0.3 |  $\sim$2420 | M17 |  $\sim$7 | T10 |  | S287
Nor OB1 | 328.0 | -0.9 |  $\sim$2780 | M17 |  |  |  |
Ori OB2 (?) | 192.6 | -11.6 |  $\sim$3240 | R66 |  |  |  |
Per OB1 | 134.7 | -3.2 |  $\sim$1830 | M17 | 8–11 | M87 |  $h$ and $\chi$ Persei |
|  |  |  |  |  |  | (NGC 869 & 884) |
Pup OB2 (?) | 244.6 | 0.7 |  $\sim$3180 | M17 |  $\sim$2 | H72 | Ruprecht 20 |
Pup OB3 | 254.0 | 0.0 |  $\sim$1460 | M17 |  $\sim$4 | W63 |  | RCW 19
Sco OB4 (?) | 352.4 | +3.4 |  $\sim$960 | M17 |  |  |  | NGC 6334, Sh2-10
Sct OB3 | 17.3 | -0.8 |  $\sim$1330 | M17 |  $\sim$4.5 | T10 |  | Sh2-50
Sgr OB4 | 12.2 | -1.0 |  $\sim$1920 | M17 |  $<$10 | T10 |  |
Sgr OB5 | 359.9 | -1.2 |  $\sim$2420 | M20 | 6–12 | S85 |  | Sh2-15
Sgr OB6 | 14.2 | +1.2 |  $\sim$1600 | M17 |  $<$10 | T10 |  |
Sgr OB7 | 10.7 | -1.5 |  $\sim$1390 | M17 | 4–5 | T10 |  |
Vela OB1 | 265.0 | -0.7 |  $\sim$1460 | M17 |  $\sim$20 | T79 | NGC 2659 |
Vul OB1 | 60.3 | +0.1 |  $\sim$1600 | M17 | 10–16 | T10 | NGC 6823 | NGC 6820
Vul OB2 (?) | 64.7 | +1.8 |  $\sim$4130 | R66 |  |  |  |
Vul OB4 | 60.5 | +0.5 |  $\sim$800 | M17 |  $\sim$10 | T80 |  |
Table 1: _continued_

References: A15 (Azimlu et al., 2015), A87 (Arnal et al., 1987), B08 (Bally,
2008), B14 (Baume et al., 2014), B19 (Berlanas et al., 2019), B59 (Blaauw et
al., 1959), B94 (Brown et al., 1994), C08 (Chen & Lee, 2008), C11 (Chen et
al., 2011), CG19 (Cantat-Gaudin et al., 2019a), D01 (DeGioia-Eastwood et al.,
2001), D18 (Damiani, 2018), DZ99 (de Zeeuw et al., 1999), F99 (Flaccomio et
al., 1999), G94 (Garcia, 1994), H72 (Havlen, 1972), H78 (Humphreys, 1978), J96
(Jordi et al., 1996), K08 (Kaltcheva & Hilditch, 2000), K18 (Kounkel et al.,
2018), K94 (Kaltcheva & Georgiev, 1994), L86 (Lozinskaia et al., 1986), M09
(Mel’Nik & Dambis, 2009), M15 (Mahy et al., 2015), M17 (Mel’nik & Dambis,
2017), M20 (Melnik & Dambis, 2020), M68 (MacConnell, 1968), M73 (Moffat &
Vogt, 1973), M87 (Maeder, 1987), P16 (Pecaut & Mamajek, 2016), R08 (Reipurth &
Schneider, 2008), R66 (Ruprecht, 1966), S06 (Smith, 2006), S18 (Santos-Silva
et al., 2018), S85 (Schild & Maeder, 1985), S97 (Sung et al., 1997), T10
(Tetzlaff et al., 2010), T76 (Turner, 1976), T79 (Turner, 1979), T80 (Turner
et al., 1980), U01 (Uyanıker et al., 2001), vG84 (van Genderen et al., 1984),
W13 (Wolk et al., 2013), W15 (Wright et al., 2015), W18 (Wright & Mamajek,
2018), W63 (Westerlund, 1963), Y20 (Yalyalieva et al., 2020), Z17 (Zari et
al., 2017), Z20 (Zucker et al., 2020).

The rich young stellar population and significant substructure across the
Orion OB1 association has lead many authors to apply clustering algorithms to
trace its star formation history. Kounkel et al. (2018) use a form of
hierarchical clustering to identifying clusters in the Orion OB1 association
from 6-dimensional spatial and kinematic data (using radial velocities),
applying the algorithm to $\sim$10,000 stars and identifying 190 groups that
they then manually combine into 5 larger structures (see Figure 4). Chen et
al. (2019) used machine learning algorithms to identify stellar groups in
Orion using Gaia DR2 astrometry, finding 21 spatially- and kinematically-
coherent groups, only 12 of which were previously known or associated with
known clusters. Zari et al. (2019) used DBSCAN to identify 17 groups in Orion,
including foreground populations.

The wealth of new, multi-dimensional data and advanced statistical techniques
for identifying clusters will be critical in the future for identifying new OB
associations and star clusters, or ascertaining the membership of known
groups. However, detailed testing of these techniques in their application to
highly complex datasets with correlated uncertainties will be necessary to
fully understand their limitations as well as any biases they introduce to the
spatial or kinematic study of the resulting groups.

##  3 Properties of individual OB associations

Figure 5: Map of the most well-studied OB associations from Table
LABEL:list_of_associations projected on an inverted Galactic plane H$\alpha$
image (Finkbeiner, 2003).

The number of known OB associations runs into the dozens, though the majority
of these have been poorly studied. The most well-studied associations are
probably Scorpius OB2 (Scorpius-Centaurus, the nearest OB association to the
Sun), Orion OB1 (the nearest OB association with active high-mass star
formation within it), Vela OB2 (one of the associations with numerous well-
studied open clusters), and Cygnus OB2 (which while distant is probably the
most massive of the known OB associations). In this section I summarise the
properties of these associations, as well as some of the less well studied
associations, highlighting their features that can help us understand the
properties and formation of OB associations.

The cataloguing of OB associations began in the 1950s by various authors
motivated to group OB stars into clusters or associations to better constrain
their fundamental properties. By the 1960s various different lists of OB
associations existed (e.g., Morgan et al., 1953; Schmidt, 1958) with some
studies even beginning to subdivide these into subgroups (e.g., Blaauw,
1964a). Ruprecht (1966) was the first to gather these different lists into a
single list, codifying the nomenclature and defining their boundaries. The
previous numbering system using Roman numerals was changed to the one we are
familiar with today that uses Arabic numbers affixed to the letters OB and
their respective constellation (for example I Ori became Ori OB1). Table
LABEL:list_of_associations provides an updated list of known OB associations
with coordinates, distances, ages and associated clusters and nebulae gathered
from the literature. Figure 5 provides an observer’s view of the distribution
of OB associations on the sky in Galactic coordinates.

The high-mass membership of these associations was expanded and refined with
improved photometry over the following decades, leading to membership lists
that have been used for many decades (e.g., Humphreys, 1978; Garmany &
Stencel, 1992). The availability of Hipparcos astrometry further refined these
membership lists for nearby ($d<500$ pc) associations (de Zeeuw et al., 1999).
At the same time attempts have been made to refine the existing lists of OB
associations using either new clustering algorithms (Mel’Nik & Efremov, 1995)
or improved astrometry for members (de Zeeuw et al., 1999).

###  3.1 Scorpius-Centaurus

Figure 6: A wide-field view of the Sco-Cen association projected onto an
inverted H$\alpha$ image from Finkbeiner (2003). Blue circles show the
classical Hipparcos members of the association identified by Rizzuto et al.
(2011), with symbol size inversely proportional to their $V$ magnitude
(brighter stars have larger symbols). The runaway star $\zeta$ Oph is also
shown as a blue circle, with a Gaia DR2 proper motion vector covering 0.5 Myr
of motion extending from it. The red contours show the spatial distribution of
Gaia-selected pre-main sequence and upper main sequence stars from Damiani et
al. (2019). The grey dashed rectangles show the historical division (Blaauw,
1946) of the association into three subgroups from de Zeeuw et al. (1999).
Black diamonds show the IC 2602 cluster and the V1062 Sco moving group (Röser
et al., 2018), while empty black circles show the locations of the $\rho$
Ophiuchus star forming region and the Lupus I–V clouds.

The Scorpius-Centaurus (Sco-Cen or Scorpius OB2) association is the nearest OB
association to the Sun at a distance of 100–150 pc (de Zeeuw et al., 1999).
Due to its proximity and $>$100 pc size the association spans over 90∘ on the
sky, extending over the constellations Scorpius, Lupus, Norma, Centaurus,
Circinus, Musca, and Crux. The association contains $\sim$150 B-type stars and
(currently) a single O-type star, the runaway O9V star $\zeta$ Ophiuchi
(Blaauw, 1964a; de Geus et al., 1989; Brown, 1998; de Zeeuw et al., 1999;
Hoogerwerf et al., 2001), as well as a rich population of lower-mass stars
(Pecaut & Mamajek, 2016; Damiani et al., 2019). The total stellar content of
the association has been estimated as anywhere from 6,000 (Preibisch &
Mamajek, 2008) to 13,000 stars (Damiani et al., 2019), with a total stellar
mass of $\sim$4,000 M⊙ (Wright & Mamajek, 2018). Figure 6 shows the
distribution of the known high- and low-mass population of Sco-Cen, as well as
notable star forming regions and young clusters in their vicinity.

The association has historically been divided into three subgroups (Figure 6)
known as Upper Scorpius (US), Upper Centaurus-Lupus (UCL) and Lower Centaurus-
Crux (LCC). Of these US is the youngest, with a median age of 11 Myr, while
UCL and LCC are older with ages of 16 and 17 Myr respectively (Pecaut &
Mamajek, 2016). Historically US has been the most well-studied due to it being
the youngest, the most compact, and the furthest from the contaminating
effects of the Galactic Plane, though this is changing thanks to data from
Gaia.

