Optical polarimetry, high–resolution spectroscopy and IR analysis of the Chamaeleon I dark cloud

Astronomy and Astrophysics Supplement Series, Jul 2018

We present optical polarimetry and high resolution spectroscopy of a sample of stars toward the Chamaeleon I dark cloud. We use our polarimetry which includes 33 stars to study the wavelength dependence of the degree and position angle of polarization.
From fits to the normalized wavelength dependence of interstellar polarization, we derive estimates of ranging from 4500 Å to 6700 Å, and PMax ranging from 3 to 8%. The values of were found to be well correlated with the IRAS 100 μm intensity, while PMax was found to increase with EB-V.
High resolution spectra of the Ca II, CH, and CH+ lines were obtained for 10 stars, which show two components of Ca II in absorption at 3.0 < 5 km , and vLSR= -3.0 km and a single strong molecular CH absorption component at 3.0 < vLSR < 5.0 km .
From our data we found, by interpretation of the various correlations between the polarimetry, photometry and IRAS fluxes, the following:
the probable presence of shocked molecular gas; a warm molecular CH component; small dust grains at the edges of the cloud, and larger grains in the central parts, which are causing the polarization.
Our results provide a consistent picture of the gas and dust content in the Cha I region, where larger grains, responsible of the starlight polarization, exist in the center of the cloud, surrounded by envelopes of warmer molecular and atomic material.

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Optical polarimetry, high–resolution spectroscopy and IR analysis of the Chamaeleon I dark cloud

Astron. Astrophys. Suppl. Ser. Optical polarimetry, high{resolution spectroscopy and IR analysis of the Chamaeleon I dark cloud? E. Covino 2 E. Palazzi 1 B.E. Penprase 0 H.E. Schwarz 3 L. Terranegra 2 0 Pomona College, Department of Physics and Astronomy , Claremont, CA , U.S.A 1 ITESRE/C.N.R. Bologna , Italy 2 Osservatorio Astronomico di Capodimonte , Via Moiariello, 16, I 3 Nordic Optical Telescope , Apartado 474, E-38700 Sta. Cruz de La Palma, Canarias , Spain 4 80131 Napoli , Italy We present optical polarimetry and high res- 1. Introduction olution spectroscopy of a sample of stars toward the Chamaeleon I dark cloud. We use our polarimetry which The Chamaeleon I dark cloud ( 11h, −76 deg) includes 33 stars to study the wavelength dependence of is one of the nearest active star formation regions the degree and position angle of polarization. (d 120 − 150 pc). Thanks to its proximity and relaFrom ts to the normalized wavelength dependence tively high galactic latitude (b −16 ), this cloud is well of interstellar polarization, we derive estimates of Max suited for investigating the composition, magnetic eld geranging from 4500 A to 6700 A, and PMax ranging from 3 ometry and IR emission of a star forming cloud. to 8%. The values of Max were found to be well correlated Observations of the Chamaeleon region were reviewed with the IRAS 100 m intensity, while PMax was found to by Schwartz (1991), who provides a particularly complete increase with EB−V . account of the determination of the stellar content, extincHigh resolution spectra of the Ca II, CH, and CH+ tion and distance toward the complex. The Cha I cloud lines were obtained for 10 stars, which show two compo- was initially discovered to harbour an unusual number of nents of Ca II in absorption at 3.0 < vLSR < 5 km s−1, H emitting stars (Henize 1954), and recent observations and vLSR= −3.0 km s−1 and a single strong molecular CH have determined that there are more than 100 associaabsorption component at 3.0 < vLSR < 5.0 km s−1. tion members (Prusti et al. 1991; Gauvin & Strom 1992; From our data we found, by interpretation of the vari- Hartigan 1993). ous correlations between the polarimetry, photometry and Additional Young Stellar Object (YSO) members of IRAS fluxes, the following: Chamaeleon have been found from both X-ray emisthe probable presence of shocked molecular gas; a sion from pointed ROSAT observations (Feigelson et al. warm molecular CH component; small dust grains at the 1993; Zinnecker et al. 1996) and deep IRAS (Assendorp edges of the cloud, and larger grains in the central parts, et al. 1990) and 2 m observations of selected elds which are causing the polarization. (Jones et al. 1985; Prusti et al. 1994). The consensus is Our results provide a consistent picture of the gas and that there are between 70 and 300 YSOs in Chamaeleon, dust content in the Cha I region, where larger grains, re- in a cloud which is approximately 4 degrees across and has sponsible of the starlight