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A&A manuscript no.
(will be inserted by hand later)
Your thesaurus codes are:
06 (02.19.1; 08.03.4; 08.06.2; 09.08.1; 09.10.1; 13.09.6)
ASTRONOMY
AND
ASTROPHYSICS
August 29, 2000
ISO Spectroscopy of PDRs and Shocks in Star Forming Regions ?
M.E. van den Ancker 1;2 , A.G.G.M. Tielens 3;4 , and P.R. Wesselius 4
1 Astronomical Institute ``Anton Pannekoek'', University of Amsterdam, Kruislaan 403, NL--1098 SJ Amsterdam, The Netherlands
2 Harvard�Smithsonian Center for Astrophysics, 60 Garden Street, MS 42, Cambridge MA 02138, USA
3 Kapteyn Astronomical Institute, Groningen University, P.O. Box 800, NL--9700 AV Groningen, The Netherlands
4 SRON, P.O. Box 800, NL--9700 AV Groningen, The Netherlands
Received ; accepted
Abstract. We present the results of ISO SWS and LWS grat�
ing scans of pure rotational lines of H 2 , as well as the infrared
fine structure lines of [C II], [O I], [S I], [Si II], [Fe I] and [Fe II]
towards galactic star forming regions. We find that intense H 2
emission is only observed in the vicinity of stars with spectral
type earlier than B4 or in the neighbourhood of embedded YSOs
associated with outflows. We interpret these lines as arising in
either a photo�dissociation region (PDR) near the early�type
stars, or a shock caused by the interaction of an outflow with
the surrounding molecular cloud material. Molecular hydrogen
excitation temperatures and fine structure line intensities were
compared with those predicted by theoretical models for PDRs
and dissociative and non�dissociative shocks and to those ob�
served towards regions known to contain PDRs and shocks.
It is shown that both shocks and PDRs show a warm and hot
component in H 2 . The observed temperatures of the warm com�
ponent appear in general too high to be explained by current
dissociative shock models. We find that the mere detection of
[S I] 25.2 �m emission with ISO is an unambiguous sign of the
presence of shocked gas. Strong [C II] 157.7 �m can only be due
to a PDR. An evolutionary scenario is suggested in which the
circumstellar material around a young star changes from being
heated mechanically by shocks into heated by radiation from
the central star through a PDR as the star clears its surroundings.
Key words: Shock waves -- Circumstellar matter -- Stars: For�
mation -- H II regions -- ISM: Jets and outflows -- Infrared: Stars
1. Introduction
The infrared emission�line spectrum of Young Stellar Objects
(YSOs) is dominated by the interaction of the central object with
the remnants of the clouds from which it formed. The intense
UV radiation generated by accretion as well as by the central star
itself causes dissociation of molecular material close to the YSO
Send offprint requests to: M.E. van den Ancker (mario@astro.uva.nl)
? Based on observations with ISO, an ESA project with instruments
funded by ESA Member States (especially the PI countries: France,
Germany, the Netherlands and the United Kingdom) and with the
participation of ISAS and NASA.
and ionizes much of the atomic material, giving rise to typical
nebular and recombination lines. The strong stellar wind, often
collimated into a bipolar outflow, will cause a shock wave as
it hits the surrounding molecular cloud, heating the post�shock
gas sufficiently to cause strong molecular and ionic emission.
Neutral clouds irradiated by far�ultraviolet (FUV) photons
are known as photodissociation regions (PDRs; see Hollen�
bach & Tielens 1999 for a comprehensive review). In these
regions, heating of the gas occurs by collisions of photoelectri�
cally ejected electrons from grain surfaces. Cooling of the gas
occurs mainly through emission in atomic fine�structure and
molecular lines, reaching intensities that should be easily ob�
servable in a wide range of astronomical objects. The gas in the
surface regions of these PDRs reaches temperatures of typically
500 K (e.g. Tielens & Hollenbach 1985), making it possible
to collisionally excite the low�lying pure rotational levels of
molecular hydrogen, provided that the region has a sufficient
density and incident FUV flux. Since the lowest ro�vibrational
lines of H 2 have energy levels of the order of 5000 K, only
these low�lying pure rotational H 2 lines, hard to observe from
the ground, can reliable probe the physical conditions of the
dominant species and thus provide us with detailed information
of the physical conditions occurring in the PDR.
The physical situation in shock waves is quite distinct: here
the molecular gas is heated by compression of a supersonic wave
moving through the gas. Shocks are usually divided into two
distinct categories: J� or Jump�shocks, and C� or Continuous�
shocks. In the J�shocks viscous heating of the neutrals occurs
in a thin shock front in which radiative cooling is insignificant,
and the post�shock gas is heated to several times 10 4 degrees
(e.g. Hollenbach & McKee 1989), dissociating all molecular
material. In these J�shocks, the physical conditions (density,
temperature) change from their pre� to post�shock values within
one mean free path, causing an apparent discontinuity. Cooling
of the post�shock gas occurs trough atomic fine�structure lines.
C�shocks are magnetized, non�dissociative shocks in which
ions and the magnetic field are compressed ahead of the shock
front and are able to heat the neutral gas (containing only a
trace fraction of ions) to a few thousand degrees. In C�shocks
the physical conditions change more gradually from their pre�
to post�shock values and cooling is mainly through radiation

2 M.E. van den Ancker et al.: ISO Spectroscopy of PDRs and Shocks in Star Forming Regions
Table 1. Properties of programme stars.
Source ff (2000) ffi (2000) Region d Ref. Sp. Type T eff Ref. Outflow AV log L=L fi M? Age
[pc] [K] [ m ] [M fi ] [yr]
AFGL 490 03 27 38.6 +58 47 00 Cam OB1 900 (1) B1--6e 18000 (13) p
� 31 3.57 � 8 --
IRAS03260+3111 03 29 10.4 +31 21 58 NGC 1333 350 (2) 12000 (14) -- � 20 2.45 � 4 --
L1489 IRS 04 04 43.0 +26 18 58 L1491 140 (3) ? Ke 4000 (15) p
� 20 0.60 ! 2 --
L1551 IRS5 04 31 34.0 +18 08 04 Tau R2 140 (3) G--Ke 5500 (15) p
� 19 1.37 � 3 --
IRAS12496\Gamma7650 12 53 16.0 \Gamma77 07 02 Cha II 180 (4) Ae 8600 (16) p
� 14 1.68 � 2.5 --
GGD 27�ILL 18 19 12.0 \Gamma20 47 31 IC 1318 1700 (5) B1 25000 (17) p
� 20 4.29 � 14 --
LkHff 225 20 20 30.4 +41 21 27 Cyg R1 1000 (6) A--Fe 9000 (18) p
� 15 3.21 � 7 --
S106 IRS4 20 27 26.6 +37 22 48 Cyg OB2 1200 (7) O6--O8e 37000 (19) -- � 14 4.62 � 20 --
AFGL 2591 20 29 24.6 +40 11 19 Cyg X 1700 (8) ! O8 40000 (20) p
� 30 4.78 � 25 --
Cep A East 22 56 18.6 +62 02 00 Cep OB3 690 (9) B: 22000 (20) p
� 62 3.65 � 9 --
Elias 3�1 04 18 40.5 +28 19 17 L1495 140 (3) A6:e 8000 (21) -- 7.83 1.32 2.0 3 \Theta 10 6
T Tau 04 21 59.1 +19 32 06 Tau R2 140 (3) K0--1e 5200 (22) p
1.39 1.22 2.5 ! 10 6
HD 97048 11 08 04.6 \Gamma77 39 17 Ced 111 180 (10) B9.5Ve 10000 (23) -- 1.27 1.50 2.5 � 3 \Theta 10 6
HD 97300 11 09 50.6 \Gamma76 36 47 Ced 112 190 (10) B9V 10500 (24) -- 1.33 1.42 2.4 � 3 \Theta 10 6
HR 5999 16 08 34.2 \Gamma39 06 18 Lupus 3 210 (10) A5--7IIIe 7900 (25) -- 0.47 1.96 3.2 5 \Theta 10 5
R CrA 19 01 53.9 \Gamma36 57 10 NGC 6729 130 (11) A1--F7ev 8600 (10) p
2.26 1.98 3.0 1 \Theta 10 6
T CrA 19 01 58.8 \Gamma36 57 49 NGC 6729 130 (11) F0e 7200 (26) p
1.64 0.88 1.5 1 \Theta 10 7
WW Vul 19 25 58.6 +21 12 31 Vul R1 440 (12) A4IV--Ve 8400 (27) -- 0.96 1.33 2.0 5 \Theta 10 6
BD+40 ffi 4124 20 20 28.3 +41 21 51 Cyg R1 1000 (6) B2Ve 22000 (17) -- 3.01 3.16 7.0 � 10 5
LkHff 224 20 20 29.2 +41 21 27 Cyg R1 1000 (6) A7e 7900 (28) -- 2.98 2.27 4.0 1 \Theta 10 5
HD 200775 21 01 36.8 +68 09 49 NGC 7023 430 (10) B2.5IVe 20400 (29) p 1.92 3.36 9.0 2 \Theta 10 7
References to Table 1: (1) Harvey et al. (1979); (2) Herbig & Jones (1983); (3) Kenyon et al. (1994); (4) Whittet et al. (1997); (5) Rodr� �guez et
al. (1980); (6) Shevchenko et al. (1991); (7) Rayner (1994); (8) Dame & Thaddeus (1985); (9) Mel'nikov et al. (1995); (10) van den Ancker
et al. (1997, 1998); (11) Marraco & Rydgren (1981); (12) Voshchinnikov (1981); (13) Snell et al. (1984); (14) This paper; (15) Kenyon et al.
(1998); (16) Hughes et al. (1991); (17) Yamashita et al. (1987); (18) Hillenbrand et al. (1995); (19) van den Ancker et al. (2000b); (20) Lada et
al. (1984); (21) Zinnecker & Preibisch (1994); (22) Kenyon & Hartmann (1995); (23) Whittet et al. (1987); (24) Rydgren (1980); (25) Tjin A
Djie et al. (1989); (26) Finkenzeller & Mundt (1984); (27) Gray & Corbally (1998); (28) van den Ancker et al. (2000a); (29) Rogers et al. (1995).
from molecular material (e.g. Kaufman & Neufeld 1996). If the
temperatures in a C�shock become sufficiently high to start to
dissociate molecules, the cooling through the molecular lines
goes down, and the shock temperature increases until it turns
into a J�shock. Shocks with a velocity larger than 40 km s \Gamma1 are
usually J�shocks, whereas slower shocks are usually of C�type
(Chernoff et al. 1982).
In this paper we will study the infrared emission�line spec�
tra of a sample of YSOs observed with the Infrared Space
Observatory (ISO; Kessler et al. 1996). In order to interpret our
observations as arising in either a PDR, C�Shock or J�shock, we
compare measured line intensities with those predicted by the�
oretical models of such regions. Since the regions surrounding
our target sources are often highly confused, and the applica�
bility of the simplified models to a complex physical situation
is by no means sure, we will also compare our new ISO obser�
vations with those of relatively clean regions known to contain
a shock or a PDR. We will show that the distinction between a
shock and a PDR spectrum can be made with relative ease and
the infrared lines are a good diagnostic tool for constraining the
physical conditions in the material surrounding YSOs.
2. Sample selection
Young stellar objects are commonly divided into classes de�
pending on their energy distribution. Class I YSOs are heavily
embedded, optically invisible, but usually infrared bright ob�
jects. Class II YSOs are already optically visible, but also still
show excess infrared emission above photospheric levels due
to dust, heated by the central star, in the direct circumstellar
environment. Intermediate�mass (2--10 M fi ) Class II YSOs are
usually referred to as Herbig Ae/Be stars, whereas their lower
mass counterparts are the T Tauri stars. In the remainder of this
paper we will study the infrared emission line spectra in the
vicinity of a sample of 10 Class I YSOs and 11 Class II sources.
They are listed in Table 1. Because of the brightness limita�
tions imposed by the instruments used, the sample is biased
towards infrared bright objects, either because they are inher�
ently anomalously luminous in the infrared or because they are
located relatively nearby.
To understand the occurrence of infrared emission lines due
to PDRs or shocks in the vicinity of the stars in our sample, it
is useful to derive their basic properties in a consistent way. We
have therefore made new estimates of the bolometric luminos�

M.E. van den Ancker et al.: ISO Spectroscopy of PDRs and Shocks in Star Forming Regions 3
Fig. 1. Hertzsprung�Russell diagram for the YSOs included in this
study. Filled and open plot symbols indicate Class I and II sources,
respectively. Also shown are theoretical PMS evolutionary tracks
(solid/dashed lines) and the birthlines for 10 \Gamma4 (upper dotted line)
and 10 \Gamma5 M fi yr \Gamma1 (lower dotted line) by Palla & Stahler (1993).
ity of each source by constructing spectral energy distributions
(SEDs), using ultraviolet to submm data compiled from litera�
ture. A fit consisting of the sum of a reddened Kurucz (1991)
model for the stellar photosphere and a spline fitted to the ex�
cesses in the infrared to submm range, extrapolated to infinity by
a blackbody curve, was made to the energy distribution for each
source. Total bolometric luminosities were then determined by
computing 4�d 2 R
F � d�, with d the distance to the source and
F � the SED fit. The results of this procedure, as well as the
distances and effective temperatures used, are listed in Table 1.
For the extincted sources the resulting values of L bol do not
depend on the adopted value of T eff .
From the results listed in Table 1 we have created the
Hertzsprung�Russell diagram (HRD) shown in Fig. 1. In this
figure, we also have plotted the pre�main sequence evolutionary
tracks and the birthline by Palla & Stahler (1993). By comparing
the location of a source in the HRD with these computations, we
can determine stellar masses and ages for pre�main sequence
stars. For single Class II sources the required parameters (lu�
minosity and temperature) can be determined fairly accurately.
For the Class I sources, the temperature is usually not well de�
termined, but since the luminosity of an intermediate�mass star
stays relatively constant during its contraction to the zero�age
main sequence, the evolutionary tracks can still be used to esti�
mate its mass. Resulting stellar masses and ages are also listed
in Table 1.
Table 2. Log of ISO observations of programme stars.
Object AOT Rev. Date JD\Gamma2450000
AFGL 490 LWS 01 623 31/07/1997 660.879
SWS 01 640 17/08/1997 6780200
SWS 02 863 27/03/1998 899.939
IRAS03260+3111 SWS 01 659 05/09/1997 696.880
LWS 01 848 12/03/1998 884.909
L1489 IRS SWS 02 654 31/08/1997 691.635
L1551 IRS5 SWS 01 660 06/09/1997 697.769
SWS 02 668 14/09/1997 705.561
LWS 01 668 14/09/1997 705.571
IRAS12496\Gamma7650 LWS 01 115 11/03/1996 153.678
SWS 02 141 06/04/1996 179.909
SWS 01 233 06/07/1996 271.331
GGD27�ILL LWS 01 148 13/04/1996 187.161
SWS 02 149 14/04/1996 187.595
SWS 01 149 14/04/1996 187.621
LkHff 225 SWS 02 142 07/04/1996 180.996
LWS 01 142 07/04/1996 181.017
SWS 02 355 05/11/1996 393.306
SWS 01 858 22/03/1998 894.889
S106 IRS4 SWS 02 134 30/03/1996 172.705
LWS 01 134 30/03/1996 172.726
SWS 01 335 17/10/1996 373.740
AFGL 2591 SWS 02 134 30/03/1996 172.577
LWS 01 142 07/04/1996 180.737
SWS 01 357 07/11/1196 395.296
Cep A East SWS 02 220 23/06/1996 258.463
LWS 01 566 04/06/1997 603.681
SWS 02 566 04/06/1997 603.706
SWS 01 843 07/03/1998 880.050
Elias 3�1 SWS 02 667 13/09/1997 704.605
SWS 01 673 19/09/1997 710.564
T Tau LWS 01 672 18/09/1997 709.675
SWS 01 680 25/09/1997 717.411
SWS 02 681 27/09/1997 718.778
SWS 02 861 25/03/1998 897.973
HD 97048 SWS 01 141 06/04/1996 179.947
LWS 01 261 04/08/1996 299.565
HD 97300 SWS 02 141 06/04/1996 179.826
HR 5999 SWS 01 289 01/09/1996 327.627
SWS 02 289 01/09/1996 327.669
R CrA LWS 01 495 25/03/1997 533.284
SWS 02 704 19/10/1997 741.322
SWS 01 704 19/10/1997 741.353
T CrA SWS 02 141 06/04/1996 179.738
SWS 01 689 04/10/1997 726.278
WW Vul SWS 01 176 11/05/1996 214.507
SWS 02 176 11/05/1996 214.541
SWS 02 865 29/03/1998 901.889
BD+40 ffi 4124 SWS 02 142 07/04/1996 181.037
SWS 02 159 24/04/1996 197.635
SWS 01 355 05/11/1996 393.283
LWS 01 768 23/12/1997 805.525
LkHff 224 SWS 02 142 07/04/1996 180.960
SWS 01 858 22/03/1998 894.920
HD 200775 SWS 02 143 08/04/1996 182.108
SWS 01 339 21/10/1996 377.670
LWS 01 339 21/10/1996 377.695
3. Observations
As part of the ISO Guaranteed Time programme on YSOs
(``YSOAEBE'', P.I. P.R. Wesselius), we obtained spectroscopic
data centered on the positions of the 21 young stellar objects
listed in Table 1. Selected infrared transitions of H 2 , [Fe I],
[Fe II], [S I], [Si II], [O I] and [C II] were scanned with the Short
(SWS: 2.4--45 �m; de Graauw et al. 1996) andLongWavelength
(LWS: 43--197 �m; Clegg et al. 1996) Spectrometers on board
ISO. Most of the lines in the SWS wavelength range were
scanned with the full grating resolution (``AOT S02''), although
in some objects some were taken from lower resolution full

4 M.E. van den Ancker et al.: ISO Spectroscopy of PDRs and Shocks in Star Forming Regions
grating scans (``AOT S01''). All lines in the LWS wavelength
range were obtained from LWS full grating scans (``AOT L01'').
After the release of all ISO data to the public domain, we
inspected the ISO data archives for any additional spectroscopy
on our 21 targets and also included those in our data�set. A
full log of the ISO observations used in this study is given in
Table 2.
Data were reduced in a standard fashion using calibration
files corresponding to ISO off�line processing software (OLP)
version 7.0, after which they were corrected for remaining fring�
ing and glitches. To increase the S/N in the final spectra, de�
tectors were aligned and statistical outliers were removed, after
which the spectra were rebinned to a lower spectral resolution
(SWS) or simply averaged (LWS).
Scanned lines and measured line fluxes for the detected
lines or upper limits (total flux for line with peak flux 3oe) for
the undetected lines are listed in Table 3. For the sources in
which AOt S01 data is available, we also report the flux in the
6.2 �m UIR band. In a number of sources the H 2 0--0 S(2) line
was not detected because of strong fringing of the continuum
in this part of the spectrum, either due to the continuum source
being spatially extended, or due to the continuum source not
being in the center of the beam. The H 2 0--0 S(6) line has a
wavelength that puts it the middle of the H 2 O � 2 absorption
band, also hindering its detection in some sources (e.g. LkHff
225; van den Ancker et al. 2000a).
Accuracies of the absolute flux calibration in the SWS spec�
tra range from 7% in the short�wavelength (! 4.10 �m) part
to � 30% in the long wavelength (? 29 �m) part (Leech et
al. 1997). The LWS absolute flux calibration is expected to be
accurate at the 7% level (Trams et al. 1997). The errors listed
in Table 3 include both the error in the line flux measurement
itself and this absolute error. Note that the applied flux calibra�
tion is based on the assumption that we are looking at a point
source. For an extended source, the diffraction losses will be
underestimated and, in particular at the long�wavelength part of
the LWS, these corrections will exceed the quoted uncertainty.
4. Molecular hydrogen emission
From the H 2 line fluxes I(J) listed in Table 3 it is possible to
calculate the apparent column densities of molecular hydrogen
in the upper J levels, averaged over the SWS beam, N(J), using
N(J) = 4�I(J)
A
�
hc , with � the wavelength, h Planck's constant
and c the speed of light. The transition probabilities A were
taken from Turner et al. (1977). Line fluxes were corrected for
extinction using the average interstellar extinction law by Fluks
et al. (1994) using the values compiled in Table 1.
A useful representation of the H 2 data is then to plot the
log of N(J)=g, the apparent column density for a given J upper
level divided by the statistical weight, versus the energy of the
upper level. The resulting excitation diagrams are shown in
Fig. 2. For a Boltzmann distribution, the points in Fig. 2 should
form a nearly straight line. The slope of this line is inversely
proportional to the excitation temperature, while the intercept
is a measure of the total column density of warm gas. Since
the A coefficients for the H 2 lines are quite small, these lines
are optically thin and, because the critical densities are low, the
excitation temperature will be close to the kinetic temperature
of the gas.
The statistical weight g is a combination of the rotational
and nuclear spin components. For this, we have assumed the
high temperature equilibrium relative abundances of 3:1 for the
ortho and para forms of H 2 (Burton et al. 1992a). From Fig. 2 it
is apparent that for most sources, the H 2 points below 5000 K
do indeed form a straight line in the excitation diagram. The fact
that the points for ortho and para H 2 lie on the same line proves
that our assumption on their relative abundances is correct.
Using the formula by Parmar et al. (1991) and the rotational
constants by Dabrowski (1984), we have fitted Boltzmann dis�
tributions to the low�lying pure rotational lines. For a number
of sources, the pure rotational lines at the higher energy levels
can be seen to deviate strongly from this fit. In these cases, we
have attempted to characterize this behaviour by fitting a second
Boltzmann distribution to the higher energy level populations.
The resulting excitation temperatures and derived column and
mass of molecular hydrogen are listed in Table 5.
In Fig. 2 we also show similar plots for six regions which
have been quoted in literature as prime examples of shock�
and PDR emission, peak 1 of the shocked H 2 emission in the
Orion BN/KL region (Hasegawa et al. 1987; Brand et al. 1988;
Parmar et al. 1994; Rosenthal et al. 2000), the shock in the old
supernova remnant IC 443 (Burton et al. 1989; Richter et al.
1995a, 1995b; Oliva et al. 1999a; Cesarsky et al. 1999), the
X�ray irradiated shock in the young supernova remnant RCW
103 (Oliva et al. 1990, 1999b), the Orion bar PDR (Burton et
al. 1990a; Parmar et al. 1991; Marconi et al. 1998), the NGC
2023 PDR (Hasegawa et al. 1987; Burton et al. 1992b; Steiman�
Cameron et al. 1997) and the PDR S140 (Timmermann et al.
1996). The ground�based observations of these regions reveal
spatial structure in the derived temperature on a scale which
is small compared to the ISO SWS beam. Basic parameters
for these regions are listed in Table 4. As can be seen from
Fig. 2, the excitation diagrams of these template regions show
in general the warm and hot component also visible in some of
our ISO sources. The results of the Boltzmann distribution fits
are again given in Table 5.
Employing predictions of H 2 emission from PDR, J�shock
and C�shock models by Burton et al. (1992a), Hollenbach &
McKee (1989) and Kaufman & Neufeld (1996), we determined
the excitation temperature T rot from the low�lying (from S(1) to
S(5)) pure rotational levels from these models as a function of
density n and either incident FUV flux G (in units of the average
interstellar FUV field G 0 = 1.2 \Theta 10 \Gamma4 erg cm \Gamma2 s \Gamma1 sr \Gamma1 ;
Habing 1968) or shock velocity v s in an identical way as was
done for the observations. For PDRs, the total infrared flux F IR
is expected to be directly correlated to G by the relation G =
0.5 F IR
=\Omega 2 ,
with\Omega the spatial extent of the PDR. We have
computed F IR for all our sources by the same procedure as was
used to compute the bolometric luminosity in Sect. 2, but in this
case only use the excess above the reddened Kurucz model to
integrate F � . From these values of F IR we then computed G,

