SMALL-SCALE STRUCTURE OF THE INTERSTELLAR MEDIUM TOWARD ρ Oph STARS: DIFFUSE BAND OBSERVATIONS

Figure 2 shows the reduced (non-telluric-corrected) spectra (I_λ versus λ) in the wavelength region covering the λλ5780 and 5797 DIBs toward ρ Oph A, B, C, and DE. The lower three panels show the ratios I_λ(B)/I_λ(A), I_λ(C)/I_λ(A), and I_λ(DE)/I_λ(A). The λλ5780 and 5797 DIBs are stronger, by 5% and 9%, respectively, toward B than A, as indicated by the residuals in the I_λ(B)/I_λ(A) plot that precisely match the positions and profiles of the respective DIBs. These two bands are also stronger toward DE than A, but in this case λ5780 exhibits a much greater change in strength of ∼20%, the change for λ5797 being only c. 7%. This is a notable finding given that the values of E_{B − V} toward ρ Oph A and ρ Oph DE are almost identical with values of 0.47 and 0.48 (Pan et al. 2004). Despite contamination of the ρ Oph C spectrum, it is apparent that both λ5780 and λ5797 are stronger toward C than A.

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Figure 3 (top panel) contains the corresponding data for λ6614 which again show this DIB to be slightly stronger toward ρ Oph B than A, and much stronger toward DE than A, in a similar fashion to λλ5797 and 5780. A significant decrease in λ6614 band strength is seen toward ρ Oph C relative to ρ Oph A. The λ6614 spectra appear to show some evidence of profile variation which may result from different carrier internal temperatures (Cami et al. 2004) or isotopic enrichments (Webster 1996).

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Measurements of the equivalent widths for λλ5780, 5797, 6614, and seven other diffuse bands are given in Table 2. Broader DIBs are omitted due to the difficulty in accurately defining their continua in échelle spectra and the increased likelihood of overlapping stellar features. The results presented here for the yellow-to-red DIBs (obtained in 2005) match closely the results obtained by Cordiner et al. (2006), who used only the 2004 data (which had slightly lower signal to noise, but was observed with better seeing), and focused on DIBs at wavelengths >5500 Å. The fact that the spectra obtained on these two epochs match so closely despite the different instrumental setups and observing conditions adds confidence to the results presented here.

Table 2. ρ Oph Diffuse Interstellar Band and Atomic/Molecular Line Measurements

	A	B	C	DE	Error
Equivalent widths (mÅ)
λ4726	99.3	97.8	112.9	67.4	4.3
λ4964	23.5	22.0	21.5	19.5	0.8
λ5780	188	197	[199]	227	2.8
λ5797	49.7	54.2	[56.0]	55.2	1.8
λ5850	29.8	31.9	35.2	25.2	1.8
λ6196	15.9	15.6	15.0	18.8	2.0
λ6203	42.0	42.5	43.0	47.2	6.0
λ6376	14.8	15.8	15.5	21.0	1.3
λ6379	27.1	28.3	28.5	34.7	1.3
λ6614	63.9	67.7	62.1	84.9	1.3
K i λ7699	88.1	89.5	77.7	90.5	0.3
Column densities and derived values
$N({\rm H\,\textsc {i}}$)/1021	4.3	...	...	4.9	1
N(H2)/1020	3.7	...	...	3.0	1
$N({\rm K\,\textsc {i}}$)/1012	1.03	1.09	0.80	1.10	0.05
$N({\rm Ca\,\textsc {i}}$)/1010	1.74	1.69	1.62	2.02	0.09
$N({\rm Ca\,\textsc {ii}}$)/1012	1.75	1.96	1.92	1.94	0.1
N(CH)/1013	2.37	2.33	1.92	2.18	0.1
N(CH+)/1013	1.59	1.43	0.70	0.75	0.06
N(CN)/1013	0.21	0.20	0.60	0.21	0.01
N(C2)/1013	4.1	2.9	3.3	4.0	1
Trot(C2)/K	51	49	46	56	1
n/cm−3	625	450	1100	425	...
Depletion (F*)	1.09	...	...	0.88	0.07

Notes. Upper part of this table gives equivalent widths of interstellar features measured toward ρ Oph stars; values in brackets indicate probable stellar photospheric line contamination. The lower part of this table gives atomic and molecular column densities (in cm⁻²) taken from Pan et al. (2004) except C₂ (from this work), H i, H₂, and depletion factor (F*) from Jenkins (2009). Gas number densities n are from the chemical model of Pan et al. (2005) and pertain to the CN-containing 2 km s⁻¹ component only. The final column gives an average (±) error estimate on the respective values.