polarization, exist in the center a mass of 700 to 1030 solar masses. Toriseva & Mattila of the cloud, surrounded by envelopes of warmer molecular (1985) have estimated the dust to gas ratio in the cloud and atomic material. to be about 0.02 by mass. We should note that the inferred star formation e ciency of the Chamaeleon cloud - like cloud at 140 pc, which extends over several degrees Walborn 1968 ). Measurements of some bright stars were on the sky into the constellation Musca. The maximum also performed with a polaroid inserted in order to deterextinction in the Cha I cloud is thought to be in the range mine the instrumental e ciency. These yielded negligible 4:0 < AV < 6:0, and the reddening of eld stars from corrections. Vrba and Rydgren (1984) behind the cloud ranges from The reduction of the polarimetric data was performed 0:05 < EB−V < 1:27, which completely samples the ex- using a revised version of the MIDAS procedure PISCO. pected extinction in the Cha I cloud. For each star the degree of polarization was determined, In this work we concentrate on the interstellar medium with values ranging from zero up to about 8 percent. The around the Chamaeleon I association, by combining new observational errors are determined from the power in optical polarimetry and optical spectroscopy of atomic the 5th and higher harmonics in the Fourier-transform, and molecular absorption features with IR data. using an algorithm which is part of the PISCO data Recent work on the polarimetry of the Chamaeleon reduction package (Schwarz 1989) . Also, the correction region has also been published (Whittet et al. 1994; for noise biasing, according to the analytical formula McGregor et al. 1994) as has a comparison between IR p = p0[1 − ( =p0)2]1=2 has been applied to the observed and UV emission and extinction (Boulanger et al. 1994). polarization, p0, using the estimated standard deviation Our work seeks to provide a uni ed approach in which of the polarization, , derived from the PISCO reduction the di erent observations provide a complementary pic- procedure (Clarke & Stewart 1986) . ture about the interrelations between dust and gas in the The list of observed stars is given in Table 1, together Chamaeleon cloud. with other relevant information. The optical polarization Where possible we will analyze the connections be- data are given in Table 2. For reference, the rst column tween the observed optical polarization, IR emission and of each table gives the list number from Whittet et al. optical absorption lines for the stars in the Chamaeleon (1987). region. Photometry suggests a very low foreground reddening (EB−V 0:04 is reported by Whittet et al. 1987, and an even smaller value of EB−V 0:007 is found by 2.2. High resolution spectroscopy Franco 1991) , which insures that optical absorption lines and the polarization of background stars of moderate distances are due to the Cha I dark cloud, and are not signi cantly a ected by foreground absorption. The background stars earlier than A0, and suitable for the spectroscopic study of the cloud were selected from the polarimetry sample. In addition, we included the stars HD 97300 and HD 97048, which are thought to be members of the Cha I association, to improve the completeness 2. Observations of our spectroscopic sample. 2.1. Optical polarimetry High resolution spectroscopy was done with the Coude Auxiliary Telescope (CAT) of the European Southern Our sample of 33 background stars was selected from the Observatory at Cerro La Silla (Chile) during two observlist of eld stars provided by Whittet et al. (1987) . Linear ing runs in 1992 March and 1994 January. The data were polarization measurements were carried out in 1992 March obtained using an RCA CCD (ESO # 9) and the CES 13 to 17 at the ESO/MPI 2.2 m telescope equipped with short camera in the blue path. This instrument con guthe two{channel photo{polarimeter PISCO (Stahl et al. ration provided a resolving power, = > 70 000 at the 1986; Schwarz 1989) . Measurements were obtained in the studied wavelengths. A few additional spectra were obCousins U BV RI system using a diaphragm of 15 arc- tained in 1994 July using the CAT with the CES long sec. The instrumental polarization compensation mode, camera, operated in remote control from Garching. The involving a rotation of a compensating phase plate unit July data have a higher resolving power = = 130 000, by 180 degrees, was systematically used. The data acqui- and were used to complete coverage of the region, and sition was performed using the two{apertures mode for obtain more detailed velocity pro les for selected stars. sky polarization compensation (the two apertures being Data reduction was performed using the IRAF system, separated by 6600 in the East{West direction). Calibration in which the standard CCD reduction packages were used in polarization, P , and position angle, , was done through to remove cosmic rays, perform bias subtraction and flatobservations of the polarization standard star HD 147084 eld the images. Spectral extraction was done using the (o Sco) (Hsu & Breger 1982) . optimal extraction algorithm in the IRAF task apsum, The uncertainty in the position angle in all cases was in which pixel values across an order of a spectrum are estimated to be less than 2 degrees, using polarized stan- weighted according to the calculated signal to noise radard stars and the internal calibration of the polarime- tio. The signal to noise was estimated using the readout ter. The instrumental polarization was measured by ob- noise and gain of the CCD. Spectra were wavelength calserving the unpolarized standard stars HD 100623, HD ibrated using a Thorium/Argon calibration lamp, and a 115617 (61 Vir) and HR 6060 (18 Sco) ( Serkowski 1974 ; linear interpolation of matched features. For each spectrum a t to the continuum was performed using a high order spline, and a correction was applied to transform the spectra to a uniform LSR velocity reference frame, by removing the velocity contributed by terrestrial and solar motions. Column densities were computed by using a Voigt pro le model, based on a version of the STARLINK curve of growth programme adapted for IRAF. For some of the spectra, line pro les were modeled using the MIDAS image processing package. Each line pro le was modeled using the minimum number of components necessary to reproduce the observed data with residual values comparable to the noise in the data. Our spectral resolution limited us to separating line components with velocity di erences of more than 1 − 2 km s−1, and we cannot rule out the possibility of very narrow blended components. However, since the turbulent velocities in molecular clouds are expected to be 1 − 2 km s−1, we believe our instrumental resolution is able to resolve the components present in the Chamaeleon cloud. This is also consistent with the determinations based on CO observations in the Cha I cloud obtained by Dubath et al. (1995) . 2.3. IR data analysis For the Chamaeleon region we obtained the latest coadded plates for all four bands of the ISSA survey. The pixel size within the plates is 1:50 which, in the case of the 100 m band, oversamples the actual instrumental resolution of the satellite, which is closer to 40. We removed zodiacal background residuals from each plate by using an iterative background removal programme which selects background points based on an analysis of the histogram of pixel values, and ts a flat background to the selected points. IRAS colors were derived using the programme skyview from IPAC, which averages pixel fluxes in a 9 pixel region centered on the coordinates of our program stars. The larger sampling region makes that the IRAS colours have a higher signal to noise ratio, and samples a region 4:500 square, with uncertainties in the IRAS flux, F , of < 0.08 mJy Sr−1 for 12 m and 25 m, and < 0.15 mJy Sr−1 for the 60 m and 100 m passbands. 3. Results and discussion 3.1. Polarimetry Table 1 presents the stellar parameters for our sample. Earlier photometric and polarimetric work in the Chamaeleon region has analyzed the stellar content of both member and eld stars (Whittet et al. 1987) , and a search of the SIMBAD database and reading of literature has produced the notes included in Table 1. Important in the discussion of any polarimetry is the possibility of intrinsic polarization from the source. We chose to observe stars which were optically determined to be eld stars, which are expected to have little or no intrinsic polarization. Three of the stars in our sample are IRAS point sources, which is probably a result of the late spectral type and bright magnitude. Recent pointed ROSAT observations (Feigelson et al. 1993; Zinnecker et al. 1996) , however, have identi ed four new YSOs which are associated with our \ eld" stars, these are indicated in Table 1. Two of them are both IRAS point sources and X-ray sources. The polarimetry and the IR emission of these stars will probably be influenced by circumstellar material. Figure 1 shows an IRAS 100 image of the Chamaeleon I cloud, with the locations of our program stars indicated. Measurements of optical linear polarization of starlight can be used to trace the geometry of the interstellar magnetic eld component, B?, in the plane of the sky. The observed polarization