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\begin{center}
{\Large HIFI Core Science Program \\
The interstellar medium}

\vspace*{0.5cm}
{\sl Draft March, 15 2003/JSt}

\end{center}

{\small commented and edited by J. Stutzki}

\newpage

{\bf Introduction}

{\bf JSt: add paragraphs to intro to give\\

- short overview of topics and their relation to other HIFI core sciene
areas\\

- give goals of document (scientific justification of each topic; HIFI
uniqueness of proposed observations; coherent time estimate; identify prep
work necessary)\\

- general comments on HIFI uniqueness: \\
--first time high spectral resolution observations throughout the FIR; \\
----for submm: first time without earth atmosphere, hence unprecedented 
    sensitivity (except for a few KAO observations); \\
----for FIR: first time at all! \\
--SOFIA is complementary: has still residual atmosphere in many places; but may
have better mapping capabilities for some strong lines; HOW DO WE DEAL WITH
 THIS IN THE PROPOSED HIFI CORE SCIENCE?\\
--other HIFI instruments: PACS: crossover resolution (equal sensitivity per
  spectral and spatial point) at about 30-100 km/s, and is going to have only
  resolution of about 100 km/s anyway. But spatial multiplexing advantage!
  SPIRE only very low spectral resolution and thus impossible to perform
  observations presented here (but both PACS and SPIRE necessary for
  complementary observations: broad band dust spectra, important lines beyond
  the spectral range of HIFI\\

- give rules for time estimate used throughout proposal (following Volker's suggestions).\\}

{\bf then go to the scientific overview}\\

The study of the interstellar medium (ISM) is critical for an understanding of
both star formation and the evolution of galaxies.  HIFI is a unique
instrument to accomplish this, since for local sources most important spectral
structures are obscured by the Earth atmosphere.

The formation of new stars in galaxies releases such large amounts of
radiative and kinetic energy in the interstellar medium that matter
continuously cycles between the dense and diffuse phases of the medium. The
chemical evolution of matter in space and the formation of stars must be
investigated within this broad context.  Gas phase abundances of elements such
as Si and Fe show that a significant fraction of matter accreted onto dust in
molecular clouds is released back to the gas within the low density phases of
the ISM. At the opposite, the detection of numerous molecular clouds at high
Galactic latitude and of large abundances of molecules in the diffuse ISM show
that the condensation of gas from the warm to cold phases is easily followed
by the formation of molecules.  Molecular abundances in these low column
density molecular clouds are not readily explained with state of the art gas
phase chemistry and alternative formation processes may be effective. For
instance, The decrease of turbulent line widths from the warm phases to atomic
clouds and from atomic to molecular clouds suggests that the condensation
process of clouds goes with a dissipation of turbulent kinetic energy. The
dissipation process may be able to greatly enhance the abundance of some
species (e.g. CH$^+$).


The proposal is organised as follows :

Section 1 deals with the ISM structure, the role of HIFI for studying physical
and chemical processes in the Warm and dense ISM is described in Section 2,
Section 3 presents observations probing the transition from atomic to
molecular gas, especially the formation of hydrides, while complex molecules
are the subject of Section 4.



\section{ISM structure}
{\bf JST: I think this section title is not very well choosen; what
  ``structure'' do we have in minde? Spatial/velocity structure? Phases of the
  ISM? we have to find some better section title!}

\subsection
{HIFI studies of the ISM phases}

The interstellar medium (ISM) exists in different phases. Traditionally one
distinguishes the warm ionized medium (WIM), the warm neutral medium (WNM) and
the cold neutral medium (CNM). All three phases play an important role in the
evolutionary cycle of interstellar matter, which overall controls the
evolution of galaxies. A full understanding of the physical and chemical
processes controling the ISM phases and their composition is thus of
fundamental importance to astrophysics.

Due to their specific excitation conditions the various phases have different
strandard tracers. The WIM is best studied in the optical and UV, in
particular through its diffuse H$\alpha$ emission, the WNM and CNM in the HI
21cm line and, for its densest component in the rotational lines of CO and
other heavy molecules. The WNM is best traces by the strong cooling lines of
C$^+$, CI, OI (for the denser parts) and other tracers such as N$^+$.  Due to
the complex spatial and kinematic structure of the turbulent ISM, high spatial
resolution and the velocity information present in the line profiles have
proven to be crucial for understanding of the ISM, allowing to separate
otherwise overlapping components along the line of sight.
 
Over the years dedicated efforts have been conducted and significant progress
has been made in the understanding of the ISM, in particular with the
HST/GHRS, COBE/FIRAS and ISO/LWS and ISOCAM instruments, and more recently
through UV absorption studies with FUSE. {\bf JST: include the Japanes CII
  balloon BICE}

{\bf JST: I suggest to add here a more detailed overview of what the
  limitations are of the present data, and where HIFI is going to improve:
e.g.:\\
FIR line observations (mostly CII, NII only from COBE and one single KAO
  observation in an external galaxy) limited to\\
(i) small telescopes (KAO, BICE, ISO) and hence low angular resolution,
  whereas comparabel HI and CO observations tell us, that there is structure
  down to the smallest scales, and that structure is relevant for
  understanding the physics and exciation\\
(ii) except for one set of pioneering [CII] KAO observations, all with
  spectral resolution of 30 km/s or poorer, i.e. only integrated line
  intensities w/r to the issues discussed below.\\ }

 However, a global description of the
structure of the ISM is missing and the overall structure of the ISM remains
elusive, in particular the volume filling factor of the various phases, the
amount of small scale structure present in each of them, the density contrasts
and perhaps most importantly the thermal pressure of each of these phases, and
its variation with the distance $z$ to the galactic plane, and at various
distances from the Galactic center.  The link between the various phases
(e.g. are they in thermal pressure balance?) is likely to be related to their
formation mechanism and has to be established. Is the hierarchy of structures
observed in the coldest component (that of molecular clouds) a fossil of
thermal/gravitational instabilities developed initially in a placental warm
medium at 10$^4$ K?  The structure of the diffuse ISM may thetrefore be
related to that of the coldest self-gravitating components, and thus
ultimately to the process of star formation.

The structure and physics of the diffuse ISM has also important bearings to 
(i) understanding the gaseous structure of one of the many screens between
the Cosmic Background (CBG) and the Sun, and therefore that of the dust
distribution which affects the continuum measurements of the CBG, 
(iii) the large observed porosity of the neutral medium to UV
photons and (iv) understanding the first steps of interstellar chemistry. 

The capabilities of HIFI will allow new important steps in the study of the
ISM thanks to its high angular resolution, its sensitivty, but most important,
its high spectral resolution: HIFI will for the first time allow to obtain
fully velocity resolved spectra across the FIR spectral range, It will thus
provide for the first time for the many important trace lines of the ISM in
its different phases the velocity information that is fully comparable to the
rich information in the extensive HI- and CO-surveys of the denses phase of
the ISM.

Within the core science program we focus on the the most important
contributions expected trough HIFI:

\subsubsection{\bf 
Thermal pressure determination from [CI] fine structure lines}

Thermal pressure determinations as derived from [CI] fine structure
lines and molecular line observations vary from tracer to tracer.
The thermal pressure derived from [CI] absorption lines in the general
{\bf JST: what is ``general'' here? which of the phases discussed above? I
  guess,it is the diffuse component.}
ISM is P$_{th}/k\sim 3 \times 10^3$ K cm$^{-3}$ (Jenkins \& Shaya 1979,
Jenkins \& Tripp 2001) 
with large excursions above this value, while recent observations of
molecules in absorption against extragalactic sources provide 
P$_{th}/k\sim 2\times 10^4$ K cm$^{-3}$ (Lucas \& Liszt 1994). 
If large pressure differences exist, shocks propagating into the 
coldest phases should be present. Small scale density structures 
are also prominent in the spectra of nearby lines of sights
(Pan et al. 2001, Andrews et al. 2001).

{\bf JSt: what we completely lack here, is an argument that says why HIFI CI
  observations are better than other possibilities in deriving the pressure,
  and thus, why it is worthwhile to spend precious HIFI time on integration on
  pretty damned weak CI lines!}
HIFI will extend the present observations of the [CI] fine structure lines to
low column densities HI clouds. In such clouds most of the gas carbon
is ionised and C is only a  minor constituent. However carbon 
abundances derived from UV absorption
studies (Jenkins \& Tripp 2001) show that the [CI] fine structure emission
lines can be measured with HIFI. Thanks to the high spectral resolution of 
HIFI we can use the information on the level population {\bf JST: what does
  the high spectral resolution have to do with being able to use the info on
  the level pop? rearrange!} , and from the
line profiles to determine the gas pressure {\bf: how dow we get the pressure
  from the line profiles?} for each of the
velocity components of the clouds which is a fundamental information
to understand the origin of clouds structure. 

{\bf JST: 
I suggest to restructure and move the details of the proposed observations
  into the individual sections, i.e. right after the following paragraph and
  have the present subsection ``proposed observations'' as a summary only;
  this avoids some hopping back and forth}\\

{\bf Proposed observations:}\\
An illustration of a possible target field is shown in Figure 
%\ref{hi} 
from
Miville-Desch\^enes et al. (2002). 
The HI emission there reveals conspicuous small scale filaments, 
the width of which has not been quite resolved by the instrument (DRAO
interferometer, resolution 1'). Such a small scale structure in the HI
emission seems to be found in many directions. 

