[CII] observations of H molecular layers in transition clouds ††thanks: Herschel is an ESA space observatory with science instruments provided by European-led Principal Investigator consortia and with important participation from NASA.
Key Words.:ISM: molecules — ISM: structure
We present the first results on the diffuse transition clouds observed in [CII] line emission at 158 m (1.9 THz) towards Galactic longitudes near 340 (5 LOSs) & 20 (11 LOSs) as part of the HIFI tests and GOT C+ survey. Out of the total 146 [CII] velocity components detected by profile fitting we identify 53 as diffuse molecular clouds with associated CO emission but without CO emission and characterized by A 5 mag. We estimate the fraction of the [CII] emission in the diffuse HI layer in each cloud and then determine the [CII] emitted from the molecular layers in the cloud. We show that the excess [CII] intensities detected in a few clouds is indicative of a thick H layer around the CO core. The wide range of clouds in our sample with thin to thick H layers suggests that these are at various evolutionary states characterized by the formation of H and CO layers from HI and C, respectively. In about 30 of the clouds the H column densities (“dark gas”) traced by the [CII] is 50 or more than that traced by CO emission. On the average 25 of the total H in these clouds is in an H layer which is not traced by CO. We use the HI, [CII], and CO intensities in each cloud along with simple chemical models to obtain constraints on the FUV fields and cosmic ray ionization rates.
The Herschel key program GOT C+ (Galactic Observations of Terahertz C) is designed to study the diffuse interstellar medium (ISM) by observing with the HIFI instrument the [CII] fine structure line emission and absorption at 1.9 THz (158 m) over a volume weighted sampling of 500 lines of sight (LOSs) throughout the Galactic disk. The GOT C+ project is described by Langer et al. (2010a) and the use of [CII] emission to detect diffuse warm “dark gas” (H molecular gas not seen by CO observations) by Langer et al. (2010b). C is a major ISM coolant, and its 158 m line is an important tracer of the properties of the diffuse atomic and diffuse molecular gas clouds. The [CII] line thus enables us to trace an important but to date poorly-studied stage in cloud evolution - the transition clouds going from atomic to molecular: HI to H and C to CI and CO (Snow & McCall 2006). These clouds have a large molecular hydrogen fraction in which carbon exists primarily as C rather than as CO (Tielens & Hollenbach 1985; van Dishoeck & Black, 1988). Transition clouds are difficult to study using the standard tracers (HI or CO) but [CII] can trace this gas.
There is growing evidence that a substantial amount of interstellar gas exists as molecular H, not traced by CO, for example: from Gamma-ray data from EGRET (e.g. Grenier et al. 2005) and Fermi-LAT (e.g. Abdo et al. 2010); and, the infrared continuum in diffuse clouds (Reach et al. 1994). Goldsmith et al. (2010) detected warm H in emission beyond the CO extent of Taurus. Wolfire et al. (2010) have modeled the molecular cloud surfaces to estimate the amount of “dark gas” in the form of molecular H in the H/C layers and find it contributes about 30 of the total mass in clouds with total A 8 mag. Here, we present direct observational evidence for the H/C layer in a number of transition clouds through the detection of an excess [CII] line emission in them. We use a sample of 53 transition clouds characterized by A 5 mag. and the presence of both HI and CO emissions but no CO. We analyze the observed [CII] intensities combined with HI and CO data to obtain an inventory of the total molecular H in different layers in transition clouds and then constrain the physical conditions by applying simple models for CO formation and photodissociation.
2 Observations and data analysis
The observations reported here were made as part of the HIFI performance verification and priority science phases. We observed the [CII] line at 1900.5469 GHz towards 16 LOSs in the galactic plane with the HIFI (de Graauw et al. 2010) instrument on the Herschel Space Observatory (Pilbratt et al. 2010). The [CII] spectra were obtained using the wide band spectrometer (with 0.22 km s velocity resolution, over 350 km s range) at band 7b and using integration of 800s to 1800s (with rms of 0.1K to 0.2K on data smoothed to 1 km s). For each target we used the Load chop (HPOINT) with a sky reference offset by 2 in latitude. The data were processed in HIPE version 3.0 using the standard pipeline for HIFI. Using a fringe fitting tool within HIPE we were able to mitigate the standing waves in band 7b (Higgins & Kooi, 2009) to sufficiently low levels to provide good baselines in the [CII] spectra (Boogert, private communication). The data presented here are in the Galactic plane at l = 337.8°, 343.04°, 343.91°, 344.78°, 345.65°, 18.3°, 22.6°, 23.5°& 24.3°; out of the plane at b=0.5 for l = 24.3°and b=1 at l=22.6°& 24.3°; at b=-0.5°& -1°at l=18.3°& 23.5°. Table 1 summarizes all the observational data used in our analysis.
