Download Star Formation in the Orion Nebula II: Gas, Dust, Proplyds and

Handbook of Star Forming Regions Vol. I
Astronomical Society of the Pacific, 2008
Bo Reipurth, ed.
Star Formation in the Orion Nebula II: Gas, Dust, Proplyds and
Outflows
C. R. O’Dell
Department of Physics and Astronomy, Vanderbilt University, Box 1807-B,
Nashville, TN 37235, USA
August Muench
Harvard-Smithsonian Center for Astrophysics
60 Garden Street, Cambridge, MA 02138, USA
Nathan Smith
Astronomy Department, University of California at Berkeley,
601 Campbell Hall, Berkeley, CA 94720, USA
Luis Zapata
Max Planck Institute for Radio Astronomy,
Auf dem Hügel 69, 53121 Bonn, Germany
Abstract.
The visually familiar Trapezium cluster is but one of three centers of
recent star formation in the Orion Nebula, with the other two still embedded in its host
molecular cloud. The Orion Nebula was produced when the hottest stars in the Orion
Nebula Cluster photoionized local gaseous material, forming an open cavity around the
Trapezium stars, with a background blister of ionized gas, then a photon dominated
region beyond that. On the near side there is a neutral veil of material. The cluster
members include many proplyds, young stellar objects that are rendered more visible
by being in or near an H II region. Their existence is an argument that the most massive
stars in the cluster formed only recently. The second-most luminous star formation
center is in the BN-KL region and is embedded in the molecular cloud, which means
that it is seen only in X-ray, infrared, and radio wavelengths. There are arguments that
it experienced a major energetic event 500–1000 years ago, producing runaway objects
and a host of expanding fingers of gas and dust. The third center of star formation,
Orion-S, lies only slightly behind the photon dominated region and produces multiple
outflows, most of which are bipolar, and are seen in molecular and ionized atomic
emission. The proximity of the Orion Nebula and its conditions of low extinction mean
that it is the richest region of coll ejecta from pre-main sequence low-mass stars.
1.
Introduction
This article extends the preceding article into additional subjects, drawing on observational results from over the energy range X-rays through radio waves. In Sect. 2 we
treat the structure of the Orion Nebula, its underlying photon dominated region (PDR),
and host molecular cloud. In Sect. 3 we consider the proplyds, that component of the
stellar population where the natal material is seen because of the newly formed star
being in or near an H II region. In Sect. 4 we discuss the embedded stars, i.e. the
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2
stars not optically visible because of the high extinction caused by these objects falling
within the PDR or the molecular cloud. In Sect. 5 we consider the outflows encountered, which range from microjets with scales of hundreds of astronomical units to large
scale outflows with scales of parsecs.
Because of the diversity of backgrounds of the contributors to these subjects, there
is an overabundance of nomenclatures and methods of indicating positions. In this
article we try to use a uniform system of designation of objects and most positions
of sources are indicated in 2000.0 coordinates. Radio results are often still reported
in 1950.0 coordinates and the reader should bear in mind that precession means that
1950.0 coordinates of the center of the Orion Nebula Cluster are 147.4s of Right Ascension smaller and 113.3′′ further south than 2000.0 coordinates. Since radio results
are often presented as Local Standard of Rest (LSR) velocities, we note that for the central region, LSR velocities are 18.1 km s−1 more negative than heliocentric velocities.
We have benefited from an early case study of this region (Goudis 1982) and the
recent reviews of the BN-KL region (Genzel & Stutzki 1989) and of the Orion Nebula
(O’Dell 2001a,b). The reader is referred to these reviews for summaries of early work
not treated in this article.
The global structure of this region is that of a molecular cloud in which an extended
period of star formation has occurred on the side facing the observer, producing the
Orion Nebula and the Orion Nebula Cluster (ONC). The recently formed hot stars of the
ONC photoionize the surrounding gas. The gas is only partially confined by overlying
neutral material, resulting in an ionization front developing on the face of the molecular
cloud. The shock produced by the ionization front creates a dense PDR between the
front and the molecular cloud. On the side of the observer there is the neutral residual
of the original confining material. Two additional centers of star formation also exist,
both are embedded in the molecular cloud. The more luminous region is associated with
the BN-KL region to the northwest from the Trapezium stars and the less luminous is
associated with the Orion-S region to the southwest. There is no evidence for a large
motion of the ONC stars (V⊙ =25±2 km s−1 , (Sicilia-Aguilar et al. 2005) and V⊙ =26.1
km s−1 with a dispersion of 3.1 km s−1 , (Fúrész et al. 2008)) with respect to the
parent molecular cloud . The molecular cloud’s velocity has been reported as V⊙ =27±2
km s−1 (O’Dell & Wen 1994) and is V⊙ = 25.8 ±1.7 km s−1 when determined from
published velocities of heavy molecules (Goudis 1982).
2.
The Orion Nebula and its Underlying Structure
The high surface brightness Orion Nebula is frequently the target of new observational
techniques and telescopes, a pattern that continues. This is due to its brightness, its
proximity, and its location in the sky (high elevation for a large fraction of the northern hemisphere winter night). The well studied bright inner few minutes of arc of the
object is commonly called the Huygens region, while associated fainter features extend beyond 10′ in several directions. It is one of the most frequently imaged parts
of the sky and its appearance is quite different depending upon the selection of filters,
as narrow-band filters can isolate emission lines from over a wide range of ionization
stages, and intermediate-band filters are often dominated by scattered light from dust
in the PDR. Arguably, the astrophysically most useful ground-based images are those
of Pogge (Pogge et al. 1992), while Hubble Space Telescope (HST) image are now
available (O’Dell & Wen 1994; O’Dell & Wong 1996; Bally et al. 2000; Henney et al.
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Figure 1. This 30′ x30′ image of the Orion Nebula is a composite of many HST
Advanced Camera for Surveys images supplemented on the edge by ground-based
images (Henney et al. 2007). The color coding is red/orange=Hα+[N II], red=nearIR F850LP and F775W filters, green=F555W V filter, blue=F435W wide blue filter.
2007) at resolutions down to better than 0.1′′ as shown in Figure 1. The radio continuum has now been mapped with a resolution of about 1.5′′ (O’Dell & Yusef-Zadeh
2000) and limited surveys to higher surface brightness have approached HST resolution
(Churchwell et al. 1987; Garay et al. 1987; Felli et al. 1993; Zapata et al. 2004a).
2.1.
Structure of the Huygens Region
The Ionized Layer. After early disputes (Osterbrock & Flather 1959; Münch 1958;
Wurm 1961; Münch & Wilson 1962) about the 3-D form of the Huygens region, the
pattern of the radial velocity becoming more blue-shifted with increasing ionization of
the nebular gas (Kaler 1967) provided the necessary proof that this part of the nebula
is a thin blister (Zuckerman 1973; Balick et al. 1974). This gas is flowing away from
the dense PDR that separates the ionized zone from the background host molecular
cloud. Ionization is dominated by θ 1 Ori C, which lies at about 0.3 pc in front of the
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ionization front. A 3-D model of the ionized gas (Wen & O’Dell 1995) shows that
the surface of the ionization front is highly irregular, being farther from the observer
to the east of the Trapezium stars (at the location of the classical Dark Bay feature),
has a local hump to the south-west of the Trapezium (at the location of the Orion-S
feature), and becomes almost perpendicular to the plane of the sky to the south-east
and gives rise to the long, nearly linear Bright Bar feature (Balick et al. 1974). The
density of ionized gas decreases from its peak electron density of about 104 cm−3 at the
ionization front as the gas accelerates away from the front (Henney, Arthur & Garcı́aDı́az 2005). The scale height of the emitting gas can be estimated from measurements
of the emission measure and electron density and shows considerable variation across
the core of the nebula (Wen & O’Dell 1995; Garcı́a-Dı́az & Henney 2007). In the
brightest part of the nebula, to the west of the Trapezium, the emitting layer is thin
(< 0.05 pc), whereas the emitting region to the east of the Trapezium has a much
greater thickness (≃ 0.3 pc), which is comparable to its lateral extent and to the distance
of θ 1 Ori C from the ionization front. The stationery shock fronts seen in front of most
of the proplyds near θ 1 Ori C (Bally et al. 2000) imply that there is a cavity formed by
the high velocity wind coming from this luminous star. The presence of narrow He I
absorption lines in the spectra of the Trapezium stars (Baldwin et al. 1991; O’Dell
et al. 1993; Wilson 1937; Adams 1937) indicates that there is low density ionized gas
in the vicinity of the center of the Orion Nebula Cluster and the [S II] measurements of
Garcı́a-Dı́az & Henney (2007) indicate a low density region lies in the southeast of the
Huygens region. The latter may be a portion of a parsec scale emission component first
posited by Deharveng (1973). Since the scale height for the distribution of the stars is
about 0.18 pc (Hillenbrand & Hartmann 1998), a fraction of the cluster stars fall within
the region of dense ionized gas and some even beyond the PDR (Herbig & Terndrup
1986).
Radial velocities have been determined across the face of the Huygens region with
spatial resolutions of a few seconds of arc and velocity resolutions of better than 10
km s−1 (Doi et al. 2004; Henney et al. 2007; Garcı́a-Dı́az & Henney 2007; Garcı́aDı́az et al. 2008). In the central portion of the Huygens region the tracers of emission
from in or near the ionization front itself ([O I] and [S II]) are at V⊙ = 25.5 ±1.5 km s−1 ,
while the emission from the higher ionization zones ([O II], [N II], [O III], H II, He I,
and[Cl III] are at V⊙ =18.2±1.4 km s−1 (O’Dell 2001a).
Properties of the Foreground Veil. Most of the extinction in the ONC stars and the
emission from the nebula occurs in layers of primarily neutral material lying in front of
the nebula. Collectively, these layers are known as the Veil and were originally discovered as 21 cm absorption lines in continuum emission from the nebula (van der Werf
& Goss 1989). The two principal components have velocities of V⊙ = 21 km s−1 and
24 km s−1 . The correlation of the 21 cm absorption column density and the extinction
is now well established (O’Dell et al. 1992; O’Dell & Yusef-Zadeh 2000). The Veil
has been studied in detail by means of optical and ultraviolet absorption lines (O’Dell
et al. 1993; Abel et al. 2004, 2006). Zeeman splitting of the 21 cm lines has been measured (Troland et al. 1989) and a thorough analysis of the conditions (Abel et al. 2006)
concludes that the magnetic field energy is greater than the thermal energy in one of
the two strong components of the Veil. Early arguments that the bulk of the extinction
occurs within the nebular gas (Gómez Garrida & Münch 1984; Münch 1985) were
later shown to be incorrect (O’Dell 2001a; O’Dell 2002). The Orion Nebula Cluster
was one of the first regions in which it was established that some stars associated with
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H II regions have an anomalous reddening curve (Baade & Minkowski 1937; Costero
& Peimbert 1970; Cardelli & Clayton 1988; Greve et al. 1994; Blagrave et al. 2007).
This means that the dust in the Veil has a different size distribution than that in the interstellar medium, probably having fewer small particles (O’Dell 2001a), whether this
is due to grain growth or destruction of small grains is undetermined.
The Background PDR. The PDR has a density of at least 105 atoms/cm3 (Tielens &
Hollenbach 1985) and apparently the dust density is correspondingly enhanced. The
PDR has a large optical depth at visual wavelengths and starlight scattered from it
accounts for the fact that the observed visual continuum is about five times stronger
than that expected from an atomic continuum (O’Dell & Hubbard 1965; Baldwin et al.
1991). One also sees the effects of scattering of emission line radiation. Red-shifted
components of the strongest emission lines (O’Dell et al. 1992; Henney 1998) arise
from the velocity difference between the scattering PDR dust and the emitting ionized
regions. It is also observed that the [O III] emission is polarized (Leroy & Le Borgne
1987). As measured by the CO and [C II] radio emission (O’Dell 2001a; Goudis 1982)
the radial velocity of the central region of the PDR is V⊙ = 28± 1.5km s−1 , placing it at
indistinguishably the same velocity as the ionization front material and the background
molecular cloud.
2.2.
