© ESO, 2012

1. Introduction

If the global scenario of the star formation history (SFH) of the Milky Way is still not fully outlined (see Wyse 2009, for a review), the recent SFH in the solar neighborhood is far from being constrained. In particular, the recent local star formation rate is poorly known because of the difficulty encountered in properly selecting young late-type stars in large sky areas from optical data alone (Guillout et al. 2009). This situation has improved thanks to wide-field X-ray observations, such as the ROSAT¹ All-Sky Survey (RASS), which allowed for efficiently detecting of young, coronally active stars in the solar vicinity (Guillout et al. 1999). Nearby star-forming regions (SFRs) have been searched for pre-main sequence (PMS) stars, and many new nearby moving groups and widely-spread young low-mass stars have been identified over the past decades based on the RASS data (Feigelson & Montmerle 1999; Zuckerman & Song 2004; Torres et al. 2008; Guillout et al. 2009, and references therein).

It has been suggested that at least some of these widely spread young stars might be associated with the so-called Gould Belt (GB; Guillout et al. 1998), a disk-like structure made up of gas, young stars, and OB associations (see, e.g., Lindblad 2000; Elias et al. 2009). However, the existence of this structure and its possible origin remain somewhat controversial (Sánchez et al. 2007, and references therein). It is also unclear whether its putative young stellar population exceeds the predictions of Galactic models, which would indicate a recent episode of star formation. On the other hand, there is no evidence of molecular material within  ~100 pc that can explain the origin of the distributed young stars as in situ star formation. These young stars may therefore represent a secondary star formation event in clouds that condensed from the ancient Lindblad ring supershell, or may have been formed in recent bursts of star formation from already disrupted and evaporated small clouds (see Bally 2008, and references therein).

An accurate census of the young population in different star-forming environments and the comparison between observations and predictions from Galactic models are therefore required to constrain the low-mass star formation in the solar vicinity.

In this work, we revisit the results by Walter et al. (2000) on the large-scale spatial distribution of 6482 RASS X-ray sources in a 5000 deg² field centered on the Orion SFR, and use optical, low- and high-resolution spectroscopic observations to investigate the optical counterparts to 90 RASS sources (listed in Table 4) distributed on a sky “strip” perpendicular to the Galactic plane, in the general direction of Orion. The complete sky coverage of the RASS allows us to perform an unbiased analysis of the spatial distribution of X-ray active stars. Our main goal is to characterize spectroscopically the optical counterparts of the X-ray sources and compare their space density with predictions of Galactic models.

In Sect. 2 we define the sky areas and describe the selection of young X-ray emitting star candidates; in Sect. 3 we present the spectroscopic observations; in Sect. 4 we characterize the young optical counterparts in terms of spectral type, lithium detection, Hα line, radial/rotational velocity, and iron abundance; in Sect. 5 we compare the spatial distribution of the young X-ray emitting star candidates with expectations from Galactic models; in Sect. 6 we discuss the results and draw our conclusions.

2. Selection of study areas and X-ray sources

Our analysis is based on the method introduced by Sterzik et al. (1995) for selecting young star candidates from the RASS sources using the X-ray hardness ratios and the ratio of X-ray to optical flux. Figure A.1 is a revisited version of Fig. 3 by Walter et al. (2000), showing the color-coded space density of the RASS sources based on the α parameter, which is related to the probability of a source to be a young X-ray emitting star. Following the Sterzik et al. (1995) selection criteria, X-ray sources with α > 0.6 are most likely young stars with a positive rate of 80%. In the figure we recognize i) the surface density enhancements corresponding to 1483 young candidates, distributed over an area much larger than the molecular gas; ii) space density enhancements not necessarily associated with previously known regions of active (or recent) star formation (e.g., at α = 5^h34^m, δ =  + 22°01′ and at α = 4^h30^m, δ =  −08°); the local enhancement at α = 5^h07^m and δ =  −03°20′ corresponds to the L1616 cometary cloud (Alcalá et al. 2004; Gandolfi et al. 2008); iii) a broad lane apparently connecting Orion and Taurus, which extends further southeastward; this wide and contiguous structure is not symmetric about the Galactic plane, but instead seems to follow the GB as drawn by Guillout et al. (1998); iv) the surface density of young star candidates drops to a background value of about 0.1 candidate star/deg² near b_II = 0°, or below that value at higher Galactic latitudes.

Fig. 1 Large-scale spatial distribution of our targets (circles) and those of Alcalá et al. (1996, 2000; triangles). Empty symbols refer to low-resolution spectra, while filled symbols represent objects observed at high resolution. The dotted lines mark the strip, the long-dashed line is the Galactic plane, while the dash-dotted lines represent the Gould Belt disk, as outlined by Guillout et al. (1998). The CO J = 1 → 0 emission contour maps by Dame et al. (2001) of the Orion, Monoceros, and Taurus molecular clouds, and the λ Orionis H ii region are also overlaid.

In order to study the global properties of the different populations close to the Orion SFR, we selected a 10° × 75° strip (see Fig. A.1) perpendicular to the Galactic plane (hence presumably crossing the GB), and in the Orion vicinity. Inside this  ~ 750 deg² strip (${\mathit{l}}_{\mathrm{II}}^{\min }\mathrm{=}\mathrm{195}°$, ${\mathit{l}}_{\mathrm{II}}^{\max }\mathrm{=}\mathrm{205}°$, ${\mathit{b}}_{\mathrm{II}}^{\min }\mathrm{=}\mathrm{-}\mathrm{60}°$, ${\mathit{b}}_{\mathrm{II}}^{\max }\mathrm{=}\mathrm{15}°$), 806 RASS X-ray sources were detected with a high-confidence level by the Standard Analysis Software System (SASS; Voges 1992). According to the Sterzik et al. (1995) selection criteria, 198 of these sources are young star candidates. Additionally, we selected a region of 10° × 10° (see Fig. A.1) centered at α = 5^h34^m and δ =  + 22°01′, where a density enhancement of young star candidates is present. We call this previously unrecognized enhancement “X-ray Clump 0534+22” (hereafter, X-Clump 0534+22). Analogously, we also identify as “X-ray Clump 0430 − 08” (hereafter, X-Clump 0430 − 08) the unknown density enhancement at α = 4^h30^m and δ =  −08°. Although X-Clump 0534+22 almost coincides in direction with the Crab nebula, it is physically unrelated to the supernova remnant.

