GIANT Hα NEBULA SURROUNDING THE STARBURST MERGER NGC 6240^*

The most remarkable feature of the extended Hα nebula is a bright, wiggled filament ("W-filament") in the western region. It is prominent in the middle panel of Figure 2 (see panel (c) of Figure 7 for the detailed structure). Although the W-filament was previously visible in the narrow-band Hα images taken by Heckman et al. (1987) and Veilleux et al. (2003), its detailed structure and surrounding complex features were only revealed in our Hα+[N ii] image. The bright part of the W-filament begins at 18 kpc west from the s-nucleus and extends to the west–southwest. The filament turns slightly west–northwest at ∼25 kpc and reaches ∼36 kpc from the s-nucleus. There is a bright cloud to the east of the bright region of the W-filament (∼12 kpc west from the s-nucleus). We assume, from a morphological point of view, this could be a part of the W-filament (Figure 6).

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On the northern side of the galaxy, there exists a complex of clouds and filaments ("N-complex": Figure 6 and panel (a) of Figure 7), which extends along the PA ∼310° from ∼25 kpc north to ∼33 kpc north–northwest of the galaxy. The unsharp masked image indicates that the N-complex is the tip of a large curved structure connected to the hourglass region (Figure 5).

A clear bubble-like structure is visible in the southeastern side of the galaxy ("SE-bubble": Figure 6 and panel (b) of Figure 7). This faint bubble was unknown before this study. The size and PA of the SE-bubble are ∼28 kpc × 38 kpc and ∼135°, respectively. On the other side of the galaxy, the N-complex and W-filament appear to form a large broken bubble whose size is comparable to the SE-bubble ("NW-bubble"). The PA of the NW-bubble axis is approximately −40°. The axes of these two bubbles are almost perpendicular to the major axis of the main body of NGC 6240. The opening angle of the SE-bubble and the NW-bubble is ∼70°–90° (Figures 5 and 6).

The western side of the Hα nebula is much brighter and has a more complex structure than the eastern side. In particular, the W-filament and its surroundings have a rich structure consisting of many filaments and blobs.

There are several bright knots ("BK1"–"BK3": Figure 6) located in a loop-like filament extending from the central nebula toward the southwest ("W-loop1": Figure 6). These knots are marginally resolved in our Hα+[N ii] image; the FWHMs of the knots are ≈10–11. W-loop1 extends almost straight toward the southwest from the center, then turns to the north at ∼23 kpc drawing a loop-like trajectory (panel (d) of Figure 7). One of the bright knots (BK3) is located at the tip of this loop.

Another loop ("W-loop2") is located at ∼35 kpc from the center (panel (e) of Figure 7). The size and shape of W-loop2 are similar to those of W-loop1, but its surface brightness is much lower than that of W-loop1.

There are two streams extending along the north–south direction on the south side of the galaxy ("S-filaments" and "SW-filaments": Figure 6; see also Figure 5). The S-filaments consist of several filaments extending from the center toward the south. The SW-filaments consist of a number of filaments distributed from the region close to W-loop2 to ∼59 kpc southwest from the s-nucleus. The extension of the S-filaments is similar to, but slightly offset from, that of the southern tidal feature, which is the southern end of the S-shaped morphology of the optical continuum emission of NGC 6240. There is no clear optical continuum counterpart to the SW-filaments.

We performed multiple aperture photometry for the Hα emission of NGC 6240. We selected the region further out than 140'' (∼70 kpc) from the s-nucleus for sky background estimation. Before performing photometry, bright stars and artifacts were masked and interpolated by linear functions around the patterns. We adopted circular apertures whose center was the brightest point of the Hα+[N ii] image corresponding to the s-nucleus. We excluded the very center (r < 1'') region in this aperture photometry procedure, because the R_C-band frame that was subtracted from the N-A-L671 frame to create the net Hα+[N ii] image was saturated at the central 1'' region of the galaxy. The contamination of the [N ii]λλ6548,6583 emission was corrected assuming ${f}_{[{\rm{N}}{\rm{II}}]6583}/{f}_{{\rm{H}}\alpha }=1$ (Heckman et al. 1990; Keel 1990; Veilleux et al. 2003) for the entire region of the nebula. In addition, we assumed that the foreground extinction is the same as the R-band Galactic extinction toward NGC 6240; A_R = 0.165 magnitude (Schlafly & Finkbeiner 2011).

