Multi-wavelength observations of a nearby multi-phase interstellar cloud
C. Nehme, C. Gry, F. Boulanger, J. Le Bourlot, G. Pineau des Forets, E. Falgarone
aa r X i v : . [ a s t r o - ph ] F e b Astronomy & Astrophysics manuscript no. 8373 c (cid:13)
ESO 2018October 24, 2018
Multi-wavelength observationsof a nearby multi-phase interstellar cloud
Cyrine Nehm´e , C´ecile Gry , , Fran¸cois Boulanger , Jacques Le Bourlot ,Guillaume Pineau des Forˆets , and Edith Falgarone LUTH, Observatoire de Paris-Meudon, Universit´e Paris 7, France Laboratoire d’Astrophysique de Marseille,OAMP, BP 8, 13376 Marseille Cedex 12, France European Space Astronomy Center, RSSD, P O Box 50727, 28080 Madrid, Spain Institut d’Astrophysique Spatiale, Universit´e Paris Sud, Bat. 121, 91405 Orsay Cedex, France Laboratoire de Radio-Astronomie, LERMA, Ecole Normale Sup´erieure, 24 rue Lhomond, 75231 Paris Cedex 05,FranceReceived / Accepted
Abstract
Aims.
High-resolution spectroscopic observations (UV HST/STIS and optical) are used to characterize the physicalstate and velocity structure of the multiphase interstellar medium seen towards the nearby (170 pc) star HD 102065.The star is located behind the tail of a cometary-shaped, infrared cirrus-cloud, in the area of interaction between theSco-Cen OB association and the Local Bubble.
Methods.
We analyze interstellar components present along the line of sight by fitting multiple transitions from a groupof species all at once. We identify four groups of species: (1) molecules (CO, CH, CH + ), (2) atoms (C i , S i , Fe i ) withionization potential lower than H i , (3) neutral and low-ionized states of atoms (Mg i , Mg ii , Mn ii , P ii , Ni ii , C ii ,N i and O i ) with ionization potential larger than H i and (4) highly-ionized atoms (Si iii , C iv , Si iv ). The absorptionspectra are complemented by H i , CO and C ii emission-line spectra, H column-densities derived from FUSE spectra,and IRAS images. Results.
Gas components of a wide range of temperatures and ionization states are detected along the line of sight. Mostof the hydrogen column-density is in cold, diffuse, molecular gas at low LSR velocity. This gas is mixed with traces ofwarmer molecular gas traced by H in the J > + must be formed.We also identify three distinct components of warm gas at negative velocities down to −
20 km s − . The temperatureand gas excitation are shown to increase with increasing velocity shift from the bulk of the gas.Hot gas at temperatures of several 10 K is detected in the most negative velocity component in the highly-ionizedspecie. This hot gas is also detected in very strong lines of less-ionized species (Mg ii , Si ii ∗ and C ii ∗ ) for which thebulk of the gas is cooler. Conclusions.
We relate the observational results to evidence for dynamical impact of the Sco-Cen stellar association onthe nearby interstellar medium. We propose a scenario where the infrared cirrus cloud has been hit a few 10 yr ago bya supernova blast wave originating from the Lower Centaurus Crux group of the Sco-Cen association.The observations provide detailed information on the interplay between ISM phases in relation with the origin of theLocal and Loop I bubbles. Key words.
ISM:structure, ISM:clouds, ISM:kinematics and dynamics, ISM:individual objects:Chamaeleon clouds,Ultraviolet:ISM, Stars:individual:HD102065
1. Introduction
Astronomical observations provide multiple perspectives onthe interstellar matter, gas and dust, that are rarely gath-ered together to study the relation between interstellar-medium phases. It is the ambition of this paper to presentsuch a study on gas and dust observed in the direction of thenearby (170 pc), lightly reddened (E(B-V) = 0.17), B9IVstar HD 102065 in Chamaeleon (Table 1). In IRAS images,the infrared cirrus (Dcld 300.2-16.9) seen in the foregroundof HD 102065 has a prominent cometary shape and an un-usually blue IRAS colors (high 12 and 25 µ m to 100 µ mbrightness ratio), indicative of a large abundance of smallstochastically-heated dust particles (Boulanger et al. 1990). Correspondence to : [email protected]
We hereafter refer to it as the Blue Cloud in reference toits blue infrared colors. While the main Chamaeleon cloudshave a comparable extent in CO surveys as in IRAS im-ages, only the center of the Blue Cloud head is detected inCO (Boulanger et al. 1998; Mizuno et al. 2001). Boulangeret al. estimate that only 10% of the cloud mass is seenin CO emission. On larger angular scales, HD 102065 lieswithin the sky area where various evidences for an inter-action between the Scorpius-Centaurus (Sco-Cen) OB as-sociation and the Local Bubble have been gathered fromH i observations (de Geus 1992), the X-ray ROSAT survey(Egger and Aschenbach 1995) and optical absorption-linespectroscopy (Corradi et al. 2004).HD 102065 is located at the same Galactic longitude andslightly lower Galactic latitude than the Lower CentaurusCrux (LCC) group of the Sco-Cen association. Hipparcos Nehm´e et al.: Multi-wavelength observations of a multi-phase cloud positions, proper motions, and parallaxes have been used toshow that the LCC and the Upper Centaurus Lupus (UCL)group of stars, presently located within the Loop I bubbleat a distance of ∼
120 pc, were located closer to the Sun5 − ∼
20 supernovae explosions associated with the LCC andUCL groups that are estimated to have occurred in thelast 10 Myr. The structure and physical state of the inter-stellar medium seen in direction of the Sco-Cen associationshould reflect the action of the supernovae explosions pow-ering and progressively shredding, and spreading away thestars’ parent cloud. This scenario has been introduced intoa numerical simulation based on the supernovae-driven ISMmodel of de Avillez (2000) by Breitschwerdt and de Avillez(2006). In the reproduction of the local interstellar mediumgenerated by this simulation (Figure 1 in Breitschwerdt andde Avillez), HD 102065 is located within the Loop I bub-ble and the line of sight to the star crosses the shell thatprovides separation from the Local Bubble.Observations characterizing the gas and the dust alongthe line of sight to HD102065 have been presented in ear-lier publications (Boulanger, Prevot & Gry, 1994 ; Gryet al, 1998). Unlike most early-type stars observed in UVspectroscopy, the position of HD 102065 corresponds to noenhancement in IRAS emission, implying that the mat-ter responsible for absorption is not heated by the star(Boulanger et al. 1994). The extinction curve produced us-ing IUE low-resolution spectroscopy, shows a strong 220 nmbump and a weak far-UV rise (Boulanger et al. 1994).Moderate-resolution spectroscopy ( λ/ ∆ λ = 20 000) ac-quired by FUSE, detected molecular hydrogen absorptionlines. These observations showed that H at a temperatureof ∼
60 K is the dominant state of hydrogen along the lineof sight to HD 102065. The presence of a smaller amountof warmer H (a few 100 K) was inferred from the H exci-tation (Gry et al. 2002). HD 102065 was also observed us-ing the Goddard High Resolution Spectrometer (GHRS)on the Hubble Space Telescope (HST). These observationsrevealed the presence of high, negative-velocity gas withunusually high Si ii excitation that was seen as the signa-ture of the dissipation of a large amount of kinetic energy(Gry et al. 1998).A number of questions remain unanswered. Why is thesmall dust abundance enhanced in the Blue Cloud, which,from the FUSE and optical spectroscopy, appears to bea typical diffuse molecular cloud? How was the negative-velocity gas accelerated? Are the low and negative-velocitygas physically connected? More generally, how can the databe interpreted within the present understanding of the in-teraction between the Loop I and the Local bubbles?In this paper, we present high-resolution ( R = 100 000)UV spectr,a obtained using the HST Imaging Spectrograph(STIS) complemented by optical absorption spectra, andCO and H i emission spectroscopy (Sections 2 and 3). Thesedata improve previous observations in spectral resolutionand spectral coverage. The multi-wavelength spectra areused to describe the physical state and velocity structure ofthe multiple gas components observed along the line of sight(Section 4). In Section 5, we discuss the cool, low-velocitygas, and in Section 6, we relate the observational resultsto evidence for dynamical interaction between the Sco-Censtars and the nearby interstellar medium, and discuss the HD 102065 α (2000) 11 43 37.87 δ (2000) -80 28 59.4 l (2000) 300.02 ◦ b (2000) -18.00 ◦ Sp. Type B9IVd (pc) 170E(B-V) 0 . ± . A v . ± . R v . ± . ,J=0) (cm − ) 2 . ± . N(H ,J=1) (cm − ) 1 . ± . N(H ,J=2) (cm − ) 2 . − . N(H ,J=3) (cm − ) 1 . − . N(H ,J=4) (cm − ) 0 . − . N(H ,J=5) (cm − ) 0 . − . N H (cm − ) 9 . f = 2 N ( H ) /N H . ± . Table 1.
