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{\Large Molecular carriers in the ISM}

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\sl{M. Gerin$^1$, E. Falgarone$^1$, T. Giesen$^2$, P. Goldsmith$^3$, 
E. Herbst$^4$, C. Joblin$^5$ J. Krelowski$^6$, W. Langer$^7$, J. Martin-Pintado$^8$,
B. Mookerjea$^2$,
D. Neufeld$^9$, J. Pearson$^7$, T. Phillips$^{10}$, M. Spaans$^{11}$,
J. Stutzki$^2$, D. Teyssier$^{11}$, C. Vastel$^{10}$}

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\vspace*{0.5cm}

{\small \noindent
$^1$ LERMA, Observatoire de Paris - France\\
$^2$ KOSMA - Cologne University- Koeln - Germany\\
$^3$ Cornell University - USA \\
$^4$ Ohio State University - USA \\
$^5$ CESR, Toulouse - France \\
$^6$ Torun - Poland \\
$^7$ JPL, USA \\
$^8$ IEM- CSIC - Madrid, Spain \\
$^9$ John Hopkins University, Baltimore, USA \\
$^{10}$ Caltech - USA \\
$^{11}$ SRON - Groningen - Netherlands }

\newpage

\centerline{ABSTRACT}

\vskip 0.2 cm
 We will carry  out a  comprehensive spectroscopic
study of molecular carriers,  probing interstellar hydrides and carbon
chain and  ring molecules.  Our investigation of  hydrides will entail
high-resolution  HIFI spectroscopy  of some  20 hydride  and deuteride
molecules   towards  18  bright   continuum  sources,   together  with
complementary observations  with PACS and  SPIRE.  The list  of target
hydrides includes  molecules containing the  elements C, N, O,  F, Cl,
and Fe.  We will detect multiple absorption components from foreground
gas  clouds of  diverse properties  that  are known  to intersect  the
selected lines-of-sight, along  with emission and absorption intrinsic
to the background continuum sources themselves. Our investigation will
provide  a wealth  of  new  information about  the  abundances of  the
interstellar  hydrides --  addressing  key puzzles  posed by  previous
observations  from the  ground and  with ISO,  SWAS, and  ODIN  -- and
leaving   an   important  Herschel   legacy   to  astrochemistry   and
interstellar medium science.  In  particular, we will address the role
of   high  temperature   chemical  reactions   in  the   formation  of
interstellar molecules,  and the question of how  such reactions might
be driven by  heating mechanisms such as shock  waves and interstellar
turbulence.  Our study  will also elucidate the role  of grain surface
reactions  in  interstellar   chemistry,  reactions  that  are  likely
critical  in  controlling   the  inventories  of  molecular  materials
available within  nascent planetary  systems. We will  investigate the
growth of carbon molecules, bridging the gap between molecules
and aggregates. FIR spectroscopy has the unique capability of enabling
the detection  of bending modes  of large carbon molecules  (chains or
rings, possibly  PAHs), securing their  spectroscopic identification.
Many of the absorption lines that are detectable only with Herschel in
the local Universe become  accessible to ground-based observatories at
moderate to  high redshift.  Thus, absorption lines  in quasar spectra
are a potential probe of molecular gas at high-z; in the submillimeter
spectral  region,  the high  sensitivity  of  ALMA  will generally  be
required.  Herschel  observations of absorption in  the local Universe
will  provide an important  benchmark for  the future  observations at
high redshift that will become possible with ALMA.   

\newpage

\section{Introduction}

Among the many  molecules  identified in the interstellar medium and 
circumstellar envelopes so far (over 120, including  
isotopologues bearing rare isotopes such 
as Deuterium, $^{13}$C, $^{18}$O and $^{17}$O, etc.) a significant
subsample has been detected in the 
diffuse interstellar medium, ranging from H$_2$ and HD (Rachford et
al. 2002) to rather
complex species such as C$_3$ (Oka et al. 2003), and
c-C$_3$H$_2$ (Lucas and Liszt 2000).

\begin{table}
\caption{\label{tab:dif-ism} Diffuse ISM species}
\begin{tabular}{cccccccc}
\hline
\hline
%\vspace*{0.1cm}
%{\bf Detected species}\\
\noalign{{ \bf Detected species}}\\
%\hline
H$_2$ & HD & H$_3^+$ & CH & CH$^+$ & CH$_3$ & OH & H$_2$O \\
NH & NH$_2$ & NH$_3$ & H$_2$S & HF & HCl \\
C$_2$ & C$_3$ & CCH & c-C$_3$H$_2$ & CN & HCN & HNC   \\
CO & CO$^+$ & HCO$^+$ & HOC$^+$ & H$_2$CO & CH$_3$OH & CH$_3$  \\
 HCS$^+$ & CS & SO & SiO\\
\hline
\end{tabular}
\vspace*{0.3cm}
\begin{tabular}{cc}
\noalign{{\bf Unknown carriers of detected spectral features  }}\\
%\hline
DIB carriers & AIB carriers (PAHs) \\ 
\hline
\hline
\end{tabular}
\end{table}

Molecules are  found in the so-called  ``Cold Neutral Medium'', 
where the kinetic temperature
ranges from a few tens to one hundred Kelvins, the gas density ranges between  
a few hundred and a few thousands cm$^{-3}$, and the velocity
dispersion is about 1 kms$^{-1}$. The average pressure
 $\sim 3 \times 10^3$ Kcm$^{-3}$, with large excursions above
the median, up to $10^5$ Kcm$^{-3}$ (Jenkins \& Tripp 2001).
 The sensitivity of the visible and FUV observations
 is such that the extinction threshold for detecting H$_2$
has been extended from $A_V \sim 0.1 mag$ up to $A_V \sim 3 mag$
(Rachford et al. 2002). Electronic transitions
of a few other molecules are  detected with very high sensitivity
in the UV/visible/infrared, the most common being CO, CH, CH$^+$, CN, 
C$_2$, C$_3$  and H$_3^+$.  Liszt and Lucas (2001 and references therein) 
have used background radio sources for detecting
many additional molecules  in the diffuse ISM, including
OH, HCO$^+$, HCN, H$_2$CO, CCH, and c-C$_3$H$_2$.
 
The standard framework of the diffuse interstellar medium chemistry 
is fairly simple  because of the rather low gas
density, which limits the efficiency of chemical reactions
for building molecules, and because the interstellar
FUV radiation efficiently destroys  molecules. In this framework,
molecules are built  from atoms and ions in 
a series of two body chemical reactions,  and the set of
reactions can be solved assuming steady
state. The whole set of
chemical reactions is called a chemical network. At the moderate
temperatures encountered in the diffuse ISM, exothermic reactions
without activation barriers
are preferred for building the chemical network. Endothermic
reactions, reactions with activation barriers, and radiative
associative reactions, which have much smaller rates, are included in
a few cases, when these reactions are the sole routes for forming the
interstellar species.
It is therefore  expected from the chemical models, that a restricted set of 
small and light molecules can be formed in the diffuse ISM, and
 that the chemistry is not efficient in building complex molecules
 (i.e. molecules with three or more heavy atoms).

However, even for low extinction lines of sight, simple  chemical
models encounter serious difficulties for reproducing  observed molecular
abundances. The failure occurs not only 
for the reactive ion
CH$^+$ -- which has a very specific chemistry, see below -- 
but also for OH and HCO$^+$ whose chemistry is 
%JST add: 
thought to be %JST
well understood, at least in  dense, shielded gas (Liszt and Lucas 1996). 
The unexpected  detection of the simple species
H$_3^+$ (McCall et al. 2003)  also does not fit in the standard framework.
The diffuse interstellar medium therefore appears to be more chemically
active than expectated based on steady-state chemical models. 
To proceed, the chemistry has to tap the large reservoir of non-thermal energy
provided, for instance, by supersonic turbulence.
Transient velocity and density structures, driving the dissipation of
turbulence, certainly
impact the chemistry, either by locally heating the gas well above
its average value, or by generating large ion-neutral drifts.
(e.g. Flower \& Pineau des For\^ets 1998; Joulain et
al. 1998).
%The strong dynamical activity creating transient velocity and density
%structures (eg Audit \& Hennebelle , 2004 astroph-0410062) 
%certainly impacts the chemistry (e.g. Joulain et al. 1998). 
Velocity resolved line profiles are a prerequisite for understanding the 
chemistry, as they provide
simultaneous constraints on the molecular abundances, physical
conditions and velocity field.

In addition,  many spectral features 
remain unidentified in all wavelength ranges. For instance, the 
carriers of the strong spectral features seen in absorption in the  
visible spectrum of mildly reddened stars, and known as the Diffuse 
Interstellar  Bands (DIBs) remain unknown despite many decades of active 
research. Large molecules made of carbon are the best candidates to  
account for the DIBs. They can be found in different forms such as 
carbon chains, rings, fullerenes or polycylic aromatic hydrocarbons  
(PAHs). The former are best detected in the far infrared via their
ro-vibrational transitions (eg Cernicharo et al. 2002), 
the latter are  good candidates to account for the so-called 
``unidentified''or aromatic infrared bands, observed in emission between 
3 and 17\,$\mu$m in many places of our Galaxy, as well as in external 
galaxies. Large carbon molecules play important roles both in the physics  
and chemistry of interstellar gas. One of the main gas 
heating process is the ejection of energetic electrons from dust
grains via the photo-electric effect. The smallest particles are
 most efficient at releasing electrons. 
The contribution of large molecules, as PAHs, is essential
for  heating the interstellar gas to observed temperatures 
(e.g. Wolfire et al. 2003). As important reservoirs of carbon and
charge carriers, PAHs impact the gas chemistry, but quantitative
estimates are difficult without spectroscopic identifications,
as the chemical processes can be very specific to individual PAHs .

The understanding of the ISM
clouds  must include the full range of sizes of ISM line carriers, from
the smallest molecules (H$_2$ and hydrides) to the carbon clusters
and PAHs.
The opening of the full far-infrared and 
submillimeter wavelength domains by Herschel will provide a unique 
opportunity for progress in astrochemistry,  because it gives access 
to  important classes of key interstellar
species, which have not been studied with heterodyne techniques before.
 The list includes the rovibrational transitions of   large 
carbon molecules and carbon clusters, either with chains, rings or
aromatic structures, which are very specific of the identity of their
carrier, together with the ground state transitions of most
hydrides. The line strength of the ground state transitions of
these light polar molecules is such that they can be probed in the far infrared
and submillimeter wavelength range with a high sensitivity, 
matching the sensitivity of optical observations : a typical
figure is a line centre opacity  of unity for a column density of
10$^{13}$ cm$^{-2}$ and a line width of 1 kms$^{-1}$.

\begin{figure}
 \resizebox{12cm}{!}{%
   \includegraphics{w49-large} }
\caption{ \label{w49} SWAS spectra obtained towards the star forming
  region W49 from Plume et al. (2004). Top : the ground state water 
line at 557 GHz, Middle :
  the fine structure line of carbon at 492 GHz, bottom, the excited
  line $^{13}$CO(5-4). Diffuse clouds along the line of sight show up
  as absorption features in the H$_2$O spectrum and as emission
  lines in the CI spectrum, but are absent in the excited
  $^{13}$CO(5-4) spectrum. Absorption spectroscopy in ground state
  transition provides a sensitive mean to probe the diffuse ISM.}
\end{figure}




\newpage

\section{Interstellar hydrides}

\subsection{The astrochemical importance of hydride molecules}

%DN Modifications
%Hydrides can be important reserviors as well as intermediaries, e.g. HF
%The reduced mass rather than the mass sets the rotational constant.  Some 
%of our target molecules (e.g. FeH) are considerably heavier than CO, so 
%it may be confusing to talk simply about "light" molecules vs. heavy

Hydrides play a central role in interstellar chemistry, both as significant reservoirs of heavy elements and as critical intermediaries in the pathways leading to more complex molecules; {\bf observations of interstellar hydrides are therefore critical in advancing our understanding of astrochemistry.}  
Unfortunately, the large opacity of Earth's atmosphere at far infrared and submillimeter wavelengths has severely limited the ability of ground-based observatories to probe hydrides containing just a single heavy element atom; thus the properties of hydrides, with ground state rotational transitions at submillimeter 
wavelengths, are less well known than those of non-hydrides -- with larger reduced masses -- that radiate in the millimeter region.  The HIFI instrument on Herschel promises to yield great advances in our understanding of hydride chemistry.

