THE LOCAL LEO COLD CLOUD AND NEW LIMITS ON A LOCAL HOT BUBBLE

The LLCC was originally discovered in the 21 cm (1420 MHz) hyperfine transition of H i, and this kind of observation remains the best way to study the morphology of the structure. Temperature measurements of ∼20 K from Crovisier & Kazes (1980) and Heiles & Troland (2003) indicate that the cloud should be largely neutral. Measured column densities below 10²¹ cm⁻² and the non-detection of OH (Crovisier & Kazes 1980) indicate that the cloud should be largely atomic and not entirely optically thick to H i. Thus, H i is an excellent tracer of the bulk of the cloud.

We observed the LLCC as part of the Galactic Arecibo L-Band Feed Array H i (GALFA-H i) survey, an ongoing survey of H i conducted with the Arecibo 305 m telescope using the Arecibo L-band Feed Array (ALFA). The GALFA-H i survey is a survey of all Galactic H i within the Arecibo observing range (−1° < δ < 38°), at a spectral resolution of 0.184 km s⁻¹, and an angular resolution near 4'. The details of the observation methods, data reduction, and properties of the GALFA-H i survey can be found in Peek et al. (2011). Observations for this region were conducted both as part of a targeted map and as part of the Turn On GALFA Survey, a commensal (simultaneous) survey conducted alongside the Arecibo legacy fast ALFA project. The targeted map was conducted in a high-speed, basket-weave or meridian-nodding mode, whereas the Turn On GALFA Survey data were collected in a low-speed drift mode. Most areas of the map were covered by both modes and have an rms noise of 0.17 K over a 0.184 km s⁻¹ channel, but a few areas were only covered with the basket-weave mode, and have an rms noise of 0.36 K per 0.184 km s⁻¹ channel. It is important to point out that the detailed analysis of the H i column density found in Section 3.1 could not be accomplished without the exquisite spectral resolution of the GALFA-H i survey. We note that while we refer to the section of cloud we are examining as the LLCC, the southern component of the LLCC is actually in Sextans.

A key result from the Infrared Astronomical Satellite (IRAS) was the discovery of the correlation between the infrared diffuse emission and the Galactic H i column density for |b| > 10° (Low et al. 1984). The conclusion reached by these authors, and others, is that there is a relatively uniform interstellar radiation field in the Galaxy and a relatively constant large-grain dust fraction in the diffuse H i. The interstellar radiation field heats the dust, which reradiates in the infrared, which in turn creates the observed correlation between H i column density and dust emission. Thus, the investigation of the large-grain dust contents of the LLCC can be conducted by examining the infrared flux in the IRAS data. IRAS observed the vast bulk of the sky for 300 days in 1983, in the infrared wavelength bands centered on 12, 25, 60, and 100 μm (Beichman 1987). These data have been reduced in numerous ways and have produced many different maps: the original SkyFlux Atlas, the IRAS Sky Survey Atlas, the maps produced in Schlegel et al. (1998), and the IRIS maps described in Miville-Deschênes & Lagache (2006). We use these most recent IRIS maps, as they are produced at 4' resolution, nearly identical to the resolution of the GALFA-H i maps. In addition to the IRIS reduction, we further clean the data set by removing point sources down to a magnitude of 1 Jy. We remove these point sources by fitting with Gaussian point-spread function and excise any point sources we cannot cleanly remove. We note that the presence of IRC+10216 in this field has made the analysis of the 60, 25, and 12 μm band impossible, as this exceedingly bright star has contaminated a large swath of the map.

Based on observations of the interstellar Na i D₂ λ5889.951 and D₁ λ5895.924 absorption toward 33 stars in the LLCC sky region, Meyer et al. (2006) were able to place a significant upper limit on the LLCC distance. Using the updated Hipparcos parallaxes of van Leeuwen (2007), this upper limit corresponds to the distances (39.9^+0.4_−0.3 and 40.5^+1.1_−1.0 pc) of the nearest stars (HD 83808 and HD 83683) exhibiting LLCC Na i absorption. The challenge in tightening this distance constraint further is finding nearer stars within the LLCC H i 21 cm emission sky envelope that are bright enough for high-resolution optical absorption-line spectroscopy. Given the paucity of such stars, another possibility is to consider nearer stars just outside the H i envelope that might be bright enough to detect very weak optical absorption at a gas column threshold more sensitive than the 21 cm observations.