{\bf JST: we lack a quantitative estimate of the line intensities for such
  clouds (note: we might use the high latitude cloud observations from AST/RO
  and the recent KOSMA observations from Heithausen's group as a guideline). } 

\subsubsection {\bf [CII]: separate the contribution from WIM and CNM}

{\bf JST: we lack a section on why it is important to separate the WIN and CNM
  contribution to the [CII] emission; who from the broader astrophysics
  community is interested in this, and why? Or is it just the [CII]
  specialist, who would like to know?}
High spectral resolution 
observations of the [CII] fine structure line will allow 
the separation of the contribution  of the
cloud phase from that of the warm ionized medium (WIM) and, together with  
H$\alpha$ and HI emission spectra, 
the thermal and non-thermal contributions to the observed line widths. 
The narrowest velocity components in HI clouds have velocity widths of 
3 kms$^{-1}$ (FWHM) which are close to the expected thermal width for 
gas at a few 
100 K. For the corresponding [CII] emission components we expect to see a
narrowing of the lines due to the higher mass of carbon. 
{\bf JST: what does it imply, if we in fact find this narrowing? What, if not? 
How about similar attemps in the optical/UV absorption work? What are their
  results?} 
Maps in the [CII] line will also help in understanding the small scale
structure of the WIM. 
{\bf JST again: why is this important? What does it tell us about the ISM,
  that we need to know? The mere fact that maps help to understand structure
  is somewhat week, I guess!}
Warm
components are presumably more smoothly distributed, but CaII absorption
lines against binary stars (Meyer and Blades 1996; Watson and Meyer
1996) suggest that small scale
structure also exists in the WIM. 

{\bf proposed observations:}
{\bf JST: note: here we actually give estimated line intensities, which we
  presently do not in the [CI] section; this has to be made coherent}

[CII] line emission for an HI cloud with a column density of
N(H) $\sim 5 \times 10^{20}$cm$^{-2}$ is 
estimated to be I(CII) $\sim 10^{-6}$ erg cm$^{-2}$ s$^{-1}$
sr$^{-1}$ (Matsuhara et al. 1997), 
{\bf JST: compare with Timmerman et al. 1998 results on L1457}
while
the WIM contribution is I$_{WIM}$(CII) = $4 \times 10^{-7}$ erg
cm$^{-2}$ s$^{-1}$sr$^{-1}$
assuming I(CII)/N(H) = $3 \times 10^{-26}$ erg s$^{-1}$ H$^{-1}$
(Gry et al. 1992 )  and all the carbon in the gas phase.
{\bf JST: drop, repeats above:
The velocity information 
will be a critical input to disentangle the relative contributions
of the CNM and WIM to the line intensity.}

\subsubsection{\bf [NII]: exclusive tracer for the WIM}
The [NII] line at 205 $\mu$m is expected to be 0.06 times weaker than
the [CII] line in the WIM and will therefore be 
%JST correctedL detected 
detectable %JST
by
HIFI only up to a few  degrees from the plane. It is an excellent
tracer of ionized gas, with no contamination from any other
phase of the ISM.
The comparison of [NII], [CII] and HI emission spectra will  permit
the investigation of  the degree of association between the so-called cold 
(or cloud)(T$<$ a few 100 K)
and warm intercloud  (T$\sim$ 8000 K) neutral phases of the interstellar
medium and the role of thermal instability in the formation of small scale
density structure within the diffuse ISM (Hennebelle \& P\'erault 2000). 
{\bf JST: this is much too vague, at least for my taste! How does the
  comparison of spectra allow to do all to investigate all this? Which feature
  of the spectra are the crucial ones?}
The
warm neutral gas does not contribute to the [CII] emission while both the 
cold and warm neutral phases contribute to the HI emission. The [NII]
observations allow subtraction of the fraction of [CII] from ionized gas. 
{\bf JST: reading the last sentence, I get completely confused! We introduce
  new termini: cold neutral (or cloud) phase, warm intercloud neutral medium;
  and then list various lines associated with various components, but not
  given a clear clue of what finally comes from what; plus: how is the
  subtraction of the fraction of [CII] from ionized gas going to work? Simply
  by identifying it through the line profile? Or by scaling it with some
  reasonable assumptions from the [NII]??}

With the gas temperature and pressure estimates deduced from
the neutral carbon lines, it might be possible to
derive the carbon abundance from the [CII] line intensity. 
{\bf JST: I regard this and the following as highly speculative, i.e. I agree
  with the ``might''. I guess, what is meant here, is that the temperature
  and/or pressure from [CI] can help to estimate the level population in
  [CII]? But his assumes that they come from the same gas. What if not? How do
  we tell?}
We will be able to 
correlate variations in the carbon depletion with local physical conditions
and dust extinction and emission properties including observations of
the dust emission spectrum with PACS and SPIRE. Such measurements
will open a completely new perspective on the cycling of carbon between
dust grains, small aromatic hydrocarbon particules, CO and atomic carbon.
Indeed, Miville-Desch\^enes et al. (2002) show evidences
for correlation between the abundance of the very small dust particles 
responsible for the mid-IR emission, and the gas velocity structure.

{\bf JST: shouldn't this go to section 1.3?}
Observations of the emission lines of the ISM as described above should 
be also performed in fields surrounding the 
lines of sight studied in absorption in the UV by the HST and FUSE. 
The absorption studies give access
to abundances and Doppler broadening for many atomic and ionic species but
they are limited to a spectral resolution of about 3 km s$^{-1}$ (HST)
and the spatial information is almost non existent.
However they provide  unique information  
on the abundances and ionization state of numerous IS 'components' 
i.e. defined by a common absorption velocity. Some of the components are 
neutral and clearly associated with the WNM, and some are at least 
partially ionized, the origin of the ionization being not always clear. 
Some of these components in the lines of sight to halo stars should 
correspond to the extended H$\alpha$ emitting layer in the Galaxy but its
identification is not straightforward, especially because of the lack of 
spatial information of UV observations and the lack of spectral resolution 
of H$\alpha$ observations.
{\bf JST: what exactly are the HIFI observations going to contribute to this?
  How do we handle the difference between the pencil and 15-30 arcsec beam?
  What observations in detail do we propose? How much time do we need?}

%\begin{figure}
%\resizebox{14cm}{!}{%
%  \includegraphics*[0.5cm,0.5cm][20.5cm,28.5cm]{ursa_hi.ps}
%}
%\label{hi}
%\caption{Channel maps for the HI 21cm emission 
%in a high latitude cirrus cloud (Miville-Desch\^enes et al. 2002). 
%The spatial resolution is 1 arcmin; these data have been obtained
%with the DRAO interferometer.}
%\end{figure}

\subsection{proposed observations}


For each source {\bf JST: which sources are we talking about? Wrap up from
  previous discussion!} we propose to map small fields using the OTF mode
in order to have a description of the small scale structure, and
to have  global information to compare with previous large scale observations
(IRAS, ISO {\bf JST add: BICE}).
For [CII] and [NII], we use a system temperature of 1800 K and adapt
the frequency resolution to the expected line width.
{\bf JST: I am not so sure that we want to use a frequ resolution adapted to
  the line width; the purpose is, to really see the detailed line profile
  structure, so that we need some other constraint on the resolution aimed at:
  thermal width, turbulent width of individual components?}
The expected [CII] line strength is 10$^{-6}$ erg cm$^{-2}$s$^{-1}$sr$^{-1}$
for high latitude clouds, and $4 \times 10^{-4}$ erg cm$^{-2}$s$^{-1}$sr$^{-1}$
for a line of sight across the Galactic Plane at $b=0.5^\circ$. 
This corresponds to line integrated 
intensities of 0.14 K kms$^{-1}$ ($\delta v=2$ km s$^{-1}$) and 
60 K kms$^{-1}$ ($\delta v=30$ km s$^{-1}$) respectively.
The ionised nitrogen line at 205 $\mu$m is about 10 times weaker than
[CII], but has broader features.

Adopting a velocity resolution of 1 $km s^{-1}$ for [CII] and 4
$km s^{-1}$ for [NII], a single pixel is detected at the
3$\sigma$ level in 900s (high latitude), 30s (Galaxy, [CII] and [NII]).
{\bf JST: is the resolution of 1 km/s really sufficient for the [CII]
  observations? We argue above, that CII should be narrower than HI by
  sqrt(mass) and quote 3 km/s as the thermal width of HI. Thus, 1km/s
  resolution is only a single resolution element for the whole, thermal [CII]
  line, certainly not sufficient, and definitely not at a S/N of 3!}

We need to observe both [CI] lines as well, covering the same
area in the 492 GHz line, and observing a few positions in the 809 GHz
line. The estimated line strength is 0.1 $K km s^{-1}$ for the 492 GHz
line in a high latitude cloud and 6 $K km s^{-1}$ across the Galaxy.
The 809 GHz line is very sensitive to the gas pressure.
In diffuse clouds, it is weaker than the 492 GHz line, by a factor of
10 approximatively. 

We adopt use a system temperature of 233 K for 492 GHz and 454 K for 809
GHz, and select a frequency resolution of 1 MHz (or 0.6 km s$^{-1}$) 
at 492 GHz.
For the high latitude fields, the requested integration time per pixel 
is 120 s for the 492 GHz [CI] line. The 809 GHz lines requires 
4 hours integration per position. 