The division of Sco-Cen into these subgroups was first defined by Blaauw
(1946), based on the positions and velocities of the brightest members, and
has since been used by many authors, including the study by de Zeeuw et al.
(1999) that defined the Hipparcos membership of the association. The revised
Hipparcos data allowed Rizzuto et al. (2011) to revisit the membership and
borders of the subgroups, favouring a more continuous distribution rather than
a division into subgroups. Recent temporal (Pecaut & Mamajek, 2016) and
kinematic (Wright & Mamajek, 2018) studies have supported this view that the
historical divisions do not necessarily reflect the real substructure of the
association (see also Villa Vélez et al., 2018; Damiani et al., 2019).

####  3.1.1 The high-mass stars

The association and its three subgroups have historically been delineated by
their bright B-type members, the first identification of which as a moving
group dates back to Kapteyn (1914), though was not fully confirmed until
Blaauw (1946). The modern census of high-mass members of Sco-Cen can be
attributed to de Geus et al. (1989), de Zeeuw et al. (1999), and since refined
by Rizzuto et al. (2011, see Figure 6). There are approximately 150 main-
sequence B-type stars across the association, the late O-type runaway star
$\zeta$ Ophiuchi (first suggested to have originated in Sco-Cen by Blaauw,
1952), and a number of evolved high-mass stars (such as the M1.5 supergiant
Antares, $\alpha$ Sco, and the B1.5III giant, $\alpha$ Lup) that define the
main-sequence turn-offs in each subgroup (de Geus et al., 1989). Many of the
LCC high-mass members have been debated (e.g., Hoogerwerf et al., 2000), with
the recent study by Rizzuto et al. (2011) adding the bright B-type stars
$\alpha$ Cru and $\beta$ Cru as members.

de Geus (1992) was the first to estimate the number of supernovae to have
exploded across parts of the association, estimating $6\pm 3$ past supernovae
in UCL, which Preibisch & Mamajek (2008) support with an estimate of $\sim$7
supernovae. Since most estimates put the total stellar content of LCC at just
over half that of UCL (e.g., Mamajek et al., 2002), this suggests $\sim$4 past
supernovae in this subgroup. For US there is one strong candidate for a past
supernova, that of the progenitor of the pulsar PSR J1932+1059, which
Hoogerwerf et al. (2001) suggest was an O5-O6 star that exploded 1.5 Myr ago
to leave this remnant and the runaway star $\zeta$ Ophiuchi. This explanation
for the origin of $\zeta$ Oph was questioned by Chatterjee et al. (2004) with
improved astrometric data. Neuhäuser et al. (2019) used Gaia DR2 astrometry to
kinematically connect $\zeta$ Oph and the pulsar PSR B1706-16, suggesting they
were ejected from a supernova in a binary system $1.78\pm 0.21$ Myr ago.
Tetzlaff et al. (2010) also identified the isolated neutron star RX J1856-3754
as probably originating from Upper-Sco $\sim$0.3 Myr ago, with an O-type
progenitor. Together these studies suggest the entire association might once
have contained $\sim$13 O-type (primary) stars.

####  3.1.2 Ages and age spreads

Figure 7: A spatially averaged ‘age map’ of Sco-Cen constructed from the ages
of 657 F/G/K/M-type pre-MS members of the association, with the luminous
early-type stars also shown. The three boxes delineate the classical
boundaries of the three subgroups US (left), UCL (centre) and LCC (right) that
have median ages of 10, 16 and 15 Myr, respectively. The very young regions of
$\rho$ Ophiuchus and the Lupus clouds, which are projected against Sco-Cen,
are masked over in this map. Figure from Pecaut & Mamajek (2016).

The Sco-Cen association exhibits a considerable range of stellar ages. The
youngest stars are typically associated with US, with stars as young as
$\sim$5 Myr at the northern end of the subgroup, though the lack of dense
molecular material or embedded stars in the association indicates that star
formation has certainly finished. The median age of US is approximately 10 Myr
(Pecaut & Mamajek, 2016), while UCL and LCC are older, with ages of 10–20 Myr
(de Geus et al., 1989; Mamajek et al., 2002; Pecaut & Mamajek, 2016).

The ages estimated for the three subgroups of Sco-Cen have generally increased
over time as evolutionary models have improved. For example, early estimates
put the age of US at approximately 5–6 Myr, as estimated from both the high-
(de Geus et al., 1989) and low-mass stars (Preibisch & Zinnecker, 1999;
Preibisch et al., 2002), and from kinematic traceback (Blaauw, 1978). Such
studies also favoured a very short age spread of, at most, 1–2 Myr, which lead
to suggestions that star formation across US had been triggered by some
external event (e.g., Preibisch & Zinnecker, 1999). More recent studies
suggest a median age of $\sim$10 Myr for US for both the early- and late-type
members of the association (Sartori et al., 2003; Pecaut et al., 2012; Pecaut
& Mamajek, 2016; Feiden, 2016). This change has followed a general increase in
pre-MS stellar ages derived from late-type stars (e.g., Bell et al., 2013)
that has been attributed to inaccuracies in the radii of pre-main-sequence
M-type stars due to either magnetic inhibition of convection (Feiden, 2016) or
starspots (Jackson & Jeffries, 2014). Sco-Cen has proved to be a valuable
calibrator for evolutionary models due to its proximity, rich low-mass
population, and observations of eclipsing multiple systems from Kepler/K2
(e.g., Kraus et al., 2015).

Figure 7 shows an age map of Sco-Cen from Pecaut & Mamajek (2016) that was
produced by spatially-averaging the ages of 657 pre-MS members of the
association calculated relative to an empirical isochrone. There are clear age
spreads of $\pm$6–7 Myr across each of the subgroups (Pecaut & Mamajek, 2016),
which are larger than previous measurements of the age spread (e.g., Preibisch
et al., 2002; Slesnick et al., 2006) due to a combination of the increased
median age (which increases all ages and age spreads) and the wider areas
studied by recent surveys compared to the small regions studied previously.
These age spreads appear to be real, as evidenced by the clear age
substructure seen in Figure 7 that suggests groups of stars exist within the
association with distinct ages relative to the age spreads across each
subgroup. This was confirmed by Damiani et al. (2019) who found that the more
clustered stars within the association were younger than the more distributed
population.

####  3.1.3 Structure and kinematics

The 3D shape of Sco-Cen is that of a highly elongated ellipsoid with
approximate dimensions of $100\times 100\times 50$ pc with the three subgroups
arranged in a row from north to south, with US the furthest away with a median
distance of 143 pc followed by UCL at 136 pc and LCC the closest at 115 pc
(Wright & Mamajek, 2018). The line-of-sight dispersion of the subgroups UCL
and LCC was first partially studied by de Bruijne (1999b) using Hipparcos
data, but it wasn’t until Gaia DR1 that the 3D internal structure of the
subgroups could be properly resolved (e.g., Wright & Mamajek, 2018; Galli et
al., 2018). The availability of Gaia DR2 facilitated more detailed structural
studies of the entire association, including the low-mass population for the
first time, and revealed a highly complex spatial structure (Damiani et al.,
2019).

The first estimation of the velocity dispersion of Sco-Cen came from de
Bruijne (1999b) who used kinematic modelling of Hipparcos data to estimate 1D
velocity dispersions of 1.0–1.5 km s-1 for each subgroup. These estimates were
supported by Madsen et al. (2002) who calculated 1D velocity dispersions of
1.1–1.3 km s-1 for each subgroup using astrometric radial velocities
calculated from Hipparcos data. The first direct measurement of the velocity
dispersion came from Wright & Mamajek (2018) who combined Gaia DR1 astrometry
with literature radial velocities to calculate 3D velocity dispersions in the
Galactic cartesian system. They found that the subgroups were highly
anisotropic, with 1D velocity dispersions varying from 0.5–2.5 km s-1 in each
dimension, thought notably lower in all three subgroups in the $V$ dimension
(the direction of Galactic rotation) than in the $U$ or $W$ dimensions. They
calculated average 1D velocity dispersions of $1.86\pm 0.21$ (US), $1.38\pm
0.21$ (UCL), and $1.21\pm 0.28$ km s-1 (LCC), which are approximately
consistent with previous studies, with the exception of US. They attribute
this disagreement to their inclusion of radial velocities and the differences
in the assumed membership of the association. Their 3D velocity dispersions
are $3.2\pm 0.2$ (US), $2.5\pm 0.2$ (UCL), and $2.2\pm 0.3$ km s-1 (LCC), for
which they calculate virial masses111The virial mass is defined as the mass
enclosed within a given radius for a system to be gravitationally bound. It is
given by a variant of the virial equation, $M_{vir}=\frac{2r\sigma^{2}}{G}$,
where $r$ is the radius of the system, $\sigma$ is its velocity dispersion,
and $G$ is the gravitational constant (e.g., Portegies Zwart et al., 2010).
that are an order of magnitude larger than the estimated stellar mass and
conclude that the three subgroups are gravitationally unbound.

Wright & Mamajek (2018) also find evidence for kinematic substructure within
each of the subgroups in the form of small groups of spatially-associated
stars that have more similar kinematics than the subgroup as a whole,
particularly in UCL and LCC. The V1062 Scorpii moving group (see Figure 6)
identified independently by Röser et al. (2018) and Zari et al. (2018) within
UCL is an example of this kinematic substructure, being a closely co-moving
group of $\sim$60 stars with a 1D velocity dispersion $<$1 km s-1 (see also
Goldman et al., 2018). These kinematic substructures appear to be related to
the spatial and temporal substructure that has already been identified in Sco-
Cen.

####  3.1.4 Expansion

Due to its proximity Sco-Cen is one of the most ideal associations to search
for evidence of expansion (studies of its expansion date back to Blaauw,
1952), but it also makes it difficult to reliably measure its expansion, for
two reasons. The first is that its large area on the sky (almost 90∘) means
that proper motions or radial velocities do not probe the same dimension
across the association, while the second and more important issue is that any
radial motion of the association towards (or away from) the observer will
cause a virtual expansion (or contraction), even when none is actually
present. Both of these issues mean that a reliable assessment of the expansion
of a nearby association requires radial velocities for some or all association
members.