is produced by the di erential extinction by non-spherical dust grains associated with the interstellar clouds along the line of sight. It is assumed in fact that elongated paramagnetic dust grains are aligned by the local interstellar magnetic eld, via the Davis{ Greenstein mechanism with their short axes parallel to the eld direction. Strong evidence that extinction by dust grains is responsible for the observed polarization comes from the existence of a correlation between the degree of polarization and the amount of reddening (Spitzer 1978) . In Fig. 2 we present a plot for the Chamaeleon I cloud, with the observed polarization vectors for our program stars, superimposed onto a plot of IRAS 100 m contours, with equatorial coordinates labeled. The magnetic eld vectors would be expected to follow the polarization vectors, and they are approximately parallel to the galactic plane. The dispersion in the values of the polarization angle is very small, typically 2 degrees or less. The complete results of the polarization measurements are presented in Table 2 for our sample of stars in the ve passbands. We have adopted the parameterisation of the wavelength dependence of the linear polarization, referred to as Serkowski's model (Coyne et al. 1974), P ( )=PMax = exp[−K ln2( Max= )] where Max is the wavelength at which the maximum polarization, PMax, is observed, and K 1:15. The value of Max for our sample ranged from 0.35 to 0.9 m and typically was found to be near Max = 0.55 m. Larger values of Max are expected in direction of dense clouds, indicating an increase in the mean grain size. A discussion of the theoretical arguments behind the Serkowski polarization model may be found in Spitzer (1978) . For selected stars we have tted the percentage polarization with a Serkowski polarization model to derive values of Max and PMax. The data points, and their best ts are presented in Fig. 3. The Serkowski model was tted only to those stars which had accurate polarization measurements in at least four passbands. We excluded measurements from our t which had polarization uncertainties P > 1:6%. To check the e ect of tting the Serkowski model with four points, we also did all the ts using both 5 and 4 1.5 1 0.5 1.5 0.5 1.5 1 1 0.5 1.5 0.5 1 0 1 2 0 1 2 0 1 2 points and, typically, the values of Max derived in the two ways do not di er by more than 100 A. For 19 stars we have tted the Serkowski polarization model with 5 points, while 9 stars were done with 4 points, since the late spectral types of the stars prevented accurate measurements of U band polarization. Seven stars were not tted at all, since both the U and B polarization were found to be inaccurate, or the star was found to be unpolarized. In Fig. 4 we present plots of the percentage polarization for each of the ve bands against the color excess EB−V . It is clear that the polarization rises for EB−V < 0:5 at a di erent slope than for the largest values of EB−V > 0:5. A similar change in slope has been observed by Whittet et al. (1991) , and it is believed to reflect a decrease in the e ciency of polarization at the highest extinctions due to either more spherical grains, or de{alignment of the grains in the densest parts of the sight lines. Figure 5 shows the maximum polarization PMax of each star, plotted against the color excess EB−V . Also shown in Fig. 5 are separate ts for the sight lines where 4 or 5 points were used to t the data to the Serkowski model. We see a gradually increasing function, with a fair amount of dispersion in the measured values of PMax. From the parameterisation of the polarization wavelength dependence it is also possible to determine Max, and we examined the behaviour of Max as a function of extinction. The values of Max increase with EB−V and a plot is presented in Fig. 6 for our sample. The increase in 0.5 1.5 Fig. 3. b) Same as in Fig. 3a 10 5 10 5 10 5 10 5 10 5 0 0.5 1 E(B-V) Fig. 4. Polarization percentage for each of the ve passbands for the polarimetric sample. Included are the best ts to the data (dashed line) and \maximum" polarization, (P=EB−V )Max = 9:0% mag−1, (solid line) from Spitzer (1978) . The di erent slope between high and low extinction sightlines is apparent 0.5 1 1.5 E(B-V) Max at large EB−V is consistent with larger dust grains occurring in denser sight lines with high extinction, perhaps as the result of ices accumulating on the grains. In Fig. 7 we present a plot of the distribution of Max, which appears to be bimodal, with a smaller population of points having larger values of Max. It appears that most of the large values of Max also coincide with the peak 100 m emission from the