CHANDRA AND SPITZER IMAGING OF THE INFRARED CLUSTER IN NGC 2071

The Chandra observation (ObsID 7417) began on 2007 November 6 at 19:40:19 TT and ended on November 7 at 15:18:03 TT. Pointing was centered near NGC 2071 IRS 1 at nominal pointing coordinates (J2000) R.A. = 05^h47^m0494, decl. =+00°22'118. The exposure live time was 67,180 s. Exposures were obtained using the Advanced CCD Imaging Spectrometer (ACIS-I) imaging array in faint timed-event mode with 3.2 s frame times. ACIS-I has a combined field of view (FOV) of ≈169 × 169 consisting of four front-illuminated 1024 × 1024 pixel CCDs with a pixel size of 0492. Approximately 90% of the encircled energy at 1.49 keV lies within 2'' of the center pixel of an on-axis point source. Our discussion here focuses on ACIS-I, but the S2 and S3 CCDs in the ACIS-S array were also enabled. More information on Chandra and its instrumentation can be found in the Chandra Proposer's Observatory Guide (POG).⁸

The Level 1 events file provided by the Chandra X-ray Center (CXC) was processed using CIAO version 3.4 ⁹ using standard science threads. The CIAO processing applied calibration updates (CALDB version 2.26), selected good event patterns, and determined source centroid positions.

Additionally, we used the CIAO wavdetect tool to identify X-ray sources on the ACIS-I array. We ran wavdetect on full-resolution images using events in the 0.3–7 keV range to reduce the background. The wavdetect threshold was set at sigthresh = 1.5 × 10⁻⁵ and scale sizes of 1, 2, 4, 8, and 16 were used. We identified 207 X-ray sources on ACIS-I down to a threshold of 5 counts, including the sources of primary interest in this work located in the NGC 2071-IR core region listed in Table 1.

The unabsorbed X-ray luminosity upper limit for non-detections is log L_X (0.3–7 keV) ⩽ 28.7 erg s⁻¹. This upper limit from the Portable Interactive Multi-Mission Simulator (PIMMS)¹⁰ assumes a 5 count on-axis detection threshold and an absorbed isothermal optically thin plasma spectrum similar to that observed for IRS 1 (N_H = 2 × 10²² cm⁻², kT = 2 keV). At this upper limit, the observation was sensitive enough to detect low-mass T Tauri stars in Orion down to ≈0.1–0.2 M_☉, based on the known correlation between L_X and stellar mass in T Tauri stars (Preibisch et al. 2005). The above upper limit should be considered as a representative value since the detection limit depends on many factors including the spectral properties of the source and extent to which the source is displaced off-axis.

At the sensitivity of our observation, we expect ≈1 extragalactic background source in a region spanning ≈1.6 arcmin², or <1 extragalactic source in the ≈40''× 40'' NGC 2071-IR region analyzed here. This estimate is based on hard-band (2–8 keV) number counts from the Chandra Deep Field (CDF) and other X-ray surveys (e.g., Brandt et al. 2001; Moretti et al. 2003) and an assumed power-law spectrum with a photon power-law index Γ = 1.4 for extragalactic sources. The above estimate should be considered as approximate because the accuracy of the log N–log S distribution for extragalactic sources based on deep-field observations when applied to lower galactic latitude star-forming regions is not well known. Getman et al. (2005) concluded that CDF number counts overestimated the number of extragalactic sources in the deep Chandra COUP observation of the Orion Nebula Cluster by a factor of ∼3–4.

CIAO psextract was used to extract source and background spectra for brighter sources, along with source-specific response matrix files (RMFs) and auxiliary response files (ARFs) used in spectral fitting. Light curves of brighter sources were also extracted using CIAO tools. We used the 3σ source ellipses from wavdetect to define the extraction regions and background was extracted from adjacent source-free regions. Background is negligible, contributing only 0.21 counts per pixel (0.3–7 keV) for the total exposure time, or less than one count in the regions used to extract source spectra and light curves. Spectral fitting and timing analysis were undertaken with the HEASOFT Xanadu¹¹ software package including XSPEC version 12.4.0. We also applied the Kolmogorov–Smirnov (KS) test (Press et al. 1992, p. 617) to unbinned photon arrival times to check for X-ray variability.

Since conventional spectral fitting is only practical for brighter sources, we used the quantile analysis approach of Hong et al. (2004) to quantify source hardness. This technique has the advantage that it can be applied to sources that are too faint for spectral fitting. The quantile method uses sorted photon event energy lists to compute the median energy E_50% and the quartile energies E_25% and E_75%. The quartile energies are defined such that one-fourth of the counts have energies below E_25%, and three-fourths below E_75%. The quantile fractions Q_x are computed from quartile energies using Q_x = (E_x% − E_lo)/(E_up − E_lo), where E_lo and E_up are the lower and upper energy bounds of the extracted events (0.3 and 7.0 keV in this work). Quantile values (or functions thereof) are then overplotted on a grid of neutral hydrogen absorption column density (N_H) and plasma energy (kT) values generated from spectral simulations based on a one-temperature (1T) APEC optically thin thermal plasma model. This provides a graphical representation of the source hardness superimposed on a (N_H, kT) grid that is akin to an X-ray color–color diagram.