Blaauw (1964b) studied proper motions and radial velocities for $\sim$20–30
members of each subgroup and concluded that their kinematics were more
consistent with linear expansion of the association than with no expansion.
Blaauw estimated an expansion age of 20 Myr based on an assumed distance of
200 pc (10 or 15 Myr if scaled to a distance of 100 or 150 pc). Blaauw (1978)
then followed this up by estimating an expansion age of $\sim$5 Myr for US.
Using radial velocities and revised Hipparcos proper motions Pecaut et al.
(2012) found a lower limit on the kinematic expansion age of 10.5 Myrs (99%
confidence), though their results are consistent with no expansion.

Figure 8: Positions versus velocities along the three Galactic Cartesian axes
$XYZ$ for each of the three subgroups of Sco-Cen (US on the left, UCL in the
centre, and LCC on the right). 1$\sigma$ error bars are shown for all sources.
The best-fit linear relationship between the plotted quantities shown as a
solid line with the best-fitting slope, $\kappa$, noted in each panel
(positive slopes imply expansion).

Using Gaia DR1 proper motions Wright & Mamajek (2018) investigated the
expansion of the three subgroups of Sco-Cen using numerous techniques to
explore evidence for expansion but could find no evidence for coherent
expansion patterns in any of the subgroups. When analysing the expansion of
the association in the 3D Galactic cartesian system (Figure 8) they found no
evidence for expansion in the $X$ and $Z$ dimensions, but notably found
significant ($\sim$3$\sigma$) evidence for expansion in the $Y$ dimension (the
direction of Galactic rotation) in all three subgroups. They suggested that
this expansion could be due to the effects of Galactic shear on the parental
molecular cloud from which the subgroups formed and that has now been
inherited by the stars of the association. Ward & Kruijssen (2018) also used
Gaia DR1 proper motions for stars across the entire Sco-Cen association to
assess evidence for expansion using multiple kinematic diagnostics, but could
find no evidence for expansion from either a single or multiple expanding
clusters.

Goldman et al. (2018) used Gaia DR2 astrometry to identify a large moving
group within LCC with approximately 1800 members, comparable in size to the
entire LCC subgroup and therefore potentially one of the larger kinematic
substructures in the region. They find significant evidence that this group is
expanding in all three dimensions with an expansion age of 8–10 Myr. This age
is consistent with the isochronal age Goldman et al. (2018) estimate, but
different from the age of $\sim$15 Myr calculated by Pecaut & Mamajek (2016)
for LCC as a whole. This would imply that this group is one of the younger
substructures within LCC, supporting the findings of Damiani et al. (2019)
that the clustered populations in Sco-Cen are younger than the diffuse
population.

Bobylev & Baykova (2020) used a sample of 700 low-mass stars selected by Zari
et al. (2018) to study the expansion of Sco-Cen. After correcting for the
influence of the Galactic spiral density wave (which previous studies of the
region have not), they measure a linear expansion rate of $39\pm 2$ km s-1
kpc-1 across the entire association. However, it is worth noting that the
magnitude of the correction applied is $\sim$50 km s-1 kpc-1, larger than the
magnitude of the expansion signal measured. If true, this implies that the
expansion of the association has been entirely hidden by the spiral density
wave perturbation, since the observed, uncorrected kinematics do not show any
evidence for expansion on the scale of the entire association. This raises
questions of whether such corrections are necessary to expose the initial
expansion of the association, and if so whether the calculations for such
corrections should be based on the initial positions of the stars, rather than
their current positions (as this study has done).

####  3.1.5 The Greater Sco-Cen Complex

The Sco-Cen association is not an isolated group of young stars as there are
many young groups of stars and small star-forming regions in the local
vicinity that appear to be related, such as $\rho$ Ophiuchus, the Lupus
clouds, and R Corona Australis (see Figure 6). Mamajek et al. (1999)
discovered a cluster of young ($\sim$8 Myr) stars in the vicinity of the B8
star $\eta$ Cha whose Galactic motion suggests that the cluster, as well
probably as members of TW Hya and $\epsilon$ Cha, probably originated in or
near the same giant molecular cloud (GMC) that formed UCL and LCC 15 Myr ago
(Mamajek et al., 2000). Going even further, Mamajek & Feigelson (2001)
consider that $\beta$ Pic and other nearby $\sim$10 Myr old stars may also
have formed in this large star-forming event.

This extended region of young stars has been referred to as either the Oph-
Sco-Cen association (Blaauw, 1991) or Greater Sco-Cen (Mamajek & Feigelson,
2001) and is usually considered to include the three OB subgroups, surrounding
molecular clouds, nearby star-forming regions such as $\rho$ Ophiuchus and the
Lupus clouds, and the outlying associations and moving groups that have
recently been uncovered. Other nearby groups or associations have been
reported as possible extensions of Sco-Cen (e.g., Eggen, 1998; Makarov &
Urban, 2000), though these have not generally stood up to further
investigation (Preibisch & Mamajek, 2008).

####  3.1.6 Feedback and effects on the local ISM

The interstellar medium surrounding Sco-Cen shows numerous large, loop-like
structures and shells that suggest significant interaction between the
association and the surrounding medium. They are most clearly seen in Hi
observations (e.g., Weaver et al., 1977; de Geus, 1992; Robitaille et al.,
2018), with small bubbles centred around US and LCC and a larger bubble
encasing the entire association and approximately centred on UCL (de Geus,
1992). The total Hi mass of these bubbles has been estimated to be $\sim
5\times 10^{5}$ M⊙ and suggested to be composed of gas swept up by stellar
winds and supernovae from the progenitor giant molecular cloud that Sco-Cen
formed from, as well as the ambient ISM (de Geus, 1992).

de Geus (1992) calculated the total kinetic energy of the bubbles to be
$10^{51}$ erg and argued that the energy input from the existing massive stars
coupled with past supernova activity across Sco-Cen would be sufficient to
produce the observed network of bubbles. Krause et al. (2018) verify this by
using a 3D hydrodynamic simulation to show that the observed massive star
population could reproduce the size and morphology of the observed
superbubbles. The presence of soft X-ray emission from the US bubble coupled
with the detection of 26Al confirms that there has been at least one
relatively recent ($\lesssim$1 Myr) supernova from within this subgroup
(Robitaille et al., 2018; Krause et al., 2018).

These supernovae may have played some role in the propagation of star
formation through Sco-Cen, an idea advocated by Preibisch & Zinnecker (1999)
to explain the small age spread measured in US at the time. This idea was
challenged by the observation of larger age spreads by Pecaut & Mamajek
(2016), with suggestions that the subgroups of Sco-Cen probably do not
represent distinct star formation events. Krause et al. (2018) put forward an
alternative model wherein the expanding superbubbles originating from one of
the subgroups triggers star formation in the next subgroup by surrounding and
squashing the primordial dense gas. Whatever the mechanism, this propagation
of star formation appears to have continued beyond Sco-Cen itself and may have
been responsible for triggering star formation in Lupus, Ophiuchus, Corona
Australis and Chameleon (Mamajek & Feigelson, 2001).

###  3.2 Orion OB1

The Orion OB association (Orion OB1) is the richest association in the Solar
neighbourhood and one of the few still closely associated with a giant
molecular cloud. It spans over 200 deg2 on the sky across the constellation of
Orion at distances between 300–450 pc (Brown et al., 1994; Zari et al., 2018).
It contains at least 56 massive stars of spectral type B2 and earlier (Blaauw,
1964a) including the bright stars $\zeta$ and $\delta$ Ori and $\theta^{1}$
Ori C, the most massive stars of the trapezium in the Orion Nebula Cluster
(ONC), part of Orion OB1d. Figure 9 shows the structure of the Orion OB1
association, its division into subgroups, and some of the notable young
clusters in its vicinity.

Figure 9: Wide-field map of the Orion constellation showing the locations of
the four subgroups of the Orion OB1 association as well as the $\lambda$ Ori
region (white ellipses and labels). Also shown are prominent star clusters
across the region (in yellow), Barnard’s Loop and the bright stars Betelgeuse
and Rigel (light blue). Background astrophotograph courtesy of Stanislav
Volskiy.

Blaauw (1964a) divided the association into four subgroups labelled $a$ to $d$
(see Figure 9). Orion OB1a is generally considered to be the oldest subgroup,
aged approximately 8–12 Myr (Brown et al., 1994), and lying north-west of
Orion’s Belt. It includes the cluster 25 Ori (Briceño et al., 2007), aged 7–10
Myr, which may be the youngest component of this subgroup. Subgroups 1b and 1c
have approximately similar ages of 2–8 and 2–6 Myr, respectively (Bally,
2008), with different authors placing them in different chronological order
(e.g., Blaauw, 1964a; Brown et al., 1994). Orion OB1b includes the stars of
Orion’s Belt as well as the cluster $\epsilon$ Ori (Kubiak et al., 2017),
while Orion OB1c includes the stars around Orion’s Sword as well as the
clusters NGC 1980 (considered by some to lie in the foreground, Alves & Bouy,
2012) and NGC 1981. $\sigma$ Ori (Walter et al., 2008) is projected close to
Orion OB1b but its age is closer to that of Orion OB1c, so is often considered
part of the latter. Finally, Orion OB1d is the youngest of the subgroups and
is associated with the molecular gas in the region (the Orion A and B
filaments). It includes the ONC, M43, NGC 1977, NGC 2024 clusters and the
Orion molecular clouds (note that while the NGC 2024 cluster is projected
against Orion OB1b and c, its very young age makes it more consistent with
OB1d, e.g., Levine et al., 2006). These subgroups have been defined almost
entirely based on position on the sky, and while important from a historical
perspective, recent work suggests that sub-dividing the association using
position, parallax and kinematics, providing a more accurate view of the star
formation history of the region (e.g., Zari et al., 2017; Kounkel et al.,
2018; Chen et al., 2019; Zari et al., 2019).

In addition to the four subgroups defined by Blaauw (1964a) there are numerous
outlying groups associated with the same star formation events that formed the
main OB association. These include $\lambda$ Ori, which while spatially offset
from the main association has an age similar to the 1b and 1c subgroups, so
appears to be related to the association (see Figure 9). Zari et al. (2017)
identified further young stars in the vicinity of $\lambda$ Ori that appear to
be associated with the H$\alpha$ bubble surrounding the main cluster,
suggesting further star formation across this area.