For the Galaxy fields, the signal is expected to be stronger 
over a larger velocity coverage than for the HLCs and we plan to decrease
the integration time by a factor of two in the plane. Since the 
lines, if they follow the latitude dependence of the emission measure,  
decrease by a factor 3 between $b=0$ and $b=2^\circ$, the integration 
time for fields at larger latitude will be similar than for HLCs.

To cover a 4'$\times$4' field in the [CII] line, 
a map of 40 $\times$ 40 positions with a
sampling of 6" is required. For
the high latitude clouds, the  spatial resolution will be degraded to
30" to be comparable to
the [CI] data. We give the integration time for
a total of 64 points. For the lines of sight through the Galaxy, we
compute the integration time for a spatial resolution of 20 arcsec
(144 independant points)

We list below the requested integration times to observe one field 
of size $M \times N$ in arcminutes, with a resolution $\delta v$ and a
final rms $\sigma$ once the angular resolution has been smoothed to
30" for the HLCs and 20" for the galaxy at $b=0$. Frequency switching
is used for the HLC fields.

\begin{itemize}
\item{High latitude [CII] : 16 hours $(M\times N /4' \times
4') ( 14 {\rm mK}/ \sigma) ^2  ( 3 {\rm km s^{-1}}/ \delta v) ^2 $}
\item{High latitude [CI] : 5 hours $(M\times N /4' \times
4') ( 12 {\rm mK}/ \sigma) ^2  ( 3 {\rm km s^{-1}}/ \delta v) ^2 $}
\item{Galaxy [CII] : 1.2 hour $(M\times N /4' \times
4') ( 60 {\rm mK}/ \sigma) ^2  ( 4 {\rm km s^{-1}}/ \delta v) ^2 $}
\item{Galaxy [NII]: 1.2 hour $(M\times N /4' \times
4') ( 60 {\rm mK}/ \sigma) ^2  ( 4 {\rm km s^{-1}}/ \delta v) ^2 $}
\item{Galaxy [CI] : 2.5 hours $(M\times N /4' \times
4') ( 17 {\rm mK}/ \sigma) ^2  ( 3 {\rm km s^{-1}}/ \delta v) ^2 $}
\end{itemize}

Adding at least one 809 GHz [CI] point, 
the total time per field is thus 25 hours (high latitude) and 9
hours (Galaxy, $b=0$). 
The total satellite time will be at least a factor of two
larger due to various overheads.


We have selected the two fields in the Ursa Major cirrus
which have been observed with ISOCAM (Miville-Desch\^enes et al. 2002).
The first
one, where no CO is detected, bears the signatures in the mid-IR 
of large abundances of small dust grains, 
the second is located at a
transition region between atomic and molecular gas.
The spectral dust properties and gas kinematics in these two fields
are different.
  
In the Galaxy, we propose to observe 4 fields at increasing Galactic
latitude to measure the scale heights of the different phases.
The position has been chosen at the tangent point, to minimize
the spread in velocity. The physical conditions
have been studied using ISO data (Verstraete et al. 1999).
We shall observe in the OTF mode for all lines of sight. We can use
the frequency switching mode for the high latitude cloud observations
where we expect narrow lines. We will investigate further the best
observing mode for the fields through the Galaxy, depending on the
instrument performances as revaled by the coming instrument tests. 
The overheads 
will likely be larger than for the high latitude cloud observation,
and we applied a factor of 2.5.   

It sums to a total of {\bf 50 hours} for the two fields in the 
Ursa cirrus, and  {\bf 36
hours} for the 4 lines of sight through the Galaxy which corresponds
to a total of about {\bf 200 hours} including the overheads.
{\bf JST: what about the FUSE lines-of-sight?}

{\bf JST: we have to make sure, that prep observations give us 
HI, CO etc. data at similar specral and spatial resolution for those fields}



\begin{table}[!h]
\caption{Position for Maps}
\begin{tabular}{lccl}
 Name & RA(2000) & DEC (2000) & Comment \\
\hline
Ursa Major &   9:45:00 & 70:25:00 & Cirrus, 2 $4' \times 4'$ fields \\
Galaxy & 18:39:16.3 & -05:42:58 & 4 $4' \times 4'$ fields with
different $b$ starting at $(l,b) =  (26.86 , 0.09)$\\
\hline
\end{tabular}
\end{table}
\vspace*{0.5cm}

%\subsection
{PDRs in diffuse gas}\\
{\bf JST: I suggest to change this subsection title. This is not at all
  specifically on PDRs, but on the chemistry of diffuse gas; hence:}
\subsection{Chemistry in diffuse gas}

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 2001) have measured the abundance of other molecules,
not easily studied in the UV/visible wavelength range. There is little
overlap so far. By observing a few lines HIFI can bring
critical information on these two types of lines of sight (bright
stars and continuum sources). As described above, the combination of
the fine structure lines of C and C$^+$ with the H and H$_2$ 
column densities (available from FUV data) gives access to the gas
pressure and carbon abundance. Since the sources are extended,
observations will be performed in frequency switching.
We shall use PACS for mapping one field of view in C$^+$
to assess the structure of the ISM in the vicinity of these particular
lines of sight. 
{\bf JST: (I) time estimate for the PACS observations! (ii) is
  the PACS data really sufficient, i.e. can we get sufficient structure
  information form integraetd profiles? Why don't we include such a field as
  one of the ones to measured in the preceeding section?}

In addition to the special direction towards $\zeta$Oph, it
is interesting to study a few more points where the CO emission
gets stronger to sample the transition from atomic to molecular gas.
Information on the line profile is important for
performing a meaningful comparison with models since 
the behaviour different velocity components, detected in e.g. sodium lines, is expected  
to be different.

The selected lines of sight are listed below.

We expect a line intensity of $\sim$ 0.1 K for the 1-0 [CI] line
and $\sim$ 1 K for [CII], similar as for the diffuse HI clouds.
{\bf JST: not true: we quote above 0.14 K km/s for a 2 km/s wide 
line from a high latitude cloud! Needs to be consistent. IN any case: based on
what assumption are these intensities estimated?}
The program therefore requires $\sim 15$ min integration time for
[CI] and $\sim 10$ minutes for [CII] per line  of sight.
{\bf JST: for what velocity resolution is this estimated?}
Stronger [CI] lines are expected towards the additional positions in
$\zeta$ oph, so we plan to spend 5 minutes only for those
positions for both [CII] and [CI].

This sums to a total of $10 \times 0.25) + (5 \times 0.1) + (10
\times 0.2) + (5 \times 0.1) = $ 5.5 hours of integration, hence a
total of {\bf 11 hours} of telescope time including overheads.


\begin{table}[!h]
\caption{Single lines of sight from FUSE and/or IRAM}
\begin{tabular}{lrrl}
Name & RA & DEC & Comment\\
\hline
$\zeta$ Oph & 16 37 09.5378 & -10 34 01.524 & Star + 5 positions ?\\
\hline
B0212+735 & 02 17 30.8134 & +73 49 32.622 & BL lac\\
B055+508 & 03 59 29.7473 & +50 57 50.162 & QSO \\
B2013+370 & 20 15 28.7 & +37 10 59 & Radio source\\
B2037+511 & 20 38 37.12 & +51 19 12.7 & QSO \\
\hline
PG0804+761 & 08 10 58.55 & +76 02 41.8 & FUSE QSO, radio quiet\\
HD73882 & 08 39 09.5344 & -40 25 09.262 & FUSE HD \\
HD34078 & 05 16 18.1497 & +34 18 44.341 & FUSE warm H$_2$ \\ 
HD185418 & 19 38 27.4827 & +17 15 26.052 & FUSE, double system \\
HD102065 & 11 43 36.9487 & -80 29 00.521 & \\
\hline
\end{tabular}
\end{table}

\newpage


XXXXXXXXXXXXXXX (JST: presently edited until here) XXXXXXXXXXXXXXXXXX

\section{Physical and Chemical processes in the Warm and dense ISM}

\subsection{PDRs as probes of the Photoelectric heating in the ISM}
The interstellar gas is efficiently heated by the photoelectric 
effect on dust grains
(visible-UV photons from stars eject energetic electrons from dust
grains which heat the gas while  thermalizing). As a consequence, the
gas temperature - which controls the chemical 
and thermal evolution of the interstellar medium - depends on the 
photoelectric heating rate. Recent theoretical 
(Bakes \& Tielens 1994, Wolfire et al. 1995, Weingartner \& Draine 2001) 
and observational (Habart et al. 2001) studies have shown 
that the photoelectric heating 
of the gas is dominated by the contribution 
of {\em small carbonaceous dust grains} (nanoparticles of radius 
less than 150 \AA, conversely the so-called 
{\em big grains} have larger sizes). In addition, the photoelectric 
heating rate depends on the local physical 
conditions (radiation field intensity, gas density:
Weingartner \& Draine 2001).


At thermal balance (heating and cooling rates are equal), the gas 
heat budget can be determined 
from observations of the major cooling lines. Moreover, in regions of 
low excitation, the 
gas heating is mostly due to the photoelectric effect on dust grains. 
This implies that abundance variations of the small carbonaceous
grains should be reflected in the 
gas heat budget as well as variations of the local physical conditions 
(see Habart et al. 2001). 

We propose to study the heat budget of interstellar gas across 
a number of low-excitation photodissociated 
interfaces in relationship to the carbon abundance. The goal is to 
follow the carbon cycle 
between the gas phase and the small carbonaceous dust particles and
its impact on the gas heating. 