cloud, based on an examination of the location of the program stars on the IRAS map of Fig. 1. Variations in Max are generally attributed to variations in the mean grain size toward the region. Some complications can arise, however, in regions with multiple clouds along the line of sight, which may have di erent values of Max, or di erent orientations of the mean magnetic eld. In the former case, signi cant deviations from the Serkowski model may be found, while in the latter case, the electric eld vector orientation may be a function of wavelength. Clarke and Al-Roubaie (1984) have modeled the e ects of various grain size distributions, and speci cally the rotation of the polarization vector I − B with the angle between two clouds along the sight line. Additionally a linear relation exists between the width parameter K and Max. The physical signi cance of the parameter K from the Serkowski model has been examined by Whittet et al. (1991) , but is presently uncertain whether K varies signi cantly with direction within the Galaxy. We have assumed a value of 1.15 for the Serkowski K parameter, and also that most of the polarization Fig. 6. Wavelength of maximum polarization, Max versus EB−V . The dotted line represents the least square ts to all data points, while the solid line gives the least square ts for only values of Max derived from more than three points ts to the Serkowski's law ( lled circles) 5000 6000 7000 results from a single component. The spectroscopic results of Sect. 3.3 suggest that there are two components of the ISM in the line of sight to the Cha I association. The second component, which appears in several of the Ca II spectra, is much weaker, however, and therefore our assumption of a single dominant component is valid. Our observed rotation of the polarization angles for a given sight line varied by less than 2o for nearly all sight lines, which is also consistent with a single dominant component of polarizing interstellar medium (Clarke & Al-Roubaie 1984) . 3.2. Analysis of IRAS infrared colors Infrared observations of nearby molecular clouds and atomic cirrus have also been used to deduce an increased small grain content from increased R(12; 100), and decreased I100 near the edges of clouds, which undergo rapid fluctuations of the number of small grains attributed variously to condensations from gas, removal from larger grains, and from slight redistributions between small (a < 30 A) and mid-sized (30 < a < 100 A) grains (Desert et al. 1990; Bernard & Boulanger 1993) . We have analyzed the values of R(12; 100) and other IRAS colours, and present the results for our program star sight lines in Table 3. The IRAS colours were also examined for systematic trends with some of our other polarimetric and spectroscopically observed quantities. Figure 8 presents a plot of percentage polarization against R(12; 100). A clear decrease in polarization is seen as R(12; 100) increases, suggesting that the processes which disrupt the grains in the cloud to produce the smaller particles, also are capable of either realigning or destroying the larger grains responsible for optical polarization. The fact that in Fig. 9 (top) we see little or no correlation between Max and R(12; 100) suggests that the process of creating the small grain population minimally a ects the grain size distribution of the polarizing grains, and therefore the small grains producing the enhanced values of R(12; 100) are probably an independent population from the polarizing grains. 5 15 10 5 10 5 10 5 10 5 0 0.1 0.2 R(12,100) 0.3 Bernard & Boulanger (1993) have modeled the vari ous contributions of PAH and other small molecules to the small grain content, and Desert et al. (1990) have determined expected IR colours with mixtures of PAH's and small grains. Boulanger et al. (1994) have determined that variations in the far UV rise of the extinction curves for stars in Chamaeleon I and II were not correlated with 12 and 25 m emission, which they interpreted as proof that the R(12; 100) fluctuations were caused by slight variations in the size distributions of mid sized and small grains having the same composition as the larger grains. Our results are consistent with this model, since only a slight trend is seen between R(12; 100) and Max which is plot ted in Fig. 9. We do see a strong correlation between I100 and Max, however, and a slight anticorrelation between R(60; 100) and Max. These results suggest that the larger grains associated with increased values of Max also have reduced grain temperatures and larger 100 m fluxes. We have also made plots of the IRAS fluxes and colours against the maximum percentage polarization Pmax, and these are presented in Figs. 10a and 10b. A de nite correlation between I100 and Pmax is seen, as is an anticorrelation between R(12; 100) and Pmax. These results strongly suggest that the larger grains which have increased values of I100 and Max also are responsible for large amounts of polarization. Colour-colour plots of R(12; 100) against I100 (Fig. 11) reinforce the trends shown in Figs. 10a and 10b, as does the plot of R(60; 100) vs. R(12; 100) (Fig. 11, bottom), which suggests that the larger values of R(60; 100), usually associated with heated grains, also give rise to larger numbers of small particles which are responsible for the increase of R(12; 100). 