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\begin{center}
{\huge \bf The dense and warm interstellar medium}\vspace{0.5cm}\\
{\LARGE \bf A Herschel Science Key Project Proposal}
\end{center}
\vspace{2cm}

{\sl Draft, 03-29-2004\vspace{1.5cm}\\Contributors:}
V.~Ossenkopf$^{1,2}$, M.~Gerin$^3$, R.~G\"usten$^4$, 
F.~Boulanger$^5$, A.~Fuente$^6$, Ch.~Joblin$^7$, Th.Klein$^4$,
C.~Kramer$^1$, W.~Langer$^8$,
S.~Lord$^9$, J.~Martin-Pintado$^{10}$, D.~Neufeld$^{11}$, 
S.~Philipp$^4$, M.~Röllig$^1$, M.~Spaans$^{12}$, J.~Stutzki$^1$,
D.~Teyssier$^2$
\vspace{2cm}\\

\noindent
$^1$ KOSMA, I. Physikalisches Institut, K\"oln -- Germany\\
$^2$ SRON, National Institute for Space Research, Groningen -- Netherlands\\
$^3$ LERMA, Observatoire de Paris -- France\\
$^4$ MPIfR, Bonn -- Germany\\
$^5$ IAS, Universit\'e Paris -- France\\
$^6$ OAN, Madrid -- Spain\\
$^7$ CESR, Toulouse - France\\
$^8$ JPL, Caltech, Pasadena -- USA\\
$^9$ IPAC, Caltech, Pasadena -- USA\\
$^{10}$ CSIC, Madrid, Spain\\ 
$^{11}$ Johns Hopkins University, Baltimore - USA\\
$^{12}$ Kapteyn Astronomical Institute, Groningen - Netherlands\\

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\section{Introduction}

Because of the combination of its extended wavelength coverage,
and its high spectral resolution, HIFI is a unique instrument for
the exploration of the physical and chemical processes controlling
the state and evolution of the interstellar medium. The physical and
chemical  processes determine the gas composition, hence its ability
to cool radiatively. In addition, the gas thermodynamic properties influence
its dynamics. High resolution spectroscopy is a powerful
mean for determining simultaneously the gas composition and its
kinematics. Emission lines at far infrared and submillimeter wavelengths,
contribute efficiently to the gas cooling, especially in moderately
warm environments at a few tenths to a few hundredths of kelvins. 

Two basic mechanisms can heat dense molecular gas to
temperatures above 100~K: the heating by radiation, mainly
in the UV, emitted by massive stars, and the dynamic impact from
violent outflows from these stars. Whereas the role of outflows
produced during the formation of stars is studied in the
HIFI key project on star-formation, we focus here on shocks
generated in later stages of the evolution of massive stars.
The most prominent example of these shocks in the dense ISM
is produced by supernova explosions in the vicinity of molecular clouds.

However, in many cases the distinction between the role of shocks and
radiation is not clear because massive stars clusters produce both
effects. Also, a complex dynamics in
all regions where the UV radiation is the main heating source 
(Gorti \& Hollenbach 2002). Even in supernova blast waves, the role
of FUV radiation can not be neglected.

In this key project we propose to perform FIR spectroscopy of molecular clouds
heated by both effects covering a broad range of parameters so that
we can test models of photodissociation regions (PDRs) and models of
shock fronts and study the combined effect of both on the chemical
and physical structure of the warm and dense ISM.

\section{PDRs}

Photon-dominated regions (PDRs) are transition zones in the interstellar
medium between a region of ionized or atomic gas characterized by a low
optical depth for UV photons and a region of high-density,
typically molecular gas with a large optical depth for UV photons. 
In PDRs the Galactic radiation field or the radiation from nearby
stars completely determines the thermal and chemical 
structure of the interstellar gas. As such, PDRs are direct
manifestations of the energy balance of interstellar gas and their study
allows one to determine how the ISM survives the presence of the stars
that it forms (Tielens \& Hollenbach 1985;
van Dishoeck \& Black 1988; Le Bourlot et al.\ 1993; Wolfire et al.\ 1993;
Spaans et al.\ 1994; Sternberg \& Dalgarno 1995; St\"orzer et al.\ 1995;
Maloney et al.\ 1996; Lee et al.\ 1996; Spaans \& van Dishoeck 1997;
Kaufman et al.\ 1999; Le Petit et al.\
2002; and Tielens \& Hollenbach 1999 and references therein).

PDRs are found close to bright O and B stars as well as cooler A and F stars.
PDRs are also associated with planetary nebulae, active galactic nuclei and
photo-evaporating planetary disks around newly formed stars as well as
red giant outflows (AGB stars). In addition, the diffuse interstellar
medium in its neutral low density phase or when condensed into clouds
is exposed the average interstellar radiation so that it 
can be understood as a low density, low radiation field PDR.
Hence, PDRs cover a wide range of astrophysical environments and trace the
overall energy balance of the atomic and molecular ISM over ranges in density
and temperature of $n_{\rm H}=10^2-10^7$ cm$^{-3}$ and $T=10-10^4$~K,
respectively. The term PDR is usually reserved for impinging radiation fields
that peak in the FUV, while cool PDRs are associated with stars of effective
temperatures of $T_{\rm eff}=6,000-30,000$ K and XDRs (X-ray dominated
regions) are formed near strong sources of X-ray emission.

The strength of the impinging UV field is one main characteristic of a PDR.
It is defined by reference to the average radiation field in
the solar neighbourhood, far from massive stars. For Galactic sources,
the radiation field may vary between $10^5$ times the average interstellar
radiation field in the vicinity of massive star forming region and less than a
fraction of it at large galactocentric distances.
In principle, almost every boundary of the molecular clouds
in the Milky Way represents a PDR.  To trace the
influence of the strength of the irradiating flux, our sample does not
only include sites in the immediate vicinity of massive young stellar
clusters but also sites at larger distances, deeper in the shielded
clouds. It also includes a few examples of PDRs subject to low
radiation fields.

\begin{figure}
\includegraphics[width=14.2cm]{horsehead-co4-3.eps}
\caption{The Horsehead Nebula is a well known example of a PDR.
The UV radiation incident from $\sigma$ Ori (to the right hand side)
induces a complex chemistry on the surface of the molecular cloud and heats
the molecular gas as indicated by the warm CO emission (Philipp et al. 
2004, in preparation).}
\end{figure}

Interstellar gas irradiated by optical, UV or X-ray photons is heated,
through photo-electric emission from
dust grains (Bakes \& Tielens 1994), and (partially) ionized and dissociated.
PDRs are bright sources of infrared dust emission because the bulk
of the impinging FUV (6-13.6 eV) flux is absorbed by dust grains and
re-radiated at longer  wavelengths. Therefore, most of the non-stellar
infrared and (sub-)millimeter CO emission from galaxies is thought to
be due to  PDRs.

They have become increasingly important as diagnostic tools of astrophysical
environments with the advent of infrared and (sub-)millimeter telescopes
and space observatories. PDRs emit fine-structure lines of [C{\sc I}]
609, [C{\sc II}] 158 and [O{\sc I}] 63 $\mu$m; rotational lines of CO; ro-vibrational and
pure rotational  lines of H$_2$; many H$_2$O lines as well as many broad mid-IR features
associated with Polycyclic Aromatic Hydrocarbons (PAHs). In PDRs the bulk of
the atomic hydrogen is converted into H$_2$ and the carbon is transformed
from its mostly ionized form into CO.

Detailed models have been constructed of PDRs over the past decades and these
have allowed us to determine, with increasing confidence: the density and
temperature structure, elemental abundances (i.e., the metallicity), the
ionization degree of the ambient medium as well as the strength of the
impinging radiation field (see Tielens \& Hollenbach 1999 for an extensive
review). Various physical
processes that play a role in PDRs explain the observed correlation between
the [CII] 158 $\mu$m (ionized) and CO 1-0 (molecular) emission lines; the
link between CO 1-0 emission and the total molecular (H$_2$) mass, i.e., the X
factor; and also the scaling of the sum of the [OI] 63 $\mu$m and the [CII]
158 $\mu$m emission (which is almost equal to the total PDR cooling rate) and
the infrared continuum luminosity of dust grains.

% On a galactic scale PDRs lead to the existence of two stable phases in the ISM,
% linked to specific temperature thresholds of the atomic coolants, that are
% roughly in pressure equilibrium. The warm (mostly) neutral medium  with a density
% of $n_{\rm H}\sim 1-10$ cm$^{-3}$ and a temperature of
% $T\sim 10^4-10^3$ K and the cold neutral medium with $n_{\rm H}\sim 100$
% cm$^{-3}$ and $T\sim 100$ K (Field et al.\ 1969).
% The hot ($106$ K) interiors of supernova
% explosions form a third phase while the cold ($10$ K) molecular
% gas out of which stars are ultimately born is formed under the influence of
% gravitational collapse of dense molecular clouds.

Finally, PDRs are regions where molecular material, necessary for the
continued formation of stars, is destroyed by radiation from already formed stars. As
such, they are regions of strong radiative feedback that plays a part
in the regulation of the overall star formation rate from individual molecular
clouds to entire galaxies by limiting the column of molecular gas that is
available for gravitational collapse (McKee 1989).

Today, models of the dynamical evolution of molecular cloud structures from
GMCs down to star-forming cores 
always try to avoid the boundary problem, because there is no 
straightforward definition of a molecular cloud boundary (Scalo et al. 1998).
With the input on the chemical behavior in the outer layers of
molecular clouds exposed to an external radiation field obtained from
PDR observations it will be possible to feed these models with a
physically justified definition of cloud boundaries.  The combined
effect of UV radiation, turbulent interstellar flows, and shock
structures created by expanding ionized regions creates a complex
chemical and physical behavior which has been studied for several
years now, but still leaves a large number of open questions
(Hollenbach \& Tielens 1999).  To understand how interstellar matter
is transformed through photodissociation from the molecular to the
atomic phase, and vice-versa, 
a complex and time dependent reaction scheme has to be
investigated. This needs the observation of the composition,
evolution, and dynamical behavior of the interstellar material in the
presence of massive star formation.

HIFI is the ideal instrument for this purpose. Its spectral range
covers the frequencies where most of the energy from PDRs is emitted
in lines. Its wide frequency coverage allows to observe many species
from the complex chemical network in PDRs under similar conditions and
its high frequency resolution allows to measure the exact velocity
structure of the interstellar material in the observed region.  Thus
HIFI can reveal the interplay between dynamics and chemistry in PDRs.
%The velocity resolution will also allow to distinguish between
%contributions from the warm ISM in the environment and locally formed
%warm and hot gas. 
With the help of theoretical models we will be able
to deduce from the observation the physical conditions of the gas
(pressure, density and radiation field intensity) and its chemical
composition.

The main aims of the HIFI observations can be summarized as follows:
\begin{itemize}
\item{}Determine the dynamical structure of cloud surfaces
\item{}Obtain an exact measure for the energy balance in PDRs
\item{}Probe the most relevant chemical processes in PDRs
\item{}Investigate the role of the dust properties on the thermal and
  dynamical properties
\end{itemize}

\subsection{The chemical structure}

In the interstellar chemistry framework, molecules are formed
gradually from atoms and ions. Chemical models describe the production
and destruction of interstellar molecules through a series of
gas-phase chemical reactions and solid phase reactions on dust grains,
many of which are photo-induced.
Many recent investigations on the chemistry of molecular cloud
surfaces concentrate on the species CO, C, and C$^+$.  They
form one of the key loops in the chemical network and provide
some of the most important cooling lines (Le Bourlot et al. 1993,
Sternberg \& Dalgarno 1995).
Low-lying rotational transitions of CO up to $J=7-6$ and two
fine-structure transitions of atomic carbon can be observed
from the ground. Carbon bearing species also play a key role
in PDR chemistry and can be uniquely observed with HIFI. We
will study the carbon cycle including CO, C, C$^+$, CH, and
PAHs. But with the large spectral coverage of HIFI it will also be
possible to scan other abundant species playing important roles in the
chemical network. By obtaining a comprehensive inventory 
of species, current chemical models of the the interstellar medium
can be tested an revised. One example clearly demonstrating the
importance of such observation is the extremely low abundance
of molecular oxygen found in observations with the SWAS and
ODIN satellites (Goldsmith et al. 2002, Pagani et al. 2003).
This was completely unexpected from most current chemical models
(Bergin et al. 2000, Spaans \& van Dishoeck 2001, Viti et al. 2001).
The observation of other species will help to resolve the puzzle.

\begin{figure}
\includegraphics[width=15.8cm]{3_transition_zones.ps}
\caption{Transition from atomic to molecular gas as a function of
the cloud depth ($A_{\rm V}$) for different species in a model of
a PDR with strong UV field. Depending on the corresponding chain
in the chemical network (H, C or N are illustrated here), the
molecular species is formed at a different depth: H$_2$ at
$A_{\rm V}$ = 0.5 mag., CO at $A_{\rm V}$ = 2 mag
and N$_2$ \& NH$_3$ at  $A_{\rm V}$ = 4 mag, corresponding to
different possible definitions of a ``molecular cloud boundary´´.}
\end{figure}

The description of the different phases of the interstellar medium as
ionized, atomic or molecular, based on the predominant form of
hydrogen
does not reflect the state of other species with different ionization
potentials. The transition from atomic carbon to molecular CO depends
on many conditions and occurs far deeper from the surface layer than the
transition from atomic to molecular hydrogen. It might therefore be problematic
to use CO as a surrogate for molecular hydrogen, especially in the presence
of a noticeable UV field and/or a low metallicity environment.
The same is true for the recombination/dissociation reaction schemes
of other species like nitrogen bearing molecules. Consequently, the
definition of a molecular cloud, has to take into account that its
boundary is different depending on the molecular tracer observed
(see Fig. 1). New
observations are needed to resolve the mutual distribution of the
transition regions seen in different molecules. The observation of
several tracers in their low-energy states by HIFI will reveal the
chemical state in the outer layers of molecular clouds exposed to a
dissociating radiation field.

A completely new field in the study of the chemical structure of
PDRs is opened by the ability of HIFI to observe many
hydrides in their lower energetic states. Hydrides convey
important chemical information since they are formed in the first
steps of chemical networks, and their abundances are therefore 
sensitive to a small number of chemical reactions.  
With Herschel (HIFI plus PACS), the oxygen hydrides OH, H$_2$O,
H$_3$O$^+$ \& OH$^+$ will be probed
extensively. The production rate of
OH and H$_2$O in PDRs depends critically on the gas temperature and 
ionization degree, as does the production rate of their main chemical 
precursor H$_3$O$^+$. The details of this process are still not
well known. As the carbon hydrides CH, CH$^+$ \& CH$_2$ initiate the carbon
chemistry, it is extremely interesting to probe these species
directly in dense environments where many other carbon species
have been detected (e.g. Teyssier et al. 2004). Such data are absolutely
necessary for testing the coherence of model predictions for molecular
abundances. The detection of large amounts of CH$^+$ and CH in
absorption against nearby stars and the observation of large
abundances of HCO$^+$ and OH in regions with no CO emission (Hogerheijde
et al. 1995, Lucas \&
Liszt 1996) indicate that diffuse interstellar clouds and the outer
regions of PDRs are more chemically active than anticipated (Lucas \&
Liszt 2000, 2001). However, Liszt \& Lucas 2002 have shown that the
cm lines of CH are no good tracer for the CH abundance, so that
HIFI observations are absolutely necessary.
In addition, CH$^+$ is particularly interesting as a potential
diagnostic for the interstellar medium, since this molecular ion is
destroyed rapidly (by collision with H, H$_2$ and other neutrals), and
may keep a memory of its formation process in its velocity
distribution and rotational excitation (Flower \& Pineau des For\^ets
1998, Black 1998). Its kinematic properties provide important
information concerning the physical conditions in the formation
regions.

Although NH$_3$ is observed as an abundant species in molecular
clouds for decades now and N$_2$H$^+$ has recently obtained much
attention due to its low depletion rate, there is still a general
lack of knowledge regarding the nitrogen-bearing part of
the chemical network corresponding to the carbon or oxygen network discussed
above. With the ability of HIFI to observe the nitrogen hydrides
in their ground-state, we face for the first time the possibility to resolve
the nitrogen part of the chemical network in PDRs.
The chemical routes towards  NH and NH$_2$ formation are as uncertain as
those for CH$^+$, but here we lack even more data.
The nitrogen chemistry is sensitive to neutral-neutral reactions with
very small barriers but hardly known reaction rates. 
For instance, the formation of N$_2$ involves two neutral-neutral
reactions, with NO as intermediate product (Gerin et al. 1992). 
NH is formed in a slightly endothermic reaction between N and NH$_2$.
Other reactions may play an important role as well.
Studying the chain of species N$^+$, NH$^+$, NH, and NH$_3$
may help to solve the puzzles of the nitrogen chemistry.
The comparison with other species will observationally reveal the
difference in the transition from atomic to molecular gas in
dissociation regions and general molecular cloud
boundaries between different tracers.

Another important part of the chemical structure in PDRs is given
by the solid phase chemistry. The formation of H$_2$ depends on the
structure of the dust surface (Cazaux \& Tielens 2004, Habart et
al. 2004). The ice mantles are expected
to be destroyed by the combined influence of the enhanced
temperature and UV radiation in PDRs thus releasing water, OH
and other molecules. In moderate UV fields the ice mantles are
chemically processed leading to complex organic compounds.
These processes clearly depend on the details of the dust
structure. Because water, CO, NH$_3$ are among the main constituents of
ice mantles, HIFI observations are ideally
suited to determine the chemical structure of the ice mantles
in the processes of their destruction.

The dust composition and size distribution determine the
attenuation of FUV radiation by the interstellar matter. 
These properties have therefore an important impact on the gas
composition and thermal balance. The fraction of nanoparticles
is a key parameter in the heating budget as the efficiency of the 
photoelectric effect is a strongly decreasing function of the particle
size. Combined with the data obtained by ground based and previous
space missions, the full opening of the far infrared wavelength range
will enable to cover completely the Spectral energy distribution, 
hence determine more accurately the dust properties in PDRs.

With HIFI observations of essential key species in the chemical
network in PDRs we will be able to resolve the chemical structure
of PDRs as molecular cloud boundaries much better than when
restricted to ground-based observations. They will help to
distinguish the role of the dust properties from the
role of the spatial configuration, including clumpiness,
for the chemical structure of a PDR (see Spaans 1996 Pety et
al. 2004, Abergel et al. 2003, Schneider et al. 2003).

\subsection{Dynamical structure}

Most models of
photon-dominated regions are static, and use a pure gas phase
chemistry, but dynamical effects
play a significant role in these regions. Störzer et al. (1997) have
shown e.g. that the modeling of the [C{\sc I}] emission from PDRs is
only possible when a non-equilibrium chemistry is taken into account.
Decamp \& Le Bourlot (2002) have shown that the dynamical transport of
species leads to a spatial distribution of molecules which
deviates by several orders of magnitude from their equilibrium
behavior. The dynamical impact from expanding H{\sc II} shells and
outflows leads to local shock structures with a special shock
chemistry and a large-scale injection of momentum (Flower \& Pineau
des For\^ets 1998). High velocity
shocks can be easily identified from the spectral
shape of the main cooling lines, but it will be more difficult to
detect shocks with low Mach numbers, which might still be responsible
for the formation of clouds in the diffuse medium (V{\'a}zquez-Semadeni
2000, Hennebelle \& P\'erault 2000, Audit \& Hennebelle 2004).

Advection
and turbulent flows may be important at the edges of molecular clouds
(Habart et al. 2001) and the ``evaporation'' of PDRs, i.e. the
formation of gas streams from the surface of the PDRs was discussed by
Gorti \& Hollenbach (2002).
Gerin et al. (1998) and Lemaire et al. (1999) measured a
velocity gradient in NGC 7023 which could trace this evaporation flow.
Moreover, it is still not clear how kinetic energy
can be fed into dense interstellar clouds. The currently available
data on the velocity dispersion in the ionized, the atomic and the
cold molecular phases, still do not provide a clear picture how the
process of cloud condensation goes with the dissipation of turbulent
kinetic energy (Wyrowski et al. 2000, Pety \& Falgarone 2000, 2003)
and how efficiently kinetic energy is transformed between the
different phases. The combined observation of the profiles of all
main cooling species will distinguish between the role of primordial
turbulence of the molecular clouds and the driving from young objects
or the radiation field. 

It is clear that the static description
of PDRs should be replaced by dynamical models including advection,
gas evaporation, shocks and turbulent flows. The measurements of line
profiles with HIFI will provide key input for the construction of such
models. Indeed, it is necessary to investigate the interplay between dynamical
processes changing the structure of the PDRs and the chemistry of the
interstellar gas and grains. For instance, the dissipation of
turbulence is an additional heating source, highly localized in
specific regions of a turbulent flow (Pety \& Falgarone 2000).
Advection of cold gas into the PDR is likely to affect the chemistry.
Reactive ions including CO$^+$ hav been detected in ground-based observations
in a few sources, namely the Orion Bar, M17, S140, NGC\,7027, and NGC\,7023
(St\"orzer et al. 1995, Black 1998, Fuente et al. 2003 and references therein),
at a high abundance level which is a good evidence for dynamical processes.
% It is well known
% that the dynamical behavior of PDRs is far from equilibrium but the
% temporal behavior of the chemical network and the energy balance
% indicates that the PDRs still form a stationary configuration.
HIFI observations of molecules like CH and CH$^+$ will be able to
measure the magnitude of these advection
processes in PDRs.  With the possibility of measuring detailed line
profiles for many species dominating the chemical network, a
quantitative assessment of the interplay between gas dynamics and PDR
chemistry will become possible.

\subsection{Energy balance}

A critical issue for the understanding of the physics and
chemistry of PDRs is their exact energy budget as a function
of depth from the ionizing front into the molecular cloud.
Here, the main heating and cooling species have to be observed
to determine their abundance and excitation and thus their
role in the total thermal budget. HIFI observable
lines provide an important measure here, but have to be complemented
by ground-based observations and existing data from other
satellite missions.

A major heating source of the interstellar gas is provided by the
photoelectric effect on dust grains. It depends mainly on the 
strength and spectral shape of the UV radiation field, of the
dust composition and the size distribution (Weingartner \&
Draine 2001). The heating rate depend also on the density structure
providing the shielding of the UV field and the coupling between dust and gas.  The radiative cooling of the gas at low temperatures is
provided mainly by molecular CO whereas at higher temperatures the
dominant cooling is provided by fine structure lines,
[C{\sc II}] at 1.9 THz, [C{\sc I}] at 492
and 809 GHz, a number of H$_2$O and CO transitions in the HIFI
spectral range, and [O{\sc I}] at 4.75 and 2.06 THz, [N{\sc II}] at 1.46 THz,
and OH at 2.5 THz in the spectral range covered by PACS.
Fixsen et al. 1999 found that even CH provides a significant contribution
to the line cooling. As predicted by models, the [O{\sc I}]
line at 4.75 THz is found to be the
most important cooling line in PDRs with strong radiation fields and
high densities, even stronger than the [C{\sc II}] line there
(e.g. Schneider et al. 2003).

In very dense regions, the main cooling of the gas
is, however, provided by the coupling to the dust which can radiate
efficiently throughout the whole FIR continuum. Ground-based
observations of the thermal dust continuum, and complementary PACS
observations are needed to estimate this contribution. The PACS
observations will trace the dust temperature structure with a
resolution not available in previous space missions. Combining them
with observations of the long-wavelength tail of the emission
by instruments like SCUBA at JCMT, SHARCII at the CSO, and LABOCA at
APEX allows to determine dust temperature and column density separately
providing a complete picture of the dust cooling. In addition
ISO and MSX data have to be used to obtain additional information
on the dust properties.

From the viewpoint of the chemical network the main cooling agents are
interrelated by some chain which may be written as C$^+$ -- C -- CO for
carbon and O -- OH -- H$_2$O for oxygen. The most important mechanism for the formation of
C$^+$ is the ionization of C, stemming from the photodissociation of
CO.  H$_2$O is photodissociated into OH and further down to O by UV
photons from the H{\sc II} region.  The details of the cooling balance
between C$^+$ and O depend on the UV intensity and the gas density
determining the photodissociation and recombination rates (Hollenbach et al.
1991). The
determination of the energy balance thus asks for a combined solution
of the chemical abundance and the temperature problem. HIFI
observations can attack this problem by measuring two transitions from
the most important cooling species and combining this information with
ground based continuum observations providing information about the
column density and the temperature of the dust.

Altogether the three problems of the chemistry, the dynamics, and the energy
balance in PDRs are tightly interrelated. The heating and cooling
of the gas drives the gas dynamics and the flow of material modifies
both the local energy balance and the distribution of species.
The temperature structure directly influences the chemical composition
of gas and dust awhile the abundances of the different components
acting as heating and cooling agents determines the energy balance.
The dynamics can be created by thermodynamics or vice-versa, and
the dynamics can trigger the chemistry.
Thus the three topics have to be investigated in an integral approach.

\section{Shock regions}

\subsection{J-Shocks and C-shocks}

Observations of the interaction of supernova blast waves with the
surrounding atomic and molecular gas provide an excellent tool to
support our theoretical understanding of molecular shocks. Evidence
for shocks in the ISM is widespread, but shocks associated with SNR
blast waves (BW) provide an uncontaminated ``laboratory'',
with a negligible role of molecular outflows (star forming regions),
ionized gas in HII regions, and atomic and/or excited molecular gas
from photo-dissociation layers.

Because -- by the nature of the excitation conditions in the
(post)shocked gas - the critical information must be derived from
FIR-spectroscopy mainly, observational constraints are rare
today. High-resolution spectroscopy with HIFI, supplemented by targeted
observations with PACS, will fill this gap.
C- and J-shocks will reveal rather distinct spectroscopic footprints:
\begin{list}{$\bullet$}{}
\item
In a {\it "jump"} (J) shock gas parameters change abruptly in a
discontinuity. All the heating (but negligible cooling) occurs in this
jump, the gas subsequently cools in the hot post-shock relaxation
layer. J-shocks are mostly dissociative.
\item In {\it "continuous"} (C) shocks, parameters vary smoothly from
pre- to post-shock conditions. Heating occurs by friction between ions
and neutrals, in an extended layer, and heating and cooling processes
compete. The temperature rise is moderate, and C-shocks are more likely
non-dissociative.
\end{list}

It is intuitively clear, that in dissociative (J) shocks the
atomic/ionic fine-structure lines will be the main coolants, while
non-dissociative (C) shocks will be prominent in the
vibration-rotation transition of H$_2$ and the high-excitation molecular
lines of CO and H$_2$O.
We ran
a set of representative models (80 km/s J- and 30 km/s C-shock,
into a $5\times10^3$ cm$^{-3}$ medium) to quantify these conclusions:
reference molecular abundances like CO, H$_2$O and OH are reduced by
factors 5 -- 10 in the J-shock. Conversely, in the slow C-shock,
atomic abundances are low(er) by 1-2 orders of magnitude. In both
scenarios, the characteristic temperature and (less) density
structure of the shock is reflected in the excitation of these
species: although not spatially resolved, its spectroscopic
footprint will allow to derive the physics of the shocks.

Reality though has never been that clear-cut, and along a given
line-of-sight and with our finite beams, the spectroscopic footprint
of a SNR/BW interaction with the surrounding cloud can easily be
confused by superposition of different shock types. Assuming a
constant ram pressure (inside the SNR bubble), density variations in
the surrounding medium will results in different types of shock (this
is observed in IC443, where the northern filament is modelled by a
fast J-type shock (v$_s$ $\sim 100$ km/s) into low/medium dense gas
(10 -- 10$^3$
cm$^{-3}$), the southern rim is represented by a slow C-type shock ($\sim 30$
km/s) ramming into a denser molecular cloud (10$^4$ cm$^{-3}$) (Burton et
al.1990, Rho et al. 2001). On smaller scale, the BW will respond
differently to the clump and interclump phases of the ISM, creating a
complex (time-dependent) scenario, in which a C-shocked dense(r) clump
is embedded in the post-(J)shocked envelope. Even more, in many cases
the shock structure may not have reached steady state, and the
temporal development must be modelled.  The H$_2$ rotational diagrams 
(e.g. Cesarsky et al. 1999, Haas et al. 2003) often imply the presence
of both hot (1200 - 2000 K) and warm (400 - 600 K) gas components. They
typically show a curvature that requires (at least) two C-shocks with
different post-shock gas temperatures.

\subsection{The OH and water problem}

% The atomic fine structure emission in these sources has generally been
% attributed to fast ($v\sub{shock} = 50-100$ km s$^{-1}$)
% dissociative J-shocks propagating in gas of moderate
% ($n[{\rm H}_2] = 10^2 - 10^3$ cm$^{-3}$) pre-shock density,
% while the excited H$_2$ and high-$J$
% CO emissions are naturally explained by slower  ($10 - 40$ km s$^{-1}$)
% nondissociative C-shocks in gas of higher ($10^4 - 10^5$ cm$^{-3}$) preshock
% density.  
% 
% While the atomic fine structure emissions and the presence of warm H$_2$ and
% CO appear to be understood relatively well, the observations of 1720 OH
% masers and the relative weakness of the H$_2$O rotational emissions remain
% puzzling.

While the physical structure of the shock will be confined by the proposed
excitation studies, HIFI will also allow to address the chemistry across
the shocked gas layer, in particular the puzzling chemistry of H$_2$O and OH,
indicated by observations of 1720 MHz OH maser emission in the post-shocked gas
but a relative weakness of the H$_2$O rotational emissions.
Several authors have discussed the 1720 MHz OH maser observations (e.g.
Wardle 1999, Lockett et al. 1999).  A
consensus has emerged that the physical and chemical properties required -- gas
temperatures in the range 50 - 200 K and OH abundances of order $10^{-5}$ --
could be attained behind a C-shock that is subject to a UV radiation field.
In this picture, H$_2$O is produced very efficiently at post-shock
gas temperatures above 400 K as a result of the neutral-neutral
reactions

\noindent\hspace*{3cm} O + H$_2$ $\rightarrow$ OH + H\\
\hspace*{3cm} OH + H$_2$ $\rightarrow$ H$_2$O + H

These reactions have activation energy barriers and are rapid at high
temperature but negligible slow at lower temperatures.  Once the
post-shock temperature drops below 400~K, however, the production route
to H$_2$O switches off, and the water is slowly dissociated to form OH.
The UV radiation fields needed to yield the required OH column densities
are modest, corresponding to $G_0 \approx 10^{-2}$ (in units of the
mean unshielded interstellar radiation field).   Such radiation fields
might result (1) from external irradiation (Lockett et al. 1999),
provided the extinction is not too high; (2) from nearby J-shocks (Snell
et al. 2003); or (3) from Lyman and Werner band photons excited by the
secondary electrons from cosmic rays (the ``Prasad-Tarafdar'' mechanism;
proposed in this context by Wardle 1999), provided the cosmic ray
ionization rate is sufficiently enhanced ($\zeta \approx 10^{-15}$ s$^{-1}$)
in the vicinity of the SNR.