3. Spectroscopic observations and data reduction

We conducted a spectroscopic follow-up of 90 young stellar candidates inside the strip and clump regions. Table 1 gives a summary of the spectroscopic observations, while the full list of the observed stars is reported in Table 4. Throughout the paper, we call “young star” an optical counterpart that shows characteristics typical of weak-lined T-Tauri stars, i.e. weak Hα emission (EW_Hα ≲ 10 Å) and strong lithium absorption.

3.1. Low-resolution spectroscopy

Low-resolution spectroscopic observations were carried out during 25–30 November 1995 and 16–21 December 1996 using the Boller & Chivens (B&Ch) Cassegrain spectrographs attached to the 1.5 m telescope of the European Southern Observatory (ESO, Chile) and to the 2.1 m of the Observatorio Astronómico Nacional de San Pedro Mártir (OAN-SPM, Mexico), respectively. Table 1 gives information on the instrumental setups and number of observed objects. The spectral resolution was verified by measuring the full-width-at-half maximum of several lines in calibration spectra. The spectra were reduced following the standard procedure of MIDAS² software packages using the procedure described in Alcalá et al. (1996).

About sixty of the 198 strip sources plus fourteen stars in the X-Clump 0534+22 direction were investigated spectroscopically at low resolution (see Table 4 and Fig. 1). The observational strategy consisted in covering the whole range of right ascension and declination each night to avoid any bias in the resulting spatial distribution of the young star candidates. The spatially unbiased sample so far observed and characterized spectroscopically by us is therefore incomplete, representing  ~40% of the total sample of potential young X-ray emitting candidates. Yet it can be used to study the strength of the lithium absorption line within the RASS young stellar sample as a function of Galactic latitude, and to trace the young stellar population in the general direction of Orion.

3.2. High-resolution spectroscopy

High-resolution spectroscopic observations were conducted in several runs in the period between October 1996 and December 1998, using the Fiber Optics Cassegrain Échelle Spectrograph (FOCES) attached to the 2.2 m telescope at the Calar Alto Observatory (CAHA, Spain). Some seventy spectral orders are included in these spectra, covering the range from 4200 to 7000 Å, with a nominal mean resolving power of λ/Δλ ≈ 30   000 (see Table 1). The reduction was performed using IDL³ routines specifically developed for this instrument (Pfeiffer et al. 1998). Details on the data reduction are given in Alcalá et al. (2000).

We also retrieved FEROS⁴ spectra for two stars of our sample (namely, 2MASS J03494386 − 1353108 and 2MASS J04354055−1017293) from the ESO Science Archive⁵. The FEROS spectra extend between 3600 Å and 9200 Å with a resolving power R = 48 000 (Kaufer et al. 1999). The data were reduced using a modified version of the FEROS-DRS pipeline (running under the ESO-MIDAS context FEROS), which yields a wavelength-calibrated, merged, normalized spectrum, following the steps specified in Desidera et al. (2001).

In summary, high-resolution spectroscopy exists for 61 stars, 33 inside the strip, 13 in the X-Clump 0534+22, 8 associated with the high space-density X-Clump 0430 − 08 (see in Fig. A.1 the spatial distribution at α = 4^h30^m and δ =  −08°), and 7 inside the strip associated to the L1616 clump. The high-resolution sample thus represents  ~70% of the known lithium stars in the strip, and can be used to verify the reliability of the lithium strength obtained from the low-resolution spectra.

In order to increase the statistics, we also included in our analysis the on-strip lithium-rich stars identified by Alcalá et al. (1996) and Alcalá et al. (2000), which were observed with the same instruments, set-ups, and observational strategy.

4. Characterization of the selected young X-ray counterparts

4.1. Near-IR color–color diagram

We identified all our targets in the Two Micron Sky Survey (2MASS) catalog. In Table 4 we list the RASS Bright Source Catalog (Voges et al. 1999) and 2MASS designations, and an alternative name. The stellar coordinates are those from the 2MASS catalog. We then examined the properties of the sources in the JHK bands using their 2MASS magnitudes (Cutri et al. 2003) looking for possible color excess. The (J − H) versus (H − K) diagram (Fig. 2) shows that our sample mostly consists of stars without near-IR excess, with the only exception of 2MASS J05122053−0255523 = V531 Ori, which was classified as a classical T Tauri star by Gandolfi et al. (2008). All stars follow the MS branch with a spread mostly ascribable to photometric uncertainties. Only 2MASS J04405981 − 0840023 departs from the sequence, maybe because of its double-lined binary nature (Covino et al. 2001).

Fig. 2 2MASS (J − H) versus (H − K) diagram for the observed sources. Squares, asterisks, and stars represent objects in the three clumps, as indicated in the legend. The circles refer to the “non-clump” stars. The solid curve shows the relation between these indexes for main sequence stars (lower branch; Bessell & Brett 1988) and giants (upper branch; Kenyon & Hartmann 1995), where the 2MASS color trasformations were used. The AV = 1 mag reddening vector is shown with an arrow. The typical 2MASS photometric errors are overplotted on the lower-left corner of the panel.

4.2. Spectral types, lithium detection, and Hα equivalent width

Spectral types were determined from the low-resolution spectra by comparison with a grid of bright spectral standard stars (from F0 to M 5) observed with the same dispersion and instrumental set-up in each observing run. The methods described in Alcalá et al. (1995) were used for the classification, leading to an accuracy of about  ± 1 sub-class in most cases. The spectral types are reported in Table 4, while their distribution is plotted in Fig. 3. The sample is composed of late-type stars with a distribution peaked around G9–K1.

Fig. 3 Spectral type distribution of the young stellar candidates.

In our low-resolution spectroscopic follow-up, the Orion stars fall basically into four categories (see Table 4): i) stars with weak Hα emission (−3 Å  ≲ EW_Hα ≲ 0 Å) and Li absorption (18 stars); ii) stars with Hα filled-in or in absorption and Li absorption (42 stars); iii) stars with Hα in emission but no Li absorption (eight stars); and iv) stars with Hα in absorption but no Li detection (four stars). In total, 60 stars with clear lithium detections were found throughout the strip and in the clumps. Practically all lithium stars have spectral types ranging from late F to K7/M0 peaking around G9 (see Table 4). The effective temperature versus spectral type relation for dwarfs by Kenyon & Hartmann (1995) was used to estimate the log T_eff values listed in Table 4.

