Free Access
Issue
A&A
Volume 660, April 2022
Article Number A109
Number of page(s) 20
Section Interstellar and circumstellar matter
DOI https://doi.org/10.1051/0004-6361/202141367
Published online 22 April 2022

© ESO 2022

1 Introduction

Massive stars have luminosities larger than 103 L, corresponding to a spectral type of B3 or earlier, and have stellar masses higher than 8 M. The formation of massive stars is far less understood than that of low-mass stars (<8 M; see reviews by Tan et al. 2014; Motte et al. 2018). The formation of massive stars differs from the formation of low-mass stars in several ways. In the former scenario, the Kelvin-Helmholtz times are much shorter due to the high luminosities of massive stars. Such stars tend to form in dense clusters and exhibit a higher multiplicity fraction (Motte et al. 2018). While accreting at high rates, massive stars that grow as large as 10 to 15 M develop extended photo-spheres and resemble red giants (Hosokawa & Omukai 2009). Recent studies have examined the formation of massive stars and its similarity to low-mass star formation by searching for ubiquitous phenomena found in low-mass star-forming regions (such as disks, jets, and outflows in the scenario of disk-mediated accretion; see Beuther et al. 2002; López-Sepulcre et al. 2010; Sánchez-Monge et al. 2013; Cesaroni et al. 2017; Purser et al. 2018; Sanna et al. 2018; Kavak et al. 2021). Massive stars, in contrast to low-mass stars, reach their main-sequence luminosity while still embedded in an accreting natal cloud of gas and dust (Hosokawa & Omukai 2009; Kuiper et al. 2011). A massive protostellar embryo heats and ionizes the gas of its surrounding envelope with extreme ultraviolet (EUV) photons (E > 13.6 eV), creating an H II region (Spitzer 1978). Young massive stars are surrounded by ultracompact H II regions of sizes <0.1 pc and densities >104 cm−3 (Churchwell 2002).

The gas in the ultracompact H II region is photoionized and heated by EUV photons, leading to an increase in gas pressure. This highly pressurized gas causes the H II region to expand until it reaches an equilibrium Strömgren sphere with a much lower gas density (Newman & Axford 1968). In the standard model of H II region evolution (Spitzer 1978), the thermal pressure of the plasma drives a D-type shock into the surrounding neutral medium; the shock sweeps up a dense, expanding shell, which traps the ionization front or photodissociation region (PDR; see reviews by Tielens & Hollenbach 1985; Hollenbach & Tielens 1997; Wolfire et al. 2003). H II regions are mainly classified on the basis of their size and internal density (Kurtz 2005), which span orders of magnitude in size (from 0.02 to 100 pc) and electron density (from 10 to 106 cm−3). In addition, H II regions are associated with interstellar bubbles due to their spherical morphology. The mid-infrared Galactic Legacy Infrared Mid-Plane Survey Extraordinaire (GLIMPSE), obtained with the National Aeronautics and Space Administration (NASA) Spitzer Space Telescope, revealed parsec-sized bubbles throughout the Galactic plane (Churchwell et al. 2006)1.

Stellar feedback results from the injection of energy, momentum, and mass into the interstellar medium (ISM) by massive stars. This feedback is a combination of ionizing radiation, radiation pressure, stellar winds, and supernovae on various spatial scales (from ~1 to ~100 pc) and dynamical timescales (from 104 to 106 yr). Without stellar feedback, the temperature of interstellar matter drops rapidly, and as a consequence of this cooling, new stars form rapidly by consuming the available gas content in the Galaxy (Kereš et al. 2009; Naab & Ostriker 2017; Lopez et al. 2014). By heating up the gas and injecting turbulence into star-forming regions, stellar feedback plays a key role in preventing this "cooling catastrophe" in the evolution of galaxies (Ceverino & Klypin 2009; Walch et al. 2012; Genzel et al. 2015).