However, the UV photodissociation invoked to explain the OH masers
cannot adequately account for the low water abundances inferred in these
molecular clouds.  Standard C-shock models predict high water
abundances, with $n($H$_2$O$)/n($CO$) \approx 3$, in the warm shocked gas,
a prediction confirmed by ISO observations of Orion, whereas upper
limits obtained toward IC443, 3C391, W28 and W44 by ISO (or, for one
position in IC443, a detection) all imply $n($H$_2$O$)/n($CO) ratios no larger
than $\approx 0.4$.  Because the neutral-neutral production routes
are so efficient, unreasonably large UV radiation fields would be needed
to suppress the water abundance by a factor $\approx$ 10 in the
warm/hot post-shock gas ($G_0 \approx 50$ is required at $T = 600$ K and
$G0 \approx 5000$ at $T = 2000$ K).  This led Snell et al. (2003) to
invoke the pre-shock freeze-out of oxygen onto icy grain mantles in
explaining the low observed water abundances.  In this picture, the
depletion of oxygen can bring the gas-phase C/O ratio close to unity;
thus the amount of gas-phase oxygen, not bound as CO, is substantially
reduced and so is the resultant post-shock water abundance.  Freeze-out
has previously been invoked (e.g. Bergin et al. 2000) to explain the low
H$_2$O and O$_2$ abundances observed in quiescent molecular clouds by the
Submillimeter Wave Astronomy Satellite (SWAS).

The freeze-out hypothesis presents additional problems, however.  Models
for the destruction of icy grain mantles in shocks (e.g. Caselli et al.
1997) suggest that shock velocities of 10 to 15 km s$^{-1}$
are sufficient to destroy icy grain mantles by means of gas-grain and
grain-grain collisions.  The shock velocities of $\approx$ 35 km s$^{-1}$
required by standard C-shock models to excite emission observed from
high-lying states of H$_2$ would therefore release large amounts of water
ice into the gas phase, violating the observational constraints on the
H$_2$O abundance.  However, Snell et al. have pointed out that shocks as
slow as $\approx$ 10 km s$^{-1}$ could yield the post-shock temperatures
$\approx$ 2000 K implied by the H$_2$ emission, provided that the
fractional ionization is enhanced sufficiently.  An increase in the
ionization reduces the ion-neutral coupling length, causing the shock
energy to be dissipated within a smaller but hotter region.  A
fractional ionization of $10^{-5}$, which could result from the
photoionization of sulphur given rather modest UV fields ($G_0
\approx 0.01$ to 0.1), would thereby allow shocks as slow as 10 km s$^{-1}$
to result in post-shock temperatures as large at 2000 K; such shocks
would leave icy grain mantles intact.

\subsection{Water observations with HIFI}

HIFI will yield the observational data needed to make further progress
in understanding the interesting physics that has been invoked to
explain the puzzles posed by previous ISO observations of water.  Its
much larger aperture and more sensitive detectors will greatly improve
the sensitivity to water vapor emission, allowing many transitions
to be observed.  Multitransition observations will allow the temperature
of the emitting water to be probed, and will indicate whether the water
abundance is uniformly suppressed at all gas temperatures.  The
temperature-dependence of the water line ratios is exhibited by the
results in Table \ref{tab:shock}. Here, we  ran a  set  of  water
excitation models  for gas  temperatures of 100,  500 and 1000K, and
easily have identified transitions for which predicted brightness vary by
more than 2  orders of magnitude, ranging from  ``undetectable'' (at 100K)
to ``easily detectable'' (1000K) [N(H$_2$O) = 10$^{16}$ cm$^{-2}$].
In the next step of model refinement, we will fold the predicted
shock physics (temperature,  density, chemistry  across the shock)  with the
molecular  excitation to  derive  at specific  predictions (vs.  shock
parameter) [PhD work].

% VO: Paragraph describing the physical conditions used in the model is missing!
\begin{table}[ht]
\begin{center}
\caption{\label{tab:shock} Predicted water line temperature for various shock temperatures}
\vspace{1mm}
\begin{tabular}{lrrrrr}
\hline
Water transitions & E$_{low}$ [K] & $\nu$ [GHz] & T$_A$ & T$_A$ &
T$_A$ \\
& & & 100K & 500K & 1000K \\
\hline
\hline
o-H$_2$$^{16}$O 1$_{1,0}-1_{0,1} $ & 0.0 & 556.9 & 1.76 & 4.75 & 5.84 \\
o-H$_2$$^{16}$O $2_{1,2}-1_{0,1}$ & 0.0 & 1669.9 & 0.26 &  1.87 &3.35 \\
o-H$_2$$^{16}$O $2_{2,1}-2_{1,2}$ & 114.4 & 1661.0 & 0.01 & 0.17   & 0.34 \\
o-H$_2$$^{16}$O $3_{0,3}-2_{1,2}$ & 114.4 & 1716.8 & 0.03 &  0.75 & 1.77 \\
o-H$_2$$^{16}$O  $3_{1,2}-3_{0,3}$ & 196.8 & 1097.4 & $<0.01$ &  0.53 &  1.30 \\
o-H$_2$$^{16}$O $3_{2,1}- 3_{1,2}$ & 249.4 & 1162.9 & $<0.01$ & 0.28 & 0.92\\
\hline
p-H$_2$$^{16}$O $2_{0,2}-1_{1,1}$ & 53.5 & 987.9 & 0.65 & 5.29 & 9.36 \\
p-H$_2$$^{16}$O $1_{1,1}-0_{0,0}$ & 0.0 & 1113.3 & 1.27 & 5.50 & 8.00 \\
p-H$_2$$^{16}$O  $2_{1,1}-2_{0,2}$ & 100.8 & 752.0 & 0.13 & 2.61 & 5.10\\
p-H$_2$$^{16}$O  $2_{2,0}-2_{1,1}$ & 136.9 & 1228.8 & 0.02 & 0.70 & 1.16
\\
\hline
\hline
\end{tabular}
\end{center}
\end{table}

It is important to note, that HIFI's high spectral resolving power
will be crucial to provide the kinematic information that is needed
to disentangle the emission of the (generally, slightly velocity-shifted)
shocked gas from the emission of the quiescent cloud. This information
has been completely unavailable from the ISO grating spectra obtained
previously, and will not be available from measurements with PACS and SPIRE.

\section{Proposed HIFI observations}

\subsection{Selection of lines}

\input{table9.inc}

For the determination of the energy budget of the ISM it is essential
to distinguish between the relative abundance of the different cooling
species and their excitation providing the total cooling rate.
Thus the main cooling species within the chemically coupled
chains C$^+$-- C -- CO and O--OH--H$_2$O are to be observed
in at least two transitions (except for C$^+$ which has only a single fine
structure transition) to allow this distinction.
Towards selected positions in the PDR or the shock ridge, prominent
in the warm gas, we propose  to derive an as  complete as necessary
inventory of the critical CO, H$_2$O and OH lines for a detailed excitation
model  of the gas  layers. Corresponding to the strongly varying conditions
across the interface regions the transitions must span an as-wide-as-possible
range  of excitation conditions. Supplementary
studies of warm  CO are under way or will be  performed with e.g. APEX
(up  to J=8-7).  HIFI will  provide the  high-excitation  J=9-8, 10-9,
13-12, 16-15  transitions that are  crucial to constrain  the hot
gas phases.  These studies will be complemented
by  ground-based  observations of  the  atomic Carbon  fine-structures
lines (partly  done already with  CHAMP/CSO, and to be  extended with,
e.g., CHAMP+/APEX), for a complete inventory of the gas phase carbon reservoirs.
A  set of suitable water lines (with selected isotopes)  will be  observed,
that  cover the  expected wide  range of excitation  conditions  in 
the PDRs and shocks. Multitransition observations of water will allow
the temperature of the emitting material to be probed, and will indicate
whether the water abundance is uniformly suppressed at all gas temperatures.
The efficiency of water as a gas coolant will be clearly established by these
observations, which will also let us probe desorption processes
induced by UV photons.
HIFI data on the thermal OH, obtained from the $^2\Pi_{1/2}  3/2^- - 1/2^+$
transition at 1834.7 GHz, will be combined with observations of OH
ground-state transitions observed with PACS or SOFIA. Finally H$_3$O$^+$
and OH$^+$ observations will complete the inventory of the oxygen chemistry.

Ammonia will trace the dense well shielded molecular gas and help us
to separate the emission from the PDR or shock front from that of
the bulk of the
molecular cloud. Combining normal radio observations of
the inversion transitions with the measurement of the rotational
ground-state transition by HIFI allows to obtain a complete excitation
analysis tracing also the cold gas.  This has to be complemented by
observations of the other abundant nitrogen hydrides to allow a
determination of this tree of the chemical network. Observations
of the lower NH$_2$ transitions should be considered in case of a key
program on particular sources but would go beyond the scope of this
proposal because the complexity and weakness of the numerous NH$_2$
transitions ask for long integration times and a sophisticated
analysis.

Carbon hydrides will be observed in their ground state as tracers of
the first steps of the molecular chemistry. The combined
observation of the ground state line of CH at 536 GHz and the
next higher transition at 1669 GHz allows to perform both an excitation
analysis and the exploitation of the higher angular resolution at the high
frequencies. 

The list of lines envisaged for the observations is summarized in
the Tables \ref{tab-lines}, \ref{tab-warmlines}, and \ref{tab-pacs-lines}.
A rough estimate for the integrated line intensities can be obtained from the
LTE relative absorption coefficients $\kappa\sub{LTE,rel}$ in columns
5 to 8 via the radiometer formula, neglecting the very weak background
radiation, by
\begin{displaymath}
I\sub{LTE}=A \times\nu \left[ \exp(A \nu/T) -1 \right]^{-1} \times \left[1- \exp(-\tau)\right]
\end{displaymath}
with
\begin{displaymath}
\tau=\kappa\sub{LTE,rel} N\sub{mol}/\Delta v
\end{displaymath}
where $A=4.80\,10^{-5}$~K/MHz, $\Delta v$ is the line width, and $N\sub{mol}$ is the
column density of the considered species.
Para- and ortho-water are treated as separate species here, because
their mutual conversion occurs only on very long timescales
in interstellar conditions (Knacke \& Larson 1991). Because current
chemical models are hardly able to accurately deduce the abundance
ratio of ortho- to para-water produced in the destruction of icy grain
mantles, the HIFI observations of both molecular ladders can be used
to determine this ratio and adjust the models.
Only the HIFI observations will
allow to measure this ratio accurately in warm interstellar regions
without strong near- and mid-infrared emission.

Only  the spectroscopic information (line  profile) will allow
us to discriminate between  the various emission components (along the
line-of-sight and  within our  finite beams, most  profiles will  be a
composite of  contributions from the cool gas inside molecular clouds
and the warm gas heated by UV radiation or shocks. Thus
these observations will be unique  for HIFI (and
cannot be derived with e.g.  PACS).  In many cases, the large
IF bandwidth of HIFI allows to observe two or more lines simultaneously
at a single receiver setting.
This may lead to a considerable reduction of the overall observing time.
The corresponding lines are marked as ``bonus lines'' in Table \ref{tab-lines}.

The observation of all lines in the tables would allow to deduce a
fairly complete chemical inventory and an accurate overview on the
energy balance in the warm and dense interstellar medium. Unfortunately,
the line estimates for the different targets obtained in the appendix
show, that the full list of all lines can only be observed in the
brightest and nearest objects. The observation of all lines in Table
\ref{tab-warmlines} is only justified in the relatively hot SNR shock
regions.

Thus the full line list represents only a common framework allowing
an easy intercomparison between different regions and the maximum
scientific output for each individual source, but for most PDRs
and SNR shocks some of the lines will remain undetectable within a reasonable
amount of integration time. We propose to observe all lines only
in the two brightest sources. Moreover, we find for all sources, that
some lines can only be observed at their expected peak position while
a third list of lines can be used to measure the spatial structure
of the emission in the different species, especially at lower frequencies.
Depending on
the yet unknown instrumental noise of HIFI in band 6 it may be necessary
to substitute the CO 16--15 observations by corresponding 14--13 or 13--12
observations. For very hot
regions produced in J-shocks it is feasible to detect even higher
lines listed in Table \ref{tab-warmlines}. Here, the 1661.007 and 1162.911~GHz
transitions of ortho-H$_2$O
should be observed simultaneously with the 1669.905 and the 1153.126 GHz
lines to obtain a more complete coverage of the excitation state.

% RG: ANMERKUNGEN/ERGÄNZUNGEN ZUR TABELLE 2:
% (4) BEIDE GROUND-STATE TRANSITONS MÜSSEN IM ISOTOP GEMESSEN WERDEN: 110-101 FEHLT

Because of the frequency range and the need for a good spectral resolution
to deduce the velocity structure within PDRs, there is no alternative
to HIFI for the observation of all these lines. For [C{\sc II}] and
a number of hydrid lines corresponding observations by SOFIA might
become possible in the next years but it is not yet clear which 
observations can be performed there on a schedule comparable to the
planning for HIFI.

To solve the energy balance problem, the HIFI observations have to
be complemented by the observation of lines at somewhat higher
frequencies by PACS observations as discussed below.
Table \ref{tab-pacs-lines} gives the parameters of these lines.

\subsection{Proposed sources}

For many warm and dense regions in the interstellar medium we expect
still important new results from observations which will be carried
out before the launch of Herschel. Thus our current knowledge on
particular sources may still considerably change during the next years.
We propose in Tab. \ref{tab_sources} a source list, based on current
knowledge, but because modifications of that list should still be
possible in the next years, reflecting new discoveries, it is actually
more important to focus on general criteria for the source selection.

\input{table11.inc}

The sources should:

\begin{itemize}

\item cover a wide range of cloud properties viz. density (10$^3$ to
  10$^7$ cm$^{-3}$), radiation field (1 to 10$^{5}$ Draine fields),
  shock velicity, dust properties, and properties of the ionization front.
The criterion for a large variety of radiation fields translates into a
range of distances of the dense regions from the illuminating source, and
into a range of galactocentric distances determining the average
external radiation field.

\item be located nearby in order to probe small linear scales.

\item have a well defined orientation with respect to the observer so
  that it is possible to analyze the stratified structure of the interface
  region. However, their orientation should not be exactly ``edge-on''
  because we want to detect a line-of-sight velocity of evaporation
  and shocked material.

\item have a relatively simple geometry with respect to the
  configuration of the illuminating sources and the molecular cloud.
  This implies that the molecular cloud should preferably not contain
  additional regions with triggered star-formation providing
  additional radiation fields and winds.
  The observations should focus on single transition regions between ionized,
  atomic and molecular gas created by the interaction between a
  well known radiation field and the molecular cloud.

\end{itemize}

% RG: DA MUSS ICH EIN BISSCHEN IN DER LITERATUR WÜHLEN; UM DIE SHOCK-QUELLEN
% ZU BESCHREIBEN.

The pre-requisite of being able to  combine the HIFI observations with
mid-infrared data on the PAH emission, needed to analyze the full carbon
cycle, biases the selection of the sources towards sources which have
been studied by ISO or MSX, i.e. mainly sources close
to the Galactic plane.

In the selection of PDRs we prefer either regions where one molecular
cloud shows both an interface to the H{\sc II} region
around the illuminating source and an interface to the general
diffuse ISM or regions where the same source illuminates two
different clouds at similar distances. This approach allows to
distinguish the impact of the radiation field from the influence
of the cloud properties.

As of today only a handful of supernova blast-wave sources has been
identified that potentially qualify for a detailed study with HIFI.
IC443 is the best studied case of a SNR - molecular cloud interaction.
Other sources, like  W28, 3C391,  and W44 should be included to trace
the influence of different physical conditions and shock velocities.

The preliminary list of sources
fulfilling most of the mentioned conditions is given in Table
\ref{tab_sources}. The table gives some rough values for 
physical parameters to allow a mutual comparison of the different
regions. The actual structure of each object is, of course, too
complex to be described by a few numbers.
A detailed description of their properties
including an overview of all existing observations is given in the
Appendix. For the sake of completeness we have also added two clouds
representing the properties of molecular clouds which are only
exposed to the average interstellar radiation field and the large-scale
velocity structure. Although they are relatively cold -- and thus
not fitting into the title of this key project -- they need to be
included for a detailed comparison between the different clouds
allowing to deduce the role of the average interstellar radiation field
and shocks in the turbulent velocity structure of the Milky Way.
It is anticipated that the availability of observational data on the
major PDR tracers will further improve during the period leading
to the launch of Herschel. 

The full list contains 12 sources which are ``classical PDRs'',
4 shock-dominated sources produced by supernova blasts, and 2 cool
clouds at average interstellar conditions. Only when all of these
sources are observed, we obtain a relatively complete coverage
of the parameter space dominating the physical processes in the warm
interstellar medium. It will then provide a good feeling
for the processes governing the common behavior and the differences
between the sources. 

For the Orion Bar we expect that aside from this proposal
the PDR will be target of a full frequency survey in an early
flight phase revealing possibly unknown lines or other
unexpected effects. This information could be used to adapt the observations
of the less prominent PDRs forming the main part of this
proposal.
\label{sect_pdr_sources}

\subsection{Observing mode}

To deduce the physical and chemical structure of the PDRs and shock
regions we do not
propose full mapping observations but to use OTF scans of two stripes
tracing the transition in the PDRs from the ionized medium into the
cold dense cloud structure for most sources. The selection of two
stripes guarantees
that the measurement is not dominated by local peculiarities while
providing information on how much the depth dependence is influenced by
local small-scale fluctuations in the density and temperature
structure.  The two stripes should basically measure the same PDR
depth structure thus being separated only by the minimum distance
needed for spatially independent measurements. Restricted by the
beam size at 540 GHz, this results in a optimum distance between the
two stripes of 40$''$. To allow a computation of line ratios
between lines taken at different frequencies, the higher
resolution of HIFI at higher frequencies has to be compensated
by measurements of additional stripes separated from the nominal
position by half a beam width at higher frequencies. These
additional measurements can be taken at much shorter integration times.
The details of the corresponding computation are described as
``Template beam synthesis observations'' by Teyssier (2000).
In case of sources where a simple geometry without peculiarities
was previously confirmed by high-resolution observations of dust
and other tracers it may be possible to restrict the observation
to a single stripe.

The length of the stripes depends on the physical length scales
of the transition regions which is clearly correlated to the
UV flux and the clumpiness of the region. We expect the emission
from different species to be restricted to different parts
within this transition region, so that we can optimize the length
of these stripes for the different species based on previous
data and chemical models. Due to the limited
observing time the observations of the weakest proposed lines will only
be possible at a few selected positions within the stripes.
The location and the length of the stripes for each transition in
the PDR has to be optimized based on previous models and observations.

In all edge-on configurations the location of the illuminating source
can be used as OFF position free of emission for most molecular species.
The [C{\sc II}] and probably also the [C{\sc I}] observations
have to be performed in frequency switching mode due to the large
extent of the emission.

% RG: KAPITEL IST FÜR DIE PDR GESCHRIEBEN. FÜR DIE SCHOCKS MÖCHTE
% ICH NICHT NOTWENDIGERWEISE AUF DIE NIEDRIGSTE AUFLÖSUNG KONVOLVIEREN.
% WENN WIR ZB PARA-WASSER MODELLIEREN, KÖNNEN WIR GUT MIT DER AUFLÖSUNG
% BEI 1 THZ ARBEITEN. SPRICH. DIE BEOBACHTUNGSZEITEN MÜSSEN SO AUSGELEGT
% SEIN; DASS AUCH DIE HOCH-FREQUENTEN ÜBERGÄNGE HINREICHENDES SIGNAL-ZU-
% RAUSCHEN HABEN (OHNE VERSCHMIERUNG).

\subsection{Complementary observations}

\subsubsection{PACS observations}

Analysis of the ISM excitation conditions requires complementary
spectral data that are covered by PACS. Though the velocity
information will be missing in the PACS data, the integrated line
intensities will provide useful constraints. Thus complementary
PACS observations are mainly needed to address the problem of the
energy balance discussed above.  Here, O and OH provide a
major cooling of the interstellar gas through their
ground-state lines which are not observable by HIFI but can
be observed by PACS.
Table \ref{tab-pacs-lines} gives the parameters of the transitions to be
observed with PACS. The PACS observations have the big advantage that the
imaging capabilities of the instrument allow to observe a field of
approximately 50$''$ diameter in a single integration, so that
only few observations are necessary to cover the stripes through the
PDRs discussed as observing modes for the HIFI observations.

In addition to the line observations PACS will be essential to
measure the dust column density and temperature structure. With
the possibility to map the cloud simultaneously at 75 and 170~$\mu$m
with the PACS bolometers we can measure the dust spatial structure
in a relatively short amount of time giving essential input on
the general morphology of the PDR/shock region and the energy
balance in the cloud.

%PACS will be able
%to measure simultaneously in the 60-90~$\mu$m and the 90-130~$\mu$m
%bands. So it will be possible to obtain observations of the O line at
%63~$\mu$m and the OH lines at 119~$\mu$m in a single observing run.

%Since nitrogen has an ionization potential (14.5 eV) higher than that of
%hydrogen, [N{\sc II}] (122~$\mu$m) arises only in the ionized gas, in contrast
%to [C{\sc II}] which may arise from both the ionized or the neutral medium.
%Thus [N{\sc II}] is crucial in separating the contribution of the different
%phases of the ISM towards cooling of the PDRs. The [N{\sc II}] line at
%205~$\mu$m falling in the HIFI regime, is rather weak and demand long
%integration times.

\subsubsection{Ground-based observations}

Complementary ground-based observational data are needed to complete
the wavelength coverage of the physical and chemical processes : H$_2$
ro-vibrational lines in the near infrared are accessible from the ground.
This gives the spatial distribution of the excited H$_2$, an important
constraint on geometry and clumpiness.
Ground-based millimeter and submillimeter telescopes provide access to the
cold dust emission from the shielded cloud interior, and to the
rotational lines of CO and other``heavy'' molecules. They measure
the column density, thermal, and dynamical structure of the cold dense
inner regions of the sources. The observation of additional molecular species
is required to provide a comprehensive inventory of the chemical
network in PDRs complementing the species observed by HIFI. This
holds in particular for sulphur-bearing species like CS and SO which
can be easily observed from the ground and provide essential clues
to the influence of sulfur on the chemistry of the hydrides observed
with HIFI. Sulfur is a key species in PDR chemistry because ionised carbon
reacts with sulfur rapidly and produces atomic carbon.
% This charge
% exchange reaction is very important for the determination of the
% onisation degree in interstellar gas.

Since shocks and PDRs represent transition regions, a combination of
observations of different tracers is required.
The molecular component of the regions will be observed using the
multiple transitions of CO (up to J=7$-$6), $^{13}$CO (J=8$-$7) and the
NH$_3$ lines tracing different physical conditions. 
The atomic hydrogen component will be traced using ground-based
observations of the 21cm line while the fine structure transitions of carbon
at $492$~GHz and $810$~GHz provide the necessary information on atomic
carbon. Many of these
complementary ground-based observations have already been performed for
known PDRs (mostly in the northern sky). With the upcoming
instrumentation on the southern sky at the new telescopes ASTE and APEX
the next years promise to provide data with a similar quality for
southern objects. A complete inventory of the available data for each
of the proposed PDRs is needed. Work is in progress in this direction.
This will allow to select the optimum position and size of the stripes
to be mapped with Herschel, and provide the necessary information on the
warm molecular gas.


Combination of the atomic and molecular data will help to constrain the
PDR densities, column densities and excitation conditions. Many of these
complementary ground-based observations have already been performed for
many of the known PDRs (mostly in the northern sky). With the upcoming
instrumentation on the southern sky at the new telescopes ASTE and APEX
the next years promise to provide data with a similar quality for
southern objects. A complete
inventory of the available data for each of the proposed PDRs is needed.
This will allow to select the optimum position and size of the stripes
to be mapped with Herschel, and provide the necessary information on the
warm molecular gas.

The column density and dust temperature structure will be determined
by dust emission observations in the mm/submm range, i.e.
by MAMBO/SCUBA/SHARC observations at wavelengths between 1200 and
350 $\mu$m.

Ground-based studies in the visible and the near UV spectrum are to be
used to characterize the properties of the H{\sc II} region in PDRs
where the geometry allows an observation of the H{\sc II} region without
strong obscuration. Observation of the strength of the different hydrogen
lines and lines from single and double ionized oxygen, nitrogen, and
sulfur provide a measure for the the temperature, the electron density,
and the chemical composition of this part of the PDRs.

A large part of these ground-based observations are currently 
being performed in various projects to study PDRs. On the time scale
of the next two years these observations have to be used to identify
a final target list. In the remaining two years before launch
possible ``holes'' in the spectral and spatial coverage of these
observations for the final targets have to be closed
to guarantee that all the ground-based observations discussed above
have been taken for the HIFI targets before the HIFI observations
start.

\subsubsection{Complementary satellite and SOFIA data}

The available data from the ISO, MSX, and Spitzer missions have to be inspected
with respect to water lines, rotational and rotation-vibrational
transitions of H$_2$, PAH emission, the mid-infrared dust spectrum,
and the general spectral energy distribution in the PDRs.
The team might be interested in submitting open time
programs to Spitzer, for completing the database.

Information about the ionized medium and the ionizing processes
should be obtained by a careful analysis of X-ray observations
performed by Chandra and MMX. A characterization of the chemical structure
and the hydrogen distribution in the H{\sc II} regions can also
be obtained from Copernicus, IUE and FUSE data.

For most of the source candidates the 557~GHz H$_2$O line was observed
by SWAS or ODIN. The resulting low spatial resolution data serve as a
guideline for optimizing the HIFI observations for each individual source.

Depending on the future progress of SOFIA and the results of
first SOFIA observations it may turn out that it will be possible
to perform the [C{\sc II}] observations, the observation of all lines
planned with PACS and several other lines with SOFIA. In the
best case this may reduce the required observing time for HIFI
by more than a factor two if a sufficient amount of observing time
on SOFIA is available. Moreover, only SOFIA will provide 
resolved line profiles of the O{\sc I} and OH lines. With this additional 
information we can then also better trace the dynamical structure of the
warm gas at the surfaces of the shocks and PDRs.

\subsubsection{Relation to other HIFI key programs}

There is a strong interrelation between this key program,
with a focus on the general physics and chemistry of PDRs
and the key program on the Orion and Sgr B2 regions including extended
observations of the Orion Bar as one of the most prominent PDRs.
We expect here, that the Orion-Bar is the target of a full
frequency survey. This may reveal unknown
important lines and can thus be used to adapt the observations
of this proposal.

There is a clear overlap of this key project and the scope of
the project on water. We restrict ourselves here, however,
to a relatively low number of water transitions in low-energy
states, to the distribution of water as a function of the
radiation field, and to the details of the process of the formation
and dissociation of water in outer layers of molecular clouds and PDRs.

The sources to be observed may agree in several cases with
sources observed in the key program on massive star formation
because young massive stars obviously produce PDRs in their
parental molecular clouds thus providing viable source candidates
for this project. Here, we can expect synergy from common
preparatory work and ground-based observations of the sources
covering both the young stellar sources and their environment.
However, there is not much overlap in the HIFI observations
because the program on star formation will concentrate on the hot
stellar cores whereas the observations discussed here are
focused on their extended environment.

\subsection{Time estimate}

Detailed time estimates are provided independently for each
source in the Appendix. The results are summarized in Table
\ref{tab_sources}.

We plan to observe 12 sources mainly heated by UV radiation and 
4 regions mainly heated by shocks. The observation of the
12 classical PDRs requires {\bf 82 h of HIFI time} and {\bf 24 h of PACS
time}. The observation of the shock regions requires
{\bf 44.5 h of HIFI time} and {\bf 9 h of PACS time}. The observation of
the diffuse clouds for comparison requires {\bf 33.5 h HIFI time}. No
PACS observations are needed here. 

The resulting total observing time amounts to {\bf 193 hours}.

% \subsection{Open questions}
%
% \begin{itemize}
%
% \item{}Do we need to include the [N{\sc II}] transition at 2.46 THz
%   here?
%
% \item{}Do we need to observe the higher NH$_3$ transition with HIFI or
%   can we rely to use complementary ground-based NH$_3$ observations
%   for the excitation analysis.
%
% \item{}Is there any need for PACS continuum observations or can we
%   obtain sufficient information here from ground-based telescopes?
%
% \item{}Do we need [C{\sc I}] measurements with HIFI or can we rely
%   entirely on ground-based observations?
% \end{itemize}

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%%%%%%%%%%%%%%%%%%%%%%%%%%%%
% macros invented from the author
%
%       UNITS
%
% Macros invented from the author
%
\def\FIR{\ifmmode {\,$\tau_{\rm FIR}$} \else $\,\tau_{\rm FIR}$\fi}
\def\tk{\ifmmode {\,$T_{\rm k}$} \else $\,T_{\rm k}$\fi}
%
%       UNITS
%
\def\mic{\ifmmode {\,\mu{\rm m}} \else $\,\mu{\rm m}$\fi}       % micron
\def\kms{\ifmmode {\,{\rm km\,s^{-1}}}                          % km s-1
	\else {\hbox{$\,$ {\rm km$\,$s$^{\rm -1}$}}}\fi}
\def\mo {\ifmmode {\,{\it M}\odot} \else $\,M$\odot\fi} % M solar
\def\lo {\ifmmode {\,{\it L}\odot} \else $\,L$\odot\fi} % L solar
\def\my {\ifmmode {\,{\it M}\solar\,{\rm yr^{-1}}}              % Msol/year
	\else {$\,M$\solar$\,$yr$^{\rm -1}$}\fi}
\def\cmm#1{\ifmmode {\,{\rm cm^{-#1}}}          % cm-1, cm-2, cm-3,....  
	\else \hbox{$\,${\rm cm$^{\rm -#1}$}}\fi}
%
%      SYMBOLS
%
\chardef\isp="10\def\i{\'\isp}		% dotless i withaccent
\def\as {\ifmmode {^{\scriptscriptstyle\prime\prime}}           % arcsec
	\else $^{\scriptscriptstyle\prime\prime}$\fi}
\def\am {\ifmmode {^{\scriptscriptstyle\prime}}                 % arcmin
	\else $^{\scriptscriptstyle\prime}$\fi}
\def\deg {\ifmmode^\circ\else$^\circ$\fi}                       % degree
\def\raw {\ifmmode\rightarrow\else$\rightarrow$\fi}             % rightarrow
\def\x {\ifmmode\times\else$\times$\fi}                         % times (x)
\def\gsim {\ifmmode {\buildrel>\over\sim}               % greater or similar
	\else {\lower.6ex\hbox{$\buildrel>\over\sim$}}\fi}
\def\lsim {\ifmmode {\buildrel<\over\sim}               % less or similar
	\else {\lower.6ex\hbox{$\buildrel<\over\sim$}}\fi}
%
\def\ra[#1 #2 #3.#4]{ #1$^{\rm h}$#2$^{\rm m}$#3$^{\rm s}$.#4}  % RA
\def\dec[#1 #2 #3.#4]{ #1\deg#2\am#3{\as}.#4}             % declination
\def\rax[#1 #2 #3]{RA: #1$^{\rm h}$#2$^{\rm m}$#3$^{\rm s}$}%RA with
%no fraction
\def\decx[#1 #2 #3]{Dec:#1\deg#2\am#3\as}          % Dec with no fraction
%
%
%       MOLECULES
%
\def\HUNO{{\sc Hi}}                   % atomic hydrogen
\def\HII{{\sc Hii}}                   % HII
\def\OI{{\sc O\,i}}
\def\CI{{\sc C\,i}}
\def\CII{{\sc C\,ii}}
\def\uch{UC\,{\sc Hii} region}           % UC HII region
\def\h2{\rm H$_2$}                      % molecular hydrogen
\def\doce{\rm CO}                       % 12C O
\def\trece{\rm $^{13}$CO}               % 13C O
\def\ocho{\rm C$^{18}$O}                % C 18O
\newcommand{\cop}{CO$^+$}               % CO(+)
\newcommand{\hocp}{HOC$^+$}             % HOC(+)
\newcommand{\sop}{SO$^+$}               % SO(+)
%
%
%        SOURCES
%
\newcommand{\g}{G29.96-0.02}
%
%        Otros simbolos
%
\def\etal{{\it et al.}}
%
%       REFERENCES
%
\def\aua{{\rm A$\,$\&$\,$A}}            % A \and\ A
\def\araa{{\rm ARA$\,$\&$\,$A}}         % Annu. Rev. A \and\ A
\def\aar{{\rm A$\,$\&$\,$AR}}           % A \and\ A Review
\def\aap{{\rm A$\,$\&$\,$A}}            % A \and\ A
\def\aas{{\rm A$\,$\&$\,$AS}}           % A \and\ A Sup.
\def\apj{{\rm ApJ}}                     % Astrophys. J.
\def\apjs{{\rm ApJS}}                   % Astrophys. J. Supl.
\def\apjl{{\rm ApJ Let}}                % Astrophys. J. (Letters)
\def\aj{{\rm AJ}}                       % Astron. J.
\def\sva{{\rm SvA}}                     % Sov. Astron.
\def\pasp{{\rm PASP}}                   % Publ. Astr. Soc. Pacific
\def\pasj{{\rm PASJ}}                   % Publ. Astr. Soc. Japan
\def\mnras{{\rm MNRAS}}                 % Monthly Notices...  
\def\aass{{\rm A$\,$\&$\,$Spa. Sci.}}   % Astrophys and Space Science
\def\spscirev{{\rm Space Sci. Rev.} }

\newpage
\appendix
\section{Sources to be observed in the HIFI
key program ``The warm and dense interstellar medium''} 

\subsection{NGC3603 \protect\footnote{M.R\"ollig, C.Kramer, KOSMA}}

\subsubsection{Introduction}
NGC 3603 is located in the Carina spiral arm at a distance of 7.7 kpc (e.g. Moffat 1983; Clayton 1986; Melnick et al. 1989; De Pree et al. 1999) in the southern hemisphere ($\delta \sim -61^\circ$). It is often considered as Galactic starburst region and therefore is an important template for extragalactic studies.  With a bolometric luminosity L$_{\mathrm bol} > 10^7$ L$_\odot$  NGC 3603 is 100 times more luminous than the Orion  Nebula and has about 10\% of the luminosity of 30 Doradus in the Large Magellanic Cloud (LMC). The only other equally luminous region in the Galaxy is the Carina nebula at 2.7 kpc distance. NGC 3603 is the only Galactic giant H II region whose massive central ionizing cluster can also be studied at optical wavelengths. The Trapezium-like system HD97950 in the cluster core contains 3 luminous hydrogen-rich WNL stars, 6 O3 stars and many other late O type stars at an age of about 2 to 3Myr.