%TABLE 1
\begin{table}[!h]
\caption{Target list and physical conditions}
\begin{center}
%\leavevmode
\begin{tabular}[h]{llccc}
\hline
%\multicolumn{5}{c}{} \\
\multicolumn{1}{l}{Target} &
\multicolumn{1}{l}{Position}            &
\multicolumn{1}{c}{$\chi$ $^a$}         & 
\multicolumn{1}{c}{$n_H$ $^b$}  & 
\multicolumn{1}{c}{$T_{\star}$$^c$}     \\
 & (J2000) & & (cm$^{-3}$) & (10$^4$ K) \\
%\multicolumn{5}{c}{} \\
\hline
%\multicolumn{5}{c}{} \\
L1721$^1$               & 16h14m30s,  & 5     & 3(3)$^\dag$ & 2.2 \\
 & -19$^{\circ}$5' & & & \\
Octopus$^{2}$           & 16h52m30s,  & 5     & 4(3)$^\P$   & 2.5  \\
 & -25$^{\circ}$45'& & & \\
Ced 201$^{3}$           & 22h13m25s,  & 10    & 2(3)$^\P$   & 1.0  \\
 & 70$^{\circ}$15' & & & \\
$\rho$ Oph. Fil.$^4$    & 16h25m55s,  & 400   & 4(4)$^\dag$ & 2.2  \\
 &-24$^{\circ}$21' & & & \\
L1630$^{5}$             & 5h41m45s,   &  1000 & 1(4)$^\P$   & 2.3  \\
 & -2$^{\circ}$15' & & & \\
Horsehead$^5$           & 5h40m55s,   & 100   & 2(4)$^\dag$ & 3.3  \\
 & -2$^{\circ}$28' & & & \\   
California Neb.$^6$     & 4h03m17s,   & 6     & 1(3)$^\P$   & 3.7  \\
 & 36$^{\circ}$28' & & & \\  
Ursa Major$^7$          & 9h45m45s,    & 1     &  100$^\P$   & 2.2  \\
 & 70$^{\circ}$22' & & & \\  
%\multicolumn{5}{c}{} \\
\hline
\end{tabular}
\end{center}
{\scriptsize 
$^a$ Intensity of the UV radiation field in units of the Habing flux at 1000 \AA, $\lambda u_{\lambda}$ = 1.2 10$^{-3}$ erg/s/cm$^2$.

$^b$ The hydrogen density of the region as derived from modelling (\dag, Habart et al. 2001, 2002) and roughly estimated ($\P$).
Notation: $1(4) = 10^4$.

$^c$ Effective temperature of the illuminating star.

References: $^1$ Habart et al. 2001, $^2$ Bernard et al. 1998, $^3$ Cesarsky et al. 2000, 
$^4$ Habart et al 2002$a$, $^5$ Habart et al 2002$b$,
$^6$ Boulanger et al. 1990, $^7$ Miville-Desch\^{e}nes et al. 2002.}

\end{table}


The HIFI instrument -- because of its high spectral resolution and 
sensitivity -- will detect the main gas cooling 
lines from  C$^0$, C$^+$ and CO without contamination 
from the background emission. 
Physical conditions of the gas (pressure, density and radiation field 
intensity) will be obtained through observations and theoretical 
modelling of the carbon cooling lines.
The hydrogen column density at each position will be constrained from 
ancillary observations of the HI-21cm line 
(Penticton, VLA) and also from the emission of the big grains.

To carry out this project, we will observe the major gas cooling lines
 (HIFI: [C$^+$]158 $\mu$m, 
[C$^0$]370,609 $\mu$m and CO(J=8-7 and up), 
PACS: [O$^0$]63$\mu$m) and compare it to 
the emission of small grains as traced by ISOCAM. 
The emission of bigger grains will be measured by PACS and SPIRE.
We will perform these observations towards regions where the emission
of small and big grains are contrasted
indicating dust abundance variations. Our target list is given below 
as well as the integration times. 
In each case, the 5$\times$5 pixels field of PACS will be centered 
on the photodissociated interface (one side of the field 
parallel to the interface) and five HIFI spectra will be taken along a 
cut in the central part of the PACS field and 
perpendicular to the interface.
For the case of the diffuse ISM (Ursa Major), we took the C$^+$ line flux of 
Gry et al. (1992). The other line fluxes were estimated from the model 
results of Wolfire et al. (1995).
The Ced 201 and Octopus lines of sight will sample regions with a
grain size distribution significantly different from that 
of the general ISM (as described, e.g., in D\'esert et al. 1990). 


\begin{table}[!h]
\caption{Integration time estimates for the heating budget section}
%\begin{center}
%\leavevmode
\begin{tabular}[h]{lccccccc}
\hline
\multicolumn{8}{c}{} \\
\multicolumn{1}{l}{Target} &
\multicolumn{1}{c}{C$^+$ $^a$}  & 
\multicolumn{1}{c}{C$^0$ $^a$}  & 
\multicolumn{1}{c}{C$^0$ $^a$}  & 
\multicolumn{1}{c}{O$^0$ $^a$}  & 
\multicolumn{1}{c}{PACS $^b$}           & 
\multicolumn{1}{c}{SPIRE $^c$}          & 
\multicolumn{1}{l}{Time $^d$} \\
 & (158$\mu$m) & (370$\mu$m) & (609$\mu$m) & (63$\mu$m) & (2 bands) & (3 bands) &  (mn) \\
\multicolumn{8}{c}{} \\
\hline
\hline
\multicolumn{8}{c}{} \\
L1721                   & 6(-3)$^e$ & 7(-3) & 6(-3) & 60$^e$ & 0.3 & 1.7 & 40 \\
%\multicolumn{8}{c}{} \\
Octopus$^{f}$           & - & - & - & - & - & - & 40 \\
%\multicolumn{8}{c}{} \\
Ced 201$^{f}$           & - & - & - & - & - & - & 40 \\
%\multicolumn{8}{c}{} \\
$\rho$ Oph. Fil.        & 0$^e$ & 7(-3) & 6(-3) & 1.5$^e$ & 0.3 & 1.7 & 3.5 \\
%\multicolumn{8}{c}{} \\
L1630$^{f}$             & - & - & - & - & - & - & 3.5 \\
%\multicolumn{8}{c}{} \\
Horsehead               & 0.2 & 0.1 & 5(-2) & 25 & 5 & 25 & 55 \\
%\multicolumn{8}{c}{} \\
California Neb.         & 0.1 & 1.8 & 0.9 & $*$ & 3 & 14 & 20 \\
%\multicolumn{8}{c}{} \\
Ursa Major              & 12 & $*$& 60 & $*$ & 12 & $*$ & 84 \\
%\multicolumn{8}{c}{} \\
\hline
\end{tabular}
%\end{center}
{\scriptsize 

All integration times are given in minutes and for a 5$\sigma$-detection. 

$^a$ The integration time in minutes. All values are given for a 5$\sigma$ detection within a resolution element 
of 0.3 km/s. The integration time for the [O$^0$]63 $\mu$m 
line with PACS is also given (175 km/s resolution element). 

$^b$ The integration time at 100$\mu$m for the PACS photometric mode. 

$^c$  The integration time at 500$\mu$m for the SPIRE photometric mode. 

$^d$ Total integration time for one line of sight. Five positions will be observed in each region except for Ursa Major 
where only 3 positions will be observed.

$^e$ The line fluxes have been measured by the gratting mode of ISO-LWS.

$^f$ Ced 201 and Octopus as well as L1630 and the $\rho$ Oph Filament have been assumed to have similar fluxes.

The asterisk ($^*$) indicates  observations which can not be performed
because they require a too large integration time. }

\end{table}


Towards most of the lines of sight, the cooling by H$_2$ may be 
important: it will estimated from 
near-IR observations of rovibrational lines and theoretical
modelling.  Observations of H$_2$ and low-$J$ CO emission are 
already available for the $\rho$ Oph filament and for the Horsehead nebula. \\

We require a total Herschel time of {\bf 20 hours} for this section amongst
which {\bf 6 hours} of the HIFI instrument.



  
\subsection{chemical tracers of PDRs and XDRs}

{\sl In addition to CH$^+$ is there any other molecule to search for
?}

\subsection{Emission lines, Photodissociation regions and shocks}
In high density warm regions, such as bright photo-dissociation
regions, the hydride lines may appear in emission. For example, CH$^+$
rotational lines have been detected in emission in the young
planetary nebula NGC7027 (Cernicharo et al. 1997). 

Most  models of photo-dissociation regions (PDRs) are static,
and use a pure gas phase chemistry, but other physical effects could play
a role in these regions. Gerin et al. (1998)
 and Lemaire et al. (1999) have
measured the line profile of the [CII] 158 $\mu$m and
H$_2$ v=1-0 S(1) 2.12 $\mu$m lines respectively, in NGC~7023. There is
a clear velocity gradient through the PDR which is not 
taken into account by models. Photodissociation and evaporation of 
molecular gas from the molecular cloud could explain this
gradient. Since this example is likely not to be unique (see
e.g. Habart et al., 2001),  the static description
of PDRs should not be valid anymore, and should be replaced by
dynamical models including advection of cold molecular into the
PDR and evaporation of gas. 
PDR are good regions to probe how  kinetic energy can be fed
into dense interstellar clouds : from observations of the carbon
recombination lines, combined with [CII] and [CI] data, 
a low level of turbulence is found in the
 C$^+$ region, which contrasts with the
larger velocity dispersion in the molecular gas (Wyrowski et al. 2000).