0.0514 0.0724 0.0383 0.0120 0.0252 0.0054 0.0573 0.0392 0.0398 0.0102 0.0386 { 0.074 0.0087 0.051 0.0077 WCaI (A) 3.3. Spectroscopic results 4. The Ca II absorption is quite uniform across the cloud in velocity, although in some sight lines a weak second The transitions from large to small grains are generally at- component of absorption appears. The second component tributed to variations in either the ambient UV radiation is strongest in the spectra at either edge of the Cha I eld or convective turbulent motions within the interstel- cloud, suggesting that this component may arise from eilar cloud. Both processes would be expected to a ect the ther a neighboring cloud, or higher velocity gas which is depletion of ices from the surfaces of grains, and to release impacting the Cha I region. highly depleted elements such as Ca into the gas phase. The observation of Ca II in absorption from background High resolution spectra of CH absorption are presented stars should therefore provide a useful probe of grain de- for the sample of stars in Fig. 13. The tted line pro les are struction or ablation from radiation or kinetic processes. included, and the values of vlsr, b, and N (CH) are also preWe present in Fig. 12 the absorption lines of Ca II for our sented in Table 4. The molecular content of clouds is gensample, in order of Right Ascension of the star. Line pro- erally found to correlate with 100 m emission, since the le models for all the spectra have been computed, and larger grains which emit 100 m are e ective in shielding the values of vlsr, b, and N (Ca II) are presented in Table the molecules from the dissociating interstellar radiation 0.5 eld. However, if transient processes exist in molecular clouds which heat small grains and can produce molecular CH, then it should be possible to see trends between N (CH) with EB−V and R(60; 100). Recent observations have found signi cant amounts of CH in the warm envelopes of molecular clouds, suggesting that both CH and CH+ may exist as a transient phase of molecular gas in some molecular clouds (Crane et al. 1995) . Figures 14a-d present stacked spectra for single sight lines, and it is clear that the CH/CH+ ratio is highly variable in the Cha I cloud. In the sections below, we discuss the data for each of the four Cha I sight lines, detailing the relationships between IRAS and spectroscopic results. Fig. 12. b) edge of the Cha I cloud, where the IRAS 100 m flux is I100 = 5:87 mJy Sr−1. The observed Ca II absorption pro le has been tted by two components giving a total column density N (Ca II) = 8:5 1011 cm−2. The molecular content for this sight line is high with N (CH) = 3:0 1013 cm−2 and N (CH+) = 1:9 1013 cm−2. The CH+ pro le was tted by a single component and appears asymmetric, which may be the result of two blended components. The extinction for the sight line is only EB−V = 0:30, and the column density of CH is one of the highest per unit of extinction, as seen in Fig. 15. In low extinction sight lines, like the one toward HD 99759, the dissociation rate of CH from the interstellar radiation eld must be balanced by an increased production of CH to account for the large molecular column densities. A production mechanism which could be attributed to shocks which collide with the Cha I cloud would be most pronounced near the edges of the dark cloud region, such as the material toward HD 99759. The large value of N (CH+)=N (CH) = 0.63 is consistent with this hypothesis, since CH+ has been thought to be produced in shocks (Allen 1994) . Figure 14b shows a stacked plot of the spectra for the HD 97048 sight line, which appears toward the Southern Fig. 14. b) Stacked spectra for the HD 97048 sight line. The absorption lines are discussed in Sect. 3.3.2. Fig. 14. d) Stacked spectra for the HD 97300 sight line. The absorption lines are discussed in Sect. 3.3.4. central part of the Cha I cloud, where the IRAS 100 m flux is I100 = 250:1 mJy Sr−1. HD 97048 is thought to