At this point, it is important to stress that i) at the Orion distance, our sample is limited to masses  >0.8 M_⊙ (Alcalá et al. 1998) because of the RASS flux limit; ii) the Li equivalent width (EW_Li) may be overestimated in low-resolution spectra because of blending mainly with the nearby Fe i line at λ = 6707.4 Å (see Sect. 4.3). This problem can be overcome by using high-resolution spectroscopy (see Sects. 3.2 and 4.3).

Fig. 4 Lithium (λ6707.8 Å) equivalent width versus log Teff for the stars in the strip (195° < lII < 205°), observed at low (left panel) and high (right panel) resolution. Three different bins of Galactic latitude were considered (see also Fig. 1). The upper envelope for the Pleiades sample, as adapted from Soderblom et al. (1993), is overplotted as a dashed line. Mean errors in $\mathit{E}{\mathit{W}}_{\mathrm{Li}}^{\mathrm{lr}}$ (~60 mÅ), $\mathit{E}{\mathit{W}}_{\mathrm{Li}}^{\mathrm{hr}}$ (~10 mÅ), and log Teff (~100 K, as obtained from the spectral synthesis) are overplotted on the lower-right corner of each panel. Open symbols represent the Alcalá et al. (1996) and Alcalá et al. (2000) results from low- (left panel) and high- (right panel) resolution spectra, respectively.

4.3. Lithium strength: low-versus high-resolution measurements

Fig. 5 Lithium abundance versus effective temperature for stars observed at high resolution (filled symbols refer to our sample, while open symbols represent the Alcalá et al. (2000) measurements). The “lithium isochrones” by D’Antona & Mazzitelli (1997) in the 2 × 106 − 1 × 108 yr range are overlaid with dashed lines. Dotted lines represent the limiting detectable log n(Li), as derived from the Pavlenko & Magazzù (1996) non-LTE curves-of-growth, assuming 10 mÅ as mean EWLi error and adding quadratically the mean contribution of the iron line at 6707.4 Å computed from the empirical relation obtained by Soderblom et al. (1993) between B − V color and EWFe (Kenyon & Hartmann 1995, tables were used to convert B − V colors into Teff). Left panel: stars located along the strip crossing the GB in the range 195° < lII < 205° (see Fig. 1) and three bins of Galactic latitude. Right panel: stars in known (L1616) and unknown (X-Clump 0430−08 and X-Clump 0534+22) groups.

For 43 stars we obtained both low- and high-resolution spectra. Based on these data, we estimated that the mean lithium equivalent width measured on low-resolution spectra ($\mathit{E}{\mathit{W}}_{\mathrm{Li}}^{\mathrm{lr}}$) is overestimated by  ~20 mÅ, with a standard deviation of  ± 60 mÅ on the average difference between the low- and high-resolution measurements. Interestingly, for several stars, the $\mathit{E}{\mathit{W}}_{\mathrm{Li}}^{\mathrm{lr}}$ matches the lithium equivalent width obtained from high-resolution spectra ($\mathit{E}{\mathit{W}}_{\mathrm{Li}}^{\mathrm{hr}}$) well and, in some cases, $\mathit{E}{\mathit{W}}_{\mathrm{Li}}^{\mathrm{lr}}$ values are underestimated (see Table 4). The good match between the $\mathit{E}{\mathit{W}}_{\mathrm{Li}}^{\mathrm{lr}}$ and $\mathit{E}{\mathit{W}}_{\mathrm{Li}}^{\mathrm{hr}}$ values can be attributed to the lithium strength of these stars, which is indeed high and also to the experience we gathered on measuring EW_Li in low-resolution spectra.

Fig. 6 Lithium abundance versus vsini for stars observed at high resolution. Symbols as in Fig. 5. Bottom panels show for each group the log n(Li) dispersion as a function of vsini at the following bins: 0–30, 30–60, 60–90, 90–120 km s-1. The values are shifted in vsini to better visualize the symbols.

In Fig. 4 (left panel), the lithium equivalent widths of the 47 stars in our sample and the 38 from Alcalá et al. (1996), observed at low-resolution and falling inside the strip, are plotted versus the logarithm of the effective temperature, for the following three bins of Galactic latitude: i)  −20° < b_II <  −10°, coinciding with the Orion complex; ii)  −30° < b_II <  − 20°, corresponding approximately to the position of the Gould Belt at those Galactic longitudes; and iii) in the two ranges  −10° < b_II < 15° and  −60° < b_II <  −30° (directions that we consider as likely dominated by field stars). Similar plots were produced for the stars observed at high resolution (see right panel of Fig. 4). The upper envelope for the Pleiades stars, adapted from the Soderblom et al. (1993) data, is also overplotted. In both panels, a spatial segregation of lithium strength can be observed, which justifies using both high- and low-resolution EW_Li. Practically all stars in the Orion region are located above the Pleiades upper envelope, as expected. The majority of these stars are indeed very young. Note that many stars, located in the region of the hypothetical Gould Belt, also have strong lithium absorption but tend to be closer to the Pleiades upper envelope. Finally, the majority of the lithium field stars are located closer to/or below the Pleiades upper envelope, and most of them seem to have an age similar to the Pleiades or older. Nevertheless, a few of these field stars have lithium strengths comparable to those of stars in Orion or the Gould Belt direction, and seem to also be very young.

4.3.1. Lithium abundance

The lithium abundance (log n(Li)) was derived from the $\mathit{E}{\mathit{W}}_{\mathrm{Li}}^{\mathrm{hr}}$ and T_eff values and using the non-LTE curves-of-growth reported by Pavlenko & Magazzù (1996), assuming log g = 4.5. The main source of error in log n(Li) is the uncertainty in T_eff, which is about ΔT_eff = 150 K. Taking this value and a mean error of about 10 mÅ in EW_Li into account, we estimate a mean log n(Li) error ranging from  ~0.22 dex for cool stars (T_eff ≈ 3700 K) down to  ~0.13 dex for warm stars (T_eff ≈ 5800 K). Moreover, the assumption of log g = 4.5 affects the lithium abundance determination, in the sense that the lower the surface gravity, the higher the lithium abundance. In particular, the difference in log n(Li) may rise to  ~0.1 dex, when considering stars with mean values of EW_Li = 300 mÅ and T_eff = 5000 K and assuming Δlog g = 1.0 dex. Hence, this means that our assumption of log g = 4.5 would eventually lead us to underestimate the lithium abundance.