 The total cluster mass is $\ge$ 4000 M$_\odot$. To the south of the cluster (1.'3 distance ) is a giant molecular cloud (often referred to as MM2) Ionizing radiation and fast stellar winds  from the starburst cluster are excavating large gaseous pillars and a bubble ($a\approx$. 30'' radius) which expands with velocities of about  50 km s$^{-1}$ into the ambient medium. Remarkable is the very low density of obscuring material in front of the central cluster (A$_V=4...5$). This makes NGC 3603 a very interesting site to study the stellar content of massive star forming regions.\\ 
N\"urnberger et al (2002) carried out CS (2-1) and CS (3-2) observations (SEST) of the molecular gas associated with NGC 3603 and 
 identified 13 molecular clumps with radii less than 0.8 pc. He also derived virial masses as well as lower limits for their H$_2$column densities.
 Two pillar like structures are identified in NGC 3603, one west and one south (east) of the central cluster. So far three proto-planetary disks (proplyd's) have been found in NGC 3603.

\subsubsection{Published Observations}
The following observations have been done so far:
\begin{itemize}
\item  IRAS HiRes Data at 12, 25, 60, 100 $\mu$m.
\item CO survey of Carina arm (e.g. Grabelsky et al. 1987, 1988, Columbia Univ. 1.2 m millimeter wavelength telescope, beamwidth 0$^\circ$.125 ), higher resolution $^{12}$CO(1-0) observations covering an area of about 6'$\times$6' reveal the spatial  distribution and clumpiness of the molecular gas right towards the south of the OB cluster (Melnick 1989).
\item CS Maps (CS(2-1) and CS(3-2))  and C$^{18}$O(2-1) data and some CS(5-4) data (SEST,HPBW $\sim 52''$ at 98 GHz)(N\"urnberger et al. 2002 A\&A 394 )
\item OH, H2O and CH3OH maser sources (Braz \& Epchtein 1983; Caswell \& Haynes 1987; Caswell 1998; Caswell et al. 1989; Caswell et al. 1995).
\item 3.4 cm radio-continuum and recombination-line measurements of NGC 3603 by De Pree, Nysewander, \& Goss (1999)
\item ATCA radio continuum at 3 and 6 cm  (M\"ucke et al 2002) of 3 proplyds
\item HST/WFPC2 observations Brandner et al 2000
\item VLT/ISAAC Data (J,H,K Band) of proplyds
\item 2MASS  data available
\item MSX observations available
\item CHANDRA observations available 
\item ISO/SWS 2.4-45.4 micron spectrum
\end{itemize}

\subsubsection{Missing Observations}
\begin{itemize}
\item CII
\item CI
\item mid-J CO
\item OI
\end{itemize}

\subsubsection{Physical Conditions}
The central cluster in NGC 3603 has a Lyman continuum Flux of 10$^{51}$ s$^{-1}$ (Kennicutt 1984; Drissen et al. 1995), about 100 times the ionizing power of the Trapezium system. The early O-type stars in NGC 3603, although more luminous than the late O-type stars in the Trapezium system, produce a disproportionate smaller rate of FUV (13.6 eV $>$ h$\nu$ $\ge$ 6 eV) photons compared to the rate of EUV photons (h$\nu$ $\ge$ 13.6 eV). The FUV flux is estimated to be $10^{52}$ s$^{-1}$. This corresponds to a Draine field of $\chi\sim  10^{6}$ at a distance of 2.4 pc. Following the nomenclature of N\"urnberger et al.(2002) the two molecular clouds closest to the central cluster are MM1 and MM2, located to the west and south-east respectively. Both of them are bordering on the wind driven cavity. 

Clump MM1 seems to be consistent with the head of the western pillar seen on optical HST images and near infrared VLT images (in contrast to M16 there is no dense molecular material detected along the ionized tail of the pillar).

 Cloud MM2 is located 2.4 pc  towards the south-east of the cluster and highly affected by the strong stellar winds and ionizing radiation. Both the high radiation and wind pressure of the OB stars may trigger the compression of the molecular gas, which is nicely confirmed by the 3.4 cm data od De Pree et al.(1999). Directly at the bottom of the eastern pillar, on the OB cluster facing side of MM2, a sparse association of IR sources (IRS9) with massive circumstellar shells is found. 

\begin{table}[ht!]
\centering
\caption{Main parameters of the two clumps}
\begin{tabular}{|c|c|}
  \hline
  % after \\: \hline or \cline{col1-col2} \cline{col3-col4} ...
   
  \hline
  Distance to ionizing source & 2.4 pc \\
  Luminosity & $2\times 10^7$ L$_\odot$ \\
  FUV flux $(13.6 eV < h\nu < 6 eV)$  & $1.2\times 10^{52} s^{-1}$ \\
  FUV field in units of Draine field &\\
  at cloud surface & $1\times 10^{6}$\\ \hline
   & R $<$ 0.4 pc\\
 MM1& M$_{vir}<0.7 \times 10^3$ M$_\odot$\\
  & $\langle N(H_2)\rangle>1.0\times 10^{23}$ cm$^{-2}$ \\ \hline
  & R = 1.0 pc\\
 MM2& M$_{vir}=1.5 \times 10^3$ M$_\odot$\\
  & $\langle N(H_2)=0.4\times 10^{23}$ cm$^{-2}$\\ 
  \hline
\end{tabular}
\end{table}

\begin{table}[ht!]
\caption{Comparison with PDR-Model results (computed with the Cologne PDR model).
 For the above parameters ($\chi=10^6$ and $M~\sim 10^3$ M$_\odot$)
the resulting column densities and radii of the Cologne PDR model are given.
\label{tab_ngc3603pdrmodel}}
\begin{center}
 \begin{tabular}{|c|c|c|c|c|c|c|c|}\hline
 &\multicolumn{7}{c|}{Log$_{10}$ n} \\
 & 3 & 3.5&4&4.5&5&5.5&6 \\ \hline
 \small $\langle$N(H$_2$)$\rangle$&2.5E22 &6.0E22 &1.3E23 &2.8E23&6.1E23 &1.3E24 &2.9E24 \\ \hline
  $\langle$N(C)$\rangle$&3.8E18 &8.8E17 &1.7E17 &2.7E16 &2.9E15 &4.2E14 &1.5E14 \\ \hline
  $\langle$N(CO)$\rangle$&3.1E18 &1.4E20 &3.4E19 &7.8E19 &1.7E20 &3.7E20 &8.0E20 \\ \hline
  $\langle$N(C$^+$)$\rangle$&7.5E17 &2.4E17 &5.8E16 &1.2E16 &2.1E15 &3.7E14 &6.9E13 \\ \hline
  $\langle$N($^{13}$CO)$\rangle$&8.1E16 &2.5E17 &5.5E17 &1.2E18 &2.6E18 &5.6E18 &1.2E19 \\ \hline
  $\langle$N(CS)$\rangle$&3.1E17 &1.8E18 &2.3E18 &1.2E18 &3.8E17 &1.1E17 &4.6E16 \\\hline
  R [pc]&1.7&1.2&0.8&0.55&0.36&0.26&0.16\\ \hline
  \end{tabular}
\end{center}

\noindent
 These values are just meant to make a quick snapshot comparison between observationally derived and theoretical values. Of course this is a very crude approach. Due to the strong stellar winds shocks may play an important role and should be considered in more detailed studies. Nevertheless the values above can be used as rough estimates.  The results from N\"urnberger et al. (2002) are 'consistent' with a surface density of $\sim 10^{3.5}$ cm$^{-3}$ for MM2 and a  higher density of $\geq 10^{4}$ cm$^{-3}$ for MM1.
\end{table}

\begin{table}[htb]
{\small
\caption{Estimated integration times in NGC3603-MM1 corresponding to a
$5\sigma$ detection.  A line width of $4.0$~km~s$^{-1}$ and a velocity
resolution of $0.7$~km~s$^{-1}$ and A=0.125 for position switched
observations was assumed for these calculations.  \label{tab_ngc3603mm1}}
\centering
\begin{tabular}{lrlrrll}
\noalign{\smallskip}\hline
  species & frequ. & column density  & $T_{B}$ & D  &T$_A$& $t_{tot}$ \\
  & [GHz] & [cm $^{-2}$] & [K] & [''] & [K]& [sec] \\
\hline
CII&        1900.545 &  2.7E+18 &      87. &      40. &      52. &      1.0\\
CH&          536.761 &  1.7E+13 &     0.38 &      40. &     0.13 &  2.6E+02\\
CH&         1656.961 &  1.7E+13 &    0.079 &      40. &    0.047 &  3.3E+04\\
CH$^+$&              835.070 &  4.8E+10 &   0.0056 &      40. &   0.0028 &  9.6E+05\\
CH$^+$&             1669.160 &  4.8E+10 &    0.012 &      40. &   0.0072 &  1.4E+06\\
NH&                  974.479 &  1.0E+13 &     0.25 &      40. &     0.13 &  4.2E+02\\
NH$^+$&             1012.524 &  8.1E+09 &  0.00054 &      40. &  0.00030 &  9.4E+07\\
NH$_3$&              572.498 &  1.6E+12 &    0.090 &      25. &    0.026 &  7.8E+03\\
NH$_3$&             1168.452 &  1.6E+12 &  0.00056 &      25. &  0.00029 &  1.0E+08\\
OH$^+$&              971.804 &  8.2E+10 &    0.031 &      40. &    0.017 &  2.7E+04\\
H$_3$O$^+$&         1655.814 &  1.6E+13 &   0.0025 &      25. &   0.0014 &  3.6E+07\\
H$_3$O$^+$&          984.697 &  1.6E+13 &    0.036 &      25. &    0.017 &  2.7E+04\\
H$_2$O&             1113.343 &  7.5E+16 &     0.77 &      25. &     0.39 &      53.\\
H$_2$O&              987.927 &  7.5E+16 &     0.25 &      25. &     0.12 &  5.7E+02\\
H$_2$O&              752.033 &  7.5E+16 &     0.32 &      25. &     0.12 &  4.0E+02\\
H$_2$O&              556.936 &  7.5E+16 &      2.5 &      25. &     0.71 &      10.\\
H$_2$O&     1669.905 &  7.5E+16 &     0.14 &      25. &    0.079 &  1.2E+04\\
H$_2$O&             1716.770 &  7.5E+16 &   0.0068 &      25. &   0.0039 &  4.3E+06\\
HDO&                 893.639 &  1.1E+12 &    0.026 &      25. &    0.012 &  4.7E+04\\
H$_2^{18}$O&                1101.698 &  1.9E+14 &     0.15 &      25. &    0.073 &  1.5E+03\\
OH&         1834.747 &  1.6E+15 &    0.083 &      40. &    0.050 &  2.5E+04\\
CO (10-9)&          1151.985 &  4.3E+18 &  6.8E-09 &      40. &  3.9E-09 & inf\\
CO (16-15)&         1841.345 &  4.3E+18 &      0.0 &      40. &      0.0 & inf\\
$^{13}$CO (10-9)&           1101.350 &  7.2E+16 &      0.0 &      40. &      0.0 & inf\\
$^{13}$CO (15-14)&          1650.768 &  7.2E+16 &      0.0 &      40. &      0.0 & inf\\
\hline
OH&         2509.988 &  1.6E+15 &      1.0 &      30. &     0.59 &  1.3E+02\\
OH&         2514.316 &  1.6E+15 &     0.89 &      30. &     0.50 &  1.8E+02\\
O&          4745.804 &  1.3E+19 &  1.2E+02 &      40. &      46. &      1.0\\
O&          2060.068 &  1.3E+19 &  1.4E+02 &      40. &      83. &      1.0\\
\hline\noalign{\smallskip}
\end{tabular}\\
\flushleft
}
\end{table}
\begin{table}[htb]
{\small
\caption{Estimated integration times in NGC3603-MM2 corresponding to a
$5\sigma$ detection.  A line width of $3.5$~km~s$^{-1}$ and a velocity
resolution of $0.7$~km~s$^{-1}$ and A=0.125 for position switched
observations was assumed for these calculations.  \label{tab_ngc3603mm2}}
\centering
\begin{tabular}{lrlrrll}
\noalign{\smallskip}\hline
  species & frequ. & column density  & $T_{B}$ & D   &T$_A$& $t_{tot}$ \\
  & [GHz] & [cm $^{-2}$] & [K] & [''] & [K]& [sec] \\
\hline
CII&        1900.545 &  2.2E+18 &      87. &      15. &      49. &      1.0\\
CH&          536.761 &  7.0E+12 &     0.19 &      40. &    0.077 &  7.8E+02\\
CH&         1656.961 &  7.0E+12 &    0.030 &      40. &    0.018 &  2.2E+05\\
CH$^+$&              835.070 &  3.2E+10 &   0.0042 &      20. &   0.0019 &  2.1E+06\\
CH$^+$&             1669.160 &  3.2E+10 &   0.0091 &      20. &   0.0053 &  2.6E+06\\
NH&                  974.479 &  4.0E+12 &    0.081 &      40. &    0.046 &  3.5E+03\\
NH$^+$&             1012.524 &  3.5E+09 &  0.00027 &      10. &  0.00010 &  7.6E+08\\
NH$_3$&              572.498 &  1.7E+11 &    0.011 &      25. &   0.0039 &  3.4E+05\\
NH$_3$&             1168.452 &  1.7E+11 &  6.8E-05 &      25. &  3.8E-05 & inf\\
OH$^+$&              971.804 &  7.5E+10 &    0.035 &      55. &    0.021 &  1.6E+04\\
H$_3$O$^+$&         1655.814 &  2.9E+12 &  0.00058 &      25. &  0.00035 &  6.1E+08\\
H$_3$O$^+$&          984.697 &  2.9E+12 &   0.0080 &      25. &   0.0042 &  4.5E+05\\
H$_2$O&             1113.343 &  2.0E+16 &     0.44 &      30. &     0.25 &  1.2E+02\\
H$_2$O&              987.927 &  2.0E+16 &     0.24 &      30. &     0.13 &  4.5E+02\\
H$_2$O&              752.033 &  2.0E+16 &     0.12 &      30. &    0.057 &  1.8E+03\\
H$_2$O&              556.936 &  2.0E+16 &      2.0 &      30. &     0.75 &      9.0\\
H$_2$O&     1669.905 &  2.0E+16 &    0.051 &      30. &    0.031 &  7.8E+04\\
H$_2$O&             1716.770 &  2.0E+16 &   0.0030 &      30. &   0.0018 &  2.0E+07\\
HDO&                 893.639 &  3.0E+11 &   0.0082 &      30. &   0.0043 &  3.5E+05\\
H$_2^{18}$O&                1101.698 &  4.9E+13 &     0.14 &      30. &    0.082 &  1.2E+03\\
OH&         1834.747 &  6.3E+14 &    0.043 &      25. &    0.026 &  9.5E+04\\
CO (10-9)&          1151.985 &  9.7E+17 &  2.7E-06 &      25. &  1.5E-06 & inf\\
CO (16-15)&         1841.345 &  9.7E+17 &      0.0 &      25. &      0.0 & inf\\
$^{13}$CO (10-9)&           1101.350 &  1.6E+16 &      0.0 &      30. &      0.0 & inf\\
$^{13}$CO (15-14)&          1650.768 &  1.6E+16 &      0.0 &      30. &      0.0 & inf\\
\hline
OH&         2509.988 &  6.3E+14 &     0.57 &      25. &     0.33 &  4.3E+02\\
OH&         2514.316 &  6.3E+14 &     0.49 &      25. &     0.28 &  5.8E+02\\
O&          4745.804 &  7.4E+18 &  1.2E+02 &      65. &      47. &      1.0\\
O&          2060.068 &  7.4E+18 &  1.4E+02 &      65. &      87. &      1.0\\
\hline\noalign{\smallskip}
\end{tabular}\\
}
\flushleft
\end{table}

\begin{figure}[!ht]
 \begin{center}
 \includegraphics[width=12cm]{proposal-cut.eps}
 \end{center}
 \caption{C$^{18}$O (2-1) data (N\"urnberger (2002)). Examples for the proposed cuts are overlayed as yellow boxes with width 40'' and 10'' respectively.}
 \end{figure}

\subsubsection{Line intensities}
 In Table \ref{tab_ngc3603mm1} and \ref{tab_ngc3603mm2} we give the estimated line intensities and the resulting integration times per position. The estimation is based on the results of spherical PDR model calculations with the Cologne PDR model. The column densities were binned in 10 temperature regions and LTE was assumed in order to calculate the intensities. 

\subsubsection{Observing strategy}
For the lines with an integration time over 10000 sec are needed, we consider that the observation is not feasible.
\paragraph{MM1}
The estimation shows, that most of the lines are too weak to observe along a cut. If we confine ourselves to observe 5 lines (CH, NH, H$_2$O 988 GHz and 752 GHz and H$_2^{18}$O) at three positions (C$^{18}$O peak, two positions $\pm$50 arcsec offset in RA) we obtain an integration time of approximately 105 minutes. Additionally the NH$_3$ line at 573 GHz should be observed at the peak position only. The estimated time for this is 130 minutes. Three lines are strong enough to propose a cut of 100'' length (20 positions with 5'' spacing). These are CII, H$_2$O 557 GHz and H$_2$O 1113 GHz. The position switching estimate for this cut is 22 minutes, or less if performed in OTF mode. The overall HIFI time for MM1 is 257 minutes. 
\paragraph{MM2}
The same observation strategy is proposed for MM2 with three differences: the cut is slightly longer (130'' , hence 26 positions), NH$_3$ is too weak to be observed at all, and the five lines given above are just observed at the MM2 peak position. The time estimates are: 57 minutes for the three line cut and 130 minutes for the peak position. The total HIFI time for MM" is 187 minutes.

\paragraph{PACS}
Corresponding cuts are proposed for the PACS lines. For MM1 this results in 3 raster points and 4 raster points for MM2. The total integration time, taking into account the smaller beam size at 63$\mu$m is 100 minutes total for both cuts and all lines.

\subsubsection{Summary}
\begin{itemize}
\item HIFI observation time of MM1, including three lines along a cut and 5 lines at two positions: 257 minutes.
\item  HIFI observation time of MM2, including three lines along a cut and 5 lines at one positions: 187 minutes.
\item  PACS line observation time, two cuts for MM1 and MM2 with a total time of 100 minutes.
\item Total observation time for NGC 3603 is 9 hours.
\end{itemize}

\subsection{MonR~2\protect\footnote{A. Fuente OAN}}

Very little is known about the chemistry of the 
PDRs associated to ultracompact (UC) \HII\ regions.
The high densities ($>$ 10$^6$ cm$^{-3}$) and kinetic temperatures 
($\sim$ 100 K) of these regions, together with the 
ionizing radiation (G$_0 > 10^5$ in units of Habing field), could drive a 
chemistry presumably different from that found in 
other PDRs. However, the PDRs associated to \uch s were
usually thought  too thin to be detected. Recent results argue
against this idea.

Mon R2, located at a distance of 850 pc, is the closest \uch\ 
to the Sun. It has an angular diameter of $\sim$ 22$''$ (0.1 pc)
(Wood \& Churchwell 1989) which is comparable to the beamwidth 
of HIFI for frequencies $>$ 1000~GHz. This provides the possibility to
spatially separate the \uch\ and the surrounding molecular cloud.
The \uch\ is highly asymmetric and has a cometary shape, 
reaching its maximum toward the exciting star Mon R2 IRS1.
The molecular emission shows an arc-like structure surrounding 
the \HII\ region, with the maximum of emission to the southeast of
IRS1. This high density arc-like structure is clearly seen
in the maps of the emission high rotational molecular lines (CO,CS,HCN,
H$_2$CO) and in the continuum emission at submillimeter wavelengths
(Giannakopoulou et al. 1997, Choi et al. 2000).

\begin{figure}[t!]
\setlength{\unitlength}{1mm}
\begin{picture}(150,80)
\special{psfile=monr2.eps
angle=0 hoffset=-20 voffset=-375 hscale=75 vscale=75}
\end{picture}
%
\caption{CS 7-6 emission map of Mon R2, taken from
Choi et al.~(2000). Equatorial coordinates are referred to 
$(\alpha, \delta)_{2000} = (6^{\rm h}07^{\rm m}46^{\rm s} 2, -6^\circ 
23' 08.3'')$. 
The size of the UC (HII) region is  indicated by the central square. 
The two positions observed in this work are marked by dots, and 
correspond to the offsets (+10'', -10'') and (0'', 0''). 
The right and left panels show the reactive ions spectra observed in this
region. }
\label{fig_monr2}
\end{figure}

\subsubsection{Previous and current work}

Recently, Rizzo et al. (2003) have detected the reactive 
ion CO$^+$ towards the peak of the \uch\  
in Mon R2 (see Fig. \ref{fig_monr2}). This ion has not been
detected in the high 
density arc of molecular emission surrounding the HII region. 
This shows that the geometry of Mon R2 is well suited to
detect the abundance gradients expected in PDRs.
Rizzo et al. (2004, unpublished data) have carried out a chemical study
of this region using the 30-m telescope. Thus far, the region has
been observed in HCN, CS, SiO,HCO, H$^{13}$CO$^+$,HC$^{18}$O$^+$,
HOC$^+$,CO$^+$, C$_2$H and c-C$_3$H$_2$. 
In order to have a more complete picture of the region we plan to
complete our study by extending our observations to other species.

\subsubsection{Physical conditions of the PDR associated to Mon R2} 
A density, $n$$\sim$ 10$^6$ cm$^{-3}$, is derived in 
Mon R2 from molecular observations  (Montalban et al. 1990, 
Giannakopoulou et al. 1997, 
Ossenkopf et al. 2001, Rizzo et al. 2003). 
Rizzo et al. (2003) have estimated the
incident UV field assuming that the observed FIR emission (L$_{\rm FIR}$) represent 
all the stellar flux, because UV, optical and IR emission is re-radiated to 
FIR.  Then a black body radiation law at the exciting star's effective temperature 
(T$_{\rm eff}$) is assumed to estimate the fraction of the luminosity radiated in
the UV range. Taking values for T$_{\rm eff}$, 
L$_{\rm FIR}$ and distance to the star of 24000 K, 3000 L$_{\odot}$ 
and 0.03 pc for Mon R2 (Henning et al.~1992), they  have obtained 
G$_0$ $\sim$ 2\ 10$^5$ for Mon R2.  The large value of G$_0$ 
and the high density make this object
an excellent candidate to be studied with HIFI.

\subsubsection{Line intensities}

The total hydrogen density in the peak of the \uch is $\sim$ 2 10$^{22}$ cm$^{-2}$. It
is difficult to estimate line intensitites because there is no chemical model published for
these physical conditions. We can do some work assuming the abundances for a
typical PDR G$_0$ = 10$^5$, n=10$^5$ cm$^{-3}$ (the same as in S106 IR) 
and using the LTE approximation. Taking into account the high densities and kinetic 
temperatures prevailing in this region, we assume a rotation temperature of 50 K.
The resulting line intensities have to be considered as upper limits because
the column density of the PDR are only a fraction of the total hydrogen column
density.

\begin{table}
\caption{MonR2 lines intensities estimates and integration times. Integration times have been estimated 
assuming $\delta$v=0.7 km/s and A=0.125 (position switching)}
\centering
\begin{tabular}{lllllr}\\ \hline
\multicolumn{1}{c}{Species} & \multicolumn{1}{c}{Freq(GHz} & \multicolumn{1}{c}{Abundance}
&\multicolumn{1}{c}{Line intensity (K)} & \multicolumn{1}{c}{t$_{int}$(s)} &
Observation type\\ \hline
 CII   & 1900.545 &   1.0E-05   &  10.9  & 0.9     &strips\\
 CH  & 536.761   &   1.5E-09   &  1.6    &   1.3   &strips\\
 CH  & 1656.961  &  1.5E-09    &  8.3    &   1.7      &strips\\
 CH$^+$ &   835.07    &    9.3E-13    &  0.007  &   $10^5$ &impossible\\
 CH$^+$   &        1669.16  &    9.3E-13     &  0.008  &  2\,$10^6$ &impossible\\
 NH          &   974.479   &    5.6E-10    &  2.9          &  0.7         &strips\\
 NH$^+$   &  1012.524   &    2.2E-13    &  0.0005     &  2\,$10^7$  &impossible\\
 NH$_3$   &  572.498     &    4.3E-11    &  0.14        &  200 &points\\
 NH$_3$    & 1168.452    &    4.3E-11    &  0.12        &  500 &points\\
 OH$^+$    &  971.804     &    5.7E-13    &  0.015      &  3\,$10^4$ &impossible\\
 H$_3$O$^+$   & 1655.814     &  1.0E-10  & 0.50        &  470  &points\\
 H$_3$O$^+$   &  984.697     &   1.0E-10  & 0.50        &   25 &points\\
 H$_2$O          & 1113.343    &   2.2E-07  &  49        &   0.003  &bonus\\
 H$_2$O          &   987.927    &   2.2E-07   & 50        &  0.003 &bonus\\
 H$_2$O          &   752.033    &   2.2E-07   &  48       &  0.003 &strips\\
 H$_2$O          &   556.936    &   2.2E-07   &  43        &  0.003 &strips\\
 H$_2$O          &  1669.905    &  2.2E-07   &  44        &  0.06  &strips\\
 H$_2$O          &  1716.770    &  2.2E-07   &  43        & 0.006 &strips\\
 HDO               &  893.639     &  1.3E-12   &  0.009 &  7\,$10^4$ &impossible\\
 H$_2$$^{18}$O  &  1101.698  & 2.E-10     &  5.3        &   0.22  &bonus\\	
 CO(10-9)           & 1151.985   & 2.6E-04    &  49        &   0.004 &strips\\
 CO(15-14)         &  1726.603   & 2.6E-04   &  0.6        &  27   &points\\
 CO(16-15)         &  1841.345   & 2.6E-04    & 0.13       &  1000 &points\\
OH                    &  1834.747   &  4.0E-10    &  0.10       &  700 &bonus\\
 $^{13}$CO(10-9)   &  1101.35   &  3.8E-06    &  5.1    &    0.26 &strips\\
 $^{13}$CO(15-14)  & 1650.768  & 3.8E-06     & 0.016 & 5\,$10^5$  &impossible\\ \hline
 OH$^a$                       &  2509.988    &  4.0E-10    &  1.9    & 150   &mapping\\
 OH$^a$                       &  2524.316     &  4.0E-10     &  1.8    & 150   &bonus\\
  O$^a$                        &   4745.804      &  1.4E-04    &  7.7   &  110 &mapping\\
  O$^a$                       &  2060.068     &  1.4E-04   &   0.6    &  420 &mapping\\ \hline
\end{tabular}
\label{tab_monr2}

\flushleft\noindent
$^a$ Assuming a linewidth of 5 km s$^{-1}$.
\end{table}

\subsubsection{Observing strategy}

For the lines where an integration time over 10000~s are needed, we consider that
the observation is not feasible. 

For all strong lines (t$_{int}$ $\lsim$ 1 sec in position-switch including the CH lines), 
it is desirable to make two strips crossing the PDR. This is the case for 12 of the
lines listed in Table \ref{tab_monr2}. Three of the lines can be considered ``bonus''
observations as they are measured simultaneously with another transition.
Using a strip length of 2.8 arcmin, the variable
beam width of HIFI at the different lines requires between 28 points (for the CII line) 
and 8 points (for CH at 537 GHz) for a fully sampled spacing. To guarantee
a clear detection we do not use the
optimistic value for the integration time of 1~s derived from the total column
density but assume that the PDR column density is only $1/10$ of the total hydrogen
column density, resulting in a standard observing time of 100~s per point
in position switch observations. When taking the higher observing efficiency of OTF
observations into account the required observing time per point is reduced to 25~s for
the long scans and to 42~s for the scans containing only a few points.
Then the two strips observed with nine frequency settings need a total observing time
of 2.8~h.

For weak lines (7 of the lines listed in Table \ref{tab_monr2}), 
two positions should be at least observed: Mon R2 IRS1 which spatially
coincides with the peak of the emission of the ionized gas and at which the reactive 
ion CO$^+$ has been detected, and at an offset (10 '',-10 '') which coincides
with the dense arc of molecular gas. At the peak positions, we expect the PDR
column density to be a relatively large fraction of the total hydrogen column density,
so that we translate the time estimates from Table \ref{tab_monr2} into
observing times of 100~s for the H$_3$O$^+$ $(0_{0,1}-1_{0,0})$ and CO(15-14) and 2000~s for all
other lines. Since the CO(16-15) and OH can be observed simultaneously, this
takes 4.6 h. The total HIFI time is thus {\bf 7.4~h}. 