Fortunately, there is a very bright (V = 1.35) star, Regulus (α Leo), at a distance of 24.3 ± 0.2 pc (van Leeuwen 2007) and a sky position that is 3° outside the LLCC 21 cm envelope. Meyer et al. (2006) obtained high signal-to-noise (S/N) spectra of Regulus in the Na i wavelength region to divide out the atmospheric absorption lines in the spectra of their program stars. These Regulus spectra reveal no evidence of any interstellar Na i absorption to high precision. Since Regulus is often used as a spectral standard for observations in other wavelength regions, we examined our archive of data obtained with the 0.9 m coudé feed telescope and spectrograph at Kitt Peak National Observatory (KPNO) over the past 15 years. Observations of Regulus in the vicinity of the interstellar Ca ii K λ3933.663 and H λ3968.468 lines were obtained with this instrumentation in 2002 November in support of a study of variable interstellar absorption toward ρ Leo (Lauroesch & Meyer 2003). Eight F3KB CCD echelle spectra spanning a total exposure time of 80 minutes were taken of Regulus with a spectrograph configuration yielding a measured velocity resolution of 3.6 km s⁻¹ at the Ca ii K wavelength. We utilized the NOAO IRAF echelle data reduction package to bias-correct, scattered light-correct, flat-field, wavelength-calibrate, order-extract, and sum the individual CCD exposures into the final net Regulus spectrum.

For the lower limit on the LLCC distance, we borrowed archival data from the California Planet Search of the nearby M-dwarf HIP 47513. This star is at a distance of 11.26 ± 0.21 pc (van Leeuwen 2007). In addition to HIP 47513, we used spectra of four additional M-dwarfs with similar features as comparison stars. Those other spectra allow us to compare the Na i D₁ line core of HIP 47513 with the cores of other M-dwarfs, and thus to place a stronger constraint on the interstellar Na i column along the HIP 47513 line of sight.

Our five M-dwarf spectra were collected with HIRES, the high-resolution echelle spectrograph on Keck I (Vogt et al. 1994) between 2005 February and 2006 December. Details are in Table 1. The spectra were originally collected with the intention of detecting extrasolar planets; see Marcy & Butler 1992 for a full description of the California Planet Search (CPS) and its goals. Because of their radial velocity measurement technique, the vast CPS archive contained only a few useful spectra of each star. The CPS places an iodine cell in the light path of its stars, which creates a forest of I₂ absorption lines for radial velocity reference. Here we make use only of the iodine-free "template" spectra, which contain no I₂.

The spectra have a resolution R = 70, 000 at λ = 5500 Å. We chose the stars in Table 1 because they were M-dwarfs with B − V within ±0.12 of HIP 47513, and they had iodine-free observations in the CPS archive. We began with 33 candidate stars that fit those two criteria, then inspected by eye three separate segments of the spectra, away from the Na iD lines. We graded candidates in how well their spectral features matched the features of HIP 47513. Four spectra were considered excellent matches, and those are the four we chose to include in this analysis.

We examine X-ray data in order to determine the provenance of the SXRB via the X-ray shadow of the LLCC (see Section 1). The Röntgensatellit (ROSAT; Truemper 1992) produced a flood of high-quality, all-sky data taken during 6 months in 1990 and 1991. Among these data were the SXRB maps. We use the second data release of these maps, refined from the original reduction to have higher resolution (12') and fidelity (Snowden et al. 1997). The SXRB maps cover 98% of the sky in three wavebands: 1/4 keV, 3/4 keV, and 1.5 keV. We focus only on the 1/4 keV (C-band) SXRB maps in this work, as H i can act as very effective absorber of photons at this energy level, even at the intermediate column densities of the LLCC. The SXRB maps were reduced so as to exclude variable local emission and the contribution of point sources. The brightness of the SXRB maps are reported in units of 10⁻⁶ counts s⁻¹ arcmin⁻² (Snowdens).

3.1.1. H i Profile Saturation and Optical Depth

Heiles & Troland (2003) measured H i absorption spectra and derived H i excitation temperatures T_x for the strong radio continuum sources 3C225a, 3C225b, and 3C237 and found T_x = 21.53 ± 1.30, 17.44 ± 1.80, 13.61 ± 0.26 K, respectively. These three low temperatures suggest that the large parts of the LLCC are very cold, ≲ 20 K, throughout. Given that we find brightness temperatures in this range for the center of the cloud, we expect large parts of the cloud to be optically thick at line center, i.e., saturated. The derivation of accurate H i column densities requires including this line saturation effect; if we neglect the saturation, then the derived values are too small.