In Fig. 5 (left panel), we show the lithium abundance as a function of the effective temperature for the stars on the strip, but coded in three bins of Galactic latitude, hence, according to their spatial location with respect to the Orion SFR and the Gould Belt. We also overplot the isochrones of lithium burning as calculated by D’Antona & Mazzitelli (1997). Three groups of stars can be identified according to the lithium content and spatial location. First, stars showing high-lithium content with ages even younger than  ~2−5    ×    10⁶ yr, mostly located on or close to the Orion clouds; second, stars with lithium content consistent with ages  ~5 × 10⁶ − 1 × 10⁷ yr, supposedly distributed along the Gould Belt, and third, stars with lithium indicating a wide range of ages, but located far off the Orion SFR or the GB.

In Fig. 5 (right panel), we show the same plot, but for the three identified young aggregates, represented by three different symbols. The lithium content in the X-Clump 0430 − 08 corresponds to an age of about 2    ×    10⁶ − 2    ×    10⁷ yr, while for the L1616 group it indicates an age of 2 − 7 × 10⁶ yr, consistent with the Alcalá et al. (2004) and Gandolfi et al. (2008) findings. Finally, the lithium content of the stars in X-Clump 0534+22 indicates a relatively narrow age range of 2 − 10 Myr, which is consistent with the age inferred from the HR diagram when adopting a distance of 140 pc (see Sect. 5.2).

4.3.2. Rotational and radial velocity measurements

Stellar rotation may affect internal mixing, hence lithium depletion. A large spread in rotation rates may introduce a spread in lithium abundance, which is observed in young clusters (Balachandran et al. 2011, and references therein). To investigate this Li behavior in the stars of our sample, rotational (vsini) velocities were derived by using the same cross-correlation method as described in Alcalá et al. (2000).

In Fig. 6 (left panel) we show the lithium abundance versus vsini for the on-strip stars, coded with three different symbols, indicating their spatial location as defined in the previous section. Despite the low statistics, the typical log n(Li) − vsini behavior observed in young clusters, i.e. a larger spread in log n(Li) for lower vsini values, is apparent for the stars projected on Orion and the Gould Belt. To assign a confidence level to this trend, a one-side 2    ×    2 Fisher’s exact test⁶ was performed (Agresti 1992). For the test, we adopted 30 km s^-1 and 2.0 dex as dividing limits in vsini and log n(Li), respectively. We found a p-value of 0.54 as chance that random data would yield the trend, indicating a correlation probability of 46%. Hence, the low-number statistics prevents a rigorous demostration of the apparent trend shown in the plots. More measurements of log n(Li) and vsini for stars projected on Orion and the Gould Belt are needed to firmly establish the preservation of Li content at high vsini ( ≳ 25 km s^-1) values. The above log n(Li) − vsini behavior is even less evident, however, for the stars flagged as “field”. These stars show a spread in log n(Li) at all vsini values. The difference between the three stellar groups is supported by the different dispersion in the log n(Li) − vsini diagram (see left-bottom panel of Fig. 6).

Fig. 7 Radial velocity distribution for the on-strip stars (left panel) and in the young aggregates (right panel), respectively.

In Fig. 6 (right panel) the lithium abundance is plotted as a function of vsini for the stars in the young aggregates. Similar conclusions can be achieved as for the young stars projected on Orion and the Gould Belt. The log n(Li) − vsini behavior in the X-Clump 0534+22 aggregate is enhanced by the star with the lowest lithium abundance and low vsini. This star, namely 2MASS J06020094+1955290, shows basically the same activity level as the other targets (see EW_Hα in Table 4 and $\log \frac{{\mathit{f}}_{\mathit{X}}}{{\mathit{f}}_{\mathit{V}}}$ values in the histograms of the top panel of Fig. 10) and is indistinguishable from the other stars in the aggregate.

Stellar radial velocities (RV) were also measured by means of cross-correlation analysis. In Fig. 7 (left panel) we show the distribution of RV for the on-strip stars, also divided into three bins of Galactic latitude as above. While the field stars show a wide RV range, the RV distribution of the stars projected on Orion and the GB is peaked at values of  ~18 km s^-1, i.e. close to the tail at low RV of the Orion sub-associations (from 19.7 ± 1.7 km s^-1 for 25 Ori to 24.87 ± 2.74 km s^-1 for the ONC, to  ~30.1 ± 1.9 km s^-1 for OB1b; see, Briceño et al. 2007; Biazzo et al. 2009) or to the Taurus-Auriga distribution (16.03 ± 6.43 km s^-1; Bertout & Genova 2006).

In Fig. 7 (right panel) the RV distribution of the stars in the young aggregates is shown. With the exception of two stars in the X-Clump 0430 − 08, likely spectroscopic binaries, the RV of these groups is in the range 10–35 km s^-1, fairly consistent with Orion or Taurus. It is worth mentioning that the RV distribution of widely distributed young stars in Orion shows a double peak (Alcalá et al. 2000), which can be explained as caused by objects associated with different kinematical groups, likely located at different distances.

4.3.3. Iron abundance

Metallicity measurements were obtained following the prescriptions by Biazzo et al. (2011a,b) and the 2010 version of the MOOG⁷ code (Sneden 1973).

After screening the sample to select suitable stars (G0-K7 stars with vsini ≲ 20 km s^-1 and no evidence of multiplicity), a total of 11 stars in our sample, plus 8 stars from Alcalá et al. (2000) were analyzed for iron abundance measurements. In Table 2 we list the final results, together with effective temperature, surface gravity, and microturbulence (the number of lines used is also given in Cols. 6 and 8). An initial temperature value was set using the ARES⁸ automatic code (Sousa et al. 2007, 2010); initial microturbulence was set to ξ = 1.5 km s^-1, and initial gravity to log g = 4.0. The effective temperatures derived using this method and from spectral types (Sect. 4.2) agree within  ~200 K (i.e.  ~1.5 spectral sub-class) on average. For the stars observed with the FOCES and FEROS spectrographs, the values of the stellar parameters closely agree. It is worth noticing that the target 2MASS J05214684+2400444 was also analyzed by Santos et al. (2008; their 1RXS J052146.7), who derived stellar parameters (T_eff = 4921 ± 59 K, log g = 4.05 ± 0.29, ξ = 1.92 ± 0.07 km s^-1, vsini = 13 km s^-1) and an iron abundance ( [Fe/H]  =  −0.07 ± 0.07) that agree well with our determinations, thus excluding any significant systematic error due to different datasets (see Biazzo et al. 2011b).