Spectroscopic and photometric observations with 
PACS are also required. The proposed OH and OI lines can be detected in
a short time (see Table \ref{tab_monr2}). The region is also a strong
source in the continuum
emission with an IRAS flux of 13~MJy at 60~$\mu$m and 20~MJy at 100$\mu$m.
Due to the large minimum required PACS integration time mapping the strip
with PACS in spectroscopic mode takes at least 9~min, for the 145 $\mu$m OI
line, 35 min are necessary. For the continuum mapping 5~min are required.
This sums up to {\bf 1~h} for PACS observations so that the whole project
can be performed in {\bf 8.4~h}.
 
\subsection{S106\protect\footnote{B. Mookerjea, M. R\"ollig, KOSMA}}

\subsubsection{Introduction}

Sharpless~106 (S106) is a highly obscured bipolar HII region at a
distance of roughly 600~pc. It is powered by a single bright IR source
designated S106~IR, with a spectral type of O7-O9, a luminosity of a few
$\sim 10^4$~$L_{\odot}$, and a line of sight extinction of A$_{v} \sim
21$ mag. Around this central source Hoddap \& Rayer (1991) found a
cluster of about $160$ stars, embedded in the molecular cloud
surrounding S106.  Observed radio emission and HI recombination lines
from the region have been suggested to arise in a strong stellar wind,
driving a shock into the surrounding molecular cloud.  Just north of the
central point source S106~IR, the nebula is bifurcated into two lobes by
a prominent dark lane aligned along a nearly east-west axis. 

Integrated intensity maps of the optically thick lines reveal an
extended cloud with a peak emission at the position of S106~IR, whereas
maps of intensities of optically thin molecular lines and mm- and sub-mm
continuum show two maxima of emission separated by 3$'$ east and
west of S106~IR.

The Cologne PDR model for the position at S106~IR corresponding to an UV
field of $10^{5}$~G$_0$, density of $10^{5}$~cm$^{-3}$ and a clump mass
of $0.01$~M$_{\odot}$ reproduces some but not all of the observations of
PDR tracers (Schneider et al. 2003) reasonably well. The inability to
reproduce all observed lines (particularly the ratio of the 2 CI
lines) holds for the Kaufman (1999) model, as shown by Schneider et al.
(2003) as well. 

An independent plane-parallel PDR model (Le Bourlot 2000, Le Bourlot
2003) using the same physical parameters was used to model the observed
line intensities at the position of S106~IR. This helps us to obtain an
estimate for the modeling uncertainty of the line predictions. The
comparison of the two intensity independent estimates in Table
\ref{tab_s106}, thus gives a feeling for the uncertainty of the current
knowledge on the PDR.

%One can clearly see that for species like CH$^+$ our knowledge is very
%limited, expressed by uncertainties of several orders of magnitude
%clearly stressing the need of the corresponding HIFI observations. 

\subsubsection{Time Estimate}

For the purpose of estimating the integration time particularly for the
species like the hydrides for which there are no previous observations,
we have derived the {relative abundances} of these species from a
spherical PDR model with the above parameters. This has been done at two
positions, one coinciding with S106~IR(0,0) and an offset position
(-2',0).

{\bf S106~IR at (0,0):} The total hydrogen column density is $2.8\times
10^{22}$~cm$^{-2}$ (Schneider et al. 2002), which corresponds to a mass
of $3.4$~M$_{\odot}$. The {excitation temperatures} for each of the
species have also been estimated from the PDR models and the main beam
temperatures are calculated {assuming LTE}, using the line
intensities corresponding to the appropriate temperatures, as given in
Table 1. For a few selected lines of C{\sc II}, CO(10-9), CO(15-14),
O{\sc I} (63~$\mu$m) and O{\sc I} (145~$\mu$m) complete radiation
transfer was performed to estimate the peak intensities. The resulting
integration times per source point are summarized in Table
\ref{tab_s106}.

{\bf Offset at (-2',0):} Schneider et al. (2003) obtained an UV field of
$100$~G$_0$ and number density of $\sim$ 10$^4$ cm$^{-3}$ for this position
using the plane parallel model by Kaufman et al. (1999). We have used
the Cologne PDR model with similar parameters and a column density 10\%
of that at (0,0). In a manner similar to that for the (0,0)
position the integration times required for this position are calculated
and summarized in Table \ref{tab_s106}.


\begin{table}[!ht]
\caption{Estimated integration times in S106 corresponding to a
$5\sigma$ detection.  A line width of $5$~km~s$^{-1}$ and a velocity
resolution of $0.7$~km~s$^{-1}$ and A=0.125 for position switched
observations was assumed for these calculations.  \label{tab_s106}}
\small{
\centering
\begin{tabular}{lrlllccccc}
\noalign{\smallskip}\hline
Species & Freq. & \multicolumn{2}{c}{Relative Abundance} & \multicolumn{2}{c}{$T_B$ {}$^\dagger$
[K]} & D &$t_{int}$ \\
  & [GHz] & at (0,0) & at (-2',0) & at (0,0) & at (-2',0) & [$'$] & [s] \\
\hline\noalign{\smallskip}
CII         & 1900.545 & 1.0E-05 & 1.9E-05 & 6.5E+01 (3.6E+1) & 5.3E+00 &  32. & 1.4E-02 \\ 
CH          &  536.761 & 1.5E-09 & 1.7E-08 & 2.0E+00 & 8.3E-01 & 7.1 & 9.2E-01 \\ 
CH          & 1656.961 & 1.5E-09 & 1.7E-08 & 6.9E+00 & 1.6E+00 & 4.3 & 1.4E+00 \\ 
CH$^+$      &  835.070 & 9.3E-13 & 2.3E-12 & 2.5E-03 (1.6E-1)& 1.3E-03 &  11 & 9.2E+05 \\ 
CH$^+$      & 1669.160 & 9.3E-13 & 2.3E-12 & 8.0E-03 (2.3E-2)& 2.5E-03 &  5.2 & 1.0E+06 \\ 
NH          &  974.479 & 5.6E-10 & 4.4E-10 & 2.2E+00 & 1.7E-01 &  3.1 & 1.4E+00 \\ 
NH$^+$      & 1012.524 & 2.2E-13 & 1.9E-15 & 4.5E-04 & 5.3E-07 &  2.4 & 3.4E+07 \\ 
NH$_3$      &  572.498 & 4.3E-11 & 3.1E-19 & 2.1E-01 & 5.6E-11 &  2.1 & 8.9E+01 \\ 
NH$_3$      & 1168.452 & 4.3E-11 & 3.1E-19 & 4.3E-03 & 5.5E-11 &  2.1 & 4.3E+05 \\ 
OH$^+$s     &  971.804 & 5.7E-13 & 1.3E-14 & 5.5E-03 & 2.5E-05 &  2.5 & 2.2E+05 \\ 
H$_3$O$^+$  & 1655.814 & 1.0E-10 & 6.7E-11 & 1.4E-02 & 1.9E-02 &  0.1 & 3.4E+05 \\ 
H$_3$O$^+$  &  984.697 & 1.0E-10 & 6.7E-11 & 4.8E-02 & 1.8E-02 &  6.2 & 2.9E+03 \\ 
H$_2$O      & 1113.343 & 2.2E-07 & 2.2E-08 & 1.6E+00 (2.0E-1)& 1.6E+00 &   & 3.0E+00 \\ 
H$_2$O      &  987.927 & 2.2E-07 & 2.2E-08 & 1.8E+00 (2.8E-2)& 3.9E-02 & 2.6 & 2.1E+00 \\ 
H$_2$O      &  752.033 & 2.2E-07 & 2.2E-08 & 1.6E-01 & 1.6E-03 & 2.6  & 2.0E+02 \\ 
H$_2$O      &  556.936 & 2.2E-07 & 2.2E-08 & 5.4E+00 (1.0E+0)& 4.9E+00 & $>15$ & 1.3E-01\\ 
H$_2$O      & 1669.905 & 2.2E-07 & 2.2E-08 & 3.8E-01 & 3.5E-01 & $>15$ & 4.5E+02 \\ 
HDO         &  893.639 & 1.3E-12 & 1.4E-12 & 1.1E-02 & 1.3E-04 & 2.6  & 5.1E+04 \\ 
H$_2^{18}$O & 1101.698 & 2.0E-10 & 2.1E-10 & 1.5E+00 & 2.0E-01 & 3.5  & 3.4E+00 \\ 
CO(10-9)    & 1151.985 & 2.6E-04 & 2.5E-04 & 1.4E+01 (2.1E+0) & 1.2E+1 &  2.5 & 1.0E+00 \\ 
CO(15-14)   & 1726.603 & 2.6E-04 & 2.5E-04 & 9.2E+00 (4.8E-4) & 1.6E+00 &  2.4 & 1.0E+00 \\ 
13CO(10-9)  & 1101.350 & 3.8E-06 & 3.9E-06 & 4.0E+00 & 4.5E-01 & 3.2  & 1.0E+00\\ 
13CO(15-14) & 1650.768 & 3.8E-06 & 3.9E-06 & 3.0E-01 & 3.0E-02 & 3.2 & 7.0E+02\\
\hline
OH          & 2509.988 & 4.0E-10 & 1.0E-10 & 1.4E-02 & 8.3E-04 & 3.0  & 5.4E+00 \\ 
OH          & 2514.316 & 4.0E-10 & 1.0E-10 & 1.4E-02 & 5.2E-04 & 2.8  & 5.8E+00 \\ 
O           & 4745.804 & 1.4E-04 & 3.5E-04 & 3.8E+00 (1.5E+2) & 9.4E-03 &  2.4 & 1.1E-05 \\ 
O           & 2060.068 & 1.4E-04 & 3.5E-04 & 2.4E+00 (5.6E+1)& 5.1E-03 &  2.4 & 4.9E-04 \\ 
\hline\noalign{\smallskip}
\end{tabular}\\
\flushleft
For the C{\sc II}, CO(10-9), CO(15-14) and O lines the peak intensities
were calculated using a full 1-D radiation transfer code.\\
For all other species, LTE intensities corresponding to the temperatures
derived from the PDR models have been used.\\
The peak intensity given here translates into line integrated
intensities by $\int T dv= 2\pi/(8 {\rm ln}2) T\sub{peak} \Delta v$,
where $\Delta v$ is 5 km/s.
$D$ is the distance from S106 IR (0,0) at which the line intensity falls
to 20\% of it's value at (0,0). It was calculated by interpolating or
extrapolating the results from model calculations for the positions
(0,0) and (-2',0) linearly.\\
$\dagger$ The numbers in brackets are results of calculations using
 Le Bourlot's PDR model.
 }
\end{table}

\subsubsection{Observing Strategy}

\paragraph{HIFI}

We propose an observation sequence including :
\begin{enumerate}
  
\item Pointed mode HIFI observation of S106~IR (0,0) requiring
  1\,hours. Here, we discarded the species for which the integration
  time exceeds 2.7\,hours.
  
\item A HIFI cut of $5'$ length. In Table\,\ref{tab_s106}, we list at
  which distance to the 0/0 position, the line strengths are expected
  to drop to 20\%. For the cut, we only select sources which are
  detectable within less than 5\,seconds of integration time at the
  0/0 position. It will take 2\,minutes per position along the cut to
  detect lines of 20\% intensity with the same SNR.  Nine of the
  species/transitions fulfill this criterium, among them CII, CH, and
  several water lines. Here, we assume that 2 parallel cuts will be
  made. For the time estimate, we further assume that cuts will be
  sampled every $10''$ (60\,positions), though in practice this will
  depend on the beam size, i.e. frequency, to obtain full-sampling.
  This results in 2\,hours of integration per line. The on-the-fly 
  is a factor of 4 more effective than position-switching, thus 1
  hours is sufficient. For the nine  lines, we thus ask for 9 hours
  of observing time.

\end{enumerate}

A total of {\bf $10$~hours} ($1+9$) will be required to observe the
  full sequence in S106 with HIFI.
 
\subsubsection{PACS time estimate}

We plan to observe the four OH and OI lines along the same cuts
planned for HIFI. As the extent of the emission in those lines is
expected to be small, only the inner 5$'$ are to be covered.  We
expect the PACS observing times to be entirely driven by the
relatively large overheads. To map the OI 63$\,\mu$m line at beam
spacing (HPBW=$4.7''$), we need 37\,min using the PACS Guesstimator
(21/10/02). Here, we take into account that the 4 settings are needed
to half-sample the data due to the PACS pixel size of $9.4''$.
Complementary observations of the weaker, but optically thin OI
145\,$\mu$m line (HPBW=$10.8''$) at beam spacing need about a factor
of 4 less time, 12\,minutes. A similar amount is need for each OH
line. We thus ask for 1.2\,hours of PACS time for all four lines.

{\bf For HIFI and PACS, we thus ask for a total of 11.2\,hours. }

\subsection{The Carina Molecular Cloud Complex (NGC3372)\protect\footnote{C. Kramer (KOSMA)}}

\subsubsection{Introduction}
 
The Carina Nebula is the nearest-by example of very massive star
formation.  In this very prominent southern star-forming region at a
distance of 2.5\,kpc, dozens of OB0-stars produce a large UV field
which, together with stellar winds, interact heavily with the material
in its parental giant molecular cloud of about $5\,10^5\,M_{\rm sun}$.
In a region about 40\,pc in diameter there are 64 O-type stars
(including 6 of the only 11 O3-type stars known in the Milky Way.)
This high concentration of the earliest-type stars is unique in the
Galaxy. The nearest region of higher concentration surrounds the
30\,Doradus region in the LMC.  The nebula extends more than 4 square
degrees on the sky between 284 and 289 degrees Galactic longitude and
between $-2$ and 1 degrees Galactic latitude. It presents a peculiar
morphological structure on all scales, clumps, filaments, arcs, and
all sorts of chemical, kinematical, and density inhomogeneities.
Massive stars are being born in several regions within molecular
condensations, located mainly to the SE and NW of the cluster Trumpler
16 (Tr\,16) which is centered in $\eta$Carinae. In the northern part
of the Carina nebula, the stellar population is dominated Tr\,14
(about $10'$ to the north of Tr\,16).  Tapia et al.  (2003) find
evidences for ionization fronts leading from Tr\,14 into the molecular
cloud, which appears to be wrapping the Tr\,14 cluster.

\subsubsection{An overview of the complete region:
CI and CO 4--3 at $3'$ resolution}

A complete survey of the $\sim150\times100$\,pc region in CO 4--3 and
CI 1--0 at a resolution of about $3'$ was presented by Zhang et al.
(2001). This AST/RO data shows approximately coextensive emission of
both tracers, and moreover large scale CI emission, also arising 
far away from UV sources as comparison with the 100$\mu$m HIRES map
shows. The observations are also compared to the high-resolution CO
1-0 MOPRA study of Brooks et al. (1998) allowing to derive estimates
of the excitation temperatures. For their ``region 2'' which covers
the Carina nebula HII region (NGC\,3372), the average peak $T_{\rm mb}$ line
temperatures are 1\,K for CO 4--3 and 0.2\,K for CI 1--0.  Average
  derived physical parameters are: $\chi=1700\,\chi_0$ ($\chi_0$ is
  the Habing field), $T_{\rm ex}=52\,$K, $n\sim10^5$\,cm$^{-3}$.

\subsubsection{Morphological components of the Carina clouds}

Going from north to south, three regions can be identified: the
northern molecular cloud, the keyhole, and the southern molecular
cloud.


\paragraph{Northern molecular cloud.}

The Northern molecular cloud contains more sites of massive star
formation than the southern cloud or the keyhole. All young stars
identified are associated with UC HII regions.  Diffuse PAH
3.29$\,\mu$m emission is found by Rathborne et al. (2002) following
the sharp edge of the northern molecular cloud. Peak intensities at
the edge towards the Car I HII region are almost 10 times higher than
in the other regions of Carina. Car I-E is interacting with the front
face of the GMC and creating a PDR seen edge-on, as is consistent with
the model proposed by Brooks et al. (2001, 2003). Car I-W is creating
a second PDR seen face-on.

A PDR analysis using the Kaufman model was recently conducted by
Brooks et al. (2003) at the position of peak CII emission. CII had
been mapped with the KAO over a $13'\times11'$ region covering Tr\,14.
This data is combined with OI $63\,\mu$m and $145\,\mu$m data from
ISO/LWS, and CO 2--1, CO 1--0 data.  From this modeling, Brooks
  et al. derive a density of $\sim10^4$\,cm$^{-3}$ and a FUV field of
  $G\ge10^4\,G_0$. The FUV field derived from the PDR analysis is
  consistent with that derived from the stellar content of Tr\,14
  assuming that the PDR lies at a distance of 2.6\,pc to the center of
  Tr\,14. It is also consistent with the IRAS $60\,\mu$m and
  $100\,\mu$m colors. Brooks et al. give an H$_2$ column density of
  $8\,10^{21}$\,cm$^{-2}$ at this position from the $-25$ to
  $-22$\,kms$^{-1}$ velocity component of $^{13}$CO.  
  
  Brooks et al. (2003, Fig.11) present a $15'$ cut from the Tr\,14
  center to the PDR in emission of dust at 100, 60, 25, and
  12\,$\mu$m, together with emission data of 3.29$\mu$m PAH, $^{13}$CO
  2--1, CI 1--0, CII, 4.8\,GHz radio continuum. Another nearby $7'$ cut
  shows the variation of CII, OI 63, OI 145, OIII 88 emission observed
  with ISO/LWS (Fig.4).

\paragraph{Keyhole nebula and $\eta$\,Car.} 

The molecular gas in the Keyhole nebula is highly structured,
consisting of several small clumps (size $10-30''$) separated both in
space and in velocity. Observations by Cox \& Bronfman (1995) at
$24''$ resolution using SEST indicates that the clouds in the keyhole
are centrally condensed, externally heated molecular globules with a
small ($<0.3$pc, $M=1-10$) and dense core ($10^5$ from CS),
surrounded by a lower density, warm envelope traced in CO and
$^{13}$CO.  Ratios of CO 2--1/1--0 reach 2.4 in clump 7 (Rathborne et
al. 2002), significantly higher than 1 expected for thermalized gas,
thus indicating external heating. The dynamics of the gas is highly
perturbed due to the exposure to the stellar winds of the massive
stars belonging to Tr\,16 including $\eta$\,Car. PAH features at
$3.29\,\mu$m are found by Rathborne et al. (2002), also in two clumps
previously not seen. And strong H$_2$ ($2.12\,\mu$m) observations is
found by Brooks (2000).  These tracers are indicative of PDRs on the
surfaces of the clumps.  Brooks (2000) has developed a geometrical
model of the keyhole nebula.  Estimates of the winds and radiation
from $\eta$Car in particular suggest that these clumps are the
photo-evaporating remains of the natal cloud.

\paragraph{Southern molecular cloud and IRAS\,$10430-5931$.}

At $25''$ resolution, the CO 2--1 map of Cox et al. (1995) shows a
sharp rim of gas lying to the south of Eta Carinae and the Tr\,16
cluster. Several IRAS sources are located along the edge, a prominent
peak in PAH 3.29$\mu$m emission (Rathborne et al. 2002) is
IRAS\,$10430-5931$.

%%\subsubsection{What is missing?}
\subsubsection{Preparatory work}

\paragraph{CI and mid-$J$ CO data.} 
The 10m APEX telescope is ideally suited to map this bright, southern
source. Supplementary large-scale maps with the 4m NANTEN2 telescope
from the same site are planned. CI data at sub-arcminute resolutions
is missing, the upper transitions has not yet been observed. Mid-$J$
CO sub-arcminute resolution observations are also missing entirely.

\paragraph{FIR line data.} 
Brooks et al. (2003) present an integrated intensity map of CII
covering Tr\,14 in addition to ISO/LWS line integrated intensities at
5 positions in the vicinity of Tr\,14.  Mizutani et al.  (2002) have
ISO/LWS data of CII, OI lines, NII, etc. which were presented at the
ISO workshop (2000) but are otherwise not yet published.

\begin{table}[!ht]
\caption{Physical properties of the three different regions.
Properties along northern cut were derived by Brooks et al. (2001). Densities
of up to $10^6$\,cm$^{-3}$ are indicated by multiline CS analysis.
Those of the cut through the keyhole nebula stem from Cox \& Bronfman (1995) and
Rathborne et al. (2002). The parameters for the Southern globule IRAS 10430-5931
were taken from Megeath et al. (1996).
} \label{tab-carina}
\centering
\begin{tabular}{lccc}
\noalign{\smallskip} \hline \noalign{\smallskip}
Region in Carina                              & $N^{\rm pk}$(H$_2$) 
                                                           & $n$ & UV-Field \\
                                              & [cm$^{-2}$]& [cm$^{-3}$] & Draine \\
\noalign{\smallskip} \hline \noalign{\smallskip}
Cut through Northern molecular cloud          & $8\,10^{21}$ & $10^{4-6}$ & $10^4$ \\
Cut through Keyhole nebula, and $\eta$Car & $\sim3\,10^{21}$ 
                                                            & $10^{3-5}$ & $>>10^5$(?) \\
Southern molecular cloud and IRAS\,10430-5931 & $2\,10^{22}$ & $5\,10^3$ & $\sim10^5$(?)\\
\noalign{\smallskip} \hline \noalign{\smallskip}
\end{tabular}\\
\flushleft
\end{table}

\begin{figure}[htb]
\centering
\includegraphics[angle=0,width=0.9\linewidth]{cuts_in_carina_brooks1998.eps}
\caption{
  A red wavelength optical image overlayed with contours of $^{13}$CO
  2--1 emission ($1.7'$ resolution using Mopra) of the inner Carina
  region (taken from Brooks et al.  1998). The most important regions
  discussed in the text are labelled. Straight lines sketch the cuts
  proposed for HIFI and PACS.
\label{fig-carina}
}
%% soffice of fig13.gif of Brooks et al. 98
\end{figure}

\subsubsection{HIFI/PACS proposal}

\paragraph{Motivation for three cuts}

We propose to conduct three cuts to sample some of the different
environments found in Carina (Fig.\,\ref{fig-carina},
Table\,\ref{tab-carina}).
  
\begin{enumerate}
\addtolength{\itemindent}{-0.5cm}
  
\item {\bf Northern molecular cloud.} A $10'$ cut is proposed in the
  northern Carina cloud, from Tr\,14 to the edge-on PDR at the outer
  part of the Car\,I-E GMC, similar to the cut presented by Brooks et
  al.  (2003, cf. Fig.11 therein). Cf. also with Figure\,11 of
  Rathborne et al. (2002) showing the region in $39\,\mu$m, 21, and
  8$\,\mu$m IR emission observed with MSX.
  
\item {\bf Keyhole nebula and $\eta$Car.} To study the very harsh
  environment in the vicinity of $\eta$Car, we also propose to conduct
  a $2'$ cut from $\eta$Car through one of the small cloudlets, e.g. through
  the small CO clump 7 (Rathborne et al.  2002) showing unusually high 
  CO 2-1/1-0 ratios.
  
\item {\bf Southern molecular cloud and IRAS\,10430-5932.} We propose
  to conduct a $5'$ cut from $\eta$Car into the southern molecular
  cloud edge at the position of the globule IRAS\,$10430-5931$
  (Rathborne et al. 2002, Fig.7) which is externally heated by
  $\eta$Car and Tr\,16 and probably internally heated by deeply
  embedded pre-main-sequence stars (Megeath et al. 1996).
\addtolength{\itemindent}{0.5cm}
\end{enumerate}

\begin{table}[!ht]
\caption{HIFI/PACS time estimate for the three envisaged cuts shown in the
above Figure.
} \label{tab-carina-time-estimate}
\centering
\begin{tabular}{lrrrr}
\noalign{\smallskip} \hline \noalign{\smallskip}
Region in Carina                              & Length & PACS-time & HIFI-time \\
                                              & [$'$]  & [hours]   & [hours] \\
\noalign{\smallskip} \hline \noalign{\smallskip}
Cut through Northern molecular cloud          & 10 & 2.1 & 5.0 \\
Cut through Clump7, Keyhole nebula, and $\eta$Car 
                                              &  2 & 0.6 & 2.0 \\
Southern molecular cloud and IRAS\,10430-5931 &  5 & 1.2 & 2.5 \\
\noalign{\smallskip} \hline \noalign{\smallskip}
Total time                                    &    & 3.9 & 9.5 \\
\end{tabular}\\
\flushleft
\end{table}

\paragraph{PACS time estimate}

We plan to observe the four OH and OI lines along the same cuts
planned for HIFI.  We expect the PACS observing times to be entirely
driven by the relatively large overheads. For the time estimate, we
assume that the OI 63$\,\mu$m line is mapped at beam spacing
(HPBW=$4.7''$). Here, we take into account that the 4 settings are
needed to half-sample the data due to the PACS pixel size of $9.4''$.
Complementary observations of the weaker, but optically thin OI
145\,$\mu$m line (HPBW=$10.8''$) at beam spacing need about a factor
of 4 less time. The same amount is need for each OH line.  The
resulting observing times are summarized in
Table\,\ref{tab-carina-time-estimate}.

\paragraph{HIFI time estimate}

The peak column density and radiation field along the Northern cut are
similar to the properties of Cepheus\,B. We can thus follow the method
used for Cepheus\,B to estimate HIFI times. A $10'$ on-the-fly cut
will take $5\,$hours with an integration time of 10\,minutes per
position every $10''$. This should allow to detect 12 species and
transitions, most notably CII, both CH lines, NH, all water lines, and
the CO and $^{13}$CO 10--9 transitions.

Peak column densities and the impinging radiation field are expected to
be higher along the Southern $5'$ cut. Here, we conservatively assume
the above properties and estimate a HIFI observing time of 2.5\,hours.

For the short $2'$ cut through one of the globules of the keyhole
nebular we expect significantly lower column densities of $10^{21}$,
i.e. $\sim1\,A_V$, though also much higher FUV fields.  Here, we use
the above abundances and kinetic temperatures, and find that
20\,minutes per position are needed to observe again 12\,species.
Twelve cuts at the twelve frequencies will thus take 2\,hours.

{\bf We ask for a total time of 13.4\,hours to observe three cuts in Carina.}

\subsection{Cepheus B\protect\footnote{C.\,Kramer, B.\,Mookerjea (KOSMA)}}

\subsubsection{Introduction}
 
The molecular cloud Cepheus B, at a distance of 730\,pc (Blaauw 1964),
is located at the edge of the HII region S155 and the Cepheus OB3
association. In the immediate vicinity of the interface region there
is a compact blister type HII region, a hot core, and a small cluster
of young stars (Testi et al. 1995, Olmi \& Felli 1998, Moreno-Corral
et al. 1993). Minchin et al. (1992) presented a CO 3-2 map covering
the core region at high spatial resolutions using the JCMT.  A CS
study at similarly high resolutions was presented by Olmi \& Felli
(1998). An overview and a first PDR analysis based on CO data was
presented by Beuther et al.  (2000) using KOSMA CO maps covering a
region of $15'\times15'$.  The PDR analysis comprises CO 2--1 and 3--2
data at 4 positions and includes isotopomeric $^{12}$CO, $^{13}$CO,
and C$^{18}$O line intensities.

\begin{table}[!ht]
\caption{PDR modeling results and observed properties
for the hot core position in Cepheus\,B 
(cf. Table 7 of Beuther et al. 2000).
% The average density is the clump density times the volume filling factor.
% The area filling factor is the number of clumps per total area times
% the clump projected area. 
} \label{tab-pdr1-cepb}
\centering
\begin{tabular}{ccc}
\noalign{\smallskip} \hline \noalign{\smallskip}
Clump density             & $6\,10^4$ & [cm$^{-3}$] \\
Clump radius              & 0.002 & [pc] \\
% Clump mass                & $10^{-2}$ & [$M_{\odot}$] \\
% Interclump medium density & 100 -- 200 & [cm$^{-3}$] \\
% average density           & $1.5\,10^3$ & [cm$^{-3}$] \\
area filling factor       & $>3$ & \\
volume filling factor     & $0.04$ & \\
\noalign{\smallskip} \hline \noalign{\smallskip}
UV-field                  & $10^4$ & Draine \\
H$_2$ column density & 5 & [$10^{21}$\,cm$^{-2}$] \\
Typical linewidth & 2.8 & [kms$^{-1}$] \\
\noalign{\smallskip} \hline \noalign{\smallskip}
\end{tabular}\\
\flushleft
\end{table}

The primary sources of UV radiation in the vicinity of the Cepheus B
molecular cloud are the O7 star HD217086 and the B1 star HD217061, at
projected distances to the hot core of $\sim1.2$\,pc and 0.7\,pc.

Table\,\ref{tab-pdr1-cepb} summarizes the results of the PDR analysis
at the hot core position (Beuther et al. 2000).
%% this is in units of the Draine field. 
%% In the Beuther paper we used: G=5.6\,10^3\,G_0 Habing field
Column densities rise by a factor of $\sim4$ deeper within the cloud.
The analysis indicates that the medium is highly clumped with very low
clump volume filling factors. The unusually high CO 3--2/2--1 ratios
in the interface region cannot be readily explained by standard PDR
models. First KOSMA CI 1--0 observations at 2 positions indicate a
CI/CO abundance ratio of 20 to 30\%.

\subsubsection{Ground based preparatory work}

$^{13}$CO and C$^{18}$O 2--1 data at $12''$ resolutions have been
observed at the IRAM 30m telescope by Ungerechts et al., covering
about the same region as the KOSMA observations.

CO 7--6 data have been mapped by C. Kramer at the 10m HHT. CI 1--0 and
CO 4--3 data have been mapped with CHAMP/CSO by D. Muders and are
currently being processed. Both these data sets are at $\sim10''$
resolutions and will thus allow a detailed comparison with future CII
HIFI data at a similar resolution. Additional data of the $^{13}$CO
6-5 line will be obtained at the JCMT to allow a better estimate of
the column densities of the warm gas component.

To obtain an overview of the region, B. Mookerjea et al. have started
to map Cepheus B with SMART/KOSMA in CI 1--0, 2--1, CO 4--3, 7--6, and
$^{13}$CO 8--7. These data will be analyzed with the help of PDR
models similar to what has been done in S106 (Schneider et al. 2003)
and W3\,Main (Kramer et al.  2004).

Preparatory ammonia observations of the inversion lines will be
conducted at the 100m Effelsberg telescope. These will allow to
derive kinetic gas temperatures and ammonia abundances.