It is easy to correct for line saturation at the position of a sufficiently strong continuum source, because then we can measure not only the emission profile but also the optical depth profile. However, it is difficult to correct for saturation throughout the cloud's area, because without a background continuum source our only indication of line saturation is the line shape.

Knapp & Verschuur (1972) attacked this problem by assuming that the intrinsic line shape is accurately represented by a single saturated Gaussian component and used least-squares fits to determine its degree of saturation and spin temperature T_x. However, we have found that the stronger H i profiles tend to be double-peaked, so this model is invalid.

Figure 1 shows an example spectrum toward the core of the cloud. Visually, the lines look double. This is confirmed by least-squares fitting of three models. The solid line is the observed profile, the dotted line the single unsaturated Gaussian fit, the dashed line the saturated Gaussian fit, and the dash-dot line the double unsaturated Gaussian fit. The dispersions of the data from each fit are indicated on the upper left of each plot. The double unsaturated Gaussian fit is the best fit, and we find this to be generally true across the highest column parts of the cloud. We must therefore consider multi-component models to accurately fit these data.

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3.1.2. Five Fitting Models

3.1.3. Procedure for Fitting Gaussians

3.1.4. Model 4 or Model 5?

Our goal is to settle on three models that apply to the whole cloud. Two of these are the single- and double-component optically thin models, which are appropriate for weak lines, where the effects of optical depth have an undetectable effect on our observed line shape. These are models 1 and 2. The third, for strong lines, is either model 4 or model 5. We choose between these by examining their σ_line statistics.

Here we compare σ_line values for various model pairs. For most of the fits model 2 has a lower σ_line than model 1 indicating that the vast majority of sight lines have double components.

By doing the same comparison between model 4 versus model 5, we find the vast majority of sight lines are better fit by model 5. Philosophically, we would have much preferred the physically motivated model 4, but we must be driven by the data. A second philosophical objection is illusory: why should the near side be optically thick and the far be optically thin? The answer is that the far side may well be optically thick; we cannot derive its optical depth because, as discussed in the context of model 3, all we have is its line shape, and this is an insufficient "handle."

We also make the same comparison between model 5 and model 2, i.e., including optical depth effects versus assuming optically thin components. We find that while a majority of points have lower σ_line for model 5, there remain a large minority of sight lines that clearly prefer model 2. This is to say that optically thin fits are useful along with the optically thick fits.

For each pixel, we choose the converged model with the lowest σ_line. Figure 2 shows an image of the cold cloud in which the brightness indicates column density and the color shows the model selected. For each pixel, the column density is calculated from the selected model.

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To generate a map of the IR flux associated with the LLCC, first we must remove the bulk of the IR flux which is associated with the background H i, beyond the local cavity. We make a map of the background H i column by integrating over the entire Galactic H i line in the GALFA-H i data and subtracting the H i flux associated with the LLCC, as found in Section 3.1, model 1. Then we fit a first-order polynomial to the relationship between the background H i and the IR flux, over the region of the map, using a standard least-squares approach. We have implicitly assumed that the background H i is largely optically thin, typically a good assumption for high-latitude warm neutral medium (WNM) where T_A ≪ T_x. We then subtract the expected IR background flux from the observed IR flux to determine the residual flux we expect to be associated with the LLCC itself. The result of this process is shown in Figure 3.

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4.1.1. Lower Limit

Ruling out Na i absorption along the line of sight to HIP 47513 requires detailed analysis for two reasons. First, HIP 47513's crowded M-dwarf spectrum hides interstellar absorption features far more effectively than smooth A-star spectra like HD 84722 in Meyer et al. (2006). Second, and more problematic, is that HIP 47513 has an LSR radial velocity of +4 km s⁻¹ (Nidever et al. 2002), indistinguishable from the cloud's line-of-sight velocity. That is to say, the cold cloud's Na i absorption features would fall right the middle of HIP 47513's deep Na i D lines.

Because of the difficulty untangling the Na i features of the cloud and the star, we compare the Na i D line core of HIP 47513 with the line cores of four additional stars with similar spectra, selected by the method described in Section 2.3 and shown in Table 1. Our goal is to place an upper limit on the Na i between the Sun and HIP 47513.