The Extended Orion Nebula
The large bounded region to the southwest of the Huygens region that can be designated
as the Extended Orion Nebula (EON) has been the subject of much less investigation
because of its much lower surface brightness. This is a elliptical section of major axes
about 31.9′ x 29.7′ with the long axis lying at about a position angle of 20o with the
Huygens region near its northeast corner. It has an irregular boundary. From a comparison of optical and radio observations Subrahmanyan et al. (2001) determined that the
density in the southwest portion of the EON is about 30 cm−3 .
Recent work (Güdel et al. 2008) has determined that two regions within the
EON contain gas at about 2x106 K. The hot region lying to the west-southwest of the
Trapezium has a density of about 0.2-0.5 cm−3 and the region to the southwest a density
of about 0.1-0.2 cm−3 . This material appears to be heated by the high velocity stellar
wind from θ 1 Ori C. Extended X-ray emission was not detected in an earlier study with
the Chandra Observatory (Townsley et al. 2003), which did cover the west-southwest
source. The absence of detection in the Huygens region may be caused by the expected
high X-ray optical depth caused by the Veil or it may be obscured by the instrumental
scattered light from θ 1 Ori C, which is very bright in X-rays.
2.3.
The Molecular Cloud
The Orion Nebula appears physically associated with the northern part of the Orion A
molecular cloud. This portion of the cloud is commonly referred to as the “integral”
or S-shaped filament (Bally et al. 1987), and the molecular and dust emission in the
S-shaped filament peaks at and is centered behind the Trapezium in the Orion Nebula.
Figure 2 illustrates the structure of this filament as traced in 850 µm dust continuum and
compared to the mid-IR emission from the nebula. The S-shaped filament is divided
commonly into four “clumps,” which are named Orion Molecular Cloud (OMC) # 14. The OMC-1 and OMC-4 clumps are the subject of this summary; the OMC-2 and
OMC-3 clumps are reviewed in the Peterson & Megeath chapter. The OMC-1 cloud is
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Figure 2. This MSX+SCUBA dust map of the Orion Molecular Cloud has the
following color coding:Red= SCUBA 850 µm, Green=MSX 14 µm, B= MSX 8 µm
(PAH) (Johnstone & Bally 1999; Kraemer et al. 2003). This figure has an angular
size of 0.9◦ by 1.2 ◦ and is centered at 5:35:23.3 -05:17:08 (J2000).
centered behind the Trapezium stars, while OMC-4 is a v-shaped group of Submillimeter cores 10′ -15′ south of OMC-1. We review here the recent work on molecular and
atomic line as well as dust continuum tracers of OMC-1 and OMC-4. Somewhat more
detailed summaries of observations of the twin sub-clumps in OMC-1, BN-KL and
Orion-S (this last feature sometimes being designated as OMC1-S), are given in Sect.
4.1. and Sect. 4.3., respectively. We again direct the reader to the review of Goudis
(1982), which provides a more exhaustive account of the wealth of observations of
Orion A during the advent of radio astronomy.
Tracers of the Molecular Gas. Very wide-field observations covering the northern
Orion A cloud in the rotational transitions of 12 CO and 13 CO were undertaken by Kutner et al. (1977), Maddalena et al. (1986), Bally et al. (1987), Castets et al. (1990),
Heyer et al. (1992), Sakamoto et al. (1994), White & Sandell (1995), Plume et al.
(2000), and Wilson et al. (2005). All of these survey clearly show the S-shaped filament of molecular gas stretching for 1◦ (∼ 10 pc) length and exhibiting a north-south
velocity gradient from VLSR = 4 km s−1 in the south to VLSR = 12 km s−1 in the north.
The first (ever) 140 GHz formaldehyde (H2 CO) map by Thaddeus et al. (1971) revealed
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what is now the familiar two peaked ridge of OMC-1 and the 4.8 GHz study of the l10 l11 line (Mangum et al. 1993) mapped the BN-KL and Orion-S regions . Complete
maps in 12 CO, 13 CO and CS of OMC-1 were first provided by Liszt et al. (1974) and in
HCN by Clark et al. (1974); all clearly established the structure of this molecular ridge.
More recent higher resolution HCN, CS and C18 O studies of the OMC-1 core include
Bergin et al. (1996) and Goldsmith et al. (1997). A fairly wide range of molecules were
used by Ungerechts et al. (1997) to examine chemical and physical variations along the
OMC-1 ridge. Most recently, the Orion A cloud has been mapped (Tatematsu et al.
2008) in N2 H+ and HC3 N.
Almost like the CO observations, surveys of the thermal emission of CS show a
more clumpy morphology toward the Orion A (Lada et al. 1991; Tatematsu et al. 1993,
1998). The CS observations also showed that the gas density of the CS cores tends
to be lower in the south of the cloud than in the north as also observed by the CO
surveys (e.g., Bally et al. 1987). Observations of formaldehyde have also confirmed
this structure toward Orion A (Cohen et al. 1983).
Ammonia: Unresolved radial velocity measures from the dense gas traced by ammonia first indicated the presence of a background ridge, sometimes referred to as the
plateau and an unresolved hot core within the OMC-1 cloud near the Orion BN-KL
region. Modeling the filling factor for these large (>1′ ) beam observations led to a
conclusion that the flux came from 1 or more small clumps of the order of 0.04pc (Barrett et al. 1977; Ho et al. 1979; Bastien et al. 1981; Ziurys et al. 1981). Observations
appearing to resolve these inferred clumps include Pauls et al. (1983) and Migenes
et al. (1989), with additional interpretation provided by Genzel et al. (1982). Murata
et al. (1990) observed ammonia at 8′′ resolution near the BN-KL region and found
filamentary structures which turn out to be the bases of much larger ammonia filaments revealed by interferometric VLA observations (Wiseman & Ho 1996, 1998) as
shown in Figure 3. Wiseman et al. found that these fingers of ammonia protrude from
the BN-KL core over scales of >0.5 pc, and display indications of having been externally heated. Given the high densities and temperatures and the subsequently large
surface brightness, the Orion molecular cloud has been a valuable target for measuring
metastable line ratios of ammonia (Barrett et al. 1977; Sweitzer 1978; Sweitzer et al.
1979; Townes et al. 1983; Hermsen et al. 1988). Additional submm ammonia transitions include detections by Schilke et al. (1992).
Pioneering far-infrared and sub-millimeter surveys of the Orion Nebula (Soifer &
Hudson 1974; Fazio et al. 1974; Werner et al. 1976; Hudson & Soifer 1976; Smith et al.
1979; Keene et al. 1982; Thronson et al. 1986) had great difficulty resolving structure
in the cloud given the complex nature of the intervening H II region. While the higher
(32′′ ) resolution 400 µm observations of Keene et al. (1982) were able to clearly resolve
OMC-1 into two peaks, recent higher resolution (∼ 12′′ ) submillimeter continuum
observations at 350 µm by Lis et al. (1998) and at 450 and 850 µm by Johnstone &
Bally (1999) were the first to show the optically thin emission from interstellar dust
and constrain the grain temperature from the Orion A cloud. Chini et al. (1997)
mapped this region at six wavelengths from 350 to 2000 µm. Arimura et al. (2004) used
balloon borne 155 µm observations compared to a CO map from Tatematsu et al. (1998)
for examining variations in the dust-gas ratio along the OMC. The higher resolution
submm continuum images revealed a remarkable chain of compact sources embedded
in a narrow (0.14 pc), high column density filament that extends over the (7 pc) length
of the S-shaped filament as originally mapped for the first time in the CO observations.
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[l]
[r]
Figure 3. Left panel) Ammonia NH3 (1,1) velocity field. Right panel) Compared
to K band image. From Wiseman & Ho (1998).
Observations at 3.6 cm (Reipurth et al. 1999) found 14 sources in the OMC-2/3 region.
The brightest region, again associated with OMC-1, contains a remarkable group of
dust filaments that radiate radially away from the two high-luminosity cores. Those
filaments may be associated with remnent pre-stellar cores in the vicinity of the BNKL and OMC-1S regions (though they are much better resolved in the NH3 data of
Wiseman & Ho (1998)).
Water vapor observations in OMC-1 with the Submm Wave Astrophysical Satellite (SWAS) include Snell et al. (2000); Melnick et al. (2000), as modeled by Ashby
et al. (2000); SWAS failed to detect molecular oxygen in OMC-1 (Goldsmith et al.
2000). The new terahertz-line molecular observations carried out in very high and dry
places have opened a new window in ground-based radio astronomy. These THz observations centered at 1-1.5 THz or 300-200 µm contain numerous high-J CO spectral
lines that have revealed very hot molecular gas (a few 100 K) and high density molecular gas, n > 106 cm−3 toward the BN-KL and OMC-1S regions (Kawamura et al. 2002;
Marrone et al. 2004; Wiedner et al. 2006). The hot molecular gas has been suggested
to be energized by radiation from the embedded massive protostar(s) rather than from
interactions with outflows.
2.4.
The Photon Dominated Region
Between the ionization front that delineates the boundary of ionized hydrogen, within
which the optical emission lines are produced, and the cool molecular clouds lies the
PDR with unique conditions. This region sees the non-hydrogen-ionizing FUV radiation field and is affected by the large pressure gradient produced by the heating of the
gas within the H II region. This means that it is a unique region. The structure of the
ionization front is dominated by θ 1 Ori C except possibly in the region of the Bright Bar
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Figure 4. Integrated intensity images of (a) CI (3 P1 - 3 P0 at 492 GHz) and (b)
CO J=1-0 (Bally et al. 1987) detected toward the Orion region. The positions of
θ1 Ori C (southeast of KL), θ2 Ori A (southeast of θ1 Ori C), and l Ori are marked
by stars. Locations of the CI intensity peaks are denoted by triangles and (left panel)
letters for both panels. Taken from Ikeda et al. (2002).
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where θ 2 Ori A’s emission is comparably important. The FUV radiation from the cooler
Trapezium stars is important in the center of the Huygens Region and all components
of θ 2 Ori are probably important near the Bright Bar.
The Large-Scale PDR. In the millimeter/submillimeter regime, as a result of the
strong UV radiation field from OB stars, the thermal emission of the fine-structure
transitions of atomic carbon (CI) will be important. Since Orion A already formed
(and is forming) OB stars located around the Trapezium region, there is strong thermal
emission of CI associated with the Orion A cloud (White & Sandell 1995; Ikeda et al.
1999; Plume et al. 2000; Ikeda et al. 2002). The morphology and velocity fields of the
emission of the CI show a good spatial correlation with the 13 CO in most of the cloud,
which we illustrate in Figure 4.
The forbidden lines of [C II] at 158 µm is a ubiquitous signature of the PDR and
has been extensively mapped at moderate spatial resolution across the face of the entire
Orion Nebula (Boreiko et al. 1988; Stacey et al. 1993; Herrmann et al. 1997). It is of
particular interest that the highest velocity resolution study (Boreiko et al. 1988) found
10
a second component in [C II] emission in the samples to the east from the Trapezium,
where optically one sees the Dark Bay feature, the portion of the Veil with the highest
column density. The sample taken at 2′ east has a velocity of 2.45 km s−1 LSR (Boreiko
et al. 1988), which means that it probably arises from a secondary PDR formed on the
illuminating stars’ side of the Veil.
Observational and theoretical studies also have shown that the CN abundance is
enhanced in photo-dissociation regions (PDRs) (Fuente et al. 1993; Sternberg & Dalgarno 1995; Rodrı́guez-Franco et al. 1998, 2001). Since excitation of CN requires very
high density, this molecule is expected to be a good tracer of the highest density interface regions between the M 42 H II region and the molecular cloud. Those observations
indicate molecular hydrogen densities of ∼ 105 cm−3 for the OMC-1 molecular ridge
and densities of ∼ 3 × 106 cm−3 for the PDR directly behind the Trapezium stars
(Rodrı́guez-Franco et al. 1998, 2001).
There have been numerous recent studies of the PDR that take advantage of the
ability to compensate for atmospheric ”seeing” in the infrared. These give the ability to
address the small scale structure in this region, the H2 2.12 µm line being particularly
useful for this purpose (Lacombe et al. 2004; Kristensen et al. 2003, 2007; Colgan
et al. 2007; Kristensen et al. 2008; Gustafsson et al. 2003, 2006a,b; Nissen et al.
2007).