M.E. van den Ancker et al.: ISO Spectroscopy of PDRs and Shocks in Star Forming Regions 5
Table
3.
Observed
and
extinction�corrected
line
fluxes
(in
10
\Gamma16
Wm
\Gamma2
).
Line
�
Beam
AFGL
490
IRAS03260+3111
L1489
IRS
L1551�IRS5
IRAS12496\Gamma7650
[�m]
[10
\Gamma8
sr]
Obs.
Ext.
corr.
Obs.
Ext.
Corr.
Obs.
Ext.
Corr.
Obs.
Ext.
Corr.
Obs.
Ext.
Corr.
H2
0--0
S(0)
28.2188
1.64
!39.9
!40.3
!5.65
!5.71
!5.59
!5.64
!20.6
!20.8
!7.89
!7.97
H2
0--0
S(1)
17.0348
1.15
!16.9
!17.2
16.6\Sigma3.6
16.9\Sigma3.6
!1.07
!1.09
!1.94
!1.97
1.11\Sigma0.36
1.13\Sigma0.37
H2
0--0
S(2)
12.2786
1.15
!52.3
!53.5
33.0\Sigma9.4
33.7\Sigma9.7
!21.1
!21.7
!16.2
!16.6
!20.1
!20.5
H2
0--0
S(3)
9.6649
0.85
3.77\Sigma1.24
3.99\Sigma1.32
21.7\Sigma5.5
22.9\Sigma5.8
!2.03
!2.15
!1.97
!2.08
2.51\Sigma0.83
2.70\Sigma0.89
H2
0--0
S(4)
8.0251
0.85
!15.9
!16.2
10.2\Sigma2.7
10.4\Sigma2.8
!2.81
!2.87
!3.04
!3.10
!2.96
!3.02
H2
0--0
S(5)
6.9095
0.85
9.62\Sigma2.43
9.72\Sigma2.45
20.1\Sigma3.8
20.4\Sigma3.8
!2.75
!2.78
!2.77
!2.80
!4.96
!5.01
H2
0--0
S(6)
6.1086
0.85
!23.0
!23.2
!50.2
!50.8
--
--
!13.9
!14.1
!12.0
!12.1
H2
0--0
S(7)
5.5112
0.85
!30.7
!31.1
!17.3
!17.6
!13.5
!13.7
!10.0
!10.1
!13.0
!13.2
H2
0--0
S(8)
5.0531
0.85
!23.1
!23.3
!7.70
!7.82
--
--
!4.78
!4.85
!22.5
!22.7
H2
0--0
S(9)
4.6946
0.85
!27.8
!28.2
!8.67
!8.82
--
--
!6.93
!7.03
!59.4
!60.5
H2
0--0
S(10)
4.4099
0.85
!9.50
!9.69
!5.07
!5.16
--
--
!6.78
!6.92
!12.8
!13.1
H2
0--0
S(11)
4.1813
0.85
!10.7
!10.9
!11.0
!11.2
--
--
!13.9
!14.2
!7.56
!7.72
[Fe
I]
(
5
D4--
5
D3)
24.0424
1.15
!24.2
!24.6
!5.06
!5.13
--
--
!3.16
!3.20
!1.47
!1.49
[Fe
I]
(
5
D3--
5
D2)
34.7135
2.01
!7.25
!7.30
!12.4
!12.5
!5.72
!5.76
!7.22
!7.27
!4.58
!4.62
[Fe
II]
(
4
F9=2
--
4
F7=2
)
17.9363
1.15
!3.44
!3.51
!2.36
!2.41
--
--
5.43\Sigma1.81
5.54\Sigma1.85
0.62\Sigma0.20
0.63\Sigma0.21
[Fe
II]
(
6
D9=2
--
6
D7=2 )
25.9882
1.15
!24.3
!24.6
!3.31
!3.35
--
--
12.6\Sigma4.2
12.8\Sigma4.2
0.75\Sigma0.25
0.76\Sigma0.25
[Fe
II]
(
6
D7=2
--
6
D5=2
)
35.3491
2.01
!7.63
!7.69
!7.30
!7.35
--
--
!26.4
!26.5
!5.15
!5.18
[S
I]
(
3
P2--
3
P1)
25.2490
1.15
!9.91
!10.0
!6.97
!7.06
!2.05
!2.08
!8.26
!8.37
!2.56
!2.59
[Si
II]
(
2
P1=2
--
2
P3=2 )
34.8152
2.01
5.64\Sigma1.86
5.68\Sigma1.87
27.2\Sigma9.0
27.4\Sigma9.1
!7.20
!7.26
20.1\Sigma6.6
20.3\Sigma6.7
5.37\Sigma1.77
5.41\Sigma1.79
[OI]
(
3
P2--
3
P1)
63.1850
16.3
222.0\Sigma15.5
222.7\Sigma15.6
1083\Sigma76
1086\Sigma76
--
--
111.2\Sigma7.8
111.6\Sigma7.8
43.0\Sigma3.0
43.1\Sigma3.0
[OI]
(
3
P1--
3
P0)
145.535
8.82
26.9\Sigma4.6
26.9\Sigma4.6
190.8\Sigma13.4
190.9\Sigma13.4
--
--
!18.7
!18.7
3.10\Sigma0.80
3.10\Sigma0.80
[C
II]
(
2
P1=2
--
2
P3=2
)
157.741
11.6
62.1\Sigma5.5
62.2\Sigma5.5
323.1\Sigma22.6
323.2\Sigma22.6
--
--
!19.1
!19.1
4.00\Sigma1.00
4.01\Sigma1.00
UIR
(C--C
stretch)
6.22
0.85
!484
!488
1327\Sigma99
1359\Sigma102
--
--
!691
!698
!416
!420
Line
�
Beam
GGD27�ILL
LkHff
225
S106
IRS4
AFGL
2591
Cep
A
East
[�m]
[10
\Gamma8
sr]
Obs.
Ext.
corr.
Obs.
Ext.
Corr.
Obs.
Ext.
Corr.
Obs.
Ext.
Corr.
Obs.
Ext.
Corr.
H2
0--0
S(0)
28.2188
1.64
!12.1
!12.3
!13.8
!14.3
!15.8
!16.7
!124.2
!125.5
!36.2
!36.9
H2
0--0
S(1)
17.0348
1.15
5.57\Sigma1.84
5.68\Sigma1.87
7.76\Sigma2.56
8.23\Sigma2.72
5.76\Sigma1.75
7.54\Sigma2.29
17.9\Sigma4.7
18.3\Sigma4.8
2.73\Sigma0.90
2.84\Sigma0.94
H2
0--0
S(2)
12.2786
1.15
!41.5
!42.5
21.8\Sigma7.2
23.3\Sigma7.7
11.8\Sigma4.9
16.1\Sigma6.6
!189.9
!194.3
12.7\Sigma4.2
13.3\Sigma4.4
H2
0--0
S(3)
9.6649
0.85
5.63\Sigma1.86
5.95\Sigma1.96
28.0\Sigma7.3
33.1\Sigma8.7
11.9\Sigma3.9
25.5\Sigma8.4
13.5\Sigma3.7
14.3\Sigma3.9
9.02\Sigma2.98
10.1\Sigma3.3
H2
0--0
S(4)
8.0251
0.85
3.35\Sigma1.10
3.41\Sigma1.13
17.1\Sigma4.6
18.1\Sigma4.8
8.16\Sigma2.69
10.5\Sigma3.5
!106.3
!108.3
13.3\Sigma3.3
13.8\Sigma3.4
H2
0--0
S(5)
6.9095
0.85
10.6\Sigma2.2
10.7\Sigma2.2
40.3\Sigma6.3
41.6\Sigma6.5
14.7\Sigma4.8
16.9\Sigma5.6
!117.4
!118.7
24.2\Sigma4.3
24.7\Sigma4.3
H2
0--0
S(6)
6.1086
0.85
!6.98
!7.06
!25.6
!26.5
!16.0
!18.9
!214.6
!217.2
8.41\Sigma2.78
8.62\Sigma2.84
H2
0--0
S(7)
5.5112
0.85
!50.2
!50.9
35.1\Sigma7.9
36.6\Sigma8.2
9.36\Sigma3.09
11.3\Sigma3.7
!157.5
!159.6
24.8\Sigma6.2
25.5\Sigma6.3
H2
0--0
S(8)
5.0531
0.85
!9.86
!10.0
8.14\Sigma2.69
8.53\Sigma2.81
!2.95
!3.64
!64.4
!65.4
4.61\Sigma1.25
4.75\Sigma1.29
H2
0--0
S(9)
4.6946
0.85
!10.9
!11.0
30.4\Sigma10.0
32.0\Sigma10.6
!5.75
!7.28
!783.4
!797.1
!11.7
!12.1
H2
0--0
S(10)
4.4099
0.85
!25.1
!25.6
!6.33
!6.70
!4.81
!6.23
!32.5
!33.1
9.79\Sigma3.23
10.2\Sigma3.35
H2
0--0
S(11)
4.1813
0.85
!44.4
!45.3
6.72\Sigma2.22
7.15\Sigma2.36
!7.75
!10.3
!30.7
!31.3
!6.41
!6.68
[Fe
I]
(
5
D4--
5
D3)
24.0424
1.15
!53.3
!54.0
!5.80
!6.05
!10.1
!12.2
!45.9
!46.5
!11.9
!12.3
[Fe
I]
(
5
D3--
5
D2)
34.7135
2.01
!5.38
!5.42
!14.5
!14.8
!59.9
!67.3
!33.9
!34.2
!47.4
!48.1
[Fe
II]
(
4
F9=2
--
4
F7=2
)
17.9363
1.15
!19.0
!19.4
!1.64
!1.74
8.94\Sigma2.63
11.9\Sigma3.5
!6.15
!6.28
2.39\Sigma0.65
2.49\Sigma0.67
[Fe
II]
(
6
D9=2
--
6
D7=2 )
25.9882
1.15
!1.84
!1.86
3.33\Sigma1.10
3.45\Sigma1.14
52.7\Sigma11.0
62.0\Sigma12.9
!62.1
!62.8
6.42\Sigma2.08
6.58\Sigma2.13
[Fe
II]
(
6
D7=2
--
6
D5=2 )
35.3491
2.01
!12.7
!12.8
!9.22
!9.43
49.4\Sigma16.3
54.6\Sigma18.0
!39.4
!39.7
!40.1
!40.7
[S
I]
(
3
P2--
3
P1)
25.2490
1.15
!3.35
!3.39
5.75\Sigma1.63
5.97\Sigma1.69
!29.6
!35.2
!59.8
!60.6
7.02\Sigma2.10
7.20\Sigma2.15
[Si
II]
(
2
P1=2
--
2
P3=2 )
34.8152
2.01
23.1\Sigma7.6
23.3\Sigma7.7
30.5\Sigma10.1
31.2\Sigma10.3
1523\Sigma251
1687\Sigma279
95.2\Sigma31.4
95.9\Sigma31.7
58.6\Sigma4.1
59.4\Sigma4.2
[OI]
(
3
P2--
3
P1)
63.1850
16.3
945.7\Sigma66.2
948.5\Sigma66.4
659.7\Sigma47.8
665.5\Sigma48.3
11253\Sigma1138
11717\Sigma1185
782.5\Sigma54.8
784.8\Sigma54.9
5348\Sigma298
5380\Sigma299
[OI]
(
3
P1--
3
P0)
145.535
8.82
189.5\Sigma17.8
189.6\Sigma17.8
!94.8
!95.0
2888\Sigma290
2912\Sigma292
225.7\Sigma25.0
225.8\Sigma25.0
!284.4
!285.0
[C
II]
(
2
P1=2
--
2
P3=2
)
157.741
11.6
294.6\Sigma20.6
294.8\Sigma20.6
625.4\Sigma46.3
626.3\Sigma46.4
1572\Sigma159
1583\Sigma160
395.3\Sigma27.7
395.5\Sigma27.7
2245\Sigma232
2247\Sigma232
UIR
(C--C
stretch)
6.22
0.85
774\Sigma58
792\Sigma59
!307
!317
2214\Sigma166
2267\Sigma170
!234
!240
!253
!259