With four stars in hand, we set about normalizing their spectra to HIP 47513. In the CPS setup, the Na i D lines fall near the edge of one of the HIRES echelle orders. The steep blaze function in that region required a somewhat sophisticated normalization procedure. Using segments of the spectrum 5 Å from the Na i line cores, we fit a fifth-degree polynomial to both of the Na i line wings.

With the Na i features normalized across all five spectra, we zoomed in on the line cores. By eye, it appeared that no cold-cloud absorption is present in the HIP 47513 spectrum (see Figure 4). We next put a quantitative limit on the amount of Na i that could be along the line of sight to HIP 47513.

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First, we averaged the four comparison spectra using a straightforward mean where each spectrum was weighted equally. Then, we subtracted the comparison spectrum from the HIP 47513 spectrum and looked for Na i absorption features in the residuals. By fitting a Gaussian to the residuals at both line cores, we find a maximum absorption-line intensity of 0.00 ± 0.15, equivalent to a Na i column density of (0.00 ± 2.1) × 10¹⁰ cm⁻².

In Meyer et al. (2006), Na i column is found in 23 stars behind the LLCC. We find that all eight stars with measured cold cloud N_HI > 2 × 10¹⁹ cm⁻² have N_NaI > 5 × 10¹¹ cm⁻². As the H i column toward HIP 47513 is ∼10²⁰ cm⁻², the measured 3σ limit of N_NaI < 6.3 × 10¹⁰ cm⁻² implies conclusively that HIP 47513 lies in front of the LLCC.

4.1.2. Upper Limit

As illustrated in Figure 5, our high S/N ≈ 900 spectrum of Regulus reveals a weak interstellar Ca ii K line at an LSR velocity of 3.1 km s⁻¹. In order to give some context for the Ca ii K line, we have also included spectra of the interstellar Na i D₁ and Ca ii K absorption toward HD 84722 in Figure 5. HD 84722 is positioned at a distance of 84.8 ± 3.8 pc (van Leeuwen 2007) behind the core of the LLCC. The Na i D₁ spectrum of HD 84722 is from Meyer et al. (2006) and the Ca ii K spectrum is the product of data obtained in 2007 March with the KPNO coudé feed. The latter spectrum was reduced in a similar manner as the Regulus spectrum and it represents the sum of six T2KB CCD spectra spanning a total exposure time of 11 hr at a velocity resolution of 1.5 km s⁻¹. The measured LSR velocities of the LLCC Ca ii K and Na i D₁ absorption toward HD 84722 are +3.9 and +3.7 km s⁻¹, respectively. The fact that the Na i absorption is much stronger than that of Ca ii toward HD 84722 is consistent with the character of the cold, dense atomic gas at the core of the LLCC.

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Detection of interstellar Ca ii absorption toward Regulus indicates that some cool interstellar matter must exist between us and Regulus. Unfortunately, Regulus is not along the line of sight directly toward the body of LLCC, but rather about 3° away. Regulus is on the side of the LLCC that seems to have a significant amount of connected cloud debris (see Figure 2), and indeed Regulus is very near a small cloudlet, which we dub the Regulus cloudlet. We expect the Regulus cloudlet is a component of the LLCC for a number of reasons. First, it is only a degree or so away from a long plume of material extending toward higher right ascension from the break between the two main LLCC clouds, and therefore is most likely a fragment of this plume. Second, it has a line width entirely consistent with the rest of the LLCC. Perhaps most tellingly, when we perform a least-squares fit to the velocities of the main body of the LLCC with a model that accounts for its velocity in 3-space, the velocity expected for the position of the Regulus cloudlet were it to be moving as a solid body with the LLCC is 3.19 km s⁻¹ LSR. This velocity is in very nice accord with what is observed for the cloudlet (see Figure 6). This also implies that the Regulus cloudlet is largely dynamically connected to rest of the LLCC, making a scenario in which the cloudlet is at a dramatically different distance rather contrived.

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The next question is whether the absorption in Ca ii seen in Regulus is associated with Regulus cloudlet. Regulus is on the very margin of the cloudlet with no obvious cloudlet CNM H i detectable directly toward the star. The CNM non-detection is consistent with the non-detection of Na i  typically associated with CNM, as in HD 84722 (Figure 5). Ca ii is frequently detected toward warmer neutral gas or even largely ionized gas, if the gas is not so hot as to deplete Ca ii into higher ionization states. Therefore, it is not surprising to detect Ca ii on the edge of an ablating piece of cloud where material could easily be heating up and mixing with the warmer, more ionized ISM. Perhaps even more conclusively, the Ca ii absorption in Regulus is at almost exactly the same velocity as the nearest part of the Regulus cloudlet (see Figures 5 and 6). We conclude that we are indeed seeing absorption of Regulus by the cloudlet, and that therefore the LLCC is closer than 24.3 ± 0.2 pc.