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The λλ5780, 5797, 5850, 6379, and 6614 DIBs have the greatest central depths of those in our sample, and are the only ones to show a significant difference in strength between ρ Oph A and B. Small differences were measured for other bands toward these stars, but they are not significant given the noise level. Our data are consistent with all DIBs being slightly stronger toward ρ Oph B than A (by c. 5%). However, for the AB pair versus C and DE, the situation is more complex. Out of all the ρ Oph sightlines, we find that λ5850 is strongest toward C, whereas λ6614 is weakest. Additionally, DE has clearly the strongest λ6196 DIB, but the weakest λ4726 (the latter DIB is 40% weaker toward DE than C and 32% weaker than AB). This wide variation in behavior of different DIBs between the closely spaced sightlines in our sample shows that the carriers are being influenced by quite subtle variations in physical/chemical interstellar conditions.

An important advance in DIB studies was the identification of a set of DIBs that have a propensity to follow in strength the column density of diatomic carbon molecules—the so-called C₂ DIBs (Thorburn et al. 2003). A spectroscopic property of the C₂ DIBs (possibly related to the structure of their carriers) is that they fall predominantly at the blue end of the spectrum (≲ 5500 Å), whereas the non-C₂ DIBs that we have measured fall at the red end (>5500 Å). Among the strongest of the C₂ DIBs is λ4964, data for which are shown in Figure 3 (bottom panel). This DIB shows no detectable difference between ρ Oph A and B but is found to be weaker toward ρ Oph DE than AB. Interestingly, we find that a total of nine out of the ten observed C₂ DIBs (λλ4726, 4734, 4964, 4984, 5176, 5418, 5512, 5541, and 5546) are measured to be weaker toward DE than AB (Smith 2006). The remaining λ5170 C₂ DIB is overlapped by a broad spectral feature, probably stellar in origin, which precludes a proper measurement of its variation. Ten out of eleven of the non-C₂ DIBs we observed in the yellow-to-red part of the spectrum (for which measured variations are larger than the errors), show the opposite trend, and are stronger toward DE than AB (data for λλ5705, 6660, and 7224 are given by Smith 2006), the single exception being λ5850. These results strongly support the association of the C₂ DIBs as a special group.

Snow et al. (2008) studied the atomic (K i, Na i, and Ca i) column densities toward ρ Oph A, C, and D and determined the three-dimensional locations of these stars with respect to the ρ Oph cloud complex. They conclude that while ρ Oph AB and D are probably in front of the dense molecular cloud and probe a more diffuse portion, ρ Oph C more likely lies embedded in a CN-rich molecular part. As discussed by Snow et al. (2008), the ratio of total-to-selective extinction (R_V) is unusually high toward ρ Oph, and the far-UV rise is quite flat. These observations are consistent with the presence of relatively large dust grains and an absence of the very smallest grains (Seab & Snow 1995). Notably, the reported E_{B − V} values for the ρ Oph stars are almost identical (Table 1), which, given the magnitude of the observed DIB variations, is consistent with previous findings that the DIBs do not arise directly from the large grains that cause optical extinction (Sarre 2006).