6 M.E. van den Ancker et al.: ISO Spectroscopy of PDRs and Shocks in Star Forming Regions
Table
3.
(Continued)
Line
�
Beam
Elias3--1
T
Tau
HD
97048
HD
97300
HR
5999
[�m]
[10
\Gamma8
sr]
Obs.
Ext.
corr.
Obs.
Ext.
Corr.
Obs.
Ext.
Corr.
Obs.
Ext.
Corr.
Obs.
Ext.
Corr.
H2
0--0
S(0)
28.2188
1.64
!4.72
!5.12
!4.67
!4.74
!5.83
!5.91
!2.33
!2.37
!4.97
!5.00
H2
0--0
S(1)
17.0348
1.15
!4.10
!4.78
1.80\Sigma0.59
1.85\Sigma0.61
!2.07
!2.12
--
--
!1.50
!1.52
H2
0--0
S(2)
12.2786
1.15
!19.9
!23.7
!9.68
!9.99
!5.40
!5.56
--
--
!20.0
!20.2
H2
0--0
S(3)
9.6649
0.85
!2.29
!3.54
6.81\Sigma1.95
7.36\Sigma2.10
!1.83
!1.97
!3.06
!3.29
!0.59
!0.61
H2
0--0
S(4)
8.0251
0.85
!3.44
!3.97
3.17\Sigma1.05
3.25\Sigma1.07
!2.20
!2.25
!2.15
!2.20
!4.22
!4.25
H2
0--0
S(5)
6.9095
0.85
!6.26
!6.79
13.6\Sigma2.6
13.8\Sigma2.7
!6.17
!6.25
!5.91
!5.99
!4.60
!4.63
H2
0--0
S(6)
6.1086
0.85
!9.12
!10.0
!8.07
!8.21
!10.2
!10.3
!9.19
!9.34
!10.6
!10.7
H2
0--0
S(7)
5.5112
0.85
!9.66
!10.7
11.0\Sigma2.6
11.2\Sigma2.7
!9.28
!9.45
!15.7
!16.0
!16.2
!16.3
H2
0--0
S(8)
5.0531
0.85
!23.1
!26.1
!4.52
!4.62
!4.53
!4.62
--
--
!12.2
!12.3
H2
0--0
S(9)
4.6946
0.85
!26.9
!30.8
2.87\Sigma0.95
2.94\Sigma0.97
!5.47
!5.59
--
--
!15.4
!15.5
H2
0--0
S(10)
4.4099
0.85
!5.66
!6.56
!3.13
!3.21
!4.99
!5.11
!7.06
!7.24
!10.7
!10.8
H2
0--0
S(11)
4.1813
0.85
!9.10
!10.7
!4.67
!4.81
!5.25
!5.39
!9.07
!9.33
!5.07
!5.13
[Fe
I]
(
5
D4--
5
D3)
24.0424
1.15
!2.13
!2.37
!1.83
!1.87
!3.02
!3.07
!1.52
!1.54
!0.89
!0.90
[Fe
I]
(
5
D3--
5
D2)
34.7135
2.01
!4.70
!4.98
!1.04
!1.05
!3.45
!3.48
!2.32
!2.35
!5.85
!5.87
[Fe
II]
(
4
F9=2
--
4
F7=2
)
17.9363
1.15
!6.74
!7.94
0.52\Sigma0.17
0.53\Sigma0.18
!1.85
!1.90
--
--
!1.43
!1.44
[Fe
II]
(
6
D9=2
--
6
D7=2
)
25.9882
1.15
!2.29
!2.51
1.85\Sigma0.61
1.88\Sigma0.62
!2.86
!2.90
!1.41
!1.43
!1.26
!1.27
[Fe
II]
(
6
D7=2
--
6
D5=2
)
35.3491
2.01
!3.86
!4.08
4.23\Sigma1.40
4.28\Sigma1.41
!4.98
!5.03
--
--
!6.98
!7.00
[S
I]
(
3
P2--
3
P1)
25.2490
1.15
!2.59
!2.86
4.09\Sigma1.28
4.16\Sigma1.30
!1.80
!1.83
!0.70
!0.71
!1.24
!1.25
[Si
II]
(
2
P1=2
--
2
P3=2
)
34.8152
2.01
!6.32
!6.70
12.6\Sigma4.2
12.7\Sigma4.2
!3.80
!3.84
!2.24
!2.26
!2.63
!2.64
[OI]
(
3
P2--
3
P1)
63.1850
16.3
--
--
234.8\Sigma23.5
235.8\Sigma23.6
19.0\Sigma3.0
19.1\Sigma3.0
--
--
--
--
[OI]
(
3
P1--
3
P0)
145.535
8.82
--
--
!14.9
!14.9
--
--
--
--
--
--
[C
II]
(
2
P1=2
--
2
P3=2
)
157.741
11.6
--
--
10.3\Sigma1.9
10.3\Sigma1.9
7.01\Sigma1.04
7.02\Sigma1.04
--
--
--
--
UIR
(C--C
stretch)
6.22
0.85
413\Sigma31
453\Sigma34
!161
!164
440\Sigma33
446\Sigma34
--
--
!293
!294
Line
�
Beam
R
CrA
T
CrA
WWVul
BD+40
ffi
4124
LkHff
224
HD
200775
[�m]
[10
\Gamma8
sr]
Obs.
Ext.
corr.
Obs.
Ext.
Corr.
Obs.
Ext.
Corr.
Obs.
Ext.
Corr.
Obs.
Ext.
Corr.
Obs.
Ext.
Corr.
H2
0--0
S(0)
28.2188
1.64
!11.4
!11.7
!6.04
!6.14
!1.78
!1.80
!6.60
!6.81
!9.51
!9.81
!6.04
!6.17
H2
0--0
S(1)
17.0348
1.15
!10.7
!11.2
3.96\Sigma1.31
4.09\Sigma1.35
!0.75
!0.77
5.37\Sigma1.77
5.69\Sigma1.88
10.1\Sigma3.3
10.7\Sigma3.5
!3.24
!3.37
H2
0--0
S(2)
12.2786
1.15
!46.0
!48.3
5.58\Sigma1.68
5.80\Sigma1.74
!9.84
!10.1
7.12\Sigma2.35
7.62\Sigma2.51
12.7\Sigma4.2
13.6\Sigma4.5
!5.32
!5.56
H2
0--0
S(3)
9.6649
0.85
!6.40
!7.25
!2.09
!2.29
!0.95
!1.01
7.74\Sigma2.36
9.15\Sigma2.79
9.81\Sigma3.12
11.6\Sigma3.7
!2.13
!2.37
H2
0--0
S(4)
8.0251
0.85
!20.3
!21.2
!3.14
!3.23
!4.34
!4.42
1.02\Sigma0.34
1.08\Sigma0.36
6.89\Sigma2.27
7.28\Sigma2.40
!2.75
!2.85
H2
0--0
S(5)
6.9095
0.85
!28.9
!29.6
!4.04
!4.11
!1.65
!1.67
2.45\Sigma0.81
2.53\Sigma0.83
11.7\Sigma2.5
12.1\Sigma2.6
!5.22
!5.33
H2
0--0
S(6)
6.1086
0.85
!19.5
!20.0
!20.1
!20.5
!5.05
!5.11
!7.96
!8.25
!6.26
!6.49
!7.65
!7.83
H2
0--0
S(7)
5.5112
0.85
!25.9
!26.7
!10.7
!10.9
!14.7
!14.9
!8.10
!8.44
!13.0
!13.5
!12.2
!12.6
H2
0--0
S(8)
5.0531
0.85
!12.3
!12.7
!6.76
!6.93
!15.7
!16.0
!18.3
!19.2
!35.5
!37.2
!7.06
!7.27
H2
0--0
S(9)
4.6946
0.85
!17.6
!18.3
!4.86
!5.00
!15.5
!15.8
!11.6
!12.2
!25.5
!26.9
!17.3
!17.9
H2
0--0
S(10)
4.4099
0.85
!35.2
!36.8
!6.10
!6.29
!7.37
!7.50
!4.17
!4.41
!4.15
!4.40
!10.3
!10.7
H2
0--0
S(11)
4.1813
0.85
!30.0
!31.5
!6.70
!6.93
!10.5
!10.7
!7.37
!7.85
!10.8
!11.5
!18.4
!19.2
[Fe
I]
(
5
D4--
5
D3)
24.0424
1.15
!5.02
!5.18
!0.68
!0.69
!1.26
!1.28
!1.72
!1.79
!1.46
!1.53
!2.05
!2.10
[Fe
I]
(
5
D3--
5
D2)
34.7135
2.01
!8.35
!8.49
!4.14
!4.19
!3.90
!3.93
!4.00
!4.09
!4.86
!4.97
!3.89
!3.94
[Fe
II]
(
4
F9=2
--
4
F7=2
)
17.9363
1.15
!3.96
!4.15
!0.89
!0.92
!0.85
!0.86
!0.97
!1.03
!1.86
!1.98
!1.10
!1.15
[Fe
II]
(
6
D9=2
--
6
D7=2
)
25.9882
1.15
1.33\Sigma0.44
1.36\Sigma0.45
!1.14
!1.17
!1.10
!1.12
0.69\Sigma0.23
0.71\Sigma0.23
!2.23
!2.31
!1.89
!1.93
[Fe
II]
(
6
D7=2
--
6
D5=2
)
35.3491
2.01
!4.76
!4.84
!4.07
!4.12
!4.19
!4.22
!4.40
!4.50
!7.70
!7.86
!3.81
!3.87
[S
I]
(
3
P2--
3
P1)
25.2490
1.15
6.78\Sigma2.05
6.97\Sigma2.10
!0.82
!0.84
!0.65
!0.65
!1.37
!1.42
1.60\Sigma0.52
1.66\Sigma0.54
!1.26
!1.29
[Si
II]
(
2
P1=2
--
2
P3=2
)
34.8152
2.01
9.38\Sigma3.09
9.54\Sigma3.15
!3.63
!3.67
!2.92
!2.94
23.2\Sigma7.7
23.7\Sigma7.8
24.2\Sigma8.0
24.8\Sigma8.2
19.0\Sigma6.3
19.3\Sigma6.4
[OI]
(
3
P2--
3
P1)
63.1850
16.3
681.4\Sigma47.6
684.9\Sigma47.9
--
--
--
--
703.6\Sigma49.5
709.9\Sigma49.9
--
--
714.4\Sigma72.8
718.5\Sigma73.2
[OI]
(
3
P1--
3
P0)
145.535
8.82
!40.0
!40.1
--
--
--
--
41.5\Sigma4.0
41.6\Sigma4.0
--
--
297.4\Sigma29.8
297.7\Sigma29.8
[C
II]
(
2
P1=2
--
2
P3=2
)
157.741
11.6
35.7\Sigma5.5
35.8\Sigma5.5
--
--
--
--
517.5\Sigma36.4
518.3\Sigma36.4
--
--
329.6\Sigma33.1
330.0\Sigma33.1
UIR
(C--C
stretch)
6.22
0.85
!270
!277
!106
!108
!192
!194
260\Sigma20
269\Sigma20
!135
!140
!132
!135

