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%\begin{document}
%\begin{center}
%{\Large The interstellar medium with HIFI}

%\vspace*{0.5cm}
%{\sl Draft 16 May, 2003}

%\end{center}

%{\small E. Falgarone , F. Boulanger} 

%\section*{Proposed outline}
%\begin{itemize}
%\item Introduction, document presentation
%\item Diffuse ISM
%        \begin{itemize}
%\item Characteristics and open questions
%\item State of the art and previous work
%\item The HIFI niche for diffuse ISM studies
% \begin{itemize}
%\item main coolants blocked by atmosphere (CII, NII lines)
%\item spectral resolution 
%\item spatial resolution
%\item sensitivity to extended sources
%\item large frequency coverage
%\end{itemize}        
%\item proposed observations
%\begin{itemize}
%\item  separation of the components (DIG, WNM, CNM) 
%\item  several hydrides/lines in absorption
% \end{itemize}
%        \end{itemize}
%\item Dense and warm ISM: PDRs and shocks
%\item Molecular complexity

%\end{itemize}



\section{The diffuse ISM}

The interstellar medium (ISM) comprises neutral and ionized
components. The neutral component is distributed  into two  thermally
stable phases, the
cold neutral medium (CNM) and the warm neutral medium (WNM) with a non
negligible fraction in evolution between the two stable phases (Heiles
and Troland 2003a,b). 
The ionized components are the hot ionized medium (HIM) at temperatures
larger than 10$^6$ K, 
%JST drop: and %JST
the diffuse ionized gas (DIG, electron densities
$n_e \sim 2$ cm$^{-3}$) and the more diluted warm ionized
medium (WIM, $n_e \sim 0.2$ cm$^{-3}$ ) likely not fully 
ionized, both at temperatures close to 10$^4$ K.

Due to their specific excitation conditions and ionization state 
the various components are traced by different lines. The DIG
is best studied through its optical emission lines, in particular 
the H$\alpha$ emission. 
%JST change: As does the WIM, 
Both the WIM and %JST
the DIG 
%JST change: contributes 
contribute %JST
to the 
[CII] line emission,  and the [NII] fine structure line emission 
is one of 
%JST change: its 
the %JST 
specific %JST change: tracer, 
tracers of the DIG, %JST
the ionization potential of 
nitrogen being larger than 13.6 eV.

The neutral phases (WNM and CNM) are traced by the HI
21cm line, both  contributing to the line in emission, the
cold phase contributing heavily to the 
%JST add: 
(self-) %JST
absorption. 
The contribution of the CNM to the [CII] line emission of the diffuse
interstellar medium is important as [CII] is the main cooling line
of the CNM (Wolfire et al. 2003).
The densest regions are  traced by the 
%JST add:
emission in the %JST
rotational lines of CO and
other heavy molecules %JST drop: in emission %JST
and the fine structure lines 
%JST add:
of %JST
[CI] and  [OI] 
(for the densest and warmest parts).  
 

\subsection{Open questions}

\begin{itemize}
\item even if not fully ionized, 
the powering and ionization of the DIG and WIM %JST change: is 
are %JST
not completely
understood. The ionized component extends as a thick gaseous disk, 
in our Galaxy and
other spirals, over several kiloparsecs. Its  
ionization/heating %JST change: sources 
source %JST
is primarily the thin layer of massive
stars. This requires a large porosity of the neutral medium to
ionizing photons (bubble, chimneys, ...) or the existence of alternative
ionizing/heating sources. 

\item thermal pressure determinations in the diffuse ISM from a survey
of the atomic carbon
fine structure lines (Jenkins \& Tripp 2001) reveals the existence of
ubiquitous wisps of high pressure material arising either from small
scale density enhancements or from pockets
of warm gas, both possibly generated by turbulence. The observed
excursions of thermal pressure are large ($\sim$ 2 orders of magnitude)

\item %JST add:
the %JST
warm components are presumably more smoothly distributed than cold ones, 
but [CaII] absorption
lines against binary stars (Meyer and Blades 1996; Watson and Meyer
1996) suggest that small scale structure also exists in the WIM.
What is at the origin of the small scale structure of the WIM ?

\item the emission line profiles of HI are broad (several
km/s). The CNM contribution is known from  absorption line
profiles when available, and is rather narrow. 
 Both turbulence and thermal broadening 
contribute to the width of the HI emission lines, but
there is no current way of disentangling these two contributions when
no absorption line is available, except by assuming the temperature 
of each phase, an assumption which can lead to serious errors.