%DN modifications
%Discussion of specific advantages of submillimeter/far-IR over optical/UV 

In diffuse clouds, hydrides can be probed by means of optical or
ultraviolet absorption line spectroscopy of background stars --
indeed, the hydrides CH and CH$^+$ were among the first three
molecules detected at optical wavelengths in the 1930's.  However,
submillimeter and far-IR observations of rotational transitions have
many critical advantages over optical/UV observations of electronic
transitions: (1) submillimeter/far-IR absorption line observations can
be carried along sightlines of large visual extinction, allowing the
chemistry of translucent and dense clouds to be probed; (2) many
hydrides have electronic transitions at UV wavelengths that are
inaccessible with current instrumentation and/or subject to confusion
with H$_2$ Lyman and Werner bands; (3) submillimeter/far-IR
transitions can be observed both in absorption and emission, providing
a more complete sampling of the molecular clouds under investigation;
(4) unlike submillimeter/far-IR observations with heterodyne
spectrometers, optical/UV observations are limited by spectral
resolution -- except for optical observations of the brightest
background sources -- and thus fail to elucidate fully the velocity
structure of the absorbing molecules (see Fig \ref{w49}). 

%DN added
%Summary of what we will do

In this Herschel key program, {\bf we will carry out a comprehensive study of some 20 hydride molecules towards 18 bright submillimeter continuum sources, leaving an important Herschel legacy to astrochemistry and interstellar medium science.}  The target molecules comprise key hydrides and deuterides of the elements, C, N, O, F, Cl, and Fe, and include both neutral species and molecular ions.
The proposed observations will simultaneously provide both absorption-line spectroscopy of intervening clouds along the sightline to the target sources, together with emission-line spectroscopy of hydrides located within the sources themselves; differential Galactic rotation will typically allow the various emission and absorption components to be disentangled kinematically.  

Many of the absorption lines that are detectable only with Herschel in the local Universe become accessible to ground-based observatories at moderate to high-redshift.  Thus, hydride absorption lines in quasar spectra are a potential probe of molecular gas at high-z; in the submillimeter spectral region, the high sensitivity of ALMA will generally be required.  {\bf Herschel observations of hydride absorption in the local Universe will provide an important benchmark for the future observations at high redshift that will become possible with ALMA} -- particularly in the case of hydrogen fluoride, discussed in \S 2.2.2 below.

%DN added
%Outline of the rest of the program and summary of common science themes

In  \S 2.2, we will discuss\footnote{The detailed discussion in \S 2.2 can be omitted by readers who wish to read an abbreviated description of the proposed key program.  \S 2.1 and \S 2.3 are intended to stand alone as a top-level summary.}  in the turn how the proposed program will advance our astrochemical understanding of molecules bearing each of the elements C, N, O, F, Cl, and Fe.  While specific puzzles will be discussed there, {\bf two general scientific questions emerge repeatedly:}

\vskip 0.2 cm
$\bullet$ {\bf What is the role of high temperature reactions in the formation of interstellar molecules, and how are such reactions driven?}  As we shall discuss below, the relative abundances of several species argue independently for the importance of chemical reactions with significant energy barriers.  Such species therefore yield an important probe of possible heating mechanisms, including shock waves and turbulence, that are of general importance to our understanding of the ISM.
\vskip 0.2 cm
$\bullet$ {\bf How do grain surface reactions affect the abundances of gas-phase molecules?} The likely importance of gas-grain chemistry has been recognized increasingly in recent years.  The question is of great and general astrophysical significance, because gas-grain processes control the composition of icy grain mantles in molecular clouds, grain mantles that participate in the process of star formation and ultimately supply the molecular inventories within nascent planetary systems. 

\subsection{The chemistry of interstellar hydrides}

%DN Moved oxygen chemistry up, since H2O is such a Herschel-unique target.
%Extended the discussion to emphasize current puzzles

\subsubsection{Oxygen chemistry}

The chemistry of oxygen-bearing molecules has been the subject of intensive investigation, particularly with the ISO, SWAS and ODIN satellite observatories, which have opened up the study of water vapor and molecular oxygen.  Recent observations have posed several key puzzles:

\vskip 0.2 cm
$\bullet$ {\bf Why is the water vapor abundance so small in cold,
  dense, quiescent molecular clouds?}  While enhanced water vapor
abundances have been detected in hot core and outflow regions (e.g.\
Boonman et al.\ 2003), in good agreement with theoretical
expectations, the cold cloud abundances fall at least one -- and
perhaps two -- orders of magnitude below the predictions of standard
astrochemical models (e.g. Bergin et al.\ 2001).  The effects of ice
mantle formation have been proposed as a possible explanation of the
low water abundances measured in the gas-phase. Some 
 water vapor is still expected to remain in the gas resulting
 from the evaporation of/ejection from mantles, and indeed observed
by ISO towards cold sources. To account for the low water vapor abundances
 observed in cold, dense clouds, another destruction mechanism
must be considered which will further decrease the gas phase water
vapor abundance, or the release from grain mantles must be inhibited
in the very dense and cold cores where water vapor abundances 
have been measured. The large abundance of atomic carbon in the same
cores, which remains not fully understood, may contribute in solving the
puzzle because atomic carbon efficiently cycles oxygen nuclei into CO as 
a result of its  reactions with OH and O$_2$.
%MG There is no definition of ``standard astrochemical model''. We may
% ask the question of ``what are the processes for releasing species from
% grain mantles ?'' here. How efficient are these processes ? 
% I agree with Tom that a significant fraction of Oxygen must remain
% in the form of atomic oxygen. Shall we try to find arguments and
% ways to estimate the abundance of gas phase O ? PACS observations
% could be useful.   
\vskip 0.2 cm
$\bullet$ {\bf Why is the molecular oxygen abundance so small?}  Here
we have only upper limits, and again the limits fall a factor $\sim
100$ below theoretical expectations (Goldsmith et al.\ 2002, Pagani et
al.\ 2003).  Again, the freeze-out of oxygen nuclei into icy grain mantles is a possible solution to the puzzle. 
\vskip 0.2 cm
$\bullet$ {\bf In diffuse clouds, why is the OH/H$_2$O ratio
  variable?}  In standard astrochemical models, where both molecules
are formed by the dissociative recombination of H$_3$O$^+$ and
destroyed primarily by photodissociation, the OH/H$_2$O ratio is
expected to show little variation.  Observations of substantial
variations from one sightline to another suggest the importance of
additional formation mechanisms for OH and/or H$_2$O (Neufeld et al.\
2002; Plume et al.\ 2004, Fig \ref{w49}).   These might include (1) formation via endothermic neutral-neutral chemistry, in a component of the gas where the temperature is enhanced by shocks or turbulence; and/or (2) formation via grain catalysis, in which oxygen atoms stick to grain surfaces, react with H to form OH and then H$_2$O and are then photodesorbed by the interstellar UV radiation field (Hollenbach et al.\ 2005).

Herschel promises to provide {\bf a critical and unique probe of water vapor in cold, absorbing clouds of various density and column densities lying along sightlines to submillimeter continuum sources.}  
Absorption features due to water lines are ubiquitous in the spectra of bright FIR and submillimeter continuum sources (e.g.\ Neufeld et al.\ 2002, Sandqvist et al.\ 2003), but thus far only the 
557 GHz $1_{10} - 1_{01}$ line has been detected at high spectral resolution.
SWAS has detected optically-thick water absorption against several sources,
so a promising target list is already in place.  Unfortunately, the SWAS
sensitivity was too poor to permit the detection of H$_2^{18}$O
along any sightline other than Sgr B2, so almost all water column density determinations
are lower limits derived from optically-thick H$_2^{16}$O lines. 
HIFI will have a big advantage both in sensitivity and because it 
will access both ortho- and para-H$_2$O.  We will target the lowest transitions of both ortho- and para-H$_2^{16}$O, the 557 GHz $1_{10} - 1_{01}$ and 1113 GHz $1_{11} - 0_{00}$ transitions respectively, as well as the analogous 548 GHz and 1102 GHz lines of H$_2^{18}$O.

Another key advantage of Herschel/HIFI will be {\bf the ability to study the molecular ions OH$^+$ and H$_3$O$^+$ in absorption along the same sightlines where water is observed.}  These species are key intermediaries in the cosmic-ray driven ion-molecule chemistry leading to OH and H$_2$O.  OH$^+$ is produced by proton transfer from H$_3^+$ to O, and subsequently undergoes hydrogen atom abstraction reactions with H$_2$ to form H$_2$O$^+$ and then H$_3$O$^+$, which ultimately undergoes dissociative recombination to form OH or H$_2$O.  Measurements of the OH$^+$ and H$_3$O$^+$ abundances will therefore probe directly the ion-molecule route to OH nd H$_2$O formation, allowing the relative importance of cosmic-ray driven chemistry to be compared with those of high temperature chemistry and gas-grain chemistry.

Using PACS in its spectrometer mode, we shall be able to complete our
investigation of oxygen chemistry with data on O (from the fine structure
lines at 63 and 145 $\mu$m) and OH (ground state transitions at 119,
79 and 53 $\mu$m and excited lines at 163 and 98 $\mu$m, see Goicoechea and
Cernicharo 2002 for OH spectra in the far infrared) at the same
spatial resolution as the HIFI data.

In the warm, dense gas within the target sources themselves, 
multitransition studies are needed to pin down the excitation conditions
and provide more accurate measurements of the
water vapor abundance.  For instance, with a multitransition study of water and its
isotopes, combining emission and absorption lines,
Comito et al.\ (2003) revisited the water abundance in the 
Sgr B2 cloud and inferred large variations of the gas phase
water abundance through the Sgr B2 molecular core.  To probe denser, radiatively-excited
regions, we will also observe the $2_{02} - 1_{11}$ and $2_{11} - 2_{02}$ lines of both para-H$_2^{16}$O
and para-H$_2^{18}$O.


%DN Revised to account for recent results with Wolfire and Schilke
%Moved up to increase prominence
%
\subsubsection{Fluorine chemistry}

The chemistry of interstellar fluorine is qualitatively different from that of any other element, because -- unlike the atom of any other element found in diffuse or dense molecular clouds -- atomic fluorine undergoes an exothermic reaction with molecular hydrogen.  Neufeld, Wolfire \& Schilke (2005) have recently presented a theoretical investigation of the chemistry of fluorine-bearing molecules in diffuse and dense interstellar gas clouds, combining recent estimates for the rates of relevant chemical reactions with a self-consistent model for the physical and chemical conditions within gas clouds that are exposed to the interstellar ultraviolet radiation field.  Over a wide range of conditions attained within interstellar gas clouds, the product of that reaction -- hydrogen fluoride -- is predicted to be the dominant gas-phase reservoir of interstellar fluorine nuclei.  {\bf Fluorine is expected to be the heavy element which shows the greatest tendency toward molecule formation; in diffuse clouds of small extinction, the predicted HF abundance can even exceed that of CO}, even though the gas-phase fluorine abundance is four orders of magnitude smaller than that of carbon.  