be HD 97048 of AV = 1:75, which is less than half of the embedded in the Cha I cloud, and the large IRAS fluxes estimated maximum extinction of AV = 5:0 for the Cha I therefore come from circumstellar material being heated cloud. The HD 97048 sight line would therefore be on the by the star. HD 97048 is associated with the reflection near side of the Cha I cloud, within a reflection nebula nebula Ced 111, and has been found to be surrounded by which consists of a mixture of radiation processed gas and additional IRAS sources, which suggests that HD 97048 dust. The CH column density for HD 97048 is slightly is a center for low-mass star formation. Recent reviews on higher than the HD 99759, at N (CH) = 3.4 1013 cm−2, the HD 97048 and HD 97300 sight lines include Assendorp which is surprising considering the large radiation eld et al. (1990), and Steenman & The (1989) . The latter have which must be present from the embedded source found that the two stars have anomalous extinction, with HD 97048. The CH+ column density for this sight line RV = 5:0, which would result in an optical extinction for is extremely low, with N (CH+) < 1:5 1012 cm−2. less favourable environment for the production of CH+. The value of the total to selective extinction ratio calculated from the wavelength dependence of the polarization is RV = 3:3, although the value of max suggested from the polarimetry is max = 5700 A, which is average for the polarimetry sample. The IRAS color index R(12; 100) = 0:04 is very small, suggesting a reduced population of heated small grains, and the value of R(60; 100) = 0:21 corresponds to one of the lowest grain temperatures of the sample. 3.3.4. HD 97300 Figure 14d shows a stacked plot of the spectra for the Fig. 15. CH column density, N (CH), is plotted against HD 97300 sight line, which appears slightly South of EB−V for spectroscopic program stars, showing enhanced HD 96675 but still in the central part of the cloud. Like N (CH)/EB−V content for sightlines toward HD 94414, HD 97048, HD 97300 is believed to be embedded in the HD 97300, HD 96675, and HD 99759 Cha I dark cloud, and therefore the IRAS colours are from heated circumstellar material. The IRAS 100 m flux is It is interesting to note that two sight lines from reported to be I100 =282 mJy Sr−1. The CH+ column the same cloud, HD 97048 and HD 99759, have nearly density for this sight line is higher than for HD 97048, with identical colour excesses of EB−V = 0:3, and yet N (CH+) = 3:3 1012 cm−2, and N (CH+)=N (CH) = 0.08, have extremely di erent ratios of N (CH+)/N (CH), with which is at least twice as high as the HD 97048 sight value of N (CH+)=N (CH) < 0:043 for HD 97048 and line, but still relatively low compared with other molecN (CH+)=N (CH) = 0:63 for HD 99759, respectively. This ular absorption sight lines. The Ca II column density is di erence may be due to the fact that HD 99759 is prefer- substantially higher than for the HD 97048 sight line, with entially sampling the edge of the cloud, where our results N (Ca II) = 8:1 1011 cm−2. The column densities of CH+ strongly suggest CH+ production is enhanced. and Ca II therefore seem to be correlated, at least for our The Ca II column density toward HD 97048 is large, limited spectroscopic sample, which again suggests that with N (Ca II) = 6.7 1011 cm−2, from a single component the CH+ appears in the warmer outer envelopes of molecof absorption. The HD 97048 Ca II absorption is more ular clouds. symmetric than the HD 99759 Ca II absorption, which is also consistent with the enhanced CH+ production toward 3.3.5. Variation of Column Densities with EB−V and HD 99759 resulting from evaporation or shock processing IRAS colours of gas at the edge of the Cha I dark cloud, which might introduce the extra component of Ca II absorption. Figure 15 is a plot of N (CH) vs. EB−V and we see a de nite increase in N (CH) with increased extinction. Also notable is the very large value of N (CH)/EB−V for the stars 3.3.3. HD 96675 HD 94414, HD 97300, HD 96675, and HD 99759. Figure 14c shows a stacked plot of the spectra for the Variations in N (CH)/EB−V have been observed preHD 96675 sight line, which appears toward a Northern viously in high galactic latitude molecular clouds and section of the Cha I cloud, where the IRAS 100 m flux is may trace shock formation of molecules (Penprase 1993; I100 = 9:2 mJy Sr−1. The Ca II column density toward Penprase et al. 1990) . HD 96675 is moderate, with N (Ca II) = 4:5 1011 cm−2, One of the more interesting results from the comparidistributed over two components as for HD 99759. The ex- son of spectroscopic, polarimetric and IRAS data involved tinction