HD 102065 “fact sheet”. A v is the visible extinc-tion, E(B-V) is the color index (Boulanger et al. 1994). H column densities from (Gry et al. 2002). The total columndensity (N H = N(H) + 2 N(H )) has been derived fromE(B-V).possible interplay between the multiple components. Themain conclusions of this paper are gathered in Section 7.
2. Ultraviolet absorption lines: STIS data.
We acquired STIS UV spectra of HD 102065 with the high-resolution MAMA Echelle gratings E140H and E230H. Weused five different grating settings covering the followingwavelength regions: 1144 - 1324 ˚A, 1322 - 1512 ˚A, and 1494- 1674 ˚A (E140H with aperture 0 . × . . × . . − . The line spread functions (LSF) are tab-ulated in the STIS Instrument Handbook. We have per-formed a double-Gaussian fit to the tabulated LSF and usethe resulting FWHM values in the profile-fitting software.The results of the double-Gaussian fit to the LSF were anarrow component with FWHM ∼ ehm´e et al.: Multi-wavelength observations of a multi-phase cloud 3Element Wavelength (˚A) f-value Element Wavelength (˚A) f-valueC I 1276.4822 0.449 10 − Fe II 1608.4510 0.58 10 − C I 1280.1353 0.229 10 − Fe II 1611.2004 0.136 10 − C I 1328.8333 0.631 10 − Fe II 2586.6499 0.691 10 − C I 1560.3092 0.128 10 − Fe II 2382.7651 0.320C I 1656.9283 0.140 10 Fe II 2600.1729 0.239C I ∗ − Mn II 2576.8770 0.361C I ∗ − Mn II 2594.4990 0.280C I ∗ − Mn II 2606.4619 0.198C I ∗ − Mg II 1239.9253 0.617 10 − C I ∗ − Mg II 1240.3947 0.354 10 − C I ∗ − Mg II 2803.5305 0.306C I ∗ − Mg II 2796.3518 0.615C I ∗ − S II 1250.5840 0.543 10 − C I ∗ − S II 1253.8110 0.109 10 − C I ∗ − S II 1259.5190 0.166 10 − C I ∗∗ − P II 1301.8743 0.127 10 − C I ∗∗ − P II 1532.5330 0.303 10 − C I ∗∗ − C II 1334.5323 0.128C I ∗∗ − C II ∗ ∗∗ − Si II 1304.3702 0.917 10 − S I 1295.6531 0.870 10 − Si II 1193.2897 0.585S I 1296.1740 0.220 10 − Si II 1260.4221 0.118 10 S I 1425.1877 0.365 10 − Si II 1190.4158 0.293S I 1473.9943 0.730 10 − Si II ∗ S I 1474.5706 0.121 10 − Si II ∗ ∗ ∗ Si II ∗ − CO A-X (2,0) : Si II ∗ − O I 1302.1685 0.519 10 − R(1) 1477.5148 1.967 10 − O I 1355.5977 0.116 10 − R(2) 1477.4786 1.574 10 − O I ∗ − Q(1) 1477.6509 1.967 10 − O I ∗∗ − Q(2) 1477.6827 1.967 10 − N I 1199.5496 0.130CO A-X (1,0) : N I 1200.2233 0.862 10 − R(0) 1509.7504 2.855 10 − N I 1200.7098 0.430 10 − R(1) 1509.6985 1.449 10 − N I 1159.8168 0.851 10 − R(2) 1509.6639 1.183 10 − N I 1160.9366 0.240 10 − Q(1) 1509.8379 1.427 10 − Ni II 1317.2170 0.775 10 − Q(2) 1509.8738 1.449 10 − Ni II 1370.1320 0.131CO A-X (0,0) : Ni II 1454.8420 0.595 10 − R(0) 1544.4515 1.598 10 − C IV 1550.7770 0.948 10 − R(1) 1544.3914 7.977 10 − C IV 1548.2030 0.190R(2) 1544.3472 6.361 10 − Si III 1206.500 1.67Q(1) 1544.5432 7.993 10 − Si IV 1393.755 0.514Q(2) 1544.5743 7.983 10 − Si IV 1402.770 0.255
Table 2.
Wavelength and f -values for the observed absorption lines. Most of those values are from Morton (2000 and2001) and Verner et al. (1994). Oscillator strengths values for S i are taken from Beideck et al. (1994). Values for O i . ii . ∼ . − . The achieved signal-to-noise ratio (hereafter SNR) inthe continuum, ranges from 23 to 41 depending on the grat-ing setting. At the positions of most interstellar lines theSNR ranges from 20 to 40, except for the lines Si ii *, C ii and C ii *, where it drops to 6 because they are located atthe bottom of strong stellar lines. Nehm´e et al.: Multi-wavelength observations of a multi-phase cloud
Figure 1.