|[km s]||[K km s]|
|[CII]||GOT C+||12||1.0||0.1 - 0.2||1,2|
|1.9 THz||Herschel HIFI|
This paper; Langer et al. (2010a); McClure-Griffiths et al. (2005) Stil et al. (2006); Pineda et al. (2010).
An example of the [CII] spectrum is shown in the top panel in Fig. 1. The [CII] intensities were corrected for main beam efficiency ( 0.63). For comparison the HI and the CO spectra are shown in the lower panel. The [CII] spectra show many velocity resolved features. All [CII] emission features show an overall correlation with the HI, though not all HI features show corresponding [CII] emission. Many [CII] features are also correlated with CO features. To separate the individual velocity components we used multiple Gaussian fitting. In the case of complex (overlapping) velocity features we used both [CII] and HI profiles together to identify the individual components. We identified a total of 146 velocity components in all the LOSs. As seen in Fig. 1 as well as in the examples shown in Langer et al. (2010b) and Pineda et al. (2010) in each spectrum we detect many velocity components. However, their identity as clouds is somewhat uncertain as the decomposition itself is not very unique and may not be reliable (e.g. Falgarone et al. 1994). Though we use the HI profile as an independent check on the features, the beam sizes (Table 1) are not modeled into the decomposition. For simplicity, here we refer to them as clouds, but in reality some of them may be for example, isolated turbulent clumps, transient fluctuations of larger structures, or superposition of extremely narrow velocity components. In view of the uncertainties, for all our quantitative analysis we do not use all of the Gaussian fit parameters. Instead we use the fitted V to locate a parcel of the gas at a certain velocity and width. The I(CII), I(HI), I(CO) intensities for each cloud were then obtained, in a consistent manner, by integrating the intensities (T) over the velocity width (V) centered at the respective V (except in a few cases which are confused by the adjacent component).
We identified 58 [CII] components as dense molecular clouds traced by their CO emission (e.g. the red Gaussian fits in Fig. 1) and these are discussed in a separate paper by Pineda et al. (2010). We regard the remaining 88 components without CO counterparts (e.g. the black Gaussian fits in Fig.1) as diffuse clouds, envelopes or transition clouds. We examined these 88 diffuse [CII] clouds by correlating them with the CO spectra. We found that 53 components have associated CO emission while the remaining 35 have no CO counterparts. These 35 clouds are labeled diffuse atomic clouds of which 29 are discussed by Langer et al. (2010b). To place our [CII] cloud samples in the context of the general interstellar clouds, in Fig. 2 we identify them in an A - FUV parameter space. We use our 3- detection limits (Table 1) for [CII], CO, CO, and CO to estimate the corresponding thresholds of A and FUV based on the calculations by Visser et al. (2009). The Visser et al. calculations use T =100 K, and n = 300 cm similar to what we use below in our analysis. Here we present results on 50 transition clouds excluding 3 for data quality and other issues.
3 Results and discussion
3.1 [CII] sample of transition clouds
In Fig. 2 we find that our [CII] sample of transition clouds are diffuse having A 3 - 4 mag. for reasonable interstellar FUV, in the range of 1 - 10 (the average FUV intensity erg cm ssr (Draine 1978)). In Fig. 3 we show a schematic of the diffuse cloud layers. In the dense cores with CO emission, the conversion of C to CO is more complete while it is partial in these diffuse transition clouds due to lack of sufficient self-shielding. All clouds contain some quantity of HI. As seen in Fig. 3, the observed [CII] emission originates from the purely atomic HI layer along with a contribution from the H/C layer, while the CO emission originates in the H/CO core. Thus estimates of H column densities using CO intensity alone entirely misses the H in the H/C (“dark gas”) layer. Therefore, a complete inventory of molecular H in the cloud requires both the [CII] and CO intensities.