\item what is the origin of the large molecular abundances observed in
the CNM, with narrow linewidths? What are the energetic processes 
able to trigger chemical reactions with activation
barriers or endothermicities of several thousands of Kelvins 
in this cold medium?

\item which components are out of thermal equilibrium, ionization, pressure
equilibrium? what is the mass fraction in each of them?

 \end{itemize}

   
\subsection{Previous works and state of the art %JST add:
of the field %JST
}

The diffuse ISM in galaxies has a major contribution to their global
emission in the main cooling lines such as [CII] or [NII]. It is one
of the most poorly known step in the evolutionary cycle of
interstellar matter from stellar interiors to dense protostellar
clouds and has therefore bearings to the evolution of galaxies.
In addition, because of its 
multiphase 
nature, the diffuse ISM
has interesting hydrodynamic properties which must be taken into
account in all calculations and simulations of interstellar gas.
For these reasons, the diffuse ISM warrants further studies 
and HIFI is particularly well suited for them, as we show below. 
 

The structure and physics of the diffuse ISM has also important bearings to 
(i) understanding the gaseous structure of one of the many screens between
the Cosmic Background (CBG) and the Sun, and therefore that of the dust
distribution which affects the continuum measurements of the CBG, 
(ii) %JST add:
investigating %JST 
the evolution of the Galaxy, since the WNM and CNM contains each 
a significant fraction of the mass of the ISM in our Galaxy and their
relative fraction is related to the star formation rate,
(iii) %JST add:
understanding %JST
the evolution of distant galaxies, since  the same species are now
probed at large redshifts, 
(iv) %JST add:
explainging %JST
the large observed porosity of the neutral medium to UV
photons and (v) understanding the first steps of interstellar chemistry. 

Over the years dedicated efforts have been conducted and significant
%JST change: progresses have been performed 
progress has been achieved %JST
in the understanding of the ISM, in
particular with the space projects HST/GHRS, COBE/FIRAS, ISO/LWS and ISOCAM
instruments, the BICE balloon,  and more recently through UV
absorption studies with FUSE. The Arecibo millenium survey
(Heiles and Troland 2003a,b) has provided a comprehensive data set
on the atomic gas through the Galaxy, while deep surveys of the ISM of
distant galaxies are now performed on 8m class telescopes
(e.g. Ledoux et al. (2002) for H$_2$ on VLT, Wolfe et al. for CII on
Keck (2003)). Kanegar \& Chengalur (2003) studied the  21cm HI line
 in Dampled Lyman Alpha (DLA) absorbers and concluded that the balance
 between CNM and WNM evolves strongly with redshift with more warm gas
in high redshift DLAs. 

This observational effort  has
been backed with theoretical work. Wolfire et al. (1995, 2003)
present comprehensive models of the heating and cooling
processes of the diffuse ISM, predicting  the properties of the stable
equilibrium phases.  Parravano et al.
(2003) discuss how the radiation from massive stars 
may modify the balance between the CNM and WNM, linking the
neutral gas properties to the global star formation rate in 
the Milky Way. The conversion from
WNM to CNM  is indeed a critical step for forming stars in the ISM.
Recent  calculations show that the condensation of CNM
clouds can be triggered by a mild compression (Hennebelle and
P\'erault 1999, 2000, Figure
\ref{fig:sim-ph}; S\`anchez-Salcedo et al. 2002). 
These 1-D 
%JST replace: work 
calculations %JST
are now being improved and
 more extensive simulations in  2D and 3D are run (Hennebelle \& 
Audit priv. com.) for studying the development of CNM structures
inside the WNM. 


\begin{figure}
\resizebox{14cm}{!}{%
  \includegraphics*[0.5cm,0.5cm][20.5cm,28.5cm]{ursa_hi.ps}
}
\caption{\label{fig:mamd-hi} Channel maps for the HI 21cm emission 
in a high latitude cirrus cloud (Miville-Desch\^enes et al. 2002). 
The spatial resolution is 1 arcmin; these data have been obtained
with the DRAO interferometer and corrected for short spacings.
Coherent structures down to the spatial and spectral resolution are
seen as long filaments. }
\end{figure}