Molecular lines. From top to bottom: CO J=2-1 emission line (Section 3.3) in the millimeter range with theSEST, two examples of the CO ultraviolet absorption bands(A-X (2,0) and A-X (0,0)) observed with STIS, and the CHand CH + visible absorption lines (Section 3.1). Histogram-style curves represent the observations normalized to thecontinuum, solid lines are the best fits, and dashed lines, theprofiles before convolution with the instrument line spreadfunction, i.e. the intrinsic interstellar profile. The letters A,B and C indicate velocity components that will be intro-duced in the spectral analysis. Here we provide an overview of the interstellar, absorptionlines present in the STIS spectra and used in the analysis.They are listed in Table 2.We identify four groups of species, classified accordingto the characteristics of their line profiles, which reflectsthe distribution and the nature of the gas traced by thegiven species. Examples of the line profiles are presented inFigures 1 to 4. The figures also show the velocity compo-nents and fits to the profiles, which are the results of theanalysis which will be described in Section 4 after presen-tation of the complementary data in Section 3.(1)
Molecules: see Figure 1. Several CO electronic A-X bands are present in the HST/STIS spectra. We use thefour following bands where CO is detected in several linesfrom rotational levels J”=0, 1 and 2 in the HST/STIS ob-
Figure 2.
Sample lines from neutral species. Histogram-style curves represent the observations (in erg cm − s − − ),solid lines are the best fits, and dashed lines, the profilesbefore convolution with the instrument line spread func-tion, i.e. the intrinsic interstellar profile. An asymmetry isevident on most profiles which are therefore best fitted withtwo components, called A and B.servations : A-X (0,0) at 1544 .
54 ˚A, A-X (1,0) at 1509 . .
57 ˚A and A-X (3,0) at 1447 . J ” = 0, 1 and 2, as illustrated in Figure 1where two bands are shown as examples.(2) Neutral species with ionization potential lowerthan that of H i : see Figure 2.Atomic carbon lines have been measured in the three fine-structure levels of the P fundamental level: J = 0 (here-after C i ), J = 1 (hereafter C i *), and J = 2 (hereafterC i **). 36 carbon lines have been detected, among which wehave selected 20 for use in the fits, eliminating lines whichare either heavily blended with another line, or lack reliablef-values. A representative subset of those lines is shown inFigure 2, with lines from the other neutral species, Fe andS. These neutral lines are narrow with a width of only a fewresolution elements ( i.e. a few km s − ), and clearly asym- ehm´e et al.: Multi-wavelength observations of a multi-phase cloud 5 Figure 3.
Sample of lines from ionized species and neutral species with ionization potential close to that of H i . Sameline style as Figure 2. Most fits have been performed with the same line-of-sight model, made of 4 components withsame velocity and same temperature for all species. There are only three exceptions: (i) An additional component ofhigh-temperature gas is included at the velocity of component 4 to fit the extended blue wing of the strong Mg ii andC ii * lines. (ii) In the O i ** profile an additional component is considered, blueward of component 4. (iii) The S ii profilesinclude an extra absorption on the red side of component 1. Nehm´e et al.: Multi-wavelength observations of a multi-phase cloud
Figure 4.
Highly-ionized species observed with STIS. Asan indication, consistent absorption line profiles for gas ata temperature of 900 000K and at the velocity of component4 are superimposed.metric indicating the presence of two velocity componentsreferred to as A and B.(3)
Low ionization species and neutrals with ion-ization potential higher than that of H i : see Figure 3.In those species the absorption extends to negative veloci-ties as low as −
50 km s − for some of the lines. The broadabsorption is resolved into four distinct components. Theyare particularly well-resolved in the lines of the heavy ele-ment Fe ii . These components are numbered 1 to 4 in orderof decreasing velocity.(4) High ionization species: see Figure 4.High ionization species Si iii , Si iv , C iv are detected in thisline of sight. Their profile is broad, and centered at negativevelocity close to the most negative velocity component seenin the other ions, Component 4.
3. Complementary observations
We present additional data sets that complement our STISobservations.
Optical spectra of the CH and CH + lines about 4300 ˚A(see Table 3), the C A Π u − X Σ +g (2 ,
0) and CN A Π − X Σ + (2 ,
0) molecular bands about 8760 and 7905 ˚ A ,were obtained in April 1990 using the Coude EchelleSpectrometer at the 3 . Spec. CH CH + Band A ∆ − X Π (0 ,
0) A Π − X Σ + (0 , (1)+R (1) R(0) λ (˚A) 4300 .
313 4232 . f .
06 10 − .
45 10 − Table 3.
Molecular absorption lines detected in the opticalspectra. Listed parameters are from Gredel et al. (1993).The spectrograph was connected to the 3 . . − . Thesampling was 2.2 pixels per resolution element, and thespectra were 1024 pixels in length. Wavelength calibrationwas performed using an observation of a Th-Ar cathodelamp. This calibration was repeated for each set-up of thespectrometer. The identification and line-center measure-ments of a set of around 20 lines in these spectra allowedus to derive the dispersion relation with an accuracy betterthan 10 m˚A. Flat-field images were acquired at the be-ginning and end of each set of observations, for the entirewavelength range.The images were corrected for dark current, and thenflat-field corrected. Spectra were extracted from the imagesby averaging 20 lines over which the flux is spread. Thespectra were normalized by dividing the original spectraby a median average showing no spectral lines from inter-stellar gas or the Earth atmosphere. Atmospheric lines wereidentified using observations of HD 106911, a bright star inChamaeleon with almost no reddening.Sections of the normalized spectra that include the CHand CH + lines, are presented in Figure 1. The C andCN spectra show no detections with 3 σ upper limits of2.7 and 1 . A , respectively. Gaussian fits to the CH andCH + lines have been used to derive the line parameters,and thereby the column density N . Because the transitionsare optically-thin, N is related to the measured equivalentwidth, W λ , by the equation: N = 1 .
13 10 W λ /λ f with W λ and λ in ˚A, N in cm − and where f is the line oscillatorstrength given in Table 3. Results are presented in Table 5.The optical absorption lines are also fitted, together withthe UV absorption lines, in the general spectral analysis de-scribed in Section 4.2.2. This, in particular, allows the CHand CH + velocity distribution to be related to the velocitystructure seen in other species. ehm´e et al.: Multi-wavelength observations of a multi-phase cloud 7 i emission A sensitive 21-cm H i southern sky survey has recently beenpublished (Bajaja et al. 2005). Its angular resolution is halfa degree and the velocity resolution 1 . − . In Figure 5, Figure 5. H i emission line in the direction of HD 102065derived from the IAR Southern Sky survey. The thick linerepresents the difference with respect to a reference positionoutside of the Blue cloud.we present the survey spectrum for the position closest toHD 102065 (thin line). The thick line is the difference be-tween this spectrum and an OFF cloud spectrum obtainedby averaging 3 spectra outside the Blue Cloud (Dcld 300.2-16.9) at l = 297 ◦ to 297 . ◦ and b = − . ◦ to − ◦ , corre-sponding to a local minimum in the 100 µ m IRAS map. Thissubtraction is supposed to remove the background emis-sion, but may also remove foreground emission that is notspatially correlated with the Blue Cloud. Within unknownvariations of the background emission between the ON andOFF cloud positions, this difference spectrum representsthe H i spectrum of the Blue Cloud. In Table 4, we list theresults of a Gaussian decomposition of this spectrum. Thecolumn densities have been calculated with the assumptionthat the H i emission line is optically-thin.In Figure 6, we show the spatial distribution of the H i emission in two-channel maps (V LSR = − − )overlaid over the IRAS 12 µ m image (Miville-Deschˆenes andLagache, 2005), and in Figure 7 a position-velocity diagramacross the Blue cloud at the Galactic latitude of HD 102065.These two figures show that the Blue Cloud, the brightIRAS 12 µ m emission on the map, is associated with gasat low velocities. The total brightness spectrum (thin linein Figure 5) shows the presence of a broad, negative ve-locity component, extending down to −
40 km s − . This gasis not spatially correlated with the infrared emission fromthe Blue Cloud. It is however observed in UV-absorptionspectra, and must therefore be located in front of the star. Long integrations have been performed with the 15m SESTtelescope on the CO J =2-1 and J =1-0 emission lines.The J =1-0 and 2-1 lines were observed in 1990 and 2001,with system temperatures of 700 K and 200 K respectively, Figure 6.