Feedback processes are divided into momentum- and energy-driven mechanisms, which have different efficiencies in terms of energy input and time ranges (Fierlinger et al. 2016). For example, feedback from supernovae can provide enormous energy input that can shape the content of galaxies on large scales (10–100 pc), but much of that energy may be expended in rejuvenating hot gas in supernova remnants rather than coupling to molecular gas. On the other hand, pre-supernova feedback is also crucial for reproducing the properties of giant molecular clouds (Fujimoto et al. 2019; Olivier et al. 2021). In the last two decades, observational studies have also demonstrated that feedback mechanisms play an important role in the dynamics of star-forming regions (Lopez et al. 2011; Naab & Ostriker 2017).

Wind bubbles produced by stars of spectral types earlier than B2 are described by Castor et al. (1975) and were subsequently studied analytically by Weaver et al. (1977). However, the expansion of the bubbles – in other words, their main driving feedback mechanism and the underlying physical process – are poorly understood but are studied by simulations that are capable of incorporating several types of feedback mechanisms individually (Walch et al. 2012; Haid et al. 2018). In the past, it has been notoriously difficult to assess the relative contribution of expansion observationally, but that is now rapidly changing with the large-scale [C II] surveys enabled by the Stratospheric Observatory for Infrared Astronomy (SOFIA)2.

Most commonly, the neutral gas in the shells that confine these bubbles is translucent to far-ultraviolet (FUV) dissociating radiation, and thus they host little CO to be detected (e.g., Goicoechea et al. 2020) because CO is readily dissociated at low AV. In addition, most stars lie in the atomic or ionized phases of the ISM and not in molecular clouds. Thus, their feedback mostly impacts atomic or ionized gas not traced by molecules, such as CO and CO-dark H2 gas (Grenier et al. 2005). To date, a few alternative tracers have been reported to probe the CO-dark H2 gas (e.g., CF+ J = 1−0 by Guzmán et al. 2012 and HF J = 1−0 by Kavak et al. 2019). However, both species produce faint emission lines and require long integration times in the various regimes of the ISM. In addition to these molecular tracers, [C II] has been proposed as a more suitable tracer because its fine-structure transition (2P3/22P1/2 at 158 μm or 1.9 THz, i.e., ∆E/kB = 91.2 K) is the main cooling agent of predominantly neutral interstellar gas (Bennett et al. 1994; Hollenbach & Tielens 1997). As this [C II] line directly probes gas exposed to the FUV photons from massive stars, the [C II] 158 μm line is an ideal tracer of many types of stellar feedback mechanisms. The [C II] line is also one of the brightest lines in PDRs, and 30% of the total [C II] emission in the Galaxy comes from dense FUV-illuminated gas (Bennett et al. 1994; Pineda et al. 2014). Moreover, velocity-resolved observations of the [C II] line are an excellent probe of the kinematic and physical conditions of extended PDR gas (Goicoechea et al. 2015; Pabst et al. 2019), in our case, bubble shells. Unfortunately, its rest-frame emission is not accessible from ground-based observatories. With the upgraded German REceiver for Astronomy at Terahertz Frequencies (upGREAT) instrument on board SOFIA, it is possible to observe this transition from the stratosphere (Risacher et al. 2018). Therefore, [C II] observations of regions with a range of massive star formation activity with stars of different spectral types will provide invaluable input for simulations of Galaxy evolution (see SOFIA/FEEDBACK Survey3; Schneider et al. 2020).

Orion's Veil ("Veil" for short) is a series of foreground layers of gas and dust that lie in front of the Trapezium stars along the line of sight toward the Orion Nebula (O'Dell 2018; Abel et al. 2019). The Veil is a unique laboratory for studying the relative effects of feedback mechanisms as its proximity allows us to resolve the bubbles in the Orion Molecular Cloud (OMC) spatially and spectrally. Recent SOFIA [C II] 158 μm observations of the Veil focusing on the large-scale emission and dynamics have shown that stellar winds have swept up the surrounding material and created the Veil shell, a half-shell of neutral gas with a mass of ~1500 M that expands at ~13 km s−1 (Pabst et al. 2019, 2020). They have also found that stellar winds are more effective in disrupting OMC-1 than photoionization, evaporation, or even a future supernova explosion. The stellar wind is shocked, creating a hot and very dilute plasma that has been observed in X-rays with Chandra (Güdel et al. 2008). The high pressure of this hot plasma has driven a shock into the environment that has swept up a dense, expanding shell of gas. In this paper, we zoom in to a specific expanding structure to the northwest of the Veil using [C II] observations. This protrusion is clearly seen in Herschel Photoconductor Array Camera and Spectrometer (PACS; 70 and 160 μm) and Spectral and Photo-metric Imaging Receiver (SPIRE; 250, 350, and 500 μm) images and in Spitzer 8 μm emission images. In this study, we investigate the origin of the protrusion using velocity-resolved SOFIA [C II] maps and compare them to the dust, CO, and polycyclic aromatic hydrocarbon (PAH) emission. Finally, we use the energetics of the protrusion to assess the driving mechanism.