There are very few estimates of the total stellar mass of the Orion OB1
association in the literature. Bally (2008) estimate that 30–100 stars more
massive than 8 M⊙ have formed across the association. Assuming a standard
Kroupa (2001) initial mass function (IMF) as formulated by Maschberger (2013)
such stars constitute approximately 0.4% of the full IMF, implying a total
stellar population of 7,500–25,000 stars. Assuming a mean stellar mass of 0.36
M⊙ (Maschberger, 2013) and an average binary fraction of 50% this equates to a
total stellar mass of 4000–13,500 M⊙.

####  3.2.1 The high-mass stars

There are a considerable number of high-mass OB stars within the Orion OB1
association. Blaauw (1964a) identify 53 OB stars of spectral type B3 and
earlier, specifically within the subgroups a, b and c. The most massive of
these are found in subgroup 1b and include $\zeta$ Ori A (O9.7Ib, $\sim$49
M⊙), $\delta$ Ori A (B0III, $\sim$45 M⊙) and $\epsilon$ Ori A (B0Ia, $\sim$42
M⊙), with masses from Lamers & Leitherer (1993). Other notable massive stars
include $\eta$ Ori (B1V, subgroup 1a, Warren & Hesser, 1978), $\iota$ Ori A
(O9III, subgroup 1c, Lamers & Leitherer, 1993), and the stars of the Trapezium
system in the ONC ($\theta^{1}$ Ori C, O6V, being the most massive of them).
The latter in particular have been the focus of numerous searches for binary
companions. $\theta^{1}$ Ori C was resolved by Weigelt et al. (1999) as a
close binary, the companion having a mass of $\sim$5 M⊙ compared to the
primary’s mass of $\sim$45 M⊙. The other O-type star in the ONC, $\theta^{2}$
Ori A (O9.5V) is known to be a triple system composed of a spectroscopic
binary with a mass ratio of 0.35 (Abt et al., 1991) and a slightly less
massive, resolved companion (approximately 3–7 M⊙, Preibisch et al., 1999).
See also the recent interferometric study by Gravity Collaboration et al.
(2018) that uncovered additional companions to many of the OB stars in Orion.

There are many runaway massive stars known in Orion, most notable of which are
AE Auriga and $\mu$ Columbae, which are moving in opposite directions at
speeds of 150 and 117 km-1, respectively (Blaauw, 1991). Hoogerwerf et al.
(2001) showed that these two stars and the massive binary system $\iota$ Ori,
were at the same location in the sky $2.6\pm 0.05$ Myrs ago, and may be the
product of a dissolved double-binary system. Another runaway, 53 Ari, may have
been ejected from subgroup 1c 7.3 Myrs ago (Brown et al., 1994) during a
supernova explosion.

In addition to runaway massive stars, the Orion association is probably the
birthplace of a number of ejected neutron stars. Pellizza et al. (2005)
studied the origin of the Geminga neutron star and estimate an origin in
either the Orion OB1 or the Cas-Tau association, while Kounkel (2020) used the
proper motion of the pulsar at the centre of the Monogem supernova remnant
(Golden et al., 2005) to suggest that it may have been ejected from the
$\lambda$ Ori cluster approximately 1.7 Myr ago.

####  3.2.2 Structure and age

The spatial, kinematic and temporal structure of Orion OB1 is highly complex.
The classical subdivision of the association into predominantly four subgroups
was based entirely on sky position, but the availability of radial velocity
data (Jeffries et al., 2006) and later from Gaia parallaxes and proper motions
have suggested a revision of this structure that takes into account the full
6-dimensional spatial and kinematic information (Kubiak et al., 2017; Kounkel
et al., 2018; Zari et al., 2019).

Figure 10: Estimated stellar ages across the Orion OB1 association from
Kounkel et al. (2018). The left-hand panel shows ages derived using
spectroscopic effective temperature and bolometric luminosities in the HR
diagram, while on the right are shown ages derived using just photometry in
the colour magnitude diagram (assigning distances according to the average
distance of stars in each group).

Kounkel et al. (2018) performed a clustering analysis of APOGEE and Gaia DR2
data of 10,248 young stars across Orion, identifying 190 groups of stars. They
associated many of these groups with either the young populations in Orion A
or B, the cluster $\lambda$ Ori, or two new progenitor groups dubbed Orion C
and D (to extend the existing naming scheme). These two new groups have
similar distributions on the sky, both broadly tracing the areas of the OB1a
and OB1b subgroups, though they have distinctly different radial velocities
and parallaxes (see Figure 4). Orion D makes up the majority of the OB1a and
OB1b subgroups, the 25 Ori cluster, and includes the bright stars $\eta$ and
$\psi^{2}$ Ori. It spans the largest area on the sky of the new groupings,
extending as far South as Rigel, Orion X (Bouy & Alves, 2015) and some of the
Orion outlying clouds (though this southern component could be associated with
Orion A). It is aged from 4–7 Myr and for the most part has a radial velocity
of 3–5 km s-1 and parallaxes of 2.6–3.0 mas ($d\simeq 360$ pc). Orion C is
mostly composed of the $\sigma$ Ori cluster, but also extends northwards to
the Orion Belt population and into the area covered by Orion OB1a. Its age
varies from $\sim$7.5 Myrs in the north down to $\sim$2 Myrs around $\sigma$
Ori and has distinctly different radial velocity (as first observed by
Jeffries et al., 2006) and parallax from the Orion B and D groups it is
projected near. Figure 10 shows the age distribution across the association,
showing its complex physical and temporal structure.

Zari et al. (2019) noted a particularly high degree of substructure in the
region towards the Belt stars, finding multiple groups of stars with different
kinematic properties, including groups towards $\sigma$ Ori, $\zeta$ Ori,
$\epsilon$ Ori, and $\delta$ Ori. These groups range in age from $\sim$4 Myr
(for the groups around $\sigma$ Ori and $\zeta$ Ori) to $>$10 Myr for the more
extended groups. The age estimates and rankings found by Zari et al. (2019)
are generally in agreement with those found by Kounkel et al. (2018), with
both studies finding a very complex and fragmentary star formation history.

Historically an age gradient was suggested to extend from subgroups OB1a and
1b to the younger OB1d group (Brown et al., 1994), however the results from
Gaia DR2 do not show evidence for a clear progression of star formation across
the association and suggest a much more complex age distribution (Kounkel et
al., 2018; Zari et al., 2019). Gaia data has also lead to the discovery of
multiple older populations in the Orion OB1 association, including a $15--21$
Myr population connected to the ASCC 20 cluster (Kos et al., 2018; Zari et
al., 2019), another similarly-aged population towards Orion’s Belt (group B8
in Zari et al., 2019), and a $\sim$17 Myr population behind the main
association ($d\sim 430$ pc) that is spatially quite extended ($\sim$90 pc in
length) but kinematically compact (Kounkel et al., 2018; Zari et al., 2019;
Jerabkova et al., 2019).

This temporal complexity isn’t limited to the widest parts of the association.
Beccari et al. (2017) found evidence for multiple pre-main sequences in the
colour-magnitude diagram of stars around the Orion Nebula Cluster, which they
argued represented multiple distinct episodes of star formation within the
cluster, with the youngest populations more concentrated toward the ONC than
the older populations. Jerabkova et al. (2019) verified the existence of these
sequences, but found indications that the older population may be slightly in
the foreground of the ONC.

####  3.2.3 Foreground populations

There is growing evidence for recent star formation in front of the Orion OB1
association. Some of the brightest stars of the Orion constellation are young
and massive stars, such as Betelgeuse ($\alpha$ Ori, M2Iab), Rigel ($\beta$
Ori, B8Iab) and Saiph ($\kappa$ Ori, B0Iab), and are found at distances of
150–300 pc (Perryman et al., 1997). These stars are too far in the foreground
to be consistent with being ejected from Orion OB1, and instead suggest recent
star formation has occurred in the foreground (Bally, 2008). This was verified
by Bouy & Alves (2015) who identified a large population of early-type stars
in the foreground of Orion and named it Orion X, suggesting they constituted a
previous generation of star formation to the main Orion OB1 association. Zari
et al. (2017) also found evidence for an older population of stars in this
direction from Gaia DR1 data, but using Gaia DR2 data Zari et al. (2018) and
Zari et al. (2019) were not able to find clear evidence for the existence of
Orion X, thus leaving the origin of the brightest stars in the Orion
constellation unresolved.

Alves & Bouy (2012) identify a rich population of young (4–5 Myr) stars just
in front of the Orion Nebula Cluster (ONC) that includes many stars associated
with $\iota$ Ori and the NGC 1980 cluster (previously considered part of Orion
OB1c, see Figure 9). They argue that NGC 1980 is a relatively massive
($\sim$2000 members) and hitherto unrecognised cluster and, combined with a
general foreground populations, these stars include many previously thought to
be part of the ONC. Bouy et al. (2014) and Zari et al. (2019) both confirmed
the existence of a foreground population, with the former estimating it to
consist of $\simeq$2600 stars with an age of 5–10 Myr. However, Da Rio et al.
(2016) and Fang et al. (2017) found that these stars have similar radial
velocities to stars in the ONC and suggested they are likely to be part of the
same star formation event. Fang et al. (2017) calculate a median age of 1–2
Myr, while Kounkel et al. (2017) estimate an age of 3 Myr, suggesting the
group is similarly-aged to the ONC. It is possible that these foreground
populations are related to the older pre-main sequence identified by Beccari
et al. (2017) towards the ONC and that Jerabkova et al. (2019) suggest may be
slightly in the foreground of the cluster.