IRAS 12 µm image (IRIS processing, Miville-Deschˆenes and Lagache, 2005) with contours of H i emis-sion at V LSR = − − (white dashed) and 5 km s − (black solid) taken from the IAR survey (30’ resolution and1 . − width velocity channels, Bajaja et al. 2005). Theposition of HD 102065 and the T Tauri star T Cha aremarked. Contours are labeled with the brightness tempera-ture. The 30 K contour is thicker. The cloud tail is roughlyat a constant Galactic longitude. The Galactic coordinatesof HD 102065 are l = 300 . ◦ and b = − . ◦ . The elec-tronic version is in color. Figure 7. H i position-velocity diagram across the BlueCloud at the Galactic latitude of HD 102065 (b=-18 ◦ ). Thisfigure shows that the two velocity components at respec-tively 0.1 and 4 . − each extends on one side of thestar position (l=300 ◦ ) where they overlap. The emission atthe OFF position (see text) has been subtracted to high-light the gas associated with the cloud.and a velocity resolution of 0 . − . The rms noise levelof the CO J =2-1 line displayed in Figure 1 is suffi- Nehm´e et al.: Multi-wavelength observations of a multi-phase cloudspecies V
LSR (km s − ) T b (K) FWHM (km s − ) b (km s − ) N (cm − )H i
21 cm (difference) 14 . ± . . ± .
15 7 . ± . . ± . . ± . . ± . . ± . . ± . . ± . . ± .
07 10 . ± . . ± . . ± . . ± . . ± .
06 10 − . ± . . ± .
07 16 ± . . ± . . ± . CO(1-0) 2 . ± . . − . ± . . ± . . ( J = 1)CO(2-1) 2 . ± .
17 3 . − . ± . . ± . . ( J = 2) I (erg s − cm − sr − ) N(C II *)C + µ m - 2 . − - - 1 . Table 4.
Results from emission line measurements. H i measurements are performed on the difference spectrum (seetext and Figure 5). The H i component at 14 . − is not detected in absorption and is probably due to backgroundgas situated behind the star HD 102065.ciently low ( ∼
10 mK) to reveal a broad and weak compo-nent (∆ v = 2 . − , T ∗ A = 36 mK ), and a marginally-detected narrower component (∆ v = 0 . − , T ∗ A =15 mK ). The antenna temperature has not been correctedfor beam coupling to the sky, because the source is muchmore extended than the telescope-beam size at this fre-quency ( ∼
22” or 0.015 pc at the distance of the cloud).This is not a critical issue because the beam efficiency was ∼ .
85 at 230 GHz. Similarly, the rms noise level of the J =1-0 line is 18 mK for a velocity resolution of 0 . − .Given the weakness of the CO lines, the assumption ofoptical thinness is justified and the column densities ofmolecules in the J = 2 and 1 levels are inferred from the lineintegrated intensities W ( CO ) = 0 . . K km s − ofthe broad and narrow components of the J = 2 − W ( CO ) = 0 . K km s − for the J=1-0 line. The respectivecolumn densities are given in Table 4. + observations The C + fine-structure transition at 158 µ m has been mea-sured towards HD 102065 with the Long WavelengthSpectrometer (LWS) of the Infrared Space Observatory(ISO), with a spectral resolution of R = 300, much lowerthan all other data presented in this paper. The datahave been processed following the LWS handbook (Gryet al 2003). The resulting line intensity I(C + λ µ m) is2 . ± . − erg s − cm − sr − , from which we estimatethe C ii * column density N(C ii *) = 1 . ± . cm − .This is a measurement of the mean column density over the80” LWS beam.
4. Analysis of the line of sight - Description of theinterstellar components
To derive the characteristics of the interstellar components,i.e. column density N , velocity v and broadening parame-ter b , we fit the observed line-profiles with theoretical ones,which are the results of the convolution of a Voigt profilewith the instrumental function (LSF). We use a profile-fitting method, instead of a traditional measurement ofequivalent width, because the velocity structure is complexas can be seen in Figures 1 to 3. Because several com-ponents are blended together in most lines, profile fittingis required to derive the characteristics of the individualcomponents. The data fitting has been performed with the use of the profile fitting software “Owens.f” developed byMartin Lemoine.The software allows several lines of the same element, aswell as lines from different elements, to be fitted together.It also allows several velocity components to be fitted si-multaneously, and the components to be included in theabsorption profiles for each element to be specified. Thesoftware therefore derives a global and consistent solutionfor a group of species. Multiple iterations are necessary tofind which species could be fitted in a consistent way withthe same velocity components. A common component fordifferent species obviously implies a common absorption ve-locity but also common physical conditions, implying con-sistent broadening parameters.The fitting software breaks the line-broadening param-eter (b-value) into thermal broadening, which depends onthe element mass, and non-thermal broadening (turbu-lence), which is the same for all elements in a component.Therefore, fitting lines from elements of different massessimultaneously, in principle enables the simultaneous mea-surement of temperature, turbulent velocity, the velocityitself, and column densities. In practice interstellar absorp-tion lines are usually found to be of two types. The firsttype displays unsaturated lines from heavy elements thathave similar b-values independently of the element mass: inthis case, the broadening is primarily non-thermal, imply-ing that T is low and can be estimated only if the analysisincludes an element whose mass is low enough to introducea significantly different b-value. Here we will consider H i emission lines. In the case of the second type of interstel-lar absorption line, the b-values vary with element massdue to a significant thermal effect, implying that T is highand can be estimated. In practice the interstellar absorp-tion lines are broader than the resolution element (for Thigher than a few thousand K), and strong lines are occa-sionally broadened considerably by saturation, leading tothe blending of components and limiting the precision withwhich broadening parameters can be derived.The quality of the fits, and the error bars on columndensities, have been computed using the ∆ χ method de-scribed, for example, in H´ebrard et al (2002). We performedseveral fits by fixing the column density of a given elementX to a different value in each fit. In all of these fits, allother parameters are not fixed and we compute the best χ for each value of column density. The plot of ∆ χ versus N ( X ) provides the 1, 2, 3, 4, 5 σ range for N (X). This isillustrated with an example in Figure 8. Examples of the ehm´e et al.: Multi-wavelength observations of a multi-phase cloud 9 Dc N(CI) A (cm −2 ) Figure 8.