FIR fine structure data of CII, OI, etc. are yet missing entirely.
Also missing are e.g. H$_2$ data of the hot core region.

\begin{table}
\caption{Lines intensities estimates and integration times for a 
3$\sigma$ detection at the hot core position. Due to its similar
radiation field and densities, we here used the abundances compiled 
for the S106 0,0 position. Integration 
times have been estimated assuming LTE at the kinetic temperatures listed
in column (4) and a velocity resolution of $\Delta$v=0.7 kms$^{-1}$, 
a line FWHM of 2.8 kms$^{-1}$, and frequency switched observations. 
Species needing more than 2 hours of integration time are labelled 
``weak'' and are discarded from the time estimate.
\label{tab_cepb_hot_core}
}
\centering
\begin{tabular}{lllrrrl}\\ 
\hline\hline
  species & frequ. & abund. & $T_{\rm kin}$ &  $T_B$ & $t_{tot}$ \\ 
  & [GHz] & & [K] & [K] & [sec] \\ 
  (1) & (2) & (3) & (4) & (5) & (6) & (7) \\ 
\hline
 CII         & 1900.545 & 0.1E-04 &  100 &     0.906 & 30 &  ok    \\ 
 CH          &  536.761 & 0.2E-08 &   50 &     0.131 & 80 &  ok    \\ 
 CH          & 1656.961 & 0.2E-08 &   50 &     0.436 & 100 &  ok    \\ 
 CH$^+$      &  835.070 & 0.2E-12 &  100 &     0.000 & 0.3E+09 &  weak  \\ 
 CH$^+$      & 1669.160 & 0.9E-12 &  100 &     0.001 & 0.6E+08 &  weak  \\ 
 NH          &  974.479 & 0.6E-09 &   50 &     0.161 & 90 &  ok    \\ 
 NH$^+$      & 1012.524 & 0.2E-12 &  100 &     0.000 & 0.2E+10 &  weak  \\ 
 NH$_3$         &  572.498 & 0.4E-10 &   50 &     0.011 & 2.7\,h &  weak  \\ 
 NH$_3$         & 1168.452 & 0.4E-10 &   50 &     0.006 & 0.9E+05 &  weak  \\ 
 OH$^+$      &  971.804 & 0.6E-12 &  100 &     0.001 & 0.4E+07 &  weak  \\ 
 H$_3$O$^+$     & 1655.814 & 0.1E-09 &   50 &     0.021 & 17\,h &  weak  \\ 
 H$_3$O$^+$     &  984.697 & 0.1E-09 &   50 &     0.027 & 0.8\,h &  ok    \\ 
 H$_2$O         & 1113.343 & 0.2E-06 &   50 &     9.394 & 1 &  ok    \\ 
 H$_2$O         &  987.927 & 0.2E-06 &   50 &    10.065 & 1 &  ok    \\ 
 H$_2$O         &  752.033 & 0.2E-06 &   50 &    11.407 & 1 &  ok    \\ 
 H$_2$O         &  556.936 & 0.2E-06 &   50 &    12.750 & 1 &  ok    \\ 
 H$_2$O         & 1669.905 & 0.2E-06 &   50 &     6.710 & 1 &  ok    \\ 
 HDO         &  893.639 & 0.1E-11 &   50 &     0.001 & 0.9E+07 &  weak  \\ 
 H$_2^{18}$O       & 1101.698 & 0.2E-09 &   50 &     0.282 & 30 &  ok    \\ 
 CO 10-9     & 1151.985 & 0.3E-03 &   50 &     7.381 & 1 &  ok    \\ 
 CO 15-14    & 1726.603 & 0.3E-03 &   50 &     0.026 & 8.3\,h &  weak  \\ 
 $^{13}$CO 10-9   & 1101.350 & 0.4E-05 &   50 &     0.272 & 40 &  ok    \\ 
 $^{13}$CO 15-14  & 1650.768 & 0.4E-05 &   50 &     0.001 & 0.6E+08 &  weak  \\ 
\hline
\end{tabular}
\end{table}

\subsubsection{HIFI and PACS observing strategy}

Here, we propose to conduct a cut from the B1 star to the cloud
interface, through the hot core, and deeper into the more
quiescent molecular cloud. This cut is based on our CO 7--6 HHT data, CO
2--1 IRAM/30m data, both at $\sim10''$ resolutions, and on the CO 3--2
KOSMA data (cf. Fig.1 in Beuther et al. 2000).

The length of the cut shall be $10'$.

\subsubsection{PACS time estimate}

We expect the PACS observing times to be entirely driven by the
relatively large overheads. To map the OI 63$\,\mu$m line at beam
spacing (HPBW=$4.7''$), we need 1.1\,hours using the PACS Guesstimator
(21/10/02). Complementary observations of the weaker, but optically
thin OI 145\,$\mu$m line (HPBW=$10.8''$) at beam spacing, will take
20\,minutes. We thus ask for 1.5 hours of PACS time.

\subsubsection{HIFI time estimate}

We used the LTE approximation assuming the kinetic temperatures listed
in column (4) of Table\,\ref{tab_cepb_hot_core}): 100\,K for the
species which are presumably be found in the hot outer PDR surface
layers and 50\,K for the other species from the warm gas further
inside the PDR. Deeper inside the cloud the gas is known to be colder,
but this gas component will not be traced by the HIFI
submillimeter/FIR transitions.

\begin{enumerate}
\addtolength{\itemindent}{-0.5cm}
  
\item We propose to conduct deep integrations at the hot core position
  to derive a chemical census of the gas in this zone
  (Table\,\ref{tab_cepb_hot_core}). This will take 1 hour, allowing to
  detect 13 of the species and transitions listed.
  
  Amongst the strongest species/transitions are H$_2$O 557\,GHz, CO
  10--9, and CII. In this case, more than 80\% of the observation time
  needs to be spend to detect the H$_3$O$^+$ line at 984\,GHz - which
  an important molecule in the chemical network. Detection of the 
  ground state ammonia transition would require almost 3\,hours.
  We instead propose to study the ammonia abundance and derive 
  kinetic temperatures from radio observations.
  
  It will be impossible to detect the ionic molecules CH$^+$, NH$^+$,
  OH$^+$. The CO 15--14 line will not be sufficiently excited to be
  detectable at the temperatures prevailing in the Cepheus\,B hot
  core.

  Note that the kinetic temperatures in the emitting regions may be
  higher than assumed here. Also, CII abundances are predicted to be
  higher by up to a factor of 5 depending on which PDR model is used.
  Thus, the predicted intensities are conservative estimates.
  
\item We plan to conduct one on-the-fly cut with an integration time
  of 10\,minutes per position, allowing to trace the variation of the
  chemical and physical conditions with high signal-to-noise ratios.
  From the CO analysis of Beuther et al. (2000) we know that the
  column densities (per $2'$ beam) vary by a factor of $\pm5$ along
  this cut. The ground-based preparatory work will show whether the
  variation is higher at higher angular resolutions. The kinetic
  temperatures are expected to drop when going deeper into the cloud
  further away from the exciting stars and the hot core. Here, we
  simply assume that the time estimate presented in
  Table\,\ref{tab_cepb_hot_core} also holds roughly for the other
  positions. Relative to the longer integration at the hot core
  position, only the H$_3$O$^+$ 984 GHz line will not be detected.
  
  It will take 5\,hours to sample 60\,positions at $10''$ along the
  cut.  We here take into account that the on-the-fly mode is a factor
  2 more efficient than frequency-switching, the mode on which
  Table\,\ref{tab_cepb_hot_core} is based to estimate the observing
  time at the hot-core position.
\addtolength{\itemindent}{0.5cm}
\end{enumerate}

\vspace*{1cm}

The overall strategy for HIFI/PACS will be reevaluated once more
ground-based data has been analyzed.  {\bf We thus ask for a total of
  $1+5$\,hours HIFI time plus 1.5\,hours PACS time, in total
  7.5\,hours.}

\subsection{NGC7023\protect\footnote{C. Joblin, CESR, M. Gerin, ENS, V. Ossenkopf, KOSMA}}

\subsubsection{Introduction}

NGC 7023 (vdB 139) is a bright reflection nebula illuminated by the 
pre-main-sequence B3Ve star HD200775. The more recent parameters for
the star are given by van den Ancker et al. (1998):
distance: 430\,pc (-90,+160), temperature: 10$^{4.31}$ and
luminosity 10$^{3.89}$ (-0.21, +0.26).
Several observations have shown that NGC 7023 is an interesting
region from a dynamical point of view, being
considered as an archetype of the interaction of a young
star of intermediate mass with its parent cloud.
There is a molecular cavity around
the star with an hourglass shape, suggesting a bipolar origin.
The cavity is also deficient in atomic gas, the largest HI
column density being located at $\sim$40" NW from the star
(Rogers et al. 1995). Fuente et al. (1998a) showed that the HI
emission is formed by two intense filaments of length $\sim$1\,pc
and adjacent to the walls of the eastern lobe of the cavity.
The HI gas is therefore concentrated along the molecular walls.

The walls of the cavity are traced by
UV fluorescent lines from H$_2$ (Lemaire et al. 1996) and PAHs
(Cesarsky et al. 1996) forming two large PDR interfaces,
nearly edge-on at the North and at the South.
Narrow band imaging in the 1-0 S(1) and 2-1 S(1) H$_2$ lines
(Lemaire et al. 1996) show that excited H$_2$ is located
in bright filaments also observed in the continuum at
2.1\,$\mu$m by Sellgren, Werner \& Dinerstein (1992).
Interferometric observations of the J=1-0 HCO$^+$ line
 by Fuente et al. (1996) reveal the same
filaments. They are located $\sim$ 50" NW and $\sim$ 70" SE
from the star and constitute the peaks of the [CII] 157.7\,$\mu$m, 
[OI] 63.2 and 145.6\,$\mu$m, [HI] 21\,cm and H$_2$ rotational lines
emission. They trace gas at densities up
to 10$^5$-10$^6$ cm$^{-3}$ embedded in a more diffuse medium,
n $\sim$ 10$^4$ cm$^{-3}$ (Fuente et al. 2000).

Lemaire et al. (1999) studied the velocity and linewidth of the
H$_2$ IR emission lines in comparison with data in CII, CI and CO
(Gerin et al. 1998; Fuente et al. 1998a) and in HI (Fuente et al.
1996). A velocity gradient from the atomic gas to the
molecular gas is clearly present. Lemaire et al. also suggested
that the density enhancement at the border of the molecular cloud
may be due to an expanding HII region but this has
been found inconsistent with current estimates of the age of
HD200775 which is in excess of 10$^6$ years (Fuente et al. 1998b).
Another possibility would be to invoke the winds driven by the
young AeBe Herbig star HD200775 (van den Ancker et al. 1998).
A dynamical behavior of the PDR (advancing dissociation
front) was also suggested by Fuente et al. (1999) to account
for the observed out-of-equilibrium ortho- to para-H2
line ratios. A kinematical view of the gas around HD200775
is therefore of great interest to disentangle the various
scenarii.

Several radio and far-IR observations of the region are available.
The north and south PDRs were mapped in CO, CI and CII by Gerin et
al. (1998). Fuente et al. (1993) studied the chemistry in NGC 7023
by mapping several molecular transitions $^{12}$CO, $^{13}$CO,
C$^{18}$O, HCO$^+$, HCN, HNC in a field 3'x3' covering the
northern PDR. They also measured CS, CN, C$_2$H, H$^{13}$CO$^+$,
H$^{13}$CN, and HN$^{13}$C at selected positions.
Using chemical equilibrium calculations, the authors showed that
the variations of the molecular abundances can be explained by
the influence of the stellar UV radiation on the chemistry if the
emission arises in a region at a visual extinction between 6 and 10 mag.
from the star.

\subsubsection{Line intensities}

A first set of line intensities as computed based on an LVG model of the PDR 
with a gas density  of $2.5\,10^4$ cm$^{-3}$, a temperature of 100~K, a
line width of 2 km/s, and a H$_2$ column density of $10^{22}$ cm$^{-2}$.
The temperature of 100 K was chosen as a typical temperature for the
warm molecular layers in a PDR. It is clear than even higher temperatures
exist but they represent a small fraction of the total gas.
On the other hand cooler layers will radiate less in the submm.
The LVG model is expected to provide reasonable estimates for the line 
intensities except for some molecules, like CH$^+$, where the cross sections
are not well known, so the accuracy of the result is not guaranteed.

An independent plane-parallel PDR model (Le Bourlot 2000, Le Bourlot 2003)
using the same physical
parameters was used to obtain an estimate for the modeling uncertainty
of the line predictions. Here, the intensities were obtained by an
approximate integration of the radiative transfer through the PDR
assuming LTE excitation. The comparison of the line predictions shows 
discrepancies up to a factor 10 stressing the need for actual observations.
The time estimates were based on the results from the LTE model.



\begin{table}[!ht]
\caption{Estimated line intensities for NGC7023.
\label{tab_ngc7023}}
\small{
\centering
\begin{tabular}{lrlllcccc}
\noalign{\smallskip}\hline
Species & Freq. & Column density & \multicolumn{2}{c}{$T_B$ [K]}
& D$^1$ &$t_{int}$ \\
  & [GHz] & [cm$^{-2}$] & LVG & LTE & [$'$] & [s]\\
\hline\noalign{\smallskip}
CII         & 1900.545 & $5\,10^{17}$&66        &81     & 5 &     1.0 \\
CH          &  536.761 & $4\,10^{13}$&          &4.6    & 3 &     1.0 \\
CH          & 1656.961 & $4\,10^{13}$&          &14     & 3 &     1.0 \\
CH$^+$      &  835.070 & $4\,10^{11}$&0.003     &0.075  & 3 &   645.9 \\
CH$^+$      & 1669.160 & $4\,10^{11}$&0.0095    &0.24   & 3 &  1843.0 \\
NH          &  974.479 & $1\,10^{14}$&          &0.76   & 5 &     7.3 \\
NH$^+$      & 1012.524 & $1\,10^{07}$&          &0.0    &   &  $>10^4$\\
NH$_3$      &  572.498 & $1\,10^{15}$&0.17      &0.73   & 1 &     4.8 \\
NH$_3$      & 1168.452 & $1\,10^{15}$&0.14      &0.022  & 1 &  $>10^4$\\
OH$^+$s     &  971.804 & $7\,10^{10}$&          &0.062  & 5 &  1100.2 \\
H$_3$O$^+$  & 1655.814 & $5\,10^{12}$&          &0.005  & 3 &  $>10^4$\\
H$_3$O$^+$  &  984.697 & $5\,10^{12}$&          &0.010  & 3 &  $>10^4$\\
H$_2$O      & 1113.343 & $1\,10^{15}$&1.1       &2.9    & 3 &     1.0 \\
H$_2$O      &  987.927 & $1\,10^{15}$&0.53      &0.32   & 3 &    43.1 \\
H$_2$O      &  752.033 & $1\,10^{15}$&0.063     &0.42   & 3 &    18.8 \\
H$_2$O      &  556.936 & $1\,10^{15}$&1.5       &4.3    & 3 &     1.0 \\
H$_2$O      & 1669.905 & $1\,10^{15}$&          &0.70   & 3 &   209.0 \\
HDO         &  893.639 & &              &$7.5\,10^{-5}$ & 3 &  $>10^4$\\
H$_2^{18}$O & 1101.698 & $2\,10^{12}$&0.0024    &0.013  & 3 &  $>10^4$\\
CO(10-9)    & 1151.985 & $5\,10^{17}$&66.4      &6.9    & 1 &     1.0 \\
CO(15-14)   & 1726.603 & $5\,10^{17}$&6.0       &0.04   & 1 &  $>10^4$\\
13CO(10-9)  & 1101.350 & $1\,10^{16}$&3.9       &0.17   & 1 &   163.2 \\
13CO(15-14) & 1650.768 & $1\,10^{16}$&0.043     &$8.1\,10^{-4}$ & 1 & $>10^4$\\
\hline   
OH          & 2509.988 & $2\,10^{14}$&0.061     &5.6    & 5 &    1.0 \\
OH          & 2514.316 & $2\,10^{14}$&0.046     &3.5    & 5 &    1.0 \\
O           & 4745.804 & $3\,10^{18}$&          &26     & 5 &    1.0 \\
O           & 2060.068 & $3\,10^{18}$&          &18     & 5 &    1.0 \\
\hline\noalign{\smallskip}
\end{tabular}\\
\flushleft
${}^1$ $D$ denotes the size of the emission across the PDR estimated from
CO observations.
}
\end{table}

\subsubsection{Observing strategy and time estimates}


The size of the emission in the
faint lines like CH$^+$, NH$^+$ can be estimated from the area of the CO(6-5)
extension (this is rather small 1$'$x1$'$). 
CI is extended there as it is in many other sources.
For strong lines, we thus expect a larger area with detectable
emission. Here, we can obviously cover more space and sample
a wider variety of physical conditions.  For CI, but also CII, one 
H2O lines it is
interesting to observe along long cuts a few arcmin (3$'$ NS to 5$'$EW) in
size to cover a wider variety of physical conditions.


The estimate shows that the lines of NH$^+$, H$_3$O$^+$, HDO, H$_2^{18}$O,
and the 15-14 transitions of the two CO isotopes
would require prohibitively long observing times, so that they
cannot be observed.
The weak lines from CH$^+$ and OH$^+$
are only to be observed at two single positions,
corresponding to the CO 6-5 peak and 2' apart to get a detection
and a rough estimate of the extent. The weak water line at 1669 GHz is
then already included as bonus line. This requires 1.8~h observing time.

The $^{13}$CO 10-9 line also needs 
a relatively large observing
time per position preventing a full OTF scan. Here,
7 positions along two NS cuts including the infrared filaments 
should be observed. This requires 0.7~h observing time.

For all brighter lines full OTF scans along the stripes with
a length of 5$'$ are possible. For the [C{\sc II}], CH, NH,
the NH$_3$ 572 GHz, the H$_2$O 1113 and 557 GHz and the CO 10-9
lines OTF scans with a sampling ever $5''$ (i.e. fully sampling
at 1.9\,THz), and every 5\,sec are adequate. With two cuts
and 6 frequency settings (bonus lines already taken into account)
this would  take 1~h. The H$_2$O 987 and 752 GHz lines are
somewhat weaker, so that the integration times should be
40 and 20~s per point leading to a total observing time for
the scans of 2~h.

The four OH and OI lines observable only with PACS are bright.
We plan to observe them along the same cuts planned for
HIFI. To guarantee some overlap in the spatial coverage of
the PACS fields we need 9 settings to cover each cut. Assuming 
half a minute integration time for each spatial setting and
frequency results in 0.6~h PACS observing time.

The total observing time is for this template source thus sums
up to {\bf 7~h}.

\subsection{IC 63\protect\footnote{M.\ Spaans (Kapteyn Astronomical Institute)}} 
\subsubsection{Introduction}

The reflection/emission nebula IC 63 is located at about 1.3 pc (projected,
20') from the illuminating B0.5p ($T_{\rm eff}\approx 30,000$ K) star
$\gamma$ Cas. The source has been mapped
in $^{12}$CO 2-1 and 3-2, $^{13}$CO 2-1, and CS 2-1 which shows a small
molecular cloud of about $1'\times 2'$ that coincides with the peak in optical
nebulosity and IRAS 100 $\mu$m emission. IC 63 is nearby, at about 230~pc, and
provides an excellent example of an edge-on PDR with a very simple, roughly
circular/spherical, geometry. Measured CO, HCO$^+$, HCN, CS and H$_2$CO line
ratios yield a (mean) temperature of $T\approx 50$ K and a (constant) density
of $n({\rm H}_2)\approx 4\times 10^4$ cm$^{-3}$. The derived H$_2$ column
density and visual extinction (Jansen et al. 1994) are
$N({\rm H}_2)=(5\pm 2)\times 10^{21}$ cm$^{-2}$ and $A_{\rm V}=6.7\pm 2.3$ mag.
More complex, volatile molecules such as CH$_3$OH, CH$_3$CN, HNCO, SO and
SO$_2$ have not been detected. The narrow line widths of the observed
transitions indicate the absence of any shocked material, i.e., IC 63 provides
a `clean' PDR.
IC 63 has been mapped as well with ISOCAM in the CVF mode
(imaging spectroscopy) between 5 and 16.5 $\mu$m (Foss\'e et al., 2000).
The emission around the 6.2 $\mu$m AIB feature, the
stretching mode of the C-C bond, follows the PDR border  and exhibits several
peaks just behind the PDR and inside the molecular tail pointing away from
$\gamma$ Cas. Also, high abundances of hydro-carbons, that are unexplained
by the UMIST95 or New Standard Model chemical networks, have been detected
by Teyssier et al. (2003) with the IRAM 30-m, SEST and
Effelsberg 100-m telescopes. The observations show that small hydrocarbons
are present up to the edge of the PDR.
This is remarkable since the strong UV radiation field should destroy
species like C$_3$H$_2$ and C$_4$H at the edge.

The strength of the FUV radiation field at the edge of the cloud is about
$I_{\rm UV}\approx 650$ times the Draine field (900 times Mathis). The actual
$\gamma$ Cas radiation field has been measured by the S2-68 and IUE satellites
and deviates substantially from the Draine and Mathis parameterizations of the
mean Galactic radiation field above 1500 \AA . A broad-band measurement with a
filter covering 912-1075 \AA\ has been made by Troy at al.\ (1975) in a rocket
experiment and can be used to constrain the important 912-1100 \AA\ range
where CO and H$_2$ dissociation occurs.
Observations of the [CI] 609 and [CII] 158 $\mu$m lines as well as multiple
high-$J$ CO transitions and its isotopes have allowed the accurate
determination of the carbon budget (Jansen et al., 1994). The carbon gas
fraction in IC 63 is only $13^{+6}_{-4}\%$ of Solar, significantly less than
in diffuse clouds and in some dense clouds. Note that IC 63 contains warm gas,
$T>500$ K, gas as indicated by ISO observations of pure
rotational H$_2$ line emission of S(0), S(1), S(3) and S(5). This temperature
is a factor of a few above the model predictions of $\sim 200$ K.

The physical and chemical characteristics of IC 63 are well established.
Its modest distance and simple geometry render it an ideal object to test
the basics of PDR physics (thermal and chemical balance). At the same time,
its sphericity deviates from the often assumed slab-like geometry in
theoretical models and allows a detailed comparison to be made.
It is not entirely clear whether the illuminating star gamma Cas lies at the
same distance as IC 63. If the star is slightly closer to us then most of
the observable surface of IC 63 is irradiated (albeit under a modest angle).
For IC 63 a very extensive database of ground- and space-based
observations already exists. Additional CSO and IRAM
observations will be obtained for the lines of CS 5-4 as well
as CO$^+$ $N=2-1$, $J=5/2-3/2$ and $J=3/2-1/2$.

\begin{table}[b]
\centering
\caption{Overview of Available Observations IC 63}
\label{table_ic63_observations}

\flushleft
{\bf IRAS}: ($20''\times 30''$ source)\\
\centering

\begin{tabular}{rrr}
\hline
wavelength ($\mu$m)&   peak surface brightness (MJy sr$^{-1}$) & color corrected flux (Jy)\\
\hline
12                 &   8.0 & 2.57\\
25                 &   9.7 & 2.73\\
60                 &   86 & 56\\
100                &   101 & 105\\
\hline 
\end{tabular}
\end{table}

\begin{table}
\addtocounter{table}{-1}
\centering
\caption{continued}
\flushleft
{\bf CSO, IRAM, NRAO, JCMT, KAO}:\\
$\Delta V$ of the order of 1 km s$^{-1}$. For multiple observations of the
same line with different telescopes, only the largest beam observations are
listed.\\
\centering

\begin{tabular}{lrrr}
\hline
species &  transition &               T$_{\rm mb}$ (K) &  I (K km s$^{-1}$)\\
\hline
CO &         1-0                   &  10.8 &              13.1\\
   &         2-1                   &  23.3 &              33.1\\
   &         3-2                   &  23.7 &             38.2\\
$^{13}$CO &  1-0                   &  0.39 &             0.34\\
          &  2-1                   &  1.64 &              2.08\\
C$^{18}$O &  1-0                   &  0.10 &              0.06\\
          &  2-1                   &  $<0.6$ &                \\
CS        &  2-1                   &  0.21   &            0.23\\
          &  3-2                   &  0.20   &            0.60 (unresolved)\\
          &  5-4                   &  $<0.07$&                \\
CN        &  1 3/2 5/2 - 0 1/2 3/2 &  0.54   &            0.42\\
          &  1 1/2 3/2 - 0 1/2 3/2 &  0.28   &            0.20\\
          &  2 5/2 7/2 - 1 3/2 5/2 &  0.39   &            0.33\\
          &  2 5/2 3/2 - 1 3/2 1/2 &  0.25   &            0.26\\
          &  2 5/2 5/2 - 1 3/2 3/2 &  0.31   &            0.32\\
HCN       &  1,1 - 0,1             &  0.16   &            0.17   (NRAO 70'')\\
          &  1,2 - 0,1             &  0.32   &            0.34      idem\\
          &  1,0 - 0,1             &  0.08   &            0.15      idem\\
          &  3-2                   &  $<0.28$&                \\
HNC       &  1-0                   &  0.22   &            0.29\\
HCO$^+$   &  1-0                   &  0.44   &            0.94\\
          &  3-2                   &  0.35   &            0.81 (unresolved)\\
          &  4-3                   &  0.38   &            0.37\\
H$_2$CO   &  2$_{1,1}$-1$_{1,0}$   &  0.14   &            0.23 (unresolved)\\
          &  2$_{1,2}$-1$_{1,1}$   &  0.25   &            0.27\\
          &  3$_{1,2}$-2$_{1,1}$   &  0.18   &            0.18\\
          &  3$_{2,1}$-2$_{2,0}$   &  $<0.24$&                \\
          &  3$_{2,2}$-2$_{2,1}$   &  $<0.08$&                \\
CH$_3$OH  &  2-1                   &  $<0.14$&                \\
C$_2$H    &  1$_{2,1}$-0$_{1,0}$   &  0.37   &            0.45\\
N$_2$H$^+$&  1$_2$-0$_1$           &  0.09   &            0.13\\
          &  1$_1$-0$_1$           &  $<0.07$&                \\
          &  1$_0$-0$_1$           &  $<0.07$&                \\
HC$_3$N   &  11-10                 &  $<0.10$&                \\
CH$_3$CN  &  8$_K$-7$_K$           &  $<0.12$&                \\
HNCO      &  6$_{0,6}$-5$_{0,5}$   &  $<0.15$&                \\
H$_2$S    &  1$_{1,0}$-1$_{0,1}$   &  0.72   &            0.44\\
SO        &  2$_2$-1$_1$           &  $<0.06$&                     \\
          &  3$_2$-2$_1$           &  $<0.12$&                     \\
SO$_2$    &  5$_{2,4}$-4$_{1,3}$   &  $<0.20$&                     \\
SiO       &  3-2                   &  $<0.10$&                     \\
          &                        &         &                     \\
C{\sc II}    & 158 $\mu$m & 7.2   &           10.2\\
C{\sc I}     & 609 $\mu$m & 2.1   &           2.2\\
CO       & 4-3        & 30.9  &           32.1\\
$^{13}$CO 6-5 &       & 2.0   &           2.1\\
\hline
\end{tabular}
\end{table}

Detailed one-dimensional (slab geometry) and two-dimensional
(prolate geometry) models have been constructed (see the description
in Jansen et al.\ 1995 and Spaans 1996) that accurately
reproduce the available observations. Table \ref{table_ic63_modelling}
presents the results of the
model calculations. The extensive data set used as a constraint is listed in
Table \ref{table_ic63_observations}.

The species/lines accessible to HIFI can be used to investigate the specific
problems of 
\begin{enumerate}
\item{}The warm layer of H$_2$ gas at the edge of the PDR (through the [CII], [OI], high-J CO and $^{13}$CO, and H$_2$O lines).
\item{}The chemistry of the radical region (through the CH and CH$^+$ lines).
\item{}The chemistry of water (through the H$_2$O, HDO, OH, OH$^+$ and H$_3$O$^+$ lines).
\item{}The electron abundance and its role in the PAH charge balance and nitrogen chemistry (through the NH$^+$ line).
\end{enumerate}

\subsubsection{Observing Strategy}

\begin{table}
\centering
\caption{ Estimated peak line temperatures for the CO 2-1 peak in IC63 and
time estimates}.
\label{table_ic63_modelling}

%\centering
\begin{tabular}{lrrrrr}
\noalign{\smallskip} \hline \noalign{\smallskip}
  species & frequ. & \multicolumn{1}{c}{column density [cm$^{-2}$]} & \multicolumn{1}{c}{$T_B$ [K]} & \multicolumn{1}{c}{size of emission} &$t_{\rm tot}$ \\ 
          &        & \multicolumn{3}{c}{} \\
  & [GHz] & \multicolumn{1}{c}{at peak} & \multicolumn{1}{c}{at peak}& [$'$]& [s] \\  
\noalign{\smallskip} \hline \noalign{\smallskip}
C{\sc II}         & 1900.545 & 1.3E+17 &   14.2   & 0.4 & 1\\
CH                &  536.761 & 1.4E+14 &   10.1   & 0.4 & 1\\
                  & 1656.961 &         &   22.3   & 0.3 & 1\\
CH$^+$            &  835.070 & 5.8E+11 &   0.011 & 0.4 & $2\,10^5$ \\
                  & 1669.160 &         &   0.024 & 0.3 & $3\,10^5$ \\
NH                &  974.479 & 3.0E+11 &   1.2  & 0.4 & 12\\
NH$^+$            & 1012.524 & 1.7E+06 &   1.1E-06 & 0.4 & Inf\\
NH$_3$            &  572.498 & 1.4E+14 &   3.3   &  0.3& 5\\
                  & 1168.452 &         &   0.62  & 0.3& 50\\
OH$^+$            &  971.804 & 3.5E+10 &   0.029 & 0.7& $1.4\,10^4$\\
H$_3$O$^+$        & 1655.814 & 9.0E+12 &   0.043 & 0.4& $9\,10^4$ \\
                  &  984.697 &         &   0.19  & 0.6& 360 \\
H$_2$O            & 1113.343 & 2.2E+16 &   6.1   & 0.6 & 1\\
                  &  987.927 &         &   4.1   & 0.6 & 1\\
                  &  752.033 &         &   9.3   & 0.6 & 1\\
                  &  556.936 &         &   12.8   & 0.7& 1 \\
                  & 1669.905 &         &   4.5   & 0.7 & 8 \\
                  & 1716.770 &         & 9.7     & 0.5 & 2\\
                  & 1153.126 &         & 1.0     & 0.5 & 15 \\
HDO               &  893.639 & 3.4E+11 &   0.12 & 0.6 & 850\\
H$_2^{18}$O       & 1101.698 & 4.3E+13 &   3.0  & 0.7 & 2\\
OH                & 1834.747 &         &  0.04  &   0.4 & $9\,10^4$\\
CO 10-9           & 1151.985 & 3.2E+17 &   0.98 & 0.3 & 20\\
\hspace*{1cm} 15-14 & 1726.603 &         &   0.01 & 0.3 & Inf\\
\hspace*{1cm} 16-15 & 1841.345 &       & 4.3E-03 &  0.3 & Inf\\
$^{13}$CO 10-9    & 1101.350 & 5.2E+15 &   0.06 & 0.3 & 5000\\
$^{13}$CO 15-14   & 1650.768 &         &   8.0E-04 & 0.3 & Inf\\
\noalign{\smallskip} \hline \noalign{\smallskip}
 OH 1- (119)      & 2509.988 & 3.5E+13 &   0.24  & 0.7 & 330\\ 
 OH 1+ (119)      & 2514.316 &         &   0.23  & 0.7 & 330\\ 
 OI 2-1 (63)      & 4745.804 & 4.9E+18 &   0.4   & 0.4 & 108 \\ 
 OI 1-0 (145)     & 2060.068 &         &   0.22  & 0.4 & 1070\\ 
\noalign{\smallskip} \hline \noalign{\smallskip}
\end{tabular}
\end{table}

The angular size of $\approx$ $1'$, edge-one geometry and large peak
brightness temperatures of selected lines allow us to perform a number of
pointings along and perpendicular to the direction of the impinging $\gamma$
Cas radiation field.
This will allow us to (1) measure the stratification that should be present in
this edge-on PDR and to identify the various chemical zones that result from
the ion-molecule chemistry and the attenuation of the FUV radiation field, and
(2) to determine the impact of the circular/spherical geometry on the thermal
and chemical balance of the nebula.