####  3.2.4 Kinematics and expansion

The youngest parts of the Orion OB1 association, particularly subgroups OB1c
and OB1d, are still spatially and kinematically associated with the molecular
clouds that they formed from. Fűrész et al. (2008) and Tobin et al. (2009)
performed wide-field spectroscopic surveys of these young stars and found a
correlation between the stellar radial velocities and that of the dense gas
they are projected against. They argue that this kinematic coherence implies
these stars are still very young and that they haven’t yet decoupled from or
dispersed the gas that they formed from. Hacar et al. (2016) and Da Rio et al.
(2017) extended this work with spectroscopic studies across the Orion A
molecular cloud and observed a correlation between stellar and gas velocities
over the 5∘ extent of the cloud. In addition to this correlation, Fűrész et
al. (2008) found evidence for kinematic substructure within the association,
in the form of spatially-coherent groups of stars with similar motions. Da Rio
et al. (2017) found that these groups are preferentially found in the low-
density regions of the cloud, while the denser ONC is less kinematically
substructured.

Most measurements of the velocity dispersion within the Orion OB1 association
have been limited to the more-easily identified young stars in and around the
ONC and the OB1c and OB1d subgroups. Jones & Walker (1988) and van Altena et
al. (1988) calculated proper motion velocity dispersions for stars in the ONC
of $2.34\pm 0.09$ and $1.49\pm 0.2$ km s-1, respectively. The smaller velocity
dispersion calculated by van Altena et al. (1988) for a sample of generally
brighter stars can be at least partly explained by an anti-correlation between
velocity and mass. This was also as found by Hillenbrand et al. (1998) who
measured a velocity dispersion of 2.81 km s-1 for stars with $M=0.1$–0.3 M⊙
and a velocity dispersion of 2.24 km s-1 for stars with $M=1$–3 M⊙. This anti-
correlation suggests partial energy equipartition within the system (i.e.,
that the massive stars are moving slower than the less massive stars). Da Rio
et al. (2017) measure a radial velocity dispersion of 2.2 km s-1, larger than
their estimated virial velocity dispersion of $\sim$1.7 km s-1, suggesting
that the system is slightly supervirial. Recently, Kuhn et al. (2019) and Kim
et al. (2019) measured the PM velocity dispersion of stars in the ONC using
Gaia DR2 and Hubble Space Telescope + Keck data, respectively. They both found
the velocity distributions to be highly anisotropic, Kuhn et al. (2019)
measuring velocity dispersions of $1.51\pm 0.11$ (RA) and $0.50\pm 0.12$ (Dec)
mas yr-1 from 48 stars, while Kim et al. (2019) measure dispersions of
$0.83\pm 0.02$ (RA) and $1.12\pm 0.03$ (Dec) mas yr-1. Kim et al. (2019)
attribute the difference in measured dispersions to differences in the sample
size and under-estimated PM uncertainties in Gaia DR2. On a larger scale,
Fűrész et al. (2008) measure a radial velocity dispersion of 3.1 km s-1 across
the Orion OB1c association.

Numerous attempts have been made to determine whether the Orion OB association
subgroups are expanding and to estimate their kinematic ages. Lesh (1968)
studied the expansion of the OB1a subgroup and measured a kinematic age of 4.5
Myr ($\pm$30%), while Blaauw (1961) estimate a kinematic age of 2.2–4.9 Myrs
for subgroup OB1b based on the motions of the three runaway stars AE Aur,
$\mu$ Col, and 53 Ari. Kounkel et al. (2018) find that their Orion D group,
which contains subgroups OB1a and OB1b as well as Orion X exhibits PMs
consistent with expansion, while the stars associated with the Orion B cloud
have PMs more consistent with contraction. Kounkel et al. (2018) also found
clear evidence for expansion for stars in the $\lambda$ Ori cluster, with a
correlation between distance from the cluster centre and radial outward
velocity that suggests a single trigger of the expansion approximately 4.8 Myr
ago.

####  3.2.5 Feedback and the Orion-Eridanus Superbubble

The Orion OB1 association is very prominently surrounded by a bright crescent
of H$\alpha$ emission known as Barnard’s Loop, which is believed to be the
brighter part of a larger bubble that extends 40∘ west and incorporates the
Eridanus Loop (Sivan, 1974; Reynolds & Ogden, 1979; Ochsendorf et al., 2015).
This is the Orion – Eridanus Superbubble, an approximately $140\times 300$ pc
bubble (Bally, 2008) with a 1–$5\times 10^{5}$ K interior that has been
detected in the UV, X-ray and in the 1.8 MeV $\gamma$-ray line that indicates
the presence of the short-lived radioactive species 26Al (Diehl et al., 2004,
implying recent supernova activity). The bubble is surrounded by an expanding
($\sim$15–20 km s-1, Reynolds & Ogden, 1979; Joubaud et al., 2019) shell of
ionized gas, which is itself surrounded by an expanding H i shell with a mass
of $\sim$300,000 M⊙ (Brown et al., 1995). Ochsendorf et al. (2015) studied the
Orion-Eridanus superbubble with multi-wavelength data and argue it is both
larger and more complex than previously thought, potentially viewed as a
series of nested shells superimposed along the line of sight, a picture that
was supported by Joubaud et al. (2019).

Bally (2008) estimate that there have been at least 10–20 SNe in the last 12
Myr that would have released $>10^{52}$ ergs of kinetic energy, sufficient to
form the bubble. Additional energy may be being injected by smaller bubbles
within the Orion-Eridanus superbubble, such as the Veil shell that surrounds
the ONC and is thought to be powered by stellar winds from $\theta^{1}$ Ori-C
(Pabst et al., 2020).

The SNe that have driven the superbubble may have contributed to the
enrichment of the association subgroups. Cunha & Lambert (1994) found that the
younger subgroups OB1c and OB1d appear to be more abundant in oxygen and
silicon with respect to stars in the older subgroups OB1a and OB1b, while
other elements such as carbon, nitrogen or iron do not show such abundance
variations. Such a pattern is consistent with being due to supernova ejecta
enriching the interstellar gas that went on to form subgroups OB1c and OB1d.
However, both Simón-Díaz (2010) and Biazzo et al. (2011) find similar
abundances of the $\alpha$-elements silicon and titanium across the OB
association, arguing against self-enrichment.

###  3.3 Vela OB2

The Vela OB2 association was first identified by Kapteyn (1914) and included
in Blaauw (1964a)’s list of associations. At a distance of between 350–400 pc
(de Zeeuw et al., 1999; Zari et al., 2018), it spans at least $15\times 15$
degrees across the constellations of Vela, Puppis and Carina (see Figure 11).
Its total mass has been estimated to be between $\sim$1300 M⊙ (Armstrong et
al., 2018) and over 2330 M⊙ (Cantat-Gaudin et al., 2019b), depending on where
the borders of the association are drawn. Due to its projection against the
Galactic Plane, its relative dearth of massive stars (compared to Sco-Cen and
Orion for example), and its older age, it has historically not received as
much attention as other nearby associations. However, in recent years and
thanks to Gaia data considerable work has helped uncover its low-mass
population (e.g., Zari et al., 2018; Cantat-Gaudin et al., 2019a).

The brightest star in the association is the Wolf-Rayet WC8 + O9I binary
system $\gamma^{2}$ Velorum, the nearest known Wolf-Rayet star, with masses of
$9\pm 0.6$ and $28.5\pm 1.1$ M⊙, respectively (North et al., 2007). Questions
have been raised over the membership of the star in Vela OB2, since
evolutionary models for the star suggest an age of $5.5\pm 1$ Myr (Eldridge,
2009), while the lower-mass pre-main sequence stars in its vicinity have ages
of $\sim$10 Myr (Jeffries et al., 2009). This has led to suggestions of mass
transfer within the binary system that might make it appear younger (Jeffries
et al., 2014), since Hipparcos astrometry suggests the star is a member of the
association (Rate et al., 2020). de Zeeuw et al. (1999) identified a further
92 members of the association, predominantly B-type stars, which Armstrong et
al. (2018) reduced to 81 after excluding photometric and astrometric
contaminants.

Figure 11: Spatial distribution of young stars across Vela OB2 and surrounding
regions from Cantat-Gaudin et al. (2019a). The stars are coloured according to
the groups identified by those authors, from I to VII (blue, pink, grey, light
green, turquoise, dark green and orange). Notable clusters in the region are
marked.

The association borders the Trumpler 10 region, previously considered an open
cluster, but redefined by de Zeeuw et al. (1999) as an association of 22
B-type stars (earliest spectral type B3V) that are slightly older than Vela
OB2 but at a similar distance. Its connection to Vela OB2 and the wider Vela-
Puppis region has recently been confirmed by Cantat-Gaudin et al. (2019a), as
seen in Figure 11.

####  3.3.1 Structure and ages

The most prominent overdensity within the association is the Gamma Velorum
cluster that surrounds $\gamma^{2}$ Vel. It was discovered by Pozzo et al.
(2000) from ROSAT X-ray observations that coincided with low-mass, pre-main
sequence stars and verified with follow-up photometric and spectroscopic
observations (Jeffries et al., 2009). The cluster and the massive binary share
a common proper motion and so are strongly believed to be associated with each
other and with the wider Vela OB2 association (Jeffries et al., 2009).
Jeffries et al. (2014) and Prisinzano et al. (2016) estimate the mass of the
cluster to be $\simeq$100 M⊙. Another cluster within the association, the P
Puppis cluster, was identified by Caballero & Dinis (2008) and later by
Beccari et al. (2018).

Using Gaia-ESO Survey spectroscopy, Jeffries et al. (2014) performed a
detailed spectroscopic study of the Gamma Velorum cluster and found evidence
for two distinct radial velocity components, one broad and one narrow, as
shown in Figure 12. The narrow velocity component has been associated with the
cluster itself, and suggests the system is in virial equilibrium, while the
broad component was attributed to the wider Vela OB2 association and seemed to
be slightly younger than Gamma Vel. Sacco et al. (2015) also identified a
broad kinematic component projected against the older NGC 2547 cluster,
suggesting the young, low-mass population of Vela OB2 was spread over a large
area in this region. Bouy & Alves (2015) found hints of this population in the
foreground of Vela OB2 and suggested an age gradient existed covering this
young foreground population, Gamma Velorum, Trumpler 10, and the older Vela
OB2 association behind it.

Figure 12: Radial velocity distribution of young stars towards Gamma Vel from
Jeffries et al. (2014). The red histogram shows the young stars towards the
cluster while the blue dotted line shows the best-fitting two Gaussian
component model (velocity dispersions noted on the figure), which takes into
account the effect of binaries on the velocity dispersions.