Model calculations predict HF column densities $\sim 10^{13} \,\rm cm^{-2}$ in dark clouds and column densities
as large as $10^{11} \,\rm cm^{-2}$ in diffuse interstellar gas clouds
with total visual extinctions as small as 0.1 mag.  Such diffuse
clouds should show HF absorption lines detectable in high
signal-to-noise HIFI spectra.  To date, the only molecules detected in
clouds of such small visual extinctions are H$_2$ and HD, observed in absorption toward hot background stars at UV wavelengths.  {\bf This raises the surprising possibility that spectroscopic observations of far-infrared continuum sources with Herschel/HIFI will reveal a component of foreground molecular gas that is observable exclusively by means of its HF absorption lines.}  

Because of the unique simplicity of the interstellar fluorine chemistry, the
abundance of HF can be used as powerful probe of the freeze-out of atoms and
molecules onto grains.  For example, the ISO detection of HF $J = 2 - 1$ in
absorption toward Sgr B2 (Neufeld et al.\ 1997) indicated an abundance 50 times 
smaller than the
cosmic abundance of fluorine nuclei, thus providing compelling evidence for
high levels of depletion in dense interstellar environments.  By contrast,
diffuse cloud observations of atomic fluorine with Copernicus implied
depletion factors of at most three.

ISO observations were limited to the $J = 2 -1$ line, which can be detected in
absorption only in regions like Sgr B2 that are subject to strong radiative
excitation.  By contrast, HIFI observations of the HF $J = 1 - 0$ line will
yield a more versatile probe of depletion that is usable in a wide variety
of environments.  By observing HF, we will be able to probe the gas phase abundance
of fluorine across the Galaxy and to detect diffuse molecular clouds of very small column density. 

Herschel observations of Galactic HF in the local Universe will provide an important benchmark for future observations at high redshift that will become possible with ALMA.
HF absorption in quasar spectra is a potential probe of molecular gas at high-z, providing a possible bridge between the UV/optical observations capable of probing H$_2$ in low column density systems and the radio/millimeter-wavelength observations that probe intervening molecular clouds of high extinction and large molecular fraction; at redshifts beyond $\sim 0.3$, HF is potentially detectable from ground-based submillimeter observatories in several atmospheric transmission windows. 


%Data on halogen hydrides are extremely scarce with two
%ground based detections of HCl (Schilke et al. 1995, Salez et
%al. 1996),
%JST add:
%one detection from the KAO (Zmuidzinas et al.\ 1995)  
%and an ISO detection of HF (Neufeld et al. 1997). 
%HIFI offers the opportunity of detecting
%HF and HCl over a wide range of physical conditions.

%DN Very minor changes in wording and a few typos correction
%
\subsubsection{Carbon chemistry}

The chemistry of the carbon hydrides has been studied extensively
by means of ground-based optical absorption line spectroscopy of 
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 the velocity structure of the 
CN electronic transitions is similar to that of  the ground 
state rotational line of carbon monoxide seen in emission at
115 ~GHz, the CH$^+$ line is broader, and has a different velocity
structure (Crane et al. 1995, Pan et al. 2001). 


\begin{figure}
 \resizebox{17cm}{!}{%
 \rotatebox{270}{%
%   \includegraphics*[6cm,0.1cm][19cm,28cm]{chchp_3} }}
\includegraphics{chchp_3} }}
\caption{\label{chchp} 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}

CH$^+$, like other extremely reactive molecules (such as H$_2^+$,
CO$^+$, OH$^+$, H$_2$O$^+$), 
is particularly interesting as a  diagnostic of 
interstellar processes, since it 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). 
{\bf Observed  CH$^+$ column densities of 10$^{12}$ to 
10$^{14}$ cm$^{-2}$ (Fig.\ \ref{chchp}) are clearly inconsistent
with steady-state low-temperature chemistry}  because the main CH$^+$  
formation  process is the endothermic reaction between C$^+$ and
H$_2$, which is very
 slow at diffuse cloud temperatures of a few $\times 10$~K.
The likely reason for this failure is the still incomplete  description
of the structure and kinematics of the ISM in steady state chemical
models, leading to
an incomplete coverage of the relevant physical 
processes. Steady state models fail for at least two reasons: (1) 
the assumption of uniform physical conditions (density and temperature)
 is clearly not valid for CH$^+$ and other species formed
in endothermic reactions; and (2) the regions where CH$^+$ is formed
must be very small, and probably transient,
because the gas physical conditions along the same lines of sight
(density $n_{\rm H} \sim \rm few \times 100 \rm \, cm^{-3}$, kinetic temperature $\sim \rm few \times 10 \rm \, K$)  
are fairly well determined with less reactive molecules such as
C$_2$, and the observed CH$^+$ column densities 
cannot be produced under those physical conditions.

Magneto-hydrodynamical shocks (i.e.\ C-type shocks; Draine 1980,
Flower \& Pineau de For\^ets 1998) can explain the observed CH$^+$ abudances, but
predict a velocity shift between neutral and ions 
($\sim$ 2 km \, s$^{-1}$ in the  comprehensive models of 
Flower \& Pineau des For\^ets 1998).  An alternative hypothesis is
that  reactive species form in 
regions  where bursts of turbulent dissipation locally heat
the gas to temperatures as large as 10$^3$K (Joulain et al.\ 1998). 
An efficient CH$^+$ formation is then obtained 
within active regions, which become warm enough
that the formation rate becomes significant. Other molecules are
formed in the same environment, notably OH, HCO$^+$ and H$_2$O. 
In order to explain the efficient formation of CH$^+$, while not overpredicting the 
OH abundance, models attempt to take proper account of the interstellar
magnetic field and transient decoupling of ions and neutrals.
 As in the case C-type shocks, ion-neutral drift has the effect of
enhancing the energy available in reactive collisions of C$^+$ and
H$_2$ whilst leaving that in collisions of O and H$_2$ unchanged.
Therefore, the formation rate for CH$^+$ is
increased by a much larger factor than that for OH. 
In the Joulain et al.\ (1998) picture, 
all species would
form with the same centroid velocity, as is observed in most
sightlines.  {\bf Whether C-type shocks, turbulent dissipation or some other heating mechanism is responsible for enhancing the CH$^+$ abundance in diffuse clouds, 
HIFI is ideally suited to detect CH$^+$ in 
this highly excited, small scale, component of the interstellar 
medium, because it will provide the key information on the abundance
dependence with velocity .}

Studying CH$^+$ formation is also important for understanding
carbon chemistry in the diffuse medium, since CH$^+$ is an intermediary in the
formation of many carbon-bearing interstellar molecules.
CH$^+$ reacts with molecular hydrogen to form CH$_2^+$, then
CH$_3^+$, and finally CH$_5^+$.
The recombination of CH$_2^+$ and CH$_3^+$ with electrons
produce CH and CH$_2$, but CH$_3^+$ 
participates in many ion-molecule reactions, leading ultimately other molecules such as
C$_2$, C$_2$H$_2$ and to the  carbon
chains and cyanopolyynes (Terzevia and Herbst 1998).  The carbon hydrides to be observed with HIFI -- 
CH, CH$_2$, and CH$^+$ -- are therefore key species for
understanding the first steps of carbon chemistry. 

%DN Expanded text provided by Eric Herbst
%
\subsubsection{Nitrogen chemistry}

Nitrogen chemistry is unique because of the importance of neutral-neutral reactions.  In order to produce nitrogen hydrides, especially ammonia, in the gas phase, it is first necessary to synthesis molecular nitrogen, N$_2$, which is produced by neutral-neutral reactions involving nitrogen atoms and other diatomics such as OH and NO.  The low temperature rates of these processes are not really known.  Once sufficient N$_2$ builds up, reaction with abundant He$^+$ ions leads to the production of the atomic nitrogen ion N$^+$.  

	The formation of ammonia starts with the weakly endothermic reaction 
N$^+$ + H$_2$ $\longrightarrow$ NH$^+$ + H,
which can be powered at low temperature by either fine-structure excitation of the nitrogen ion or by the ortho spin modification of H$_2$, the lowest rotational state of which ($J=1$) lies 170 K above the ground ($J=0$) rotational state.  The availability of either source of energy is uncertain.  Subsequent hydrogen atom-transfer reactions with molecular hydrogen lead to NH$_2^+$ and NH$_3^+$.  The final reaction in the series, NH$_3^+$ + H$_2$ $\longrightarrow$ NH$_4^+$ + H,
is known to be slow at room temperature and to decrease initially in rate as the temperature is reduced below 300 K.  But, at temperatures near 100 K, the rate begins to increase and lower temperatures lead to an even higher rate.  Since ammonia is formed by dissociative recombination of the NH$_4^+$ ion, its abundance may well be temperature dependent.  

In summary, the nitrogen chemistry through ammonia is still somewhat uncertain and more observational data on nitrogen hydrides are clearly needed for a better understanding.  The data on nitrogen hydrides are extremely scarce
up to now with only a few reported detections (Crawford and Williams
1997, Cernicharo et al. 2000).  Moreover, there is also the possibility that nitrogen hydrides can be produced on the surfaces of grains and desorbed into the gas via non-thermal desorption mechanisms such as photodesorption.  The detection of NH in diffuse clouds is often ascribed to a surface formation mechanism, but this assignment is very uncertain.

%DN New section for FeH and HCl
%
\subsubsection{HCl and FeH}

While the program proposed here focuses primarily on the hydrides of first-row elements, HIFI presents an opportunity to study the heavier hydrides HCl and FeH. 

HCl is expected to be an important major reservoir of gas-phase chlorine, and thus measurements of its abundance serve as a vauable probe of depletion -- as in the case of HF.  Under favorable atmospheric conditions, observations of HCl $J = 1-0$ are possible from ground-based observatories; to date, only two ground-based detections of HCl have been reported (Schilke et al. 1995, Salez et al. 1996).
Severe atmospheric absorption always limits the precision of HCl $J = 1-0$ flux measurements from the ground, and the stability of a space-based heterodyne system will be important in doing quantitative science.

FeH has yet to be detected in the interstellar gas, and its rotational transitions are inaccessible to ground-based telescopes.  The chemistry of interstellar iron has yet to be investigated theoretically, although the recent millimeter-wave detection of FeO (Furaya et al.\ 2003) provides a strong motivation for such investigations.  Bar-Nun et al.\ (1980)
have suggested that FeH would be a likely signpost of gas-grain chemistry, based upon laboratory experiments that suggest that $\rm Fe + H \rightarrow FeH + 1.7\,eV$ proceeds rapidly on nonmetallic grain surfaces.





%DN New section provided by Eric Herbst
%
\subsubsection{Deuterium fractionation and interstellar deuterides}

The detection of deuterium isotopic fractionation in the interstellar medium is of great interest, because the {\bf fractionation is a strong indicator of the physical conditions in any source where it is observed.}   Once the chemistry proceeds to the point where the $\rm HD/H_2$ abundance ratio reaches the elemental value of $\sim 10^{-5}$, ion-molecule reactions produce very large abundance ratios of deuterated to normal species for most heavy interstellar molecules with at least one hydrogen atom.  This effect, known as fractionation, occurs in the gas phase via exchange reactions, the most important of which is 
$\rm H_3^+   +   HD   =$   $\rm H_2D^+   +   H_2$.
Analogous systems exist for the ions CH$_3^+$ and $\rm C_2H_2^+$.  Because of zero-point vibrational energy differences and the Pauli Exclusion Principle, the left-to-right reaction for the $\rm H_3^+ + HD$  system  is exothermic by 230 K.  Thus, in a cloud much colder than this, the backwards reaction becomes very slow. If there were no other destruction pathways for $\rm H_2D^+$, the predicted $\rm H_2D^+/H_3^+$ abundance ratio would exceed unity at 10 K (Roberts, Herbst \& Millar 2003; Walmsley, Flower \& Pineau des Forets 2004).  However, $\rm H_2D^+$ is destroyed by reactions with electrons and heavy species such as CO.