Example of a ∆ χ curve, here for C i in com-ponent A. Each point corresponds to a fit performed withthe C i column density in component A fixed to the valuegiven in abscissa. ∆ χ is rescaled by dividing it by the 95%-confidence level for this fit (rescaled ∆ χ = 1 . σ -uncertainty is given by ∆ χ = 3 × .
41 = 12 . i . The common asymmetry of the absorption lines of the neu-tral species (Figure 2) suggests the existence of two velocitycomponents in all of these lines. When fitting simultane-ously C i , C i *, C i **, S i and Fe i , a best fit is achievedfor two components, that we call component A and com-ponent B, for the following LSR velocities: V LSR = 3 . . − . They may be converted into heliocentric veloc-ities by adding 10 . − .The b-values derived in the fits, b A = 1 . ± . b B = 1 . ± . − are not significantly different for thethree species with different masses. This is strong supportto the fact that broadening is primarily non-thermal.Interestingly, the velocities identified for the two maincomponents in the H i Blue Cloud emission spectrum (seeTable 4) are identical to those of A and B within the errorbars. The b-values are also similar to those expected in coolgas, with the H i b-values being slightly higher than thatof S i and C i due to the mass difference. The main com-ponents seen in emission are therefore likely to be relatedto the two main components observed in absorption. In contrast, the CO lines do not accept a two-componentsolution. The singularity of the CO component is evi-denced by the small width of the CO lines, with b CO =1 . ± . − , similar to that of component A or B alone.This width is significantly lower than the velocity differ-ence between A and B, 3 . − , toward which value b CO would tend to increase if CO was distributed over the twocomponents. When fitting the CO lines in a two-componentmodel together with all well-behaved (i.e. not heavily sat-urated) neutral lines, unavoidably in the resulting fit mostof the CO column density is found in one of the compo- nents, and the resulting theoretical profile appears slightlyshifted relative to the observations. This component cannotbe identified with either A or B, since its velocity shift issignificantly higher than the uncertainty of the wavelengthcalibration in a single STIS sub-spectrum. Alternatively,when fitting the neutrals and CO lines altogether in a one-component model, the fit to the neutral lines is significantlydegraded. The line characteristics (central velocity and linewidth) of CO in emission agree with that derived from UVabsorption lines. This agreement validates the UV wave-length calibration and confirms that CO, unlike H i andthe atomic species C i and S i , is found at a single velocity.We infer that CO is detected primarily in one com-ponent at a velocity intermediate to that of A and B,V LSR
CO = 2 . − . We will call this component C (seeFigure 1 illustrating the line fit for two CO bands).Further information is provided by the CH and CH + absorption line profiles. We performed the profile fittingof the CH and CH + lines with Owens together with theSTIS data of the species C i , S i and CO. We made severalfits with CH and CH + either using one or two components,and with a velocity shift as a free parameter to take into ac-count uncertainties in the absolute wavelength calibrationof both the UV and optical spectra. The best results havebeen obtained if (i) CH + is present in both components Aand B and (ii) most of the CH absorption is in one compo-nent, corresponding to the CO component C. This solutionimplies that the optical spectra are shifted by about 0.3km s − relative to the STIS data, a value consistent withthe wavelength calibration uncertainty.We conclude that the bulk of the CO and CH moleculesare in a single component (C) at a different velocity to thebulk of the neutral atoms carbon, sulfur, iron and hydrogenand the CH + molecules, which appear to be distributedover two well-separated velocity components (A and B).Of course it is not excluded that some neutrals are presentat the velocity of CO, and alternatively that some CO ispresent in the components A and B, at undetectable levels.We believe that the three components A, B and C are, infact, physically related. Within this assumption, the threevelocity components would reflect the particular velocitystructure associated with local abundance variations withina single cloud . i . Almost all ions and neutrals with an ionization potentialslightly higher than H i : Mg ii , Mg i , Mn ii , P ii , Ni ii ,C ii , N i and O i , can be fitted altogether with a uniquesolution for the line of sight structure : component 1 at 2 . ± . − , component 2 at − . ± . − , component 3at − . ± . − and component 4 at − . ± . − (velocities are LSR, see Table 5).The low-velocity component, component 1, correspondsto the gas detected in neutral and molecular lines. In theunsaturated ion lines this component can be fitted ei-ther with two components at the velocity of A and B,or with a single component at the intermediate velocityV LSR = 2 . ± . − without affecting the columndensity results or the quality of the fits (monitored by the χ ). However when adding stronger lines to the fit, it hasbeen impossible to find a good fit to the spectra in com-ponent 1 with the column density provided by the unsatu- rated lines and the A+B velocity structure. This suggeststhat the structure of this component is more complex andmay indicate the presence of warmer gas at the velocity ofcomponent 1. Even with a low column density, in stronglines warm gas can contribute a large part of the equivalentwidth due to a larger b-value. For the ions, we thereforequote a unique column density, obtained using the unsatu-rated lines when they exist, for the low velocity component1 alone.The components structure is summarized in Table 5. Once the velocity structure is defined, column densities canbe derived following the method described in section 4.1.In this respect, not all of the lines are equally useful.Fe ii is the most favorable ion in our sample for tworeasons. First, several Fe ii lines are available with a largerange of f-values from 0.0014 to 0.32 (see Table 2). Stronglines allow detection of weak components, and weak linesenable non-saturated profiles of strong components to beanalyzed. Another advantage of Fe ii lines is that, due toits high mass, Fe is less affected by thermal broadening thanlighter elements; therefore in warm gas Fe ii absorption linesare narrower and the components better-resolved as seen onFigure 3. The errors in N(Fe ii ) are therefore due only tothe combination of noise and uncertainty in the continuumplacement.We achieve our poorest quality analyses for ions forwhich only strong lines are available. This is the case herefor C ii , C ii ∗ and Si ii , for which all components are satu-rated and for which we can derive only lower limits.When only one weak line for a given element is detected,only the strongest component 1 is detected and measuredand we get column-density upper limits for all other com-ponents. This is the case for both Ni ii and P ii .For some elements, we have both a very weak line,which is only detected in the strongest component, allowingmeasurement of the column density, and very strong linesthat are saturated in all components. Combining weak andstrong lines gives upper and lower limits for the weak com-ponents. This is the case for Mg ii λ i λ . × and 10 times weaker thanthe strong lines Mg ii λ i λ ii is anintermediate case where the three S ii lines are saturatedin component 1 and at least one S ii line is unsaturatedin the other weaker components. The uncertainty in theircolumn density is due to blending between components anduncertainty in the continuum placement.As mentioned in Section 2.2, Si ii * and O i ** are de-tected primarily at the velocity of component 4, whilethey are hardly detectable in components 1, 2 and 3.Nevertheless O i * is detected in the four components.Column densities resulting from absorption line fittingare provided in Table 5. In this section, we compare column densities derived fromemission and absorption spectra.The total hydrogen column density can be es-timated from the color excess, and the standard ratio N (H tot ) /E (B − V) = 5 . × cm − mag − (Bohlin et al. 1978). This provides a measurement of N (H tot ) = 9 . × cm − . It can also be derived in ad-dition using H column densities measured in FUSE ab-sorption spectra (Table 1), and the H i column densityestimated from the Blue Cloud H i emission spectrum : N (H tot ) = N (H)+2 N (H ) = 9 . × cm − excluding the14 . − background component (Table 4). The quanti-tative agreement between emission and absorption resultsfor column densities, velocities, and line widths, supportsthe idea that the two H i low velocity components observedin the Blue Cloud H i emission spectrum correspond tocomponents A and B observed in absorption.One can also note that there is a consistency betweenthe C + column density derived using ISO-LWSdata, andthe lower limits measured from the saturated absorptionlines.CO column densities were measured both in emissionand absorption. In absorption, they can be directly com-pared to the H column densities derived from FUSEspectra (Gry et al., 2002). This comparison enables themeasurement of a total CO abundance N (CO) /N (H ) =1 . +0 . − . − .It was interesting to find that the CO column densitieswere larger when determined in emission than in absorp-tion, by a factor 3 in the J = 1, and by more than a factorof 7 for the J = 2 level. Because the spatial resolutionsof the observations are very different, this could imply theexistence of small-scale structure of the CO emitting gaswithin the SEST beam. On the other hand, part of theCO emission could be background emission from the en-velope of the nearby Chamaeleon III cloud, which is ob-served over the same velocity range (Mizuno et al. 2001).Based on measurement of the CO excitation, we considerthat this second possibility is however unlikely. The col-umn density ratio between the J = 2 and 1 levels, derivedfrom the emission spectrum provides an excitation temper-ature T ex = 2 . N ( J = 2) /N ( J =1) = 0 . , isnegligible when compared to radiative excitation, hence thelocal density is much lower than the critical density of the2-1 transition n H = 10 cm − . However, if the gas was as-sociated with the Cham III cloud, we expect that radiativeexcitation alone could produce a higher T ex . As noted in Section 4.1, information on the temperatureof the gas in each component follows from the ability tofit simultaneously lines from several elements of differentmasses, which allows the nature of the broadening to bedetermined.