The surface brightness profile of the Hα emission is given in Figure 8. It is clear that the Hα surface brightness distribution indicates the core-halo structure of the nebula, which is similar to the case of the soft-X-ray nebula (Huo et al. 2004; Nardini et al. 2013), whereas the core is more compact in Hα.

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The total Hα flux of NGC 6240 was derived by integrating the fluxes in the regions at which the Hα surface brightness is greater than 2σ of the sky background. We assumed that the Hα surface brightness within the central 1'' region, which we excluded in the aperture photometry, is the same as that at 1'' away from the s-nucleus. As mentioned above, we assumed that ${f}_{[{\rm{N}}{\rm{II}}]6583}/{f}_{{\rm{H}}\alpha }=1$ for the whole region of the nebula. This ratio tends to increase toward the outer part of the Hα nebula, reaching ≳1.4 around the W-filament (Veilleux et al. 2003). Thus our method may underestimate the contribution of the [N ii] emission in some parts of the nebula. However, since most of the bright regions of the nebula have almost constant ${f}_{[{\rm{N}}{\rm{II}}]6583}/{f}_{{\rm{H}}\alpha }\sim 1$ (Keel 1990; Veilleux et al. 2003), the variation of this ratio does not significantly affect the estimation of the total Hα flux. The resultant total Hα flux ${f}_{{\rm{H}}\alpha }$ is ≈1.3 × 10⁻¹² erg s⁻¹ cm⁻². The Hα luminosity ${L}_{{\rm{H}}\alpha }$ is ≈1.6 × 10⁴² erg s⁻¹.

We derived the root-mean-square (rms) electron density ${n}_{{\rm{e}},{\rm{H}}\alpha }$(rms) of the Hα nebula assuming Case B recombination. We did not, however, apply an extinction correction to this calculation. The surface brightness of ${\rm{H}}\alpha \;{\mathrm{SB}}_{{\rm{H}}\alpha }$ is defined as

$$\begin{eqnarray}&&{\mathrm{SB}}_{{\rm{H}}\alpha }={\alpha }_{{\rm{B}},{\rm{H}}\alpha }\cdot h{\nu }_{{\rm{H}}\alpha }\cdot {\rm{EM}},\end{eqnarray} \tag{ 1 }$$

where ${\alpha }_{{\rm{B}},{\rm{H}}\alpha }$ and EM are the photon production coefficient of Hα in Case B recombination (Osterbrock & Ferland 2006) and the emission measure, respectively. Assuming a fully ionized pure hydrogen gas (n_e ≈ n_p; n_p is proton density) and a constant density distribution along the line of sight, the EM can be written as ${\rm{EM}}={n}_{{\rm{e}},{\rm{H}}\alpha }^{2}l$, where l is the line-of-sight length of the nebula. In the case of spherically symmetric nebula, $l=2\sqrt{{R}^{2}-{r}^{2}}$, where R and r are the radius of the nebula and the projected distance from the nebula center, respectively.

Thus, the ${n}_{{\rm{e}},{\rm{H}}\alpha }$(rms) is given by

$$\begin{eqnarray}&&{n}_{{\rm{e}},{\rm{H}}\alpha }({\rm{rms}})=\sqrt{\displaystyle \frac{{\mathrm{SB}}_{{\rm{H}}\alpha }}{{\alpha }_{{\rm{B}},{\rm{H}}\alpha }\cdot h{\nu }_{{\rm{H}}\alpha }\cdot l}}.\end{eqnarray} \tag{ 2 }$$

The radial profile of the ${n}_{{\rm{e}},{\rm{H}}\alpha }$(rms) distribution is shown in Figure 9.

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The mass of the Hα emitting gas ${M}_{{\rm{H}}\alpha }$ is given by

$$\begin{eqnarray}&&{M}_{{\rm{H}}\alpha }=\int {f}_{{\rm{v}},{\rm{H}}\alpha }^{1/2}\cdot {n}_{{\rm{e}},{\rm{H}}\alpha }({\rm{rms}})\;{m}_{{\rm{H}}}\;{dV},\end{eqnarray} \tag{ 3 }$$

where ${f}_{{\rm{v}},{\rm{H}}\alpha }$ and m_H are the volume filling factor of the Hα emitting gas and the mass of a hydrogen atom, respectively.