M.E. van den Ancker et al.: ISO Spectroscopy of PDRs and Shocks in Star Forming Regions 7
Table
4.
Observed
(Obs.)
and
extinction
corrected
(EC)
data
for
shocks
and
PDRs
from
literature.
Numbers
in
brackets
indicate
the
errors
in
the
last
digits.
Orion
BN/KL
IC
443
RCW
103
Line
Obs.
EC
Beam
Ref.
Obs.
EC
Beam
Ref.
Obs.
EC
Beam
Ref.
[10
\Gamma3
erg
s
\Gamma1
cm
\Gamma2
sr
\Gamma1
]
[10
\Gamma
8
sr]
[10
\Gamma3
erg
s
\Gamma1
cm
\Gamma2
sr
\Gamma1
]
[10
\Gamma
8
sr]
[10
\Gamma3
erg
s
\Gamma1
cm
\Gamma2
sr
\Gamma1
]
[10
\Gamma
8
sr]
H2
0--0
S(1)
17.03
�m
3.3(6)
4.0(7)
0.01
(1)
--
--
--
--
0.11(3)
0.12(3)
1.15
(13)
H2
1--0
S(1)
2.122
�m
8.4(8)
18.0(17)
0.06
(2)
0.637(2)
1.172(3)
0.05
(7)
0.386(12)
0.544(17)
0.11
(14)
H2
2--1
S(1)
2.248
�m
0.85(10)
1.67(20)
0.06
(2)
0.0547(5)
0.0937(9)
0.05
(7)
0.047(10)
0.064(14)
0.11
(14)
[S
I]
25.25
�m
3.9(6)
4.4(7)
2.45
(3)
--
--
--
--
!
0.01
!
0.01
1.15
(13)
[Si
II]
34.82
�m
6.1(6)
6.6(6)
3.27
(3)
0.26(3)
0.27(4)
2.01
(8)
1.4(2)
1.4(2)
2.01
(13)
[O
I]
63.18
�m
56(20)
58(21)
2.15
(4)
2.26(19)
2.31(19)
2.51
(9)
0.44(5)
0.44(5)
15.3
(13)
[O
I]
145.5
�m
--
--
--
--
--
--
--
--
0.037(8)
0.037(8)
15.3
(13)
[C
II]
157.7
�m
3.9(5)
3.9(5)
7.24
(5)
0.39(8)
0.39(8)
2.61
(10)
0.54(6)
0.54(6)
15.3
(13)
AV
[
m
]
10.0
(1)
8.0
(11)
4.5
(14)
d
[pc]
430
(6)
1500
(12)
6600
(15)
Orion
Bar
NGC
2023
S140
Line
Obs.
EC
Beam
Ref.
Obs.
EC
Beam
Ref.
Obs.
EC
Beam
Ref.
[10
\Gamma3
erg
s
\Gamma1
cm
\Gamma2
sr
\Gamma1
]
[10
\Gamma
8
sr]
[10
\Gamma3
erg
s
\Gamma1
cm
\Gamma2
sr
\Gamma1
]
[10
\Gamma
8
sr]
[10
\Gamma3
erg
s
\Gamma1
cm
\Gamma2
sr
\Gamma1
]
[10
\Gamma
8
sr]
H2
0--0
S(1)
17.03
�m
0.110(5)
0.121(6)
0.06
(16)
0.022(4)
0.022(4)
1.15
(21)
0.020(6)
0.022(6)
1.15
(25)
H2
1--0
S(1)
2.122
�m
0.078(13)
0.114(19)
0.15
(17)
0.060(2)
0.067(2)
0.92
(22)
0.16(2)
0.22(3)
0.03
(26)
H2
2--1
S(1)
2.248
�m
0.008(4)
0.011(6)
0.15
(18)
0.016(2)
0.018(2)
0.92
(22)
--
--
--
--
[S
I]
25.25
�m
!
0.13
!
0.14
2.45
(10)
--
--
--
--
--
--
--
[Si
II]
34.82
�m
9.0(5)
9.3(5)
5.51
(19)
0.20(11)
0.20(11)
3.10
(23)
0.030(9)
0.032(9)
2.01
(25)
[O
I]
63.18
�m
54(5)
55(5)
2.15
(5)
3.62(10)
3.63(10)
2.61
(23)
0.33(10)
0.33(10)
15.3
(27)
[O
I]
145.5
�m
4.2(5)
4.2(5)
6.72
(5)
0.23(4)
0.23(4)
4.42
(23)
--
--
--
--
[C
II]
157.7
�m
3.4(5)
3.4(5)
7.24
(5)
0.68(4)
0.68(4)
4.22
(23)
0.35(11)
0.35(11)
15.3
(27)
AV
[
m
]
5.0
(20)
1.4
(24)
4.0
(25)
d
[pc]
430
(6)
475
(24)
910
(28)
References
to
Table
4:
(1)
Parmar
et
al.
(1994);
(2)
Brand
et
al.
(1988);
(3)
Haas
et
al.
(1991);
(4)
Crawford
et
al.
(1986);
(5)
Stacey
et
al.
(1993);
(6)
Warren
&
Hesser
(1978);
(7)
Richter
et
al.
(1995a);
(8)
Oliva
et
al.
(1999a);
(9)
Burton
et
al.
(1990b);
(10)
Haas
(1997);
(11)
van
Dishoeck
et
al.
(1993);
(12)
Burton
et
al.
(1988);
(13)
Oliva
et
al.
(1999b);
(14)
Oliva
et
al.
(1990);
(15)
Leibowitz
&
Danziger
(1983);
(16)
Parmar
et
al.
(1991);
(17)
Burton
et
al.
(1990a);
(18)
Hayashi
et
al.
(1985);
(19)
Haas
et
al.
(1986);
(20)
Assumed;
(21)
This
paper;
(22)
Hasegawa
et
al.
(1987);
(23)
Steiman�Cameron
et
al.
(1997);
(24)
Jaffe
et
al.
(1990);
(25)
Timmermann
et
al.
(1996).
(26)
Carr
(1990);
(27)
Emery
et
al.
(1996);
(28)
Crampton
&
Fisher
(1974).

8 M.E. van den Ancker et al.: ISO Spectroscopy of PDRs and Shocks in Star Forming Regions
Fig. 2. H2 excitation diagrams for programme and comparison stars. Shown are apparent columns of H2 in the pure�rotational (0--0) transitions
(filled dots), 1--0 transitions (open squares), 2--1 transitions (filled triangles), 2--0 transitions (open stars), 3--2 transitions (open triangles), 3--0
transitions (filled squares), 4--3 transitions (open diamonds), and 4--1 transitions (open crosses). Observational errors are smaller than the size
of the plot symbol. The Boltzmann distribution fits are plotted as dashed lines. The solid lines show the sum of both thermal components for
each source.

M.E. van den Ancker et al.: ISO Spectroscopy of PDRs and Shocks in Star Forming Regions 9
Fig. 3. Comparison of observed H2 rotational temperatures to theo�
retical relation between continuum fluxes and T rot (PDR models) or
summed intensity in all observed lines and T rot (shocks). Sources
which show PAH emission are plotted as squares. Plot symbols are
filled for Class I sources. The arrows show the direction of beam
dilution.
assuming a beam filling factor of one. The resulting G�values,
listed in Table 6, were used to plot our observations directly in
the parameter space of the PDR models (Fig. 3). Note that if
the extent of the PDR is smaller than the beam size, as might
be expected for some sources, our data points in Fig. 4 may be
shifted along the direction indicated by the arrows. However,
for PDRs the resulting T rot does not depend much on G, so this
does not introduce an additional uncertainty in our analysis.
Table 5. Results of H2 Boltzmann distribution fits.
Object T rot N rot T vib N vib M(H2 )
[K] [cm \Gamma2 ] [K] [cm \Gamma2 ] [M fi ]
AFGL 490 1400 1.9 \Theta 10 18 -- -- 2 \Theta 10 \Gamma4
IRAS03260+3111 600 1.2 \Theta 10 20 -- -- 2 \Theta 10 \Gamma3
IRAS12496\Gamma7650 540 9.1 \Theta 10 18 -- -- 5 \Theta 10 \Gamma4
GGD 27�ILL 650 4.2 \Theta 10 19 -- -- 1 \Theta 10 \Gamma2
LkHff 225 720 6.7 \Theta 10 19 2200 1.7 \Theta 10 18 8 \Theta 10 \Gamma3
S106 IRS4 560 8.7 \Theta 10 19 3800 3.6 \Theta 10 17 8 \Theta 10 \Gamma3
AFGL 2591 390 2.2 \Theta 10 20 -- -- 2 \Theta 10 \Gamma2
Cep A East 730 5.0 \Theta 10 19 5200 3.0 \Theta 10 17 3 \Theta 10 \Gamma3
T Tau 440 1.5 \Theta 10 19 1500 2.1 \Theta 10 18 4 \Theta 10 \Gamma5
T CrA 770 2.7 \Theta 10 19 -- -- 1 \Theta 10 \Gamma4
BD+40 ffi 4124 490 4.8 \Theta 10 19 -- -- 5 \Theta 10 \Gamma3
LkHff 224 610 7.0 \Theta 10 19 -- -- 8 \Theta 10 \Gamma3
Orion BN/KL 390 6.1 \Theta 10 21 3300 1.7 \Theta 10 19 2 \Theta 10 \Gamma3
IC 443 610 6.5 \Theta 10 20 3400 1.5 \Theta 10 18 2 \Theta 10 \Gamma1
RCW 103 400 y 1.6 \Theta 10 20 3600 5.5 \Theta 10 17 9 \Theta 10 \Gamma1
Orion Bar 450 1.3 \Theta 10 21 4500 9.1 \Theta 10 16 2 \Theta 10 \Gamma3
NGC 2023 570 1.4 \Theta 10 20 6000 5.7 \Theta 10 16 4 \Theta 10 \Gamma3
S140 480 1.9 \Theta 10 20 8700 1.4 \Theta 10 17 2 \Theta 10 \Gamma2
y Assumed temperature.
For the shocks, a direct comparison between theory and
observation is more cumbersome. We have decided to plot the
results of the T rot fits from the J� and C�shock models versus the
total flux observed in all lines (Fig. 3). Since [O I] 63.2 �m, [C II]
157.7 �m and [Si II] 34.8 �m are expected to be the dominant
coolants in the lines we study here, we only plot this quantity in
sources where these three lines were measured. It is also given
in Table 6. For shocks with velocities in excess of 20 km s \Gamma1 ,
Lyff is also an important coolant, which we have neglected here.
The total luminosity of the cooling in a shock is proportional
to nv 3
s , so for the shocks plotted in Fig. 3 the ordinate will no
longer directly correspond to a single parameter in the models,
and instead of the simple lines for the PDRs, we get a grid of
models with different shock parameters. We are also left with
the additional uncertainty that the shock will likely not fill the
entire beam, so our data points must be shifted along a vector
indicated in the figures to take this into account.
Comparing the excitation temperature of the template
sources listed in Table 5 with the ones plotted in Fig. 2, we
note that for the relatively clean PDRs S140 and the Orion
Bar, the derived excitation temperature falls within the range
predicted by PDR models. Although the Tielens & Hollen�
bach (1985) models do not predict a detailed ro�vibrational H 2
spectrum, competing PDR models (e.g. Black & van Dishoeck
1987; Draine & Bertoldi 1996) also show the same hot compo�
nent which we observe in these PDR template sources. This hot
PDR may reflect the combined effects of UV�pumped infrared
fluorescence and the presence of a very warm, but thin, surface
layer.