For components A and B, the similarity of the b -values forlines from heavy elements of different atomic masses (seeSection 4.2) suggests that their widths are dominated bynon-thermal motions. To estimate the gas temperature, wecompared heavy-element line-widths to the widths of theassociated H i line emission. ehm´e et al.: Multi-wavelength observations of a multi-phase cloud 11Component A C BV LSR (km s − ) 3.8 ± . ± . ± .
1T (K) 85 ±
65 - 270 ± NT ( km s − ) 1.5 ± . ± . ± . J = 0 3 . ± . . (13)CO J = 1 2 . ± . . (13)CO J = 2 < . . ± . + ± ± i ( P ) 3 . ± . . (14) 1 . ± . . (14)C i * ( P ) 6 . ± . . (13) 2 . ± . i **( P ) 1 . ± .
15 (13) 3 . ± . . (12)S i . ± . . (12) 2 . ± . i . ± . < .
81 (10)Component 1 (A+B) 2 3 4V
LSR (km s − ) 2 . ± . − . ± . − . ± . − . ± .
2T (K) - 8 500 ± ± ∼
20 000 − ≥
200 000Fe ii . ± . . ± . . (13) 1 . ± . . ± . ii . ± . . (15) 1 . . . . . . ii > . . ± . . (13) 1 . ± . . ± . . (14)P ii . ± . < . < . < . ii . ± . . (13) 1 . ± . . ± . . ± . ii > . ii * > . ii > . ii * < . < . < . > . ii . ± < . < . < . i . ± . . (12) < . < . ± i . ± . . . . . . i . ± . . . . . . . iii < . iii > iv . ± . iv . ± . Table 5.
Interstellar component characteristics derived from the absorption lines in the UV and optical spectrum. Theions absorption in component 1 may be distributed over the velocities of components A and B however the fits have beenperformed with one component.
Velocities:
The listed uncertainties on V LSR apply to relative velocities. The absoluteuncertainty from the precision of the STIS wavelength calibration is about 1 . − , absolute velocities have howeverbeen confirmed to about 0 . − by the good agreement between component C and the CO emission in the heterodynedata, which have an absolute precision of 0 . − . Temperature:
When possible, temperatures are derived from thecombination of b -values from species of different masses. The precise temperatures for components A and B have beenderived under the assumption that they can be identified to the two components seen in emission in the Blue Cloud H i spectrum. Column densities (in cm − ): Numbers in parentheses are powers of 10. All error bars are 3 σ uncertainties,except for the intervals given by the arrows, indicating and upper limit derived from a faint line and a lower limit derivedfrom a strong saturated line. Si ii , C ii and C ii * profiles are saturated for all components, implying that only a lowerlimit could be derived for the whole absorption.Although the beam of the H i emission observa-tions is several orders of magnitude larger than the pen-cil beam of UV absorption observations (0 . ◦ vs milli-arcseconds), we assume that the samplings of the veloc-ity field are similar. By computing spectra produced bynumerical simulations of mildly compressible turbulence,Pety and Falgarone (2000) showed that one-dimensionalsampling of a turbulent, homogeneous field is similar tothat performed by an extended beam. This behavior of thelinewidths is indeed found in mm-line observations of COin emission and absorption against extragalactic sources(Liszt and Lucas 1998). Once the H i linewidths are in-cluded in the analysis, the temperatures, yielded formallyby the b -value differences between H i and heavy neu- trals, are 85 ±
65 K for A and 270 ±
100 K for B. Forboth components, the non-thermal velocity contribution tothe linewidths is ∼ . − . We note that these tem-peratures apply to the atomic species, and not necessar-ily to CO and CH because these molecules appear to bekinematically-decoupled from the atomic species. Components 2 to 4 belong to the second type of interstel-lar absorption lines, and T is estimated using the UV-linefitting procedure. Component 2 has a temperature typicalof warm diffuse gas (7 000–10 000 K), while in Component3 and 4 the thermal broadening is clearly dominant, with temperatures of 13 000 ± Indeed the strong Mg ii lines show the presence of addi-tional high temperature gas at approximately 200 000 K,moving at a velocity close to that of component 4. Thepresence of high temperatures is well-illustrated by the cleardifference between the blue wings of the Mg ii profile andthe N i profile in Figure 3. While the relatively steep bluewing of the N i line can be fitted with a temperature close to20 000 K for component 4, the extended blue wing in Mg ii requires the addition of a hot component at the same ve-locity, of a temperature ∼
200 000 K, and a column densityabout 10% of that of the total column density of component4. If this extended blue wing was a damping wing due to acomponent with very high column density, this componentwould create a significative absorption in the faint Mg ii line at 1240 ˚A, at least as strong as that of component 1.However, no additional component is detected significantlyin Mg ii λ ii * lines compared to the width of other lines ofmoderate strength. The FWHM of the Si ii * line, shown inFigure 3, is ∼
18 km s − , corresponding to ∼
200 000 K. Inthe C ii * line, the blue wing is extended, up to LSR veloc-ities of −
50 km s − . This may correspond to the presenceof very high temperature gas at a velocity close to that ofcomponent 4, as in the case of Mg ii .The high-ionization species Si iii , Si iv and C iv havebeen detected in the spectra at the velocity of component 4.They are broad and the temperature implied by their widthis at least 200 000 K. A temperature close to 10 K, is infact consistent with the observations (see figure 4 for anillustration), although the large width of the lines could bedue to the presence of several components at a tempera-ture of a few 10 K. OVI is unfortunately not available inabsorption in the FUSE spectra, because the stellar contin-uum is too weak at these wavelengths.