We estimated the volume filling factor ${f}_{{\rm{v}},{\rm{H}}\alpha }$ by the following procedure. First, considering the close spatial correlation between the Hα nebula and the soft-X-ray-emitting hot gas (see Section 4.4 below), we assumed that the Hα gas and X-ray gas are locally coexisting. Then we considered the following two cases for the coexistent state: (1) a simple thermal pressure equilibrium case and (2) a ram pressure balance case.

In case (1), we calculated the local electron densities ${n}_{{\rm{e}},{\rm{H}}\alpha }({\rm{local}})$ of the Hα emitting cloud, assuming that the equation ${n}_{{\rm{e}},{\rm{X}}}({\rm{rms}})\cdot {f}_{{\rm{v}},{\rm{X}}}^{1/2}\cdot {T}_{{\rm{X}}}={n}_{{\rm{e}},{\rm{H}}\alpha }({\rm{local}})\cdot {T}_{{\rm{H}}\alpha }$ holds over each radial zone in which Nardini et al. (2013) provided the averaged electron densities ${n}_{{\rm{e}},{\rm{X}}}$ and temperatures T_X of the hot gas individually. Here, ${f}_{{\rm{v}},{\rm{X}}}$ is the volume filling factor of the hot gas. Because the optical emission line filaments of NGC 6240 show a clear shock excitation nature in its emission line ratios (Heckman et al. 1990; Keel 1990), we assumed that ${T}_{{\rm{H}}\alpha }={10}^{5}$ K (typical temperature of optical shock excitation gas Shull & McKee 1979) for all of the zones in this case.

Case (2) is the one in which shock excitation is responsible for the Hα emission gas and X-ray emitting hot gas. We calculated the ${n}_{{\rm{e}},{\rm{H}}\alpha }({\rm{local}})$, assuming the ram pressure balance equation (Spitzer 1978), ${n}_{{\rm{e}},{\rm{X}}}({\rm{rms}})\cdot {f}_{{\rm{v}},{\rm{X}}}^{1/2}\cdot {V}_{{\rm{s}},{\rm{X}}}^{2}\;={n}_{{\rm{e}},{\rm{H}}\alpha }({\rm{local}})\cdot {V}_{{\rm{s}},{\rm{H}}\alpha }^{2}$, where ${V}_{{\rm{s}},{\rm{X}}}$ and ${V}_{{\rm{s}},{\rm{H}}\alpha }$ are the shock velocities at the shock front between the hot gas and warm gas, respectively. The shock velocity required to produce X-ray emitting gas with a temperature of ∼10⁷ K is ∼1000 km s⁻¹ for an adiabatic shock (Hollenbach & McKee 1979). Heckman et al. (1990) found that the averaged [O i]λ6300/Hα and [S ii]λλ6717+6731/Hα emission line intensity ratios of the optical filaments of NGC 6240 are ∼0.3 and ∼0.8, respectively. These values are consistent with shock models with the shock velocity of ∼100 km s⁻¹ (e.g., Shull & McKee 1979). Hence, we assume that ${V}_{{\rm{s}},{\rm{X}}}$ and ${V}_{{\rm{s}},{\rm{H}}\alpha }$ are ∼1000 km s⁻¹ and ∼100 km s⁻¹, respectively.