The impact  of such dynamical phenomenon on the structure and chemistry
of interstellar gas is not well understood as most
models focus on either the chemistry or the dynamics separately.
With the possibility of measuring detailed line profiles
for all  species dominating the gas cooling, a quantitative
assessment of the role of gas dynamics will become possible.
Conclusions based on bright photo-dissociation regions will
be valid for the diffuse interstellar gas as a whole. For instance, 
the dissipation of turbulence is an additional heating source,
highly localised in specific regions of a turbulent
flow (Pety \& Falgarone 2000). Measurements of the main cooling lines, either
from the ground (see below)  or with Herschel
will help to quantify the impact of this  additional heating
source on the properties of diffuse interstellar gas.
In addition, advection of cold gas into the PDR is likely
to affect the chemistry. The ice mantles are expected
to be destroyed by the combined influence of the enganced
temperature and UV radiation. Since water is the main constituent of
grain mantles,  HIFI is ideally suited to reveal the presence and
measure the magnitude of advection processes in PDRs. 

The target list includes potential
sources covering a range of physical conditions (density, radiation
field, shocks) to search for emission lines as well.
Lines as weak as 1 K can be mapped with OTF maps using  circa 60sec
integration time per position. Since the velocity dispersion is small,
all lines can be observed in FSW mode, with no loss of efficiency due
to telescope slewing. The extension of all sources is a few minutes.
For a $2' \times 2'$ map at 15 arcsec sampling, the total number of
points is 64, so the total integration time is 3840s. Such a map can
be performed in less than  two hours including overheads.

We shall search for a restricted number of lines since the physical
conditions for emission are more restrictive than for absorption. The
sources have been selected as dense PDRs with different illumination
conditions G$_0 = 100 - 10^5$, with and without ionisation fronts. 
We propose to search for CH and CH$^+$ as tracers of the first steps
of the chemistry. Water and H$_3$O$^+$ are interesting for two reasons
described above : i) as H$_2$O cooling is important in dense regions, 
and ii) as
water is the main constituent of the grain mantles, we want to search for
evidences of the  photodesorption of grain mantles. Both the
spatial and spectral information are needed for this purpose.
Finally, ammonia will trace the dense well shielded molecular gas
and help us to separate the emission from the PDR  from that of
the bulk of the molecular cloud.

Note that the reactive ion
CO$^+$ has been detected in two sources, the Orion Bar and NGC~7023
(Black 1998 and references therein), which is a good eveidence
for dynamical processes at least for these two sources.
We have also included the young supernova remnant IC443 as a prototype
of shocked region in the ISM.

 
For 6 sources and 9 maps per source ([CII], CH, CH$^+$, H$_2$O (2 ground
state lines, ortho species at 1.7 THz and para species at 1.1 THz) and
H$_2^{18}$O (same 2 lines),  H$_3$O$^+$, NH$_3$) it sums to a total 
of {\bf 6$\times$7$\times$2 = 94 hours} of telescope time.

{\sl Ground based data on the spatial distribution of the
submillimeter C and CO lines should be obtained in the coming years to
set the optimum size an position of the regions to be mapped with
Herschel, and provide the necessary information on the warm molecular
gas. There is no need to perform these observations with Herschel as
the C and CO lines are bright and accessible from the ground.}


\begin{table}[!h]
\caption{Photo dissociation regions and shocks }

 \begin{tabular}{lccl}
Name & RA(2000) & DEC(2000) & Comment \\
\hline
{\bf PDRs}\\
NGC~7023 & 21 01 36.92 & +68 09 47.8 &  Origin on HD~200775 \\
Orion-Bar &05 35 22.90 & -05 24 57.8 &  Origin on $\theta_2$
Ori (HD~37041)\\
NGC~7027 & 21 07 01.7 &  +42 14 10 & Young PNe\\
Horsehead nebula & 05 40 59.0 & -02 27 30 &  \\
IC~63 & 00 59 01.37 & +60 53 17.8 &  \\
\hline
{\bf Shocks}\\
IC~443-G  & 06 16 43.0  & 22 32 38.0 & \\
\hline
\end{tabular}
\end{table}

{\bf Complementary data}   
Finally, the spatial structure of the continuum sources must be known
to interpret properly the absorption data. Since the sources are
bright, they  can be observed almost instantaneously. We
have chosen an integration time of 5 minutes
for each instrument. For a total number of 8 sources,
the required time for photometry is 
{\bf $7 \times 0.12 \times 2 = $  2 hours}.


\newpage
\section{Transition from atomic to molecular gas}

        Due to the large opacity of earth atmosphere to
far infrared and submillimeter radiation, light interstellar molecules,
with rotational transitions  below 1mm, are less well
known than heavier molecules, radiating in the millimeter region.
Though hydrides such as OH, CH, CH$^+$ have been known as 
interstellar molecules
since the  1940's and 1950's, their production is not yet completely
understood. In fact, the abundance has been measured over a
much smaller range of physical conditions than for e.g. CO or HCN.
Also the kinematic properties which provide  very important
information concerning the physical conditions, are poorly known 
for species observed through their electronic transitions.
In some exceptional cases, data with very high spectral resolution 
have been obtained, in which  CN, CH and CH$^+$ present different line
profiles : while CN electronic transitions have a similar velocity 
structure than  the ground 
state rotational line of carbon monoxide seen in emission, the CH$^+$
line is broader, and has a different velocity structure.

Chemical models describe the production of interstellar molecules
through   a series of two-body chemical reactions in the gas
phase, coupled with solid phase reactions on dust grains. Molecular
hydrogen is the most important molecule formed on the dust grain surfaces,
but similar processes are efficient to produce hyrogenated species. 
In the interstellar chemistry framework, molecules are formed 
gradually from atoms and ions.
Small light molecules convey  very important chemical information
since they are formed in the first steps of chemical networks, and are 
therefore sensitive to a small number of chemical reactions.
Models of interstellar sources have been devised in the past
twenty years, which describe the
interstellar medium with an increasing degree of complexity, but
are still far  simpler than the structure revealed by the data.
The simplest models are steady state, with a single density and
temperature. More sophisticated models deal with time dependent
effects, density structure, solve the thermal balance etc., but still
fail to reproduce the observed abundance of important species unless
 they introduce a coupling between the chemistry and the gas dynamics.

\begin{figure}
\label{chchp}
 \resizebox{17cm}{!}{%
 \rotatebox{270}{
   \includegraphics*[6cm,0.1cm][19cm,28cm]{chchp_3.ps} }}
\caption{The column densities of CH$^+$ (left) and CH(right)
as a function of the total hydrogen column densities. Data points
are taken from (Crane et al. 1995, Gredel 1997).}
\end{figure}
 
The likely reason for this failure is the still incomplete  description
of the structure and kinematics of the ISM in these models,  leading to
an incomplete coverage of the relevant physical and chemical
processes. In fact, diffuse interstellar clouds appear more active
chemically than anticipated, given their low gas density and average
temperature (Lucas \& Liszt 2000, 2001).
This unexpected chemical activity 
has been known for years, after the detection of large amounts of
CH$^+$ and CH in absorption against nearby stars. 
Large abundances of molecules like HCO$^+$ and OH 
have been observed recently in regions with no CO emission
nor absorption (see Figure 2 from Lucas \& Liszt 1996).
 Contrary to CH, which shows
a tight correlation with the gas column density, CH$^+$ is
less tightly correlated with total gas column density, as shown in 
Fig. \ref{chchp}.
The observed  level of CH$^+$ column densities (10$^{12}$ to 10$^{14}$
cm$^{-2}$ are difficult
to explain because the formation paths of molecules such as CH$^+$ 
 require triggering of endothermic reactions 
or overcoming activation barriers. Their formation 
therefore requires the presence of energetic events
and, whatever these events are, HIFI is ideally suited to detect lines in 
this highly excited, probably small scale, component of the interstellar 
medium. Magneto-hydrodynamical shocks (Draine 1980,
Flower \& Pineau de For\^ets 1998) fulfill these criteria, but
predict a velocity shift between neutral and ions 
($\sim$ 2 kms$^{-1}$ in the recent comprehensive models by 
Flower \& Pineau des For\^ets 1998).  An alternative would be that             
these reactive species form in 
regions in atomic clouds where bursts of turbulent dissipation locally heat
the gas to temperatures as large as 10$^3$K (Falgarone et al. 1995;
Joulain et al. 1998 A\&A 340 241 ). In that case all species would
form with the same centroid velocity, as observed in most
lines of sight.

 CH$^+$, as other reactive molecules (H$_2^+$, CO$^+$, OH$^+$, 
H$_2$O$^+$etc.), 
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).


HIFI is ideally suited for such studies since it gives access 
to the ground state
rotational lines of most interstellar hydrides, and most importantly,
provides the critically needed high spectral resolution. The sensitivity
is such that  we shall be able to probe the abundance of interstellar
hydrides  in a large range of
physical conditions, which overlaps the domain probed by
electronic transitions for e.g. CH and CH$^+$.
Studying hydrides, HIFI will be able to investigate the role of 
chemistry on dust grains and 
small aromatic hydrocarbons as well.     
Dust observations from mid-IR to sub-mm wavelengths (IRAS, COBE and ISO)
show that significant dust growth by gas accretion and/or coagulation 
occurs even within translucent clouds where ices are not observed.
The accretion of gas atoms on grains may be followed by the formation
of simple molecules such as hydrides on grains in the same way as H$_2$. 
Such a formation process has been
proposed to explain the detection of NH in the diffuse medium
(Crawford \& Williams 1997, Goiecochea et al. 2000).