It should be possible to do a grid, fully sampling at 1.9 THz (5''), for
the strongest transitions: [CII], CH, H$_2$O and H$_2$18O, NH, NH$_3$ and
CO 10-9 (all lines less than 50 s), and to do one long
integration of all lines with integration times less than $10^4$ seconds
on the peak of the CO 2-1 emission.
The three PACS line observations require 25 min. A single integration on the CO 2-1
peak takes 1.8 hr (with the bulk going into the $^{13}$CO 10-9 line).
The grid, covering IC 63 between the CO 2-1 peak and the exposed edge
(about 50 pointings), takes 1.7 h. The total requested time for this source is
{\bf 4.0 h} and would drop to 2.2 h when the $^{13}$CO 10-9 line 
is skipped from the list of single integrations on the CO 2-1 peak.

\subsubsection{Overview of Line Predictions}

One-dimensional and two-dimensional (Jansen et al., 1994, Spaans 1996) models
have been constructed for IC 63 using a plane-parallel slab and ellipsoidal
geometry. These reproduce successfully the available observations. The latest
UMIST database release has been used to update these models and detailed model
predictions are presented in Table \ref{table_ic63_modelling}. The model
predictions are based on an
extension of the 2D prolate PDR model of Jansen et al.\ (1995) to a 3D prolate
ellipsoid PDR, using the CSO CO 2-1 map as the cloud's projected shape. The
model parameters are as given above and the beam is located on the (0,0)
position of the CO 2-1 map of Jansen et al.\ (1994).

%
We assumed a D/H abundance ratio of
$1.5\,10^{-5}$ as found in the ISM by Myers et al.\ 1996.
For the HIFI lines, we used the assumptions of
0.7\,km/sec resolution (1\,MHz at 500\,GHz), line width of 2.5\,km/sec,
position switched observations, including an overhead of 50\%.
In case of the PACS lines, the channel resolution of 180\,km/sec is assumed. 

The listed diameters are the widths of the emission regions while the height
is in all cases 0.6' with a small uncertainty due to the prolate shape of the
cloud.  The total integration times were computed for
5$\sigma$ detections at the central position.

\subsection{Rho Ophiuchi filament\protect\footnote{D. Teyssier SRON}}
%, L. Verstraete IAS, E. Habart Arcetri}}

\subsubsection{Introduction}

The Ophiuchi molecular complex is one of the nearest star formation
region (distance 135$\pm$15\,pc, Habart et al. 2003). While its
so-called ``North'' area is considered inactive, the West counterpart
contains two main massive cores, L1688 and L1689, heated on large
scales by the Sco-OB2 association. We concentrate here on the L1688
complex, which hosts numerous pre-stellar cores (e.g. Motte et
al. 1998), as well as hundreds of YSOs (Greene \& Young 1992, Bontemps
et al. 2001), very likely responsible for the strong luminosity in the
infrared.

In the mid-IR image of the dust emission made with the ISO camera (LW2
filter tracing the PAH emission, Abergel et al. 1996), the western
edge of the L1688 cloud is delineated by a long and bright interlaced
filaments of width 1\am~($\sim 0.03$\,pc), located at the edge of the
dense molecular gas as traced by its $\rm ^{13}CO(1-0)$ emission
(Loren 1989). The $\rho$ Ophiuchi West ($\rho$ Oph-W) interface,
centered at $\alpha$ $\sim$ 16h 25m 57s and $\delta$ $\sim$
-24\deg20\am50\as~ (Epoch 2000), has an edge-on geometry with respect
to the B2 star HD147889 which lies 9\am~away from the interface. The
B2 star HD147889, which appears to be the center of a more-or-less
spherical cavity, is in this region the main heating source (Liseau et
al. 1999).

The $\rho$ Oph-W interface is moderately excited ($\chi=400$) and its
density gradient has been derived from the PAH and from the
H$_2$(1--0) S(1) line emission (Habart et al. 2003). The density at
the peak of the H$_2$ and PAH emission is 2$\times 10^4$
cm$^{-3}$. Modeling has shown that the H$_2$ emitting material has a
depth of about 0.2 pc along the line of sight. The density rises by a
factor 2 within the molecular region where the CO emission comes
from. The low ortho-to-para ratio as well as the ratio of rotational
to ro-vibrational lines of H$_2$ suggest that advection motions of
fresh molecular gas across the interface are important in this region
(Habart et al. 2003).

On large scales, the mid-IR emission shape is well reproduced by the
CO and isotopes maps obtained by Lada \& Wilking (1980) and Wilking\&
Lada (1983). Their data revealed a complex velocity structure and, in
particular, significant self-absorption in the $^{13}$CO species. On
these scales, $^{12}$CO data are compatible with $A_{\rm V}$ in the
range 50-100\,mag. Hydrocarbon molecules have also been mapped in
this source by Teyssier et al. (2004), revealing a rich carbon
chemistry at the PDR interface. As in the case of the Horsehead
nebula, the HIFI observations will highlight the interplay between the
chemistry, the dynamics and the dust grains.


\subsubsection{Line intensities}

As for the Horsehead, most line intensities have been estimated at LTE
for a column density of $10^{22}$ cm$^{-2}$, a line width of 2 km/s
and the abundances given in Table \ref{tab-time-rophw}. The excitation
temperatures and abundances have been estimated from a PDR model with
$\chi=60$ and $n=10^4$ cm$^{-3}$. When available measured intensities
have been used.


%%%%%%%%%%%%%%%%%%%%%%%%%%%% FIG ROPH %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%
\begin{figure}
 \begin{center}
 \includegraphics[width=8cm]{roph_strat.eps}
 \caption{Positions aimed at being observed in the $\rho$ Oph-W PDR, overlaid over the ISOCAM map around 7\,$\mu$m (Abergel et al. 1996). The black squares indicate the positions to be observed in deep integration, while the white squares stand for positions where stronger lines are probed (see text). They correspond to points also observed by ISO-SWS. An arrow shows the direction of the illuminating star HD147899.}
 \label{roph obs strategy}
 \end{center}
\end{figure}
%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% 



%%%%%%%%%%%%%%%%%%%%%%%%%%%%% TAB ROPH %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%
\begin{table}[!ht]
\caption{Time estimate for the $\rho$ Oph center position. We assume 
FWHM=1 km/s and a velocity resolution of 0.3\,km/sec, $A=0.125$ for
position switched observations (respectively $A=0.25$ for frequency
switched observations, only considered for C[II]) and a SNR of 10 (5) for strong(weak) lines. The OH and O lines are
observed by PACS. Column (1) is the species, column (2) gives the
frequency of the transition, column (3) gives the abundance, column
(4) gives the excitation temperature, column (5) the brightness
temperature and column (6) the observing time in seconds with the achieved SNR in parenthesis.} \label{tab-time-roph}
\centering
\label{tab-time-rophw}
\begin{tabular}{lrrrrr}
\noalign{\smallskip} \hline \noalign{\smallskip}
  species & $\nu$ & [X/H] & $T_{ex}$ & $T_B$ & Time (SNR)\\
  & [GHz] & & [K] & [K] & [sec.] \\
\noalign{\smallskip} \hline \noalign{\smallskip}
 C{\sc II}      & 1900.545 & 	     & 50   &  50${}^a$ 	& 2.7 (10)\\
 C		&  492.161 & 1.4e-05 &      &  10${}^a$		& 0.7 (10) \\
		&  809.342 & 1.4e-05 &      &  10${}^a$	        & 1 (10)\\
CH            	&  536.761 & 7.2e-09 & 10   & 1.9 		& 19 (10) \\
	       	& 1656.961 & 7.2e-09 & 10   & 0.017		& ... \\
 CH$^+$   	& 835.070  & 3.6e-09 & 10   & 0.74 		& 195 (10) \\
		&1669.160  & 3.6e-09 & 10   & 0.006 		& ... \\
 NH   		& 974.479  & 5.4e-09 & 10   & 0.44		& 163 (5) \\
 NH$_3$   	& 572.498  & 2.4e-07 & 10   & 1.9 		& 21 (10) \\
 		& 1168.452 & 2.4e-07 & 10   & 0.21 		& 877 (5) \\
 H$_3$O$^+$   	&  984.697 & 4.2e-10 & 10   & 0.10 		& 1800$^b$ (5) \\
 	   	& 1655.810 & 4.2e-10 & 10   & 0.01 		& ... \\
 H$_2$O  	& 1113.343 & 5.0e-07 & 10   & 0.26 		& 542 (5)  \\
  		&  987.927 & 5.0e-07 & 10   & 0.42		& 727 (10) \\
		&  752.033 & 5.0e-07 & 10   & 0.24 		& 430 (5) \\
                &  556.936 & 5.0e-07 & 10   & 2.0		& 18 (10) \\
  		& 1669.905 & 5.0e-07 & 10   & 0.03 		& ... \\
 H$_2{}^{18}$O  & 1101.698 & 1.0e-09 & 10   & 0.26 		& 542 (5) \\
                & 547.676  & 1.0e-09 & 10   & 1.9 		& 19 (10) \\
CO       	& 1151.98  & 3.0e-04 & 50   & 27.3		& 0.2 (10) \\
 		& 1841.34  & 3.0e-04 & 10   & 0.15 		& ... \\
 OH  		& 2509.988 & 3.2e-09 & 12   & 7.8e-10 (!)	& ... \\
 \noalign{\smallskip} \hline \noalign{\smallskip}
 OH${}^f$ 	& 2509.988 & 3.2e-09 & 12   & 6.2e-15${}^c$	& 3173 (5) \\
 OH${}^f$  	& 2514.316 & 3.2e-09 & 12   & 6.0e-15${}^c$	& 3410 (5) \\
  O  		& 4745.804 & 1.7e-04 & 25   & 2.9e-14${}^{a,c}$ & 375 (8) \\
  O  		& 2060.068 & 1.7e-04 & 25   & 3.1e-13${}^{a,c}$ & 108 (50)\\
\noalign{\smallskip} \hline \noalign{\smallskip}
\end{tabular}
\flushleft
${}^a$ measurements or inferred from Habart et al. 2003.
${}^b$  The LTE excitation temperature is probably underestimated
and we assume it can be detected in 0.5 hour.
${}^c$ in erg cm$^{-2}$ s$^{-1}$, assuming the PACS resolution of
180\,km/s and a 5$\sigma$ sensitivity of
2.8e-15\,erg cm$^{-2}$ s$^{-1}$ in one hour,
${}^f$ Observed simultaneously
\end{table}
%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%

%\begin{table}[!ht]
%\caption{
%\label{tab-flux-rophw} Measured cooling line intensities in $\rho$ Oph-W}
%\centering
%\begin{tabular}[t]{ccccc}
%\noalign{\smallskip} \hline \noalign{\smallskip}
%C{\sc II}$^a$ & C$^a$ & O$^a$ & O$^a$ & $^{13}$CO (3-2)$^a$ \\
%158$\mu$m & 609$\mu$m & 63$\mu$m & 146 $\mu$m & 907 $\mu$m \\
%\noalign{\smallskip} \hline
%\hline \noalign{\smallskip}
%3 & 0.5 & 0.6 & 0.1 & 0.2   \\
%\noalign{\smallskip} \hline \noalign{\smallskip}
%\end{tabular}
%\flushleft
%$^a$ in 10$^{-4}$ erg cm$^{-2}$ s$^{-1}$ sr$^{-1}$
%\end{table}


\subsubsection{Observing strategy and time estimates}%

The observing strategy is similar to that for the Horsehead nebula, except that much less positions will be probed (see Fig.~\ref{roph obs strategy}). While strong lines ([CII], [CI], CH, CH$^+$, H$_2$O, H$_2^{18}$O, NH$_3$, CO) will be observed along 7 positions crossing the PDR, the weaker lines will be restricted to three locations:
a position on the molecular core, one position centered on the filament and a position away from the filament towards the HD147889 star. All these positions correspond to points previously probed by ISO-SWS (Habart et al. 2003). The required time for each molecules are indicated in Tab.~\ref{tab-time-rophw}. For each of the three deep integration position, {\bf 1.5\,h} and {\bf 1.0\,h} are requested for HIFI and PACS respective observation. In strong lines, only 6 minutes per points are required. In total, this amounts to {\bf 7.9\,h}.

\subsection{Ced 201\protect\footnote{M.~Spaans (Kapteyn Astronomical Institute)}}

\subsubsection{Introduction}

The reflection nebula Ced 201 is at a distance of 420~pc and is illuminated
by a B9.5 ($T\sub{eff}\approx10,000$ K) star. Its angular extent is about $
1'\times 1'$ in $^{12}$CO 2-1 and $^{13}$CO 2-1 line emission (Kemper
et al. 1999) and is roughly circular on the sky. The star is located in front
of the cloud, leading to a face-on PDR. Because the illuminating radiation
field has a color temperature of only 10,000 K, Ced 201 is a so-called cool
PDR (Spaans et al. 1994). In a cool PDR, the temperature of
the gas is lower because the photo-electric heating rate is decreased through
the weaker FUV radiation field, while the abundances of molecules are enhanced
with respect to classical PDRs. In cool PDRs the resulting photo-electric
heating rate has a much larger contribution from grain species with low
ionization potentials and negative charges, i.e., PAHs. As such, cool PDRs are
ideal to study the role of PAHs in the photo-electric heating process.
In this light it should also be noted that Cesarsky et al. (2000) have
studied the evolution of carbonaceous dust in Ced 201 through
spectrophotometric imaging in the mid-IR. The spectra of Ced 201 indicate that
a transformation of very small carbonaceous grains into the carriers of the
Aromatic Infrared Bands (AIBs) is taking place under the influence of the
radiation field of the illuminating star and/or of shock waves.

Observations of HCO$^+$ and CS yield a density of $n({\rm H}_2)\approx 5000$
cm$^{-3}$ and a gas temperature $T\approx 40$ K. Kemper et al.\ (1999) find
that Ced 201 is flattened (oblate), by a factor of $\sim 2.4$, along the line
of sight with an H$_2$ column density
of $1.6\times 10^{21}$ cm$^{-2}$, or about 2 mag of extinction.
[CI] 609 $\mu$m, CO and ISO [CII] 158 $\mu$m observations constrain
the carbon budget (Kemper et al., 1999) to be about $20\pm 8$\%
of Solar, less than the typical diffuse cloud value. Pure rotational H$_2$ line
emission, $S(3)$ and $S(5)$, obtained with ISO, indicates warm, 330 K, gas
that is currently unexplained and seems to indicate
an additional heating source (other than photo-electric). If the gas is located
in a single layer, then this layer is only 0.05 mag thick. A shock origin
appears likely and would require a C-shock of about 7 km s$^{-1}$ in a gas
of pre-shock density $10^4$ cm$^{-3}$. The enhancement of
the radiation field is I$_{\rm UV}$=200 in units of the Draine field between
13.6 and 2 eV for a separation of the star and the cloud of 0.4 pc.

\begin{table}
\centering
\caption{Overview of Available Observations Ced 201}
\label{table_ced201_observations}
\flushleft
{\bf JCMT}: ($v_{\rm LSR}=-4.9$ km s$^{-1}$)\\
\centering

\begin{tabular}{lrrr}
\hline
species   &transition   &I (K km/s)  &T$_{\rm peak}$ (K)\\
\hline

HCO$^+$     &1-0          & 1          & 0.15\\
            &3-2          &            & $<0.08$\\
CS          &2-1          &            & $<0.03$\\
$^{12}$CO   &2-1          & 21.3       & 23.2\\
            &3-2          & 25.3       & 33.1\\
$^{13}$CO   &2-1          & 5.4        & 6.9\\
            &3-2          & 3.9        & 5.5\\
C$^{18}$O   &2-1          &            & $<0.09$\\
C{\sc I}    &609 $\mu$m   &            & $<1.3$\\
\hline
\end{tabular}
\medskip

\flushleft
{\bf ISO}:\\
\centering

\begin{tabular}{lrrr}
\hline
species   &transition     &I (W m$^{-2}$) & \\
\hline

C{\sc II}       &158 $\mu$m     &3.80(-15)  &(beam 90'')\\
O{\sc I}        &63 $\mu$m      &$<8(-15)$  &(beam 90'')\\
H$_2$       &S(1) 17 $\mu$m &1.10(-16)  &(beam 14''x27'')\\
H$_2$       &S(3) 9 $\mu$m  &5.61(-17)  &(beam 14''x20'')\\
\hline
\end{tabular}
\end{table}

The observational characteristics of Ced 201 are well established
(ground-based, ISO). The low color temperature of the illuminating radiation
field and its simple geometry makes it a very suitable object to test the
physics of cool PDRs.
Detailed one-dimensional (slab geometry) and three-dimensional
(oblate spheroid geometry) models have been constructed (see the description
in Spaans et al.\ 1994 and Kemper et al.\ 1999) that accurately
reproduce the available ground-based and ISO observations, with the exception
of the warm 330 K gas that is required to reproduce the S(3) and S(5) pure
rotational lines of H$_2$. Table \ref{table_ced201_modelling} presents the results of the model
calculations. The extensive data set used as a constraint is listed in
Table \ref{table_ced201_observations}.

\begin{table}
\caption{Line predictions and observing time estimates for Ced 201}
\label{table_ced201_modelling}
\centering
\begin{tabular}{lrrrr}
\noalign{\smallskip} \hline \noalign{\smallskip}
  species & frequ. & \multicolumn{1}{c}{column density [cm$^{-3}$]} & \multicolumn{1}{c}{$T_B$ [K]} &$t_{tot}$ \\
          &        & \multicolumn{3}{c}{} \\
  & [GHz] & \multicolumn{1}{c}{at peak} & \multicolumn{1}{c}{at peak}& [s] \\
\noalign{\smallskip} \hline \noalign{\smallskip}
C{\sc II}         & 1900.545 & 3.7E+17 & 11.2   & 2 \\
CH                &  536.761 & 5.3E+13 & 1.16   & 30 \\
                  & 1656.961 &         & 2.31   & 35 \\
CH$^+$            &  835.070 & 8.7E+09 & 0.007 & $4\,10^5$ \\
                  & 1669.160 &         & 0.014 & $9\,10^5$ \\
NH                &  974.479 & 1.2E+12 & 0.19  & 450 \\
NH$^+$            & 1012.524 & 8.8E+06 & 6.6E-07 & Inf \\
NH$_3$            &  572.498 & 1.9E+11 & 0.02   & $8\,10^4$ \\
                  & 1168.452 &         & 2.0E-04  &  Inf \\
OH$^+$            &  971.804 & 1.0E+09 & 1.0E-05 & Inf \\
H$_3$O$^+$        & 1655.814 & 6.5E+11 & 0.004 & Inf \\
                  &  984.697 &         &  0.03  & $2\,10^4$\\
H$_2$O            & 1113.343 & 2.0E+15 &   0.46   & 70 \\
                  &  987.927 &         &   0.42   & 95 \\
                  &  752.033 &         &   0.18   & 640 \\
                  &  556.936 &         &   0.84   & 50 \\
                  & 1669.905 &         &   0.18   & 5200 \\
                  & 1716.770 &         &   0.17  & 5000 \\
                  & 1153.126 &         &   0.02  & $4\,10^4$ \\
HDO               &  893.639 & 2.9E+10 &   0.012  & $1\,10^5$ \\
H$_2^{18}$O       & 1101.698 & 3.9E+12 &   0.20   & 350 \\
OH                & 1834.747 &         &   2.0E-4 & Inf \\
CO 10-9           & 1151.985 & 2.1E+17 &   0.017 & $5\,10^4$ \\
\hspace{1cm}15-14 & 1726.603 &         &   0.002 & Inf \\
\hspace{1cm}16-15 & 1841.345 &         &   8.2E-4 & Inf \\
$^{13}$CO 10-9    & 1101.350 & 3.2E+15 &   3.0E-04 & Inf \\
$^{13}$CO 15-14   & 1650.768 &         &   3.8E-05 & Inf \\
\noalign{\smallskip} \hline \noalign{\smallskip}
 OH 1- (119)      & 2509.988 & 6.4E+11 &   0.0024  &  Inf\\
 OH 1+ (119)      & 2514.316 &         &   0.0018  &  Inf \\
 OI 2-1 (63)      & 4745.804 & 1.9E+18 &   0.034   &  2100 \\
 OI 1-0 (145)     & 2060.068 &         &   0.018  &  $4\,10^5$ \\
\noalign{\smallskip} \hline \noalign{\smallskip}
\end{tabular}
\end{table}

The species/lines accessible to HIFI can be used to investigate the specific
problems stated below. It should be noted in this that the likely presence of
a weak, 5-10 km/s, shock along the line of sight adds to the need to observe
this source already on purely energetic grounds alone.

\begin{enumerate}
\item{}The warm layer of H$_2$ gas at the edge of the PDR (through the [CII], [OI], high-J CO and $^{13}$CO, and H$_2$O lines).
\item{}The chemistry of the radical region (through the CH and CH$^+$ lines).
\item{}The chemistry of water in the presence of a soft radiation field (through the H$_2$O, HDO, OH, OH$^+$ and H$_3$O$^+$ lines).
\item{}The electron abundance and its role in the PAH charge balance, for low ionization potentials, and nitrogen chemistry (through the NH$^+$ line).
\end{enumerate}

\subsubsection{Observing Strategy}

Because of its face-on geometry and roughly circular shape, the extents of
the various emission zones are all comparable to one another, and equal
to 0.5'x0.5', since edge effects influence all tracers similarly.
Thus it is modest prudent to do a single long pointing on Ced 201 at the
(0,0) position of the CSO CO 2-1 map with shorter additional pointings, a cross
with four pointings to the north, south, west and east (16 in total), in the
lines of C$^+$ and CH only to explore edge
effects in the radical region caused by the geometry of the PDR.
Only the [OI] 63 $\mu$m line is doable with PACS and takes about 35 min. The
single pointing, including all lines with integration times less than $10^4$
seconds, takes 3.3 hr. The cross of pointings takes approximately 18 min.
The total requested time for this source is then {\bf 3.9 hr} but can be
reduced to {\bf 1.6 hr} when the H$_2$O 1670 and 1717 GHz lines are
omitted.

Additional JCMT and NRAO observations will be performed of
Ced 201 in the lines of CS 2-1, HCO$^+$ 3-2, C$^{18}$O 2-1
and [CI] 609 $\mu$m, to improve on existing limits.

\subsubsection{Overview of Line Predictions}

The model predictions are based on the 3D oblate ellipsoid PDR model of
Kemper et al.\ (1999) with the model parameters given above.
The species, transition and frequency are given in columns 1 and 2.
We assumed a D/H abundance ratio of
$1.52\,10^{-5}$ as found in the ISM by Myers et al. (1996).
%
For the HIFI lines, the assumptions are:
0.7\,km/s resolution (1\,MHz at 500\,GHz), line width of 2.5\,km/s,
position switched observations, including an overhead of 50\%.
In case of the PACS lines, the channel resolution of 180\,km/s is assumed.
The size of the emitting region was assumed to be $0.5'$ for
all lines. The integration times were computed for
 5$\sigma$ detections at the central position.

\subsection{Rosette Molecular Complex\protect\footnote{R. Simon (KOSMA)}}
 
% General

The Rosette Molecular Complex (RMC) is a Giant Molecular Cloud of moderate
density ($\sim 10^3$ cm$^{-3}$) exposed to a comparably low flux of UV radiation
($\sim$100 G$_0$). The RMC is associated with the \HII\ region NGC 2237/2246,
whose optical appearance is responsible for the naming of the region. The 17 OB
stars of the central OB cluster NGC 2244 form the core of the Mon OB2
association and have a total Lyman-$\alpha$ luminosity of $3.8 \times 10^5
L_\odot$ (Cox et al.  1990).  The photometric distance to the Mon OB2
association is estimated to be 1.6 kpc (Turner 1976).  In the immediate vicinity
of the stars, the stellar winds created a cavity. To the south-east, an edge-on
ionization front, the 'Monoceros Ridge' (Blitz \& Thaddeus 1980), marks the
boundary of the cavity at the interface of the \HII\ region to the molecular
cloud.

% Molecular cloud

The full size of the RMC, extending further to the south-east of the OB
association, was first revealed by a $^{12}$CO $J=1-0$ map at $8'$ angular
resolution (Blitz \& Thaddeus 1980). Within a projected area of about 2000
pc$^2$, the gas has a mean large scale \h2\ density of 30 cm$^{-3}$ and a total
mass of $10^5M_\odot$.  The molecular complex shows active star formation,
including at least 34 IR sources of which AFGL 961 is a prominent outflow
source. The total IR luminosity of the complex is $1.1\times10^6 L_\odot$,
estimated by summing over the four IRAS bands (Cox et al. 1990).

High angular resolution observations ($\sim 1'$ down to $12''$) in molecular
lines of CO isotopomers have been carried out with the KOSMA 3m and IRAM 30m
telescopes and combined with [\CII] 158 micron observations ($55''$ resolution)
obtained at the KAO (Schneider et al. 1996, 1998a/b).  Typical clump densities
range from $10^3$ to $10^5$ cm$^{-3}$ (larger scale KOSMA clumps to small clumps
resolved in the IRAM 30m data). While mean \h2\ column densities are a few
$10^{21}$ cm$^{-2}$ (Williams, Blitz, \& Stark 1995), the column density of the
individual clumps in the RMC reaches values of $10^{22}$ cm$^{-2}$ (obtained
from IRAM 30m CO data). The gas temperature, as gleaned from CO observations, is
around 20 K, although likely to be warmer in the \CII\ emitting material at the
UV irradiated clump surfaces and in the interclump medium.

The average density of 300 cm$^{-3}$ from the CO observations yields a high
(30-300) clump to interclump density contrast in the molecular material. In this
case, the UV radiation can penetrate deep into the cloud to create Photon
Dominated Regions (PDRs) on the clump surfaces, which give rise to the observed
ubiquitous \CII\ 158 micron line emission (peak intensity $5\times10^{-4}$ erg
s$^{-1}$ sr$^{-1}$ cm$^{-2}$).  Weak extended \CII\ emission of $10^{-4}$ erg
s$^{-1}$ sr$^{-1}$ cm$^{-2}$ is detected deep in the cloud at a distance of 25
pc from the central OB cluster.  Altogether, the observed intensity
distribution is consistent with the picture of a clumpy molecular cloud, exposed
to a UV flux of around 200 G$_0$ at the interface decreasing logarithmically (as
predicted by PDR models for a weak UV illumination) to a few G$_0$ in the more
shielded parts of the molecular cloud.  The density of the \CII\ emitting gas is
around $10^4$ cm$^{-3}$, compatible with the densities derived for the clumps.

\subsubsection{Ground based preparatory work}

At the KOSMA 3m telescope, we started a campaign to observe the \CI\ lines at
490 and 810 GHz and the CO 4-3 line at 460 GHz. Preliminary results from the
winter 2004 season show that both \CI\ 1-0 and CO 4-3 show strong emission
towards the molecular peak in the RMC. The CO 4-3 line intensities drop from
$\sim 10$ K in the $1'$ KOSMA beam at the molecular cloud peak to $4-5$ K in the
interface region. The high intensities of this mid-$J$ CO line detected in a
rather large beam point towards the presence of warm ($T>100$ K) gas.

\subsubsection{Time estimate and observing strategy}

\begin{figure}[t]
 \begin{center}
 \includegraphics[width=16cm]{rosette-fig.eps}
 \end{center}
 \label{rosette-figure}
 \caption{Contour plot of the velocity integrated CO $J=3-2$ intensity (obtained with
   the KOSMA 3m telescope) overlaid to a grey scale plot of the 158 micron
   [\CII] emission from the KAO (taken from Schneider et al. 1998). The
   Monoceros ridge is indicated by the arrows, which point away from the OB
   cluster. The proposed cut runs north-east from the interface to the peak of
   CO and \CII\ emission indicated by the triangle.}
 \end{figure}
 
The intensities and corresponding integration times in the table below have been
estimated for a column density of $10^{22}$ cm$^{-2}$ and the abundances and
kinetic temperatures given in Table\ref{table_rosette}. The abundances are from
the Cologne PDR model with a UV field of 100 G$_0$ and a number density of
$10^4$ cm$^{-3}$.

\begin{table}[!ht]
\centering
\caption{Estimated integration times for the Rosette molecular
cloud corresponding to a $5\sigma$ detection. A line width of 5 km s$^{-1}$, a
velocity resolution of 0.7 km s$^{-1}$, and $A=0.25$ for position switched
observations was assumed.}
\label{table_rosette}
\flushleft
%{\bf JCMT}: ($v_{\rm LSR}=-4.9$ km s$^{-1}$)\\
\centering

\begin{tabular}{lrccccr}
\hline
  species & freq. & abund. & $T_{\rm kin}$ &  $T_B$ & $t_{\rm tot}$ \\
  & [GHz] & & [K] & [K] & [sec] \\
  (1) & (2) & (3) & (4) & (5) & (6) & (7) \\ \hline
 \CII         & 1900.545 & 0.2E-04 &  300 &     1.447 & 0.3E+02 &  ok    \\
 CH        &  536.761 & 0.2E-07 &  100 &     0.658 & 0.9E+01 &  ok    \\
 CH        & 1656.961 & 0.2E-07 &  100 &     3.758 & 0.5E+01 &  ok    \\
 CH$^+$    &  835.070 & 0.2E-11 &  300 &     0.000 & 0.1E+09 &  weak  \\
 CH$^+$    & 1669.160 & 0.2E-11 &  300 &     0.001 & 0.1E+09 &  weak  \\
 NH          &  974.479 & 0.4E-09 &  100 &     0.071 & 0.1E+04 &  ok    \\
 NH$^+$      & 1012.524 & 0.2E-14 &  300 &     0.000 &     INF &  weak  \\
 NH$_3$      &  572.498 & 0.3E-18 &  100 &     0.000 &     INF &  weak  \\
 NH$_3$       & 1168.452 & 0.3E-18 &  100 &     0.000 &     INF &  weak  \\
 OH$^+$     &  971.804 & 0.1E-13 &  300 &     0.000 & 0.2E+12 &  weak  \\
 H$_3$O$^+$   & 1655.814 & 0.7E-10 &  100 &     0.007 & 0.1E+07 &  weak  \\
 H$_3$O$^+$   &  984.697 & 0.7E-10 &  100 &     0.008 & 0.1E+06 &  weak  \\
 H$_2$O       & 1113.343 & 0.2E-07 &  100 &     8.267 & 0.1E+00 &  ok    \\
 H$_2$O       &  987.927 & 0.2E-07 &  100 &     4.509 & 0.3E+00 &  ok    \\
 H$_2$O       &  752.033 & 0.2E-07 &  100 &     6.012 & 0.1E+00 &  ok    \\
 H$_2$O       &  556.936 & 0.2E-07 &  100 &     6.764 & 0.8E-01 &  ok    \\
 H$_2$O       & 1669.905 & 0.2E-07 &  100 &     8.831 & 0.8E+00 &  ok    \\
 HDO         &  893.639 & 0.1E-11 &  100 &     0.000 & 0.2E+09 &  weak  \\
 H$_2$$^{18}$O & 1101.698 & 0.2E-09 &  100 &     0.118 & 0.5E+03 &  ok    \\
 CO(10-9)     & 1151.985 & 0.3E-03 &  100 &    14.092 & 0.4E-01 &  ok    \\
 CO(15-14)    & 1726.603 & 0.3E-03 &  100 &     4.697 & 0.3E+01 &  ok    \\
 $^{13}$CO(10-9)   & 1101.350 & 0.4E-05 &  100 &     1.578 & 0.3E+01 &  ok    \\
 $^{13}$CO(15-14)  & 1650.768 & 0.4E-05 &  100 &     0.113 & 0.5E+04 &  ok    \\
\hline
 OH       & 2509.988 & 0.1E-09 &  300 &     0.001 & 0.7E+03 &  ok    \\
 OH       & 2514.316 & 0.1E-09 &  300 &     0.001 & 0.8E+03 &  ok    \\
 \OI       & 4745.804 & 0.3E-03 &  300 &     0.833 & 0.2E-03 &  ok    \\
 \OI       & 2060.068 & 0.3E-03 &  300 &     1.333 & 0.2E-02 &  ok    \\ \hline
\end{tabular}
\end{table}

The intensities of the \CII\ KAO data (the only line in the list where we can
compare the predictions of the model with observations) range between 3 and 15 K
assuming the above given mean and peak intensities in a $55''$ beam and a line
width of 5 km s$^{-1}$.  The model prediction for this line thus falls short by
at least a factor of 2 compared to the observations.