We then derived ${f}_{{\rm{v}},{\rm{H}}\alpha }$ by ${f}_{{\rm{v}},{\rm{H}}\alpha }={({n}_{{\rm{e}},{\rm{H}}\alpha }({\rm{rms}})/{n}_{{\rm{e}},{\rm{H}}\alpha }({\rm{local}}))}^{2}$ for both cases. In case (1), ${f}_{{\rm{v}},{\rm{H}}\alpha }$ is given by

$$\begin{eqnarray}{f}_{{\rm{v}},{\rm{H}}\alpha } & \approx & {10}^{-4}\cdot \left(\displaystyle \frac{{f}_{{\rm{v}},{\rm{X}}}}{{10}^{-2}}\right){\left(\displaystyle \frac{{T}_{{\rm{H}}\alpha }}{{10}^{5}{\rm{K}}}\right)}^{2}{\left(\displaystyle \frac{{T}_{{\rm{X}}}}{{10}^{7}{\rm{K}}}\right)}^{-2}\\ & & \times {\left(\displaystyle \frac{{n}_{{\rm{e}},{\rm{H}}\alpha }({\rm{rms}})}{{10}^{-1}{\mathrm{cm}}^{-3}}\right)}^{2}{\left(\displaystyle \frac{{n}_{{\rm{e}},{\rm{X}}}({\rm{rms}})}{{10}^{-2}{\mathrm{cm}}^{-3}}\right)}^{-2}.\end{eqnarray} \tag{ 4 }$$

In case (2), we obtain

$$\begin{eqnarray}{f}_{{\rm{v}},{\rm{H}}\alpha } & \approx & {10}^{-4}\cdot \left(\displaystyle \frac{{f}_{{\rm{v}},{\rm{X}}}}{{10}^{-2}}\right){\left(\displaystyle \frac{\eta }{10}\right)}^{-4}\\ & & \times {\left(\displaystyle \frac{{n}_{{\rm{e}},{\rm{H}}\alpha }({\rm{rms}})}{{10}^{-1}{\mathrm{cm}}^{-3}}\right)}^{2}{\left(\displaystyle \frac{{n}_{{\rm{e}},{\rm{X}}}({\rm{rms}})}{{10}^{-2}{\mathrm{cm}}^{-3}}\right)}^{-2},\end{eqnarray} \tag{ 5 }$$

where $\eta ={V}_{{\rm{s}},{\rm{X}}}/{V}_{{\rm{s}},{\rm{H}}\alpha }$.

The volume filing factor ${f}_{{\rm{v}},{\rm{X}}}$ is not well known, but the soft X-ray would be strongly enhanced at the upstream of the shocks in the wind (Strickland & Stevens 2000; Cooper et al. 2008, 2009) and the ${f}_{{\rm{v}},{\rm{X}}}$ of such shocked gas clouds is considerably smaller than one. Strickland & Stevens (2000) found that the typical ${f}_{{\rm{v}},{\rm{X}}}$ of the X-ray gas in a superwind is of the order of 10⁻². Assuming that ${f}_{{\rm{v}},{\rm{X}}}={10}^{-2}$, ${T}_{{\rm{H}}\alpha }={10}^{5}$ K, and η = 10, we found that the resultant volume filling factors are orders of 10⁻⁴–10⁻⁵ for both cases (Table 3). Although these values are much smaller than the typical f_v of H ii regions or planetary nebulae in our Galaxy (∼10⁻²–10⁻³; e.g., Osterbrock & Ferland 2006), they are comparable to the f_v derived for outflow/wind gas by the observations of nearby starburst galaxies and AGNs (Yoshida & Ohtani 1993; Robinson et al. 1994; Cecil et al. 2001; Sharp & Bland-Hawthorn 2010).

Table 3.  Electron Densities and Volume Filling Factors

r Rangea	${n}_{{\rm{e}},{\rm{X}}}$ (rms)b	TXb	${n}_{{\rm{e}},{\rm{H}}\alpha }$ (rms)c	${f}_{{\rm{v}},{\rm{H}}\alpha }$	${f}_{{\rm{v}},{\rm{H}}\alpha }$
(kpc)	(cm−3)	(K)	(cm−3)	case (1)d	case (2)e
0−1.7	4.5 × 10−2	1.3 × 107	2.4 × 10−1	1.6 × 10−5	2.9 × 10−5
1.7−3.4	2.3 × 10−2	1.1 × 107	1.2 × 10−1	2.2 × 10−5	2.7 × 10−5
3.4−7.4	6.7 × 10−3	8.7 × 106	5.2 × 10−2	7.9 × 10−5	5.8 × 10−5
7.4−11	3.5 × 10−3	8.4 × 106	2.8 × 10−2	8.9 × 10−5	6.5 × 10−5
11−15	2.4 × 10−3	7.8 × 106	2.0 × 10−2	1.1 × 10−5	7.1 × 10−5
15−20	2.0 × 10−3	7.7 × 106	1.4 × 10−2	8.5 × 10−5	4.9 × 10−5
20−27	1.6 × 10−3	8.1 × 106	1.1 × 10−2	7.5 × 10−5	4.7 × 10−5
27−39	1.4 × 10−3	7.2 × 106	0.7 × 10−2	5.2 × 10−5	2.5 × 10−5

Notes.