The description of the line of sight that emerges from thedirect analysis of the spectral data is complex. The bulk ofthe gas is at low LSR velocity. Neutral low-potential atomsand most of the H i are distributed over two componentsseparated by ∼ − (called A and B) that appear inthese species to have temperatures of at most a few hundredK, and are associated with CH + . The molecules CO and CHhowever, are within a single component of similar b-value( b ∼ . − ) at a velocity intermediate between that ofA and B.Other components exist at high negative velocities withrespect to the bulk of the gas. Negative-velocity gas is seenin the absorption lines of ions and atoms with ionizationpotentials higher than that of H i , as well as in emissionin H i . The temperature of these components ranges from typical temperatures of diffuse atomic gas ( ∼
10 000 K) tosomewhat higher temperatures (comparable to 20 000 K),and increases with velocity offset from the bulk of the gas,in addition to the gas excitation state. The highest velocitycomponent, while evident for moderately-ionized species,of temperature approximately 20 000 K, is associated withhighly-excited, highly-ionized hot gas of temperatures upto several 10 K.
5. The cool gas and its traces of warm gas
The bulk of the gas detected along the line of sight, is cooldiffuse molecular gas at low velocities. We comment on itsproperties within the framework of diffuse clouds studies.Diffuse molecular clouds are often modeled as extended,homogeneous, low-density structures with properties regu-lated by the progressive extinction of the UV field by dust.In a companion paper (Nehm´e et al. 2007), we show thatmost observations of the cold, low-velocity gas, namely theH abundance and ortho-to-para ratio, the CO abundance,the C i abundance and excitation, and the 158 µ m C ii lineemission fit with model computations for a diffuse molecu-lar cloud of a gas density of 80 cm − , illuminated by a UVradiation field close to the mean Solar Neighborhood value.The gas temperature, derived from the balance betweencooling and heating processes, ranges from 60 to 80 K.Combining the gas density and gas temperature we calcu-late a thermal pressure, P/k ∼ − K, which is withinthe range of values measured by Jenkins (2002), for cloudswithin the local Bubble. While the physical conditions in-ferred from the model probably apply correctly to the bulkof the gas, we hereafter discuss two points that show thelimitations of the model’s simplicity.The first point is the apparent kinematic separationof molecular and atomic species in the cool gas, with thedistribution of atomic species and CH + over two velocitycomponents A and B, and the molecules CO and CH ob-served at a single intermediate value of velocity. As noted inSection 4.2, we believe that these components are relatedand therefore treat them as a single cloud in the model-ing. However, the observed velocity structure of the cloudcannot be accounted for by the photon-induced chemistrymodel where most of the C i , CO and H are predicted tooccupy the same volume (see figures 3 and 4 in Nehm´e etal. 2007).The second point is the existence of traces of warmmolecular gas, as indicated by a CH + absorption line, and apopulation of J > rotational levels that exceed modelvalues by almost two orders of magnitude. Falgarone etal. (2005) determined an average Galactic-fraction of warmH , in cool diffuse gas, expressed as the ratio of the col-umn density of H molecules in levels J > N (H ∗ )per magnitude of gas sampled: N (H ∗ ) /A V = 4 × cm − . The warm H predicted along the line of sight toHD 102065, with this average Galactic fraction, would thenequal N (H ∗ ) = 2 . × cm − , compared to the range 1-3 × cm − given in Table 1. The ratio of the CH + tothe total hydrogen column densities towards HD 102065 N (CH + ) /N H = 1 . × − , is comparable to the valuescalculated in previous observations by Crane, Lambert &Scheffer (1995) and Gredel (1997). The formation of CH + in diffuse gas involves the endothermic reaction C + + H → CH + + H (∆ E/k = 4640 K), and has to be triggered bydeposition of supra-thermal energy either in MHD shocks ehm´e et al.: Multi-wavelength observations of a multi-phase cloud 13 (Flower & Pineau des Forˆets 1998), or in large velocityshears (Joulain et al. 1998).The chemical patterns triggered by the local depositionof non-thermal energy involve several other endothermicreactions, including CH destruction, which produces neu-tral carbon. According to the models explored in Joulainet al. (1998), the column density of C i produced inthe warm chemistry N (C w ), scales with that of CH + as N (C w ) ≈ N (CH + ) or N (C w ) ≈ × cm − . This cor-responds to less than 10% of the total C i column den-sity N (C) = 5 . × cm − detected in components Aand B (see Table 5). The domain of parameters influenc-ing the out-of-equilibrium chemistry is difficult to explore,and this is only an estimate provided by existing models.The production of CO is enhanced by the warm chemistrypatterns, due to enhanced production in the warm gas ofOH, HCO + and CH +3 , all chemical precursors of CO. It ispossible that the observed heterogeneity of CO abundancesdiscussed above, originates in such processes.