The full extent and complex spatial and kinematic structure of Vela OB2 and
the surrounding region has only become fully apparent after the release of
Gaia data. Armstrong et al. (2018) used Gaia DR1 photometry to identify low-
mass stars across Vela OB2, uncovering an extended population with
concentrations around Gamma Velorum and a distributed population over a wider
area that has considerable spatial substructure. Beccari et al. (2018) applied
a clustering analysis to Gaia DR2 astrometry across this area to verify some
of these over-densities and identify new possible clusters or substructures
within the association (see also Damiani et al., 2017; Cantat-Gaudin et al.,
2019b).

Cantat-Gaudin et al. (2019a) performed a wide-area study of young stars across
the Vela and Puppis regions, selecting young stars using a combination of
photometry and astrometry and then using an unsupervised classification scheme
to divide their sample into separate populations, finding 7 distinct groups as
shown in Figure 11. Three of these populations (labelled III, V and VII by the
authors) fall within the classical boundaries of Vela OB2 (de Zeeuw et al.,
1999), while two (I and II) fall within the Trumpler 10 region. The other two
populations (IV and VI) are found to the west of Vela OB2 towards the
constellation Puppis. Population IV is the most populous of all the groups and
contains the open clusters NGC 2547, NGC 2451B, Collinder 135, Collinder 140
and UBC 7, while population VI contains the open cluster BH 23. Cantat-Gaudin
et al. (2019a) estimate a range of ages for these populations, from $\sim$10
Myrs for population VII (containing Vela OB2 itself) through to 45–50 Myrs for
Populations I and II in the Trumpler 10 region. The authors find no evidence
for the spatial gradient with age found by Bouy & Alves (2015) and instead
argue for multiple episodes of star formation across the region with a complex
star formation history.

####  3.3.2 Kinematics and expansion

Jeffries et al. (2014) calculated radial velocity dispersions of $0.34\pm
0.16$ and $1.60\pm 0.37$ km s-1 for the two kinematic components they
identified towards Gamma Vel. The first component corresponds to the Gamma
Velorum cluster itself, and is consistent with the cluster being in virial
equilibirum (given a cluster mass of $\sim$100 M⊙). The second component
constitutes some part of the Vela OB2 association and indicates that it is
unbound. Franciosini et al. (2018) studied the correlation between radial
velocity and parallax of this component and found a gradient that they
interpret as a sign of expansion.

Cantat-Gaudin et al. (2019a) measure the 3D expansion in all 7 groups they
identify across Vela–Puppis and find evidence for expansion in all of them,
predominantly along the Galactic $X$ axis (towards the Galactic centre),
though the majority do also show expansion in $Y$ and $Z$. Notably, the
expansion rates they measure appear non-isotropic, in that the expansion is
occurring at different rates along the different axes, with differences up to
a factor of 10 between the expansion rates (one of the groups also shows
significant evidence for contraction along the $Z$ axis). This may be due to
physical and kinematic substructure within the groups that has yet to be
resolved or it may be that the expansion of these groups has been non-
isotropic since they started expanding. Due to the anisotropic expansion it is
not immediately clear of the legitimacy of calculating kinematic ages for
these groups but they vary from $\sim$20 to $\sim$50 Myrs, approximately in
line with the evolutionary ages for each population.

Armstrong et al. (2020) performed a spectroscopic survey of young stars in
Gamma Vel and the region of Vela OB2 in its vicinity, confirming the youth of
their targets from the presence of lithium, and separating cluster and
association members based on both their positions and kinematics. They
combined spectroscopic radial velocities with Gaia DR2 proper motions to study
the structure and 3D dynamics of their targets, finding strong ($5-7\sigma$)
evidence that the association is expanding in all three dimensions. The
expansion rates measured are non-isotropic, varying by more than a factor of
two between axes.

####  3.3.3 Feedback and formation

Sahu (1992) discovered a giant ($\sim$10∘) ring-like structure in IRAS images
of Vela and Puppis, part of the Gum nebula. At a distance of $\sim$450 pc it
is believed to be connected to the Vela OB2 association and is also
discernible in maps of interstellar extinction and neutral hydrogen (Testori
et al., 2006). The massive stars $\eta$ Puppis (O4I) and $\gamma^{2}$ Velorum
are thought to be responsible for powering the shell (Reynolds, 1976). Sahu
(1992) estimate the energy output of the Vela OB2 stars from both stellar
winds and supernovae, and find it to be consistent with that needed to produce
the shell, while Higdon & Lingenfelter (2013) find that only 10% of the
ionizing radiation from the Vela O stars is necessary to power the shell.
Cantat-Gaudin et al. (2019b) argue that the distribution of stars in Vela OB2
forms a ring-like structure that follows the Vela shell and may have formed by
triggering as part of the shell’s expansion.

###  3.4 Cygnus OB2

Figure 13: Spatial distribution of known high-mass stars in the vicinity of
Cygnus OB2 from Wright et al. (2015) and Berlanas et al. (2018). Blue dots
show known O-type stars, red dots show B-type stars, white triangles show
Wolf-Rayet stars and orange dots show suspected OB stars (photometrically
identified, primarily from Comerón et al., 2002, 2008). The white circle shows
the 1 deg2 area studied by Wright et al. (2015), which is centred on the
‘trapezium’ of O stars known as Cyg OB2 #8. It is clear however that the
density of both O and B-type stars is non-zero across the entire Cygnus
region. The background image is a mid-IR 12 $\mu$m image (Wright et al.,
2010a) showing emission from the surrounding Cygnus X giant molecular cloud,
with prominent radio sources from Downes & Rinehart (1966) labelled.

The Cygnus OB2 association is probably the most massive association known in
our galaxy, with at least 50 O-type stars and a total mass of $\sim$16,500 M⊙
(Wright et al., 2015). While it is more distant than the majority of the other
associations discussed here, it has received a comparable level of attention
thanks to its impressive contingent of massive stars and its position within
the massive Cygnus X star forming region that spans almost $\sim$200 pc
(Reipurth & Schneider, 2008). Figure 13 shows the distribution of known OB
stars in and around Cygnus OB2, showing a more concentrated distribution than
many other OB associations due to its mass and relative youth. It was first
identified as a group of early-type stars by Münch & Morgan (1953), labelled
as an O-association by Johnson & Morgan (1954, who obtained some of the first
spectral types for the O stars in the region) and given the name VI Cygni by
Schulte (1956). As the second of nine OB associations listed by Ruprecht
(1966) it became Cygnus OB2 and recognised as the massive association that we
know now by Reddish et al. (1967).

The distance to the association was originally estimated to be about 1.8 kpc
(Torres-Dodgen et al., 1991; Massey & Thompson, 1991), but was brought down to
$\sim$1.4 kpc by Hanson (2003) based on a revised effective temperature scale
for OB stars. This was found to be in good agreement with radio parallax
measurements towards multiple masers within Cygnus X that suggested a typical
distance of 1.4 kpc for the star forming region (Rygl et al., 2012). A recent
study by Berlanas et al. (2019) used Gaia DR2 parallaxes to constrain the
distance to the most massive members of the association, suggesting that the
main Cyg OB2 association was actually at 1.76 kpc, with a smaller, foreground
population of massive stars at 1.35 kpc.

####  3.4.1 The high-mass stars

A considerable amount of attention has been focussed on the high-mass stars
within Cyg OB2 since the association provides one of the largest groups of
O-type stars in our galaxy. The first detailed study of the high-mass stellar
content of Cyg OB2 was performed by Massey & Thompson (1991), who
spectroscopically identified 42 O-type and 26 B-type stars. The association
however is highly obscured by the foreground Cygnus Rift, with the majority of
its members having an extinction of $A_{V}=4$–7 mag (Wright et al., 2015), and
therefore further work to expand the census of the association was carried out
in the infrared (e.g., Knödlseder, 2000; Comerón et al., 2002). Knödlseder
(2000) performed a photometric study of the association that suggested it
might include $>$100 O-type stars, and while Comerón et al. (2002) and Hanson
(2003) were able to verify many of these spectroscopically, their estimates of
the high-mass stellar content in the association could not reach these levels.

The current estimate of 55–60 O-type stars (e.g., Wright et al., 2015;
Berlanas et al., 2018) still makes it comparable to many of the most massive
young clusters in our galaxy and yet Cyg OB2 has the advantage that it is
significantly closer and not as crowded, making source confusion less of a
problem (e.g., Caballero-Nieves et al., 2014). Studies of the high-mass
initial mass function in Cyg OB2 have found evidence for the slope to be both
flatter (Massey & Thompson, 1991) and steeper (Kiminki et al., 2007) than the
canonical Salpeter value, though more recent studies that take a full census
and account for the ageing of the population suggest it is consistent (see
Section 4.8 and Wright et al., 2015).

The recent census of Cyg OB2 by Wright et al. (2015, see Figure 13) puts the
total number of O-type stars currently in the association at $\sim$52, though
new members are still being discovered (e.g., Berlanas et al., 2018, 2020).
This includes two stars in the rare O3 spectral class (Walborn, 1973, 2002),
the massive main sequence O5.5V star MT516, three Wolf-Rayet stars, and the
B-type supergiant Cyg OB2 #12, one of the most luminous and heavily reddened
stars known. Comerón & Pasquali (2007) have also identified a massive ($70\pm
15$M ⊙) runaway O4If star whose bow shock suggests it was ejected from Cyg OB2
$1.7\pm 0.4$ Myr ago. If this star was ejected when its binary companion
exploded as a supernova (Blaauw, 1961) it also provides evidence that Cyg OB2
has already seen its first supernova (for a discussion of evidence for other
past supernovae in Cyg OB2 see Wright et al., 2015). Berlanas et al. (2020)
performed a spectroscopic survey of the O-type stars in and around Cyg OB2,
expanding the census of Wright et al. (2015) and extending it to the south-
west of the association where an older population of massive stars had
previously been identified (Comerón & Pasquali, 2012).