10 M.E. van den Ancker et al.: ISO Spectroscopy of PDRs and Shocks in Star Forming Regions
The only shock template source for which we were able
to derive a reliable T rot , Orion BN/KL peak 1, agrees well
with the value expected for a J�shock. Note however that all
three of our template shock sources also show evidence for a
hot H 2 component, which is not predicted by J�shock models.
This probably means that even in these relatively clean template
cases, we do not observe a single J�shock, but a combination of
one or more J� and C�shocks. Possibly this reflects the presence
of shocked clumps of gas where the J�shock corresponds to the
region around the point of impact while the C�shock emission
originates from oblique shocks further away.
The situation for the programme sources is distinctly less
clear�cut. Except for GGD 27�ILL and IRAS 03260+3111,
which appear somewhat warmer in pure rotational H 2 than
predicted by the models, all sources which show PAH emission
and might therefore be expected to be dominated by PDR emis�
sion have excitation temperatures within the range predicted by
the PDR models. Small differences between observations and
the models may reflect the fact that these observations spatially
resolve the PDR while the models refer to integrated (i.e. aver�
aged) values. Of the remaining sources, only AFGL 2591 has
an excitation temperature compatible with the J�shock models.
All are compatible with excitation temperatures expected for
C�shocks. However, in the next section we will show that the
observed atomic fine�structure spectrum in these sources does
require the presence of a J�shock. This could either mean that
the H 2 has a stronger C�shock component than the template
sources, or that the J�shock models of Hollenbach & McKee
(1989) underestimate the temperature in the post�shock gas,
where the H 2 emission originates in a J�shock. We will come
back to this in the discussion and argue that the latter explana�
tion is more likely.
5. Fine structure lines
In order to interpret our observations as arising in either a PDR,
a C�Shock or a J�shock, we compare our results with the line
fluxes predicted by theoretical models of such regions (Tielens
& Hollenbach 1985; Hollenbach & McKee 1989; Kaufman &
Neufeld 1996). Since the Infrared Space Observatory offers the
first possibility to confront the infrared spectrum predicted by
these models with a wide range of astronomical objects, and the
physics behind these models is still not completely understood,
there are significant uncertainties in such a procedure. There�
fore we also again compare our results with observations from
literature of the template sources containing shocks and PDRs.
In Figs. 4--6 we have plotted the absolute line intensities
predicted either against the G�values computed from the in�
frared luminosity (PDRs) or against the sum of all observed
lines fluxes (shocks), computed in the previous section. Note
that since the quantities plotted on both axes of these plots de�
pend on the beam filling factor, the points may be expected to
shift along the arrows shown in the plots in case the PDR or
shock is smaller than the beam size.
By comparing the predicted line fluxes for the PDR, J�shock
and C�shock models with our observations in Figs. 4--6 we
can directly draw several interesting conclusions. The predicted
intensity in the [Fe I] and [S I] lines is orders of magnitudes
higher in the shock models than in the PDR models, because
in the latter nearly all iron and sulfur will be ionized in the
heated surface layer of a PDR. [Fe I] was not detected in any
source within our sample. However, [S I] 25.3 �m emission was
detected for seven sources. Comparing the location of these line
fluxes in Fig. 4 with the PDR model prediction we see that the
observed [S I] fluxes are too high to be compatible with a PDR
origin, even for PDRs as dense as 10 6 cm \Gamma3 . This means that
the detection of [S I] emission with ISO can only mean that a
shock is present. Also note that the detection of [S I] and PAHs
appear mutually exclusive in our sample (Table 6), suggesting
that these may be a fast and relatively straightforward way to
distinguish PDRs and shocks.
By definition, C�shocks cannot produce ionized species like
[Fe II], [Si II] and [C II]. The last species may still be detected in
such regions because it will also reach observable strengths in
PDRs illuminated by weak FUV fields (c.f. Fig. 4), such as the
diffuse galactic background emission (G � 1). Since the [C II]
line fluxes listed in Table 3 are not corrected for background
emission, we cannot draw conclusions from just the detection
of [C II]. However, the mere detection of [Si II] or [Fe II] means
that either a PDR or a J�shock must be present within the ISO�
SWS beam. For most sources in which we have detected [Si II]
or [Fe II] emission, it can be well explained by either the PDR
or the J�shock models (Figs. 4--5). In only one case, S106 IRS4,
is the observed [Si II] 34.8 � line intensity clearly too strong to
be explained by either PDR or J�shock models. However, van
den Ancker et al. (2000b) found that the infrared spectrum of
this source is dominated by its associated H II region, which
will also contribute to the [Si II] and [Fe II] lines studied here.
Although both the PDR and the J�shock models predict the
presence of observable amounts of [C II] emission, the expected
line strength in PDRs of moderate to high density (� 10 3 cm \Gamma3 )
is more than an order of magnitude higher than in J�shocks.
Most of the sources with observed [C II] 158 �m emission have
strengths that are too high to be explained by the J�shock mod�
els. For these sources we can therefore immediately infer that
they must be due to a PDR. For the remainder of the sources,
the [C II] 158 �m line strength becomes comparable to the one
expected for the background, so a direct comparison of those
values to J�shock models will yield inconclusive results.
Intense [O I] 63.2 �m emission can be producedby PDRs, as
well as J� and C�shocks. The observed [O I] 63.2 �m intensities
are all within the range that can be reproduced by the models.
The [O I] 145.5 �m line can also be produced by the PDR and
C�shock models; the J�shock models predict less [O I] 146 �m
emission than is observed. Also, in nearly all the cases we study
here, the PDR or C�shock model parameters found from the [O I]
63.2 and 145.5 �m lines do not agree with each other; either
the [O I] 63.2 �m line appears too weak, or the [O I] 145.5 �m
line appears too strong. Such a discrepancy is not unique to the
sources we study here. It has also been observed in spatially
resolved PDRs (Liseau et al. 1999), so it cannot be caused by
the different LWS beam sizes at 63 and 146 microns. One pos�

M.E. van den Ancker et al.: ISO Spectroscopy of PDRs and Shocks in Star Forming Regions 11
Fig. 4. Comparison of observed fine�structure line intensities and continuum flux to the PDR models. Sources which show PAH emission are
plotted as squares. The large arrows show the direction of beam dilution.

12 M.E. van den Ancker et al.: ISO Spectroscopy of PDRs and Shocks in Star Forming Regions
Fig. 5. Comparison of observed fine�structure line intensities to the J�shock models. Sources which show PAH emission are again plotted as
squares. The large arrows show the direction of beam dilution.

M.E. van den Ancker et al.: ISO Spectroscopy of PDRs and Shocks in Star Forming Regions 13
Fig. 6. Comparison of observed fine�structure line intensities to the
C�shock models. Sources with PAH emission are plotted as squares.
The large arrows show the direction of beam dilution.
sible explanation would be that [O I] 145.5 �m is blended with
the CO J=18--17 line. In view of the strength of other CO lines
in our LWS spectra, this line is in general expected to be too
weak to dominate the line flux, and thus cannot account for the
observed anomaly. Another, more promising, explanation may
be that the [O I] 63.2 �m line intensity has been diminished by
absorption by cool foreground material. [O I] 63.2 �m absorp�
tion has indeed been found in the line of sight towards a few
sources (Poglitsch et al. 1996; Kraemer et al. 1998), lending
credence to the suggestion that this may also be the case in our
sources. Note that if this suggestion is correct, this means that
the observed [O I] emission in all our sources must come from
either PDRs or C�shocks, since the [O I] 145.5 �m line fluxes
are too high compared to the J�shock models.
Physical parameters for each star's circumstellar environ�
ment, estimated from the different measured species, as well as
the likely underlying physical mechanism, are listed in Table 6.
6. Notes on individual objects
AFGL 490 (IRAS 03236+5836) is a luminous infrared source,
powering a strong bipolar CO outflow with an extent of �
100 00 (Lada & Harvey 1981). The central peak in AFGL 490
is believed to be an embedded intermediate�mass YSO and the
driving source for the molecular outflow (Minchin et al. 1991;
Haas et al. 1992). Ground�based H 2 1--0 S(1) images (Davis et
al. 1998) show several knots of shocked molecular gas in the
AFGL 490 outflow as well as diffuse H 2 emission, suggesting a
composite origin for the emission lines detected with ISO. The
temperatures derived from the H 2 0--0 S(3) and S(5) lines can
only be explained by a non�dissociative shock. However, the
detection of strong [Si II] and [O I] emission show that either a
PDR or a fast J�shock must be present as well.
IRAS 03260+3111 was found to be a candidate for a class
of objects making the transition between Class I and II YSOs
by Magnier et al. (1999). Strong PAH�emission in the SWS
spectrum suggests that it has a significant PDR contribution.
The lines detected confirm this hypothesis: The [Si II], [O I] and
[C II] are strongly suggestive of a PDR origin. The temperature
of the H 2 is somewhat higher than that expected for a PDR, but
this may be influenced by an apparent discontinuity between the
0--0 S(2) and S(3) lines, where SWS uses apertures of different
sizes.
L1489 IRS (IRAS 04016+2610) is a Class I source in an
extended molecular gas core in the Taurus star forming re�
gion. It powers a low�velocity molecular outflow adjacent to
the source (Myers et al. 1988; Hogerheijde et al. 1998) and
has been suggested as the powering source of a more extended
molecular outflow (Terebey et al. 1989; Moriarty�Schieven et
al. 1992). Recent HST NICMOS images of the source reveal a
bright near�infrared nebula crossed by a dark lane, interpreted as
an optically thick disk seen in silhouette (Padgett et al. 1999).
Ground�based images in the H 2 1--0 S(1) line reveal strong
H 2 emission from the HH 360a knot in the molecular outflow
(G� omez et al. 1997). No emission lines were detected with ISO,
suggesting that the shocked material in the region is confined
to these small knots.
L1551 IRS5 (IRAS 04287+1801) is one of the best stud�
ied embedded YSOs. The bipolar molecular outflow emanating
from L1551 IRS5 was the first to be discovered (Snell et al.
1980), and it remains a text�book example of a Class I source.
It illuminates the reflection nebula HH 102, and drives both a
strong molecular outflow as well as a highly collimated jet, as
traced by a number of Herbig�Haro objects with high proper
motions (Mundt et al. 1991; Staude & Els� asser 1993). Interfer�
ometry of L1551 IRS5 at mm wavelengths (Keene & Masson
1990) reveals the source to consist of two distinct components:
an envelope with radius � 12 00 and a compact structure of �

14 M.E. van den Ancker et al.: ISO Spectroscopy of PDRs and Shocks in Star Forming Regions
Table 6. Results of line analysis.
Object FIR Gav: PAH \Sigma I n [cm \Gamma3 ] Type
[W m \Gamma2 ] [G0 ] [erg cm \Gamma2 s \Gamma1 sr \Gamma1 ] H2 [Fe I] [Fe II] [S I] [Si II] [O I] [C II]
AFGL 490 1.4 \Theta 10 \Gamma10 5.5 \Theta 10 3 -- 4.1 \Theta 10 \Gamma4 -- -- -- -- ? � 10 3 ? � 10 3 -- Shock?
IRAS03260+3111 7.4 \Theta 10 \Gamma11 2.9 \Theta 10 3 p
2.4 \Theta 10 \Gamma3 ? � 10 6 -- -- -- ? � 10 5 ? 10 3 ? 10 2 PDR
L1489 IRS 6.2 \Theta 10 \Gamma12 2.4 \Theta 10 2 -- -- -- -- -- -- -- --
L1551 IRS5 3.9 \Theta 10 \Gamma11 1.5 \Theta 10 3 -- 1.7 \Theta 10 \Gamma4 -- -- ? � 10 4 -- � 10 3 ? � 10 3 -- Shock?
IRAS12496\Gamma7650 4.3 \Theta 10 \Gamma11 1.7 \Theta 10 3 -- 1.4 \Theta 10 \Gamma4 -- -- � 10 2 -- � 10 3 ? � 10 3 -- Shock?
GGD 27�ILL 2.2 \Theta 10 \Gamma10 8.3 \Theta 10 3 p
1.5 \Theta 10 \Gamma3 ? � 10 6 -- -- -- � 10 4 ? � 10 3 ? 10 2 PDR
LkHff 225 5.3 \Theta 10 \Gamma11 2.0 \Theta 10 3 -- 3.4 \Theta 10 \Gamma3 � 10 4 -- ? � 10 4 � 10 4 � 10 4 ? � 10 4 -- Shock
S106 IRS4 9.1 \Theta 10 \Gamma10 3.5 \Theta 10 4 p
2.2 \Theta 10 \Gamma2 � 10 6 -- ? 10 5 -- -- � 10 4 ? 10 3 PDR
AFGL 2591 6.7 \Theta 10 \Gamma10 2.6 \Theta 10 4 -- 1.9 \Theta 10 \Gamma3 � 10 4 -- -- -- ? 10 3 � 10 3 ? � 10 3 PDR?
Cep A East 3.0 \Theta 10 \Gamma10 1.1 \Theta 10 4 -- 6.9 \Theta 10 \Gamma3 ? � 10 4 -- � 10 4 ? � 10 4 � 10 4 ? � 10 4 -- Shock
Elias 3�1 3.3 \Theta 10 \Gamma11 1.3 \Theta 10 3 p
-- -- -- -- -- -- --
T Tau 2.2 \Theta 10 \Gamma11 8.5 \Theta 10 2 -- 7.7 \Theta 10 \Gamma4 ? � 10 4 -- � 10 3 � 10 3 � 10 3 � 10 3 -- Shock y
HD 97048 1.4 \Theta 10 \Gamma11 5.5 \Theta 10 2 p
4.0 \Theta 10 \Gamma4 -- -- -- -- -- � 10 2 -- PDR
HD 97300 1.1 \Theta 10 \Gamma11 4.3 \Theta 10 2 p
-- -- -- -- -- -- --
HR 5999 3.0 \Theta 10 \Gamma11 1.2 \Theta 10 3 -- -- -- -- -- -- -- --
R CrA 1.8 \Theta 10 \Gamma10 6.9 \Theta 10 3 -- 5.7 \Theta 10 \Gamma4 -- -- ? � 10 3 � 10 3 � 10 3 ? � 10 3 -- Shock y
T CrA 1.4 \Theta 10 \Gamma11 5.3 \Theta 10 2 -- -- -- -- -- -- -- --
WW Vul 2.5 \Theta 10 \Gamma12 9.4 \Theta 10 1 -- -- -- -- -- -- -- --
BD+40 ffi 4124 9.0 \Theta 10 \Gamma12 3.4 \Theta 10 2 p
1.1 \Theta 10 \Gamma3 ? � 10 4 -- � 10 5 -- -- ? � 10 4 � 10 3 PDR
LkHff 224 8.9 \Theta 10 \Gamma12 3.4 \Theta 10 2 -- -- -- -- -- -- -- Shock y
HD 200775 6.7 \Theta 10 \Gamma11 2.6 \Theta 10 3 -- 1.2 \Theta 10 \Gamma3 -- -- -- -- ? � 10 4 ? � 10 3 ? � 10 3 PDR
y Probably not connected to optical source.
0: 00 3, presumably an accretion disk of a similar size as the solar
system. Diffuse H 2 1--0 S(1) emission in the L1551 region was
observed by Yamashita & Tamura (1992), and later shown to
have a patchy structure by Davis et al. (1995), which may be
due to gas which is shock�heated by the strong wind from IRS5.
No compact H 2 knots similar to those observed in the surround�
ings of other YSOs were found in the large�scale surroundings
of L1551 IRS5 (G� omez et al. 1997), suggesting the presence
of shocks sufficiently powerful to destroy the dust required to
re�form molecular hydrogen in post�shock gas. [Fe II], [Si II]
and [O I] emission was detected by ISO. In view of the previous
discussion it seems likely that this will be shock excited.
IRAS 12496\Gamma7650 (DK Cha) is an embedded active young
stellar object located in the Chamaeleon II dark cloud. It is asso�
ciated with a weak molecular outflow (Knee 1992) and possibly
drives a collimated jet (Hughes et al. 1991). It is associated with
the Herbig�Haro objects HH 52--54, some of which are also in�
cluded in the ISO beam. If the H 2 ortho�/para�ratio in these
objects is smaller than the high�temperature equilibrium value
of 3, as suggested by ISO spectroscopy of HH 54 (Neufeld et
al. 1998), our estimate of the H 2 mass (Table 5) will be too low.
The ISO [O I]/[C II] ratio suggests a strong shock component
to be present in the infrared emission�line spectrum. In view of
the detected [Fe II] and [Si II] emission, this shock must be of
J�type.
GGD27 (IRAS 18162\Gamma2048) is a small reflection nebula
containing six point�like infrared sources, interpreted as a small
group of massive YSOs (Yamashita et al. 1987; Stecklum et
al. 1997). The object GGD 27�ILL (Aspin et al. 1991) is the
probable illuminating source of the nebula and might also power
the bipolar outflow seen in CO (Yamashita et al. 1995). The
emission lines observed with ISO appear dominated by PDR
emission rather than trace a shock driven by the outflow. The
ISO emission�line spectrum is also indicative of a PDR origin,
although the H 2 temperature falls outside the range predicted
by the PDR models.
LkHff 225 (IRAS 20187+4111) is an embedded source
in the BD+40 ffi 4124 region. The source was shown to be a
triple system oriented north�south by Aspin et al. (1994). It is
associated with a H 2 O maser source and drives a CO outflow
(Palla et al. 1994). The detection of [S I] emission shows that
a shock must be present in the region. Van den Ancker et al.
(2000a) studied the ISO spectra of LkHff 225 in more detail
and concluded that the emission line spectrum can be well
explained as arising in the combination of a C�shock produced
by a slow (� 20 km s \Gamma1 ) outflow and a J�shock. In addition to
this, we found absorption due to solid water ice and to gas�phase
H 2 O, CO and CO 2 in the line of sight towards LkHff 225, with
unusually high gas/solid ratios. The extreme [O I]/[C II] ratio
found at the position of LkHff 225 suggest that a PDR might
be present throughout the region as well.
S106 IRS4 (IRAS 20255+3712) is a massive young stellar
object believed to be the powering source for the well�known
bipolar nebula S106. Apart from the central source IRS4 (Gehrz

M.E. van den Ancker et al.: ISO Spectroscopy of PDRs and Shocks in Star Forming Regions 15
et al. 1982), S106 also contains an embedded cluster of about
160 stars, as well as an expanding ring of molecular material
(Hodapp & Rayner 1991; Loushin et al. 1990; Staude &Els� asser
1993). Van den Ancker et al. (2000b) made a more detailed study
of the ISO SWS and LWS spectra of S106 and concluded that
the emission line spectrum could be well explained as arising
from the superposition of an extended H II region with a density
of about 10 3 cm \Gamma3 surrounding an O6--8 star, and a PDR of
much higher density (10 5 --10 6 cm \Gamma3 ). Since it is implausible
that such densities would exist throughout the PDR, the PDR
must be clumpy in nature.
AFGL 2591 (IRAS 20275+4001) is one of the rare ex�
amples of a relatively isolated massive young star. Located in
the obscured Cygnus X region, it is invisible in the optical,
although it is one of the brightest YSOs at infrared to submm
wavelengths. It is associated with a powerful bipolar molecular
outflow (Lada et al. 1984) and is surrounded by an envelope of
which part might be in free fall collapse onto the star (van der
Tak et al. 1999). Tamura &Yamashita (1992) discovered a bipo�
lar outflow�like structure extending over 90 00 in ro�vibrational
H 2 emission, suggesting a shock origin for the spectral lines
detected with ISO. This would also explain the absence of PAH
emission. The [O I]/[C II] flux ratio is more suggestive of a PDR
origin, though.
Cep A East (IRAS 22543+6145) is the eastern lobe of an
energetic, complex molecular outflow (Bally & Lane 1982),
believed to be powered by the deeply embedded source HW�2
only visible at radio wavelengths (Hughes &Wouterloot 1984).
Cep A East is the site of a fast jet, as well as a number of
Herbig�Haro like objects, whereas the western lobe of the Cep
A outflow contains a number of distinct shells (Hartigan et al.
1996; Goetz et al. 1998). ISO SWS observations of the western
lobe were discussed by Wright et al. (1996), who modelled the
emission line spectrum as the combination of several C�type
shocks with a planar J�shock of 70--80 km s \Gamma1 . Van den Ancker
et al. (2000b) discussed the entire SWS and LWS spectra of Cep
A East in detail and concluded that the continuum radiation at
near�infrared wavelengths is dominated by emission from the
embedded source IRS 6A (Casement &McLean 1996), whereas
another, more luminous, componentonly becomes visible in the
far�infrared. We explain the infrared emission line spectrum as
the superposition of a 20 km s \Gamma1 C�shock and a 60 km s \Gamma1
J�shock arising in a dense (10 6 cm \Gamma3 ) medium.
Elias 3�1 (IRAS 04155+2812) is an embedded luminous
object which is believed to be a young Herbig Ae star. It il�
luminates a cometary�shaped reflection nebula, in which an
east�west oriented elongated structure is visible at near�infrared
wavelengths, interpreted by Haas et al. (1997) as two lobes of
a cone�like structure. One feature that distinguishes Elias 3�1
is that it is one of the only three sources known that shows the
3.43 and 3.53 �m infrared emission bands, as well as the other
infrared emission features usually attributed to PAHs (Schutte
et al. 1990). No infrared emission lines were detected.
T Tau (IRAS 04190+1924) is a young binary system, con�
sisting of an optically visible K0--1e classical T Tauri star with
an embedded (A V = 17: m 4), more luminous, companion. Both
stars are surrounded by a compact nebula as well as spatially
separated arc�shaped cloud 30 00 to the west. The ISO SWS
observations presented here only include the compact nebula,
whereas the LWS observations also include the more extended
cloud. Both the optically visible and embedded components of
the T Tau system drive separate molecular outflows, whereas
the embedded source might also power a string of Herbig�Haro
objects as far away as 38 arcminutes (1.55 pc) from the source
(Edwards & Snell 1982; van Langevelde et al. 1994; Reipurth
et al. 1997). Van den Ancker et al. (1999) and Spinoglio et al.
(2000) studied the complete ISO SWS and LWS spectra of T
Tau in more detail and came to the conclusion that the observed
molecular and atomic fine structure emission line spectrum can
be well explained as arising in the superposition of a C�shock
with a speed of � 35 km s \Gamma1 and a much faster (� 100 km s \Gamma1 )
J�shock arising in a fairly dense (5 \Theta 10 4 cm \Gamma3 ) medium.
HD 97048 (IRAS 11066\Gamma7722) is a fairly typical late B�
type Herbig Ae/Be star located in the Chamaeleon I dark cloud.
Wesselius et al. (1984) discovered the source to have a strong
infrared excess, which was interpreted in terms of a flattened
shell or disk of dust particles heated by the central star by
Th� e et al. (1986). One feature that distinguishes HD 97048 is
that it is one of the only three sources known that shows the
3.43 and 3.53 �m infrared emission bands, as well as the other
infrared emission features usually attributed to PAHs (Schutte
et al. 1990). Although extended mid�infrared emission around
HD 97048 has been reported from multi�aperture photometry
(Prusti et al. 1994), unpublished N band imaging by the authors
with the TIMMI instrument at the ESO 3.6 m telescope show
the source to be point�like (diameter ! 1 00 ), suggesting that the
extended emission might be limited to the PAH bands. Only
weak [O I] and [C II] emission was detected. If we assume that
this emission comes from the same region as the continuum
emission, it is compatible with an origin in the surface layer of
a circumstellar disk surrounding HD 97048, acting as a PDR.
HD 97300 (IRAS 11082\Gamma7620) is a B9 star located in the
Cha I complex. The absence of Hff emission and infrared excess
at wavelengths shortwards of 5 �m suggests that it is a relatively
evolved object within the group of Herbig Ae/Be stars (Th� e et
al. 1986). HD 97300 is located near (or is seen projected on) the
reflection nebula Ced 112, in which Siebenmorgen et al. (1998)
discovered a ring of PAHs with ISOCAM. No emission lines
were detected with SWS.
HR 5999 (IRAS 16052\Gamma3858) is one of the best studied
Herbig Ae stars. In the visual, it shows irregular large�amplitude
photometric variations due to variable amounts of circumstel�
lar extinction (Tjin A Djie et al. 1989; Th� e et al. 1996). HR
5999 is located in the pre�main sequence instability strip, and
also shows regular pulsations similar to those observed in the
evolved ffi Scuti stars (Kurtz & Marang 1995; Marconi & Palla
1998). In the infrared, HR 5999 is a fairly isolated point source
(Siebenmorgen et al. 1997), suggesting the star has cleared most
of its wide stellar environment. This is in agreement with the
non�detection of infrared emission lines with ISO.
R CrA is a well studied Herbig Ae star at the apex of a
cometary nebula in the Corona Australis cloud. In the visual, it