The Bright Bar. As indicated in Sect. 2.1, the Bright Bar feature of the optical image
of the Orion Nebula is the result of the H II ionization front curving up to be almost
perpendicular to plane of the sky. Since the nebular emission is concentrated towards
that ionization front, the amount of emitting material along the line of sight is greater
and the nebula is brighter there. The same will be true when examining the PDR emission that occurs to the southeast from the bright bar. This is a unique opportunity to
examine and test theories and models for the PDR, because one can hope to spatially
resolve there the different zones within the PDR whereas when looking at the rest of
the Orion Nebula one sees the PDR as sampled through all the zones.
The Bright Bar has been examined in a rewarding series of observations of increasing spatial resolution and wavelength coverage producing detailed studies of a wide
variety of PDR atoms, ions, molecules, and particles: HD (Wright et al. 1999), NH3
(Larsson et al. 2003; Batrla & Wilson 2003), H2 (Parmar et al. 1991; van der Werf
et al. 1996; Luhman et al. 1997, 1998; Habart et al. 2004; Allers et al. 2005) CN
(Simon et al. 1997), CS (Omodaka et al. 1986), C+ (Tauber et al. 1995; White &
Sandell 1995; Wyrowski et al. 1997), CO (Omodaka et al. 1994; Störzer et al. 1995),
Heavier Molecules (Sellgren et al. 1990; Hogerheijde et al. 1995; Fuente et al. 1996;
Young Owl et al. 2000; Lis & Schilke 2003), Silicates (Cesarsky et al. 2000), and
PAH (Roche et al. 1989; Giard et al. 1994; Bregman et al. 1994; Jansen et al. 1995;
Sloan et al. 1997; Kassis et al. 2006).
This wealth of information has led to a good first–order model for the PDR behind
the optical Bright Bar feature. In its simplest form it agrees with the expectation of
conditions driven by increasing optical depth to photons of less than 13.6 eV energy
(Tielens & Hollenbach 1985; Tielens et al. 1993), although time dependent effects may
be important (Bertoldi & Draine 1996; Störzer & Hollenbach 1998a). In order to fit the
observations, it is necessary to assume that the cloud material is quite clumpy (Tauber
et al. 1994; Gorti & Hollenbach 2002) and there are features that are in conflict with the
simplest (static) models (Marconi et al. 1998; Walmsley et al. 2000). Dynamic models
(Henney, et al. 2005) indicate that the narrow [N II] emission spike that is seen at the
11
Bright Bar is indicative of important advective effects in the local ionization front and
by extension in the PDR. The Bright Bar feature is probably a relic of the conditions
within the host molecular cloud, rather than something that has been initially created
by photoionization. This means that the PDR here is probably the best laboratory for
testing models of the Orion Nebula’s PDR.
3.
Proplyds
The Orion Nebula presents a unique opportunity for studying young stellar objects
(YSOs), i.e. young stars still partially surrounded by their natal material. One opportunity arises from the fact that when such an object is directly illuminated by an
ionizing star, such as θ 1 Ori C, all or a portion of the gaseous component will be ionized and that material will be visible in the same emission lines that one sees from the
nebula. The second opportunity occurs because most of the nebular emission arises in
the background, so that one can see the dust component in extinction against the nebular emission. When these YSO’s are all or partially photoionized then there can be
a large pressure excess and the material will be lost through photo-evaporation, thus
potentially destroying the envelope. The unique nature of their method of discovery
and study, together with the issues of survival imposed by their local environment has
led to designating these objects as a subclass of the YSOs and the name “proplyd” has
been applied (O’Dell & Wen 1994) in the first paper clearly describing these objects.
Because these objects are subject to discovery at different times by different methods,
a position-based system of nomenclature (O’Dell & Wen 1994) is usually adopted. In
this designation system the central nebula is divided into boxes of 0.1s in Right Ascension and 1′′ in Declination (Epoch 2000), and the first digits of the Right Ascension and
Declination, which are shared by all central region objects, are dropped (i.e. 5:35 -5:2).
This means that the proplyd lying at 5:35:17.67 -5:23:41.0 becomes 177-341. Although
this method avoids the confusion resulting from different designations in various publications, it has the disadvantage of being dependent upon exactly how and where within
the proplyd that the position is calculated, so that variations in the last digit are sometimes encountered. Table 1 gives a listing of various catalog designations of prominent
proplyds for which results were published prior to the adoption of a uniform system of
designation.
3.1.
Discovery and Subsequent Observations
Discovery. Identifying the discovery of the first proplyd depends upon how one wants
to define “discovery”. The first record as a star of an object subsequently identified as
a proplyd was for LV 2 (167-317). This object was the first star noted within the enclosed space of the Trapezium and was discovered on the first night (1888 January 7)
of operation of the Lick Observatory 36 inch refractor by Alvan G. Clark (Sheehan
1995). Subsequent visual and photographic surveys of the ONC included many objects
now known to be proplyds. A key work (Laques & Vidal 1979) discovered six unresolved emission line “stars” near the Trapezium in the high ionization [O III] 5007
Å line. These objects were among the compact bright thermal radio sources that were
discovered in Very Large Array studies of the inner Huygens region (Garay et al. 1987;
Churchwell et al. 1987). Among several possible interpretations of the radio objects,
the subsequent correct identification was made (Churchwell et al. 1987).
12
Table 1.
Catalog Numbers of the Varied Designation Proplyds
Designation LV HST MS VLA J&W Lada Robberto HHM
158-338
5
32
14
79
4
158-323
5
46
11
488
87
158-327
6
4
47
13
86
159-350
3
52
9
499
161-324
4
58
8
98
6
163-317
3
63
7
512
7
167-317
2
73
6
524
116
8
168-328
1
74
5
118
2
170-317
2
4
121
177-341
1
97
1
558
30
182-413
10
O
28
183-405
16
588
244-440
756
Note–Sources: LV (Laques & Vidal 1979), HST (O’Dell et al. 1993), MS (McCaughrean & Stauffer
1994), VLA (Felli et al. 1993), J&W (Jones & Walker 1985), Lada (Lada et al. 2000), Robberto
(Robberto et al. 2005), HHM (Hayward et al. 1994)
It was with the advent of the HST that their true nature was revealed (O’Dell
et al. 1993). The initial discovery with the spherical aberration affected WF/PC camera
(O’Dell et al. 1993) was soon confirmed with the purer images of the aberration-free
WFPC2 (O’Dell & Wen 1994) as shown in Figure 5 and by seeing-corrected groundbased imaging (McCullough et al. 1995). The tell-tale feature was a bright ionized cusp
facing a nearby massive ionizing star, usually with a visible low-mass central star, and
sometimes a dark central region. In a few cases there was no photoionized outer feature
and the object would appear only in silhouette against the nebular background and it
was quickly identified (McCaughrean & O’Dell 1996) that these were proplyds lying
within the foreground Veil. Hydrogen is neutral within the Veil components, which
means that ionizing photons do not penetrate, sparing any proplyd located therein the
complications of photoionization. By now there are at least 150 known emission line
proplyds and 15 silhouette-only proplyds (O’Dell & Wong 1996; Bally et al. 1998,
2000; O’Dell 2001c; Smith et al. 2005a) with some objects having both extinction and
emission characteristics. At least are known to form a binary (Graham et al. 2002).
The most useful measure of their size is the distance between the tips of their bright
cusps. This is about 0.15′′ (about 65 AU) for those closest to θ 1 Ori C and increases
with distance from that ionizing star (O’Dell 1998). A more recent paper (Vicente &
Alves 2005) studying the size distribution did not employ the full resolution version
of the HST images. Although they are concentrated towards the Trapezium, where
observational selection effects favor their detection (O’Dell & Wong 1996), one of the
most interesting objects has been found outside the Huygens region (Bally et al. 2006).
Most of these emission line observations have been in filters isolating the strongest
emission lines from the nebula (Hα, [N II], and [O III]) but additional images in [O I]
have been particularly useful (Bally et al. 1998).
Since the central stars are of low mass and effective temperature, many of them do
not appear in intermediate width visual images, but are revealed in near infrared images
(Lada et al. 2000; Hillenbrand & Carpenter 2000; Lucas & Roche 2000; Muench et al.
2001; Lada et al. 2004). The cause of the X-ray emission from proplyds is still not
understood (Kastner et al. 2005), however, some of them are significant sources of Xrays and this property is especially useful in looking at proplyds obscured by the PDR
13
Figure 5. This 15.9′′ x15.9′′ image made with the HST WFPC2 at 0.0455′′ pixel−1
is color coded with red=[N II] 6583 Å, green=Hα 6563 Å, and blue=[O III] 5007 Å
(Bally et al. 1998). The silhouette proplyd in the upper left is 183-405, the brightest
proplyd near the center is 182-413 and the fainter object in the lower right is 183-419.
(Prisinzano et al. 2008). One must go to quite long wavelengths to clearly discriminate
thermal emission from the dust component and this has been attempted with increasing
success (Hayward et al. 1994; Mundy et al. 1995; Hayward & McCaughrean 1997;
Bally et al. 1998; Robberto et al. 2002; Smith et al. 2005b; Robberto et al. 2005;
Williams et al. 2005; Williams & Andrews 2006; Eisner & Carpenter 2006).
A key paper (Chen et al. 1998) reported on observations of two proplyds in the
2.12 µm line of H2 , producing particular good images of 182-413. In both cases the
H2 emission was seen to arise from the surface of the inner disk that was previously
only seen in extinction. This object was also seen to be peculiar in [O I] 6300 Å images (Bally et al. 1998). This emission line is usually an excellent way of tracing the
location of an ionization front because collisionally excited [O I] demands the presence
of neutral oxygen (which has essentially the same ionization potential as hydrogen)
and also abundant electrons heated by photoionization. The combination of circumstances are only found within the ionization front. The outer part of 182-413 shows
bright [O I] as expected from the proplyd’s ionization front, but the inner dark disk,
which is seen edge-on, is sheathed in [O I] emission as shown in Figure 6. It was later
shown (Störzer & Hollenbach 1999) that this radiation could also be produced by the
photo-dissociation of OH molecules, thus confirming that the gas in the inner disks is
molecular.
14
Figure 6. This 5.6′′ x10.3′′ image of proplyds 182-413 (left) and 183-419 (right)
was made from HST WFPC2 0.0455′′ pixel−1 images with red depicting Hα and
blue the [O I] 6300 Å emission (Bally et al. 2000). In 182-413 the inner molecular
disk is seen nearly edge-on and a faint perpendicular bipolar jet is seen, which forms
a local disturbance when it passes through the dense ionization front on the left. The
[O I] emission near the ionization front is caused by collisional excitation while that
surrounding the inner disk is produced by the photodissociation of OH.
Spectra. Because the proplyds emit their radiation in essentially the same emission
lines as the nebula, obtaining their spectra demands accurate correction for the nebular contribution. Early ground-based observations with Fabry-Perot (de la Fuente et al.
2003) and an echelle spectrometer (Henney et al. 1997) succeeded in obtaining usable proplyd spectra, although they were uncertain because of the large correction for
background radiation. Later ground-based echelle spectroscopy with larger telescopes
under conditions of better seeing produced superior spectra of the strongest lines in
four proplyds (Henney & O’Dell 1999) and then many lines in proplyd 167-317 (Vasconcelos et al. 2005). 167-317 was also the subject of an early HST Faint Object
Spectrograph study at moderate spectral resolution (Walsh & Rosa 1998) as were the
objects 158-327 and 159-350 (Bally et al. 1998). The most complete study of 167-317
is a high resolution, long-slit set of observations (Henney et al. 2002) of the [C III]
doublet at 1907-1909 Å using HST’s Space Telescope Imaging Spectrometer (STIS),
which produced spatially resolved density-dependent line ratios across the object and
its microjet.
3.2.
Physical Models for the Proplyds
The basic nature of the proplyds in the ONC was obvious in the HST discovery paper
(O’Dell et al. 1993) and had been anticipated from the radio observations (Churchwell
et al. 1987). Although an early paper (Henney et al. 1996) showed that the form of the
bright cusps facing θ 1 Ori C could be produced from the interaction of the slow wind
from the stellar accretion disks interacting with the fast wind from θ 1 Ori C, this model
fell aside when it was appreciated that these cusps represented local ionization fronts
(O’Dell & Wen 1994) whose surface brightnesses scaled as expected (McCullough et al.