Figure 14: Proper motion vector map for X-ray and spectroscopically selected
young stars towards the centre of Cyg OB2 (including 16 O-type stars and 34
B-type stars) from Wright et al. (2016). Dots show the current positions of
the stars with the vectors showing their proper motions, colour-coded based on
their direction of motion to highlight the kinematic substructure (colour
wheel in the top-right corner illustrates this). Kinematic outliers were
removed due to their large motions relative to most stars in the sample. The
grey box shows the border of the region studied.

Thanks to the proximity of the association many of these stars are known to be
in either close spectroscopic (SB1/2) or visual binary systems and have had
their properties studied in considerable detail (e.g., Kiminki et al., 2009;
Kobulnicky et al., 2014; Caballero-Nieves et al., 2014). This includes the
massive O3If+O6V binary Cyg OB2 #22 (Walborn, 2002), the colliding-wind binary
Cyg OB2 #8A (O6I+O5.5III, De Becker et al., 2004), and the quadruple system
and contact-binary Cyg OB2 #5 (O7I + O6I + O9V + B, Dzib et al., 2013). Based
on their accumulated sample, Kiminki & Kobulnicky (2012) estimate an upper
limit for the intrinsic binary fraction in Cyg OB2 to be $90\pm 10$%, with
$\sim$45% having companions close enough to interact with during their
lifetimes. Caballero-Nieves et al. (2020) performed a high angular resolution
survey of 74 O and early B-type stars in Cyg OB2 and found that 47% of targets
had a resolved companion with a $<$1% chance of not being a companion to the
primary star. They noted a particularly large number of targets with wide
(100–100,000 AU) binary companions. Including the known spectroscopic binaries
from Kobulnicky et al. (2014), Caballero-Nieves et al. (2020) calculated lower
limits to the multiplicity fraction and companion frequency of $0.65\pm 0.05$
and $1.11\pm 0.13$, respectively.

Hanson (2003) noted that many of the OB stars discovered at that time appeared
to be older than the previously-accepted age of $\sim$2 Myrs for the
association, and that these stars might represent a ‘halo’ population
distributed over a wider area than the stars studied by Massey & Thompson
(1991). Recent studies have widened the area surveyed around Cyg OB2,
uncovering many more massive stars over a wider area that are typically older
than those in the centre of the association (e.g., Comerón et al., 2008;
Comerón & Pasquali, 2012; Berlanas et al., 2018). While some of these may be
related to the nearby Cygnus OB9 association, it does suggest that massive
star formation has been occurring in Cygnus long before the formation of Cyg
OB2 and that an older population of massive stars now exists across the area
(Comerón et al., 2008, 2016).

####  3.4.2 Age and age spread

Early estimates of the age of the high-mass stars in Cyg OB2 suggested an age
of $2\pm 1$ Myr, based on the presence of a well-defined upper main sequence
and O5V and O5.5V stars (Massey & Thompson, 1991; Hanson, 2003). The presence
of numerous evolved giant and supergiant stars in the area was seen as
contamination from a non-coeval population within the region (Hanson, 2003).
Later studies of the A-type star population across the region (Drew et al.,
2008) and X-ray studies of the low- and solar-mass population in relatively
compact areas of the sky (Wright et al., 2010c) suggested ages of 5–7 and 3–5
Myr, respectively, implying an older population.

Later studies of the high-mass stars in Cyg OB2 have not distinguished between
any young ‘central’ group and an older more dispersed population (Comerón &
Pasquali, 2012). The large census of OB stars by Wright et al. (2015) showed
that when stellar ages for the high-mass stars are calculated for more-
realistic rotating stellar models (rather than the classical non-rotating
models) the ages of these stars shift from 1–4 Myr to 2–6 Myr. Furthermore,
when the sample is limited to stars in the mass range of 20–35 M⊙ (at which
masses no stars should have exploded as supernovae within 6 Myr, Ekström et
al., 2012), then the star formation history appears broadly constant over the
last $\sim$6 Myr, suggesting more or less continuous star formation over this
time period across the association. Combining new spectroscopic, Gaia
parallaxes, and by comparison with evolutionary models, Berlanas et al. (2020)
studied the age distribution of the O-type stars in Cyg OB2 and found evidence
for two possible star-forming bursts at $\sim$3 and $\sim$5–6 Myr.

####  3.4.3 Structure and kinematics

In their photometric study of Cyg OB2, Knödlseder (2000) identify the
association as spherically symmetric and suggest it to be considerably more
massive than previously suspected, with a total mass of (4–10) $\times 10^{4}$
M⊙. They argued that Cyg OB2 should not be classified as an OB association,
but as a young globular cluster or super star cluster, but the size of the
association ($\sim$30 pc) weakens the argument it is an analogue of the super-
star clusters found in other galaxies. A number of studies have found the low-
(Drew et al., 2008; Vink et al., 2008) and high-mass (Comerón et al., 2008)
stellar population extends further from the association than previously
recognised, particularly towards the south.

Wright et al. (2014b) performed the first quantitative study of the 2D
structure of Cyg OB2 by using statistical diagnostics to measure substructure
and mass segregation in the association, finding considerable physical
substructure but no evidence for mass segregation. Both of these are
indications that the association is not dynamically evolved (as mixing and
two-body interactions destroy substructure and can lead to mass segregation
through energy equipartition). Gaia DR2 parallaxes allowed Berlanas et al.
(2019) to extend this work to the third dimension, revealing that the known OB
stars constitute at least two separate groups along the line-of-sight. Lim et
al. (2019) studied the spatial distribution of high-mass members of Cyg OB2
and suggest the association is comprised of two dense clusters and a lower-
density halo, though the clusters do not appear to be kinematically distinct
and there is no evidence for such clusters in the spatial distribution of low-
mass stars (e.g., Wright et al., 2014a).

Kiminki et al. (2007) performed the first kinematic study of Cyg OB2,
measuring radial velocities for 146 OB stars over a 6 year period, allowing
them to identify binaries and measure systemic velocities. They derived a
radial velocity dispersion of $8.03\pm 0.26$ km s-1 (see Kiminki et al., 2008)
from the systemic velocities of the systems studied (and therefore no further
correction due to binarity was applied, noting that this was larger than found
for most (typically less-massive) OB associations and open clusters.

Wright et al. (2016) conducted a proper motion study of $\sim$900 low- and
solar-mass stars in Cyg OB2, finding dispersions of $13.0^{+0.8}_{-0.7}$ and
$9.1\pm 0.5$ km s-1 in RA and declination, respectively. In addition to
producing a 3D velocity dispersion of $17.8\pm 0.6$ km s-1 (implying that the
association is gravitationally unbound) this also revealed the velocity
dispersion to be significantly non-isotropic. Wright et al. (2016) searched
for evidence of expansion in Cyg OB2, but could find no indication of a global
expansion pattern. They also found that their proper motions exhibited
significant kinematic substructure, as shown in Figure 14, which they argued
echoed the physical structure already observed (Wright et al., 2014b) and also
implied that the association was dynamically unevolved. They estimated that
the kinematic subgroups they observed were close to virial equilibrium and had
typical masses of $\sim$40–400 M⊙.

####  3.4.4 Feedback and the Cygnus Super-Bubble

The Cygnus Super-Bubble, first detected by Cash et al. (1980), is a large
($18^{\circ}\times 13^{\circ}$) soft X-ray emitting region that is roughly
centred on Cyg OB2 (see also Bochkarev & Sitnik, 1985). The bubble is believed
to be a single massive shell, powered by both Cyg OB2 and possibly the other
OB associations in the region. Murray & Rahman (2010) associate one of their
radio free-free emission sources with Cyg OB2 and suggest the source is
powered by the O stars in the association, similar to the bubble.

Whether the Cygnus super-bubble is a real, single structure has been a matter
of debate for many decades. Uyanıker et al. (2001) studied the Cygnus super-
bubble and argued that the X-ray emission observed was actually due to a
superposition of separate regions at different distances. Albacete Colombo et
al. (2018) traced the diffuse X-ray emission across Cyg OB2 that arises from
stellar winds from the O-type stars in the association, which showed evidence
for the interaction of these winds with the surrounding interstellar medium,
and suggest that at least this emitting region is constrained to the area
immediately around the association. Kimura et al. (2013) have however shown
that the large-scale X-ray emitting regions in the super-bubble have similar
absorbing columns and temperatures that suggests they originate from a single
entity and therefore that the bubble is one large structure. Taken together
these results imply that if the Cygnus super-bubble is a single structure with
a single origin then it might not be related to Cyg OB2 but to a previous
generation of massive stars in the Cygnus complex.

Both Cash et al. (1980) and Kimura et al. (2013) have estimated the energy
within the bubble to be $\sim$1052 erg. Cash et al. (1980) estimate, based on
the radiative criterion of Cox (1972), that if the $\sim$450 pc diameter shell
were produced by a single, very-powerful supernova then this would require an
injection of $10^{54}$ erg to produce such a structure, larger than currently
observed in the bubble. Instead they suggest that if the energy were injected
steadily over a longer period by a series of 30–100 SNe, then the energy
required would be lower. Such a number of recent SNe would not be impossible
over the entire Cygnus complex, though it is considerably more than could have
originated from Cyg OB2 itself (Wright et al., 2015). Kimura et al. (2013)
suggested that this level of energy could instead be injected by 2–3 Myr of
mechanical wind energy from the massive stars in Cyg OB2, or alternatively
from a hypernova, the explosion of a massive star composed mainly of carbon
and oxygen, that Iwamoto et al. (1998) have hypothesised could release a
sufficient amount of energy. Higdon & Lingenfelter (2013) estimate that the O
stars in Cyg OB2 could produce an ionizing flux of $10^{51}$ s-1 (40 times
larger than from the Orion OB1 association), irradiating the bubble and the
surrounding Cygnus X molecular cloud.

###  3.5 Other prominent or well-studied associations

####  3.5.1 Perseus OB2

Figure 15: Density map of probable members of the Perseus OB2 association from
Belikov et al. (2002a) with the main area of the association shown with a
rectangle. The greyscale shows the surface density of members while the black
solid lines show contours of CO emission. Also shown are the outline of the
California Nebula (dashed line), the position of the young embedded cluster IC
348 (plus sign) and the approximate division within the association of the
subgroups $a$ and $b$.