The paper is organized as follows. In Sect. 2, we describe the observations of [C II], 12CO, and 13CO as well as dust emission. In Sect. 3, we derive observational results on the general morphology, emission features, and stars (young stellar objects and early O, B, and A stars) in the Veil. Section 4 contains a detailed analysis of the morphology, the expanding shell and its velocity, and calculations of the kinetic energy of the protrusion. Finally, we discuss whether or not the Veil is breached at the location of the protrusion in Sect. 6.

thumbnail Fig. 1

Integrated intensity [C II] 158 μm map. Left: integrated (between −5 and +14 km s−1) intensity [C II] 158 μm map of the OMC observed by the upGREAT receiver on board SOFIA. The positions of NGC 1977, Trapezium stars, M42, M43, and the Orion Bar PDR are labeled. The green box shows the extracted region from the map, including the area of interest for this study, the protrusion. Right: close-up view of the protrusion. The bright ridge of emission is the edge of the expanding Veil shell. Faint emission extends well beyond this shell – the protrusion of which has a multi-component structure. More details about the kinematics and components of the protrusion are given in Sect. 3.

2 Observations

2.1 C II observations

The observations were conducted with SOFIA, which is an airborne observatory project of NASA and the German Aerospace Centre (DLR). SOFIA is a modified Boeing 747-SP airplane that carries a telescope with a diameter of 2.7 m in the rear fuselage (Young et al. 2012). By flying up to 45 000 ft, SOFIA makes it possible to observe at frequencies blocked by the atmosphere from the ground. A large part of the spectrum at far-infrared (FIR) frequencies (1-10 THz) becomes accessible. At the same time, a few molecular species (H20, O3) in the Earth's atmosphere still block FIR radiation at certain frequencies (Risacher et al. 2016).

The data were collected with the upGREAT instrument on board SOFIA (Risacher et al. 2018) for the Large Program of the C+ SQUAD led by A. G. G. M. Tielens. upGREAT is a heterodyne array receiver with 21 pixels. At the time of the observations, it comprised two 7-pixel low-frequency arrays (LFAs) plus one 7-pixel high-frequency array (HFA). The two 7-pixel sub-arrays with hexagonal layouts are designed for the LFA receiver with dual-band polarization. They cover the 1.83−2.07 THz frequency range, where the [C II] 158 μm and [O I] 145 μm lines can be found. The other hexagonal 7-pixel array is located in the HFA that covers the [O I] 63 μm line. The upGREAT instrument uses local oscillators to achieve very high spectral resolution (vv = 107). An area of about 1 square degree in Orion was surveyed in the [C II] 1.9 THz line (cf. our Fig. 1; Pabst et al. 2019). The native spectral resolution of the map is about 0.04 km s−1. The final data were resampled to 0.3 km s−1 to achieve a better signal-to-noise ratio. The final rms noise (in Tmb) is 1.14 K in 0.3 km s−1 velocity channels. The spatial resolution of the map is 16, which corresponds to 0.03 pc at the distance of Orion, 414 pc4 (Menten et al. 2007). The data cube is made at local standard of rest (LSR) velocities between −50 and +50 km s−1. The [C II] emission mostly appears between −10 and + 15 km s−1 in the entire cube. More detailed information about the observations has been given in Pabst et al. (2019).