is again similar to that of the previous two sight the di ering behaviour of N (Ca II) and N (CH) with IRAS lines, with EB−V = 0:31, and with a substantial CH col- colours and max. Figure 16 presents a plot of the column umn density of N (CH) = 3.1 1013 cm−2, one of the high- densities N (CH) and N (Ca II) against R(60,100), which is est per unit of extinction, as seen in Fig. 15. The CH+ col- commonly considered to be a good indicator of grain temumn density is very weak, with N (CH+) < 3:2 1012 cm−2. perature. Values of N (CH) are seen to rise steadily with It turns out that the HD 96675 sight line has a large molec- R(60; 100), suggesting strongly that the warmer grains ular content, yet a small ratio of N (CH+)=N (CH) and are somehow producing additional CH. At the same time, N (Ca II)/N (CH). The high value of I100 =EB−V may hint Ca II appears to decrease with grain temperature, which at a di erent value of RV for this sight line, or may re- is surprising since much of the gas phase Ca II is expected flect a substantial population of large grains which would to result from the disruption (and therefore heating) of shield the CH from dissociation, but perhaps provide a grains at the edges of clouds. Another interesting relation is found be tween N (CH)/I100 and max, which is shown in Fig. 17. N (CH)/I100 decreases with max, while N (Ca II)/I100 is unchanged. One interpretation of this result is that the CH production may be more favourable on smaller grains, which have decreased values of max. Further observations are needed to test this possibility. 4. Conclusions and summary We have found many consistent relationships between the grain and gas phase diagnostics of the interstellar medium of the Cha I cloud. The warmer interstellar material, based on the IRAS colour index R(60; 100) has more molecular CH, while CH+ shows up toward sight lines at cloud edges with enhanced values of R(12; 100) and strong Ca II absorption. The large column densities of CH and the high values of N (CH)/EB−V for the stars in our sample suggest a possible additional production mechanism of CH which outpaces the photodissociation of CH at cloud edges. The Ca II line pro les toward Cha I show a consistent velocity structure, where a single strong component at vLSR = 4:5 km s−1 is seen along all sight lines. Several sight lines exhibit a weak component of Ca II at vLSR = −1:0 km s−1. Our observations suggest that the Ca II absorption lines trace a single dense cloud and an Fig. 17. Plots of normalized densities N (CH)/I100 and N (Ca II)/I100 against Max, derived from the polarimetry. The Max, values are considered an indicator of grain sizes, and therefore it appears that CH is produced more rapidly in regions with smaller mean grain sizes additional warmer cloudlet with a di erent velocity. The detections of the strongest CH+ absorptions appear in sight lines which have the additional Ca II component, suggesting that a possible kinematic disruption of the Cha I cloud plays an important role in CH+ production. Additionally, we observed that the polarization of the stars was less e cient at higher extinctions with EB−V > 0:5, which we attribute to either more spherical dust grains or increased de-alignment of the grains in the cores of the dense, high extinction sight lines. The IR diagnostics of R(12; 100) and R(60; 100) were consistent with our observations of polarization and molecular absorption. We found a strong anticorrelation between R(12; 100) and the maximum percentage of polarization PMax, suggesting that the processes responsible for the production of the small grain population with strong 12 m emission is also responsible for disrupting the alignment of the grains in the clouds. A lack of correlation between R(12; 100) and max suggests that if small grains are responsible for the enhanced values of R(12; 100), they are a population which is separate from those responsible for most of the polarization of starlight. If the small grains do arise from the larger grains of the Cha I sight lines as Boulanger & Gry (1994) have suggested, they must evaporate or split from the larger polarizing grains non-destructively in order to maintain the size distributions which indicate large average grain sizes. Acknowledgements. We wish to thank the European Southern Observatory for the allocation of time for this project. 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E. Covino, E. Palazzi, B. E. Penprase, H. E. Schwarz, L. Terranegra. Optical polarimetry, high–resolution spectroscopy and IR analysis of the Chamaeleon I dark cloud, Astronomy and Astrophysics Supplement Series, 95-109, DOI: 10.1051/aas:1997295