\subsection{Proposed observations : survey for absorption lines}

In the diffuse ISM, HIFI is able to search for absorption features
due to hydrides. The selection criteria for the sources should be
the availability of a strong far infrared continuum and the presence
of interstellar clouds along the line of sight of the continuum
source. The sample is built with targets already studied by ISO and
SWAS (SgrB2 Neufeld et al. (1997), W49N Caux et al. (2002)) 
as well as HII regions studied at mm wavelenghs. In addition,
we searched for a few specific lines of sight bright in the visible
and the far infrared to allow a meaningful comparison with previous work.
The selected sources are located at different galacto-centric distances 
and have different properties to allow us to sample
a wide variety of physical conditions.

>From published ground based and space data, we expect continuum 
intensities of hundreds of
janskies in the HIFI domain. High spectral resolution data (1 MHz) can
be obtained with a S/N ratio for the continuum larger than 10 in
a few minutes of integration for most sources. The average time spent
on each source is estimated to be 15 minutes per tuning. 
The most intense lines for interstellar hyrides are listed in Table
\ref{linelist}. The third column gives an estimation of the
line opacity for a total column density of 10$^{13}$cm$^{-2}$ and a
line width of 1 kms$^{-1}$. The most promising lines
are the ground state lines of CH (2 tunings), 
CH$^+$ (1), NH (1), NH$_2$ (1), NH$_3$ (1),
NH$_2$D (1), H$_2$O (8 lines for $^{16}$O, $^{18}$O and $^{17}$O isotopes) 
HDO (1), D$_2$O (1) H$_3$O$^+$ (1), HF (1) and DF (1), HCl (1).

The total
number of lines per source is 20 including deuterated species and isotopes.
With 17 sources, the total time is {\bf 20$\times$17$\times$0.25 = 85 hours}. 


\begin{table}[!h]
\caption{Absorption line sources}

\begin{tabular}{lccccl}
Name & RA & Dec & F100  & F350 & Comment \\ 
& 2000 & & (Jy) & (Jy) & \\
\hline
RAFGL2136 & 18 22 29.1 & -13 29 46 & 5696 & 100 & \\
RAFGL2591 & 20 29 24.9 & +40 11 21 & 5721 & 100 & \\
RAFGL 490 & 03 27 38.4 & +58 47 04 & 785 & 45 & \\
W33A & 18 14 39.9 & -17 51 59 & 6310 & 400 & \\
W3-IRS5 & 02 25 40.7 & +62 05 52 & 42100 & 200 & \\
W3-H$_2$O & 02 27 04.6 & +61 52 25 & 10600 & 498 & \\
S140-IRS1 & 22 19 18.4 & +63 18 45 & 13460 & 400 & \\
RAFGL7009S & 18 34 19.7 & -05 59 44 & 2136 & 200 & \\
NGC7538-IRS1 & 23 13 45.38 & +61 28 10.7 & 14100 & 301 & Confused IRS1-IRS9\\
NGC6334-IRS1 & 17 20 00.8 & -35 57 33 & 19800 & 1263 & extended
emission, IRAS source 1 arcmin \\
W49N & 19 10 13.2 & +09 06 12 & 36200 & & \\
W51 & 19 23 43.90 & +14 30 34.8 & 26800 & & IRAS 1 arcmin \\
Sgr B2 & 17 47 20.4 & -28 23 07 & 66600 & 3962 & SgrB2(M)\\
 Sgr A$^*$ & 17 45 39.95 & -29 00 28.2& 23300& & \\
GCS-3 & 17 46 14.9 &  -28 49 37 &  410 & Visible data possible \\
Cyg OB12 & 20 32 41.1 & +41 14 28 & 131 &  Visible data available\\
\hline
IRAS05389-6908 & 05 38 36.7 & -69 06 29 & 2490 & & LMC, 40'' from Sk -69 243\\ 
\hline
\end{tabular}

Fluxes from IRAS point source catalog, Sandell (1994), van der Tak et
al. (2000), Goldsmith et al. (1990)
\end{table}

{\bf Complementary data}   
Finally, the spatial structure of the continuum sources must be known
to interpret properly the absorption data. Since the sources are
bright, they  can be observed almost instantaneously. We
have chosen an integration time of 5 minutes
for each instrument. For a total number of 17 sources,
the required time for photometry is 
{\bf $17 \times 0.12 \times 2 = $  4 hours}.




\newpage


\section{Molecular complexity}
The inventory of the  molecules present in the interstellar medium
is fairly incomplete. For example, the identification of the
carriers of most Diffuse 
Interstellar Bands (DIBs) is still pending. 
The full opening of the far infrared
and sub-millimetre domain with HSO will offer
the opportunity to search for many new  species in the interstellar medium
(e.g the contributions J. Goicochea \& J. Cernicharo in the Toledo meeting,   
present detailed calculations of the expected spectra 
for  some molecules).

\begin{figure}[!h]
  \begin{center}
\resizebox{16cm}{!}{%
\rotatebox{270}{
\includegraphics*[14.5,0.1cm][20.5cm,27.5cm]{comp-chim1.ps}
}}
 \end{center}
\caption{ Prediction of the abundance of hydrocarbons having one (black)
 to seven (purple) carbon atoms  using two different
chemical networks : the New Neutral
Neutral chemical model from E. Herbst and the UMIST95 data rate.
The calculations are performed for a moderate
gas density (n$_H$ = $2 \times 10^4$ cm$^{-3}$),  illuminated by a mild UV radiation
(G$_0$ = 100). Carbon chains are shown
with full lines, C$_n$H species with dashed lines and C$_n$H$_2$ 
species with dotted lines. There is a clear tendency to have
C$_n$ species more abundant than the corresponding C$_n$H
species, more abundant than the corresponding C$_n$H$_2$ species.
(E. Roueff, private communication).  }  
\label{gerin:fig4}
\end{figure}

 \subsection{carbon chains and cycles}
It is  expected that  some large molecules, bridging the gap between the
known gas phase species and the Polycyclic Aromatic Hydrocarbons
(PAHs), survive in the diffuse ISM. Large molecules are good candidates for some
Diffuse Interstellar Bands, though no definitive identification
has been made yet. The strength of these absorption
features  require quite high abundances for their carriers
(Herbig 1995, 2000). Carbon is generally thought 
to be a major constituent
of the DIBs carriers, therefore, carbon chains and
other hydrocarbons have been suggested as possible carriers
(e.g. Thaddeus et al. 1995).
Carbon chains (linear or cyclic) are also very good 
candidates for Herschel since they have low lying vibrational modes
in the far infrared. Hence absorption measurements in the far infrared
will allow to detect symmetric species which escape
detections in the radio domain. For example, the C$_3$ radical has been
recently detected towards the Galactic Centre by ISO
 (Cernicharo et al. 2000). 

It is possible to use gas phase chemical models to
obtain predictions for the abundances of some large
molecules.  Though the chemical networks have not
been as extensively tested for these large species as
they are for smaller species, these models provide
some inputs to drive the observational effort. For example,
we show in Figure \ref{gerin:fig4} predictions for
a PDR (Le Bourlot et al. 1993) obtained with two different chemical networks,
the "New Neutral Neutral" chemical
network from E. Herbst (Terzevia et al. (1998, 2000)) and the UMIST95
network. From this model, as from preliminary observations on
small hydrocarbons (Foss\'e, 2001 and in prep., Liszt and Lucas 2001) hydrocarbons
may have a significant abundance in photon dominated regions.
There is a clear pattern  with C$_n$ species
being more abundant than C$_n$H species, which are themselves
more abundant than C$_n$H$_2$ species. C$_n$ species cannot
be detected in the radio domain as they have no permanent dipole
moment, but could be seen in the  diffuse ISM with Herschel through
their ro-vibrational far infrared spectrum.

As the ro-vibrational frequencies are not accurately known at
present, it is difficult to set a particular frequency range and
integration time. This search should be combined with the line
surveys.
The best target is probably one of the strong continuum sources in
the SgrB2 molecular clouds, either SgrB2N or SgrB2M, since these lines
of sight cross several clouds along the line of sight with
different chemical properties. SgrB2 is listed already in the
hydrides section. The search for carbon chains and cycles will
be part of the line surveys and does not require additional
time given the current knowledge. 

\subsection{Search for low energy bending modes of PAHS}

Polycyclic aromatic hydrocarbons (PAHs) have been proposed as
the carriers of the Aromatic IR (AIR) bands at 3.3, 6.2, 7.7, 8.6,
11.3 and 12.7 $\mu$m since about twenty years. Interstellar PAHs
absorb UV photons and they cool down by emitting in their vibrational
modes (L\'eger et al. 1989, Allamandola et al. 1989).
The identification of interstellar PAHs has motivated a lot
of observational, laboratory, quantum chemical and modelling studies.
Although progress has been made, identification of a specific PAH
has remained elusive. Part of the difficulty
is due to the fact that all PAHs have aromatic CC and CH bonds
which give IR modes close to the AIR bands.
One idea to circumvent these difficulties is to look for the
rotational structure of the low-frequency IR bands of PAHs.
Indeed, PAHs are known to have flopping modes which
involve the deformation of the whole skeleton, some of them being
located in the THz domain. Detecting such a rovibrational structure
would help a lot in identifying the carriers. For a symmetric rotor,
the rotational constants are inversely proportional to N$_C^2$ and 
the spacing between the 2 rotational lines would therefore gives
an immediate idea on the size of the rotor. The position of the whole
band (the Q branch) is on the other hand specific of the molecule
itself. Finally, another interest is the characterisation of
the rotational temperature of PAHs which would have strong impact
in the models (cf. Rouan et al. 1992).