\begin{figure}
\resizebox{14cm}{!}{%
  \includegraphics*[0.5cm,0.5cm][20.5cm,28.5cm]{mod-hi-hp.ps}}
\caption{\label{fig:sim-ph} Monodimensional MHD simulation of a 
multiphase medium. The
gas is in the WNM initially.  Different epochs
are shown when CNM structures  have developed following a mild compression.
For each frame, the top spectrum 
shows the predicted HI emission line, the middle one
the HI absorption line and the bottom one the [CII] emission line.
Notice the narrowing of the line between HI and [CII]. The CNM
structures appears as sharp peaks in [CII], with corresponding
HI absorption features. However, they are not so easily identified in 
emission because of the significant contribution from the WNM.
 The horizontal axis is the velocity in km/s, the vertical axis the antenna
temperature in K. (Hennebelle \& P\'erault 2000) }
\end{figure}


Despite the wealth of theoretical and observational
effort, a global description of the structure of the ISM 
is missing and the overall structure of the ISM remains elusive, 
in particular the volume filling factor of the
hot, warm and cold phases, the
amount of small scale structure present in each of them, the density
contrasts and perhaps most
importantly the thermal pressure of each of these phases, and its
variation with the distance $z$ to the galactic plane.

The relation  between the various phases 
%JST moved up form below: 
 is likely to be related to their formation
mechanism and 
%JST change: has 
needs to be established. %JST %JST
%--
%JST add:
This includes aspects such as: %JST
how are %JST change; they 
the various phases %JST
embedded?  Are they in thermal pressure balance? 
Is the hierarchy of structures observed in the coldest component (that of
molecular clouds) a fossil of thermal/gravitational
instabilities developed initially in a placental warm medium at 10$^4$K?
%  --
The interrelation between the ISM phases  is
still missing,  because the available information, especially in the
FIR domain of the cooling lines, is still too
scarce. While the 
%JST add: 
line integrated %JST
 [CII] intensity -- but not its line profile -- has
been measured over a wide range of environments (with the KAO, ISO, BICE,
COBE...), information on [NII], which traces the
diffuse ionized gas, is available from the large scale COBE/FIRAS survey
(Bennett et al. 1994), and for one single nearby 
%JST change: galaxy  
galaxy, M82 %JST
(KAO data on M~82 by Petuchowski et al. 1994). 
The far infrared surveys have been performed with small
telescopes %jst chagne:, 
and %JST
thus provide information on large scale only,
while  ground based data on HI and CO tell us that small scale
structure is a generic property of the diffuse ISM, which impacts
all diagnostic methods on the physical conditions (Fig. \ref{fig:mamd-hi}). 
%JST change: Finally, most 
Allmost all %JST
of the previous FIR data lack spectral resolution
and %JST add:
thus %JST
do not give any information on the line profile, while the
rich information in the extensive HI and CO surveys has been fully
used in the past years (Fig. \ref{fig:mamd-hi}, Heiles and Troland 2003a,b
for HI, Bensch et al. (2001) for CO)
%JST add:
and has been shown to be crucial for a proper interpretation of the
observations. %JST 
Furthermore, it is seldom the case that
both high spectral and spatial resolution information on the same
species are available, except for CO and HI.
 For instance, the high spectral resolution HST data on atomic carbon
allow to measure accurately the gas pressure, but only in a limited
sample of lines of sight towards massive stars (Jenkins and Tripp
2001). Ground based observations of atomic carbon are limited
to the denser gas in molecular clouds and do not have the
sensitivity to detect the faint emission expected from diffuse gas : the
sensitivity of the large scale surveys of the atomic carbon fine
structure line by the Mt Fuji telescope is typically 0.4 K (Kamegai et
al. 2003) while emission has been detected at the 0.1 K level in
the diffuse cloud towards $\zeta$ Oph with the KAO (Keene et
al. 1987).



%- ISO-LWS: CII/OI line ratios difficult to interpret because of possible
%self-absoprtion of CII lines (LWS on LMC, SMC)

%- HST: NaI/CaII line ratios 
%High spectral resolution observations show that b(NaI)$«$ b(CaII)
%although m(Ca)=40 and m(Na)=23 and components. At same velocity,
%therefore  presumably coexisting

%- the SIII line seems to carry a lot of energy even in ELD ($I=24$eV)
%Signature of energetic phenomena?