In Figs. 8 and A.2 we show the iron abundance distribution for the stars with strong lithium absorption (16 stars representing the young population) and those without it (three stars representing the old field population). Two results are noticeable: i) the stars with high lithium content show a distribution with a mean [Fe/H]⁹  =  −0.02 ± 0.09, consistent with the value of  − 0.01 ± 0.03 for open clusters younger than 150 Myr and within 500 pc from the Sun (see Biazzo et al. 2011a, for details); ii) stars with no lithium absorption line show [Fe/H] values that are within the metallicity distribution of field stars in the solar neighborhood ( [Fe/H]  =  −0.10 ± 0.24; Santos et al. 2008). This would imply that the young stars in the Orion vicinity have solar metallicity, consistent with the distribution of the Galactic thin disk in the solar neighborhood (Biazzo et al. 2011b).

5. Discussion

5.1. RASS and Galactic models

How many young, X-ray active stars are expected inside the strip? A comparison of the number counts with predictions from Galactic models provides the basis for a quantitative analysis of source excesses (or deficits) in order to understand their origin. Because of the strong dependence of stellar X-ray emission on age, an X-ray view of the sky preferentially reveals young stars (ages  ≤ 100 Myr), in contrast to optical star counts, which only loosely constrain the stellar population for ages  ≤ 10⁹ yr. A Galactic X-ray star count modeling starts by adopting a Galactic model, including assumptions about the spatial and temporal evolution of the star formation rate and the initial mass function, and uses characteristic X-ray luminosity functions attributed to the different stellar populations. These models are able to predict the number of stars per square degree N(>S) with X-ray flux  >S, taking into account the dependence on Galactic latitude, spectral type, and stellar age. An elaborate Galactic model, including kinematics, is the evolution synthesis population model developed at Besançon (Robin & Crézé 1986), which computes the density and the distribution of stars as a function of the observing direction, age, spectral type, and distance.

Fig. 8 [Fe/H] distribution for stars observed with the FOCES spectrograph and representing the young population (dashed histogram) and the field population (filled histogram). The bar represents the mean value of  [Fe/H]  =  −0.01 ± 0.03 obtained for open clusters younger than 150 Myr within 500 pc from the Sun (see Biazzo et al. 2011a, and references therein).

Our X-ray synthetic model is based on the Besançon optical model and has been first applied to the analysis of the RASS stellar population by Guillout et al. (1996). Motch et al. (1997) successfully used this model in a low Galactic latitude RASS area in Cygnus and found a good agreement between observations and predicted number counts using the “canonical” assumption of a uniform and continuous star formation history in the solar vicinity. We note, however, that following the publication of Hipparcos results, the stellar density in the solar neighborhood was revised (lowered) in the Besançon model, thus propagating model predictions in the X-ray population. The apparent disagreement between observed and predicted number count (by  ~20%) can in fact be explained by the population of old close binaries (RS CVn-like systems), as suggested by Favata et al. (1988) and Sciortino et al. (1995). RS CVn systems for which the high magnetic activity level results from the synchronization of rotational and orbital periods can mimic young active stars and contaminate the young star population detected in soft X-ray survey (Frasca et al. 2006; Guillout et al. 2009). The eight stars with Hα emission, but no Li absorption identified by us (cf. Sect. 4.2) may represent this type of objects.

In Fig. 9, we compare the RASS stellar counterparts (taken from the Guide Star Catalog) in our field with the current X-ray Galactic model predictions using a cumulative distribution function log N(>S) − log S. We selected two fields: one centered on the Galactic plane (b_II = 0° ± 10°, 190° < l_II < 220°), and the other, southern, includes Orion and a section of the Gould Belt (b_II =  −20° ± 10°, 190° < l_II < 220°). The deviation from a power-law function for both data distributions at a count rate of  ≈ 0.03 cnts/s is related to the completeness limit of the RASS at this value. Comparing these two distributions, it is noticeable how the source density in the area containing Orion exceeds that of the Galactic plane by about 30 − 60% (at the highest count rates even by a factor of two). Considering that the population of old active binaries is not yet taken into account in our X-ray model, the predictions closely agree with the RASS data around the Galactic plane and with optical counterparts from the Guide Star Catalog. This model predicts a surface density of 0.37 stars/deg² for the Galactic plane at the RASS completeness limit (see Table 3). The source excess in the Orion field can be attributed to additional, probably younger, X-ray active stars generated by recent, more localized star formation that is not included in the Galactic model. In fact, the difference in surface density between the regions in Orion + Gould Belt and the Gould Belt alone is 0.14 ± 0.04 stars/deg² for the RASS stellar counterparts and of 0.21 ± 0.07 stars/deg² for the stars with lithium detection, i.e. a significant number of young stars in the general direction of Orion is evident that did not necessarily originate in the star formation complex (see Table 3, where the errors were computed from Poisson statistics). We note that our results are not influenced by the density of active stars: from our spectroscopic identifications, we observed 60 stars in the strip out of 198 young star candidates resulting from the Sterzik et al. (1995) selection criteria. Eight are likely active stars because of their Hα emission but no Li detection (see Sect. 4.2). Therefore, inside the 750 deg² strip area, we estimate a surface density of only 0.03 stars/deg² as due to active stars.

Encouraged by the success of our X-ray star count modeling in reproducing the background counts and in revealing the excess of sources associated with Orion, we then performed a more detailed analysis of the RASS sources located in the strip shown in Fig. A.1, including the available information on spectroscopic identification and lithium abundance. Our goal is to intercompare average source densities of RASS-selected young stars and to constrain the age distributions in different parts of the strip. Therefore, we divided the strip in three subareas, one is a 200 deg² large field centered on the Galactic plane and with  −10° < b_II <  +10°, the adjacent 100 deg² area between  −20° < b_II <  − 10° contains the northern parts of the Orion complex, and the third 100 deg² area between  −30° < b_II <  − 20° is formally unrelated to the Orion molecular clouds but contains a significant part of the Gould Belt in that direction.

Fig. 9 Cumulative log N(>S) − log S distributions for RASS data in the Orion vicinity. A field centered on the Galactic plane is shown with a dotted line. The dashed line includes Orion and a section of the Gould Belt. The solid line corresponds to the prediction of our X-ray models at lII = 200° and zero Galactic latitude. The RASS completeness limit of 0.03 cnts/s is marked with the double arrow.