1995; O’Dell 1998) with their distance from θ 1 Ori C.
The Standard Model. The widely accepted model (Sutherland 1997; Johnstone et al.
1998; Henney & Arthur 1998; Richling & Yorke 1998; Störzer & Hollenbach 1999;
15
Figure 7. This figure (Henney & O’Dell 1999) illustrates the major features of
the standard model for proplyds. The ionizing star is to the right. The white arrows
show the flow of gas photodissociated by non-ionizing FUV that reaches the inner,
molecular disk. The shadowed region behind the ionized hydrogen bright cusp is
illuminated by Lyman continuum photons produced by surrounding nebular gas. The
appearance of an individual proplyd will depend upon the orientation of the inner
disk, whether or not there is an additional ionizing star nearby, and if the proplyd is
shielded from all ionizing photons by being in or beyond the Veil (in which case the
object will appear as a silhouette proplyd).
Kessel et al. 1998; Nguyen et al. 2002) posits an inner accretion-disk of molecular
material that surrounds a low-mass pre-main sequence star. This disk sees FUV radiation only, i.e. radiation of less energy than the 13.6 eV needed to photoionize hydrogen.
This is because the photodissociated molecular gas that is heated and slowly driven off
the inner disk forms an extended atmosphere that is optically thick to Lyman continuum radiation. The latter means that the inner atmosphere will be surrounded by a local
ionization front, which will be brightest in the direction of the dominant ionizing star
and will have a fainter, comet shaped zone behind it that is photoionized only by scattered Lyman continuum radiation. The general form of this model is shown in Figure
7 (Henney & O’Dell 1999). The surface brightness and geometry of the ionized cusps
(O’Dell 1998) indicate that they have electron densities of about 1-10x105 cm−3 , while
16
the [C III] ultraviolet doublet also indicates (Henney et al. 2002) that peak densities
of 106 cm−3 are reached. The trend for increasing size with distance from θ 1 Ori C
is probably the result of a lower column density of hydrogen being necessary to absorb
the lower Lyman continuum flux at greater distances from the source.
This general model explains the form and surface brightness of the bright cusps.
The appearance of the inner disk of material (most easily seen in extinction) is primarily
determined by the orientation of the rotation axis of the disk. When it lies in the plane
of the sky, the disk is a thin dark silhouette, and when it lies along the line of sight,
the disk appears as a circular silhouette. If the proplyd is located within or on the
observer’s side of one of the layers of the foreground Veil, then it will not see any
photons of greater than 13.6 eV, as these will have been blocked by the neutral material
in the Veil, and we will only see the object in silhouette (as frequently seen) or in the
products of photodissociation, which has yet to be done.
Grain Growth in the Neutral Disks? The largest silhouette proplyd is 114-426, which
is seen almost edge-on and has an apparent diameter of over 2′′ . Projected onto a portion
of the Orion Nebula that is relatively smooth, it presents a natural target for trying to
determine if the dust particles in this proplyd are different from those in the general
interstellar medium. This was first done using HST images having resolution as good
as 0.07′′ (30 AU) (Throop et al. 2001). The technique was to compare the difference
in the extinction along the edge of the proplyd at wavelengths of 6563 Å and 1.87 µm
(McCaughrean et al. 1998). Although this technique is very sensitive to the correction
for the different point spread function of the two different cameras employed, it appears
that the extinction is even more grey than the grey extinction already known to exist in
the ONC stars. A subsequent ground-based study (Shuping et al. 2003) at 4.05 µm
indicates that the object at that wavelength is only marginally smaller than in visible
wavelengths, suggesting that the extinquishing particles are larger than 1.9 µm but not
much greater than 4 µm. Later a ground-based study (Shuping et al. 2006) of the
8-12 µm spectra near the silicate emission feature showed that feature in seven of the
eight sources observed. However, the profile is significantly flattened compared with
the general interstellar medium dust, again arguing that grain growth has occurred.
3.3.
The Particularly Well Studied Proplyd LV 2 = 167-317
Arguably the best studied proplyd is 167-317 (LV 2). While still thought to be simply
a star, it was shown that there was a high velocity redshifted outflow (Meaburn 1988;
Meaburn et al. 1993) with a peak velocity at about 120 km s−1 relative to the systemic
velocity. Subsequent ground-based spectroscopy (Massey & Meaburn 1995) indicated
that there was a monopolar jet extending about 2′′ to the southeast. This jet and other
features are essentially high ionization sources, which is not surprising since the object
is close to and obviously photoionized by θ 1 Ori C, as shown in HST images of the
Trapezium (Bally et al. 1998). Figure 8 is an original image derived from HST WFPC2
[O III] exposures and shows the main features of the object. The jet orientation is now
known (Henney et al. 2002) to be towards a position angle (PA) of 120◦ . The peak
emission is at a velocity of 123 km s−1 relative to the cluster velocity (Henney et al.
2002). There is an opposite-pointing counter-jet with an intensity peak at about -134
km s−1 relative to the cluster velocity (Henney et al. 2002). Velocity images of this
region (Doi et al. 2004) show that there are redshifted and blueshifted features extending up to 13′′ along the axis of the jet and counter-jet, collectively being designated
17
Figure 8. The bright object near the center of this HST WFPC2 [O III] 5007 Å
0.0455′′ pixel−1 image is the proplyd 167-317 (LV 2), the central portion being
shown at a higher brightness level for clarity. θ1 Ori C is in the direction PA=215◦
from the center of the bright cusp. The “stand-off” bowshock of LV 2 is displaced
1.74′′ towards θ1 Ori C from the cusp. The high velocity jet is faintly visible as
a nearly vertical feature labeled as ”High-ionization jets” and ”Low-ionization jet”
in this figure prepared by W. J. Henney. The bright linear features are diffraction
patterns from the secondary mirror support of the HST.
as Herbig–Haro (HH) 726. 167-317 has also been the subject of a deep ground-based
spectroscopic study (Tsamis et al. 2007) which succeeded in detecting faint recombination lines of C II and O III.
In addition to the original high resolution radio surveys of the inner Huygens region, the object has been imaged at 6 cm wavelength using the MERLIN radio interferometer and an elliptical resolution of about 0.05′′ by 0.13′′ (Henney et al. 2002).
Comparison with Hα images of the object indicates important differences in appearance (Henney et al. 2002), which are probably due to blending of the thermal emission
of the bright cusp and emission from the jet, together with a variation in electron temperature along the cusp.
167-317 was the target of a STIS spectrometer program with the HST, where the
spatial resolution was about 0.05′′ and velocity resolution about 3 km s−1 (Henney
et al. 2002). The slit was oriented along a line towards θ 1 Ori C, thus providing
spectra across the bright cusp, the “stand-off” shock at 1.74′′ towards θ 1 Ori C, and
the tail region in the opposite direction. The spectra were centered on the [C III] doublet
(actually, the longer wavelength component is only semi-forbidden), which is density
dependent. It was found that the peak electron density in the cusp was about 106 cm−3 ,
diminishing rapidly in the direction of θ 1 Ori C and trailing off to 5x104 cm−3 at the
18
limit of visibility in the tail direction. The “stand-off” shock has a density of about
104 cm−3 , which is consistent with the results of a hydrodynamic simulation (Garcı́aArredondo et al. 2001, 2002).
The STIS spectra are nearly ideal for determining the mass-loss rate from this
proplyd since one can derive from them the gas density at each position and the velocity
of that gas, thus diminishing (although not eliminating) the dependence upon assumed
models. This effort (Henney et al. 2002) yields a mass-loss rate of 8.2x10−7 M⊙ yr−1
with an estimated uncertainty of 20%.
3.4.
Survival against Photo-Evaporation
The same processes that render the Orion Nebula visible and cause mass to flow from
the underlying molecular cloud, through the PDR, and out through the ionized layer will
also be operating in the proplyds. This was recognized when Churchwell originally
posited the correct interpretation of the high resolution radio observations (Churchwell et al. 1987) . Using the available photo-evaporation models and an assumption of
masses like other young stellar objects, he first presented the lifetime problem, in his
case arguing that the proplyds should be destroyed in a time much less than the age of
θ 1 Ori C, yet they are still there. We’ll refer to this as the lifetime “conundrum”. Does
this mean that we do not have good mass-loss rates? Is the ionizing star quite young?
Do we not have good masses for the disks? Can it be that the proplyds have been hidden
from the stellar Lyman continuum radiation until recently? The conundrum remains.
Determination of Disk Masses. Although the bright ionized surface of proplyds is
the most visible portion of these objects, it contains (O’Dell & Wen 1994) only about
10−5 M⊙ of material. It is basically only a way-point as material leaves the innerdisk through photodissociation. What one needs to know are the masses of the inner,
molecular disks where most of the mass resides.
The first estimates for the neutral regions (O’Dell & Wen 1994; McCaughrean &
O’Dell 1996) were for silhouette proplyds, where it was assumed that these objects had
the same gas to dust ratio as the general interstellar medium and the column density
of material was calculated point by point from the visual region optical depth and integrated over the entire dark object. The results were masses of 3x10−7 – 2x10−3 M⊙ ,
but it was recognized that these were likely to be much lower than the actual values as
light from the central stars precluded determination of optical depths in the inner-most
regions, where most of the material must lie, and the method does not give results in
those regions of large optical depth.
The more hopeful method of mass determination is from measurement of the thermal radiation from the proplyd dust (although this method still has the short-coming of
needing to make the assumption of a gas to dust ratio). The observational problem is
significant because one needs high spatial resolution to discriminate proplyd emission
from the irregular background emission from the PDR (which has a similar temperature) and from the thermal continuum of the ionized gas. The earliest attempts (Mundy
et al. 1995; Hayward & McCaughrean 1997; Bally et al. 1998) obtained either very
low masses or only upper limits. More credible results were produced from a study
(Williams et al. 2005; Williams & Andrews 2006) at 880 µm wavelength with the
Sub-Millimeter Array. In that case 23 proplyds lay within the 32′′ primary beam of
the system and the spatial resolution was 1.5′′ . Four of the five sources detected were
proplyds with thermal dust emission (the fifth proplyd’s emission was from the ionized
19
portion). The masses derived were 1.3–2.4x10−2 M⊙ and the 18 non-detected sources
gave upper limits of 8x10−4 M⊙ . A recent study (Eisner & Carpenter 2006) investigated an even wider region at 3 mm wavelength with the Owens Valley Millimeter
Array. Among the sub-set of objects that were detected at both 3 mm and the infrared,
they found six proplyds with masses of 0.13-0.39 M⊙ . By stacking the images of the
other proplyds in the field of view, they found that the average flux was detectable at a
3σ confidence level and corresponded to a disk mass of 0.005 M⊙ . They argue that the
fraction of high mass disks is no different than in the Taurus region, where there is no
reason to expect disk destruction. The most recent study, at 1.3 mm (Eisner et al. 2008)
with a resolution of about 0.60′′ x 0.69′′ and covering the central 2′ x 2′ field yielded
a contrasting result. Dust masses for 33 detected sources (which included 11 known
proplyds) ranged from 0.01 to 0.5 M⊙ and stacking of the 225 known near-IR cluster
member images implied an average disk mass of 0.001M⊙ . They concluded that the
percentage of stars in Orion surrounded by disks more massive than 0.01 M⊙ is substantially lower than in Taurus and that there is marginal evidence for mass depletion
nearer θ 1 Ori C.
Models. Mass loss from the proplyds occurs in two zones. The outer zone is the
local ionization front, where the gas is photoionized and heated, then freely expanding
because its local pressure is much greater than the ambient nebula. The processes in
this zone are driven by Lyman continuum radiation, often called the EUV. The density
distribution in this freely expanding zone is approximately exponential, which lead to
an unnecessary attempt to explain how this ionized atmosphere could resemble a stable
atmosphere (O’Dell 1998). This interpretation of the near exponential decay of density
disappeared with the first measurements of the expansion velocity. The second zone of
mass-loss is at the outer boundary of the inner molecular disk, where photodissociation
of molecules occurs and that gas then expands at a velocity of only a few km s−1 ,
eventually reaching the local ionization front. This inner zone is driven by energetic
photons of less than 13.6 eV, the FUV. Since the time for material to flow from the
inner to the outer zone is short as compared with the predicted lifetimes, the mass-loss
rates in the two zones must be in approximate equilibrium. In the case of the silhouette
proplyds only the FUV driven processes occur.