%- The uniformity of the ratio SII/H$\alpha$  suggests that the bulk of
%ionization cannot be provided by HII regions 

%- H$\alpha$: scaleheight of DIG $>$ 2 kpc in NGC891 
%and 1 kpc in our Galaxy (Rand (pas retrouve ses papiers); Haffner with WHAM) 

%- existence of molecular gas at high z detected in CO (Thaddeus)
%1/3 of the HI thick disk

%- HST/GHRS  Fitzpatrick and Spitzer (1997)  Inextricably mixed ionized
%and neutral species. 

%- little HeII in the WIM therefore ionized by stars later than O8
%but these provide just 22\% of ionizing photons
%These stars could expliain the diffuse ionization if 60--75\% of the photons 
%escape the immediate vicinity of the ionizing  stars and travel
%through the disk and halo

%- COBE/FIRAS
 
\subsection{The HIFI %JST drop: niche for 
diffuse ISM studies}

The capabilities of HIFI will allow new breakthroughs thanks to its
high angular resolution, its high frequency resolution (velocity
being the only means accessible to identify different gaseous components on
a same line of sight) and its sensitivity. 
%JST drop:
%Some aspects of the foreseen
%investigation possibilities are described in the following. %JST
Within
the core science program, we focus on the most important contributions
expected through HIFI.


As described above, the degree of intrication of the
stable phases, and the fraction of matter currently in transition 
between the WNM and CNM are two ``hot questions'' for the diffuse ISM.
HIFI offers tremendous capabilities for both points :

The comparison of high spectral resolution 
 [CII] and HI emission spectra will  permit
the investigation of  the degree of association between the so-called cold 
(T$<$ a few 100 K)
and warm   (T$\sim$ 8000 K) neutral phases of the interstellar
medium and the role of thermal instability in the formation of small scale
density structure within the diffuse ISM (Hennebelle \& P\'erault 2000). The
warm neutral gas does not contribute to the [CII] emission while both the 
cold and warm neutral phases contribute to the HI emission (Fig. 
\ref{fig:sim-ph}). 

As [NII] is a tracer of ionised gas only, the [NII] observations
will provide for the first time a
clear picture of the diffuse  ionised 
gas, and will
 allow subtraction of the fraction of [CII] from 
diffuse ionized gas.
 Together with   H$\alpha$ and HI emission spectra, 
the thermal and non-thermal contributions to the observed line widths
of [CII] and [NII] can be determined. 
The narrowest velocity components in HI clouds have velocity widths of 
$\sim 3$ kms$^{-1}$ (FWHM) which are close to the expected thermal width for 
gas at a few  100 K. It is thus difficult to disentangle the
thermal and turbulent contributions to the line profiles, and only
upper limits to the temperature can be obtained (Heiles and Troland
2003a,b). By comparing [CII] and HI emission, we expect to see a
narrowing of the lines due to the higher mass of carbon and properly
separate the thermal and turbulent contributions. 


With its exceptional brightness
sensitivity, HIFI will extend the present observations of the 
[CI] fine structure lines to
 HI clouds, thus providing an independent measurement of the
thermal pressure which can be combined with the
column density and  velocity structure available from HI, 
and with cooling radiation from [CII].
Thermal pressure determinations as derived from [CI] fine structure
lines and molecular line observations vary from tracer to tracer.
The thermal pressure derived from [CI] absorption lines in the diffuse
ISM is P$_{th}/k\sim 3 \times 10^3$ K cm$^{-3}$ (Jenkins \& Shaya 1979,
Jenkins \& Tripp 2001) 
with large excursions above this value, while  
observations of
molecules in absorption against extragalactic sources provide 
P$_{th}/k\sim 2\times 10^4$ K cm$^{-3}$ (Lucas \& Liszt 1994). 
If large pressure differences exist, shocks propagating into the 
coldest phases should be present. Small scale density structures 
are also prominent in the spectra of nearby lines of sights
(Pan et al. 2001, Andrews et al. 2001). 

With the gas temperature and pressure estimates deduced from
the neutral carbon lines, it might be possible to
derive the carbon abundance from the [CII] line intensity. We will be able to 
correlate variations in the carbon depletion with local physical conditions
and dust extinction and emission properties including observations of
the dust emission spectrum with PACS and SPIRE. Such measurements
will open a completely new perspective on the cycling of carbon between
dust grains, small aromatic hydrocarbon particules, CO and atomic carbon.
Indeed, Miville-Desch\^enes et al. (2002) show evidences
for correlation between the abundance of the very small dust particles 
responsible for the mid-IR emission, and the gas velocity structure.