We extracted the [C II] observations within the green box from the map presented in Fig. 1. The map is centered on an arbitrary point, а = 05h34m17.77s, δ = −05°20′ 03.89′′ (J2000), and covers the entire protrusion to the northeast of the Veil (Fig. 1).

thumbnail Fig. 2

Images of Orion's protrusion at different wavelengths and angular resolutions. The observed transition or frequency is given for each panel. The [C II],12CO (2−1), and 13CO (2−1) observations are integrated between −5 and +14 km s−1.

2.2 Molecular gas observations

We used new 12CO J = 2−1 (230.5 GHz) and 13CO J = 2−1 (220.4 GHz) line maps taken with the Institut de Radioas-tronomie Millimétrique (IRAM) 30 m telescope. These data are part of the Large Program "Dynamic and Radiative Feedback of Massive Stars" (PI: J. R. Goicoechea). This project uses the old CO HEterodyne Receiver Array (HERA) and the new Eight MIxer Receiver (EMIR) observations of the Orion Nebula. Goicoechea et al. (2020) describes how the old HERA and the new EMIR CO maps were merged. The most recent data relevant to this study were acquired in 2020. We extracted the region indicated in Fig. 2 from the original CO cubes. The line intensities are presented in main-beam temperature (Tmb) for both CO observations. In order to compare with the velocity-resolved [C II] map, we smoothed the 12CO (2−1) and 13CO (2−1) data to the angular resolution of the SOFIA [C II]maps, 16′′. The average rms noise level in these maps is 0.20 K in 0.41 km s−1 velocity channels. A more detailed description of the CO observations can be found in Goicoechea et al. (2020).

2.3 Ionized gas observations

We used the Hα images of the calibrated ESO/Digitized Sky Survey 2 (DSS-2) image obtained at the ESO/Max Planck Gesellschaft (MPG) 2.2-meter telescope at La Silla (Da Rio et al. 2009). The Orion Nebula was observed on two different nights with the same observing strategy, with 0.238′′ per pixel. After combining the dithered exposures, the final map was created after trimming to the overlapping area. In the final map, the surroundings of the Trapezium stars are saturated, but no saturation is seen in our region of interest. We extracted the same region as indicated in Fig. 1 to trace ionized gas with the Hα map within the protrusion. The trimmed Hα map we used is given in Fig. B.2.

2.4 Far-IR photometric observations

We used the archival Herschel images of the dust thermal emission for comparison to the [C II] data, and in particular used this to estimate the mass of dust (and gas) associated with the protrusion. The OMCs were observed as part of the Gould Belt Survey (André et al. 2010) in parallel mode using the PACS (Griffin et al. 2010) and SPIRE (Poglitsch et al. 2010) instruments on board Herschel. We used the photometric images of PACS at 70 μm (beam FWHM of 5.6′′), 100 μm (beam FWHM of 6.8′′), and 160 μm (beam FWHM of 10.7′′) and those of SPIRE at 250 μm (beam FWHM of 18.1′′) and 350 μm (beam FWHM of 25.2′′). Because of the limited angular resolution, we refrained from using the longest wavelength SPIRE band at 500 μm in the comparison of the dust emission with the SOFIA [C II] emission. Inspection of the 350 μm map reveals that omission of the 500 μm data does not compromise our analysis. We give more details about the model for fitting the Herschel fluxes and the results of the spectral energy distribution fitting in Sect. 4.

A comparison between the [C II] and Herschel maps shows that the shorter wavelengths have almost the same morphology, which clearly represents FUV-heated warm dust in the protrusion (see Fig. 2). However, faint emission, which could be physically connected to the protrusion itself, appears to the northwest of the protrusion. This component is also visible in most of the maps in Fig. 2. Unfortunately, our [C II] observations do not cover this component.

2.5 Mid-IR observations

We also made use of the Wide-field Infrared Survey Explorer (WISE) map of the Extended Orion Nebula5 (see also Fig. B.4). This map shows emission at 3.4 μm and 4.6 μm, both of which come mainly from hot stars, as well as 12 μm emission and 22 μm emission from relatively cooler objects, such as the dust in the nebulae. The field of view of the image is 3° × 3°, which covers the Veil and the extended emission coming from the dust. We trimmed the map to highlight a few striking jet-like structures that are present near the protrusion to the northeast of the Trapezium cluster.