We use IR observations of the NGC 2071-IR region obtained by the Spitzer Space Telescope to aid in X-ray source identification and to determine YSO classes using mid-IR colors. Spitzer's IRAC mapped the NGC 2071-IR region as part of Spitzer Orion Survey (GTO program ID 043; Megeath et al. 2005b; Allen et al. 2007; Flaherty & Muzerolle 2008; S. T. Megeath et al. 2009, in preparation). Further information on IRAC can be found in Fazio et al. (2004). The IRAC data were obtained in four channels (3.6, 4.5, 5.8, and 8.0 μm) using the high dynamic range (HDR) mode, acquiring a short 0.6 s exposure plus a longer 12 s exposure at each pointing position. Four dithers were obtained at each position.

The reduction and analysis of the complete data set will be discussed in detail by S. T. Megeath et al. (2009, in preparation). In brief, individual basic calibrated data (BCD) frames were mosaicked with custom software to create separate mosaic images for the four IRAC channels. Point sources were identified using PhotVis version 1.10 (Gutermuth et al. 2004, 2008). Using the source positions determined from the mosaics, photometry was measured for each point source on the individual BCD images using the IDL module aper.pro (Landsman 1993). Magnitudes were measured in a circular aperture of radius r = 2 pixels centered on the source, using a background annulus of r = 2–6 pixels. A gain correction depending on pixel position was applied to each magnitude. The final magnitudes are the median magnitudes returned from the four dithers. Magnitudes of those sources in the NGC 2071-IR region discussed here are given in Table 2, along with magnitude zero points and aperture correction factors.

Spitzer MIPS data at longer wavelengths 24–160 μm are also available (Flaherty & Muzerolle 2008; S. T. Megeath et al. 2009, in preparation). Because the NGC 2071-IR region is confused and heavily saturated, we only provide here the MIPS 24 μm photometry for IRS 7. MIPS data were reduced using the procedure given in Megeath et al. (2009).

IRS 1 was identified as a bright 10 μm source by Persson et al. (1981). It is also clearly visible in the 2MASS K_s image (Figure 1) and in Spitzer IRAC images (Figure 2). A very bright source at or near the position of IRS 1 saturates the MIPS 24 μm images. IRS 1 is likely the dominant contributor to the mid/far-IR luminosity toward the NGC 2071-IR region, which was determined to be L_tot = 520 L_☉ (d = 390 pc) by Butner et al. (1990). This luminosity suggests that IRS 1 may be an embedded early-type star, as discussed further in Section 5.1. But, its true nature is still somewhat of a mystery.

IRS 1 is clearly detected in all four IRAC channels and its SED rises sharply toward longer wavelengths over the 3.6–8.0 μm IRAC range (Figure 5). It is the brightest IR source in the NGC 2071-IR core region at 8.0 μm (Table 3). Spitzer's angular resolution of FWHM ≈ 17 (3.6 μm) is unable to fully resolve VLA 1 located 22 north of IRS 1. But, the IRAC images do show what appears to be faint emission extending northward of IRS 1, providing a clue that VLA 1 may be weakly contributing to the mid-IR flux measured inside the r = 24 aperture centered on IRS 1. Higher angular resolution mid-IR observations will be needed to confirm that VLA 1 is indeed a mid-IR source.

The IRAC colors of IRS 1 are characteristic of a class I protostar (Figure 4). Adopting the YSO classification scheme of Gutermuth et al. (2008), class I sources are defined by [4.5] − [5.8] > 1, where brackets denote magnitudes at the enclosed wavelength (μm). IRS 1 has [4.5] − [5.8] = 1.67 (Table 3), so the class I criterion is satisfied. A similar conclusion is reached using the YSO classification schemes of Allen et al. (2004), Hartmann et al. (2005), and Megeath et al. (2005a).

Chandra detects an X-ray source that is offset by only 042 from the near-IR position of IRS 1 (=2MASS J054704.77+002142.8). This offset is within Chandra positional uncertainties and we thus associate X-ray source CXO J054704.75+002142.9 with the IR source IRS 1. The IRAC 3.6 μm detection of IRS 1 has a centroid position J054704.76+002143.0, which is offset by only 018 from the Chandra X-ray position. This Chandra detection confirms that IRS 1 is an X-ray source, as suspected on the basis of the previous lower spatial resolution XMM-Newton observation (S07).