3.2 [CII] in the HI/C layer
In Fig. 4a we plot the [CII] intensities against the HI intensities for all 50 transition clouds. The error bars in Fig. 4a ṟepresent the 1- uncertainties in the respective measured intensities. In spite of the large scatter we note a lower bound to I(CII) that increases gradually with I(HI) which is consistent with the [CII] emission expected from a HI/C layer (Fig. 3). For quantitative analysis of [CII] emission from the HI/C layer we use the following steps (see also discussion in Langer et al. 2010b):
i) The observed I(CII) is regarded as the total = + , where and are the emissions originating from the HI/C and H/C layers respectively with no [CII] emission from the CO emitting core (Fig. 3).
ii) Use the HI intensity, I(HI) to estimate the HI column density, N(HI) = I(HI) cm.
iii) Use N(HI) to estimate the C column density in the HI/C layer, X(C)N(HI), where X(C) = n(C)/n(HI) is assumed to be 1.5 10.
iv) Using the above we calculate I(CII) f((n(HI),T) , where the function f accounts for the excitation conditions for density n(HI) and temperature T (see Langer et al. 2010b). Then we can express it in terms I(HI), as f(n(HI), T)I(HI)
We find that on average the form of I(CII) versus I(HI) can be fitted by a straight line obtained for density n(HI) 200 cm at temperature T 100K as shown in Fig. 4a. In more massive clouds with HI intensities greater than 1000 K km s Wolfire et al. (2010) estimate n(HI) 50 -150 cm and T 70 -80 K. However, all our [CII] clouds are less massive with HI intensities 600 K km s. In the present analysis we assume n(HI) 200 cm and T 100K which seem to describe best the contribution to the [CII] intensity from the HI/C layer.
3.3 [CII] in the H/C layer
Having estimated arising from the HI/C layer we can now calculate the [CII] excess arising from the H/C layer as = - . In Fig. 4b we show this excess plotted against . Since all the carbon in the CO emitting region is converted to CO we do not expect to see any correlation. However, in spite of the large scatter we do note a lower bound to the excess that increases gradually with as seen by the straight line fit in Fig. 4b. This suggests the presence of a C layer surrounding the CO emitting core as shown in the schematic in Fig. 3. The straight line in Fig. 4b is an approximate fit to the lower bound to the [CII] intensities as a function of and it corresponds roughly to [CII] intensities for clouds with a H/C layer containing a H column density 15 of the H in the CO core; that is, = 0.15 (see below). Therefore this line may be regarded as representing the “nominal” diffuse CO clouds which contain a small H envelope around the CO core. However, the clouds with large [CII] excess well above this line could represent a sample of clouds in transition with larger H envelopes and relatively smaller CO cores. We can now use this excess and the observed to estimate the H column densities in the H/C and CO layers respectively. For CO we use the phenomenological relationship (c.f. Dame et al. 2001):
In the H/C layers the C excitation is by H molecules and we can use the [CII] excess shown in Fig. 4b to derive the column density as follows:
i) Use the to calculate the C column density, in the H/C layer, as a function of density n(H) and temperature (T). Here we assume a higher density of n(H) 300 cm than in the HI layer and a temperature T 100K.
ii) Use this N(C) column density to estimate N(H)= N(C)/2X(C), where X(C) = 1.5 10. Thus we get H column density as a function of excess I(CII) for the above assumed n(H) and T (see Langer at al. 2010b),
|Cloud||I(C II)||I(HI)||I(CO)||N(H) in C||N(H) in CO||[A (C/CO)]||[FUV ]|
A corresponding to the C/CO layer. External FUV radiation field derived for two cosmic ray ionization rates.
In Fig. 5 we show the distribution of the ratios of the H column density traced by [CII] to that traced by CO. In Table 2 we list a few diffuse clouds showing a large H layer around the CO emitting core. A majority of the clouds have . In 15 clouds the is 50 or greater than . In this sample of 50 transition clouds, on average, 24 of the total H column density is in the H/C layer which is not traced by CO. Although these estimates are only approximate, they show a likely scenario in the transition cloud structure. Lower densities ( 100 cm) and/or lower temperatures ( 50K) will increase the N(H) in the H/C layer (required to account for the observed I(CII)) by factors of 2-3, while higher density (500 cm) will decrease the N(H) by a factor of 2. However, at higher densities the temperature is likely to be 100 K and the required N(C) and N(H) will be larger.