\subsection{Proposed observations}

We summarize below the main tracers of the different phases in the 
diffuse ISM, and how they are used :

\begin{itemize}
\item {H$\alpha$ and [NII] are tracers of the ionized gas exclusively, 
the diffuse  ionized gas and HII
regions along the line of sight } 
\item {[CII] is a tracer of the  CNM and also the DIG} 
\item {[CI] is a tracer of dense structures
  within atomic clouds molecular (some of which are also detected in CO)} 
\item {HI in emission is a tracer of CNM and WNM}
\end{itemize}

%\subsubsection{Goals}

%- separate the CNM and WIM contributions to the [CII] line. 
%To do this, we propose to deduce the DIG contribution (intensity and velocity
%profile) from the [NII] line using a model for the WIM providing the
%[NII]/[CII] line ratio

% from a comparison of the HI and [CII] emission of the CNM (narrow
%components only) separate the
%thermal and turbulent contributions to the linewidths (the turbulent
%contribution being the same, the thermal contribution being inversely
%proportional to $\sqrt{m}$). 

%- if [CI] is detected in the CNM, test the ionization stage of carbon
%and using both fine structure lines, determine the
% thermal pressure measurements with two CI lines.

\subsubsection{Targets}

%\subsubsection{CII and NII line emission}

We have selected two kinds of targets for our purpose. For studying
the global properties of the diffuse ISM in the Galaxy, we have
selected a few lines of sight through the Galactic Plane minimising
the contribution from HII regions. We propose to study the variation
of the [CII] and [NII] spectra as a function of Galactic latitude
for these lines of sight to determine the scale height of the various
phases probed by the HIFI observations. Typical column densities and
corresponding line
intensities  are listed in Table \ref{tab:diff-ism}.

While the observations along Galactic lines of sight will provide
information on  global scale, HIFI can also probe the small scale
structure thanks to its high spatial resolution. We propose to
observe the same spectral lines in nearby high latitude clouds (HLC), where
both the CNM and WNM are thought to coexist from previous
HI observations (Miville-Desch\^enes et al. 2002).  The expected
intensity in a high latiture HI cloud  with a total column density of 
N(HI) = 10$^{21}$ cm$^{-2}$ are listed in Table \ref{tab:diff-ism}.



  \begin{table*}[h!]
    \begin{center}
      \caption{\label{tab:diff-ism} Predictions of line intensities in
      the diffuse ISM}
      \begin{tabular}{ccccc}
        \hline
         & unit  & HLC & HLC & Galactic LOS ($b=0.5^{\circ}$)\\
         & & DIG & CNM & CNM+DIG\\
          \hline
N(HI) & cm$^{-2}$ & -- & 10$^{21}$ & $ 5 \times 10^{22}$\\
C/H &           & -- & $2 \times 10^{-6}$ & -- \\
N(CI)$^a$ & cm$^{-2}$  & -- &$2 \times 10^{15}$ & -- \\
$\Delta v$$^d$ & km/s & -- & 3 & 30\\
I(CII)$^b$ & erg s$^{-1}$ cm$^{-2}$ sr$^{-1}$ &$4 \times 10^{-7}$ & 
$2.4 \times 10^{-6}$ & $4 \times 10^{-4}$\\
I(NII)$^{b,c}$ &erg s$^{-1}$ cm$^{-2}$ sr$^{-1}$& $2.4 \times 10^{-8}$
         &  & $2.4 \times 10^{-5}$ \\
W(CI) 1-0 & K km s$^{-1}$ & -- & 0.1 & 6 \\
W(CI) 2-1 & K km s$^{-1}$ & -- & 0.01  & 0.6  \\
W(CII)& K km s$^{-1}$  & 0.06 & 0.35 & 58 \\ 
W(NII)& K km s$^{-1}$ & 0.007 & -- & 7 \\
T(CI)$^e$ 1-0  & K &  -- & 0.04 & $\sim$ 0.1 \\
T(CI)$^e$ 2-1  & K &  -- & 0.004 & $\sim$ 0.01 \\
T(CII)$^e$ & K& $<$ 0.02 & 0.14 & $\sim$ 2\\
T(NII)$^e$ & K & -- & -- & 0.23 \\
        \hline
      \end{tabular}
 \end{center}
$^a$ Jenkins \& Tripp (2001)
$^b$ Timmermann et al. (1998), Reynolds (1992), Bennett et al. (1994),
      Gry et al. (1992)
$^c$ I(NII) = 0.06 I(CII) \\
$\Delta v$ is the total line width. We expect a combination of broad
      and narrow features for the  Galactic lines of sight.\\
$^e$ T is the peak intensity of an individual velocity component \\
  \end{table*}