In Fig. 10 we compare the number of RASS sources that have stellar counterparts from the Guide Star Catalog and count rates  ≥ 0.03 cnts/s (as indicated by the thin-lined histogram) with the numbers of these sources predicted by the Galactic model (indicated by star symbols with statistical error bars). The histograms show the X-ray to optical flux ratio distributions. The thick-lined histogram indicates the RASS sources that have been selected as young star candidates according to the Sterzik et al. (1995) criteria. A subsample of those (hatched histogram) were observed spectroscopically and classified according to their lithium absorption strength. The dark gray histogram denotes objects where lithium absorption has been found, and the solid histogram refers to high-lithium stars. We note that the “identified” subsample is by now only complete to 40 − 70% depending on the area. We can draw the following main conclusions:

• 1.

  The total number and the flux ratio distribution of the RASSsources agree well with the Galactic model for the Galactic planefield. High-lithium sources represent about 1/3 of the observedsources. Extrapolating to all RASS-selected young starcandidates, we expect a total of about 15 high-lithium sources inthis 200 deg² large Galactic plane field. The Besançon model predicts 72 X-ray sources with ages younger than 150 Myr in this area. The high-lithium sources are expected to be the youngest within our sample, and we can roughly estimate their characteristic age to be about 15/72 × 150 Myr = 30 Myr (assuming continuous star formation). This age is consistent with their spectroscopic signature.

• 2.

  A large RASS source excess (~100%) is present in the Orion field. The fraction of high-lithium stars among identified sources is now about 9/10. In this 100 deg² field, we expect to find a total of 38 high-lithium stars once all young candidates have been observed.

• 3.

  A large RASS source excess (50%) is also present in the field that includes the Gould Belt, preferentially at flux ratios $\log \frac{{\mathit{f}}_{\mathit{X}}}{{\mathit{f}}_{\mathit{V}}}\mathit{>}\mathrm{-}\mathrm{2}$. The fraction of high-lithium stars among the observed sources is about 2/3. We expect eventually to find a total of about 17 high-lithium stars in the RASS-selected subsample.

These extrapolations to the total number of expected high-lithium sources should be reliable enough, since the spectroscopic identifications were performed on statistically representative subsamples of RASS-selected young star candidates. It is also possible that some of the high-lithium stars present among RASS sources have been missed by the Sterzik et al. (1995) selection criteria. Hence, the extrapolated number of high-lithium X-ray sources in the respective areas is a lower limit.

The results of this section are summarized in Table 3, where we report the estimated surface density of RASS sources and of high-lithium RASS-selected sources in the three investigated sky areas. While the Galactic plane source densities agree well with the Galactic X-ray star count model, the source density excesses in the other two directions indicate additional populations of high-lithium stars younger than 150 Myr.

Fig. 10 Detailed comparison of three fields in the strip perpendicular to the Galactic plane between lII = 195° and lII = 205°, as shown in Fig. 2. The three areas include: the Galactic plane, the Orion SFRs, and the Gould Belt, respectively. Histograms are cumulative, not stacked, i.e., each histogram includes the stars in all subcategory histograms.

5.2. The X-ray clumps 0534+22 and 0430−08: two new young stellar aggregates?

These two groupings were first recognized as surface density enhancements of X-ray emitting young star candidates (see Fig. A.1), and then classified as young star clumps based on the spectroscopic follow-up. In particular, from lithium abundance measurements, the stars in the X-Clump 0534+22 have an age of  ~5–10 Myr, while those in the X-Clump 0430−08 of  ~2–20 Myr. Inspection of recent CO maps (see Fig. 1) seems to indicate that the stars in the X-Clump 0534+22 follow the spatial location of some of the Taurus-Auriga dark nebulae (namely L1545, L1548, and L1549; Kenyon et al. 2008). Therefore, the X-Clump 0534+22 might be an extension of the Taurus-Auriga association toward the Galactic plane, although the average RV (~22 ± 8 km s^-1) of this clump is higher than the mean RV for Taurus, but within the quoted uncertainty. Therefore, it deserves additional kinematic studies. We attempted to estimate the stellar luminosities of the X-Clump 0534+22 stars by assuming the distance to Taurus to be (d ~ 140 pc; Kenyon et al. 1994), and adopting the Tycho magnitudes (V_Tycho; Egret et al. 1992), a solar bolometric magnitude of ${\mathit{M}}_{\mathrm{bol}}^{\mathrm{\odot }}\mathrm{=}\mathrm{4.64}$ (Cox 2000), and the temperatures listed in Table 4. The resulting HR diagram is shown in Fig. 11. The error in log L/L_⊙ is estimated considering an uncertainty of 0.1 mag in V_Tycho. The range of ages  ~3–30 Myr for the stars in the X-Clump 0534+22 appears to be consistent with the estimates from the lithium abundance (see Fig. 5). Adopting instead the Orion distance (~400 pc; Bally 2008) would place the stars close to or above the birthline in the HR diagram, in contrast with their lithium content and lack of IR excess.

Fig. 11 HR diagram for the stars in the X-Clump 0534+22. The pre-main sequence evolutionary tracks of Palla & Stahler (1999), together with their isochrones and birthline are shown by solid, dotted, and dash-dotted lines, respectively. Labels indicating the track masses and the isochronal ages (from 1 to 100 Myr) are also overlaid.

5.3. The variety of populations in the Orion vicinity

Based on the morphology and surface density distribution of the X-ray selected young star candidates, the spectroscopic analysis, and the comparison of RASS number counts with Galactic model predictions, we argue that the analyzed widespread stellar sample consists of a mixture of three distinct populations:

• 1.

  The clustered population comprises the dense regions( ≳ 5–10 stars/deg²) associated with sites of active or recent star formation (e.g., OB1a, OB1b, OB1c, λ Ori, L1616). These clusters contain the highest fraction of lithium-rich stars. The number counts result considerably exceed the Galactic model predictions. The stars in these clusters are, on average, among the youngest in our sample and their low spatial dispersion allows us to locate their birthplace. The RASS limit in X-ray luminosity at M ~ 0.8   M_⊙ is  ~log L_X = 30.5 erg/s; as a consequence, with some caution, we can extrapolate the X-ray luminosity function of RASS sources to lower luminosities (e.g.,  ~29 erg/s), and estimate the total number of X-ray emitting young stars in the Orion complex to be on the order of a few thousand.

• 2.