6. HD 102065 gas components within the localbubble context
In this section, we place the HD 102065 observations ina broader context by relating them to the interaction be-tween the nearby interstellar medium, and the Sco-Cen OBassociation, the Local and the Loop I Bubbles.Based on an extensive study of colour excess versus dis-tance for a large sample of stars within 294 ◦ < l < ◦ and − ◦ < b < ◦ , Corradi et al. (1997) concluded thatthe local density of matter is low until an extended in-terstellar dust feature at 150 ±
30 pc from the Sun associ-ated with the Chamaeleon, Musca and Southern Coalsackdark clouds. The same authors complemented their in-vestigation with NaI absorption spectra. Absorption linesat negative velocities are detected towards many stars in-cluding the stars closest to the Sun (Corradi et al. 2004),which is indicative of a low column density (N H ∼ a few10 cm − ), nearby ( d <
60 pc) sheet of gas, outflowingfrom the Scorpius-Centaurus stars with mean radial ve-locity V
LSR = − − . The presence of such a sheetwas first proposed by Egger and Aschenbach (1995) to ac-count for the soft X-ray shadow seen in the ROSAT All SkySurvey maps.In the northern sky, the Loop I shell extends up to NorthPolar radio spur. In the southern sky, a broad H i filamentat b = − ◦ and V LSR ∼ −
20 km s − marks the outer ex-tension of gas swept up by the 10-15 Myr old LCC andUCL subgroups (de Geus 1992).In the Argentine H i survey (Kalberla et al. 2005and Bajaja et al. 2005), H i gas at velocities V LSR < −
20 km s − , down to −
50 km s − , is visible over 60 ◦ in lon-gitude about the Sco-Cen sub-groups (see Figure 9). The se-lection in radial velocity focusses on gas streaming towardsthe Sun at the center of Loop I. We note also that much ofthe negative velocity gas close to the Upper-Scorpius (US)group that still contains O stars must be photo-ionized andthus not seen in H i . The H i survey data show that neg-ative velocity gas is clumped and scattered over a widerange of negative velocities, which cannot be accounted forwithin the simple picture of an expanding shell about a lo-calized group of massive stars. We propose instead that weare looking at dispersed fragments of the parent cloud of the LCC and UCL stars accelerated to a range of velocitiesby supernovae blast waves. Since the stellar groups havemoved with respect to the Sun (de Zeeuw et al. 1999), theexplosions have occured at different locations and distancesfrom the Sun fragmenting and spreading the stars’ parentcloud between the Local and Loop I bubbles. This viewaccounts for the stream of negative-velocity gas observedclose to the Sun, including the local cloud about the Sun,detected at negative velocities, streaming from the generaldirection of the Sco-Cen association (Ferlet 1999). Our viewof the relation between the nearby interstellar medium andthe Sco-Cen stars, the Local and Loop I bubbles buildson the original scenario proposed by Breitschwerdt et al.(2000). However, the observations suggest a more complexpicture where there is not a continuous sheet of swept-upgas separating the two bubbles, but instead clouds of coldand warm gas spread out between the two bubbles. Figure 9. H i emission at negative velocities − < V LSR < −
20 km s − in direction of the Sco-Cen asso-ciation. H i observations from Bajaja et al. (2005) andKalberla et al. (2005). The position of HD 102065 and theUS (Upper-Scorpius), UCL (Upper Centaurus Lupus) andLCC (Lower Centaurus Crux) sub-groups of the OB associ-ation are marked. The black area indicates the region wherethe selected range of negative velocities overlaps with theGalactic rotational velocities of distant gas. The electronicversion is in color.We believe that the diffuse molecular cloud in the fore-ground of HD 102065, is such a cloud. As discussed byMizuno et al. (2001), there is strong evidence that Dcld300.2-16.9 (the Blue Cloud) is not part of the Chamaeleonmolecular clouds complex. A distance of 70 ±
15 pc fromthe Sun, is inferred from the parallax of the T-Tauri starT Cha, which coincides in position and velocity (Covino et al. 1997) with the cloud CO emission. This places thecloud at a distance of ∼
30 pc from the center of the LCCsub-group (de Geus 1992). Its cometary structure, with atail extending in the opposite direction to that of the LCCstars, suggests that it could have been shaped by a super-nova explosion from one of the stars. With a total mass of ∼
100 M ⊙ (including the tail, as calculated by Boulangeret al. (1998) and corrected assuming a distance of 70 pc),the cloud is too massive to have been globally acceleratedby a distant explosion but its structure suggests that gasablated from the surface by the blast wave is streaming out-wards. In the H i data, the low velocity components A andB are connected to the cloud head but extend in differentdirections along the cloud tail (Figure 6). This may implythat they represent both sides of the gas flow, that theyencompass the cloud core, and that the difference in radialvelocity is the result of a difference in angle between thestreaming direction, and the line of sight. At the surface ofthe gas flow, the CO and CH molecules are photodissoci-ated, explaining why they are not detected in componentsA and B, but instead in a single component C, which wouldcorrespond to the internal part of the gas flow. To accountfor the 3 km s − velocity difference between components Aand B within a reasonable angle difference ( ∼ ◦ ), the gasstreaming velocity must be ∼
10 km s − . The tail length ∼ ◦ (2 . ∼ − × yr.The HD 102065 observations were originally motivated bythe strong mid-IR emission of the foreground cloud inter-preted as an unusually high abundance of small dust grains(Boulanger et al. 1990). This distinct characteristic may berelated to the position of the cloud within the Local Bubble,and its plausible interaction with a supernova blast wave.The interaction between interstellar clouds and shockwaves travelling in their surrounding medium, has been in-vestigated using numerical simulations. Images presentedin Nakamura et al. (2006) help to picture the complexity ofthe interaction that ablates matter from the cloud and gen-erates a turbulent flow streaming away from the cloud. Thefact that the C iv and Si iv absorption lines peak at neg-ative velocities qualitatively agrees with this picture whereone expects the low density gas in the turbulent mixinglayer to be entrained by the flow of hot gas. The discrep-ancy between gas temperature and ionization state evidentfor some of the Mg ii , Si ii ∗ and C ii ∗ observed at a tem-perature 2 × K, implies that dynamical mixing of hotand warm gas, on timescales shorter than that necessary toreach ionization equilibrium, is taking place.
7. Conclusion
High resolution spectroscopic observations are presentedand used to characterize the physical state and velocitystructure of the multiphase interstellar medium observedtowards the nearby star HD 102065. These spectroscopicobservations provide detailed information on the kinemat-ics and physical state of the gas, along the line of sight. – Gas is observed over a wide range of velocities, with coldmolecular gas at low velocities, and warmer and lowerdensity gas at negative velocities. Most of the hydrogencolumn density is in cold diffuse molecular gas at lowLSR velocities. – The spectra show three distinct components at nega-tive velocities as low as V
LSR = −
20 km s − . The ex-citation and temperature of the species detected in thenegative velocity components, increase with the velocityseparation from that of the diffuse cloud. H i emissionmaps show that the negative velocity gas extends over ∼
200 deg . It is fragmented and spreads over velocitiesdown to −
50 km s − . – A striking result of the data analysis is the detection ofgas out of equilibrium in both the low and the negativevelocity gas. At low velocities, we find that the coldmolecular gas is mixed with traces of warmer moleculargas where the observed CH + must be formed and whichis also traced by the observed excess population of H in the J > ii , Si ii ∗ and C ii ∗ ) is observedfar out of ionization equilibrium at a temperature of2 × K.We have set the observational results within the sce-nario proposed by Maiz-Apellaniz (2001) and Bergh¨oferand Breitschwerdt (2002) that connects the origin of theLocal Bubble to supernovae explosions from early starmembers of the oldest Sco-Cen groups. – The gas at low LSR velocities is observed to be asso-ciated with the Blue Cloud (Dcld 300.2-16.9) locatedwithin the local bubble at a distance of ∼ ∼ − × yr) supernova ex-plosion from one of the LCC stars. The line of sightto HD 102065 goes across the cloud tail. The velocitystructure of the low velocity gas may be understood asthe two sides of a flow of gas ablated from the cloud headDcld 300.2-16.9. The large abundance of small dust par-ticles inferred from the IRAS colors of the Blue Cloudtail, possibly results from the cloud interaction with thesupernova blast wave. – The negative velocity gas must have been acceleratedby a supernova explosion that occurred in the Sco-CenOB association. Our interpretation of the Blue Cloudshape would imply that it is physically connected withthe lower velocity gas. The column density of Si ii ∗ at −
20 km s − might be a signature of an ongoing shock.A complete diversity of diffuse interstellar-medium compo-nents, is present along the line of sight. Our analysis of bothlow and negative velocity, hot, warm and cold gas compo-nents, stresses the complex dynamical interplay betweenthe interstellar medium phases, in agreement with the pre-dictions of numerical simulations (Audit and Hennebelle2005, de Avillez and Breitschwerdt 2005). We started thiswork by considering an apparently unusual cloud, but nowbelieve that the questions we are now asking, have signifi-cance for the general study of diffuse, interstellar matter inthe Galaxy. Acknowledgements. ehm´e et al.: Multi-wavelength observations of a multi-phase cloud 15