%\subsubsection{Carbon fine structure lines}

We propose to observe the same regions as [CII] and [NII] in the
 fine structure lines of atomic carbon. Because the
%JST add:
2-1 line, originating from %JST
excited %JST add:
levels, %JST
is very faint for diffuse gas, a factor of ten weaker than the ground
state %JST add:
1-0 %JST
line ,  we propose to observe a single position
in this line, choosen so as to maximise the signal.


\subsubsection{ Observing modes and integration times}

  \begin{table*}[h!]
    \begin{center}
      \caption{ \label{tab:int-diff} Integration times for diffuse ISM
      positions }
      \begin{tabular}{ccccc}
        \hline
         &  & HLC & HLC & Galactic LOS ($b=0.5$deg)\\
         & & WIM & CNM & CNM+DIG+WIM\\
          \hline
Resolution & km/s & -- & 0.5 & 0.5 (CII), 5(NII)\\
rms [CI] 1-0 & mK & -- & 5 & 10 \\
rms [CI] 2-1 & mK & -- & 2 & 2 \\
rms [CII] & mK& -- & 15 & 200 \\
rms [NII] & mK& -- & -- & 20\\
Time/pos [CI] 1-0 & hr  & -- & 1  & 0.5\\
Time/pos [CI] 2-1 & hr & -- & 9 hr & 9 hr\\
Time/pos [CII] & hr & -- & 3 & 0.05 \\
Time/pos [NII] & hr & -- & -- & 0.4 \\    
        \hline
      \end{tabular}
    \end{center}
  \end{table*}


We will adapt the observing mode to the expected line profiles.
We propose to observe the high latitude clouds  in Frequency Switching 
mode since we expect narrow line profiles (a few kms$^{-1}$ or 25 MHz
at most for [CII]). The profiles for Galactic lines of sight
are expected to be broader and cover
a wider velocity range, with narrow features from the CNM 
superimposed on a broad pedestal of a few tens of km/s %JST change: wide.
width. %JST
Because the Galactic emission is extended, position switching mode
would be extremely time consuming. Therefore, we propose to use the
cold load as reference. The base line can then be measured in an additional 
position, selected as having a low HI column density.  
Table \ref{tab:int-diff} gives the target noise level (one tenth of the
expected brightness temperature Tab. \ref{tab:diff-ism}), and
spectral resolution for each line, and the corresponding observing
time.

In order to get information on the spatial structure, we propose
to map with PACS a field of $4' \times 4'$ for each HLC position,
 which fits nicely the 
ISOCAM field of view, and one single IRAS pixel. We shall then select 
$\sim 5$ positions in each field for deep spectroscopic observations.

 Using the current time estimator, we need $\sim 15$ hours to map [CII] at the
$5\sigma$ level with PACS for one field. We propose to observe  three fields,
in high latitude clouds, selected as having different grain
properties, and also with and without CO emission. 
This amounts to  {\bf 45 hours} for the PACS observations. 
To observe 5 positions in each field, [CII], [CI](1-0) and
[CI](2-1) we need  {\bf 30 hours} in addition per field, so 
we require {\bf 90 hours} for the HIFI observations of high latitude clouds.
 

For the Galactic lines of sight, we shall sample 2 directions, one towards
the molecular ring, one in the outer galaxy, and for each direction 20
latitudes. The total is therefore 
{\bf 20 ( 0.05 +  0.4 +  0.5 + 9) = 200 hours}

\begin{table}[!h]
\caption{\label{tab:diff-pos} Position for diffuse ISM studies}
\begin{tabular}{lccl}
 Name & RA(2000) & DEC (2000) & Comment \\
\hline
Ursa Major &   9:45:00 & 70:25:00 & Cirrus, 3 $4' \times 4'$ fields \\
Inner Galaxy & 18:39:16.3 & -05:42:58 & 20 positions with
different $b$ starting at $(l,b) =  (26.86 , 0.09)$\\
Outer Galaxy & & & 20 positions with
different $b$ \\
\hline
\end{tabular}