  The broad lane that apparently connects the Orion and Taurus SFRs might contain late-type Gould Belt members and/or stars belonging to large star halos around Orion and/or Taurus, possibly originated from an earlier episode of star formation. We call this component the dispersed young population (density of  ~0.5−5 stars/deg²). This population appears to be unrelated to any molecular cloud. The number counts in these areas are in significant excess compared to Galactic model expectations. A high fraction of these sources have strong lithium absorption features rather typical of stars with age  ≲ 20 Myr. They also show an iron abundance distribution consistent with that of nearby open clusters and associations younger than 150 Myr. Recent analysis including parallax information from Hipparcos and Tycho suggests that these stars are uniformly distributed along the line of sight, as in the disk-like structure proposed by Guillout et al. (1998) for the GB, and do not pile up at the supposed outer edge of the Gould Belt (Bally 2008). On the other hand, the interpretation of large star halos around Orion and/or Taurus as result of an earlier episode of star formation, e.g. about 20 Myr ago, would require an initial velocity dispersion of nearly 10 km s^-1 to explain the extension of these halos, whereas the value typically observed for young associations is less than about 2 km s^-1, although some authors claim in situ star formation in turbulent cloudlets (Feigelson 1995).

• 3.

  The uniformly distributed, widespread population has a density  ≲ 0.5 stars/deg² near the Galactic plane and is dominated by stars with a lithium abundance compatible with ZAMS or older ages. The number counts agree with standard Galactic models. However, a few of these stars appear to be of PMS nature, and explaining their origin remains challenging.

Analogous studies based on RASS data and spectroscopic effort are ongoing in the southern hemisphere and within the SACY¹⁰ project (Torres et al. 2006). Preliminary results seem to show two populations: i) an old one associated with evolved stars (similar to what we call the “widespread population”); ii) a young one corresponding to  ~10−100 Myr old associations (e.g., ϵ Cha, Argus, Carina, etc.).

6. Conclusions

We analyzed the young stellar content within the ROSAT X-ray All-Sky Survey on a scale of several thousand square degrees in the general direction of the Orion star-forming region. The study of this young stellar component, through spectroscopic follow-up of a subsample of stars and comparison with Galactic model predictions, leads us to the following conclusions: The X-ray selected young star candidates consist of a mixture of three distinct populations.

• 1.

  The youngest clustered population comprises the denseregions associated with sites of active or recent star formation,where the number counts far exceed the Galactic modelpredictions.

• 2.

  The young dispersed population of late-type stars, whose number counts are in significant excess with respect to Galactic model expectations. We cannot firmly establish whether this population represents the late-type component of the Gould Belt or originated from distinct episodes of star formation.

• 3.

  The widespread population uniformly distributed and dominated by field (ZAMS or older) stars, where the number counts agree with standard Galactic models.

In addition, two new young stellar aggregates, namely the “X-ray clump 0534+22” and the “X-ray clump 0430−08”, were singled out. These aggregates deserve future investigation for their complete characterization.

Finally, the future Gaia¹¹ mission will provide trigonometric parallaxes and radial velocities for targets in the general direction of Orion with a precision of a few parsecs at the Orion distance. Knowing the distances to individual stars will permit us to place them on the HR diagram with unprecedented precision, allowing us to firmly establish the origin of the widespread population of young stars on a Galactic scale.

Online material

Appendix A: Large-scale spatial distribution of the targets

Fig. A.1 Large-scale spatial distribution of 1483 candidate young stars (small dots) in the Orion and Taurus-Auriga associations. The X-ray sources have been selected from the ROSAT All-Sky Survey. Their space density is color-coded, where the parameter α represents the discriminant probability; the higher α, the more probable is the occurrence of a source corresponding to a young star. Circles refer to sources observed spectroscopically in this work, while triangles represent targets previously analyzed by Alcalá et al. (1996) and Alcalá et al. (2000). The symbol size increases with the lithium equivalent width EW (0,  <150, 150–300,  > 300 mÅ). The dashed line coincides with the position of the Galactic plane, while the dash-dotted line represents the Gould Belt, as given by Guillout et al. (1998). The sky strip (195° < lII < 205°,  − 60° < bII < 15°) selected by us across the Gould Belt in the Orion vicinity is also shown together with the area selected on the X-Clump 0534+22 (see Sect. 4). The regions of enhanced density of young star candidates and of high-lithium stars southward of the Galactic plane coincide with the location of the Gould Belt midplane. North is up and East to the left.

Fig. A.2 Large-scale spatial distribution of all targets observed at low and/or high resolution. Circles refer to sources observed spectroscopically in this work, while triangles are targets analyzed by Alcalá et al. (1996) and Alcalá et al. (2000). The symbol size increases with the lithium equivalent width (0,  < 150, 150–300,  > 300 mÅ). The arrows refer to stars for which we measured the iron abundance (downward for  [Fe/H]  ≤  −0.14, horizontal for  −0.14 <  [Fe/H]  ≤ 0.06, upward for  [Fe/H]  > 0.06; the central bin corresponds to the average metallicity of our targets  ± 1σ). Dotted lines represent the strip, long-dashed line is the Galactic plane, while dash-dotted lines represent the Gould Belt disk, as outlined by Guillout et al. (1998). The CO J = 1 → 0 emission contour maps by Dame et al. (2001) of the Orion, Monoceros, and Taurus molecular clouds, and the λ Orionis H ii region are also overlaid.

Acknowledgments

This research made use of the SIMBAD database, operated at the CDS (Strasbourg, France). We acknowledge the use of the ESO Science Archive Facility. K.B. acknowledges the financial support from the INAF Postdoctoral fellowship. We thank Fedor Getman for software counsel. We thank the anonymous referee for his/her careful reading and useful comments and suggestions. Last but not least, K.B. gives special thanks to AM-LP-MZ for their moral support during the preparation of the manuscript, and J.M.A. & E.C. are grateful to MD, RNF, and LMS for encouragement to pursue this work.

References

All Tables

All Figures

Fig. 1 Large-scale spatial distribution of our targets (circles) and those of Alcalá et al. (1996, 2000; triangles). Empty symbols refer to low-resolution spectra, while filled symbols represent objects observed at high resolution. The dotted lines mark the strip, the long-dashed line is the Galactic plane, while the dash-dotted lines represent the Gould Belt disk, as outlined by Guillout et al. (1998). The CO J = 1 → 0 emission contour maps by Dame et al. (2001) of the Orion, Monoceros, and Taurus molecular clouds, and the λ Orionis H ii region are also overlaid.