By now there is a plethora of models for the two zone process (Johnstone et al.
1998; Störzer & Hollenbach 1999; Scally & Clarke 2001; Matsuyama et al. 2003;
Clarke 2007) and even mass-loss from star-star encounters (Olczak et al. 2006).
Although there is general agreement about the basic physics, the models differ in detail
and one cannot say that the definitive model has been constructed. This is primarily
because the inner zone is the most important and the distance of Orion is such that we
don’t have strong observational constraints on the properties of the inner disks. The
mass-loss rate slows with diminishing inner disk size, thus extending their lives.
Mass-loss Rates. Given the uncertainties of the models, it remains wise to look to
observations as a guide. Fortunately, at high spectral resolution one can discriminate
emission of the outflow from the proplyd’s ionization front from that of the background
nebula. The best ground-based study used the Keck Telescope HIRES spectrograph to
study four proplyds (Henney & O’Dell 1999) and found a mass-loss rate of the best
studied object (177-341) of (9±4) x 10−7 M⊙ yr−1 and for the other three objects (8±5)
x 10−7 M⊙ yr−1 (after nearly balancing post-publication corrections (Henney 2001;
Henney et al. 2002)). The method employed was to construct theoretical models of the
20
proplyd ionized gas that would satisfy both the emission line velocity profile integrated
over the bright cusp and the appearance of the object on HST images. The much more
useful study of 167-317 (Henney et al. 2002) determined the density distribution from
lines whose velocity profiles were known and spatially resolved the bright cusps. The
model used in the interpretation of the spectra was also confined to agree with the
appearance of the object, and a mass-loss rate of 8.2 x 10−7 M⊙ yr−1 was found with
an estimated uncertainty of 20%. The results of these detailed spectroscopic studies
are only slightly higher than the mass-loss rates estimated for 27 additional proplyds
from only their calibrated HST images (Henney & Arthur 1998). For the total of 31
inner proplyds that have been studied, the mean mass-loss rate is 3.3 x 10−7 M⊙ yr−1
(Henney et al. 2002). Study of the binary proplyd 168-326 (Graham et al. 2002;
Henney 2002) shows that mass loss occurs not only from the proplyd heads but also
by transonic mass loss from the sides of the tails.
Is there actually a conundrum and if so, what is the way out? The age of the ONC
is addressed in the preceding chapter in this Handbook so that only a brief summary
is needed here. There is evidence from runaway stars (Hoogerwerf et al. 2001) that
massive star formation occurred about 2.5 Myr ago. Certainly star formation has occurred over a period of several million years, but observations show that most stars
are less than a million years old (Hillenbrand 1997; Herbig & Terndrup 1986). When
one tries to use pre-main sequence positions of low mass stars to derive ages the times
become very uncertain because of uncertainties in the models assumed. There are arguments (Palla & Stahler 1999) that the rate of star formation in the ONC has accelerated
and evidence (Hillenbrand 1997) that the stars towards the middle of the cluster are
younger.
Infrared excess observations (Hillenbrand et al. 1998) indicate that up to 90% of
the ONC stars have retained inner (less than 0.1 AU) circumstellar disks. There is no
incompatibility of this result with the fact that the cluster age estimates indicate that the
observational mass-loss rates cannot be sustained to the present because the mass-loss
rates should slow as the circumstellar material is reduced to only the inner disk (Clarke
2007).
However, the conundrum remains that in the innermost region of the ONC (≤30′′
from θ 1 Ori C) (O’Dell & Wong 1996) 66% of the lower mass stars reveal themselves
on HST images as proplyds. This means that they are in a phase of high mass-loss
rate. The 3 mm wavelength study indicated that proplyds in this region have masses
of 0.005–0.39 M⊙ , with many more at the low end of this range. Combining these
masses with the mean mass-loss rate in the preceding section means that the nominal
time over which these objects can have their present basic structure is 1.5 x 104 –1.1 x
106 years, with most of them constrained to the short time period. How can this be? It
must certainly be the case that the formation of a massive star capable of photoionizing
a large region will terminate nearby star formation, which means that this star will be
the youngest in a cluster with an extended range of star formation.
It is possible that we are now in the right few tens to one hundred thousand years
to have observed the most massive stars form in this very close cluster. Smith et al.
(2005b) argue from the concentration of proplyds near θ 1 Ori C that this hot star and
the nearby proplyds are quite young. Models indicate that photo-evaporation of proplyd material should rapidly decrease as the disks become smaller leaving longer-lived
compact disks. There is an observed broader distribution of unresolved IR visible disks
which they associate with an earlier epoch of star formation while the extended material
21
of the prolyds indicate their youth. This is a good argument that there is a quite recent
epoch of star formation within the larger ONC and that this recently formed group has
shed its natal material and is now optically visible.
Quantitatively explaining the revealing of this putative newest region of star formation is difficult. Four of the five proplyds for which the inclination angles are known
from studies of their internal kinematics are beyond θ 1 Ori C (Henney et al. 2002).
This raises the possibility that the proplyds are not centered on the Trapezium stars,
meaning that the geometry is quite complex. At some point the main ionization front
must have moved from very near θ 1 Ori C to its present position. The current distance
between θ 1 Ori C and the ionization front of about 0.3 pc would be traversed in 105
years at a velocity of 3 km s−1 . Current observations show no relative motion of the
ionization front and the cluster stars to within a few kilometers per second, so there is
no incompatibility between the required and observed relative velocities. However, the
complex problem of the motion of the ionization front around a new star in a geometry
like the Orion Nebula has not been solved. This means that we cannot say if the current
structure of the nebula is compatible with the Trapezium stars and the proplyds being
coeval and very young.
The evidence from runaway stars (Hoogerwerf et al. 2001) and the presence of
separate massive star formation in the BN-KL region and the Orion-S region all indicate
that massive star formation has occurred in several centers and at several times. It may
be that the Trapezium grouping of massive stars is but the most recent and visible of
such events.
4.
Secondary Star Formation Centers and Their Outflows
Although the ONC is the most populous star formation center in the immediate vicinity
of the Orion Nebula, it is clearly not the only such center. However the other two known
centers are within or beyond the dense dust and gas of the PDR and suffer from high
extinction. Except for when an outflow feature penetrates the PDR, these regions are
best seen in X-ray, infrared, and radio wavelengths. Figure 9 is an image of the thermal
emission from regions of the OMC containing the two embedded secondary regions of
recent star formation discussed in the remainder of this section.
4.1.
The Becklin-Neugebauer and Kleinmann-Low Region
The brightest infrared sources in the Orion Nebula region are the Becklin-Neugebauer
(BN) (Becklin & Neugebauer 1967) and the Kleinmann-Low (KL) (Kleinmann & Low
1967) objects, which are referred to in common as the BN-KL region. The Genzel &
Stutzki (1989) review provides a superb tabulation of earlier observations of the BN-KL
region and it is not necessary to try to supersede that review here. The BN-KL region is
located about 1′ northwest of the Trapezium and is the richest molecular line emission
region in the entire Orion Nebula. This region has a bolometric luminosity of about 105
L⊙ .
Molecular Emission and Compact Sources. The molecular line surveys carried out
by single-dish and interferometer instruments toward the BN-KL region have shown
a very rich variety of lines, see Figure 10 (e.g. Sutton et al. 1985; Blake et al. 1987;
Schilke et al. 1997, 2001; Comito et al. 2005; Beuther et al. 2004, 2005, 2006). Those
surveys have found about 2200 lines in a range 70 to 345 GHz (Schilke et al. 1997) and
22
Figure 9. SCUBA 850 µm map of the central Orion Nebula (Johnstone & Bally
1999) compared to interferometric 2.7mm continuum observations obtained with the
BIMA array (blue contours; J. Di Francesco, in preparation). Continuum sources of
> 5σ are designated with green star symbols.
Figure 10. Submillimeter Array (SMA) spectral line observation in the 850 µm
band toward Orion-KL from Beuther et al. (2005). The panel shows a vectoraveraged spectrum in the uv-domain on a baseline of 21m.
23
recent Odin satellite observations have revealed 280 spectral lines from 38 molecules in
the 486-492 GHz and 541-577 GHz bands (Olofsson et al. 2007; Persson et al. 2007).
Studies of the molecular line emission toward BN-KL site have allowed identification
of several chemically and kinematically different gas components within this region.
1. A broad-line gas component originating from at least two molecular outflows,
extended roughly over scales > 104 AU and centered in the vicinity of IRc2 (Genzel
& Stutzki 1989). The presence of a high velocity outflow was first discovered in CO
and H2 S emission with a southeast-northwest direction (Wilson et al. 1970; Thaddeus
et al. 1972). This outflow is best traced by sulfur-bearing species (H2 S, SO, or SO2 , see
Wright et al. 1996). A second lower velocity outflow in the northeast-southwest direction was later discovered by Plambeck et al. (1982). The outflow-tracing species show
a characteristic triangular line shape (usually indicated as a plateau profile) with broad
wings of up to 100 km s−1 in the CO features (see Rodrı́guez-Franco et al. 1999a,b;
Wirström et al. 2006). This outflow is best traced by thermal SiO and H2 O maser
emission, as well as some H2 bow shocks (e.g., Genzel & Stutzki 1989; Blake et al.
1996; Chrysostomou et al. 1997; Stolovy et al. 1998)
2. The so-called Compact Ridge, a compact region whose main chemical feature is the high abundances of complex oxygen-bearing species, such as CH3 OH or
HCOOCH3 (see Johansson et al. 1984; Wright et al. 1996). Molecular lines radiating
from the Compact Ridge show relatively narrow widths (∼ 10 km s−1 ). It has been
reported to have warm temperatures (of the order of 100 K, Wright et al. 1996; Liu
et al. 2002).
3. The Extended Ridge, which is made up of quiescent gas which extends NE to
SW through Orion BN-KL and is characterized by VLSR in the range 7-10 km s−1
(Johansson et al. 1984; Blake et al. 1996), line widths of 2-4 km s−1 , and gas kinetic
temperatures of 50-60 K (Minh et al. 1990; Askne et al. 1984). Chemically this region
is characterized by standard gas phase, ion-molecule chemistry, with an abundance of
carbon-rich species (e.g. CS, CN, CCH), and a lack of oxygen-rich molecules.
4. Finally, lying along our line of sight, very close to IRc2 and at a projected
distance of about 500 AU from source I (Beuther et al. 2006), are clumps of very
dense (106 cm−3 ), warm (130-335 K, Wilner et al. 1994; Wright et al. 1996; Wilson et al. 2000) material known as the Hot Core, first detected in the inversion lines of
NH3 (Morris et al. 1980). In general, this source shows unusually high abundances of
nitrogen-bearing species, possibly induced by the strong outflows located close to this
region (see Blake et al. 1987; Wright et al. 1996).
Recently, infrared and radio maps have revealed a more complicated view of the
BN-KL region, resolving the enigmatic source I from the hot molecular core (Figure 11
and Figure 12), and finding an isolated new protostellar source SMA1, that is a strong
line emission source (Beuther et al. 2005).
Menten & Reid (1995) and Gezari et al. (1998) proposed that the main heating
source in this molecular region is radio source I. However, the lack of an IR counterpart
to source I has led some authors to invoke large intrinsic foreground extinction toward
source I. Wright et al. (1996), Blake et al. (1996) and Chandler & Wood (1997) showed
that the dust and molecular peak emission coincide neither with radio source I nor with
IRc2. The main peak of continuum emission at 1.3 mm is 1′′ east from source I and is
associated with the hot core, as seen in Figure 12. Furthermore, the distribution of the
HC3 N J = 24-23 line suggests that source I is the dominant energy source in the region
and that there is no evidence of internal heating within the molecular hot core (Blake
24
Figure 11. This composite figure (Shuping et al. 2004) shows a Keck 12.5 µm
image in gray, overlaid with ammonia emission as green (Wilson et al. 2000), H2 O
masers as filled circles (Genzel et al. 1981; Gaume et al. 1998), and OH maser clusters as open circles (Johnston et al. 1989). Red and blue indicate doppler shifts
relative to VLSR = 5 km s−1 .
et al. 1996). However, based on a model for the heating of the Hot Core, Kaufman
et al. (1998) concluded that it is internally heated by young embedded stars. de Vicente
et al. (2002) also proposed that it is being internally heated, most likely by one or more
young massive protostars, hence presenting the ideal conditions for an extremely active
gas-phase chemistry. The highest resolution images suggest that it is a massive disk
around a YSO (Reid et al. 2007).