HIFI can have a very unique role in these studies regarding
the identification of "large" carbonaceous molecules such
as PAHs (but also carbon chains), through the rovibrational
structure of their low frequency bending modes. This programme
requires both I) a large spectral coverage,
and II) a good sensitivity.
\begin{itemize}
\item{I) The extend of the rovibrational structure depends on the
rotational constants (i.e. the molecular size) as well as on the
rotational temperature.
For instance, this extend is 105 and 26 GHz with a line spacing
of 667 and 42 MHz for C$_{24}$H$_{12}$ and C$_{96}$H$_{24}$ respectively. Note that
the line spacing assumes no anharmonic effects
(no vibration-rotation coupling). 
Whereas the rotational structure depends mostly on the molecular
size, the position of the band (Q branch) is very specific of the
identity of the molecule.}
\item{II) the low-frequency bending bands have weak
Einstein coefficients (A = 10$^{-4}$ s$^{-1}$ N$_C$$^{-1}$ or lower) and the best
strategy is to search for them in emission  where the UV/visible
radiation field is sufficiently high to pump them. Indeed, the bands are involved
quite efficiently during the last steps of the cooling of the
molecules.}
\end{itemize}

Two observation strategies can then be proposed :\\
1- We have a precise idea on the molecule we want to look for and
its spectrum is well characterised in the laboratory. 
2- We are ready for something new and certainly we do not
want to miss anything important (and for very long!). The identification
of these large carbonaceous molecules would really help us a lot
in our understanding of interstellar chemistry.

Though the first strategy  requires a more limited spectral coverage
and therefore allows a deeper integration, it is clear that
HIFI is a unique instrument to make decisive progress given its
ability to cover a broad spectral
range with high sensitivity but also because the  the low
frequency bending modes are expected to be more frequent and
more intense in the high frequency range (THz). The intensity of
these bending modes can be estimated as follows.
The emission in the AIR bands  is estimated to be
$I_{IR}=1.8 \times 10^{-21}G_{O}$\,Wcm$^{-2}$ in the
Herschel beam. According to recent calculations
(Joblin et al. 2002), the emission in the flopping modes
$\nu_{f}$ of PAHs can represent 0.1\% of the total IR
emission (in the best cases). This lead to a total
emission of $I(\nu_{f})=9.0 \times 10^{-19}$\,Wcm$^{-2}$
for the Orion Bar ($G_{O}=5 \times 10^{4}$).
Assuming a PAH with 1/30 of the total abundance of PAHs,
this leads to an intensity of $3.0 \times 10^{-21}$ Wcm$^{-2}$ emitted in
the low frequency IR band of this specific PAH. Now there are
about 400 lines per rotational structure (Trot=100K)
and the width of the individual lines
is assumed to be 3~MHz. This leads at 1 THz to an intensity of
1.46 10-20 W cm$^{-2}$ micron$^{-1}$, i.e. 3.1 Jy or 6 mK per line.
This estimation also does not take into account the
effects of anharmonicity that have to be included since
the emission in $\nu_{f}$ arises from several $v$ levels.
The numbers provided here should be used with caution
till more laboratory data are available.

To estimate the required observing time, we performed two calculations :\\
-- For the SIS mixers, we assume a typical system temperature
of 500K and a frequency resolution of 2.5 MHz. Then individual lines
are detected at the $3 \sigma$ level after 5 hours integration time.
The observations shall be performed in position switching, and cover
about 24 GHz, or 4 contiguous tuning.\\
-- For band 6, we use a system temperature of 1800 K and
a frequency resolution of 6 MHz. Then 10 hours are needed for
a single 4 GHz band.
As a whole, this program requires {\bf $\sim 60$} hours of
telescope time to reach the required S/N in two or three sources.
The sources will be selected from the list of bright PDRs,
chosing the position expected to maximise the 
PAH signal, and the efficiency of observation.












\newpage

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\newpage
\appendix

\section{Line frequencies for the ISM program}

Total number of lines for a complete survey : 27 \\
Total number of lines for the most abundant species with H,
C, N or O : 14\\ 

Other lines outside HIFI range :\\
H$_2$D$^+$  1.370146 THz \\
OH \& $^{18}$OH 2.5099 THz, 2.494695 2.494734 2.498994 2.499031\\
CH 2.00679 THz\\
OD 1.390613 THz, 1.391994 \\

Missing species with no frequencies yet :\\
\begin{itemize}
\item OH$^+$
\item{ H$_2$O$^+$}
\item{ H$_2$DO$^+$, HD$_2$O$^+$ and D$_3$O$^+$}
\item{NH$_3$D$^+$ and other NH$_4^+$ isotopes} 
\item{ $^{13}$CH$^+$ and CD$^+$}
\item{ $^{13}$CH and CD}
\item{ NH$^+$ and ND$^+$}
\item{ ND}
\item{ CH$_2$D$^+$ and otherCH$_3^+$ isotopes}
\item{D$_2$H$^+$}
\item{CH$_2^+$}
\item{H$_2$O$_2^+$}
\item{ Carbon chains and PAHs, including the positive ions and
radicals (e.g. C$_4$H)}
\item{ What else ?}
\end{itemize}


\begin{table}[!h]
\caption{List of lines}
\label{linelist}
\begin{tabular}{lrrl}
Name & Frequency   & $\tau$/1 kms$^{-1}$ & Comment\\
     & (GHz) & 10$^{13}$cm$^{-2}$ & \\
\hline
C & 492.161 & & \\
NH$_2$D & 494.454 & 0.62 & \\
NH$_2$D & 494.455 & 1.04 & \\
NH$_2$D & 494.457 & 0.21 & \\
CH &   532.724 & 0.25 & small difference JPL/Cologne \\
 & 532.793 & 0.10 & \\
 & 536.761 & 0.27 & \\
 & 536.796 & 0.11 & \\
H$_2$$^{18}$O-ortho & 547.676 & 1.80 & \\ 
H$_2$$^{17}$O-ortho & 552.021 & 1.80 & \\ 
H$_2$O-ortho & 556.936 & 1.78 & \\ 
$^{15}$NH$_3$ & 572.113 & 2.43 & \\
NH$_3$ & 572.498 & 2.46 & \\
D$_2$O & 607.3496 & 4.03 & \\
H$^{37}$Cl & 624.964 & 0.48 & \\
           & 624.978 & 0.72 & \\
           & 624.988 & 0.24 & \\
H$^{35}$Cl & 625.901 & 0.48 & \\
    & 625.919 & 0.72 & \\
    & 625.932 & 0.24 & \\
\hline
DF & 651.099 & 3.88 & \\
H$_2$S-ortho & 736.034 & 0.49 & \\
\hline
C & 809.342 & & \\ 
CH$^+$ & 835.071 & 1.17 & $\nu$ not accurate \\
HD$^{18}$O & 883.189 & 3.49 & \\
HDO & 893.639 & 3.54 & \\
D$_2$O       & 897.947 & 1.89 & \\
NH$_2$ & 952.542& 0.27 & many weaker lines \\
       & 952.550 & 0.10 & \\
       & 952.562 & 0.17 & \\
       & 952.572 & 0.20  \\
       & 952.573 & 0.10 &\\
       & 952.577 & 0.33 & \\
&      952.578 & 0.57 & \\
&      952.628 & 0.13 & \\
    & 959.426 & 0.16 & \\
    & 959.512 & 0.20 & \\
    & 959.526 & 0.12 & \\
    & 959.562 & 0.16 & \\
\hline
\hline
\end{tabular}
\end{table}

\begin{table}[!h]
Table \ref{linelist} continued

\begin{tabular}{lrrl}
\hline
NH & 974.462 & 0.12 & Many weaker lines \\
   & 974.470 & 0.21 & \\
   & 974.479 & 0.63 & Strongest multiplet\\
   & 999.973 & 0.17 & many weaker lines \\
H$_3$O$^+$ & 984.697 & 0.09 & \\
H$_2$$^{18}$O-para & 1101.698 & 4.00 \\
H$_2$$^{17}$O-para & 1107.167 & 4.00 \\
H$_2$O-para & 1113.343 & 4.00 \\
\hline
HF & 1232.476 & 3.91 & \\
\hline
N$^+$ & 1470 & & \\
H$_3$O$^+$&          1655.813 & 1.10 &      \\
H$_2$$^{18}$O-ortho& 1655.868 & 1.80 & \\
H$_2$$^{17}$O-ortho& 1662.464 & 1.80 & \\
CH$^+$ & 1669.59  & &  uncertain $\nu$ \\
H$_2$O-ortho &            1669.905 & 1.78 \\ 
C$^+$ & 1902.346 & & \\  
\hline 
\end{tabular}
\end{table}

\vspace*{5cm}

\newpage


\section{Deuterated hydrides}
The presence of  deuterium  in  the universe  is  an  anomaly  
that can  only  be
explained by Big Bang theories. It cannot be created any more, but can
only burn inside stars causing a drop of its abundance with respect to
hydrogen as function of time. The measurement of the
deuterium abundance relative to hydrogen is an important constraint
on primordial nucleosynthesis. Due to its high spectral resolution
and sensitivity,  HIFI can measure accurately the deuterium
fractionation for interstellar hydrides.                 
Deuterated hydrides  are chemically important in the early stages
of star formation, whereas the deuterated  form of  some hydrides  
play a  unique role  in cosmology.
In cold dense prestellar cores,  the abundances of deuterated
species are  enhanced over the gas phase D abundance by chemical
processes. This large deuterium fractionation  may
be kept through the cycle of star- and planet formation, and is in
fact used to confirm  the  primordial origin of some meteorites.  