For most of the species, the predicted intensities are too low to detect them in
a cut (indicated by {\it weak} in the table), with the exception of the strong
\CII, CH, water, CO, and oxygen lines.

{\bf HIFI:} We propose a fully sampled (i.e., $5''$ at 1.9 THz (\CII)) cut
through the interface, the Monoceros ridge, in the 10 lines predicted to be
strongest in the PDR gas.  The proposed cut is $20'$ long and follows the
direction of the UV photons away from the central stellar cluster. It runs 
through a number of clumps prominent in the IRAM CO and KAO \CII\ data towards
the brightest CO clump associated with the embedded IRAS source IRAS06314+0427
(marked by a triangle in Fig.~\ref{rosette-figure}). Such a cut will take 20
minutes per frequency setting assuming 5s integration time per point including
overheads. A total of {\bf 3.5 hours} observing time is needed to complete the HIFI
cuts.

{\bf PACS:} We propose the same cut to be carried out with PACS in the \OI\ 
lines at 63 and 145 microns. The cut then will be beam-sampled at 63 micron and
take 2.2 hours per frequency setting. The calculation is based on the Line
Mapping mode described in the PACS Guesstimator and is done for the minimum
integration time including overheads. A total of {\bf 4.4 hours} is needed to complete
the PACS cuts.

 We propose to obtain in addition to the cuts two pointed deep integrations
 towards a CO/\CII\ clump at the interface and towards the molecular cloud peak.
 For the deep integrations, we propose 30 minutes per point and three
 frequencies (NH at 974 GHz and H$_2$$^{18}$O at 1101 GHz with HIFI and both OH
 lines at 2510 GHz simultaneously with PACS) resulting in {\bf 3 hours} observing
 time.
 
 The observing time to complete the proposed Rosette observations amounts to {\bf 11
 hours}.


\subsection{S140\protect\footnote{C. Kramer, M. R\"ollig , V. Ossenkopf (KOSMA)}}
 
\subsubsection{Introduction}

The HII region S140 lies at the south-west edge of the molecular cloud
L1204. It is illuminated by the B0V star HD211880 at a distance of
about $7'$ to the cloud. The distance to the star was measured to be
910\,pc. An embedded cluster of infrared sources (IRS1 to 3, Evans et
al. 1989) lies at a distance of about $75''$ to the HII region. Their
luminosities correspond to those of ZAMS B1 to B2 stars.

S140 is a well studied region. Observations of CO and its isotopomers
were carried out by Koester (1998), Johnen (1992), Plume et al.
(1994), Zhou et al.  (1994).  Observations of the CI 1-0 transition
were published by Keene et al.  (1985), Hernichel et al.  (1992),
Minchin et al. (1993, 1994), Plume et al. (1994). Observations of CS
were conducted by Johnen (1992).  Observations of CII and OI
(63\,$\mu$m) were done by Emery et al. (1996). Yongsun \& Youngschul
(1995) mapped several transitions and isotopomers of CO, CS, HCO$^+$,
HCN, SO, SO$_2$, OCS, and HNCO.
C$^{18}$O 2--1 observations indicate a clumpy medium, the mass of the
molecular cloud is estimated to be $840\,$M$_\odot$ (Johnen 1992, Kramer et
al. 1998).
The strength of the impinging UV field is of the order
$10^2$ in units of the Draine field at the position of IRS1 
(Keene et al. 1985, Emery et al.  1996, Timmermann et al. 1996).
The geometry is more or less (Spaans et al. 1997) edge on.
%MG
Large scale maps of the environment of S140 can be found
in Li et al. (2002, ISO data) and Ridge et al. (2003, CO maps).
%MG

Table\,\ref{tab-ii} lists the spectral line observations which were
used as input for modeling S140 as a PDR (cf.  Table 19 of Koester 1998).

\begin{table}[!ht]
\caption{Observed integrated intensities at the position of IRS\,1
($\alpha$=22:17:41.1, $\delta$=63:03:42 J1950). 
The numbers in brackets are the observed linewidths in kms$^{-1}$. 
} \label{tab-ii-s140}
\centering
\begin{tabular}{cccccccc}
\noalign{\smallskip} \hline \noalign{\smallskip}
Position & [CII]$^{c,e}$ & [OI]$^{c,e}$ & [CI]$^{d,f}$ & $^{13}$CO$^{d,g}$ &
$^{13}$CO$^{d,h}$ & $^{12}$CO$^{d,h}$ & $^{12}$CO$^{d,i}$ \\
         &                &              &             & 2--1      &
3--2      & 3--2 & 2--1 \\ 
\noalign{\smallskip} \hline \noalign{\smallskip}
(0,0) & 3.3 & 3.2 & 1.3 & 0.33(3.8) & 2.8(4.4) & 9.7(10.0) & 1.8(8.0) \\ 
\noalign{\smallskip} \hline 
\end{tabular}\\
\flushleft
$^c$ -- units of $10^{-4}$ erg s$^{-1}$ cm$^{-2}$ sr$^{-1}$,
$^d$ -- units of $10^{-6}$ erg s$^{-1}$ cm$^{-2}$ sr$^{-1}$,\\ 
References: $^e$ Emery et al. (1996); $^f$ Plume et al. (1994); $^g$ Johnen (1992);
Schneider (1998); $^h$ Schneider et al. (1995), $^i$ Margulis \& Lada (1985).
\end{table}

Koester (1998) used a two component clumpy PDR model based on the
spherical PDR code described by St\"orzer et al. (1996). The clumpy
model models the emission from an interclump medium with low volume
densities and many small and dense clumps.  The UV field is assumed to
be 125 times the Draine field.  The observed line intensities are
reproduced by clump masses between $10^{-2}$ and $10^{-3}$M$_\odot$ and
clump densities between $5\,10^5$ and $5\,10^6$\,cm$^{-3}$
(Table\,\ref{tab-pdr-s140}).  Such high densities are independently
indicated by high-resolution observations of Jaffe et al. (1989) and
Minchin et al. (1993). See also ammonia observations of Ungerechts et
al. (1986) who derive a density of $5\,10^5$\,cm$^{-3}$. Density
estimates based on C$^{34}$S (Snell et al. 1984, Mundy 1984) and
H$_2$CO ($3\,10^5$\,cm$^{-3}$, Blair et al.  1978) are consistent with
the NH$_3$ value.

\begin{table}[!ht]
\caption{PDR modeling results for the center position of S140 
(adapted from Table 20 of Koester 1998).
The average density is the clump density times the volume filling factor.
The area filling factor is the number of clumps per total area times
the clump projected area. 
} \label{tab-pdr-s140}
\centering
\begin{tabular}{ccc}
\noalign{\smallskip} \hline \noalign{\smallskip}
Clump density             & $10^6$ & [cm$^{-3}$] \\
Clump radius              & 0.0018 & [pc] \\
Clump mass                & $10^{-2}$ & [$M_{\odot}$] \\
Interclump medium density & 100 -- 200 & [cm$^{-3}$] \\
average density           & $1.5\,10^3$ & [cm$^{-3}$] \\
area filling factor       & 2.6 & \\
volume filling factor     & $7\,10^{-4}$ & \\
\noalign{\smallskip} \hline \noalign{\smallskip}
UV-field                  & 100 & Draine \\
H$_2$ column density$^a$ & 21 & [$10^{21}$\,cm$^{-2}$] \\
\noalign{\smallskip} \hline \noalign{\smallskip}
\end{tabular}\\
\flushleft
$^a$ The H$_2$ column density is estimated from
$^{13}$CO 3--2 integrated intensities using LTE formulae and $T_{\rm ex}=31\,$K
from the $^{12}$CO temperature.
\end{table}

%% for (0,0) - output of n13co.f: 
% Calculation of Tex, tau etc.. Outputfile = calc.out
%
% xoff  yoff Tex12co32 tau13co32 ftex   N(13co)            N(H2)      Av
%   "     "    [K]               [E15] [E15/cm^2]          [E21/cm^2] [mag]
%
%   0.0   0.0 30.59     1.03     0.43   45.242    21.264    22.718
%
%% for (15',0')
%xoff  yoff Tex12co32 tau13co32 ftex   N(13co)            N(H2)      Av
%   "     "    [K]               [E15] [E15/cm^2]          [E21/cm^2] [mag]
%
%   0.0   0.0 14.82     0.45     0.65    5.159     2.425     2.590


% ckr: 16.5.03:
\begin{sidewaystable}
\caption{Estimated peak line temperatures at two positions in S140.
Total integration times for 5$\sigma$ detections are given
for the center position.
%
%
The species, transition and frequency are given in columns 1 and 2.
Abundances are derived from the OBSPM model by J.\,LeBourlot and
from the KOSMA model by R\"ollig \& Ossenkopf.
For the HDO abundance, we assumed a D/H abundance ratio of
$1.52\,10^{-5}$ as found in the ISM by Myers et al. 1996.
%
Column 9 gives the size of the emitting region, i.e. the distance
from 0,0 at which peak line temperatures have dropped
to a fifth. (But note that the NH temperature rises with distance!)
%
The time estimates were performed using the KOSMA peak brightness 
temperatures given in column 6.
%
For the HIFI lines, the assumptions are:  
0.7\,km/sec resolution (1\,MHz at 500\,GHz), line width of 2.5\,km/sec,
position switched observations, including an overhead of 50\%.
% This is the width of the optically thin C$^{18}$O 2--1 line observed at the IRAM 30m 
% telescope with $10''$ resolution (Johnen 1992). 
% Optically thick lines are significantly broader (cf. Table\,\ref{tab-ii-s140}).
%
%
In case of the PACS lines, the channel resolution of 180\,km/sec is assumed. 
%
%
} \label{tab-ii}
\centering
\begin{tabular}{lrrrrrrrrr}
\noalign{\smallskip} \hline \noalign{\smallskip}
% dv = 0.7 a=0.25 sigfak=5.0
% FWHM = 2.6
  species & frequ. & \multicolumn{3}{c}{abundances} & \multicolumn{3}{c}{$T_B$ [K]} & D &$t_{tot}$ \\ 
          &        & \multicolumn{3}{c}{} \\
  & [GHz] & \multicolumn{2}{c}{at (0,0)} & at (15,0) & \multicolumn{2}{c}{at (0,0)} & at (15,0) & [$'$] & at (0,0) \\ 
  &       & OBSPM & KOSMA & KOSMA & OBSPM & KOSMA & KOSMA & & [s] \\  
\noalign{\smallskip} \hline \noalign{\smallskip}
C{\sc II}         & 1900.545 & 1.3E-06 & 1.7E-06 & 0.6E-07 &   8.7 & 5.3 &    0.001 &  1.7   & 4\\
CH                &  536.761 & 1.6E-09 & 2.5E-09 & 0.4E-08 &   3.5 & 11.4 &    0.13 &  2.7  & 1\\
                  & 1656.961 &         &         &         &   2.3 & 23.6 &    0.041 & 2.0  & 1\\
CH$^+$            &  835.070 & 2.7E-13 & 1.9E-13 & 0.2E-12 &   0.008 & 0.005 &    $6\,10^{-7}$ & 1.6 & $>10^5$\\
                  & 1669.160 &         &         &         &   0.014 & 0.004 &    $3\,10^{-7}$ & 1.6 & $>10^5$\\
NH                &  974.479 & 9.5E-11 & 1.6E-10 & 0.6E-06 &   0.90 & 2.4 &    0.007 & 2.1  & 1\\
NH$^+$            & 1012.524 & 4.8E-19 & 1.8E-16 & 0.8E-14 &   0.000 & $9\,10^{-7}$ & $8\,10^{-12}$ & 1.5 &  $>10^5$\\
NH$_3$            &  572.498 & 2.1E-08 & 1.4E-10 & 0.6E-10 &   3.1 & 0.18 &    0.002 & 2.7  & 78 \\
                  & 1168.452 &         &         &         &   0.53 & 0.002 &   $2\,10^{-5}$ & 2.7  & $>10^5$ \\
OH$^+$            &  971.804 & 2.0E-13 & 2.8E-13 & 0.3E-12 &   0.004 & $4\,10^{-5}$ &    $6\,10^{-6}$ & 9 & $>10^5$\\
H$_3$O$^+$        & 1655.814 & 4.3E-11 & 2.2E-10 & 0.3E-09 &   0.020 & 0.002 &  0.003 & $>15$  &$>10^5$\\
                  &  984.697 &         &         &         &   0.11 & 0.016 &  0.020 & $>15$  & 17000\\
H$_2$O            & 1113.343 & 2.0E-08 & 3.9E-07 & 0.2E-04 &   4.0 & 2.4 &    1.5 &  $>15$  & 1\\
                  &  987.927 &         &         &         &   4.9 & 1.0 &    0.44 & $>15$  & 5\\
                  &  752.033 &         &         &         &   7.1 & 0.76 &    0.21 & $>15$  & 6\\
                  &  556.936 &         &         &         &   9.5 & 3.4 &    4.5 &  $>15$  & 1\\
                  & 1669.905 &         &         &         &   1.5 & 1.5 &    0.32 & $>15$  & 46\\
HDO               &  893.639 & 2.4E-13 & 5.9E-12 & 0.0E+00 &   0.087 & 0.012 &   0.02 & $>15$ & 27000\\
H$_2^{18}$O       & 1101.698 &         & 8.2E-10 & 0.5E-07 &   4.0 &   0.31 &  0.35 & $>15$& 51\\
CO 10-9           & 1151.985 & 4.6E-05 & 2.8E-04 & 0.3E-03 &   0.014 & 0.023 & $5\,10^{-5}$& 2.0 & 9600\\
CO 15-14          & 1726.603 &         &         &         &   0.000 & 0.003 & $9\,10^{-6}$ & 2.1  & $>10^5$\\
$^{13}$CO 10-9    & 1101.350 & 7.7E-07 & 4.2E-06 & 0.4E-05 &   0.000 & $4\,10^{-4}$ & $9\,10^{-7}$ & 2.0 & $>10^5$\\
$^{13}$CO 15-14   & 1650.768 &         &         &         &   0.000 & $5\,10^{-5}$ & $1\,10^{-7}$ & 2.1 & $>10^5$\\
\noalign{\smallskip} \hline \noalign{\smallskip}
 OH 1- (119)      & 2509.988 & 1.4E-11 & 0.1E-08 &         &   0.24 & 0.28 &    0.0014  &      2.3    & 108 \\ 
 OH 1+ (119)      & 2514.316 &         &         &         &   0.23 & 0.22 &    0.0013  &      2.3    & 174 \\ 
 OI 2-1 (63)      & 4745.804 & 1.1E-05 & 0.1E-03 &         &   0.4 & 0.79  &   $2\,10^{-4}$ &    1.7   & 3 \\ 
 OI 1-0 (145)     & 2060.068 &         &         &         &   0.22 & 0.17  &   $2\,10^{-5}$  &  1.7    & 780 \\ 
\noalign{\smallskip} \hline \noalign{\smallskip}
\end{tabular}
\end{sidewaystable}


\subsubsection{Modeling the emission}

Two PDR models have been used to calculate the chemical abundances:
the OBSPM model of J.LeBourlot and the KOSMA spherical model.  The
resulting abundances are shown in Table\,\ref{tab-ii}.  They do agree
well for some prominent species like [CII], but show discrepancies of
up to a factor of 300 for minor species like NH$^+$. This comparison is
meant to give a first rough idea of the uncertainties inherent in
current state-of-the-art models. This model comparison still needs to
be supplemented by more observational data. However, it is out of the
scope of this proposal, to discuss this comparison in any detail.

A full-fledged radiative transfer calculation to derive the emitted
intensities can be done only for a few species like [CII], [CI], CO,
which are predicting to have strong lines.
However, this is not easily done for the minor species of which
partition functions are much less well known. LTE integrated
intensities were derived using the LTE formulae accounting for optical
depth effects given above. Bright lines  with more than 1\,K peak temperature are
those of water, H$_2^{18}$O, %MG, 
CH, CH$^+$, NH and maybe NH$_3$.

\paragraph{Center position:}

For the center position, we used the total column density, clump
density, and UV-field of 100 listed in table\,\ref{tab-pdr-s140}, to run the
two PDR models. 
Table\,\ref{tab-ii} lists the resulting observing times for the S140
center position. The [C{\sc II}], CH, and the water lines are relatively
bright with line temperatures of a few K. The weaker lines of NH$_3$,
NH, and H$_2^{18}$O are also well observable with reasonable integration times.
If the Le Bourlot PDR model is used, also the H$_3$O$^+$ line at 984 GHz
and the HDO line are detectable within less than 300~s.
The lines of CH$^+$, NH$^+$, OH$^+$, and the two CO isotopes, however,
would require prohibitively long observing times, so that they are
ignored for this PDR. Summing up the observations of all detectable lines
at the center position results in a total observing time of 15\,minutes.

\paragraph{Offset position:}
A second estimate of intensities and observing times was made for an
offset position at (15',0) well inside the cloud where the UV field
has presumably dropped to 1 and the column density is only 10\% of the
peak column density at the center position. 

The resulting abundances and peak line temperatures are also shown in
table\,\ref{tab-ii}. Surprisingly, the NH line is predicted to become
brighter than at the 0,0 position. All
other lines show a drop in intensities relative to the center
position. Compared to the center position the observation of
HDO is not expected to be detected any more and most other lines need a
considerably longer integration time for detection. any signal Other bright lines with more than 1\,K peak temperature are
those of water, H$_2^{18}$O, and CO.any signal Other bright lines with more than 1\,K peak temperature are
those of water, H$_2^{18}$O, and CO.The CII line is predicted to have
only 1\,mK, much less than the observed values which correspond to
2.6\,K (Koester et al.). That may indicate that the UV field is
actually stronger than assumed at this offset position. 
Excluding the very weak lines, the remaining lines need a total
integration time of 125\,minutes.

% in stark contrast to the observations which only show a drop by a
% factor of 5.5 (Koester et al.) which would translate into peak
% temperatures of 2.6\,K. This indicates that the UV field is probably
% significantly higher than estimated at the $15',0'$ position. A longer
% cuts is probably needed to sample the regions of low UV fields of
% about 1.
%MG
Observations of the complete set of lines for the two positions
(center and 15'E) therefore require a total of {\bf 140 minutes = 2.3 hrs}
%MG

\paragraph{Rough stepping for 5 transitions:}

The NH$_3$ 572 GHz, the H$_3$O$^+$ 984 GHz, the H$_2$O 1669 GHz,
the HDO, and the
H$_2^{18}$O lines need relatively long integration times per point,
so that full OTF scans are too time consuming. Here, the two 
stripes observed for other species in OTF are to be covered by 10
points extending over 5' in the central region.
This requires {\bf 4.3~h} observing time.

\paragraph{Cut for 8 transitions:}

It takes 15\,minutes, to conduct one cut of $15'$ length, sampling ever
$5''$ (i.e. fully sampling at 1.9\,THz), and every 5\,sec. We plan to
conduct 2 cuts per frequency for the 8 brightest transitions. Taking
into account that in one case two lines can be observed with one
LO setting we arrive at a total observing time of {\bf 3.5~h}.

\paragraph{PACS observations:}

We plan to observe the four OH and OI lines along the same cuts
planned for HIFI. As the extent of the emission in those lines is
expected to be small, only the inner 5$'$ are to be covered.  We
expect the PACS observing times to be entirely driven by the
relatively large overheads. To map the OI 63$\,\mu$m line at beam
spacing (HPBW=$4.7''$), we need 37\,min using the PACS Guesstimator
(21/10/02). Here, we take into account that the 4 settings are needed
to half-sample the data due to the PACS pixel size of $9.4''$.
Complementary observations of the weaker, but optically thin OI
145\,$\mu$m line (HPBW=$10.8''$) at beam spacing need about a factor
of 4 less time, 12\,minutes. The same amount is need for each OH
line. We thus ask for 1.2\,hours of PACS time for all four lines.

\paragraph{Total observing time: 11.3\,hours.}
The center together with the offset position together will take
2.3\,hours. The rough stepping needs 4.3 hours and the OTF cuts
take 3.5 hours. PACS observations add another 1.2 hours. Thus ask for a total of
11.3\,hours for Herschel observations of this template PDR.  (Note that a
factor of 2 overhead relative to the pure radiometer formula is
already included.)

\paragraph{} An improved understanding and modeling of this template PDR,
will obviously lead to refinements of this preliminary, initial
time estimate.

\subsection{The Horsehead nebula\protect\footnote{M. Gerin LERMA, D. Teyssier SRON}}
%, L. Verstraete IAS, E. Habart Arcetri}}

\subsubsection{Introduction}

In the visible, the Horsehead nebula (B33, Barnard 1919) appears as a
dark patch seen in silhouette against the ionized H$\alpha$ emission
emanating from the IC434 H{\sc ii} region. The large-scale maps
reported by e.g.  Maddalena et al. (1986, CO(2--1)) or Lada et
al. (1991, CS) show its connection to the L1630 cloud, which belongs
to the Orion B molecular complex (see e.g. Abergel et al. 2002 for an
overview of the region).  In the scenario proposed by Reipurth \&
Bouchet (1984), the B33 dark cloud is assimilated to an early Bok
globule emerging from its parent cloud (L1630) via the eroding incident
radiation field emitted by the close-by O9.5V star $\sigma$Ori. The
Horsehead nebula thus corresponds to a condensation resisting the
illumination and it presents on its western edge a PDR. Assuming a
distance of 400\,pc (Anthony-Twarog 1982), and considering a projected
distance to $\sigma$Ori of 0.5$^\circ$, the incident radiation field
is of order $\chi \sim 60$ (Zhou et al. 1993, Abergel et
al. 2003). With its small distance, moderate size (4\am$\times$4\am)
and active chemistry, the horsehead nebula is a moderately excited
PDR. Optical plates of this nebula shows emission streaming away from
the interface suggesting that photoevaporation of the cloud skin is
underway.  The photodissociated interface as seen in the mid-IR by
ISOCAM (6\as~resolution) in filters around 7 and 15\,$\mu$m (thin
filament of 10\as~wide or 0.02 pc) appears homogeneous down to scales
below 0.01 pc (Abergel et al. 2002, 2003). This very narrow filament
suggests that the PDR is seen close to edge-on ($<$ 10\,degrees).

In the CO and isotope maps of Kramer et al. (1996), the nebula is
heavily diluted in the 2\am~beam and only appears as a small intrusion
of L1630 into the H{\sc ii} region. These large-scale data already
indicated mean densities of order $3\times10^{4}$\,cm$^{-3}$ (Zhou et
al. 1993 from CS lines) and column densities compatibles with $A_{\rm
V}\sim 3$ (Kramer et al. 1996). Observations at higher spatial
resolution (10-20\as) were only reported very recently.  Pound et
al. (2003) presented BIMA CO(1--0) and discussed in details the
possible origin and evolution of the dark cloud. They reveal a complex
velocity structure but their data which could be affected by
non-negligible optical depth effects. Abergel et al. (2003) report
IRAM 30-m maps obtained in the ($J$=1--0) and ($J$=2--1) transitions
of CO, $^{13}$CO and C$^{18}$O. They infer densities compatible with
the values reported at larger scale, and show clear selective
photo-dissociation of the considered species at the PDR edge. Zhou et
al. (1993) also present [CII] observations obtained with the
KAO. Teyssier et al. (2004) and Pety et al. (2004) have obtained maps
of hydrocarbon molecules in the PDR, showing that this source harbors
a rich carbon chemistry, in close connection with the PAH
distribution. Large scale and high resolution submillimeter maps of
the horsehead nebula can also be found in Teyssier (2002 - PhD thesis,
dust continuum at 1200, 850 and 450\,$\mu$m) and Lis et al. (2003)
(350\,$\mu$m continuum, C and CO(4-3), with CHAMP at the CSO), showing
that large extinction regions are concentrated in the UV illuminated
edge, and in a dense filament joining the PDR and the L1630 cloud.  A
near infrared image (H$_2$ v=1-0 S(1)) is shown in Habart (2002, PhD
thesis). These data, together with molecular emission of CO and
isotopes, has allowed to derive the density profile at the PDR border,
with a total hydrogen density n$_{\rm H}$ of 2$\times$10$^4$\,cm$^{-2}$ at the H$_2$ and PAH
peak, and a peak density of 10$^5$\,cm$^{-2}$ in the molecular cloud
area (Habart et al., in preparation).

%%%%%%%%%%%%%%%%%%%%%%%%%%%% FIG HH %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%
\begin{figure}
 \centering
 \includegraphics[width=9cm,angle=270]{hh_strat.eps}
 \caption{Positions aimed at being observed in the Horsehead nebula, overlaid over an integrated CO(3--2) map (Gerin priv. comm.). In white squares are indicated the three positions to be probed in deep integrations, while the black squares indicate the position of the two cuts. An arrow shows the direction of the illuminating star $\sigma$\,Ori.}
 \label{HH obs strategy}
\end{figure}
%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% 

It will therefore be interesting to pursue the study of the gas
chemistry and dynamics in this region with new chemical
tracers. Further on, the study of the relationship of the gas species
with the dust emission will highlight the dominant gas-grain couplings
and their impact on the gas dynamics.


\subsubsection{Line intensities}

Line intensities have been estimated at LTE for a column density of
$10^{22}$ cm$^{-2}$, a line width of 1 km/s and the abundances given
in Table \ref{tab-time-hh}. The excitation temperatures and
abundances have been estimated from a PDR model with $\chi=60$ and
$n=10^4$ cm$^{-3}$.

\subsubsection{Observing strategy and time estimates} 

In order to study the stratified emission and advection motions, we
propose to observe two cuts across the interface, 4\am~long with
10\as~pixel spacing, so 24 positions each. These cuts are illustrated
on Fig.~\ref{HH obs strategy}. They cross two characteristic peaks in
this source (Teyssier et al. 2004), namely the IR peak
($\alpha$= 05h40m53.7s, $\delta$=-02\deg28\am00\as, J2000), located at
the edge of the PDR and coincident with, among others, emission peaks
in H$_2$ and in several hydrocarbon molecular lines, and the CO peak
($\alpha$=05h40m53.7s, $\delta$=-02\deg26\am50\as, J2000), located
more inside the cloud, about 1\am~north to the IR peak. These cuts
will be tilted by 14\deg~clockwise with respect to the equatorial
frame in order to be aligned with the direction of the radiation field
emitted by $\sigma$\,Ori. In order to span all possible phases and
physical condition across the ionized, neutral and molecular layers,
the cut will extend from about 30\as~outwards of the cloud, to
3.5\am~inwards the molecular cloud. In the case of the IR-peak cut,
the observed positions will also probe an isolated emission
peak located well inside the molecular cloud, and seen both in
continuum and molecular emission (e.g. Lis et al. 2004).
 
Observations of weaker lines will be restricted to 3 positions, the IR
peak, the CO peak and the cold dust peak ($\alpha$=05h41m07s,
$\delta$=-02\deg27\am30\as, J2000) as e.g. traced by continuum at
1.2\,mm. In all cases, an OFF position located 5 arcmin west will be
used for Position Switched observations. C[II] shall be observed in
Frequency Switching due to the expected very extended emission.

%%%%%%%%%%%%%%%%%%%%%%%%%% TAB HH %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%
\begin{table}[!ht]
\caption{Time estimate for the Horsehead IR peak position. We assume 
FWHM=1 km/s and a velocity resolution of 0.3\,km/sec, $A=0.125$ for
position switched observations (respectively $A=0.25$ for frequency
switched observations, only considered for C[II]), and a SNR of 10 (5) for strong(weak) lines. The OH and O lines are
observed by PACS. Column (1) is the species, column (2) gives the
frequency of the transition, column (3) gives the abundance, column
(4) gives the excitation temperature, column (5) the brightness
temperature and column (6) the observing time in seconds with the achieved SNR in parenthesis.}
\centering
\label{tab-time-hh}
\begin{tabular}{lrrrrr}
\noalign{\smallskip} \hline \noalign{\smallskip}
  species & $\nu$ & [X/H] & $T_{ex}$ & $T_B$ & Time (SNR)\\
  & [GHz] & & [K] & [K] & [sec.] \\
\noalign{\smallskip} \hline \noalign{\smallskip}
 C{\sc II}     	& 1900.545 &         & 50   	&  10.0${}^a$ 	& 6 (10) \\
 C		&  492.161 &         &    	&  10.0${}^b$ 	& 0.7 (10) \\
		&  809.342 &         &    	&  10.0${}^c$	& 1 (10) \\
CH            	&  536.761 & 7.2e-09 & 10   	& 1.7 		& 24 (10) \\
	       	& 1656.961 & 7.2e-09 & 10   	& 0.01 		& ... \\
 CH$^+$   	& 835.070  & 3.6e-09 & 10  	& 0.73 		& 200 (10) \\
		&1669.160  & 3.6e-09 & 10  	& 0.004 	& ... \\
 NH   		& 974.479  & 5.4e-09 & 10   	& 0.43 		& 172 (5) \\
 NH$_3$   	& 572.498  & 2.4e-07 & 10   	& 1.9 		& 21 (10) \\
 		& 1168.452 & 2.4e-07 & 10   	& 0.21 		& 878 (5) \\
 H$_3$O$^+$   	&  984.697 & 4.2e-10 & 10   	& 0.07 		& 1800$^d$ (5)\\
 	   	& 1655.810 & 4.2e-10 & 10   	& 0.01 		& ... \\
 H$_2$O  	& 1113.343 & 5.0e-07 & 10   	& 0.25 		& 542 (5)  \\
  		&  987.927 & 5.0e-07 & 10   	& 0.4 		&  181 (5) \\
		&  752.033 & 5.0e-07 & 10   	& 0.2 		& 880 (5) \\
                &  556.936 & 5.0e-07 & 10   	& 2.0 		& 18 (10) \\
  		& 1669.905 & 5.0e-07 & 10   	& 0.03 		& ... \\
 H$_2{}^{18}$O  & 1101.698 & 1.0e-09 & 10   	& 0.25 		& 654 (5) \\
                & 547.676  & 1.0e-09 & 10   	& 1.7 		& 24 (10) \\
CO       	& 1151.98  & 3.0e-04 & 50   	& 27.3		& 0.2 (10) \\
 		& 1841.34  & 3.0e-04 & 10   	& 0.1		& ... \\
 OH  		& 2509.988 & 3.2e-09 & 12   	& 7.0e-10 (!)	& ... \\
 \noalign{\smallskip} \hline \noalign{\smallskip}
 OH${}^f$  	& 2509.988 & 3.2e-09 & 12   	& 5.9e-15${}^e$ & 3419 (5) \\
 OH${}^f$  	& 2514.316 & 3.2e-09 & 12   	& 5.8e-15${}^e$ & 3686 (5) \\
  O  		& 4745.804 & 1.7e-04 & 25   	& 1.1e-13${}^e$ & 108 (8) \\
    		& 2060.068 & 1.7e-04 & 25   	& 1.1e-12${}^e$ & 108 (176) \\
\noalign{\smallskip} \hline \noalign{\smallskip}
\end{tabular}
\flushleft
${}^a$ Zhou et al. (1993)
${}^b$ Lis et al. (2004) 
${}^c$ Assumed from Lis et al. (2004) 
${}^d$ The LTE excitation temperature is probably underestimated
and we assume it can be detected in 0.5 hour.
${}^e$ in erg.cm$^{-2}$.s$^{-1}$, assuming the PACS resolution of
180\,km/s and a 5$\sigma$ sensitivity of
2.8e-15\,erg.cm$^{-2}$.s$^{-1}$ in one hour,
${}^f$ Observed simultaneously
\end{table}

\vspace{2mm}
Time estimates for each molecule are given in Tab.~\ref{tab-time-hh}. For 24 positions and two cuts in the brightest lines ([CII], [CI], CH, CH$^+$,
H$_2$O, H$_2^{18}$O, NH$_3$, CO) we need a total of {\bf 4.9 hours}.  %%1 position is 365 sec. for all lines
The deep integration (3 positions) require {\bf 3.7 hours}.  %%1 position is 4394 sec. for all lines
PACS observations add another {\bf 3.2 hours}.   %%1.06h per position
The total is therefore {\bf 11.8 hours}.