4.2.
The Wide-Angle BN-KL Outflow
Optical and Infrared Observations. The molecular outflow associated with the BNKL region is arguably the most spectacular outflow source within a kpc of the Sun.
The outflow is almost completely invisible at optical wavelengths, but in the near-IR,
its system of multiple “fingers” or bullets provide a spectacle in narrowband H2 and
[Fe II] images of the region (Figure 13). These fingers were first reported by Taylor
et al. (1984) and then by Allen & Burton (1993), but have been studied with narrowband
IR imaging and long-slit spectroscopy several times since then with increasing image
25
Figure 12. SMA 865 µm map of Sources “I” and “n” in Orion-KL , from Beuther
et al. (2004).
quality (Schild et al. 1997a,b; Stolovy et al. 1998; Schultz et al. 1999; Kaifu et al.
2000; McCaughrean & Mac Low 1997). The H2 outflow is oriented along a southeastnorthwest axis, perpendicular to the extension of the hot molecular core seen in NH3 by
Wilson et al. (2000). The northwest lobe is brighter and Doppler shifts indicate that it is
tilted toward us (Scoville et al. 1982; Beck et al. 1982; Beck 1984; Salas et al. 1999). It
appears to originate near radio source I at the center of the cluster of high-mass YSOs
embedded in OMC-1 (Menten & Reid 1995; Greenhill et al. 1998).
A number of the “fingers” have at their ends low ionization shock features that
are seen in the visual window, indicating that they occur just on the observer’s side of
the PDR, for if they were deeper inside the PDR, extinction would prevent their being
seen and if they were well inside the H II zone they would be photoionized. These
optical features have been well studied for both their velocities and their proper motions
(Hu 1996; Doi et al. 2002; Graham et al. 2003; Doi et al. 2004; López et al. 2005).
Determination of the spatial motion of the feature designated as HH 201 indicate that
the source of the outflow is about 0.2 pc beyond the ionization front (Doi et al. 2004).
Proper motions of several optically visible HH objects as well as IR knots seen in H2
and [Fe II] indicate a common explosive origin about 1000 yr ago assuming linear
motion (Doi et al. 2002), or significantly less if they have been decelerated (Lee &
Burton 2000). This coevality and the wide angle morphology are important clues that
the BN-KL outflow is very different from the highly-collimated and relatively steady
jets associated with low-mass star formation.
The wide-angle molecular outflow has apparently cleared out a large cavity that allows us to see scattered near-IR continuum light from the central massive stars (Werner
et al. 1983; Morino et al. 1998; Simpson et al. 2006). Yet, our line of sight to the cen-
26
Figure 13. This composite infrared image (127”x180”, ≃0.28x0.39 pc) shows
the spatial relationship between warm dust at a few hundred K (colored orange)
and molecular hydrogen (blue) in the brightest region of the Orion Nebula. The
warm dust is seen in a mosaic of ∼12 µm images taken with the T-ReCS camera on
Gemini South, see Smith et al. (2005b), and the H2 2.12 µm line was imaged with
a narrow-band filter with NIC-FPS on the 3.5m ARC telescope (image courtesy J.
Bally; composite image made by N. Smith).
tral engine is highly obscured even at thermal-IR wavelengths as we would be looking
through a hot molecular core (there is a tight anti-correlation between escaping mid-IR
emission and the spatial extent of NH3 from the hot core; see Shuping et al. (2004). It is
likely that this dark and dense material along our line of sight is part of a toroidal cloud
that pinches the waist of the larger bipolar outflow. Cunningham et al. (2005) have
shown that a wide-angle bipolar outflow such as this can arise from a spherical wind
expanding into a rotating and collapsing envelope. Other evidence for bipolar outflow
is found in polarization sensitive infrared images (Tamura et al. 2006).
Radio Observations. Well before the high-resolution near-IR images in H2 emission
available today, the OMC-1 core was known to host a powerful molecular outflow seen
in its high velocity CO line wings (Kwan & Scoville 1976). The molecular outflow
is accompanied by a complex series of masers seen in OH (Norris 1984; Cohen et al.
27
2006), H2 O (Genzel et al. 1981; Gaume et al. 1998), and SiO (Greenhill et al. 1998;
Doeleman et al. 1999) on various size scales from about 1350 AU down to a few AU
near source I. These masers hold important clues to the fast and slow molecular outflow components in OMC-1, but their geometric interpretation can be complicated. For
example, Gaume et al. (1998) interpret the masers associated with source I as originating in an expanding rotating disk, signifying an outflow oriented in line with the larger
BN-KL outflow. Greenhill et al. (2004), on the other hand, prefer the source I masers
to originate in a jet oriented orthogonal to the larger outflow. In any case, given the putative recent gravitational encounter (see below), the current orientation of the outflow
from source I on small scales may be unrelated to the orientation of the larger BN-KL
outflow.
Origin of the Outflows. The kinetic energy of the BN-KL molecular outflow is roughly
4×1047 ergs (Kwan & Scoville 1976). Proper motions point toward a single catastrophic event that occurred during the last millenium in order to power this outflow.
Interestingly, the motions of the three most massive stellar sources in OMC-1 indicate
a recent close encounter that may have provided the necessary trigger. Proper motions
measured in 3.6 and 6 cm radio continuum data obtained with the VLA spanning 15 yr
reveal that sources BN, n, and I are all runaway stars from OMC-1, and that all three are
moving away from a common point where they were located ∼500 yr ago (Rodrı́guez
et al. 2005; Gómez et al. 2005). Significantly, this location corresponds to the center
of the BN-KL outflow. Furthermore, the kinetic energy in these runaway stars is about
2×1047 ergs (Gómez et al. 2005), which is similar to the observed kinetic energy of
the outflow. It seems an unlikely coincidence that the center of the outflow also happens to be the point where the three most massive stars were all located within a few
hundred AU of one another, that the ages for these events agree, and that the amount
of energy involved in each is comparable. Instead, it seems plausible to infer a causal
relationship between them (Bally & Zinnecker 2005). Bally & Zinnecker also discuss
the possibility that the energy source was a binary merger event. A contrasting view
(Beuther & Nissen 2008) is that while the low velocity outflow originates in source I,
the high velocity flow originates in source SMA1 with a dynamic lifetime of about 103
yrs. Whatever the exact cause, the BN-KL outflow may be an example of rare ejection events that accompany catastrophic gravitational encounters in the dense clustered
environments where massive stars form.
The BN-KL large scale outflow is essentially unique within the Orion Nebula region and in other nearby well-studied star formation regions. This indicates that the
driving event is associated with the formation of massive stars, a property missing in
these other regions.
There may be evidence for an earlier similar event. On scales larger than the H2
fingers, emission from NH3 shows a series of large filaments (Figure 3) extending to
the northwest that may be the walls of cavities carved during ancient eruptions from
OMC-1 (Wiseman & Ho 1996, 1998). These cannot be linked with the event posited
in the preceding paragraph, but argue for much earlier similar expulsion of material.
The NH3 features align with the dusty fingers imaged by Johnstone & Bally (1999)
and these features may simply be reflections of primordial conditions within the host
molecular cloud. An argument has been made linking the large HH 400 shock feature
found to the southeast of the Huygens region with an earlier (5×104 yr) event in the
BN-KL region (Bally et al. 2001), but it is more likely that this object originates from
the Orion-S region (Henney et al. 2007).
28
4.3.
Orion-S
The third center of star formation near the Orion Nebula is the Orion-S region, which
is like the BN-KL region in that it too suffers considerable optical extinction, but it has
multiple nearly co-located centers of atomic and molecular outflows. It has a bolometric
luminosity of ∼ 104 L⊙ (Drapatz et al. 1983; Mezger et al. 1990), a factor of 10 less
than the BN-KL region. It is different from the BN-KL region in lying very close to or
within the PDR, so that many of its outflows produce optically visible features and in
this sense the region is easier to study. The region lies in a local rise in the surface of the
H II ionization front (Wen & O’Dell 1995), which probably means that the ionization
front is slowly eroding the host molecular cloud, lagging behind where the underlying
density is greatest. The region has recently been shown to have a hot core of H13 C+ O
within the larger OMC-1 cloud, being source 112 in the study of Ikeda et al. (2007).
The region has been surveyed at 10 µm (Smith et al. 2004) and in the L band (Lada
et al. 2004). There are two groupings of molecular emission, a northern group with
large-scale outflows and a southern group with hot cores.
Figure 14. This SMA 1.3 mm continuum image of the southern-most region of
Orion-S (Zapata et al. 2005) is overlaid with the CH3 CN[124 -114 ] integrated emission (white contours) of the hot molecular cores 139-409 and 134-411. The scale
bar indicates the 1.3 mm emission in mJy beam−1 . The yellow rhombi indicate the
positions of the compact 7 mm continuum radio binaries associated with hot cores.
The green rhombus and triangle denote the much more uncertain positions of the
source FIR 4 (Mezger et al. 1990) and the millimeter source CS 3 (Mundy et al.
1986), respectively. The blue and red crosses indicate the position of the blue- and
red-shifted H2 O maser spots, respectively, reported by Gaume et al. (1998). Note
that the masers associated with the hot molecular core (134-411) show a large velocity gradient, going from -20 to +45 km s−1 . The green ”X” symbols indicate
the position of the 3 mm BIMA continuum sources reported by Eisner & Carpenter
(2006). This figure is from Zapata et al. (2007).
29
Figure 15. This annotated 43′′ x43′′ image of the Orion-S region is a composite of
HST WFPC2 images with red=[S II], green=[N II], and blue=[O III] (Henney et al.
2007). The white ellipse indicates the OOS region discussed in the text. The squares
and circles represent the positions of various radio and infrared compact sources,
with the white circle representing the 1.6 and 2.2 µm sources from Hillenbrand &
Carpenter (2000); the red circles representing the 10 and 20 µm MAX study of Robberto et al. (2005), with the brightest source indicated by a filled circle; the black
squares representing the 1.3 cm sources of Zapata et al. (2004b); the red squares
representing the sources seen at both 1.3 cm and 1.3 mm; and the blue squares representing the sources seen only at 1.3 mm (Zapata et al. 2005). FIR 4 is the 1.3 mm
source of Mezger et al. (1990), and CS 3 is from Mundy et al. (1986). CO flows
(Zapata et al. 2005) are shown as contours of the redshifted (red) and blueshifted
(blue) components. The SiO flows (Zapata et al. 2006) are shown in the same way
except that pink and light blue are used.
Radio Molecular Emission Lines. The Orion-S region is a rich molecular line emission region within the Orion A molecular cloud (Ziurys et al. 1981; Mundy et al. 1986;
McMullin et al. 1993; Zapata et al. 2007). The southern region contains two highly
obscured hot cores (Zapata et al. 2007) associated with the compact millimeter and
centimeter continuum sources: 134-411 and 139-409 (Zapata et al. 2004a,b) and surrounded by many water masers (Gaume et al. 1998). Furthermore, these two compact
hot cores show high density tracers, like CH3 CN, CH3 OH and HCOOCH3 and shock
tracers like CS, SO, SO2 indicating the early formation of massive or intermediate-mass
stars (Figure 14).
Molecular Outflows from Orion-S. Many high and low velocity molecular outflows
have been detected toward this region (Figure 15). The first is a low-velocity (10 km
s−1 ) bipolar SiO (5-4) outflow with a length of ∼ 30′′ in the northeast (blueshifted) –
30
southwest (redshifted) orientation, and it is centered at the position 5:35:12.8 −5:24:11
(J2000) (Ziurys et al. 1990). A second quite extended (3′ ), collimated, low-velocity
(5 km s−1 ) bipolar CO outflow, oriented NE (blueshifted) – SW (redshifted), has been
reported by Schmid-Burgk et al. (1990), and its center has been associated with the
continuum radio source 134-411 (Zapata et al. 2004a,b). However, recent SiO molecular observations show that 137-408 seems to be the exciting source (Zapata et al. 2006).