The reactions of $\rm H_2D^+$ with heavy species serve to spread the
deuteration around and cause typical abundance ratios between singly
deuterated isotopomers and normal species (e.g. DCO$^+$/HCO$^+$) in
the range 0.01 -- 0.10.  For regions at particularly high densities
($\rm 10^{6-7}\,cm^{-3}$) and low temperatures, heavy species are
significantly depleted onto grains within rather short periods (1000
yr) and then the $\rm H_2D^+/H_3^+$ abundance ratio can indeed exceed
1.  Moreover, subsequent exothermic reactions with HD produce $\rm
D_2H^+$ -- detected by Vastel et al.\ 2004 with a $\rm
p-D_2H^+/o-H_2D^+$ ratio $\sim 1$ as predicted (Roberts et al. 2003,
Flower, Pineau des For\^ets \& Walmsley 2004) --
and even $\rm D_3^+$.  Reactions between the deuterated ions and heavy species still in the gas lead to multiply-deuterated heavy neutrals such as 
$\rm D_2CO$ and ND$_3$ (Roueff et al.\ 2005). When comparing the degree of
deuterium fractionation of H$_3^+$ isotopologues with models, care
must be taken of the molecule symmetry properties which depend in a
complex way on the H$_2$ ortho/para ratio (Flower et
al. 2004). Unfortunately only one symmetry state of each H$_2$D$^+$ and
D$_2$H$^+$ are accessible from the ground, in windows with rather poor
atmospheric transmission. Given the importance of these species for
the propagation of D fractionation in all interstellar molecules,
 accurate determination of their abundances is a must. 
In addition, dissociative recombination reactions with deuterated ions produce atomic deuterium, which, upon sticking to the surfaces of interstellar grains, is able to deuterate species such as CO into deuterated isotopologues of formaldehyde and methanol.  Such deuterated isotopologues remain in the solid phase unless some event, such as star formation, causes surface temperatures to rise and accelerates evaporation. 
	
Although deuterated species typically possess rotational transitions
at lower frequencies than the normal isotopomers, several important
transitions (ground state lines) 
are found at frequencies uniquely accessible to HIFI.  Some
deuterated molecules of interest for HIFI include HDO, D$_2$O,
NH$_2$D, DF, and $\rm D_2H^+$ (2 lines from both
symmetry states).


%DN Minor changes from here on

\subsection{Proposed observations: an absorption line survey}

\subsubsection{
%JST change: why HIFI
HIFI specific goals %JST
}

Experience with SWAS and ODIN has shown that {\bf absorption line studies remain
the most reliable method for measuring the column densities of interstellar
hydrides.}  Because they are almost always subthermally-populated at
interstellar densities, hydrides show emission line strengths that are
proportional both to the poorly-known gas density and to
the collisional rate coefficients.  Furthermore, theoretical calculations suggest that the latter are often
very different for excitation by ortho- and para-H$_2$, so the hydride
abundances derived from emission line observations depend upon what is
assumed for the H$_2$ ortho-to-para ratio;  radiative transfer effects can
also be problematic because of the large optical depths.  Absorption line observations
largely circumvent these difficulties.

With its broad wavelength coverage, HIFI can access key
hydrides participating in interstellar chemistry with high
sensitivity and high spectral resolution. The list 
includes the carbon hydrides CH, CH$_2$, and CH$^+$, the nitrogen
hydrides NH, NH$_2$, NH$_3$, NH$^+$, and the oxygen hydrides
H$_2$O, H$_3$O$^+$ and OH$^+$, with their most abundant isotopologues
including the deuterated species.  
OH can be observed from the ground, and also
with PACS at moderate spectral resolution. 

{\bf HIFI is ideally suited for such studies, since it gives access 
to the ground state
rotational lines of those hydrides, and -- most importantly --
provides  high (and critically-needed) 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 CH and CH$^+$. As shown in Figure
\ref{chchp}, typical CH and CH$^+$ column densities for diffuse clouds lie
in the range 10$^{12}$ to 10$^{14}$ cm$^{-2}$. Such column
densities will give rise to strong absorption features,
with opacities of order unity for a column density of 10$^{13}$
cm$^{-2}$ and a linewidth of 1 km s$^{-1}$.

Although very few hydride lines can be observed from the ground in the
submillimeter region, most of them will be accessible with SOFIA, except 
for the water and CH$^+$ lines
which are unique to HIFI. However, HIFI is more sensitive than
SOFIA because of the 
%JSt change: superior 
larger %JST
telescope
size, the better receiver sensitivity, and the lack of telluric features, 
which could overlap with the pattern of interstellar absorption features
under investigation.
%JSt drop (does not apply for non-CII obs on SOFIA): 
%Also the advantage of having a multi-pixel cameras on SOFIA can not
%counterbalance the better operating conditions since we do not
%need any spatial information for this absorption survey.


%JST change:{\sl 
\subsubsection { %JST
Selection of lines}

The line list is given in Table \ref{tab:lines-h}. 
Even for absorption lines, the derivation of column densities is
dependent upon the excitation conditions. 
While the assumption of a low excitation temperature is
valid for most diffuse clouds, it can lead to severe errors
in the derived column density if the line is formed in very dense or radiatively-pumped
gas.  To circumvent such errors, we propose to search for absorption
from excited levels for selected species that we expect
to produce the strongest features: H$_2$O, H$_3$O$^+$,
CH$^+$, CH, and NH$_3$.

To illustrate the sensitivity of the proposed observations, we list in 
Table \ref{tab:lines-h} the opacity of the transition for a column
density of 10$^{13}$ cm$^{-2}$ and a line width of 1 km s$^{-1}$ in the absence of 
collisional excitation (i.e.\ with $T_{ex} = 2.73$ K). 
The molecular data have been extracted from the JPL
and the Cologne spectroscopic databases\footnote{In cases where the dipole moment
is not known, these assume a value of 1 Debye.}
and from the literature (CH$_2$ from Michael et al.\ 2003)
Whenever possible, we have tried to optimize the number of lines
observed with a single tuning by combining lines in the upper and
lower side bands.

%DN 
%NH+, ND, FeH, and $\rm D_2H^+$ added.  Still need to add tau for 
%these species and check tau for p-H2O 987 GHz and NH3 1168 GHz
%(I'm skeptical it could be as large as this.)  Darek has suggested
%the excited state NH3 lines at 1214.8 and 1215.2 GHz in lieu of the
%1168 GHz line. Frequencies needed for ND.
%
% Add ND. 3 bands are possible, need to choose the best one based
% on compromise between line strength and possible line confusion
% OK for changing the NH3 lines

\begin{table}
\caption{\label{tab:lines-h} Selection of hydride lines}
\begin{tabular}{lcrccrl}
\\
\hline
\hline
Molecule & & Frequency & $\tau$$^a$ & Priority$^b$ & E$_l$ & Comment\\
& & (GHz) & & & (K) & \\
\hline 
CH$^+$ & 1 -- 0 & 835.071 & 1.17 & 3 & 0 & $\nu$ not accurate\\
CH$^+$ & 2 -- 1 & 1669.16\phantom{0}  & & 3 & 40 & $\nu$ not accurate\\
CH & $^2\Pi_{3/2}\,\,\, 1,2 - ^2\Pi_{1/2}\,\,\, 1,1$  & 536.761 & 0.27 & 2 & 0 & \\
CH & $^2\Pi_{5/2}\,\,\, 2,3 - ^2\Pi_{3/2}\,\,\, 1,2$& 1656.961 & 0.08 & 2 & 26 & \\
CH$_2$ para& $1_{1,0}-1_{0,1}$ & 1910.000 & & 2 & 0 & $\nu$ not accurate   \\
\hline
H$_2$$^{18}$O-ortho &  $1_{1,0} - 1_{0,1}$ & 547.676 & 1.80 & 3 & 0 & \\
H$_2$O-ortho &  $1_{1,0} - 1_{0,1}$ & 556.936  & 1.80 & 3 & 0 & \\
H$_2$O-para & $2_{1,1} - 2_{0,2}$ & 752.033 & & 2  & 89 & \\
H$_2$O-para & $2_{0,2} - 1_{1,1}$ & 987.927 & & 3 & 53 & \\
H$_2$$^{18}$O-para & $1_{1,1} - 0_{0,0}$ &1101.698 & 4.00 & 3 & 0 &\\
H$_2$O-para  & $1_{1,1} - 0_{0,0}$ & 1113.343 & 4.00 & 3 & 0 & \\
H$_2$$^{18}$O-ortho& $2_{1,2} - 1_{0,1}$ &1655.868 & 1.80 & 3 & 0 & \\
H$_2$O-ortho &  $2_{1,2} - 1_{0,1}$ &  1669.905 & 1.78 & 3 & 0 & \\
OH$^+$ & $^3\Sigma^- \,\,\,1,2,5/2 - 0,1,3/2$ & 971.804 & 3.2 & 3 & 0 & \\
H$_3$O$^+$ & $0_{0,1} - 1_{0,0}$ &  984.697 & 0.3 & 3 & 7 & \\
H$_3$O$^+$&  $1_{1,1} - 1_{1,0}$   &      1655.814 & 1.10 &  3 & 0 & \\
HDO & $1_{1,1}-0_{0,0}$ & 893.639 & 3.54 & 1 & 0 & DFG$^c$ \\
D$_2$O & $1_{1,1}-0_{0,0}$& 607.350 & 4.03 & 1 & 0 & DFG   \\
D$_2$O  &   $2_{1,1}-1_{0,1}$  & 897.947 & 1.89 & 1  & 0 &  DFG \\
\hline
NH$^+$ & $^2\Pi_{1/2}\, J =3/2 - 1/2$ & 1012.5, 1019.2  & & 3 & 0 & Complex h.f.s.\\
NH & $^3\Sigma^- \,\,\, 1,1/2 - 0,1/2$ & 974.479 & 0.63 & 3 & 0\\
ND & &  491.934 &  & 1 & 0 & Complex h.f.s.\\
ND & &   522.077  &  & 1 & 0 & Complex h.f.s.\\
ND & & 546.128 & &   1 & 0 & Complex h.f.s.\\
NH$_2$ & &      952.578 & 0.57 & 1 &   0 & Complex h.f.s.\\
NH$_3$ & $1_0 - 0_0$ & 572.498 & 2.46 & 2 & 0 & \\
NH$_3$ & $2_0 - 1_0$ & 1214.859 &  & 2 & 23 & \\
NH$_3$ & $2_1 - 1_1$ & 1215.245 &  & 2 & 28 & \\
NH$_2$D & & 494.454 & 0.62 & 1  & 0 & DFG \\  
\hline
%DN omit H2S 
%H$_2$S-ortho & & 736.034 & 0.49 & 1 & 0 & \\
FeH & $^4\Delta_{7/2} \,J = 9/2 - 7/2$ & 1411.2\phantom{00} & & 3 & 0 & $\Lambda$ doublet\\
H$^{37}$Cl & $1 - 0$ & 624.978 & 0.72 & 1 & 0 & DFG \\ 
H$^{35}$Cl & $1 - 0$ & 625.919 & 0.72 & 1 & 0 & DFG \\
HF & $1 - 0$ & 1232.476 & 3.91 & 3 & 0 & \\
DF & $1 - 0$ & 651.099 & 3.88 & 1  & 0 & DFG \\
\hline
$\rm D_2H^+$-ortho & $1_{11} - 0_{00}$ & 1476.606 & & 2 & 0 & \\
$\rm D_2H^+$-para  & $1_{10} - 1_{01}$ & 691.660  & & 2 & 0 & DFG\\
\hline 
\hline 
\end{tabular}