The Perseus OB2 association was first identified by Blaauw (1944) and is part
of a series of star formation events in Perseus that includes multiple young
clusters in the Perseus molecular cloud and the quiescent California Nebula
(for a review see Bally, 2008). While it is amongst the closest OB association
to the Sun ($d=296\pm 27$ pc, de Zeeuw et al., 1999) it has significantly
fewer massive stars compared to Sco-Cen, Orion or Vela, and therefore has not
received as much attention. de Zeeuw et al. (1999) identified 17 B-type
members of the association, including $\zeta$ Per (B1Ib) and 40 Per (B0.5V),
the two brightest members of the association. The runaway O7IIIe star $\xi$
Per is thought to have been ejected from the association following a supernova
explosion (Blaauw, 1961), which may have played a role in powering a shell
observed to be surrounding the association (Sancisi et al., 1974). The total
mass of the association has been estimated as $\sim$6000 M⊙ (Belikov et al.,
2002a).

Belikov et al. (2002b) and Azimlu et al. (2015) have identified many hundreds
of candidate low- and intermediate-mass members of the association, tracing
out its distribution from the classical region where the high-mass stars can
be found out as far as the nearby California Nebula and the Auriga dark cloud
(see Figure 15. The association is also visible in the pre-MS density map
produced by Zari et al. (2018), but not in their upper MS density map due to
the low number of massive stars. While almost spherical with a diameter of
$\sim$40 pc there is evidence for multiple subgroups, both structurally
(Herbig, 1998; Belikov et al., 2002a) and from the age distribution of stars
(Azimlu et al., 2015). Belikov et al. (2002a) also observed an anti-
correlation with the distribution of dust clouds in the almost co-spatial
Perseus molecular cloud, which may suggest an extinction bias or that the
association is excavating a cavity within the molecular cloud.

The association is approximately co-spatial with the Perseus molecular cloud
and its embedded clusters IC 348 and NGC 1333 (Ortiz-León et al., 2018), and
these are often considered part of Per OB2 given that they each contain two
young B-type stars. Azimlu et al. (2015) trace a population of young stars
aged $\sim$1 Myr within these star forming regions, as well as older
populations aged 1–5 and $>$5 Myr, with the latter concentrated towards the
north-west of the association.

There have been very few kinematic studies of Per OB2. Steenbrugge et al.
(2003) calculated a radial velocity dispersion of $\sim$3.9 km s-1 from the
known high-mass stars, but this was not corrected for the influence of binary
motions, which will inflate the dispersion. Blaauw & Morgan (1953) and Lesh
(1969) used proper motions from photometric plates to assess evidence for
expansion within the association, calculating an expansion age of $\sim$1.3
Myr, but Belikov et al. (2002a) could find no evidence for expansion from
Hipparcos data.

####  3.5.2 Carina OB1

Despite being distant, the Carina OB1 association is prominent thanks to its
very large number of OB stars (e.g., Massey & Johnson, 1993; Walborn, 1995)
and bright Hii regions. The brightest of these is NGC 3572, the Carina Nebula,
which sits broadly at the centre of the $4\times 4$ degree area covered by the
association’s members (Humphreys, 1978). The total mass of the association is
thought to exceed $2\times 10^{4}$ M⊙ (Preibisch et al., 2011). The
association includes several massive star clusters (Trumplers 14, 15, & 16,
NGC 3293 & 3324, Bochum 10 & 11, and Collinder 228222Collinder 228 may appear
as a separate cluster, but it is believed to be part of Trumpler 16, divided
from it by an obscuring foreground dust lane. & 232), as well as a distributed
population of young stars (Feigelson et al., 2011). The association, and the
Carina Nebula itself, sit at a distance of $2350\pm 50$ pc (Smith, 2006). The
edge of the association is also projected against the massive star cluster
Westerlund 2, though this is significantly more distant.

There are at least 70 O-type stars across the association, the majority of
which are found in the dense clusters of the Carina Nebula. Of these, Trumpler
16 is by the wealthiest, home to multiple very massive stars such as $\eta$
Carina (a luminous blue variable), the very early-type stars HD 93205
(O3V+O8V) and Tr16-244 (O3/4 If), as well as many late O-type stars (Smith &
Brooks, 2008). Trumplers 14 and 15 are also relatively massive and contain
$\sim$16 O-type stars between them, the latter containing the O2 supergiant HD
93129A (Walborn et al., 2002). The association contains three late-type WN
Wolf-Rayet star members (WR22, WR24 & WR25, Crowther et al., 1995), and Rate
et al. (2020) suggest that WR18 and WR23 may be possible members of the
association as well. In addition to these there are a number of candidate OB
stars awaiting spectroscopic confirmation (e.g., Povich et al., 2011; Mohr-
Smith et al., 2017).

There are stars and clusters in the association with ages ranging from 1–10
Myrs. The youngest parts of the association are centred around the Carina
Nebula, including the clusters Tr 14, 15 & 16, aged approximately 1, 5, and 3
Myrs, respectively (Smith & Brooks, 2008; Hur et al., 2012). Trumpler 14, the
youngest of the clusters, is kinematically associated with the molecular gas
in its vicinity, consistent with its young age (Kiminki & Smith, 2018). Beyond
the Carina Nebula, there are some young clusters such as NGC 3324 (1–2 Myr,
Preibisch et al., 2014), as well as older clusters like NGC 3293 (8–10 Myr,
Preibisch et al., 2017).

Kiminki & Smith (2018) performed a radial velocity survey of the Carina
Nebula’s O-type stars and measured a velocity dispersion of $<9.1$ km s-1 (not
corrected for binaries), consistent with those measured for other similarly-
sized and substructured OB associations such as Cyg OB2. The association, and
particularly the well-studied young stellar population in the vicinity of the
Carina Nebula, is highly substructured, both spatially (Feigelson et al.,
2011; Buckner et al., 2019; Reiter & Parker, 2019), kinematically (Kiminki &
Smith, 2018), and temporally (see previous paragraph). Mel’nik & Dambis (2017)
measured the expansion of the association for the first time, deriving a
kinematic age of $\sim$8 Myrs along both proper motion axes, consistent with
the maximum age of stars in the association.

####  3.5.3 Canis Major OB1

Ambartsumian (1949) was the first to propose the existence of an association
of early-type stars in Canis Major, which is thought to contain over 200
B-type stars and a few O-type stars (Gregorio-Hetem, 2008) at a distance of
$\sim$1.2 kpc (Zucker et al., 2020). The region is dominated by the large H ii
region Sh2-296 (The Seagull Nebula), as well as numerous bright-rimmed clouds
and an expanding shell of emission nebulae that Reynolds & Ogden (1978)
suggest was produced by stellar winds or supernovae (Fernandes et al. 2019
suggest that Sh2-296 is part of a 60 pc diameter shell that surrounds CMa OB1
that itself may be nested within a larger, 140 pc superbubble that is visible
in H$\alpha$).

Within the boundaries of the association can be found the CMa R1 association
and at least 19 open clusters (e.g., Dias et al., 2002; Kharchenko et al.,
2013), with the 7.6 Myr open cluster NGC 2353 (Fitzgerald et al., 1990)
thought to be the nucleus of the association (Clariá, 1974). The brightest and
earliest stars in the association are HD 54662 (O7III) and HD 54879 (O9.5V),
as well as many early B-type stars (Clariá, 1974). There have been multiple
runaway stars connected to the association, including HD 54662 (Herbst &
Assousa, 1977, O7V), HD 57682 (Comeron et al., 1998, O9V) and HD 53974
(Fernandes et al., 2019).

Recent surveys have begun to uncover the low-mass population of the
association. Fischer et al. (2016), Santos-Silva et al. (2018) and Pettersson
& Reipurth (2019) used infrared, X-ray and H$\alpha$ observations,
respectively, to uncover hundreds of YSOs across the association and in nearby
star forming regions. Fischer et al. (2016) find the spatial distribution of
young stars in the association to be highly clumpy, with multiple groups and
clusters within it. Santos-Silva et al. (2018) find that the age distribution
of the young stars is bimodal, with an older population aged $\sim$10 Myr and
a younger population localised to the regions of molecular gas, in agreement
with the discovery by Sewiło et al. (2019) of a new star forming region within
the association with $\sim$300 very young stars within it.

####  3.5.4 Monoceros OB1 and OB2

The Mon OB1 and OB2 associations are projected close to each other in the
constellation Monoceros at distances of $\sim$600 and $\sim$1200 pc,
respectively. They each include a bright H ii region, which is the youngest
part of each association, the Cone Nebula, the young cluster NGC 2264 and the
R association Mon R1 in Mon OB1, and the Rosette Nebula and its central
cluster NGC 2244 in Mon OB2. Of the two associations Mon OB2 is richer and
more massive, with at least 13 O-type stars (e.g., Johnson, 1962; Humphreys,
1978), including 5 in the cluster NGC 2244 that are responsible for
illuminating and shaping the Rosette Nebula (Wareing et al., 2018). These
stars represent the most recent epoch of star formation in Mon OB2, with
evidence for other subgroups with ages of $\sim$5 and $\sim$15 Myr (Turner,
1976; Blitz & Thaddeus, 1980). Mon OB1 is less massive, its most massive
member being the O7V star S Monoceros in the young young NGC 2264 cluster,
though the region also shows evidence for multiple subgroups of different ages
(Flaccomio et al., 1999). There is also evidence for an older population of
stars across Mon OB1 and a possible indication of post supernova activity from
the presence of six arc-like molecular structures around the association
(Oliver et al., 1996).

####  3.5.5 Scorpius OB1

The Sco OB1 association is located on the near side of the Sagittarius spiral
arm at a distance of $\sim$1.5 kpc (van Genderen et al., 1984; Sung et al.,
1998; Yalyalieva et al., 2020), and includes the massive young cluster NGC
6231 at its core, the open cluster Trumpler 24, and the Hii regions
G345.45+1.50 and IC 4628. The association contains a large number of OB stars,
with 28 OB giants and supergiants listed in the census of Humphreys (1978),
including two of the most luminous OB stars known, HD 151804 (O8Iaf) and HD
152236 (B1.5Ia), as well as multiple Wolf-Rayet stars.