The X-ray light curve of IRS 1 (Figure 9) shows low-amplitude fluctuations about the mean at the ±2σ level but no large flares. The KS test gives a probability of constant count rate P_const = 0.11. Thus, low-level variability may be present but is not proven at high confidence levels.

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IRS 1 is clearly detected in the hard-band ACIS-I image (Figure 7) but its emission is not as hard as that of VLA 1 or IRS 3 as gauged by median photon energy (Table 1). The X-ray spectrum of IRS 1 (Figure 10) is absorbed below 1 keV but the absorption is not as high as that of VLA 1. The ACIS-I spectrum of IRS 1 shows an emission line from the Si xiii He-like triplet complex (1.839–1.865 keV), implying thermal emission. Maximum power for the Si xiii triplet is emitted at log T_max = 7.0 (K) so hot plasma is clearly detected. The spectrum also reveals a faint (7 counts) fluorescent emission line at ≈6.4 keV from neutral or low-ionization stages of Fe (Figure 11), as anticipated from XMM-Newton spectra. This line forms in "cold" material in the vicinity of the star that is irradiated by the hard X-ray source and is discussed further in Section 5.3.

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We attempted to fit the ACIS-I spectrum of IRS 1 with various emission models. All models included a fixed-width Gaussian line at 6.4 keV to reproduce the faint fluorescent Fe line. A simple 1T APEC solar-abundance optically thin plasma model underestimates the flux in the 1.7–1.9 keV range and is unacceptable in terms of χ² statistics. This range includes the Si xiii He-like triplet lines at 1.839, 1.854, and 1.865 keV, which cannot be individually distinguished at the spectral resolution of ACIS-I. An acceptable 1T fit can be obtained by allowing Si to be overabundant by a factor of ∼3 with respect to its solar value. The apparent Si overabundance could either be a real physical effect due to factors such as grain destruction (e.g., via grain–grain collisions in shocks) or an unphysical consequence of fitting the spectrum with an overly simplistic 1T model.

An acceptable solar-abundance fit can be obtained with a two-temperature (2T) APEC model that includes both a cool and hot plasma component (Table 3). Some further improvement in the 2T fit can be obtained by allowing Si to be overabundant, but this is not required since the solar-abundance fit is already statistically acceptable (χ²_red = 1.01). Because of the heavy absorption below 1 keV, the temperature and emission measure of any cool component are quite uncertain. The 2T APEC model summarized in Table 3 uses a fixed temperature kT₁ = 0.55 keV for the cool component. This temperature results in a minimum value of the C-statistic (Cash 1979) when fitting unbinned spectra. However, there is very little change in the C-statistic for values in the range kT₁ = 0.25–0.70 keV. The 2T model yields a best-fit temperature for the hot plasma component kT₂ = 3.3 keV, but the true temperature could be higher since the 90% confidence upper limit is not tightly constrained by the data. Using the conversion N_H = 2.22 × 10²¹A_V cm⁻² (Gorenstein 1975), the solar-abundance 2T APEC model gives A_V = 16 [12–20] mag, where square brackets enclose the 90% confidence range. The slightly different conversion N_H = 1.6 × 10²¹A_V cm⁻² (Vuong et al. 2003) yields A_V = 22 [16–28] mag. The standard conversions used above are reliable for low-extinction sources, but some studies have questioned their accuracy for deeply embedded YSOs (Winston et al. 2007).

The above APEC model assumes that the X-ray emission arises in an optically thin thermal plasma and is typically used to model coronal emission in magnetically active late-type stars. However, other emission processes might be operating in IRS 1. In particular, X-rays could be produced in shocks associated with its powerful outflow. Thus, we tried to fit the IRS 1 spectrum with an isothermal plane-parallel shock model (VPSHOCK). The fit statistic for VPSHOCK (χ²_red = 1.16) is not quite as good as the 2T APEC model and the fit converges to shock temperatures kT ≈ 2–3 keV (Table 3). Such shock temperatures are realistic but do require high-velocity flows, as discussed further in Section 5.2.1.

In summary, a 1T APEC thermal plasma model is able to reproduce the Chandra spectrum of IRS 1 provided that Si is allowed to be overabundant. An acceptable solar-abundance fit can be obtained with a 2T APEC model, but it does not tightly constrain the temperature or emission measure of the heavily absorbed cool component. This model requires hot plasma at kT₂ ≈ 3 keV and an absorption column density equivalent to A_V ≈ 16–22 mag.

VLA 1 is a compact 3.6 cm radio continuum source revealed in high-resolution VLA A-configuration images obtained by Carrasco-González et al. (2007). It was also marginally detected in previous 5 GHz VLA images shown in Snell & Bally (1986). The 3.6 cm radio source is located 22 north of IRS 1, which equates to a projected separation of ≈860 AU at d = 390 pc.