We can use N(H) in the H/C layer and N(HI) in the HI layer derived from I(CII) and I(HI) to evaluate A in the cloud up to the C/CO transition layer. We define the C/CO transition layer as an inner cloud boundary where X(C) X(CO) = X(C)/2. We can now solve for the ratio of external FUV to density (/n(H)) balancing the photodissociation and the CO formation rates at the C/CO transition layer. We derive analytical photodissociation rates for the attenuation and the self-shielding which are consistent with those given by Lee et al. (1996). In the warm regions (in all our chemical modeling and analysis we use T 100 K) the H + OI chemistry dominates CO production over C + H (which dominates for T 35 K). Therefore, for the CO formation rates we use a simple chemical network incorporating the H + OI chemistry by extending the approach discussed by Nelson & Langer (1997) for a CO core surrounded by a warmer tenuous C envelope. In our calculation we use the reaction rates given by Glover et al. (2010). The results for a few clouds are listed in Table 2; the last three columns list A up to the C/CO transition layer and the external FUV, , in units of Draine radiation field (). (Though the solution to the chemical modeling was obtained as /n(H) here we give only as we have assumed n(H)= 300 cm in our analysis of [CII] intensities). In nearly half of our sample the clouds have very low FUV, in the range of 0.01 to 0.1. Such low values seem less likely in the ISM; it has been suggested that the [CII] in the ISM originates from clouds exposed to FUV, 10 (Cubick et al. 2008). Using PDR models Pineda et al. (2010) find that in a [CII] sample of dense molecular clouds the majority have = 1 - 10 . Furthermore it may be noted that the H + OI chemistry used here is sensitive to the cosmic ray (CR) ionization rate. Above we used the standard CR ionization rate (c.f. Shaw et al. 2008). However, there is recent evidence of much higher rates in the outer layers (low A) of clouds (Shaw et al. 2008; Indriolo et al. 2007 & 2009). We find that using 40 increases the derived value of the FUV substantially as shown in the last column in Table 2. At least two of the clouds with high FUV values (G337.82+0.00V-127 and G337.82+0.00V-118) are near the supernova remnant G337.8-0.1, about a shell radius from its boundary at V -122 kms (Caswell et al. 1975), and thus may be consistent with our results for higher value for CR ionization. However, for the cloud G345.65+0.00V-120, though this LOS passes near a HII region (G345.645+0.010), no enhanced radiation feature is observed at this V (Caswell & Haynes, 1987).
Our preliminary analysis assumes optically thin HI and C emission. We do not take into account the different beam sizes used in the observations. We do not include the gas traced by CI in the C/CO transition zone. Nevertheless, the results of our simplified approach show a definite statistical trend for the presence of a majority of “nominal” diffuse clouds with a thin H layer and a significant fraction of clouds with a thick H layer without any accompanying CO.
We have observed [CII] line emission in 16 LOSs towards the inner Galaxy and detected 146 velocity resolved [CII] components. We identify 53 of these components that are characterized by the presence of both HI and CO but no CO emission as transition clouds in which the conversion of C to CO is partial and a large fraction of carbon exists as C mixed with H in a “dark gas” layer surrounding the CO emitting core. Our results show that [CII] emission is an excellent tool to study transition clouds in the ISM, in particular as a unique tracer of molecular H which is not easily observed by other means. In about 10 of the clouds the H column density traced by the [CII] emitting layer is greater than that traced by CO emission. On average 25 of the H in these clouds is in the H/C layer which is not traced by CO. Finally our estimates of the FUV field indicate the CR ionization is likely much larger than the standard value in the outer layers, consistent with recent determinations from chemical abundances in diffuse regions.
Acknowledgements.We thank the referee for suggestions. This work was performed by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration. The Mopra Telescope is managed by the Australia Telescope, and funded by the Commonwealth of Australia for operation as a National Facility by the CSIRO.
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