Finally, there is a high-velocity bipolar CO outflow, with a length of 0.07 pc (0.5′ )
and velocities of -140 km s−1 toward the northwest and 88 km s−1 to the southeast
(Rodrı́guez-Franco et al. 1999a,b). The exciting source of this high-velocity outflow
was proposed to be 20′′ north of the 1.3 mm continuum source FIR 4 (Rodrı́guez-Franco
et al. 1999b); but, more recently, Zapata et al. (2005) using interferometric observations
resolved the CO thermal and continuum emission and found that the exciting source of
this outflow is the infrared source 136-359. However, the bolometric luminosity of
this source appears to be far too low to account for the powerful molecular outflow.
Most recently, Zapata et al. (2006) have found a cluster of hidden compact outflows
in the Orion-S region using silicon monoxide observations toward this region. Those
outflows are very compact and show a rich variety of morphologies and velocities. Furthermore, they also found that some of these outflows seem to be the base of powerful
Herbig-Haro jets and large-scale molecular flows that emanate from a few arcseconds
around this zone. Some of these sources may reveal themselves in polarization sensitive infrared observations that have the signature of bipolar flows near embedded stars
(Hashimoto et al. 2007).
Formaldehyde Absorption Lines in Orion-S. Johnston et al. (1983) and (Mangum
et al. 1993) have resolved a region of H2 CO absorption in Orion-S in the vicinity of
the SiO and CO outflows discussed in the previous section. The presence of H2 CO
in absorption means that there must be a cloud of very cold molecular gas and also a
background source of continuum. This brings into question the standard model for this
region (that Orion-S simply represents a rise in the main ionization front caused by a
density concentration in the underlying molecular cloud). Explanation of the H2 CO
absorption demands that there is another bright ionization front that we do not see
optically and this layer lies beyond a cold dense molecular cloud. It may even be that
the sources driving the SiO and CO outflows lie in this same cloud. If this model is
correct, in the direction of Orion-S there would first be the foreground veil, then the
open cavity including most of the cluster stars, then the main ionization front, and
beyond that both a dense molecular cloud, a second major ionization front and finally
the host OMC.
4.4.
Optical Outflows from Orion-S
There are multiple outflows seen to emerge from the vicinity of Orion-S. The reason that
one sees so many optical features is the fact that some of the outflow sources lie only
a few hundredths of a parsec behind the H I ionization front of the Orion Nebula (Doi
et al. 2004). Because of the bright background of the nebula, all but the most prominent
outflows have been discovered with the HST. The first HH objects discovered (Cantó
et al. 1980; Meaburn 1986) were HH 202, HH 203, and HH 204 (M42-HH2, HH3, and
HH4 in the original designations), all of which show high blueshifted velocities (O’Dell
& Wen 1994). The less obvious outflows revealed by their proper motions and Doppler
shifted emission are HH 269, 507, 528, 529, 530, 605, 606, and 625 (c. f. Figure 16
31
Figure 16. This annotated 143′′ x168′′ image of the Orion-S region is similar to
Figure 15, except that the field of that figure is outlined in red and the locations
of several HH objects are indicated. Thin yellow lines indicate H2 emission from
Subaru images (Kaifu et al. 2000) and high velocity flow indicated by enhanced He I
10830 Å emission (Takami et al. 2002; Doi et al. 2004). The large central cross
indicates the reference position of 5:35:13.6 -5:24:00 (J2000).
and Figure 17). This census of HH features has resulted from studying a combination
of their morphology in HST images, proper motions measured in multi-epoch HST
images, and ground-based studies of the Doppler-shifted emission using long-slit spectrographs or Fabry-Perot imaging (O’Dell & Wen 1994; O’Dell et al. 1997a,b; Bally
et al. 2000; Rosado et al. 2001; Doi et al. 2002; O’Dell & Doi 2003; Smith et al. 2004;
Doi et al. 2004; Henney et al. 2007; Garcı́a-Dı́az & Henney 2007). A few of these
features, like HH 529 and HH 202, are also bright in thermal-IR emission from dust
that is entrained (Smith et al. 2005b), as well as near-IR emission from [Fe II] λ12567
and He I λ10830 (Takami et al. 2002).
All of these optical features are blueshifted, even those that seem to be the result of
bipolar outflows on the basis of their images and proper motions. Initially this apparent
peculiarity was thought simply to be due to the selection effect of our only seeing the
results of embedded source outflows that are coming towards the observer and breaking
through the PDR. However, this cannot work for explaining the blueshifted bipolar
features. The solution probably lies in the deflection mechanism suggested by Cantó &
Raga (1996), who demonstrated that a collimated jet feature passing through a medium
with a lateral density gradient will be deflected in the direction of decreasing density.
These conditions are satisfied within the PDR, so that even an initially redshifted jet
32
Figure 17. This drawing depicts the 1650′′x1825′′ field survey by HST’s ACS
(Henney et al. 2007) and labels the major features of the Orion Nebula and its largescale outflows. The diverging red arrows indicate flow away from the BN-KL source.
The dashed yellow line traces the path of the HH 625 components and the dashed
black line the aligned HH 529–HH269 features. The red H2 symbols indicate compact H2 sources (Stanke et al. 2002). The outermost shock features are shown in
blue for clarity.
entering it can be deflected into being blueshifted. This mechanism is aided by the
fact that Orion-S coincides with a locally convex region of the main ionization front
(Wen & O’Dell 1995), giving more of the initially redshifted jets the opportunity to be
deflected.
All of the outflows appear to originate from the same general region of Orion-S,
and they branch out in multiple different directions. The deflection mechanism of Cantó
& Raga (1996) cited in the previous paragraph compounds the situation, because it can
not only alter the radial velocity of the flow, but can also alter the direction in the plane
of the sky if there is a density gradient in that plane. In addition, once the outflow has
broken out into the H II region, it can be deflected by global flows of the ambient gas
(Masciadri & Raga 2001). The rule of thumb for finding the sources is to trace the flow
back as far as possible. The most up-to-date attempt at this has been done using both
optical and radio features (Henney et al. 2007). In that study it is shown that although
33
some outflow sources are well identified (e. g. the HH 625 feature arises from the same
source that produces the blueshifted CO molecular outflow), most remain unidentified.
O’Dell & Doi (2003) argued that many of the brightest HH features originate in a region
they designate as the OOS (optical outflow source) on the assumption that the HH 269
and HH 529 features are two sides of the same bipolar flow. Oppositely moving features
with known proper motions narrow the location to a small ellipse, shown in Figures 15,
16, and 17, which has only a few infrared sources, none of which are particularly bright,
extending the under-luminosity puzzle first encountered when trying to explain the CO
outflow. Given the possibility of deflection along the light of sight by the gradient in
density in the PDR, one can conclude that the molecular flows and their associated
optical features are the results of bipolar flow from multiple sources. Although it is still
ambiguous as to which embedded sources power specific individual outflows, the large
number of IR and radio sources discovered in Orion-S is probably adequate to account
for the multiple outflows (Smith et al. 2004; Zapata et al. 2005).
Figure 17 shows a large number of features lying outside of the Huygens region
of the Orion Nebula that may be associated with the outflows from the Orion-S region (Henney et al. 2007). Because of the deflection mechanisms arising from pressure
gradients (Cantó & Raga 1996) and lateral winds (Masciadri & Raga 2001) these associations are much more speculative that those posited for HH objects found in the
Huygens region.
5.
Outflow Features Associated with Optically Exposed Objects
The Orion Nebula region contains the highest concentration of known YSO outflows
in the sky. This is a testament to the extreme youth and clustered nature of recent star
formation in Orion, as well as its proximity, making it a powerful laboratory for investigating the intimate relationship between accretion and outflow in the early lives
of stars. It also shows the most diverse range of outfow properties among nearby starforming regions, from the smallest microjets associated with individual low-mass stars,
up to the wide-angle explosive outflow from the massive young stars in the BN-KL
region. These YSO outflows can be grouped into three distinct phenomena that are
conveniently represented by the outflows from three separate regions: the brighter embedded OMC-1 region with its explosive BN-KL outflow, highly collimated outflows
and classic Herbig-Haro features from embedded sources in Orion-S, and jets from exposed sources in the ONC. The first two groups are covered in the preceding section
and the third in this section.
Because the driving sources of these outflows are within the ionized cavity of the
Orion Nebula, their material is exposed to the EUV radiation from θ 1 Ori C and they are
frequently referred to as being “irradiated”. Irradiated objects represent a phenomenon
that is distinctly different from typical HH objects in quiescent regions like Taurus. In
Taurus-like regions the HH objects seen at visual wavelengths are dominated by forbidden lines emitted only in the cooling zones behind shocks in the flow ( Reipurth & Bally
2001 provide a general review of HH objects).This means that only the portion of the
outflow that has recently passed through a shock can be seen at any given time, and the
analysis of the emitting gas is subject to non-equilibrium shock physics. In irradiated
jets, on the other hand, all the jet material is photoionized and rendered visible, and
its optical emission can be analyzed with simpler diagnostics applied to H II regions
(Bally & Reipurth 2001). Several irradiated jets like these were also seen in the σ Ori
34
region (Andrews et al. 2004; Reipurth et al. 1998) and in the Pelican Nebula (Bally &
Reipurth 2003). Shocks are still important in the hydrodynamics of the outflows, but
the excitation and emission mechanisms are dominated by photoionizing radiation.
The observed phenomena associated with irradiated objects in Orion can be grouped into three basic classes: 1) “microjets”, which are highly-collimated jets that can be
traced very close to their point of origin, 2) the larger classical HH flows, often seen to
be an extension of these microjets, and 3) LL Ori objects, which contain a curved shock
front pointing back toward θ 1 Ori C. All three of these phenomena are represented in
the image of HH 502 (plus LL5 and LL6) in Figure 18 (Bally et al. 2006).
Figure 18. HST/ACS image of the HH 502 flow in the Orion Nebula, including the
main HH 502 flow, the HH 502 microjet emerging from the proplyd that surrounds
the star V421 Ori, as well as two LL Ori objects in the field. This color image has
red=[N II], green=Hα, and blue=[O III] (from (Bally et al. 2006), with permission).
Subcomponents of the HH 502 jet are labeled, and the arrow points in the direction
toward θ1 Ori C.
5.1.
Microjets
Microjets are highly-collimated outflows that are typically seen within about 1′′ of the
driving star, which in many cases is at the center of a silhouette disk or a bright proplyd.
These microjets are orders of magnitude smaller than classical HH objects, and are
analogous to the microjets from T Tauri stars (e.g. Solf & Böhm 1993). Because of
their small size, all the microjets in Orion have been discovered only in the past decade
in high resolution narrowband images obtained with HST (e.g., Bally et al. 2000, 2006;
35
Bally & Reipurth 2001; Smith et al. 2005a). Typically, when the microjet originates
within a proplyd, the part of the jet that propagates within the neutral interior of the
proplyd is brightest in [O I] λ6300, whereas this emission disappears and the jet is seen
in Hα when it breaks through the proplyd ionization front (Bally et al. 2000).
Although the jets are bipolar, many of them appear one-sided in HST images,
or have a strong brightness asymmetry. For the apparently one-sided jets, where the
fainter lobe is lost in the glare of the background nebula, both outflows can be seen in
high resolution spectra where the Doppler-shifted jet material is separated from the
nebula (Bally et al. 2000). Bally & Reipurth (2001) and Reipurth & Bally (2001)
discuss possible mechanisms for this asymmetry, which include extinction of one half
of the flow by a disk, radiative excitation of only the side facing θ 1 Ori C, a density
gradient in the surrounding medium, or intrinsic differences in the opening angle or
flow speed in each direction of the jet. In contrast Henney et al. (2002) argue that
at the sub-arcsecond scale the jets are excited by shock interactions with the ambient
gas, differences of density and velocities producing the asymetries. Jets with strong
brightness asymmetries usually show asymmetry in the outflow speed as well (Bally
et al. 2000; Bally & Reipurth 2001). Some of the most striking examples of one-sided
microjets from proplyds in Orion are HH 514 in the Trapezium (Bally et al. 2000),
HH 526 and 527 located south of the Orion bar (Bally et al. 2000), and HH 668 in
M43 (Smith et al. 2005a).