The Chandra X-ray position is in very good agreement with the radio position of VLA 1 (Table 1). Due to its close proximity to IRS 1, VLA 1 was not spatially resolved from IRS 1 by XMM-Newton. We are not aware of any published high-resolution IR images which show that VLA 1 is an IR source. But, Spitzer images (Figure 2) reveal faint northward extension from IRS 1 that may be due to VLA 1.

The KS test shows that the X-ray emission of VLA 1 is likely variable with P_const = 0.02, though the light curve shows no large amplitude flares (Figure 9). The ACIS-I spectrum is heavily absorbed with little or no X-ray emission detected below 2 keV. The spectrum can be acceptably fitted with a 1T APEC model at kT = 2.2 keV (Table 2). No significant Fe line emission is detected in the 6.4–6.7 keV range. The inferred absorption column density log N_H = 22.94 gives an equivalent extinction A_V = 39 [22–69] mag using the Gorenstein (1975) conversion. Based on X-ray spectral fits, VLA 1 is more heavily absorbed than IRS 1 and Figure 8 indicates that it is the most heavily absorbed X-ray source in NGC 2071-IR. VLA 1 could thus be in an earlier evolutionary stage than IRS 1. Because of the high absorption, the intrinsic (unabsorbed) X-ray luminosity of VLA 1 is quite uncertain. But, the value inferred from the 1T APEC model (Table 2) is consistent with that of a ∼solar-mass pre-main-sequence star or protostar, based on the known correlation between L_X and stellar mass in low-mass YSOs (Preibisch et al. 2005; Telleschi et al. 2007a).

Previous studies have concluded that IRS 2 is likely a YSO (W93; T07) and Aspin et al. (1992) argued that it is driving a bipolar outflow. IRAC images show some indication that the source is extended. IRAC colors (Figure 4) indicate that IRS 2 is either a class I protostar or a heavily reddened class II object. However, some caution is warranted in interpreting IRAC colors. High-resolution 3.6 cm radio continuum observations reveal two closely spaced components A and B separated by ≈14 with the brighter component A at radio position VLA J054705.367+002150.5 (C. Carrasco-González et al. 2009, private communication). Spitzer cannot resolve such a close pair and it is thus possible that the IRAC colors are the superposition of two sources.

Chandra detects a very faint X-ray source (4 net counts) at an offset of only 027 from the brighter radio component A (Table 1). Because of the good positional coincidence with the radio source, this weak X-ray detection is likely real but is of marginal significance. There are too few counts to do any substantive X-ray analysis but the high median photon energy E₅₀ = 5.2 keV suggests that this source is exceptionally hard or very heavily absorbed.

The IR source IRS 3 is about 1 mag fainter at K band and at 3.6 μm and 4.5 μm than IRS 1. Its K-band emission is highly polarized (W 93; T07). A 1.3 cm radio continuum source lies near IRS 3 (Torrelles et al. 1998), but the near-IR peak is offset slightly to the northeast relative to the VLA radio position (T07). As Table 1 shows, the weak X-ray source detected by Chandra is in excellent positional agreement with the 1.3 cm radio source detected by Torrelles et al. (1998). The nearest 2MASS source lies 23 from the X-ray peak. This offset is significant and supports the conclusion of T07 that the near-IR emission is nebulosity. The object detected with Chandra and the VLA is likely the star that illuminates the near-IR nebula. Spitzer IRAC images show an IR source near IRS 3 (Figure 2). The IRAC position as measured in the 4.5 μm (ch2) image is offset by only 042 from Chandra. Thus, it appears that the self-luminous source is also detected by Spitzer. We do not have complete IRAC photometry for IRS 3 because it lies in the bright wings of IRS 1.

The KS test shows no significant X-ray variability. The distribution of event photon energies (Figure 12) shows clearly that this is a hard absorbed source, and the energy quantile plot (Figure 8) substantiates this. There were too few counts to analyze a binned spectrum but we did attempt to fit the unbinned spectrum with a 1T APEC thermal plasma model using C-statistics (Cash 1979). The absorption was held fixed at log N_H = 22.79 cm⁻², based on the estimated extinction A_V = 28 mag (W93). The fit converges to a high but uncertain temperature kT ≳ 10 keV, as also indicated by the quartile plot (Figure 8).