In the text

Fig. 2 2MASS (J − H) versus (H − K) diagram for the observed sources. Squares, asterisks, and stars represent objects in the three clumps, as indicated in the legend. The circles refer to the “non-clump” stars. The solid curve shows the relation between these indexes for main sequence stars (lower branch; Bessell & Brett 1988) and giants (upper branch; Kenyon & Hartmann 1995), where the 2MASS color trasformations were used. The AV = 1 mag reddening vector is shown with an arrow. The typical 2MASS photometric errors are overplotted on the lower-left corner of the panel.

In the text

	Fig. 3 Spectral type distribution of the young stellar candidates.
In the text

Fig. 4 Lithium (λ6707.8 Å) equivalent width versus log Teff for the stars in the strip (195° < lII < 205°), observed at low (left panel) and high (right panel) resolution. Three different bins of Galactic latitude were considered (see also Fig. 1). The upper envelope for the Pleiades sample, as adapted from Soderblom et al. (1993), is overplotted as a dashed line. Mean errors in $\mathit{E}{\mathit{W}}_{\mathrm{Li}}^{\mathrm{lr}}$ (~60 mÅ), $\mathit{E}{\mathit{W}}_{\mathrm{Li}}^{\mathrm{hr}}$ (~10 mÅ), and log Teff (~100 K, as obtained from the spectral synthesis) are overplotted on the lower-right corner of each panel. Open symbols represent the Alcalá et al. (1996) and Alcalá et al. (2000) results from low- (left panel) and high- (right panel) resolution spectra, respectively.

In the text

Fig. 5 Lithium abundance versus effective temperature for stars observed at high resolution (filled symbols refer to our sample, while open symbols represent the Alcalá et al. (2000) measurements). The “lithium isochrones” by D’Antona & Mazzitelli (1997) in the 2 × 106 − 1 × 108 yr range are overlaid with dashed lines. Dotted lines represent the limiting detectable log n(Li), as derived from the Pavlenko & Magazzù (1996) non-LTE curves-of-growth, assuming 10 mÅ as mean EWLi error and adding quadratically the mean contribution of the iron line at 6707.4 Å computed from the empirical relation obtained by Soderblom et al. (1993) between B − V color and EWFe (Kenyon & Hartmann 1995, tables were used to convert B − V colors into Teff). Left panel: stars located along the strip crossing the GB in the range 195° < lII < 205° (see Fig. 1) and three bins of Galactic latitude. Right panel: stars in known (L1616) and unknown (X-Clump 0430−08 and X-Clump 0534+22) groups.

In the text

	Fig. 6 Lithium abundance versus vsini for stars observed at high resolution. Symbols as in Fig. 5. Bottom panels show for each group the log n(Li) dispersion as a function of vsini at the following bins: 0–30, 30–60, 60–90, 90–120 km s-1. The values are shifted in vsini to better visualize the symbols.
In the text

	Fig. 7 Radial velocity distribution for the on-strip stars (left panel) and in the young aggregates (right panel), respectively.
In the text

	Fig. 8 [Fe/H] distribution for stars observed with the FOCES spectrograph and representing the young population (dashed histogram) and the field population (filled histogram). The bar represents the mean value of  [Fe/H]  =  −0.01 ± 0.03 obtained for open clusters younger than 150 Myr within 500 pc from the Sun (see Biazzo et al. 2011a, and references therein).
In the text

Fig. 9 Cumulative log N(>S) − log S distributions for RASS data in the Orion vicinity. A field centered on the Galactic plane is shown with a dotted line. The dashed line includes Orion and a section of the Gould Belt. The solid line corresponds to the prediction of our X-ray models at lII = 200° and zero Galactic latitude. The RASS completeness limit of 0.03 cnts/s is marked with the double arrow.

In the text

	Fig. 10 Detailed comparison of three fields in the strip perpendicular to the Galactic plane between lII = 195° and lII = 205°, as shown in Fig. 2. The three areas include: the Galactic plane, the Orion SFRs, and the Gould Belt, respectively. Histograms are cumulative, not stacked, i.e., each histogram includes the stars in all subcategory histograms.
In the text

	Fig. 11 HR diagram for the stars in the X-Clump 0534+22. The pre-main sequence evolutionary tracks of Palla & Stahler (1999), together with their isochrones and birthline are shown by solid, dotted, and dash-dotted lines, respectively. Labels indicating the track masses and the isochronal ages (from 1 to 100 Myr) are also overlaid.
In the text

Fig. A.1 Large-scale spatial distribution of 1483 candidate young stars (small dots) in the Orion and Taurus-Auriga associations. The X-ray sources have been selected from the ROSAT All-Sky Survey. Their space density is color-coded, where the parameter α represents the discriminant probability; the higher α, the more probable is the occurrence of a source corresponding to a young star. Circles refer to sources observed spectroscopically in this work, while triangles represent targets previously analyzed by Alcalá et al. (1996) and Alcalá et al. (2000). The symbol size increases with the lithium equivalent width EW (0,  <150, 150–300,  > 300 mÅ). The dashed line coincides with the position of the Galactic plane, while the dash-dotted line represents the Gould Belt, as given by Guillout et al. (1998). The sky strip (195° < lII < 205°,  − 60° < bII < 15°) selected by us across the Gould Belt in the Orion vicinity is also shown together with the area selected on the X-Clump 0534+22 (see Sect. 4). The regions of enhanced density of young star candidates and of high-lithium stars southward of the Galactic plane coincide with the location of the Gould Belt midplane. North is up and East to the left.

In the text

Fig. A.2 Large-scale spatial distribution of all targets observed at low and/or high resolution. Circles refer to sources observed spectroscopically in this work, while triangles are targets analyzed by Alcalá et al. (1996) and Alcalá et al. (2000). The symbol size increases with the lithium equivalent width (0,  < 150, 150–300,  > 300 mÅ). The arrows refer to stars for which we measured the iron abundance (downward for  [Fe/H]  ≤  −0.14, horizontal for  −0.14 <  [Fe/H]  ≤ 0.06, upward for  [Fe/H]  > 0.06; the central bin corresponds to the average metallicity of our targets  ± 1σ). Dotted lines represent the strip, long-dashed line is the Galactic plane, while dash-dotted lines represent the Gould Belt disk, as outlined by Guillout et al. (1998). The CO J = 1 → 0 emission contour maps by Dame et al. (2001) of the Orion, Monoceros, and Taurus molecular clouds, and the λ Orionis H ii region are also overlaid.

In the text