Crovisier & Kazes (1980) and Heiles & Troland (2003) found temperatures for the LLCC ranging from 13 K to 22 K for positions toward strong radio continuum sources. To determine temperatures for the rest of the cloud, we turn to our saturated line fits from Section 3.1.

4.2.1. Deriving the Spin Temperature for Model 5

We use the standard radiative transfer equation to derive the spin temperature T_x for the near-side absorbing component. Model 5 consists of four independent components that contribute to the measured brightness temperature T_B(V).

• 1.  
  The continuum background, which consists of the cosmic microwave background (2.8 K) plus the Galactic continuum background. We estimate the latter from the 408 MHz Haslam et al. (1982) survey by applying a brightness-temperature spectral index of −2.6 to their measured value of ∼15 K; this gives 0.6 K at 1420 MHz. Thus, we take the total continuum brightness temperature to be T_bkgnd = 3.4 K.
• 2.  
  The contribution from background H i, which we call the WNM contribution because the profiles are much wider than those of the cold cloud. This is velocity dependent. We represent it by the symbol T_WNM(V). Numerically, it is represented by the fitted polynomial mentioned in Section 3.1.
• 3.  
  The LLCC's more distant optically thin component, which we represent with subscript 0. Its brightness temperature in the absence of absorption is

  $$\begin{equation} T_{B, 0}(V)= T_0 \exp \left[ {-} {(V-V_0)^2 \over 2 \delta V_0^2} \right]. \end{equation} \tag{ 1 }$$
• 4.  
  The closer optically thick component, which we represent with subscript 1. Its brightness temperature in the absence of absorption is

  $$\begin{equation} T_{B, 1}(V)= T_{x,1} (1 - \exp (-\tau _1 (V))), \end{equation} \tag{ 2 }$$

  where

  $$\begin{equation} \tau _1 (V)= \tau _{1,0} \exp \left[ {-} {(V-V_1)^2 \over 2 \delta V_1^2} \right]. \end{equation} \tag{ 3 }$$

Putting all these together, the emergent brightness temperature T_B(V) is

$$\begin{equation} T_B(V) = T_{B, 1}(V) + [ T_{{\rm bkgnd}} + T_{\rm WNM}(V) + T_{B, 0}(V) ] \exp [-\tau _1 (V)]. \end{equation} \tag{ 4 }$$

Our measured profiles are frequency-switched, which means that we measure the quantity T_B(V) − T_bkgnd. For each pixel, this is the fitted function; we derive the spin temperature T_{x, 1} from the fitted parameters.

Figure 7 shows the histogram of derived spin temperatures T_{x, 1} for the near component. It peaks near 20 K and a secondary peak near 26 K, and is well-confined within the range 15–30 K. Some anomalously low points with very low spin temperatures are presumably the result of spurious fits.

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Figure 8 shows the image of T_{x, 1}. Regions of high column density tend to be warmer. For weak profiles with low column density, T_{x, 1} is often indeterminate, as shown by the gray.

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Given the lower and upper limits on LLCC distance derived in Section 4.1.1 and Section 4.1.2, we can determine a range for a variety of other quantities that pertain to the cloud. We determine the mass by simply integrating the H i column over the area of the cloud, and physical size can be determined from angular size. If we make the very rough assumption that the cloud is as thick as it is wide we can determine a density for the cloud, using a fiducial peak column density of 2.5 × 10²⁰ cm⁻². Using the temperature range found in Section 4.2, we can determine a range of pressures. We can determine a rough lifetime for the cloud by dividing the fiducial scale, in this case the cloud width, by the collapse velocity (see Section 4.3.1). Results are shown in Table 2.