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Visual inspection of the unbinned spectrum shows a weak buildup of counts near 6.4 keV, hinting that a weak fluorescent Fe line may be present. The IRS 3 event list reveals five events within the 6–7 keV range with event energies 6.11, 6.20, 6.42, 6.53, and 6.69 keV. The 6.42 and 6.53 keV events could be due to fluorescent Fe. We added a Gaussian line component to the 1T APEC model at a line energy of 6.4 keV, fixing the line width at FWHM = 120 eV, corresponding to the ACIS-I instrumental width at this energy (Chandra POG, Figure 6.8). With the Gaussian line included, the C-statistic C = 133.1 is slightly less than without the line (C = 138.2), and the fit gives log N_H = 22.79 cm⁻², kT = 10 keV, and an unabsorbed flux F_X,unabsorbed (0.5–7.5 keV) = 1.6 × 10⁻¹⁴ erg cm⁻² s⁻¹. This equates to log L_X = 29.5 erg s⁻¹ (d = 390 pc), which is typical of low-mass YSOs in Orion (Preibisch et al. 2005). Thus, the event list and spectral fits provide some marginal evidence for fluorescent Fe, but a deeper exposure would be needed to make a compelling case for a 6.4 keV line detection.

The high-resolution VLA 1.3 cm radio continuum images of IRS 3 obtained by Torrelles et al. (1998) show elongation in the ≈N–S direction, and elongated morphology is also visible in the more recent 3.6 cm VLA images of Carrasco-González et al. (2007). The extension was interpreted as a thermal radio jet, adding support to the belief that IRS 3 is driving an outflow (Section 1). The orientation of H₂O maser spots approximately perpendicular to the jet axis and maser velocity data led Torrelles et al. (1998) to conclude that the masers trace a compact rotating molecular disk surrounding a low-mass (∼1 M_☉) YSO. Such a disk viewed in a nearly edge-on geometry would act as a strong absorber of soft X-ray photons and would provide a natural explanation for the high N_H (Figure 8) and lack of detected low-energy X-ray photons below ∼2 keV in IRS 3 (Figure 12).

Tamura et al. (2007) concluded that IRS 4 is an IR cluster member because it is surrounded by a compact reflection nebula and is associated with faint H₂ emission. The faint X-ray source detected by Chandra is in good positional agreement with 2MASS and shows no significant variability. It has the lowest median and mean photon energy of the six X-ray sources detected in the cluster core. There are insufficient counts available for spectral analysis. About half of the detected photons have energies below 2 keV. Quantile analysis shows that this source has lower absorption than the other X-ray detections in NGC 2071-IR (Figure 8). Near-IR colors (Figure 3) are consistent with a star having intrinsic cTTS color but viewed through substantial reddening (A_v ≈ 13 mag). IRAC colors (Figure 4) are consistent with a reddened class II YSO but the SED (Figure 5) is nearly flat in the IRAC bands and is more suggestive of a flat-spectrum source.

The near-IR morphology of this IR source changes with time and W93 concluded that the near-IR emission could be from a jet, a moving clump of shocked gas, or nebulosity. IRS 5 is located between the two nearby bright mid-IR sources IRS 1 and IRS 2. We were not able to identify a definite point source at the IRS 5 position in the IRAC images, and thus cannot provide any reliable IRAC photometry. No significant X-ray emission was detected so we find no evidence of a young star and our results are thus consistent with the nonstellar classification proposed by W93.

Near-IR observations have shown that IRS 6 is a binary with a 2'' separation (T07). We detect a variable X-ray source that is offset to the northwest of the 2MASS position by 168 at position angle PA = 315° (Table 1). This offset is significant and implies that the 2MASS near-IR source is not the same source detected in X-rays. Spitzer images show an IRAC source at an offset of 119 from the Chandra position. IRS 6 has similar near-IR and mid-IR colors to IRS 4, and both the near-IR and mid-IR colors of IRS 6 are consistent with a reddened class II object (cTTS).

The spectral fit of IRS 6 with a 1T APEC thermal model (Table 2) suggests a rather hot source with plasma temperature kT = 3.85 keV and moderate absorption N_H ∼ 10²² cm⁻². Similar values are inferred from the quantile plot (Figure 8), but both kT and N_H have rather large uncertainties due to the low number of X-ray counts. The absorption from the IRS 6 spectral fit (Table 2) gives an equivalent extinction A_V = 5.6 [1.7–10.4, 90% conf.] mag (Gorenstein 1975 conversion), which is less than IRS 1 and VLA 1.

On the basis of its relatively high kT and variability, we conclude that the X-ray source lying 168 from the 2MASS position is a magnetically active star. Because this system is a close near-IR pair, higher spatial resolution observations are needed to confirm the IR colors and more accurately determine the relative positions of the two near-IR sources and the X-ray source.

Tamura et al. (2007) argued on the basis of polarization angle that circumstellar material is present around IRS 7. IRAC colors support this conclusion, indicating either a class I protostar or a heavily reddened class II source (Figure 4). The object is visible in MIPS 24 μm images and the SED continues to rise out to 24 μm (Figure 5). Its [4.5] − [24] color is consistent with a protostar (Megeath et al. 2009). We detect no X-ray emission at or near IRS 7, so it is either intrinsically X-ray faint or its emission is reduced to undetectable levels by high absorption (e.g., a nearly edge-on disk or dense surrounding envelope).