\subsection{Diffuse interstellar clouds -- $\zeta$ Oph, HD 34078\protect\footnote{M.Gering (LERMA)}}

The chemical properties of diffuse clouds such as those lying along
the line of sight to the bright O9 star $\zeta$ Oph are well known
from a continuous effort on visible and UV spectroscopy during the
past 40 years. Using the same technique  in the radio
domain, with distant radio sources as background objects, Liszt and
Lucas (1996 to 2004) have measured the abundance of  molecules,
not easily studied in the UV/visible wavelength range, but there is little
overlap so far between the lines of sight studied in the UV/visible or in
the radio domains. While absorption studies in the UV and visible 
are extremely powerful
 in detecting ionic, atomic or molecular species in the various ISM phases,
 they are limited in spectral resolution -- the best spectral
 resolution for HST data is
 about 3 kms$^{-1}$ -- and provide no information
 on the  spatial structure.  Nevertheless,  absorption studies
provide  unique information
on the abundances and ionization state of numerous IS ``components''
i.e. defined by a common absorption velocity and Doppler broadening parameter.
Typically, some of the components are
neutral and clearly associated with either the CNM or the WNM, and
some are at least
partially ionized, the origin of the ionization being not always clear.
The components with the largest gas column density are mostly
molecular. Because of the broad wavelength coverage, and 
the high sensitivity of
the FUV to visible high resolution spectroscopy, a wealth of data on
the physical conditions and molecular abundances is now available for
selected lines sight. For instance, the HI and H$_2$ column densities
from either Copernicus or FUSE data (see Rachford et al. 2002) for
FUSE, gas phase elemental abundance from the STIS surveys (.g. Cartledge
et al. 2003, 2001 for O and Kr, Sofia et al. 2004 for C), and
sensitive molecular column densities (H$_3^+$, CH, CN, C$_2$, C$_3$,
OH, ...) visible and near-IR spectroscopy. The physical conditions
are derived from these extensive data sets using state of the art
PDR models. However, the density structure and source geometry,
which are known to be important for the model results, are
not constrained by these observations. To
 progress further, information on the spatial structure is critical.  
Previous attempts of comparing radio and UV data have used
ground based telescopes for
CO and CN. The outcome can be summarized as follows : despite the
difference in spatial resolution ($\sim 20''$ for emission lines
versus $\le 0.1''$ for absorption), column densities 
derived from radio data are similar to UV column densities 
(see e.g. the work on the diffuse clouds towards $\zeta$ Oph by
 Langer et al. (1987) \& Keene et al. (1987)).  In addition to the spatial
structure  and the line profile, radio data are useful for
constraining the excitation conditions (temperature, density or a 
combination of these parameters). Since HIFI can
observe important species (C, C$^+$) also sampled in the 
UV, we wish to take this opportunity and gather additional information 
on well studied lines of sigh to be able to combine
efficiently  the information acquired on  powerful ground based
and  space telescopes with the future data from HIFI.
The selected diffuse clouds correspond to the low end both
in density and radiation field of PDR parameter space. 

We have selected two lines of sight with among the most 
complete information (H and H$_2$ from FUV telescopes, high spectral 
resolution HST data with STIS or GHRS, and molecular lines
from ground based telescopes.
The  CO rotational lines at 2.6mm and 1.3mm  have been detected too.

\begin{table}[!ht]
\caption{Diffuse interstellar clouds}
\centering
\begin{tabular}{lrrl}
\hline
Name & RA(2000) & DEC(2000) & Comment\\
\hline
$\zeta$ Oph & 16 37 09.5378 & -10 34 01.524 & Star + 1 positions \\
HD~34078 & 05 16 18.1497 & +34 18 44.341 & FUSE warm H$_2$ \\
\hline
\end{tabular}
\end{table}


\subsubsection{$\zeta$ Oph}

 $\zeta$Oph (HD~149757) is a runaway O star (escaped from the Sco OB2
association) located a the distance of 
about 140 pc near the Ophiuchus dark cloud complex (Blaauw 1961, Draine
1986). It is surrounded by an extended HII region, S27, of $\sim 7^\circ
\times 10^\circ$ ($\sim 18 \times 26$ pc
pc) nicely traced on the H$\alpha$ plate (Sivan 1974). Extinction
features can be identified on this H$\alpha$ plate, some of which
coinciding with the dark clouds complexes L~156 and  L~204.
The CO emission from these  cloud complexes has been mapped
at large scale and extends over about one
degree north  and 3 degrees south of $\zeta$ Oph (Tachihara et
al. 2000), the star itself lies in minimum of CO emission. A few
arcminutes south of $\zeta$ Oph, the CO
emission is moderately bright $\sim 1 - 5 $K and exhibits
a complex velocity structure (Fig \ref{fig:zeta-oph}). 

The diffuse clouds along the line of sight have been
studied so extensively at all wavelengths,  that all molecules
detected in absorption in the UV-visible-near IR wavelength range 
have been detected. previous studies have also been
performed at millimeter wavelengths by Langer et al. (1987), Kopp et
al. (1996),  Liszt (1997). The ground state fine structure line of atomic carbon 
has also been detected with the KAO (Keene et al. 1987).
The CO emission from the diffuse clouds has been mapped
at large scale (Fig. \ref{fig:zeta-oph}), and extends over about one
degree north  and 3 degrees south of $\zeta$ Oph (Tachihara et al. 2000), the star itself lies in
minimum of CO emission.  The structure can be understood in
the framework of PDR models where the increased line intensity away
from the star traces the larger CO column density further
away from the source of dissociating radiation (e.g. Kopp et al. 1996).

\begin{figure}
\begin{center}
\resizebox{12cm}{!}{
\rotatebox{270}{
\includegraphics{zoph_chan.eps}}}
\end{center}
\caption{\label{fig:zeta-oph}CO(2-1) channel maps of the diffuse
clouds south of $\zeta$ Oph obtained with the CSO with 30" spacing.
The coordinates are given in B1950. $\zeta$ Oph is
at 16h34m24s,-10$^\circ$28'03". (D. Lis, priv. comm.)}
\end{figure}

With a typical extinction of 1 magnitude,  line width of 
0.5 -- 1 kms$^{-1}$ and gas density of $n_H \sim 1000 cm^{-3}$ as
derived from the FUV and visible data,
we expect C and C$^+$ column densities of 10$^{15}$  and 
10$^{17}$ cm$^{-2}$ respectively, and line
intensities of $\sim$ 0.1 K for the 1-0 [CI] line (Keene et al. 1987)
and $\sim 0.3$K for [CII] (Ingalls et al. 2002).

Since the source is extended, observations will be performed in 
frequency switching. We propose to sample the spatial
structure by mapping two fields, one centered on $\zeta Oph$ itself
and one further away, having stronger CO emission. 
We have chosen  to map $1' \times 1'$ around $\zeta Oph$ and
to perform a cut of $0.5' \times 3'$ for the CO position.
 This represents 
36 points for [CII], and 9 for the ground state  carbon line. We shall
smooth the [CII] data by binning 4 positions together. The time
estimates takes the spatial smoothing into account.
For the weak
excited carbon line at 809 GHz, we shall observe
the central point only.
The program therefore requires  integration time per
position of $\sim 15$ minutes for [CI](1-0)  (20 mK rms using 2 mixers) 
and $\sim 15$ minutes for [CII] (11 mK rms, 1 mixer ) with a
frequency resolution of 200 kHz and 1 MHz respectively (0.15
kms$^{-1}$). The foreseen integration time for the excited [CI] line
is one hour (rms 8 mK).
The total time request is therefore ,
$2(9 \times 0.25 + 36 \times 0.25 + 1)  = ${\bf 24.5 hours} for the [CII] and [CI]
observations.  

\begin{table}[!ht]
\caption{Observing time for $\zeta Oph$ clouds}
\centering
\begin{tabular}{lrrrrrr}
\hline
Line & $\delta \nu$ & $\delta v$ & $\sigma$ & time/position & Positions & total 
time \\
     & MHz & kms$^{-1}$ & mK & minutes & & hours \\
\hline
C($^3P_1 - ^3P_0$) & 0.2 & 0.12 & 20 & 15 & 9 & 2.25 \\
C($^3P_2 - ^3P_1$) & 0.32& 0.12 & 8  & 60 & 1 & 1.0 \\
C$^+$ & 1.0 & 0.16 & 100 & 15 & 36 & 9.0 \\
\hline
{\bf Total per field} & & & & & & {\bf 12.25} \\
\hline
\end{tabular}
\end{table}

\subsubsection{ HD~34078}

HD~34078 (or AE Aur) is a reddened runaway O star with a large
transverse velocity (103 km/s). The line
of sight to this target drifts through the foreground molecular cloud at
a rate of 17 AU/yr (HD 34078 is located at 500pc while the cloud is
assumed to lie at 400 pc). The source has been selected by Boiss\'e
and collaborators for studying the small scale density
structure using FUV H$_2$ lines. By comparing H$_2$ absorption lines
in  spectra taken at several epochs this methods enables,
 for the first time, the study in a
very direct way of the {\it density} structure in the foreground gas over the
range 5 - 50 AU (parallel observations are also done in collaboration
with Steve Federman to study the structure in CH and CH$^+$: Rollinde,
2003: PhD Thesis; Rollinde et al. 2003). 

The latest results, based on 5 FUSE spectra taken over 3
years, indicate that fluctuations of N(H$_2$) are no larger than 5\%,
i.e. points towards very little density structure. The observed
properties for this molecular gas (temperature, chemical composition,
etc) suggests that the intervening
cloud is far away from HD~34078 and unaffected by it.
The FUSE observations have also revealed an unexpected large amount
of highly excited H$_2$, which can be understood as the result of the
presence of dense gas (n $\simeq 10^4 cm^{-3}$) located in a bow shock
around HD~34078 and strongly illuminated by this bright O star (cf the
scenario proposed by van Buren \& McCray 1988). This
is consistent with the detection by IRAS of an IR excess associated
with HD~34078 (van Buren et al. 1995). Detailed
modeling of the PDR 
resulting from this interaction successfully reproduces the observed 
H$_2$ excitation diagram (Le Petit et al. 2001, Le Petit, 2003 (PhD
Thesis), Boiss\'e et
al. 2004, in preparation).
This second component accounts for the highly excited H$_2$ but
represents a small fraction of the total H$_2$ gas on the line of
sight. The bulk of the H$_2$ is, in our model, comprised in a
"classical" translucent cloud located far away from HD~34078 and
illuminated by the general interstellar radiation field. This
quiescent material accounts for (most of) the J = 0 and J = 1 H$_2$ 
absorptions and for the CH, CH$^+$, CN (Rollinde et al. 2003), 
CO (IUE: Mac Lachlan \& Nandy 1984) features.

\begin{figure}
\begin{center}
\resizebox{9cm}{!}{
\rotatebox{270}{
\includegraphics{map-hd34078.eps}}}
\end{center}
\caption{\label{fig:hd34078}12CO(2-1) map towards HD34078 obtained
with HERA on the IRAM 30m telescope (P. Boiss\'e, priv. comm.)
HD3~4078 lies at the map center. Both velocity components are
seen in the visible CH absorption spectrum.}
\end{figure}

There are however some observed properties of the gas that are not
well explained by this two-component model (large CH/H$_2$ ratio, absence of
velocity shift between the two components ...) and which suggests that 
the foreground cloud may in some way be related to HD~34078. To
investigate this possibility, a sensitive CO(2-1) map
has been obtained with HERA on the IRAM-30m telescope (Fig. \ref{fig:hd34078}).
It indicates that: 
i) CO emission is easily detectable ($T_A^*$ = 0.5 K), ii) two
velocity components are present in good agreement with those seen in
high resolution CH visible observations (Rollinde et al. 2003) and iii)
A complex velocity and intensity structure tracing the interaction
of the star with the diffuse cloud.

\paragraph{Proposed Observations} 

We propose to map the interaction region in [CII] and [CI], and
to perform a deep integration in CH, CH$^+$ and H$_2$O. Because the extend
is very small, we propose to cover $40" \times 40"$, or 4 points
in [CI] and 16 in [CII]. We add one hour integration for each
CH, CH$^+$ and H$_2$O, which will help constraining the
combined PDR+shock model. Observations will be performed in 
frequency switching. The line strength is expected to be
similar as for the $\zeta$ Oph line of sight, $\sim 15 minutes$ for
[CI] and [CII]. This sums to a total
of $4 \ times 0.25 + 16 \times 0.25 + 3 + 1 = 9 hours$ of observations.

\begin{table}[!ht]
\caption{Observing time for HD~34078}
\centering
\begin{tabular}{lrrrrrr}
\hline
Line & $\delta \nu$ & $\delta v$ & $\sigma$ & time/position & Positions & total 
time \\
     & MHz & kms$^{-1}$ & mK & minutes & & hours \\
\hline
C($^3P_1 - ^3P_0$) & 0.2 & 0.12 & 20 & 15 & 4 & 1.0 \\
C($^3P_2 - ^3P_1$) & 0.32& 0.12 & 8  & 60 & 1 & 1.0 \\
C$^+$ & 1.0 & 0.16 & 100 & 15 & 16 & 4.0 \\
CH $^2\Pi_{3/2} - ^2\Pi_{1/2}$  & 0.2 & 0.12 & 10 & 60 & 1 & 1.0 \\
H$_2$O $1_{10}-1_{01}$ & 0.2 & 0.12 & 10 & 60 & 1 & 1.0 \\
CH$^+$ 1 - 0 & 0.4 & 0.14 & 7.4 & 60 & 1 & 1.0 \\
\hline
{\bf Total} & & & & & & {\bf 9.0} \\
\hline
\end{tabular}

\end{table}

The total combining $\zeta$ Oph and HD34078 is therefore {\bf 33 hours}.

\subsubsection{Preparatory observations :}

As the selected targets have been extensively observed in the past,
the necessary information is available in the literature. This
includes : \\

\begin{itemize}
\item{Ground based CO (mm), CH, CH$^+$, C$_2$, C$_3$, OH, H$_3^+$
(visible, near IR)}
\item{Space FUV H$_2$ (FUSE, Copernicus), C, C$^+$, CO (HST STIS or GHRS)}
\end{itemize}

An observing program has been started with HERA on the IRAM 30m
telescope.

\subsection{SNR/Molecular Cloud Interaction: IC443, W28, 3C391,
and W44\protect\footnote{R. G\"usten, S. Philipp, D. Neufeld}}

\subsubsection{Introduction}

In the previous sections we have argued, that we have reasonably
good theoretical understanding how a SNR/blast wave may be interacting
with the parent molecular cloud, but solid observational evidence
is widely lacking. We propose to carry out HIFI observations of
H$_2$O, high-J CO and thermal OH emissions, supplemented by targeted
FIR observations with PACS, towards four molecular clouds that are
interacting with a SNR: IC443, W28, 3C391 and W44.

These sources all exhibit common observational characteristics
(Reach and Rho 1998, 2000, Snell et al. in prep.), like prominent atomic
fine structure emissions from a variety of neutral, and singly- and doubly-
ionized species (including CII, NII, NIII, OI, OII, OIII, SiII, PII, and
FeII), emission from high-excitation rotational ($J$ up to 15) and
rovibrational ($v=1$) states of H$_2$ and CO (e.g. $J = 16-15$),
maser emission in the 1720 MHz lambda-doubling transition of OH, and notably,
reduced H$_2$O abundances [n(H$_2$O)/n(CO) $\le 0.4$] relative to other
regions in which the presence of warm H$_2$ has been attributed to shock
waves.

\begin{figure}[!ht]
\setlength{\unitlength}{1mm}
\begin{picture}(200,140)
\special{psfile=ic443_co21_h2.ps
angle=-90 hoffset=0 voffset=440 hscale=78 vscale=78}
\end{picture}
\caption{Velocity-integrated emission of CO(2-1), outlining the
interaction of the SNR blast wave with the surrounding
molecular gas (Philipp et al., in prep). The southern rim can be
characterised by a slow C-shock. Superimposed the integrated H$_2$
from Burton et al., which nicely traces the interaction layer
between the blast wave and the surrounding molecular gas.}
\label{ic443}
\end{figure}

The atomic fine structure emission in these sources has been
attributed to fast ($v_{sh} = 50-100$ km s$^{-1}$) dissociative J-shocks
propagating in gas of moderate ($n[{\rm H}_2] = 10^2 - 10^3$ cm$^{-3}$)
pre-shock density, while the excited H$_2$ and high-$J$
CO emissions are naturally explained by slower ($10 - 40$ km s$^{-1}$)
nondissociative C-shocks in gas of higher ($10^4 - 10^5$ cm$^{-3}$) preshock
density.  The H$_2$ rotational diagrams (e.g. Cesarsky et al. 1999,
Haas et al. 2003) imply the presence of both
hot (1200 - 2000 K) and warm (400 - 600 K) gas components. The
curvature requires (at least) two C-shocks with different
postshock gas temperatures.
While the atomic fine structure emissions, the warm H$_2$ and
CO appear to be understood relatively well, the observations of the
1720~MHz OH masers and the relative weakness of the H$_2$O rotational
emissions remain puzzling.

\begin{figure}[!ht]
\setlength{\unitlength}{1mm}
\begin{picture}(200,120)
\special{psfile=cesarsky_fig.ps
angle=0 hoffset=-20 voffset=-15 hscale=78 vscale=78}
\end{picture}
\caption{(from Cesarsky et al. 1999) integrated emission of the H$_2$
S(2)-S(7) rotational lines with representative spectra (box-smoothed
to 18$"$x18$"$). The excitation plot reveals a wide range of excitation
conditions, with the bulk of the warm H$_2$ at temperatures at a few
100 K, and 10-20\%
of the gas at elaborated temperatures exceeding 1000 K (in a
two-component decomposition).}
\label{cesarsky}
\end{figure}

The observational data obtained by ISO on H$_2$O, OH and CO in these
sources is severely sensitivity-limited and  beam
diluted. Indeed, emission from warm water vapor has only been
detected reliably (i.e. at S/N ratios $\ge 3$) in a single transition
(the $4_{14}-3_{03}$ line at 113 $\mu$m)
at a single position (clump C in IC443) in one of the four sources to be
observed.  In addition, the $1_{10} - 1_{01}$ 538 $\mu$m transition has
been detected by SWAS towards three positions (C, B, and G) in IC443,
and the $2_{12} - 1_{01}$ 179 $\mu$m transition has been detected by ISO towards
IC443C and 3C391.  However, because these transitions lie so close to
the ground-state of ortho-water (with upper state energies $E_u/k$ of 80
and 27~K respectively), they are effectively excited in gas as cool as
30~K and do not yield a specific probe of the warm gas behind a C-shock.
Thus our present knowledge of the water abundances in the warm and hot
H$_2$ components is based entirely upon upper limits on several
higher-excitation transitions and the single detection of $4_{14} - 3_{03}$
(towards IC443C).

\subsubsection{Typical Characteristics of the sources: IC443 \& W28}

{\bf IC443} is in the Gem OB1 association,
located close to the Galactic Plane (G 189.1+3.0) at an assumed
distance of 1.5 kpc. In the shock-tracing IR images of e.g. Rho et
al., the source appears as an incomplete shell of filaments (ca. 20
arcmin or 9 pc in diameter). The supernova explosion took place
ca. 10$^4$ yrs ago. Numerous studies have been reported in the
literature, from X-rays to visible, infrared and radio
wavelengths. The interaction of the blast wave with the surrounding
gas is documented by the observation of the rotational-vibrational
transitions of H$_2$ (Burton et al.) and by a wide range of molecular
mm-transitions revealing the perturbed kinematics (e.g. van Dishoeck
et al. 1993).

In Fig. \ref{ic443} we present an image of the CO(2-1) emission, superposed on the
integrated H$_2$ from Burton et al., nicely tracing the interaction layer
between the blast wave and the surrounding molecular gas. Along this rim,
the general characteristics in the spectra are wide wings, reflecting the
perturbed kinematics of the shock-accelerated gas. Towards clump G (west rim,
see the figure), Cesarsky et al. (1999) have spectro-imaged with ISOCAM
(6" pixel) the  H$_2$  S(2)-S(7) rotational lines (Fig. \ref{cesarsky}), revealing clumpy
gas with a wide range of temperatures (bulk temperatures of a few 100 K "co-exist"
with much higher temperatures, exceeding 10$^3$ K).

\begin{figure}[ht]
\begin{picture}(200,320)
\special{psfile=w28_channel_a_fgl1.ps
angle=0 hoffset=-15 voffset=-15 hscale=78 vscale=78}
\end{picture}
\caption{Channel map of CO(2-1) on SNR W28 (channel at v=-2.5 km/s, integrated for
2.5 km/s) . Clearly visible the filamentary and clumpy structures of the molecular
gas and the bright emission from regions of strong interaction with the SNR blast wave.
The inserted zoom-in shows a 75$''\times$110$''$ map of CO(3-2) (Frail \& Mitchell
1998) towards the W28F OH-Maser region. The locations of the OH-masers are indicated by circles.
The CO(3-2) spectra were integrated from -10 to 40 km/s.}
\label{w28}
\end{figure}

Prominent {\bf W28 (G6.4-0.1)}, located in a complex region
of the Galactic plane, belongs to the class of SNRs known as "thermal
X-ray composites" with a "mixed morphology" of shell
appearance at radio wavelengths with centrally condensed thermal X-ray
emission. X-ray emission is also reported towards segments of limb-brightened
shell in the northeast
and northwest, where at the encounter of the SNR shock with the molecular
material over forty 1720 MHz OH masers were found, lining up along
the brightest synchrotron features. This interesting correlations between the molecular
gas, synchrotron emission and OH-masers is displayed in Fig. \ref{w28}.

\subsubsection{Proposed Observations}

\begin{table}[ht]
\caption{Predicted antenna temperatures and
integration time requirements for IC443 clump C}
\vspace{3mm}
\label{table_ic443_model}
\centering
\begin{tabular}{lccrc}
\hline
Transition&$\nu$&$\lambda$ & $T_{A}^*$ $\Delta v$ &Integration time$^{**}$\\
&(GHz)&($\mu$m)& (K km/s) &\\
\hline
H$_2$O  $1_{10} - 1_{01}$&556.9&538.3&0.68$^{*}$&A\\
H$_2$O  $2_{11} - 2_{02}$&752.0&398.6&0.44$^{*}$&B\\
H$_2$O  $2_{02} - 1_{11}$&987.9&303.5&0.92$^{*}$&A\\
H$_2$O  $1_{11} - 0_{00}$&1113.3&269.3&0.55$^{*}$&B\\
H$_2$O  $3_{12} - 2_{21}$&1153.1&260.0&1.33&A\\
H$_2$O  $2_{12} - 1_{01}$&1669.9&179.5&1.06$^{*}$&A\\
H$_2$O  $3_{03} - 2_{12}$&1716.8&174.6&0.95&C\\
CO J = 10 -- 9&1152.0&260.2 &8.84&Observed with
H$_2$O  $3_{12} - 2_{21}$\\
CO J = 16 -- 15&1841.3&162.9 &2.74&Observed with
OH $^2\Pi_{1/2} 3/2 - 1/2$\\
OH $^2\Pi_{1/2} 3/2 - 1/2$&1834.7&163.4&0.65$^{***}$& C\\
\hline
\end{tabular}
\flushleft
$^{*}$transition for which significant additional emission is expected from
cooler gas unassociated with the shock

$^{**}$Integration times assumed: \\
\hspace{1cm} A: 10 min in IC443 (3 positions) and W28; 20 min in W44 and 3C391 \\
\hspace{1cm} B: 15 min in IC443 (3 positions) and W28; 60 min in W44 and 3C391\\
\hspace{1cm} C: 60 min in IC443 (3 positions) and W28 \\
\hspace{1cm}
\hangindent=30pt
[{\small Taking the H$_2$ S(3) line
intensity as a probe of the total column density of warm and hot shocked
gas, we compute the factor by which the line flux estimates
tabulated for IC443C should scaled. Since the H$_2$ S(3) line
intensity was not measured by ISO at positions C and B in IC443, we used
the H$_2$ $v = 1-0$ intensity instead.}]

%\hangindent=0pt
$^{***}$ Value predicted if the OH/CO abundance ratio is the same as that
inferred from observations of 3C391

{\small Note: as secondary spectral line targets, we will observe line emission from
the molecules CH$^+$ ($J = 2 - 1$ at 1669 GHz) , NH, NH$_3$ ($1_0$ -- $0_0$ at 572
GHz), and H3O$^+$ ($0_{0,1} - 1_{0,0}$  at 985 GHz) simultaneously with the
primary target lines listed above.}
\end{table}


The proposed observing strategy shall target at representative cuts in
shock-probing transitions across the interaction zones
as outlined by e.g. the IR data of H$_2$ and [FeII]. A more complete mapping
of selected clumps will be
possible only in the strongest the fine-structure lines like [OI] and [CII].

Because the spectroscopic information (line  profile) will be crucial
 to discriminate between the various emission components (along the
line-of-sight and  within our  finite beams, most profiles will  be a
composite of contributions from the cool(er) pre-shocked  gas and the
shocked gas proper), these observations will be unique for HIFI. Observations
with PACS of critical FIR lines (like OI) will complement our data base.

Based on ISO data, the [CII] line is sufficiently strong to perform on-the-fly mapping.
With 2 sec per dump and 2 coverages per core, we need 4 hours of HIFI integration time for
all 6 clumps (assumed size: 3$\times$2 arcmin). Mapping of the [OI, 63$\mu$m] transition with PACS
with minimum integration time (20 sec) requires 1.1 h per clump (same size of coverage as before, beam spaced),
plus an additional 0.4 h for targeted measurements of the [OI, 145$\mu$m], hence 9 hours in total.
Similarly, the CO(10-9) is predicted to be strong, and a fast OTF map (2 sec/dump) will outline the
extend of the warm, shocked CO emission (2 hours); towards the peak targeted observations of the
CO(16-15) line will envelope our excitation studies (20 min each, 2 hours total).
\\
Wrt. to water, the observing strategy will be to map the two lowest
ground-state transitions of both ortho- {\bf and} para-H$_2$O, but
then - as predicted best-suited from the detailed models - we will
focus on either the para- or the ortho-water excitation ladder to
derive shock physics (this will include the isotope transition of
the resp. ground-state transition):  \\
(a) in order to probe the extent of the water emission in IC443, we will
carry out raster maps in the bright 557 and 1113
GHz ground-state transitions along the shocked ridge,
covering an angular extent of ca. 3$\times$1 arcmin.
With 6$\times$2 (557 GHz) and 12$\times$4 (1113 GHz) points and 3 min total integration time each, we
achieve a rms of 30 (50 mK, resp.). The total additional time required is 3 hr (per core).
\\
(b) targeted observations: in Table \ref{table_ic443_model} we present a list of target transitions,
along with CO and H$_2$O line strengths predicted for IC443, position C.
The line strength predictions were obtained from the molecular shock model
proposed by Snell et al. (2003), in which a slow (10 km/s)
C-shock propagating in gas of enhanced ionization yields the 2000 K gas
needed to explain the H$_2$ rovibrational emission but leaves icy grain
mantles intact.  The observed H$_2$O ($4_{14}-3_{03}$)/CO ($J=16-15$) line
ratio $\approx$ 0.6 requires a postshock $n$(H$_2$O)/$n$(CO) abundance
ratio of 0.4, almost an order of magnitude below the value expected
without oxygen depletion.  The H$_2$ $v = 1-0$ map of Richter (1995)
was used to estimate and correct for the degree of beam dilution
in the ISO observations, given the large ISO beam diameter $\approx
75''$ (FWHM).

We propose to observe the same set of target transitions in two other
positions of luminous H$_2$ $v = 1 - 0$ emission within the southern  rim of
IC 433: positions B and G in the notation of (Huang et al. 1986).
Neither of these positions has been observed with ISO, but
we expect the line emission to be comparable to that at position C,
given the relative strength of the H$_2$ vibrational emissions at these
positions.  We suggest to observe one position each in 3C391,
W44, W28.  For these sources, no water
lines other than the $2_{12}-1_{01}$ 179 $\mu$m transition have been detected;
the upper limits on the H$_2$O ($4_{14}-3_{03}$) / CO(J=16-15) line
ratios are comparable to the measured value in IC443C.
\\[2ex]
Based upon the proposed observing times in the table of target lines,
HIFI observations of these sources will require a grand total of {\bf 45 hr}
with HIFI plus additional {\bf 9 hr} with PACS.

\subsubsection{Preparatory Work}
The distribution of the medium warm CO (means: $J=4-3$) has been completely mapped with CHAMP at the CSO,
the higher submm CO transitions are scheduled for CHAMP+/APEX. Similarly, the neutral carbon
fine-structure lines will be measured with APEX. Towards IC443G a dedicated high-resolution interferometric study
is currently performed with the Plateau-de-Bure Array. Supplementary programs have been launched with Spitzer.
SOFIA could potentially provide the FIR-spectra with adequate spectral resolution.

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\end{document} 