$^a$ Optical depth for a linewidth of 1 km s$^{-1}$ and 
column density of 10$^{13}$ cm$^{-2}$ \\
$^b$ Priority runs from 1 (lowest) to 3 (highest)\\
$^c$ DFG Line detectable from the ground under excellent atmospheric conditions \\
(but severe absorption limits the precision of flux measurements)  \\


\end{table}



%JST change: {\sl 
\subsubsection { %JST
Selection of sources}

To maximize the sensitivity to small column densities of
absorbing material, we propose to search for hydrides absorption lines
using bright submillimeter continuum sources as background sources.
Our observations will also probe emission from hydrides intrinsic
target sources themselves.
The selection criteria for the sources are
the availability of a strong far infrared continuum and the presence
of interstellar clouds along the line of sight, as evidenced
by existing data (on HI, H$_2$CO, CS, H$_2$O, OH and other species.)
The sample includes  targets studied by ISO 
(W49N, Vastel et al.\ 2000; Caux et al.\ 2002; W33A; and AFGL7009S, Dartois et al.\ 2000), by 
SWAS (Sgr B2, Neufeld et al.\ 1997; W51, Neufeld et al.\ 2002; W49, Plume et al.\ 2003; G34.3+0.1; G10.47+0.03), and by ODIN (Sgr A, Sandqvist et al.\ 2003); as well as HII regions studied at centimeter, millimeter and submillimeter wavelengths
(Fish et al.\ 2003, Mueller et al.\ 2002). In addition,
we selected a few specific sightlines that are bright both at optical
and far-IR wavelengths to allow a meaningful comparison with previous
work; these are  Cyg OB2 - 12 (McCall et al.\ 2002), GCS-3 (Moneti et
al.\ 2001, Goto et al.\ 2002), and IRAS05389-6908 in the LMC (Tumlinson et al.\ 2002).
  
The selected background sources are listed in Table \ref{tab:abs-sources}.
We expect at least one absorbing cloud along each line-of-sight.
Since the target HII regions are located at different galactocentric
distances, a sample with $\sim 18$ sources thus enables us to 
cover a wide variety of physical conditions in the interstellar
medium -- from the most diffuse gas to dense clouds. 

%DN Added G10.47+0.03 (SWAS source) 
%   and W49N-HD, IRAS 15502, IRAS 17455, W3, DR21 (Charlotte Vastel's suggestions)
%   (==> ~ 6 sources to be removed later to keep total the same?)
%
%   SWAS position entered for G34.3+0.1
%
%   Several fluxes/distances entered with info provided by Charlotte
%
%   Charlotte has different positions for W49N, G005.88-0.39 (W28A), Sgr B2, W33A
%   (not entered pending further discussion)
%
%   TO DO:
%
%   Need to coordinate with Bergin/Lis on Sgr B2
%   Need to consider effect of different beam sizes for flux estimates
%   Need to fine tune optimum positions
%   Need to add positions Tom has up his sleeve
%   Need to prioritize and then remove positions
%

\begin{table}[ht!]
\caption{\label{tab:abs-sources} Absorption line sources}
\begin{tabular}{lllrrc}
\\
\hline
\hline
Name & R.A. & Dec & F$_{100\mu\rm m}$  & F$_{350\mu\rm m}$ & D  \\ 
& (J2000) & (J2000) & (Jy) & (Jy) & (kpc) \\
\hline
W3     & 02 25 43.92 & \phantom{--}62 06 01.5 & 42100 & 30 & 2.4 \\
W3(OH) & 02 27 04.7 & \phantom{--}61 52 26 & 10600 & 400 & 2.4 \\
IRAS 15502--5302 & 15 54 06 & --53 11 36 & 12300 & & 8.5 \\
NGC6334--I & 17 20 53.35 & --35 47 01.5 & 19800 & 1130 & 2.0 \\
Sgr A$^*$ & 17 45 39.95 & --29 00 28.2& 23300& & 8.6 \\
GCS--3 & 17 46 14.9 &  --28 49 37 &  410 & & 8.6\\
%Sgr B2 & 17 47 20.4 & --28 23 07 & 66600 & 3962 & 8.6     \\
IRAS 17455--2800 & 17 48 41.6 &  --28 01 38.3 &  5390 & & 8.6\\    
G005.88--0.39 (W28A) & 18 00 30.4 & --24 04 00 & 26800 & 950 & 3.8 \\
G008.67--0.36 & 18 06 18.9 & --21 37 35 & 5220 & 650 & 4.8 \\
G10.47+0.03  & 18 08 38 & --19 51 50 & 10160 & 533 & \\
G10.62--0.39 (W31C) & 18 10 28.7 & --19 55 49 & 21400 & 1110 & 4.8 \\
W33A & 18 14 39.9 & --17 51 59 & 6310 & 350 & 4.0 \\
RAFGL7009S & 18 34 19.7 & --05 59 44 & 2136 & 200 & 2.0 \\
W43S & 18 46 03.5 & --02 39 26 & 11700 & 360 & 7.4 \\
G34.3+0.1 & 18 53 19 & \phantom{--}01 14 46 & 32500 & 1125 & \\
W49N & 19 10 13.2 & \phantom{--}09 06 12 & 36200 & & 11.5 \\
W49N--HD & 19 10 17.5 & \phantom{--}09 05 59.5 & 66200 & & 11.5 \\
W51 & 19 23 43.90 & \phantom{--}14 30 34.8 & 26800 & 280 & 5.5 \\
Cyg OB2--12 & 20 32 41.1 & \phantom{--}41 14 28 & 131 &  &       \\
DR21(OH) & 20 39 00.9 & \phantom{--}42 22 38 &   & 800 & 1.0 \\
DR21 & 20 39 02.04 & \phantom{--}42 19 30.6 & 22200 &  & 1.0 \\
NGC 7538--IRS1 & 23 13 45.38 & \phantom{--}61 28 10.7 & 14101 & 150 & 2.8 \\
IRAS 05389--6908 (LMC) & 05 38 36.7 & --69 06 29 & 2490 & & 50   \\ 

\hline
\hline
\\
\end{tabular}

Fluxes from IRAS point source catalog, Sandell (1994), van der Tak et
al. (2000), Goldsmith et al. (1990), Mueller et al. (2002), Sandell (2000).

\end{table}




%JST change: {\sl 
\subsubsection { %JST
Observing mode and time needed }

From published ground based and space data, we expect continuum
intensities of hundreds of Jy in the HIFI domain. 
 Table \ref{tab:abs-sources} lists fluxes at 350 $\mu$m. With a
telescope sensitivity of 450 Jy/K, these fluxes correspond to
continuum level (SSB) from 0.24 to 2.7 K (9 K for Galactic Center sources).
We will observe
in double beam switching mode. 
%JST comment: I do not understand this observing mode; but I guess it is too
%premature to argue about it %JST
%We plan to  spent 1 minute with fast chopping
%to detect the continuum level, and then turn to a slower chopping rate
%to detect the narrow features. 
With a uniform velocity resolution of
0.3 kms$^{-1}$, we can detect the weakest sources in the list 
(100 Jy at 350 $\mu$m) at the 10 $\sigma$ level in 15 minutes
observing time. 
%JST add:
A signal-to-noise ratio of 10 on the continuum level gives a 3 $\sigma$ detection for absorpion with $\tau=1$. %JST
We will
achieve a higher sensitivity towards the strongest sources, allowing us
to probe considerably smaller optical depths. The
sensitivity is rather uniform in bands 1 through 5, because the rise in the system 
temperature with frequency is typically compensated by the larger continuum
fluxes.  For band 6, we will devote 30 minutes to the weakest sources
and observe a restricted number of lines.  An observing time of 15 minutes gives an
excellent sensitivity (10 $\sigma$ or more) for the strong sources.
The average observing time spent is therefore 15 minutes per tuning for strong
sources and 22 minutes for the weak ones.
%DN
%The number of tunings will have to be refined to account for the additional
%species added.  We need to redo some of the number here and bring everything 
%into accord with the time available
We will observe the full set of lines (21 tunings) in 6 sources and
a restricted set of lines (10 tunings for the strongest
water, CH$^+$, CH, NH and H$_3$O$^+$) in the remaining 12 sources.
%DN This added by Tom, who has identified a very promising source
% 
In addition, we will devote 15 hours to a spectral line survey of the 
source TBD.

The time required for this investigation of hydrides
is {\bf $21 \times 6 \times 0.25$ = 31.5 hours} for
the strongest sources, {{\bf $10 \times 12 \times 0.375$ = 45}
hours} for the weaker sources, and {\bf 15 hours} for TBD, summing to a total of {\bf 92 hours} of HIFI time.\\




\subsubsection{Complementary data with PACS and SPIRE}   

In order to widen the wavelength coverage, and provide access to additional
interstellar species -- albeit with a lower spectral resolution --
we propose to perform a spectral scan of the brightest sources with
PACS. We will be able to detect ground state lines of OH(119$\mu$m,
possibly the 79 and 53 $\mu$m)
 CH(150$\mu$m) and maybe ortho-CH$_2$ (127 $\mu$m), 
fine structure lines of oxygen and ionized
carbon, and many  spectral features from water, carbon clusters
and other species.  We require {\bf 2.5 hours} per source for
this spectral scan. 
%Finally, the spatial structure of the continuum sources must be known
%for a good understanding of the absorption data. Since the sources are
%bright, they can be observed rapidly. We
%have chosen an integration time of 5 minutes
%for both PACS and SPIRE.  
For the highest priority  sources,
the required time for  low resolution spectroscopy is 
{\bf $6 \times 2.5 = $  15 hours}.
\newpage



\section{Molecular complexity} 

 Large molecules made of carbon are the best candidates to  
account for the unidentified spectral features in the spectra
of bright stars (The DIBs) and as carriers of the 
emission bands in the infrared.
They can be found in different forms such as 
carbon chains, rings, fullerenes or polycylic aromatic hydrocarbons  
(PAHs). The latter are the best candidates to account for the so-called 
``unidentified''or aromatic infrared bands observed in emission between 
3 and 17 $\mu$m in many places of our Galaxy as well as in external 
galaxies. Large carbon molecules play an important role in  
interstellar chemistry, building a bridge between small gas-phase  
species and carbonaceous grains. Following the major advances made by
ISO and Spitzer, the opening of the full far-infrared and 
submillimeter wavelength domains by Herschel will provide a unique 
opportunity to understand this bridge. Indeed many of the large 
carbon molecules have rovibrational transitions in this spectral domain 
which are very specific of the identity of their carbon skeleton. 
Full spectral line surveys of template sources 
(star forming regions, hot cores, 
carbon stars, planetary nebulae) could reveal some of them but  
likely not many because of the moderate sensitivity at particular
frequencies. Large molecules, 
even abundant, have a large partition function which make them more 
difficult to detect. The systematic exploration must therefore be 
completed by deeper searches at more specific frequencies to uncover 
weaker features which would otherwise escape detection in line surveys. 
Since Herschel will provide this opportunity to detect large  
carbon molecules, it is fundamental to devote some time to this topic. 
This requires a good strategy and some preparatory work both  
experimental and theoretical. We provide here the status and 
major guidelines for this topic. 
 