16 M.E. van den Ancker et al.: ISO Spectroscopy of PDRs and Shocks in Star Forming Regions
shows strong photometric variations, while its visual absorption
line spectrum is also variable, suggesting it may in fact be
a composite system (Bibo et al. 1992; van den Ancker et al.
1998). Several strong molecular outflows containing knots of
shocked gas have been detected in the vicinity of R CrA, but
recent high�resolution continuum and molecular mappings of
the region have shown that the driving sources of these are the
embedded sources IRS7 and HH100�IR rather than the optical
Herbig stars in the region (Harju et al. 1993; Anderson et al.
1997). The detection of [S I] shows that a shock must be present
within the SWS beam. Most likely, this is due to the outflow of
IRS7.
T CrA is a late�type Herbig star located in the tail of a
cometary nebula in the Corona Australis cloud. It shows strong
photometric variations in the visual, probably due to variable
amounts of circumstellar extinction (Bibo et al. 1992). Based on
optical spectro�astrometric measurements, Bailey (1998) sug�
gested T CrA to be a binary system. Ward�Thompson et al.
(1987) reported the presence of a jet associated with T CrA, but
more recent papers suggest that this may in fact be part of the
east�west molecular outflow from the source IRS7 in the region
around R CrA (Canto et al. 1986; Harju et al. 1993; Anderson
et al. 1997). The H 2 emission detected by ISO is suggestive of
a C�shock which may be attributed to this same outflow.
WW Vul (IRAS 19238+2106) is a prototypical Herbig Ae
star showing large amplitude photometric variations due to vari�
able circumstellar extinction (Friedemann et al. 1993), as well
as evidence for the presence of evaporating solid bodies (Grinin
et al. 1996). It shows a large infrared excess, indicative of the
presence of a dusty circumstellar disk or envelope. WW Vul
is located in relative isolation, which is confirmed by our non�
detection of infrared emission lines.
BD+40 ffi 4124 (V1685 Cyg) is the most massive member
of a small cluster of young stars, commonly known as the
BD+40 ffi 4124 group (Hillenbrand et al. 1992). It is the ionizing
source of a low�density (� 10 2 cm \Gamma3 ) H II region, surrounded
by a dense (? 10 5 cm \Gamma3 ) PDR (van den Ancker et al. 2000a).
The lines detected with ISO are indicative of PDR emission.
In view of the luminosity of BD+40 ffi 4124 and the strong [C II]
emission detected at the position of LkHff 225, it seems likely
that the PDR is more extended than the ISO beam and may
in fact produce observable PDR emission lines throughout the
optical reflection nebula.
LkHff 224 (V1686 Cyg) is a fairly typical Herbig Ae star
located in the BD+40 ffi 4124 region. In the optical it shows large
variations in brightness due to variable amounts of circumstellar
extinction (Wenzel 1980; Shevchenko et al. 1991). One remark�
able aspect about LkHff 224 is the complete absence of a 10 �m
silicate feature, either in absorption or emission in its spectrum
(van den Ancker et al. 2000a). These authors also studied the
infrared emission line spectrum of LkHff 224 and concluded
that it is most likely due to non�dissociative shock produced by
a slow (� 20 km s \Gamma1 ) outflow arising from LkHff 225.
HD 200775 is the illuminating star of the well�known re�
flection nebula NGC 7023 (IRAS 20599+6755). A biconical
cavity of � 20 00 diameter surrounds the star, outside of which
the nebulosity shows a highly filamentary structure (Rogers et
al. 1995; Lemaire et al. 1996; Gerin et al. 1998). Observations
with ISOCAM and ISOPHOT revealed the nebulosity to show
strong emission bands due to Polycyclic Aromatic Hydrocar�
bons (Cesarsky et al. 1996; Laureijs et al. 1996). The SWS
observations shortward of 12.0 �m only include the central star
and the cavity, whereas the SWS data at longer wavelengths
as well as the LWS spectrum also cover part of the PDR. The
[O I]/[C II] ratio clearly reflects this PDR nature. Fuente et al.
(2000) have analyzed ISO SWS observations centered on other
parts of NGC 7023 and concluded that they arise in a PDR with
temperatures in the range 300--700 K, and with an ortho/para
ratio of H 2 of 1.5--2. Our observations suggest the PDR to have
a density of � 10 5 cm \Gamma3 , much higher than the modelling re�
sults of the NGC 7023 PDR by Chokshi et al. (1988). Most
likely the lines we have detected are dominated by a few of the
high density clumps reported by Martini et al. (1997).
7. Discussion and conclusions
In the previous sections we have seen that based on the observed
H 2 and atomic fine�structure lines we can make a distinction
between PDRs and shocks with relative ease, and derive the
physical properties in the emitting region, provided that key
lines like [S I], [O I] and [C II] and the low�lying pure rotational
molecular hydrogen lines were detected. The mere presence
of [S I] is a sure way to recognize a shock, whereas the [O I]
63.2 �m/[C II] 157.7 �m ratio is able to distinguish PDRs of
average to high density from shocks. From Table 6 we note
that both shocks and PDRs are observed near Class I YSOs,
whereas Class II YSOs only produce PDR emission. Note that
the three Class II sources in whose vicinity we have detected
[S I] emission, T Tau, LkHff 224 and R CrA, are all located
near Class I sources driving a strong bipolar outflow. We asso�
ciate the observed [S I] emission in these three sources with the
embedded instead of the optical star.
An interesting question which we have not yet addressed in
the previous sections is whether the heating efficiency assumed
in the theoretical models is confirmed by the observations pre�
sented here. This efficiency will determine the absolute line
fluxes, and could also influence the detailed structure of PDRs
and shocks. For PDRs, the model parameter which determines
the efficiency is the ratio of gas heating to the far�ultraviolet ab�
sorption rate of dust grains and PAHs, ''. On theoretical grounds,
it is expected that the maximum allowed value of '' would be
3--5% (Bakes & Tielens 1994). This efficiency decreases when
the PAHs are charged up and hence decreases with increasing G
(ionization rate) and decreasing n e (recombination rate; Bakes
& Tielens 1998). Since for most sources in our sample the lines
which we have detected include the major PDR coolants [O I]
63.2 �m, [C II] 158 �m and [Si II] 34.8 �mwe are able to directly
derive a value of '' for our PDR sources. For this reason, we
have plotted the summed intensity in all detected lines against
the values of G derived from the infrared continuum flux de�
rived listed in Table 6. From this plot (Fig. 7), in which we also
show lines of constant '', it can be seen that most of our PDR

M.E. van den Ancker et al.: ISO Spectroscopy of PDRs and Shocks in Star Forming Regions 17
Fig. 7. Comparison of the sum the fluxes of all observed lines with G
values as derived from the infrared continuum flux. Plots symbols have
the same meaning as in Fig. 3. Lower limits are shown for objects in
which one or two of the major coolants ([O I] 63 �m, [C II] 158 �m,
[Si II] 35 �m) were not observed. Also shown are lines of equal '' (ratio
of gas heating to far�ultraviolet absorption rate of grains and PAHs).
Fig. 8. Comparison of flux emitted in [S I] 25.3 �m (top) and 6.2 �m
PAH band (bottom) with estimated visual extinction towards the asso�
ciated source. Plot symbols have the same meaning as in Figs. 3--6.
sources (as identified from the associated PAH emission) have ''
of a few times 10 \Gamma3 . This estimate is nearly independent of the
beam filling factor. Although somewhat lower than assumed in
most PDR models, this value agrees well with those observed
towards PDRs in different astrophysical environments (see e.g.
Hollenbach & Tielens 1999).
As a final note, let us now wonder how to put our main
result, that both shocks and PDRs are observed near Class I
YSOs, whereas Class II YSOs only show PDR emission, in
context. In the previous sections we have seen that the detec�
tion of [S I], indicative of shocks, and PAH�emission, indicative
of PDR activity, appear mutually exclusive. To illustrate this
mutual exclusion of [S I] and the PAH emission and to illustrate
its dependence on the evolutionary state of the YSO, we have
created plots of the luminosity in the [S I] 25.3 �m line and in
the 6.2 �m PAH band, normalized to the bolometric luminosity,
against the visual extinction A V for the sources in our sample
(Fig. 8). We can see here that the Class I sources for which we
have detected [S I] emission, LkHff 225 and Cep A East, are
also among the sources with higher values of A V . However, ap�
parently the mere embeddedness of a source does not guarantee
that strong [S I] is present, as evidenced by some of the upper
limits plotted in Fig. 8.
PAH emission was observed in both Class I (IRAS
03260+3111, GGD 27�ILL, S106 IRS4) and Class II
(BD+40 ffi 4124, HD 97048) sources. The range in fraction of the
total luminosity coming out in the 6.2 �m C--C stretch mode
does not appear different between the two groups, although the
cooler stars in both groups appear to have a higher LPAH=L bol
than the hotter stars. The values of A V used in Fig. 8 are those
for the associated continuum source; the lines do not necessar�
ily have to suffer from the same amount of extinction. Because
a young star is expected to clear its circumstellar environment
as it accretes mass, these values of A V may be taken as a
rough indication of evolutionary status. Geometry is however
also important in interpreting these results as exemplified by
S106 IRS4 where our direct view of the YSO is inhibited by an
optically thick circumstellar disk but the gas along the poles is
photoionized by the star (van den Ancker et al. 2000b).
It is interesting that the most embedded source in our sam�
ple, Cep A East, also has the highest [S I] luminosity, and that
all Class I sources which show PAH emission are only moder�
ately embedded. This is exactly what one would also expect:
as a star has just formed, it is heavily embedded and any PAH
emission arising in a PDR close to the star will be completely
obscured. As long as the young star accretes matter it will also
drive a bipolar outflow, which can power a shock as it inter�
acts with the ambient medium. As this process proceeds, the
surroundings of the star are cleared and the UV radiation from
the newly formed star can escape to illuminate a PDR. The in�
tensity of the PDR emission is mainly determined by its spatial
extent. As the material in the wide circumstellar environment
is slowly cleared by the still strong stellar wind and the effect
of photoevaporation, the extent of the PDR and hence the in�
tensity of the PDR lines will gradually decrease until the only,
weak, source of PDR emission will become the surface of a cir�
cumstellar disk, acting as a PDR (e.g. the case of HD 97048).
With the tools presented in this paper, the study of such PDRs
offers the promise of a fundamental better understanding of the
processes occurring in protoplanetary disks when far�infrared
spectrographs with higher spatial resolution become available.
Acknowledgements. The authors would like to thank the SWS and
LWS IDTs for their help with the observations. Drs. David Hollen�
bach and Michael Kaufman are gratefully acknowledged for shocking

18 M.E. van den Ancker et al.: ISO Spectroscopy of PDRs and Shocks in Star Forming Regions
discussions. MvdA acknowledges financial support from NWO grant
614.41.003 and through a NWO Pionier grant to L.B.F.M. Waters.
This research has made use of the Simbad data base, operated at CDS,
Strasbourg, France.
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