To trace the FUV-illuminated surface of PDRs, we used the Spitzer 8 μm image (see Fig. 2). The full width at half maximum (FWHM) of the point spread function is 1.9′′ at 8.0 μm. As in all observations, we extracted the same region from the 8 μm image for further analysis.

3 Results

Figure 2 shows the integrated intensity maps of the protrusion. The protrusion is clearly seen in Herschel PACS 70 and 160 μm and SPIRE 500 μm images. We show three representative dust emission maps in Fig. 2 that trace the emission of dust heated by the Trapezium stars to ~40 K. We also used the 12CO and13CO J = 2−1 observations to identify CO molecular gas exposed to intense FUV radiation. To confirm the location of PDRs, we overlay the Spitzer 8 μm emission produced by PAHs on the [C II] map in the right panel in Fig. B.1. We see that the [C II] emission has a similar distribution as the 8 μm emission map at the bottom and along the arm-like structure of the protrusion. We also compared the Hα emission with [C II] to trace the ionized gas emission within the protrusion. The outlines of the protrusion are also quite apparent in Hα.

Perusal of the individual channel maps (see Fig. B.7) reveals that the protrusion is particularly noticeable in the LSR velocity range of −3 to +8 km s−1 in the [C II] observations. It is clearly offset from the main [C II] emission associated with the OMC–1 core at ≳9 km s−1. Unlike the [C II] map, the protrusion does not appear in the 12CO J=2–1 velocity channel maps (e.g., see the –0.8 km s–1 channel map in Fig. B.7) associated with the boundary of the Veil. On the other hand, 12CO J = 2–1 shows a protrusion-like structure at higher velocities (12–13 km s–1) than those of the OMC–1 core (see Fig. B.7). More details about the origin of this high-velocity CO emission toward our protrusion are given in Sect. 4.3.

4 Analysis

4.1 Morphology of the protrusion

Our observations (see Fig. 2) reveal expanding bow-shaped cavities in the northwestern part of the Veil. The inside wall of these cavities is ionized, as shown by the Hα emission, and the [C II], 8 μm, and 70 μm emission trace the surrounding PDR. First, we explored the protrusion itself, and later the ionizing star(s) and the origin of the protrusion. We fitted the elliptical structure of the limb-brightened shell in the channel map at 12 km s–1 with a least-square fit to estimate the size and expansion velocity (see Sect. 4.2) of the protrusion. We find that the size of the protrusion is 1.3 ± 0.1 pc from the Veil boundary in the northwestern direction. The minor and major axes of the model are 0.5 ± 0.1 and 1.3 ± 0.1 pc, respectively. The thickness of the shell derived is 0.1 ± 0.05 pc. We assume an elliptical geometry to calculate the energetics of the protrusion in Sect. 4.4 because the channel maps suggest an elliptical morphology.

We examined the channel maps and determined the size of the expanding structure to be 1.3 pc in the southeast-northwest direction and 0.5 pc in the northeast-southwest direction. This ellipsoidal morphology is already quite apparent from the 8 μm and 70 μm dust emission maps. While morphologically the structure resembles a half-cap in the plane of the sky, perusal of the position-velocity (pv) diagrams shows that in all crosscuts the structure starts and ends at the cloud velocity (+9 km s–1), even in the southeast-northwest direction (cf. crosscut 23 in Fig. 3). The observed pv diagrams are reasonably well fitted by a coherent half-ellipsoidal shell.

4.2 Expansion velocity and timescale

Guided by the velocity channel maps, we quantified the characteristics of the protrusion in [C II] pv diagrams. We created pv diagrams along 30 diagonal crosscuts, which are shown in Fig. 3. We illustrate the results with two pv diagrams (crosscuts 8 and 23 in Fig. 3). The other pv diagrams are presented in Figs. B.9 and B.10 and support the analysis presented here. Both pv diagrams in Fig. 3 reveal two arc-like structures, which are the telltale signs of two half-bubbles, both expanding only toward us.