 
\subsection {Carbon chemistry} 
 
Most interstellar molecules, especially the largest, 
have a carbon skeleton. The identification of these species has been 
done mainly in the radio domain (cm, mm and submm) and is thus heavily 
biased towards polar molecules. With the increased sensitivity of 
visible and infrared detectors, a few symmetric molecules including  
carbon clusters have been securely identified in the interstellar 
medium (C$_2$, C$_3$, possibly C$_4$ and maybe C$_5$ or C$_6$, 
Thorburn et al. 2003,  
Oka et al. 2003, Cernicharo et al. 2002, Giesen et al. 2001,
Goicoechea et al. 2004). Thorburn et 
al. (2003) have shown that the 
strength of some DIBs is correlated with the presence of carbon 
clusters, supporting the hypothesis proposed by Herbig (1995)  
that their carriers are indeed  
reactive organic molecules (carbon clusters, chains or 
cycles). Further support for the presence of still unknown carbon 
chains is provided by the  
excellent correlation between CCH, c-C$_3$H$_2$ and C$_4$H both in the 
diffuse ISM and in PDRs (Liszt \& Lucas 2000, Teyssier et al. 2004,
Pety et al. 2005) and their significant abundances. 
 
By combining abundances of hydrocarbons detected in various 
interstellar environments, Gerin  et al. (2003) and Teyssier et 
al. (2004) have proposed a first sketch of the carbon budget in 
the interstellar medium (see Table \ref{tab:carbon}). This analysis 
shows that, while the inventory of species with two and three C-atoms 
seems to be fairly complete, our knowledge of heavier  
species with more carbon atoms is still fragmentary. It is very likely 
that many molecules with four or more carbon atoms remain to be
indentified in various 
regions of the interstellar medium.  
 
 \begin{figure}[h!] 
 \begin{center} 
\resizebox{12cm}{!}{% 
%\rotatebox{270}{ 
%\includegraphics*[14.5,8.1cm][20.5cm,25.5cm]{fig-new-hifi} 
\includegraphics{fig-new-hifi} 
} 
 \end{center} 
\caption{\label{gerin:fig5}  Prediction of the abundance of 
hydrocarbons having one to 
eight carbon atoms obtained with a PDR model. The gas density is 
$2 \times 10^{4}$ cm$^{-3}$ and the radiation field is 100 times 
more intense than the average interstellar radiation field in the
solar neighborhood. Abundances 
relative to H$_2$ are given at a cloud depth of $\sim 2$ mag. We 
used two chemical rate networks : the New Standard Model (NSM) (left) 
and the UMIST95 rate file (right). 
(E. Roueff, E. Herbst private communication).  }   
\end{figure} 
 
\begin{table} 
\caption{\label{tab:carbon}  Carbon Budget in the Interstellar Medium} 

\begin{tabular}{ll|cc|c} 
\hline 
C atoms & Molecules & Abundance / H$_2$ &  
Fraction of  & Abundance / H$_2$ \\ 
& & PDR and/or & total C & Dark clouds \\ 
& & diffuse ISM & & \\
\hline 
\hline 
1 & C$^+$, C, CO & $2.6 \times 10^{-4}$ & 1 & $8.0 \times 10^{-5}$ \\ 
2 & C$_2$, C$_2$H, ... & $8 \times 10^{-8}$ &  $6.1 \times 10^{-4}$ &  
$1 \times 10^{-8}$\\ 
3 & C$_3$, C$_3$H$_2$, HC$_3$N, CH$_3$CCH,... & $5 \times 10^{-9}$ &   
$5.8 \times 10^{-5}$ & $1.2 \times 10^{-8}$ \\ 
4 & C$_4$, C$_4$H,... & $2 \times 10^{-9}$ &   
$3.1 \times 10^{-5}$ & $1.7 \times 10^{-8}$ \\ 
5 &  C$_5$H, HC$_5$N... &  &  & $1.2 \times 10^{-9}$ \\ 
6 & C$_6$H, ... & $1.1 \times 10^{-10}$ &   
$2.5 \times 10^{-6}$ & $4.2 \times 10^{-10}$ \\ 
7 & HC$_7$N, ... &  &   & $5.5 \times 10^{-10}$ \\ 
8 & C$_8$H, ... &  & & $1.1 \times 10^{-11}$ \\ 
\hline 
$\sim 50$ & PAHs & & 0.1 & \\ 
\hline 
\end{tabular}  
 
We use a gas-phase carbon abundance of $1.3 \times 10^{-4}$ . 
\end{table} 
 
% These numbers should be revised using the most recent ISM C/H by
% Sofia et al. 1.6e-4
%O/H from Cartledge et al. 2004 ApJ 613, 1037, Jensen et al. 2005 ApJ
%619, 891
%and Ueda et al. 2005 ApJ 620 274
%

 
 
We have used gas-phase chemical models to 
obtain predictions for the abundances of some large 
molecules.  Though 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.  
We show in Figure \ref{gerin:fig5} predictions from 
a PDR model (Le Bourlot et al. 1993, Le Petit et al. 2001) obtained with two 
 different chemical networks, the last revision of the "New Standard Model "  
from E. Herbst (Terzevia et al. (1998, 2000), and priv. com. ) and the UMIST95 
network (Le Teuff et al. 1999). More neutral-neutral reactions 
participating in the synthesis of carbon chains are included in 
the  NSM network compared to the UMIST95 rate file. Though hundreds of 
chemical species and thousands of reactions are taken into 
account, these rate files do not consider isomers. This may affect 
the predictions for some species since the reactivity 
of isomers could be different, particularly between cyclic and 
linear isomers (see discussion in Turner, Terzevia and Herbst 2000,
and Foss\'e etal. 2001). 
However, at this stage and given the scarce information both 
on the observational and on the theoretical side, no isomers are
included. 
 
Reported abundances of selected hydrocarbons in either the diffuse ISM
 (Liszt and Lucas
2000) or PDRs (Teyssier et al. 2004, Pety et al. 2005) are even higher than
model predictions. Therefore, carbon clusters and other  hydrocarbons 
 have a significant abundance in photon dominated regions. 
Models predict 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. Even though actual models
fail in reproducing the observed abundances, their prediction is
useful as a first order approximation. As
C$_n$ species cannot 
be detected in the radio domain because they do not have a permanent dipole 
moment, they are  first priority targets  with HIFI which can
access their ro-vibrational lines. 
 
In fact,  model predictions of carbon chains and clusters
abundances should  be viewed 
as lower limits of the actual abundances. In all sources they studied, 
Teyssier et al. (2004) and Pety et al. (2005) report 
larger abundances for both  
C$_3$H$_2$ and C$_4$H, by at least one order of 
magnitude,  than model predictions. To explain this discrepancy,  
they propose an additional production mechanism not relying on gas phase chemistry. The tight spatial association of both CCH and C$_3$H$_2$ with the 
mid-IR emission features (likely produced by PAHs), suggests that 
the photothermodissociation (Verstraete et al. 2001, Allain et al.1996, 
Le Page et al. 2003) of some PAHs releases small carbon clusters and 
radicals in the gas phase (C$_2$, C$_3$,..), which can then 
participate in the carbon chemistry. The smooth spatial evolution of the
profiles of the Aromatic Infrared Bands (AIBs) -- attributed to PAHs --
across PDRs analysed by Rapacioli et al. (2004) is a further
indication of the evolution of grain properties in regions illuminated
by FUV radaiation.
 
While more laboratory work is needed to quantify the production 
rate of carbon clusters from PAHs, order of magnitude calculations 
show that this is a valid hypothesis. The sole detection of carbon 
chains with a large abundance gives good prospect for future work. 
Also, the necessity for a better understanding of the role of 
PAHs in the interstellar medium in terms of gas heating, charge 
balance, and now chemistry is further strengthened. 
 
 
 
 \subsection{Search for new carbon chains} 
 
It is therefore expected that some large molecules, bridging the  
gap between the known (small) gas phase species and the Polycyclic Aromatic  
Hydrocarbons (PAHs), survive in the diffuse ISM. 
Though no definitive identification 
has been made yet, large organic molecules including carbon chains and 
rings (Thaddeus et al. 1995, Oka et al. 2003) are good 
candidates for some Diffuse Interstellar Bands. Given typical values 
for oscillator strengths,  reported equivalent widths of these absorption 
features  require quite high abundances for their carriers 
(Herbig 1995, 2000).  
 
Herschel, and especially HIFI, is perfectly fitted for bringing new 
information on carbon chains and cycles, and related reactive species, 
since it gives access with unprecedented sensitivity to the frequencies 
where low lying vibrational modes can be detected.  Hence deep 
observations  in the far infrared and the submillimeter will allow the 
detection of  symmetric species which escape detection in the radio 
domain. After the first tentative detection by Van Orden et al. (1995) 
using KAO, the C$_3$ radical has been clearly detected recently  towards 
the Galactic Centre by ISO (Cernicharo et al. 2000) and the KAO (Giesen 
et al. 2001). Cernicharo et al. (2002) report a tentative detection of 
C$_4$ towards a variety of sources while Goicoechea et al. (2004)
report the detection of a new band in NGC~7027, possibly due to C$_5$
or C$_6$. So the primary targets for 
investigation are the pure linear carbon clusters.  
 
%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 can be combined with the line surveys, 
%but additional deep  observations will certainly be necessary for 
%achieving the best sensitivity. 
 
\subsubsection{Estimated Line Intensities} 
 
Except for C$_3$, the lowest bending modes of carbon chain molecules 
have not  been measured in the laboratory yet. For the time being, we 
base our estimate of the expected line intensities on a combination of 
IR-laboratory data and ab-initio calculations. With more THz-laboratory 
data becoming available, the estimates given below will be updated and 
become more precise, allowing to arrive at a more precise final 
observing plan for HIFI.  
 
Our analysis includes not only pure carbon chains, but also carbon 
chains with hydrogen, oxygen and nitrogen atoms attached to the end. For 
each species, we take the present best estimate for the transition 
frequencies between the ground state and first excited level of the lowest 
energy vibration mode together with the induced dipole moment for this 
transition. Following standard assumptions (thermalized excitation in 
the rotational levels of the ground state), we have calculated 
the rotational quantum number for the strongest (i.e.\ highest opacity) 
transition, and the column density at which this transition reaches an 
optical depth of unity, for a set of tempratures ranging from cold to 
moderately hot gas.  
 