Our assumption that the cloud's thickness along the line of sight is comparable to its width (i.e., the cloud is tubular), is at odds with the result from Heiles & Troland (2003) that this cloud is thin along the line of sight (i.e., sheet-like). Heiles & Troland (2003) assumed the cloud was outside the local cavity, and thus 4–9 times more distant than we now know it to be. This distance assumption gave an overestimate of the width of the cloud and thus an overestimate of its aspect ratio. Heiles & Troland (2003) also measured the thickness at low-column (N_H ≃ 3 × 10¹⁹ cm⁻²) sight lines to 3C225a and 3C225b, about 10 times lower column than the thickest part of the cloud in the present analysis. This low column density gave a lower expected thickness for a given measured temperature and assumed pressure, further exacerbating the inconsistency with our work. If the cloud were 10 times flatter along the line of sight than we have assumed, the derived pressures would be 10 times higher, far out of the standard ISM pressure range (e.g., Wolfire et al. 2003). This result validates our present picture of the cloud as more tubular than sheet-like.

4.3.1. Kinematics of the Two Velocity Components

We found that two velocity components usually provide better H i line fits than a single velocity component. Model 5 allows us to discriminate between the velocities of the two components, because the optically thick foreground component is closer. Accordingly, comparing their velocities allows us to determine whether the components are approaching or receding from each other.

There is an overwhelming tendency for the front component to have higher velocities than the rear. This indicates that the components are approaching one another, typically with a velocity difference ∼0.4 km s⁻¹.

Figure 9 images the velocity difference between the two components, with color indicating velocity difference and lightness the total profile area, i.e., the sum of the two component areas. We only show pixels which had a successful two-component fit, where the two components were not degenerate. Generally, the edges of the cloud (where the intensity is low and the line-of-sight thickness is small) have larger relative velocities. In these portions of the cloud, the cloud is getting thinner with time by ∼0.4 pc Myr⁻¹. If we assume the two components are not more separated than the width of the cloud, we find the collision time to be approximately 1 Myr.

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To determine the large-grain dust content of the cloud we investigate the relationship between the H i column density and the 100 μm emission. Figure 10 shows the relationship between these two quantities over the cloud area. It is clear from this plot that there is a very narrow and tight linear relationship between H i column and IR flux. This serves as an important confirmation that the fitting procedures used in Section 3.1 to determine the saturated H i column densities are not wildly incorrect. For comparison we show the relationship between the IR flux and the H i column under an optically thin assumption, model 1 (Figure 10, contours). These two quantities are clearly not linearly correlated, confirming that an optically thin assumption for the H i in the LLCC is wrong.

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The cloud has a 100 μm to H i column value of 48 × 10⁻²² MJy sr⁻¹ cm². This is somewhat lower than the result from Schlegel et al. (1998) for low column density areas of 67 × 10⁻²² MJy sr⁻¹ cm². It is also lower than the standard Galactic values found in Boulanger & Perault (1988), about 100 × 10⁻²² MJy sr⁻¹ cm², although consistent with their measurement toward Auriga. Similarly, it is lower than the average value found for Galactic clouds measured in Peek et al. (2009), but consistent with the values found for the L1, L7, and L8 clouds found in that work. We conclude that this cloud has either a lower-than-expected overall dust grain density, or a population that has somewhat larger grains than typical for the ISM that maintain a somewhat lower temperature. In either case the cloud is only marginally anomalous in IR production via grains.

In recent years much pioneering work has been done on the closest known interstellar clouds to the Sun (e.g., Redfield & Linsky 2008; Frisch et al. 2009; Redfield & Falcon 2008). The core results from these investigations is that the Sun resides within the CLIC, a collection of clouds within about 15 pc of the Sun, and that these CLIC clouds move as solid bodies past the Sun in roughly the same direction. These clouds have typical column densities near 10¹⁸ cm⁻² with temperature near 10⁴ K. Many authors (Audit & Hennebelle 2005; Vázquez-Semadeni et al. 2006; Heitsch et al. 2006) have conjectured that CNM clouds are formed by the convergence of flows of warm neutral gas. RL08 pointed out that the LLCC is coincident on the sky with the faster moving Gem cloud and the slower LIC, Auriga, and Leo CLIC clouds and suggest that the collision of these clouds may be responsible for the formation of the LLCC.

Our new lower limit on the LLCC of 11.26 ± 0.21 pc seems to rule this out. All of the clouds suggested by RL08 to be in collision to form the LLCC are detected to be, at least in part, closer than 11.1 pc (the LIC envelops the Sun; Gem is detected within 6.7 pc, Leo within 11.1 pc, and Aur within 3.5 pc). The CLIC clouds certainly span some range of distances, and the extent to which they are entangled is still largely unknown; it is possible that the LLCC exists at some more complex interface between these clouds. That said, the simple picture in which the LLCC is at the interface of two independent colliding clouds seems unrealistic given the new distance lower limit on the LLCC.