IRS 8 and 8a are two faint closely spaced IR sources separated by 16 in R.A. and 07 in decl., with IRS 8a lying northeast of IRS 8 (W93; T07). Both objects are >1.5 mag fainter at K than the other core sources (T07). The closest 2MASS positional match to IRS 8 has K_s = 13.54 (2MASS J054704.45+002138.8) and there are no other 2MASS sources within 5''. IRS 8 is associated with a near-IR nebula and was classified as a self-luminous object, probably an embedded star, by W93. IRS 8a is thought to be a shocked molecular hydrogen peak (W93).

The IRAC images reveal what appears to be extended structure or nebulosity in the NE-SW direction near IRS 8. A faint emission peak is visible close to the IRS 8 position at 3.6 μm (Table 2) and the peak shifts slightly to the northeast in the 4.5 μm image, which is sensitive to shock emission. The close pair IRS 8/8a is not clearly resolved at IRAC's resolution, but the slight shift of the 4.5 μm emission peak toward IRS 8a may be an indication that it is a stronger shock-excited source. Because of the nebulous morphology in IRAC images, we were not able to obtain any reliable IRAC photometry. There is no X-ray emission at or near the 2MASS position of IRS 8 or at IRS 8a, and consequently no X-ray evidence of an embedded star.

The heavily reddened IRAC colors of IRS 1 and its rising SED in the IRAC bands, along with its bright (saturated) detection in MIPS 24 μm images, strongly suggest that it is a class I protostar. Based on its association with a radio continuum source and luminosity considerations, W93 concluded that IRS 1 is most likely a B0–B5 star. Under the assumption that the 6 cm radio continuum emission of IRS 1 arises in a photoionized compact H ii region at a distance of 500 pc, Snell & Bally (1986) found that a single B2 star would be needed to produce sufficient Lyman continuum ionizing flux to account for the radio emission. Adjusting their distance downward to 390 pc (Anthony-Twarog 1982) would require a ∼B3 star. However, we note that the volume emission measure of the radio continuum source computed from their 6 cm radio flux density S_6 cm = 7.9 mJy is 2 orders of magnitude below that typically found for optically thin compact H ii regions (Equation (3) of Skinner et al. 1993). Thus, the radio continuum emission from IRS 1 could very well involve processes other than free–free emission from a compact H ii region.

The more recent high-resolution 3.6 cm VLA A-configuration images of IRS 1 obtained by Carrasco-González et al. (2007) reveal a complex source morphology for IRS 1. A compact core is present along with fainter emission extending in the E–W direction. More interestingly, compact knots are seen expanding away from IRS 1 in an approximate E–W direction with high velocities of up to ∼600 km s⁻¹, as determined from radio proper motions. These results, along with the positive radio spectral index, suggest that the radio emission includes a contribution from an extended ionized wind or jet. A possible analog is NGC 7538-IRS1, whose free–free emission is dominated by a highly collimated ionized jet (Sandell et al. 2009).

Using far-IR data obtained with the Kuiper Airborne Observatory (KAO), Butner et al. (1990) determined the total luminosity toward the NGC 2071-IR region to be L_tot = 520 L_☉ (d = 390 pc). This is about 15% greater than the value determined by Harvey et al. (1979) when normalized to a common distance of 390 pc. Because of the rather large ≈30'' KAO beam size, this total will include any contributions from other sources near IRS 1 and should thus be considered an upper limit for IRS 1 itself. If due to a single star, a value L_tot = 520 L_☉ would correspond to a ZAMS star of mass ∼5 M_☉ (Siess et al. 2000) and ZAMS spectral type ∼B4 (Thompson 1984). But, if the L_tot value is interpreted as an upper limit for IRS 1, a lower mass and later spectral type are possible. The above estimates are based on an assumed distance d = 390 pc. If IRS 1 lies at the far edge of the Orion OB1 association then a larger distance d ∼ 500 pc is possible (Section 1). In that case, one infers a slightly higher mass of ∼6 M_☉ and ZAMS spectral type of ∼B3–B4.

The X-ray luminosity and temperature of IRS 1 are within the broad range observed in young B-type stars. The study of 20 early-type stars in the Chandra Orion COUP sample by Stelzer et al. (2005) included 12 B stars ranging in spectral type from B0.5 to B9. Their X-ray luminosities spanned a wide range from log L_X = 29.4–32.4 erg s⁻¹, and a hot plasma component at typical temperatures kT_hot ≈ 1–3 keV was present in all B stars. Different B stars with similar spectral types showed dramatic differences in L_X and no support was found for earlier claims that L_X correlates with L_bol in OB stars. Using L_bol = 520 L_☉ for IRS 1 and the range of L_X values from the different models in Table 3, we obtain log(L_X/L_bol) = −5.7 ± 0.5. This is consistent with values obtained for the COUP B-star sample, but some caution is needed in interpreting this ratio for IRS 1 since its L_bol could be dominated by disk emission.