A body of data for atoms and small molecules toward the ρ Oph stars is given in Table 2, where the column densities quoted are in general the sum over a number of blended velocity components. Unfortunately, the widths of the DIBs preclude their association with any particular component. We measured equivalent widths for 23 lines of C₂ in each of our spectra (in the range 7700–8900 Å, which covers the (2–0) and (3–0) C₂ vibrational bands). Total C₂ column densities were calculated using the oscillator strengths of Bakker et al. (1997); the best-fitting rotational temperatures (T_rot(C₂)) were obtained using B. McCall's C₂ calculator (at http://dib.uiuc.edu/c2/), and are given in Table 2. All four lines of sight show absorption due to CN at around 2 km s⁻¹, with a common Doppler b ≈ 0.8 km s⁻¹.

Apart from CH⁺ which has a larger total column density toward ρ Oph A than B, the errors on the total column densities are too large to permit any clear distinction to be drawn between the properties of the gas toward ρ Oph A and B. However, if one considers only the main 2 km s⁻¹ component, CH, CH⁺, K i, and Ca ii are all 10%–20% stronger toward B than A. Our high signal-to-noise K i equivalent width measurements also show more neutral potassium toward B than A. We infer that the column of DIB carriers increases toward ρ Oph B in concert with the species noted above. It is difficult, however, to know whether this is due to variations in the gas density, lower atomic (and DIB) depletions, changes in ionization equilibrium, or due to the geometric structure of the cloud.

Our observations toward ρ Oph C, as well as those of Pan et al. (2004) and Snow et al. (2008) show a relative underabundance of K i and Ca i compared with the other three sightlines. Consistent with its (∼3 times) greater CN column density, Snow et al. (2008) concluded that ρ Oph C is more deeply embedded in the dense, CN-rich material of the background (star-forming) molecular cloud. This relatively high density material should result in more depletion, which is the most likely explanation for the low K i and Ca i column densities. The lower λ6614 strength toward ρ Oph C may also be related to this depletion. The fact that the variations in the "red" DIB strengths are only on the order of a few percent for C compared with AB and DE shows that their carriers are, in general, not closely associated with CN. This result is in accordance with the study of λλ5797 and 5780 by Weselak et al. (2008). By contrast, the "blue" λ4726 DIB is very strong toward C, which confirms the association of its carrier with molecular material traced by C₂ (Thorburn et al. 2003).

The greater strengths measured for λλ5797 and 5850 toward ρ Oph C are consistent with the assignment of these DIBs to Krełowski & Walker (1987)'s group III, which is a family of diffuse bands that tend to have relatively narrow profiles. These "KWIII" DIBs (in contrast to those in the KWII group) are generally observed to be strongest in better shielded, denser diffuse clouds where the UV radiation field is weaker and small molecules are more abundant (Cami et al. 1997).

The sightline toward ρ Oph DE is relatively more diffuse than C, as indicated by the lower CN column density. It has a lower depletion factor (F*; Jenkins 2009) than A, indicative of a lower gas density. The 2 km s⁻¹ component toward DE also has the lowest modeled density (Pan et al. 2005), consistent with it being the more diffuse sightline. The greater C₂ rotational temperature measured toward DE may reflect the increased (photo-electric) heating expected to result from a stronger UV radiation field pervading this more diffuse gas. As shown in Figure 2, the amount by which the λ5780 DIB strength increases between A and DE (∼20%), is greater than the amount by which λ5797 increases (∼10%). This pattern is consistent with the hypothesis (Krełowski & Walker 1987; Cami et al. 1997) that the carrier of the λ5780 DIB favors more diffuse environments than λ5797. The strengths of other DIBs in the same (KWII) group as λ5780 (λλ6196, 6203 and 6614) show a similarly large enhancement (∼20%) toward DE compared to A/B. With the striking exception of CH⁺ (−26%), CH, Ca i, K i, and Ca ii show increases in column density in the range 15%–30% (for the main 2 km s⁻¹ component) from A to DE, which is analogous to their increase from A to B discussed earlier, and could plausibly be related to a lower degree of depletion in the gas toward DE.

The relative weakness of the C₂ DIBs toward DE is consistent with this being the more diffuse (and therefore, more molecule-poor) sightline of our sample. Our C₂ spectra are, however, of insufficient signal-to-noise to discern differences in the C₂ column densities between the ρ Oph sightlines; dedicated future C₂ observations will be required to confirm this.