These microjets are seen close to their central driving source, and represent our
best view of the initial conditions of the jet, before it is decelerated by the surrounding
material in the nebula. Consequently, microjets are fast, with Doppler shifts of typically
50–200 km s−1 (Bally et al. 2000). In cases where multi-epoch HST images exists,
proper motions can be measured with temporal baselines as small as just 1 yr (Bally
et al. 2000; Smith et al. 2005a; Bally et al. 2006). Some microjets show clear signs of
clumps in the outflow, even within less than 100–200 AU from the driving source; the
HH 668 microjet is a prime example with new clumps emerging on a 1 yr timescale
(Smith et al. 2005a). When inferred from Hα emission measures or [S II] ratios, typical
electron densities range from 103 to 105 cm−3 . The corresponding mass-loss rates in
the microjets are a few ×10−9 to as much as ∼10−6 M⊙ yr−1 (Bally et al. 2000; Smith
et al. 2005a; Bally et al. 2006). If the photo-evaporation timescales for the proplyds are
of the order 105 yr, only the most massive of these outflows are tracing accretion phases
that are still adding significantly to the central star’s mass.
5.2.
Larger HH Flows and Bent Jets
Outward along the projected axis of a microjet, one often finds chains of HH objects
out to 30′′ or more from the central star, signifying that microjets are only the brightest
and densest inner parts of more extended and time-variable bipolar outflows. A good
example is HH 502 shown in Figure 18 (Bally et al. 2001, 2006), while other prominent
examples are HH 540 from the Beehive proplyd (181-826), HH 668 from the binary
proplyd (253-1536) in M43, and HH 667 from an edge-on disk in the outer Orion
Nebula (Bally et al. 2005; Smith et al. 2005a). Many other examples of proplyds with
microjets and larger bipolar outflows have been noted as well (Bally et al. 2000; Smith
et al. 2005a; Bally et al. 2006). In many of these cases, proper motions of features seen
in HST images confirm that they are associated with the same driving sources as the
microjets along the same apparent axis.
36
At large distances from the driving sources, outflows are susceptible to significant bending by external forces. This is particularly apparent within the interior of the
Orion Nebula, where the jets interact with bulk flows of plasma and strong EUV radiation. Consequently, the bent jets tend to be C-shaped bends, rather than the S-shaped
bends that are more common in quiescent regions like Taurus where a precessing central source usually causes kinks or bends along the jet. The C-shaped bending can
potentially arise from two physical mechanisms. One is the simple ram pressure of a
side wind, when the jet source is emersed in a large scale plasma flow (e.g., Masciadri
& Raga 2001). Another possibility is the rocket effect; if part of the jet is neutral and
bathed in a strong UV radiation field from one side, the photoevaporation of the neutral
jet can push the remaining part of the jet in the opposite direction (Bally et al. 2006)
5.3.
LL Orionis Objects
The most extreme examples of the external bending of HH jets are the so-called LL Ori
objects, named for the prototype LL Orionis, which is located a few arcminutes southwest of the Trapezium. The parabolic nebula around LL Ori, whose apex points almost
east, was first noted by Gull & Sofia (1979). Fourteen additional objects with similar
parabolic or C-shaped morphology were later seen in narrowband HST images obtained with the WFPC2 camera (Bally et al. 2000; Bally & Reipurth 2001), all with
their apexes pointing back toward the center of the Huygens region. Bally et al. (2006)
noted a number of additional examples in Orion, seen in a larger survey with the ACS
camera. Several LL Ori objects are clustered southwest of the Trapezium near LL Ori
itself (Figure 19), while another group is found south of the Bright Bar. The LL Ori
objects are different than the stationary wind-wind collision shocks around many of the
inner-most proplyds in the Trapezium seen in [O III], Hα, and thermal dust emission
(Bally et al. 1998, 2000; Smith et al. 2005b; Hayward et al. 1994), because those represent the shock between the slow flow of ionized gas away from the proplyd and the
high velocity stellar wind of θ 1 Ori C.
Initially, the LL Ori objects were interpreted as wind-wind collision fronts; their
parabolic morphology implying that a wide-angle stellar wind from a young star or
evaporating photoionized gas encounters a large-scale bulk flow of plasma coming from
the center of the Orion Nebula. In that case, a cluster of LL Ori objects would signify
a particularly strong local bulk flow of plasma in that part of the nebula. However,
more recent HST/ACS observations (Bally et al. 2006), with the ability to measure
proper motions compared to earlier WFPC2 images (although it should be noted that
the inter-camera proper motions derived for LL Ori do not agree with those previously
published from intra-WFPC2 comparisons (Bally et al. 2000)), have shown that most
LL Ori objects in Orion also contain collimated jets that lie along the long dimension
of the shock. The LL Ori bows themselves are stationary, like the shocks in the similar
objects near θ 1 Ori C. Some of the most intriguing cases are HH 505 (LL2) and HH 876
(LL6; see Fig. 18), while LL Ori itself (LL1) is associated with the HH 888 jet. Some of
the LL Ori objects also contain proplyds around their central stars; LL5 is an example
(Fig. 18). These HH jets and proplyds associated with LL Ori objects mean that instead
of a collision between the bulk H II region and a spherical stellar wind, the LL Ori
bows may also be affected by the photoevaporative proplyd flow or the bipolar jets
themselves.
37
Figure 19. This 79′′ × 85′′ HST ACS image prepared by John Bally was made
with 0.05′′ pixels and is color coded with red=[S II] 6717+6731 Å, green=Hα 6563
Å, and blue=[O III] 5007 Å . The bipolar jet from the brightest star, LL Ori is lost in
the overexposure of the stationary shock, which faces the brightest part of the Orion
Nebula, rather than θ1 Ori C. The bipolar jet is curved and one sees a series of shocks
that it forms, indicating that the flow from the jet is highly irregular.
Table 2.: Outflows in the Orion Nebula
Origin
Name
BN/KL
BN/KL
BN/KL
BN/KL
BN/KL
BN/KL
BN/KL
BN/KL(?)
OMC-1 S
OMC-1 S
OMC-1 S
OMC-1 S
OMC-1 S
OMC-1 S
OMC-1 S
OMC-1 S
OMC-1 S
OMC-1 S
α(2000)
δ(2000)
Comment
(5h )
(–5◦ )
H2 fingers
...
...
(many)
HH201
35 11
21 54
M42-HH1, finger
HH205-210
35 12
20 34
M42-HH5-10 fingers
HH601
35 14
20 29
HH602
35 13
20 35
HH603
35 11.5
21 07
HH604
35 11.6
21 37
HH400
35 34
31 55
large bow
HH202
35 11
23 00
M42-HH2
HH203/204
35 22
25 10
M42-HH3/4, crosses the bar
HH269
35 08
23 45
HH528
35 18
25 00
HH529
35 16
23 55
HH530
35 12
24 10
HH605
35 17
23 26
HH606
35 15.5
23 04
HH625
35 12
23 30
Schmid-Burgk
35 12
24 45
molecular jet
38
α(2000)
δ(2000)
Comment
(5h )
(–5◦ )
OMC-1 S
FMO
35 13
23 50
molecular jet
ONC
HH44
35 16
10 27
Large HH object
ONC
HH384
35 26
09 23
chain
ONC
HH502
35 28
29 30
microjet, proplyd, bent jet
ONC
HH503
34 49
31 46
microjet
ONC
HH504
35 04
29 26
microjet, one-sided, binary
ONC
HH505
34 41
22 42
LL 2
ONC
HH506
35 47
10 29
complex
ONC
HH508
35 16
23 07
ONC
HH510
35 11
23 27
microjet, proplyd, one-sided
ONC
HH511
35 13
22 47
microjet
ONC
HH512
35 16
25 33
microjet, proplyd
ONC
HH513
35 16
22 36
microjet, bipolar HH jet
ONC
HH514
35 17
23 37
proplyd HST2
ONC
HH515
35 17.6
25 43
microjet, proplyd
ONC
HH517
35 18
24 13
microjet, proplyd 182-413
ONC
HH519
35 20
25 04
microjet, proplyd, binary
ONC
HH520
35 20
25 06
microjet, proplyd
ONC
HH521
35 20.6
24 46
microjet, proplyd 197-427
ONC
HH522
35 24
23 34
ONC
HH523
35 18
23 27
ONC
HH524
35 24
24 40
ONC
HH525
35 25
24 36
bipolar
ONC
HH526
35 25
24 57
microjet, proplyd
ONC
HH527
35 28
24 58
microjet, proplyd, one-sided
ONC
HH532
35 46
09 53
ONC
HH533
35 26
09 22
chain
ONC
HH535
35 19
11 40
complex
ONC
HH536
35 18
12 40
ONC
HH537
35 01
14 07
bright bow
ONC
HH538
35 33
13 09
bow
ONC
HH539
35 37
11 43
ONC
HH540
35 18
28 26
microjet, beehive proplyd
ONC
HH540A
35 19
31 05
bow, jet from beehive
ONC
HH541
35 06
33 30
ONC
HH558
35 15
31 07
ONC
HH559
35 24
28 00
ONC
HH560
35 30
30 25
microjet, proplyd
ONC
HH561
35 20
30 39
microjet, MY Ori
ONC
HH667
35 21.6
09 39
microjet, silhouette disk
ONC
HH668
35 25.3
15 36
microjet, binary proplyd
ONC
HH725
35 15.5
23 38
microjet, proplyd
ONC
HH726
35 16.7
23 17
microjet, proplyd 167-317
ONC
HH873
35 20
23 59
complex
ONC
HH874
35 23
28 38
proplyd
ONC
HH875
35 31.4
28 16
LL 5
ONC
HH876
35 33
30 22
LL 6
ONC
HH877
34 36
21 46
ONC
HH878
35 02
26 36
C-symmetric
ONC
HH879
35 16
32 59
V1504
ONC
HH880
35 09
31 49
LM Ori
ONC
HH881
35 08
32 44
ONC
HH882
35 06
33 35
ONC
HH883
35 10
28 23
V484 Ori, proplyd
ONC
HH884
35 11
30 35
ONC
HH885
35 06
29 22
proplyd
ONC
HH886
34 47
26 05
microjet, proplyd
ONC
HH887
35 27
10 07
ref. nebula
ONC
HH888
35 05
25 20
LL Ori
Note—because of their extended nature, many coordinates for HH jets are approximate.
Origin
Name
39
6.
Summary and Conclusions
The proximity of the Orion Nebula and its associated cluster allows one to see phenomena that are not observed in more distant similar regions. It is a good prototype
for young stellar clusters with associated optically visible H II regions because the
observational selection effect is to see those clusters formed near the surface and on
the observer’s side of the parent molecular cloud. This geometry means that the nebula is largely in the background of the hot luminous stars, which means that one can
observe the placental material from star formation as proplyds. In the proplyds the
material surrounding the lower mass stars is seen in emission (gas component) through
photoionization or in silhouette (dust component) against the background nebular emission. Many of the stars show collimated outflows on scales as small as 100’s of AU and
as large as tenths of a parsec. These jets and their associated shocks formed in nebular
gas are quite different from those observed in less populous young clusters because of
external photoionization by the hottest stars. The presence of circumstellar material
around the proplyds near the hottest star (θ 1 Ori C) in the cluster argues that this star has
formed only recently or that it has only recently broken out of its surrounding nebular
material.
The eponymous Orion Nebula Cluster is but one of three centers of star formation
in this region. The second most luminous center is near the BN-KL region and the
geometry of outflow from it argues that it is buried about 0.2 pc behind the nebula’s
ionization front. There is evidence of a major energetic event about 500-1000 years
ago which produced both high velocity motion of certain compact radio sources and
also a large-scale flow of material that has produced highly structured features seen as
fingers pointing away from the source. The third center of star formation is the Orion-S
region, where one sees numerous bipolar molecular and atomic outflows. This region
is probably only a few hundredths of a parsec behind the nebula’s ionization front and
is located in a rise in the nebula’s surface that probably reflects the higher density there
in the parent molecular cloud.
Acknowledgements CRO wishes to acknowledge partial support during the preparation of this paper to HST grant GO 10967 and NS to grant GO 10421. We thank J. Di
Francesco for data in advance of publication and W. J. Henney for numerous comments
on a draft of this paper.
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