On the basis of the above comparisons, we conclude that a B-star classification for IRS 1 is plausible, but a mid-to-late B star is more likely than an early B star based on the luminosity determined by Butner et al. (1990). If IRS 1 is indeed an embedded intermediate-mass B star, it is tempting to speculate that it might be a Herbig Ae/Be star progenitor. Herbig Ae/Be stars are optically revealed intermediate-mass pre-main-sequence stars which in many cases show high-temperature X-ray plasma that is likely of magnetic origin. But, the question of whether their X-ray emission is intrinsic or instead due to coronal late-type companions remains open. For previous studies of X-ray emission from Herbig Ae/Be stars, the reader is referred to Damiani et al. (1994), Zinnecker & Preibisch (1994), Preibisch & Zinnecker (1996), Skinner & Yamauchi (1996), Skinner et al. (2004), Hamaguchi et al. (2005a), Stelzer et al. (2006, 2009), and Telleschi et al. (2007b).

Given that the true nature of IRS 1 is not well known and that it may be an embedded B-type star, some further discussion of the origin of its X-ray emission is warranted. Very few massive embedded stars have been detected in X-rays and little is known about their X-ray emission during early evolutionary stages prior to becoming optically visible. We thus consider several possible emission mechanisms below.

5.2.1. Shocked Winds and Outflows

5.2.2. Accretion Shocks

5.2.3. Magnetic Activity

5.2.4. Magnetically Confined Wind Shocks

5.2.5. Summary of Emission Mechanisms

A fluorescent Fe line at 6.4 keV can be produced by photoionization when "cold" material containing neutral or weakly ionized iron is irradiated by a nearby hard X-ray continuum source. Photons with energies above 7.11 keV are needed to eject a K-shell electron. The K-shell vacancy is filled by another electron such as an L-shell electron, resulting in fluorescent Fe line emission at 6.4 keV. The line is in fact a doublet consisting of two closely spaced lines at 6.391 keV and 6.404 keV, but these cannot be distinguished at the energy resolution of the ACIS-I CCDs. The above photoionization process is discussed in more detail by George & Fabian (1991) and Kallman et al. (2004). Analysis of the fluorescent Fe line can potentially provide information on the properties of the absorber and its location relative to the hard continuum source.

The fluorescent Fe line detected in IRS 1 is faint (Figure 11), consisting of seven events distributed in the 6.298–6.536 keV energy range, with a median energy 6.395 keV and mean energy 6.419 keV. Three of the seven photons arrived in the first half of the observation. A KS test applied to the arrival times of the seven events shows no significant variability and gives a probability of constant count rate P_const = 0.35. The line flux is 7.2 (±1.2) × 10⁻¹⁵ erg cm⁻² s⁻¹. This is about a factor of 5 less than measured during the 2005 March XMM-Newton observation (S07). The line equivalent width in the ACIS-I spectrum is uncertain because the continuum flux near 6.4 keV is very faint. But, our flux estimates give values EW ⩾ 2.4 keV, consistent with the previous XMM-Newton estimate (S07).

The detection of fluorescent Fe line emission in IRS 1 is a rare occurrence in YSOs, but certainly not unique. Tsujimoto et al. (2005) examined 1616 X-ray sources detected in the deep Chandra COUP observation of the Orion Nebula Cluster and identified seven sources with excess emission at 6.4 keV and line equivalent widths EW ⩽ 0.27 keV. All seven X-ray sources showed flare-like variability and high line-of-sight absorption N_H ≳ 10²² cm⁻², as also seen in IRS 1. The broad fluorescent Fe line during a superhot flare in V1486 Ori (COUP source 331) was analyzed by Czesla & Schmitt (2007). The 6.4 keV line has also been detected in other YSOs including the class I objects YLW 16A (Imanishi et al. 2001) and Elias 29 (Giardino et al. 2007) in ρ Oph, and the class I source R CrA X_E (= IRS 7B; Hamaguchi et al. 2005b).

The fluorescent Fe line detection in IRS 1 is interesting in three respects: (1) it was detected in the apparent absence of any large-amplitude flares (but low-level variability may be present), (2) there is no accompanying line at 6.7 keV from the highly ionized Fe xxv complex, and (3) the line has a very large equivalent width that is difficult to account for by illuminated disk models.