While the distance of the LLCC does seem to separate it somewhat from the CLIC, the velocity of the LLCC may imply that the CLIC and the LLCC are more related. Haud10 fit the LRCC (the ribbon of clouds that contains the LLCC) with a traversing, rotating, expanding, ring. It is unclear if this model has physical significance, as such structures aren't typically found in dynamic models of the ISM, but the results are striking. The overall velocity of the ring is 37.4 km s⁻¹ toward l = 213°, b = −111 in the Haud10 fit (when converted to the heliocentric reference frame), which is nicely consistent with the velocities of the CLIC clouds found in RL08 (Figure 11). It is possible that this consistency is a coincidence, but it seems more likely that the CLIC clouds and the LRCC are related to each other in some way, or to a parent cloud population.

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As we outlined in the Introduction, there is significant controversy over the origin of soft X-rays seen coming from all directions in the sky. In the local hot bubble theory, soft X-rays at low and intermediate latitudes emanate from a pervasive hot gas filling the local cavity, with the excess of soft X-rays at high latitude coming in part from the halo. In the SWCX theory, the isotropic component of soft X-rays comes from the very nearby (∼100 AU; Stone et al. 2005) interaction of the solar wind with the surrounding ISM. We can differentiate between these theories by looking for absorption of these X-rays by the LLCC.

In the direction of the LLCC the wall of the local cavity is between 100 pc and 150 pc away (Meyer et al. 2006), and we have demonstrated that the LLCC is between 11.26 pc and 24.3 pc away (Section 4.1). We have determined the H i column of the LLCC rather carefully (Section 3.1), and the local bubble wall has a column density of ∼3 × 10²⁰ cm⁻² in the direction to the LLCC. For a reasonable spectral energy distribution over the 1/4 keV band of ROSAT, an optical depth of 1 corresponds to an H i column of 1.05 × 10²⁰ cm⁻² (Snowden et al. 1997). To determine the amount of flux coming from in front of the LLCC we fit the equation

$$\begin{equation} I_C\left(\alpha, \delta \right) = I_{C, \rm foreground} + I_{C, \rm background} e^{-\tau _C\left(\rm H\,{ \mathsc {i}}\left(\alpha, \delta \right)\right)}, \end{equation} \tag{ 5 }$$

where I_C(α, δ) is the flux detected toward the cloud in the SXRB 1/4 keV (C-band) map and $\tau _C({\rm H}\, \mathsc{ i}(\alpha, \delta))$ is the optical depth to C-band X-rays caused by the LLCC as a function of position on the cloud. I_{C, foreground} and I_{C, background} are the unknown X-ray flux from in front of and behind the cloud, respectively. To do this we use a standard least-squares fit. The standard errors derived from a least-squares fit are correct under the assumption that the deviation from the model is Gaussian and uncorrelated. Neither of these assumptions is true in our case, since there are obvious structures in the SXRB 1/4 keV maps in this direction that are far from random, and will generate strong, unmodeled errors. To account for this we use the displacement mapping technique described in Peek et al. (2009), which determines the errors by examining the fluctuation in the result as the cloud is placed in different positions on the nearby sky. We then use this fluctuation to better estimate the errors. Using this method we find $I_{C, \rm total} = I_{C, \rm background} + I_{C, \rm foreground} = 708 \pm 11$ Snowdens and I_{C, foreground} = 398 ± 38 Snowdens (see Figure 12).

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To constrain models for the origins of the soft X-rays, we combine the above result with the results toward the intermediate latitude (b = −34.04) molecular cloud MBM 12 in Kuntz et al. (1997) and Kuntz & Snowden (2000). In this analysis the authors found that MBM 12 was at the back of the local bubble, as it shows very little, if any, shadowing in the C-band, is at a distance of ∼90 pc (Sfeir et al. 1999), and with a C-band flux of 347⁺⁵₋₁₀. Any model we construct for the origin of the C-band X-rays must account for these two measurements.

5.2.1. Model A: Local Hot Bubble and Halo Emission

5.2.2. Model B: Solar Wind Charge Exchange and Halo Emission

5.2.3. Model C: Local Hot Bubble, Solar Wind Charge Exchange, and Halo Emission