Here  we  propose  to  observe  the  fundamental  transitions  of  the
following  (deuterated-)  hydrides  by  means of  their  ground  state
transitions: NH$_3$, NH$_2$D, HF, DF,  H$_2$O, H$_2$$^{18}$O, HDO, 
HD$^{18}$O, D$_2$O, OD, H$_2$D$^+$ (the latter two transitions 
lie currently in the gap between HIFI Band
V and band VI).

This program consists of three parts:

\begin{itemize}
\item{I){\bf Is the HDO abundance conserved in the formation 
process of low-mass, solar-like, stars ?} \\
 The [HDO]/[H$_2$O] abundance   ratio is rather similar for different
stages  in   the  star formation    cycle  from  molecular   clouds 
 via   YSOs  to
planets. Observations  yield: $3.0  \times 10^{-4}$ in comets 
 (Hale-Bopp, Haley,
Hyakutake), $2.4 \times 10^{-4}$ in hot cores and high-mass YSOs,
$ 8 \times 10^{-4}$ in low-mass
YSOs (Stark  et al. 2002) and $ 1.6 \times 10^{-4}$ in the Earth's
oceans. These  values lie all
closely together  and suggest  that the ratio  may be constant  in the
cycle of star formation. The fundamental question is then where and when is
the HDO/H$_2$O ratio determined? There are two favoured scenarii (Aikawa
\& Herbst, 1999, 2001): 1) Deuteration from ion-molecule chemistry in the 
protoplanetary disk;   2) The deuteration is determined 
by gas phase and/or solid phase chemistry
 in the parent  molecular cloud before star formation sets in, 
 i.e. in  the dense  ISM. The
deuterium enrichment  is propagated  in e.g. HDO  through ion-molecule
and neutral-neutral reactions in the  gas-phase as well as solid state
reactions  with  atomic  deuterium  on grains.  To  investigate  these
scenarii, we propose to study  the fundamental transitions of HDO and
H$_2$$^{18}$O in  objects spanning the full range of
star formation : ranging from  dark cores without star  formation, in
absorption against  background continuum sources,  in Class 0,  I, II,
III YSOs,  in the solar  system (planets, comets).  Since (deuterated)
water is already the subject of many studies with HIFI we 
do not propose any integration time yet. The project should be
coordinated later with the star formation and comet projects.
It is important to ensure that a sufficient integration time 
is spent on these
lines so that a uniform  sensitivity is reached for all objects.

The main lines for this search are the ground  state lines of
HDO , H$_2$O (2 lines, ortho and para) and H$_2$$^{18}$O(2 lines ortho
and para)}


\item{II) {\bf Determining the
galactic D/H  gradient by  means of observations  of the  ground state
transitions of  HF and DF}.\\
 The  J=2-1 rotational transition  of HF has
been discovered with ISO-LWS in absorption  toward SgrB2N by Neufeld
et al. (1997).  As for the other hydrides, we  most likely 
will detect  lines in absorption in
cold regions against a  background submillimeter continuum e.g. in the
dense warm envelopes  around YSOs. The simple chemistry presented by
Neufeld et al. (1997) predicts that HF is formed by neutral-neutral
reactions between fluorine and H$_2$ or H$_2$O. Since HF is a very
stable molecule, not easily destroyed in the ISM, it is the main
fluorine reservoirs. Reactions with deuterium have not been
studied but the [DF]/[HF] fractionation ratio is expected to
depend o a very limited set of chemicalreactions. Since the
fractionation for molecular hydrogen and water is very low,
the deuterium fractionation for HF is expected to be very low also,
hence to track accurately the elemental Deuterium abundance.   
  By measuring the  [DF]/[HF] ratio  we will
have  a new  tool to  determine the  D/H ratio  in the  galaxy 
to be combined with HD measurements on the same sources.
{\sl This program should be coordinated with PACS to ensure a
good wavelength coverage of the most interesting sources.}
}
\item{II){\bf Studying  the  basics  of  HDO  and  D$_2$O  chemistry}  
by  means  of
observations  of the  fundamental  lines  of HDO,  D$_2$O,
H$_2$$^{18}$O fundamental  lines  toward  high  mass   star  forming
regions,   combined  with H$_3$O$^+$  and H$_2$DO$^+$.  }
\end{itemize}

The integration time for the sections relevant to
the interstellar mediun  has been included in the main hydrides section.


{\sl The following sources should be looked at in the deuterium lines. Most are
well known low mass star forming regions, to be checked with the 
low mass star proposal :
IRAS16293-2422, VLA1623, NGC1333IRAS2, NGC1333IR4A, NGC1333IR4B, L1157, 
L1551-IR5, L1489, TMR1, L483, L1527. }



\newpage




\section{The $^{12}$C/$^{13}$ abundance gradient in the Galaxy}

The fine structure line of $^{13}$C$^+$ can be observed simultaneously
with the $^{12}$C$^{+}$ line at 1.9 THz. Since little fractionation 
of $^{13}$C is expected in regions where the carbon
is ionized (Keene et al. 1998), the detection of the faint
$^{13}$C$^+$ lines simulatneaously with the strong $^{12}$C$^+$ line
permits to obtain a good measurement of the elemental
$^{12}$C/$^{13}$C abundance ratio.  The weak $^{13}$C$^+$ features
appear at velocity offsets of  -63, 11, and 65 km s$^{-1}$ relative to
the velocity of the $^{12}$C$^+$ line. Signal to
noise ratios in the C$^+$ line will be large ($>100$) in the bright star forming
regions, providing a good signal-to-noise ratio in the $^{13}$C$^+$
lines as well.  

In the Galaxy, previous
work using molecules have found a smooth gradient from the Galactic
Center to the solar neighborhood. Using HIFI, a more accurate and less
biased determination of this abundance gradient can be obtained.  
About fifteen sources sampling the Galactic plane have been selected in the list 
of bright HII regions osberved with the ISO spectrometers (Peeters et
al. 2002). For all sources the [CII] line flux as seen by ISO-LWS is
larger than 10$^{-17}$Wcm$^{-2}$ and the source is not spatially resolved by the
instrument. So we expect to detect a line intensity
larger than $100$ Kkms$^{-1}$  for the $^{12}$C$^+$ line. The
$^{13}$C${^+}$ features are expected at the $\sim 1$Kkms$^{-1}$ level
with a line width of a few kms$^{-1}$, and can be detected after a few
hours of integration. The exact figure depends largely on the
observing mode since the sources are located in fairly crowed
regions of the sky. Assuming a maximum of 4 hours per source,
it sums to a total of {\bf 4 $\times$ 15 = 60 hours} of
HIFI time for this section.

\begin{table}[!h]
\caption{Compact HII regions for the $^{12}$C/$^{13}$C Galactic
gradient}
\begin{tabular}{lcccc}
Source & RA(2000) & DEC(2000) & Distance & [CII] \\
\hline
SgrC & 17 44 35.6 & -29 27 29.3 & 0.3 & 31.6\\
IR01420 & 01 4 39.6 & 64 12 02.1 &   9.5 & 23.5\\
IR10589 & 11 00 59.8 & -60 50 27.1 &  8.0      & 14.3\\
IR12073 & 12 10 00.3 & -62 49 56.5 &   10.8      &24.3\\
IR15502 & 15 54 06.0 & -53 11 36.4 &   4.6        & 13.7\\
IR16128 & 16 16 39.3 & -51 16 58.3 &  5.5        & 19.4\\
IR17160 & 17 19 26.1 & -37 10 53.8 &    3.0       & 16.3\\
IR17455 & 17 48 41.5 & -28 01 38.3 &   0.2     & 12.8 \\
IR18116&  18 14 35.2 & -16 45 20.6 &   4.3     & 19.0\\
IR18317&  18 34 24.9 & -07 54 47.9 &    4.5     & 18.4\\
IR19442& 19 46 20.1 & 24 35 29.4&      7.6   & 17.9\\
IR19598&    20 01 45.6 & 33 32 43.7 &    9.8    & 14.2\\
DR21&  20 39 00.9 & 42 19 41.9 &    8.6    & 12.8\\
IR23030& 23 05 10.6 & 60 14 40.6 &   11.4    &12.0\\
IR23133& 23 15 31.4 & 61 07 08.5 &  11.7     & 12.9\\
\end{tabular}
{\small
The Distance is given relative to the Galactic Center in kpc.\\
The [CII] flux is given in 10$^{-18}$Wcm$^{-2}$ (Peeters et al. 2002)}
\end{table} 
\vspace*{5cm}







\end{document}