Positions are still preliminary
\end{table}




\subsection{Diffuse interstellar clouds}

The chemical properties of diffuse clouds such as those lying along
the line of sight to the bright O9 star $\zeta$ Oph are well known
from a protracted effort with visible and UV spectroscopy during the
past 40 years. Using the same technique  in the radio
domain, with distant radio sources as background objects, Liszt and
Lucas (1996 to 2001) have measured the abundance of  molecules,
not easily studied in the UV/visible wavelength range, but there is little
overlap so far between the lines of sight studied in the UV/visible or in
the radio domains. While absorption studies in the UV and visible 
are extremely powerful
 in detecting ionic, atomic or molecular species in the various ISM phases,
 they are limited in spectral resolution -- the best spectral
 resolution for HST data is
 about 3 kms$^{-1}$ -- and provide no information
 on the  spatial structure.  Nevertheless,  absorption studies
provide  unique information
on the abundances and ionization state of numerous IS ``components''
i.e. defined by a common absorption velocity and Doppler broadening parameter.
Typically, some of the components are
neutral and clearly associated with either the CNM or the WNM, and
some are at least
partially ionized, the origin of the ionization being not always clear.
Some of these components in the lines of sight to halo stars should
correspond to the extended H$\alpha$ emitting layer in the Galaxy but its
identification is not straightforward, especially because of the lack of
spatial information of UV observations and the lack of spectral resolution
of H$\alpha$ images.

To progress further, information on the spatial structure is critical.  
Previous attempts of comparing radio and UV data have used
ground based telescopes for
CO and CN. The outcome can be summarised as follows : despite the
difference in spatial resolution ($\sim 20''$ for emission lines
versus $\le 0.1''$ for absorption), column densities 
derived from radio data are similar to UV column densities 
(see e.g. the work on the diffuse clouds towards $\zeta$ Oph by
 Langer etal. (1987) \& Keene et al. (1987)).  In addition to the spatial
structure  and the line profile, radio data are useful for
constraining the excitation conditions (temperature, density or a 
combination of these parameters). Since HIFI can
observe important species (C, C$^+$, N$^+$, ...) also sampled in the 
UV, we wish to take this opportunity and gather additional information 
on well studied lines of sigh to be able to combine
efficiently  the information acquired on  powerful ground based
and  space telescopes with the future data from HIFI.

To be most effective, we have selected lines of sight with the most 
complete information (H and H$_2$ from FUV telescopes, high spectral 
resolution HST data with STIS or GHRS, and at least CH and CH$^+$ from 
ground based telescopes (preferably also C$_2$ and CN).
The ground state
CO rotational line has been detected in most cases.
% we wish to include lines of sight sampling both the 
%neutral and ionised gas such as PG0804+761. This particular lines of 
%sight crosses an intermediate velocity cloud with a ionisation
%fraction of 20\%, and detected molecular hydrogen. The ground state
%CO rotational line has been detected in most cases.

%JST {\sl 
\subsubsection { %JST
%JST change: Why HIFI
HIFI unique goals %JST
}

  HIFI combines
in a unique way  high spectral and spatial resolution
which can not  be obtained by any other ground based or space telescope. 
For instance, though the Carbon fine structure lines 
 can be observed from the ground, the sole 
detection reported for diffuse clouds  
have been obtained 
by space projects, SWAS (Li et al. 2003) and the KAO (Keene et
al. 1987) because 
of the low line intensity. The superior sensitivity of 
HIFI over SOFIA and ground based telescopes 
will therefore be fully exploited for  these observations. 
 As described above, the combination of 
the fine structure lines of C and C$^+$ with the H and H$_2$ 
column densities (available from FUV data) gives access to the gas 
pressure and carbon abundance for each velocity component. Herschel's  
mapping capabilities can be fully exploited to unravel the small scale  
structure of these interstellar 
clouds
%JST change:, determine 
by determining the %JST
scale length for each cloud separately. 
Information on the line profile is critical for 
performing a meaningful comparison with models since 
the behaviour of  different velocity components, detected in e.g. 
sodium lines, can be different. 

%JST change:{\sl 
\subsubsection { %JST
Proposed Observations} 

Since the sources are extended, observations will be performed in 
frequency switching. We propose to sample the spatial
structure by mapping small fields around each background source. The
exact mapping  strategy must be adapted to each source to gain the
full benefit of
HIFI. Since we have a preliminary list of sources
only, we have chosen the same area, $1 \times 1$ arcmin or 0.03 pc for
nearby clouds (100 pc distance) as a first guess. This represents
36 points for [CII], and 9 for the ground state  carbon line. For the weak
excited carbon line at 809 GHz, we shall observe
the central point only.