Tables \ref{tab-pure-carbon}-\ref{tab-oxy-carbon} present the best 
estimate transition frequencies, the dipole moments and the column 
densities of these species so calculated for which the optical depth is 
equal to unity. A linewidth of $\Delta v = 10$~km~s$^{-1}$ (so that 
$\frac{\Delta\nu}{\nu}$=$\frac{\Delta v}{c}$=$\frac{1}{30000}$) is 
assumed for these calculations %and the column density for which $\tau$=1 
%is calculated using: 
 
%\begin{equation} 
%N_{low}(\tau=1)~=~\left(\frac{A}{8\pi c \nu^3}\frac{\nu}{\Delta\nu}\left[ 
%1-exp(-E_{low}/kT)\right]\right)^{-1} 
%\end{equation} 
 
 
%where $A$ is Einstein's coefficient, $E_{low}$ is the energy of the 
%lower level and $T$ is the temperature in K. 
 
\begin{table}[h!] 
\begin{center} 
\caption{\label{tab-pure-carbon}Pure Carbon Chains} 
\begin{tabular}{lcccccc} 
\hline 
\hline 
Molecule&C$_3$&C$_9$&C$_{10}$&C$_{11}$&C$_{12}$&C$_{13}$\\ 
\hline 
Frequency [THz]& 1.901 & 1.469 & 1.274 & 0.974 & 0.854 & 0.672 \\ 
Dipole Moment [D] &0.43&0.24&0.24&0.25&0.25&0.26\\ 
N$_{l}$($\tau$=1 at 10~K)[cm$^{-2}$] &1.40E+15&3.88E+15&3.77E+15&3.42E+15&3.38E+15&3.12E+15\\ 
N$_{l}$($\tau$=1 at 50~K)[cm$^{-2}$]&1.40E+15&3.56E+15&3.57E+15&3.34E+15&3.34E+15&3.04E+15\\ 
N$_{l}$($\tau$=1 at 100~K) [cm$^{-2}$]&1.20E+15&3.59E+15&3.55E+15&3.28E+15&3.30E+15&3.02E+15\\ 
N$_{l}$($\tau$=1 at 500~K)[cm$^{-2}$] 
&1.15E+15&3.55E+15&3.56E+15&3.29E+15&3.28E+15&3.02E+15\\ 
\hline 
\hline 
\end{tabular} 
\end{center} 
\end{table} 
 
\begin{table}[h!] 
\begin{center} 
\caption{\label{tab-nitro-carbon}Chains of the type NC$_n$N} 
\begin{tabular}{lcccccc} 
\hline 
\hline 
Molecule & NC$_6$N & NC$_8$N & NC$_{10}$N & NC$_{12}$N & NC$_{14}$N & 
NC$_{16}$N \\ 
\hline 
Frequency [THz]& 1.874 & 1.205 & 0.836 & 0.615 & 0.474 & 0.375 \\ 
Dipole Moment [D] &0.18&0.18&0.18&0.19&0.19&0.2\\ 
N$_{l}$($\tau$=1 at 10~K)[cm$^{-2}$] & 
6.47E+15&6.56E+15&6.45E+15&5.87E+15&5.79E+15&5.24E+15\\ 
N$_{l}$($\tau$=1 at 50~K)[cm$^{-2}$] & 
6.36E+15&6.29E+15&6.32E+15&5.70E+15&5.72E+15&5.16E+15\\ 
N$_{l}$($\tau$=1 at 100~K)[cm$^{-2}$] 
&6.51E+15&6.39E+15&6.32E+15&5.68E+15&5.65E+15&5.14E+15\\ 
N$_{l}$($\tau$=1 at 500~K)[cm$^{-2}$] 
&6.35E+15&6.36E+15&6.32E+15&5.67E+15&5.65E+15&5.10E+15\\ 
\hline 
\hline 
\end{tabular} 
\end{center} 
\end{table} 
 
\begin{table}[h!] 
\begin{center} 
\caption{\label{tab-hydro-carbon}Chains of the type HC$_n$H} 
\begin{tabular}{lccccc} 
\hline 
\hline 
Molecule & HC$_8$H & HC$_{10}$H & HC$_{12}$H & HC$_{14}$H & HC$_{16}$H \\ 
\hline 
Frequency [THz]& 1.865 & 1.202 & 0.839 & 0.621 & 0.477 \\ 
Dipole Moment [D]&0.12&0.12&0.12&0.12&0.12\\ 
N$_{l}$($\tau$=1 at 10~K)[cm$^{-2}$] &1.51E+16&1.46E+16&1.43E+16&1.45E+16&1.43E+16\\ 
N$_{l}$($\tau$=1 at 50~K)[cm$^{-2}$]&1.48E+16&1.45E+16&1.44E+16&1.44E+16&1.44E+16\\ 
N$_{l}$($\tau$=1 at 100~K)[cm$^{-2}$] &1.47E+16&1.42E+16&1.43E+16&1.43E+16&1.43E+16\\ 
N$_{l}$($\tau$=1 at 500~K)[cm$^{-2}$]&1.41E+16&1.43E+16&1.42E+16&1.42E+16&1.42E+16\\ 
\hline 
\hline 
\end{tabular} 
\end{center} 
\end{table} 
 
\begin{table}[ht!] 
\begin{center} 
\caption{\label{tab-oxy-carbon}Chains of the type OC$_n$O} 
\begin{tabular}{lcccc} 
\hline 
\hline 
Molecule & OC$_3$O & OC$_4$O & OC$_5$O & OC$_6$O\\ 
\hline 
Frequency [THz]& 0.545 & 2.938 & 1.439 & 1.739 \\ 
Dipole Moment [D]&0.06&0.06&0.06&0.06\\ 
N$_{l}$($\tau$=1 at 10~K)[cm$^{-2}$] &6.63E+16&6.62E+16&5.85E+16&6.17E+16\\ 
N$_{l}$($\tau$=1 at 50~K)[cm$^{-2}$] &5.96E+16&5.85E+16&5.92E+16&5.76E+16\\ 
N$_{l}$($\tau$=1 at 100~K)[cm$^{-2}$] &5.72E+16&5.70E+16&5.81E+16&5.77E+16\\ 
N$_{l}$($\tau$=1 at 500~K)[cm$^{-2}$] &5.68E+16&5.76E+16&5.71E+16&5.68E+16\\ 
\hline 
\hline 
\end{tabular} 
\end{center} 
\end{table} 
 
Tables \ref{tab-pure-carbon}-\ref{tab-oxy-carbon} shows that the different species reach unit optical depth at 
column densities of around $10^{15}$-10$^{16}$~cm$^{-2}$. An additional 
constraint results from the excitation of the levels: the 
ro-vibrational transitions have very high critical densities 
%(GIVE NUMBER), 
so that the population will be confined to the vibrational 
ground state.  The rotational levels will be excited close to thermal 
population because of the lack of purely rotational transitions for 
these non-polar molecules.  The most promising scenario for detecting 
carbon chains through the ro-vibrational transitions of their low 
bending vibrations is thus in absorption against a bright  
background continuum source, to be selected among  the target list
for hydrides (see Table \ref{tab:abs-sources}).
 
\subsubsection{Observations} 
 
For 2 km/s wide velocity bins, Herschel/HIFI reaches an rms of about 
24~mK  at 1.5 THz (Band 6) in a total observing time of 3600 secs and 
the same rms is achieved in 300 secs at lower frequencies (Bands 1-5). 
With typical continuum brightnesses of the background and/or embedded 
sources of a few K to %JSt change: 100 K 
about 10 K %JSt
(see table \ref{tab:abs-sources}) and aiming at a S/N of $10$, 
absorption lines with an optical depth of a few percents
 are thus detectable.  In comparison to the total 
H$_2$-column densities of the absorbing layers, ranging from a few 
$10^{21}$ up to a few $10^{24}$ cm$^{-2}$, relative abundances for the 
carbon chain species of $5 \times 10^{-8}$ in the weaker background, low 
column density sources, and down to well below $10^{-10}$ for the 
strongest, high column density sources would be required to clearly 
detect the absorption within 1 hr total observing time, and 
correspondingly lower for longer integrations. 
 
This makes HIFI/Herschel uniquely sensitive for the detection of such 
carbon chain molecules. The high spectral resolution allows to resolve 
out the absorption line. PACS will be used to fully cover the far
infrared wavelength range towards a few sources. The same PACS observations
will be used for analysing both the narrow lines from
interstellar hydrides and the more complex rotational pattern from 
C clusters and C chains.

% where the other, lower spectral resolution 
%instruments such as PACS would see a smaller line/continuum contrast, 
%outweighting the slightly higher sensitivity per spectral and spatial 
%resolution element; their spatial multiplexing is in this case no 
%advantage, as the absorption is confined to the small and bright 
%continuum background source. 
 The frequency range covered by Herschel 
from space is unique; only SOFIA may compete in some parts of the 
spectrum without residual atmospheric opacity from stratospheric 
altitude, although the larger telescope of Herschel will result in a 
stronger signal from the continuum background source. 
 
We propose to use about {\bf 30 hrs}  of HIFI core science observations 
{\bf to perform deep, 1 to a few hrs long, integrations at selected 
frequencies} (yet to be determined, following the progress of 
THz-laboratory spectroscopy until Herschel launch) for the carbon chains 
listed in Tables \ref{tab-pure-carbon}-\ref{tab-oxy-carbon} and in  
sources with bright continuum background selected from the list in Table 
\ref{tab:abs-sources}. These sources cover a variety of astrophysical 
environments, ranging from hot cores immediately around embedded, young 
high mass stars, to the cold envelope not affected yet by the
birth of the massive star, to unrelated foreground layers exposed
to FUV radiation. We propose to observe few C$_3$ ro-vibrational lines 
towards the same sources which will provide useful information on the
excitation degree of the pure C chains. As C$_3$ can be detected in
the visible and in the far infrared, we shall use this species for
combining the information obtained from visible and far infrared
spectroscopy.

 Given the present uncertainties about the 
detailed chemistry producing carbon chains it would be premature to 
constrict the search to specific types of sources.
Note that absorption lines from 
interstellar molecules, and noticeably from
$c$-C$_3$H$_2$ and $l$-C$_3$H$_2$ have been
detected towards the selected sources (Cernicharo et al. 1999).
%MG add 
 C$_3$ is thought to be the most abundant molecule with 3 carbon
atoms in the diffuse ISM (see Table \ref{tab:carbon}).
     By detecting  C$_3$ ( and possibly other carbon clusters)
in the same interstellar clouds responsible
for $c$-C$_3$H$_2$ absorption, the proposed observations will provide
new information on the carbon budget.
%MG
  
 
%JST added:
%remark: we also have to consider, whether fluorescense through optical/IR
%pumping can cause sufficiently bright emission lines in dense, warm PDRs
%etc. %JST


%NEEDS SOME MORE SPECIFIC SELECTION OF SOURCES/LINES? 
 

 
 
% \begin{thebibliography}{}
%\bibitem[]{}Boonman A.M.S. et al., 2003, A\&A 406, 937.
%\bibitem[]{} Michael E.A., Lewen F., Winnewisser G., Ozeki H., Habara
%  H., 2003, ApJ 596, 1356.
%\end{thebibliography}

\include{hifi-ref}

 \end{document}


%\subsubsection{Preparatory work} 
 
%Lab. Work 
 
%Theoretical Calculations 
 
%Prep. Observations (CCH and C3H2 and C4H) 
%in absorption for selecting sources. 
 
 
 
 
\subsection{Search for low-energy bending modes of PAHs} 
 
\subsubsection{PAHs in the ISM} 
 
Polycyclic Aromatic Hydrocarbons (PAHs) have been proposed as 
the carriers of the Aromatic IR  Bands (AIBs) at 3.3, 6.2, 7.7 8.6, 
11.3 and 12.7 $\mu$m since about twenty years (e.g. L\'eger \& Puget 
1984). Interstellar PAHs absorb UV photons and cool down by emitting  
in their vibrational modes (L\'eger et al. 1989, Allamandola et 
al. 1989). Space observations in the thermal infrared  
with the Infrared Space Observatory (ISO) have shown that the features 
are ubiquitous in the Milky Way and in external galaxies and that 
they represent about 10\% of the total carbon   
(see the review by Boulanger et al. 2000).  
For galaxies, the absence of the AIBs in 
their spectrum is a clear sign of the presence of an Active Nucleus (Seyfert 
galaxy or QSO) (Genzel et al. 1998), which could have otherwise escaped 
detection -- Active galaxies are fairly rare and usually recognized 
with their visible spectrum but the method does not work in 
heavily reddened sources --  
 
PAHs are important players in the physics and chemistry of the 
interstellar medium. They  participate in the 
charge balance (Bakes and Tielens 1998), in the chemistry  
including a possible participation to the formation of molecular 
hydrogen ( Cassam-Chena\"\i\ et  al. 1994,  Bauschlicher 1998, 
Joblin et al. 2000) and of carbon clusters by photodissociation 
(Joblin 2003, Le Page et al. 2003). But the most 
critical role of PAHs is likely to be their contribution 
to the heating of the gas (Bakes and Tielens 1994, Weingartner 
and Draine 2001). One of the most effective gas heating mechanism 
is the photoelectric effect on solid particles or macro molecules 
(PAHs, dust grains) 
whose efficiency is heavily dependent on the particule size and composition. 
Because of their small size and large combined area, PAHs give 
an important contribution to the heating of the gas. Strong  
observational evidences have been obtained with ISO 
(Helou et al. 2000 at the galaxy scale, Habart et al. 2001, 2002 
for individual PDRs). 
 