The absence of any detectable large flares in IRS 1 is unusual because the fluorescent Fe lines in almost all of the YSOs mentioned above occurred during flares. Analysis of the V1486 Ori data showed that its broad fluorescent Fe line appeared during the rise phase of a flare (Czesla & Schmitt 2007). An interesting exception is Elias 29, which showed a 6.4 keV line during a large flare. But the line persisted after the flare had decayed, remaining over a timescale much longer than the radiative lifetimes of fluorescent Fe transitions. Giardino et al. (2007) thus concluded that the X-ray flare was not directly responsible the post-flare 6.4 keV line. Instead, they proposed that the line was produced by collisional ionization from an unseen population of non-thermal electrons.

The 6.4 keV line in IRS 1 is similar to Elias 29 in the sense that it does not appear to be directly associated with a large flare. The line was present in both the XMM-Newton spectrum obtained in 2005 March and in the more recent Chandra spectrum, but no large flare was detected in either observation. If a large flare was involved, then it must have occurred prior to the start of the observations, or it escaped detection (as could occur if the flare occurred on the back side of IRS 1 and was occulted).

Unless large X-ray flares occurred and escaped detection, the existing data indicate that the 6.4 keV line in IRS 1 is persistent and does not correlate with large flares. In that case, some mechanism needs to be operating nearly continuously to ionize the absorbing material. At present, the ionizing source and the ionization mechanism are not known. The obvious candidate for the ionizing source is IRS 1 itself. But, in terms of hardness, its Chandra spectrum is rather benign, showing no large flares, no Fe xxv line and a weak continuum above ∼7 keV. However, we cannot rule out the possibility of substantial hard emission (kT ≳ 20 keV) above the energy range accessible to Chandra and XMM-Newton. Nevertheless, the intriguing possibility does remain that the hard ionizing source is not IRS 1 and is not directly detected in our X-ray observations.

However, before discounting IRS 1 as the ionizing source, two factors should be noted. First, low-level variations appear to be present in the Chandra light curve of IRS 1 (Figure 9). These variations may be a signature of persistent low-level flaring. If that is the case, then quasi-continuous hard energy release by repeated small flares may play a role in photoionizing the absorber. Second, IRS 1 is thought to drive a powerful outflow and also shows evidence of a radio jet (Carrasco-González et al. 2007). If a population of non-thermal electrons is produced in the shocked jet (which may be magnetically collimated) they could play a role in collisional ionization of the absorber. Collisional ionization has been invoked as a means of explaining the 6.4 keV line in Elias 29 (Giardino et al. 2007) and in the active binary system II Peg (Osten et al. 2007). If a population of non-thermal particles is present, they may imprint a signature on the radio continuum emission. Thus, further radio observations of IRS 1 at lower frequencies where non-thermal emission can dominate over thermal (free–free) emission could be informative.

The large equivalent width EW ⩾ 2.4 keV of the IRS 1 fluorescent line is a noteworthy feature of its X-ray spectrum. This value is extreme for a YSO, exceeding even the maximum value EW = 1.4 keV observed for V1486 Ori during the rise phase of its superhot flare (Czesla & Schmitt 2007). These values conflict with theoretical models of centrally illuminated cold disks, which predict EW ≈ 0.1–0.2 keV (George & Fabian 1991). The EW of the 6.4 keV line is proportional to the column density $N_{\rm H^{\prime }}$ in the fluorescing material (Tsujimoto et al. 2005). If taken literally, the large EW of IRS 1 would imply a high absorber column density $N_{\rm H^{\prime }} \gtrsim 10^{24}$ cm⁻² (S07). Since this value is much larger than the line-of-sight photoelectric absorption inferred from the X-ray spectrum (Table 3), one is led to the conclusion that the fluorescing material lies off the line-of-sight. But, the above conclusion rests on the assumption that the continuum measured in the X-ray spectrum has undergone negligible absorption. If the hard irradiating source is not IRS 1 itself and is occulted or very heavily absorbed, then some or all of the intrinsic hard continuum will be missing from the observed X-ray spectrum. In that case, the EW measurement would be based on a depressed continuum and thus overestimated, along with $N_{\rm H^{\prime }}$.

In summary, our main conclusions with respect to the fluorescent Fe line in IRS 1 are (1) the line is present even in the absence of the Fe xxv line (6.67 keV) and any large flare detections, and evidently does not directly depend on large-amplitude X-ray flares for its existence, (2) the line flux is variable, showing a decrease between the XMM-Newton and Chandra observations taken 2.6 years apart, but no significant variability during the 67 ks Chandra observation, (3) the line EW is much larger than predicted by centrally illuminated disk models, suggesting either that the photoionization models do not accurately reflect the line formation process and that other mechanisms (e.g., collisional ionization) could play a role, or that the observed continuum is not the same as the intrinsic continuum of the photoionizing source (e.g., due to the effects of absorption or occultation), and (4) the hard ionizing source may have so far escaped X-ray detection.