^aAveraging radial range. ^bNardini et al. (2013). ^crms electron density of Hα emitting gas calculated by emission measure of the Hα nebula, assuming a spherical symmetric geometry of the nebula. ^dThe volume filling factor of Hα gas calculated assuming a simple thermal pressure equilibrium with ${f}_{{\rm{v}},{\rm{X}}}={10}^{-2}$ and ${T}_{{\rm{H}}\alpha }={10}^{5}$ K. ^e ${f}_{{\rm{v}},{\rm{H}}\alpha }$ calculated assuming a ram pressure equilibrium with ${f}_{{\rm{v}},{\rm{X}}}={10}^{-2}$ and $\eta ={V}_{{\rm{s}},{\rm{X}}}/{V}_{{\rm{s}},{\rm{H}}\alpha }=10$.

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Under the assumption of spherically symmetric nebula, we obtained ${M}_{{\rm{H}}\alpha }\sim 5\times {10}^{8}\;{M}_{\odot }$ using the ${f}_{{\rm{v}},{\rm{H}}\alpha }$ derived above. However, the assumption of spherical symmetry is too simple for the geometry of the emission line nebula. Hence, we next considered a more realistic geometry for the Hα nebula; it has a filamentary structure and the emission line filaments occupy a fraction of ζ (ζ < 1) along the line of sight of the spherical nebula. In this case, the rms electron densities are increased by a factor of ${\zeta }^{-1/2}$. Since we think that the X-ray gas and Hα gas spatially coexist, the line-of-sight occupation factor ζ is the same for both gases. Thus the factor ζ vanishes in the calculation of the volume filling factor (see Equations (4) and (5)). On the other hand, the total volume of the nebula is decreased by a factor of ζ. As a result, the ${M}_{{\rm{H}}\alpha }$ alters by a factor of ζ^1/2. Then we obtained ${M}_{{\rm{H}}\alpha }\sim 5\times {10}^{8}\cdot {\zeta }^{1/2}\;{M}_{\odot }$ for this partly occupied filamentary nebula case.

We found that the overall morphology of the Hα nebula is dominated by the radially extended filamentary structure (Figure 5). We also identified a large-scale bipolar bubble (SE- and NW-bubbles) along with prominent loop structures (W-loop1 and 2). These morphological characteristics indicate that most of the Hα nebula consists of radially outflowing gas—the starburst superwind.

The close spatial coincidence between the Hα nebula and the extended soft-X-ray gas supports this idea. Many common features, including the SE-bubble, the N-complex, the W-filament, and W-loop2, are observed in the extended ionized gas of NGC 6240 in both Hα and X-ray. A similar morphological correlation between the optical emission line region and the soft-X-ray gas has been widely observed in starburst galaxies (Strickland et al. 2000, 2004; Smith & Wilson 2001; Cecil et al. 2002; Veilleux et al. 2003; Tüllmann et al. 2006). Such good spatial coincidence has been interpreted as an indication that the hot and warm phase gases are heated and excited via the same mechanism—probably from shocks in the outflowing gas—over the entire region of the nebula. Numerical simulations of superwinds have supported this scenario (Strickland & Stevens 2000; Cooper et al. 2008, 2009).