Given the multiple roles of PAHs in the physics and chemistry 
of the interstellar medium, it is frustrating that none 
has been directly identified yet. The identification of 
interstellar PAHs has motivated a lot of observational, 
laboratory, quantum chemical and modelling studies. 
Although ISO has released a lot of high quality data, 
the identification of one specific PAH has remained elusive. 
Part of the difficulty is due to the fact that all PAHs 
have aromatic CC and CH bonds 
which IR modes are all very similiar to the AIBs. The analysis 
of the shape and intensity of the AIBs gives  clues on the 
physical processes, but very few on the individual molecules. 
Only limits on the size  and some indication on the stucture 
have been obtained (Hony et al. 2001, Verstraete et al. 2001, 
Pech et al. 2002, Peeters et al. 2002). Also, the identification is 
particularly difficult because interstellar PAHs are thought 
to have larger sizes than the species commonly studied in the 
laboratory (50 carbon atoms versus 24 carbon atoms for coronene, one 
of the largest laboratory PAH). Dedicated experiments are 
needed (and indeed have been built) for studying in the 
laboratory good candidates for interstellar PAHs (Joblin et 
al. 2002a). Spectroscopic work has also been performed, in the hope 
of finding specific spectral features in the visible, and eventually 
spectral lines corresponding to known DIBs. This has remained 
unsuccessful which is not completely suprising considering that 
electronic transitions are very specific and the number 
of species to be considered is tremendous (for a given carbon 
skeleton, various charge states, various hydrogenation states ...).  
 
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 and therefore are poorly 
sensitive to the number of hydrogens at the periphery. Some of these 
bands are located in the THz domain. The possibility of observing 
such a band and resolving its rotational structure 
would help a lot in identifying its carrier. For a symmetric rotor, 
the rotational constants are inversely proportional to N$_C^2$ (N$_C$ 
is the number of carbon atoms in the molecule) 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 size and 
shape of the carbon skeleton. The identification of some of 
these large carbonaceous molecules 
would really help us a lot in our understanding of interstellar chemistry, 
of heating processes in the interstellar medium and of the overall 
life cycle of interstellar matter. 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). 

 
Herschel can have a very unique role in these studies regarding 
the identification of large carbonaceous molecules such 
as PAHs, through the rovibrational structure of their low frequency 
bending modes. This programme requires both I) a large spectral coverage, 
and II) a good sensitivity. No other ground based or space telescope 
can combine these two requirements as efficiently as Herschel. 
 
 
\begin{itemize} 
 
\item{I) The extent of the rovibrational structure depends on the 
rotational constants (i.e. the molecular size) as well as on the 
rotational temperature. 
For instance, this extent is 105 and 26 GHz with a line spacing 
of 667 and 42 MHz at T$_{rot}$=100\,K for C$_{24}$H$_{12}$ and 
C$_{96}$H$_{24}$ respectively.  Note that the line spacing 
assumes no anharmonic effects and no vibration-rotation coupling.  
Whereas the rotational structure depends mostly on the molecular 
size, the position of the band (Q branch) is specific of the 
identity of the molecule.} 
 
\item{II) the low-frequency bending bands have weak 
Einstein coefficients ($A \leq 10^{-4}$ s$^{-1}$ N$_C$$^{-1}$ )  
and the best strategy is to search for  emission lines  in regions 
where the UV/visible radiation field is sufficiently high. Indeed,  
these low frequency  bending modes are involved quite efficiently 
during  the last steps of the cooling of the molecules after  
absorption of  UV photons (Joblin et al. 2002).} 
\end{itemize} 
 
\subsubsection{Proposed observations} 
 
Two observation strategies can then be proposed :\\ 
-- 1 -- \ We have a precise idea on the molecule(s) we want to search for and 
its(their) spectrum(a) is(are) 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!). 
 
Though the first strategy requires a more limited spectral coverage 
and therefore allows a deeper integration, it is clear that, given the 
available information, we  need to perform deep searches with little 
a-priori knowledge. 
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  low 
frequency bending modes are expected to be more frequent and 
more intense in the high frequency range (THz).  
 While SOFIA can also observe in the same wavelength range 
with slightly lower sensitivity, the data will always be contaminated 
by numerous atmospheric features due to minor species, which cannot 
all be predicted accurately. This residual contamination  
will limit severely the possibility of detecting the PAH bands, 
composed of many weak features. 
 
The intensity of these bending modes can be estimated as follows. 
The emission in the mid-IR AIBs  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 UV excitation, as
traced by the IR emission, in the best cases. This leads to a total 
emission of $I(\nu_{f})=9.0 \times 10^{-20}$\,Wcm$^{-2}$ 
for the Orion Bar ($G_{O}=5 \times 10^{4}$). 
For one specific 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 FIR band of this particular molecule. Now there are 
about 400 lines per rotational structure (The rotational temperature 
is about $T_{rot} = 100$K),  and the width of the individual lines 
is assumed to be 3~MHz (1 kms$^{-1}$ at 1 THz). With this simple calculation, 
we predict  an intensity of 
1.46 10$^{-20}$ W cm$^{-2}$ micron$^{-1}$, i.e. 3.1 Jy or 6 mK per line at 
1 THz. 
This crude estimation  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 (see below). 
 
 
We describe the observing strategy  adapted for a deep search 
with some (but not very accurate) a priori knowledge.  
We shall use both PACS and HIFI to optimise the chances of 
identifying some interstellar PAHs.
More calculations have to be done with our model before exact
positions are chosen, but  the 
the best position to search for PAHs will be in PDRs with a high UV flux 
(because the energy emitted in the far-IR is proportional to 
the energy absorbed in the UV). We have however to be careful 
to minimise the contribution of the dust emission (to increase 
the line/continuum ratio)
 
To estimate the required observing time on HIFI, we performed two 
 calculations, assuming  position switching mode.\\ 
-- For the SIS mixers, we assume a typical system temperature 
of 360K at 1 THz and a frequency resolution of 1.0 MHz. Then individual lines 
are detected at the $2 \sigma$ level after 12 hours of 
observations. This noise level is correct since we want to 
detect a pattern with a few hundred lines and not a single line. We 
 shall therefore make use of the redundancy to extract the signal from 
the noise.  
%(CAN THIS BEEN DONE AT A LOWER S/N RATIO ? THIS WOULD 
%SAVE A LOT OF TIME !  ). 
 
-- For band 6, we use a system temperature of 2000 K and 
a frequency resolution of 6 MHz. Then 20 hours are needed for 
a single 2.4 GHz band. 
 
Assuming a set of three frequencies, two in band 5 and one in band 6,
the total  HIFI time is {\bf 64 hours}.
 
We shall use PACS (or SPIRE) to search for the Q branch. This 
feature is predicted to be more intense than individual lines, 
and significantly broader. The best spectrometers 
to detect the Q branches are therefore the medium resolution 
spectrometers. Both spectrometers can be used in a scanning mode. 
 Using PACS time estimator, the spectral range from 100 to 210 $\mu$m
can be scanned down to a flux level of 10 Jy at SNR of 5 in {\bf 5
  hours}. We shall use SPIRE during the same integration
time to search for a Q branch in the 400 - 200 $\mu m$ wavelength
interval.

 
 
As a whole, this program requires {\bf $\sim$ 74 hours} of 
telescope time to reach the required S/N in one source. 
The source will be selected from the list of bright PDRs, 
chosing the position expected to maximise the  
PAH signal, and the efficiency of observation. 
 
 
 
\subsubsection{Preparatory work} 
 
\underline{Laboratory work :}\\ 
Two programmes are currently in progress:\\ 
-(i)- Photodissociation of PAHs and carbon clusters with the 
PIRENEA experiment.\\ 
The PIRENEA set-up has been built to approach in the laboratory the  
physical conditions of the interstellar medium: long isolation time scales and 
very low temperatures. It is devoted to the study of the photophysics 
and chemistry of large molecules and nano-sized particles. 
Recent results show that the photodissociation of PAHs leads to various 
types of carbon clusters which photodissociate by loosing C$_3$, 
C$_2$ and C (Joblin 2003). More experiments are planned to derive 
fragmentation rates to be used in chemical models of photodissociation 
regions. This will enable to quantify the production of small carbon 
clusters by the photodissociation of PAHs via large carbon clusters.\\ 
-(ii)- Spectroscopy of the flopping modes of gas-phase PAHs \\ 
This work is performed by two research teams (M. Vervloet \& O. Pirali at 
the LPPM-CNRS in Orsay and K. Demyk, G. Wlodarczak and L. Constantin 
at the PhLAM- Lille University).  The first team records spectra in  
absorption or in emission of gas phase PAHs using a high resolution 
Fourier transform interferometer (resolution up to 0.002\,cm$^{-1}$). 
This enables to record easily the spectra over a very braod spectral range 
which is necessary to locate the position of the low-frequency bands 
(the flopping modes). Then the 
second team can use this information to record a very high resolution 
spectrum using a gas cell coupled with a THz spectrometer. 
The spectra of a few molecules: pyrene C$_{16}$H$_{10}$, 
phenanthrene C$_{14}$H$_{10}$ and naphthalene C$_{10}$H$_{8}$ have 
been successfully measured with the FT interferometer. Unfortunately, 
none of these species have transitions in the spectral range reachable by the 
THz spectrometer. New attemps are currently made with chrysene 
C$_{18}$H$_{12}$ which has bands in the THz range according to 
quantum chemical calculations. 

 
\underline{Theoretical calculations :}\\ 
Theoretical calculations are also in progress and search for  
new collaborations is still continuing. Several calculations 
are necessary:\\ 
-(i)- ab-initio calculations of the far-IR spectra of PAHs 
to provide guidelines for the experiments.\\ 
-(ii)- calculation of the ro-vibrational structure (position and 
intensity of the rotational lines) \\  
-(iii)- calculation of the intensity of the band emitted in astrophysical 
environments. A model has been done to calculate this intensity 
as a function of the UV flux (Joblin et al. 2002b). It however 
requires as input some molecular data in particular the list 
of vibrational modes and the Einstein coefficients for the 
active IR-FIR modes. This information is available for a 
limited number of species.  Other important molecular data is 
the anharmonicity of the flopping modes which allows to quantify 
the contribution from the hot bands in the band profile. 
This anharmonicity has never been measured nor calculated. 


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%Observations : \\ 
%-- The question is : how to choose the best position for the 
%deep integrations ? The beam will be 10 to 20'', 
% similar to IRAM, larger than CAM. 
 
%ISOCAM images + radio observations + near-IR images (H2) ? 
%Are mm/submm images in the continuum needed ? This may provide 
%a way to compute the continuum level in the submm. 
%D. Teyssier has images of the horsehead nebula at 1200,850, 450 
%and 350 microns . Can use this source as a template ? 
 
%Source models ?  
%\\ 
%\\ 
%CJ: More calculations have to be done with our model. Probably 
%the best position is in PDRs with a high UV flux 
%(because the energy emitted in the far-IR is proportional to 
%the energy absorbed in the UV). We have however to be careful 
%to minimise the contribution of the dust emission (to increase 
%the line/continuum ratio).  
 
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