A first selection of possible lines of sight is listed below. We 
have chosen  the  
archetype of diffuse interstellar clouds,  $\zeta$Oph, 
 HD~102065 observed with the  HST, FUSE and ESO, and the 
intermediate velocity cloud PG0804+761. 
We propose to sample one or two more positions in the $\zeta$Oph clouds,  
where  CO emission 
is stronger to sample the transition from atomic to molecular gas
(Kopp et al. 1996, Liszt 1997).
 

With a typical extinction of 1 magnitude,  line width of 
0.5 -- 1 kms$^{-1}$ and gas density of $n_H \sim 100 - 1000 cm^{-3}$ as
derived from the FUV and visible data,
we expect C and C$^+$ column densities of 10$^{15}$  and 
10$^{17}$ cm$^{-2}$ respectively, and line
intensities of $\sim$ 0.1 K for the 1-0 [CI] line (Keene et al. 1987)
and $\sim$ 0.5 K for [CII] (Gerin et al. 2001)
The difference between these lines of sight 
and the cirruses studied above is the larger radiation field 
because of the loose association with massive stars, and the 
somewhat larger gas density, which both increase the line  intensity.


The program therefore requires  integration time per
position of $\sim 15$ minutes for [CI](1-0)  (5$\sigma$ detection) 
and $\sim 10$ minutes for [CII]  with a
frequency resolution of 200 kHz and 1 MHz respectively (0.15 kms$^{-1}$).


\begin{table}[!h]
\caption{Observing time for diffuse clouds}
\begin{tabular}{lrrrrrr}
\hline
Line & $\delta \nu$ & $\delta v$ & $\sigma$ & time/position & Positions & total time \\
     & MHz & kms$^{-1}$ & mK & minutes & & hours \\
\hline
C($^3P_1 - ^3P_0$) & 0.2 & 0.12 & 20 & 15 & 9 & 2.0 \\
C($^3P_2 - ^3P_1$) & 0.4 & 0.15 & 10 & 30 & 1 & 0.5 \\
C$^+$ & 1.0 & 0.16 & 100 & 10 & 36 & 6.0 \\
\hline
{\bf Total} & & & & & & {\bf 8.5} \\
\hline
\end{tabular}

\end{table}



\begin{table}[!h]
\caption{Single lines of sight from FUSE and/or IRAM}
\begin{tabular}{lrrl}
\hline
Name & RA(2000) & DEC(2000) & Comment\\
\hline
$\zeta$ Oph & 16 37 09.5378 & -10 34 01.524 & Star + 2 positions \\
%\hline
%B0212+735 & 02 17 30.8134 & +73 49 32.622 & BL lac\\
%B055+508 & 03 59 29.7473 & +50 57 50.162 & QSO \\
%B2013+370 & 20 15 28.7 & +37 10 59 & Radio source\\
%B2037+511 & 20 38 37.12 & +51 19 12.7 & QSO \\
\hline
PG0804+761$^a$ & 08 10 58.55 & +76 02 41.8 & FUSE QSO, radio quiet\\
%HD~73882 & 08 39 09.5344 & -40 25 09.262 & FUSE HD detection \\
%HD~34078 & 05 16 18.1497 & +34 18 44.341 & FUSE warm H$_2$ \\
%HD185418 & 19 38 27.4827 & +17 15 26.052 & FUSE, double system \\
HD102065$^b$ & 11 43 36.9487 & -80 29 00.521 & FUSE, HST\\
\hline
\end{tabular}

$^a$ Richter et al. (2003) $^b$ Rachford et al. (2002), Gry et al. (2002)
\end{table}


For 5 fields towards 3 background sources,
  it sums to a total of HIFI time of {\bf 42.5 hours}. 


%HOW MANY SOURCES CAN WE AFFORD. A SURVEY WOULD 
%BE VERY INTERESTING AS PART OF A KEY PROGRAM ?  
 
%DO WE WANT PACS OR RATHER COMPLETE WITH GREAT ON SOFIA ??? 
 
 

%{\sl Preparatory observations :}

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

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


 
%BETTER CO MAPS CAN BE OBTAINED WITH HERA AT IRAM.
%SHALL WE PROPOSE A NEW PROGRAM ? 






%\end{document}



