Heckman et al. (1990) obtained the optical spectra of two of the bright knots (BK1 and BK2), and found that they are active star-forming regions. Although there are no spectral data for BK3, we suggest that it is also a star-forming region because the continuum counterparts of the bright knots that show a blue color (B − i' (AB) ∼ −0.3) are clearly seen in our broadband images (Figure 12).¹⁰

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It appears that the bright knots are interacting with the outflowing filament, and the gas in the filament is intercepted by these knots, creating W-loop1. Tail-like structures are also observed in the broadband images, in particular, the tails extending to the southwest are clearly visible in the B-band image (Figure 12). Indications of the tails are also marginally seen in the i'-band image. This feature is recognizable in a Hubble Space Telescope public image of NGC 6240.¹¹ The head-tail morphology of the knots suggests that the outflow gas and condensed gas clouds in the main stellar disk interact with each other and that its ram pressure and/or shock may induce active star formation. The bright knots are also visible in the UV images supporting the view that these knots are the sites of active star formation. In all probability, the young stellar population observed in the tails of the bright knots was also formed by this process.

The S-filaments are well traced by the UV emission. There are also faint UV counterparts for the SW-filaments. Because these filaments have either a weak or no spatial correlation with the X-ray emission, we consider that the S- and SW-filaments are not closely related to the starburst winds. These features are more likely to be star-forming tidal tails.

The UV emission associated with the optical filaments may partly be the precursor emission of a shock within the nebula. The extended UV emission associated with a superwind has been observed in nearby starburst galaxies including M82 and NGC 253 (e.g., Hoopes et al. 2005). In these galaxies, most of the UV emission is thought to be from the stellar continuum scattered by dust entrained in the wind (Hoopes et al. 2005; Hutton et al. 2014), because it is too bright to be provided by shock-heated or photoionized gas in the wind. It is well known that NGC 6240 is a dusty system, thus it is possible that the extended UV emission of NGC 6240 is mostly derived from light scattered by dust.

To estimate the contribution of dust scattering, we compared the UV intensity and Hα intensity at the W-filament. We measured the UV and Hα intensities in a 195 × 105 rectangular region at the W-filament (the center of the photometry region is located 54'' west and 5'' north from the s-nucleus), and determined that I_NUV = 2.0 × 10⁻¹⁴ erg s⁻¹ cm⁻², I_FUV = 9.3 × 10⁻¹⁵ erg s⁻¹ cm⁻², and ${I}_{{\rm{H}}\alpha }=2.7\times {10}^{-15}$ erg s⁻¹ cm⁻² (assuming that ${f}_{[{\rm{N}}{\rm{II}}]6583}/{f}_{{\rm{H}}\alpha }=1.4$, Veilleux et al. 2003). The average Hα surface brightness in this region was 1.7 × 10⁻¹⁷ erg s⁻¹ cm⁻² arcsec⁻². Both the ${I}_{\mathrm{NUV}}/{I}_{{\rm{H}}\alpha }$ and ${I}_{\mathrm{FUV}}/{I}_{{\rm{H}}\alpha }$ ratios—7.4 and 3.4, respectively—are slightly smaller than, but almost consistent with, values in the outflow filaments in M82 and NGC 253 at the same Hα surface brightness level (Hoopes et al. 2005). The ratios are, however, still larger than the values predicted by the shock model. This indicates that the UV emission of NGC 6240 is too bright to be attributed to shock-heated or photoionized gas emission in the wind. Thus, some part of the UV emission is possibly light scattering by dust.

Alternatively, this emission may reflect extended star-forming activity within the filaments. In this case, gas compression by the interaction of the outflowing gas and the ISM in the main stellar disk may have induced star formation. Nardini et al. (2013) found that the X-ray hot gas of NGC 6240 has a spatially uniform sub-solar metal abundance whose abundance pattern is consistent with Type-II supernova enrichment. They discussed the age of the nebula using the thermal velocity dispersion estimated from the temperature of the hot gas, and concluded that the whole NGC 6240 nebula could not be formed by the recent activity of the nuclear starburst or AGNs. They suggested that extended continuous star formation from the early phase of the merging event would be responsible for the formation of the outer nebula. The extended UV emission and its spatial coincidence with the Hα emission might support this scenario.

Most of the Hα emission of the Hα nebula, however, probably originates from shocks in the outflow because the extended ionized gas shows a large [N ii]/Hα intensity ratio ([N ii]/Hα ≳ 1; Keel 1990; Veilleux et al. 2003). Detailed polarimetry and spectroscopy of the extended nebula are necessary to qualitatively understand the contribution of dust scattering and local star formation to the Hα nebula.