Inspection of all pv diagrams reveals two expanding shells. We fitted these two arc-like structures in the pv diagram with a least-square fit over the chosen positions. The expansion velocity (Vexp) of the first shell (yellow dashed line in Fig. 3) is Vexp = 6 ± 0.2 km s–1, and the second (white dashed line in Fig. 3) is Vexp = 12 ± 0.2 km s–1, which indicates the maximum expansion velocity of the outer shell. We fitted the pv diagrams (numbers 8 and 23) that represent the maximum expansion of the protrusion using a simple bubble model (see Fig. 3). The emission at VLSR = +9 km s–1 (i.e., the dashed green line) seen horizontally in both diagrams arises from the Orion cloud itself.

When we take a closer look at the [C II] channel maps in Fig. B.7, we find two spatial components between –5 and + 14 km s–1. The first component appears from –3 to +5 km s–1. The second component is identified between +6 and +14 km s–1 (see Fig. 3). We did not see the expanding shells in the CO channel maps or pv diagrams (see also Fig. B.8) and only detected an emission feature at VLSR +2.7 km s–1 (Globule #10 of Goicoechea et al. 2020).

The classical way to calculate the expansion timescale (texp) for structures moving perpendicular to the line of sight is to use the ratio between the size of the outer shell and the maximum expansion velocity (size/vexp) (see also Beuther et al. 2002; Maud et al. 2015). box, we estimated the expansion velocity as 12 km s–1 using pv-diagram fit results. Using this expansion velocity and size (1.3 pc), we derive ~1.1 × 105 yr, which is ~50% of the expansion timescale (texp) of the entire Veil shell (Pabst et al. 2019, 2020).

4.3 Components of the protrusion

In Fig. 4, we show the comparison of [C II] 158 μm, 12CO J = 2−1, and 13CO J = 2−1 spectra at 12 positions, covering the protrusion, to find substructures of the protrusion. 12CO and 13CO always have a similar profile but at different brightnesses. The CO lines typically show two emission components (at +7 and +13 km s−1) at several positions (3, 4, 5, and 12) that correspond to the bottom of the protrusion. The velocity separation between the two CO peaks varies between 1 and 3 km s−1. The peaks in both CO isotopologues show small shifts (2−3 km s−1) to higher or lower velocities. The absence of these velocity peaks in the [C II] line profiles indicates that the CO emission is associated with structures deeper in OMC-1 that are not exposed to FUV radiation.

In contrast, the [C II] line shows a different behavior than CO, with the exception of position 2. In addition to the OMC, which emits predominantly at VLSR = +9 km s−1 (i.e., the dotted red line in Fig. 4), we identified two other components on the [C II] emission. To investigate the origin of these components, we integrated the [C II] emission between −5 and +3 km s−1 (the first component is shown in blue), +3 and +12 km s−1 (the second component, the OMC itself, is shown in green), and +12 and +15 km s−1 (the third component is shown in red). It should be noted that the first component shifts to somewhat higher and lower velocities and that part of the profile of the first emission structure may be confused by emission of the OMC–1 core surface that dominates the total emission. We are therefore not able to use a fixed integration range for this component. The integration range was assumed based on positions 1 and 12 in Fig. 4. Using integrated intensity maps, we created a three-color, red-green-blue (RGB), map of our protrusion using [C II] and 12CO cubes, which is shown in Fig. 5. In the [C II] RGB map, relative to the background OMC–1 core, the protrusion and the other structure are moving toward us at 9 km s−1. Together with the OMC, the blue component moving toward us is associated with the smaller (in size) expanding shell that we identified in the pv diagrams in Fig. B.9. The presence of a red component at velocities (13 km s−1) higher than that of the OMC–1 core suggests that there is a backward extension of the Veil shell. It is possible that the Veil shell on the rear side is tilted with respect to the background of the OMC–1 core, and sticking out of it, allowing for extension away from us. We do note, though, that the extension of the Orion Bar in the M42 H II region is also quite prominent in this red channel, and, in that case, this velocity behavior could be related to complex morphology/velocity structures within the H II region or at the PDR/H II edges. In the 12CO RGB map, we track several components with velocities that differ from those of [C II]. We conclude that the limb-brightened shell of the protrusion observed in [C II] does not contain CO and that the CO emission is associated with the molecular cloud in the background.

In addition, the cavity (red-shifted emission in the vicinity of positions 8 to 11) was identified with an expanding shell identified in Combined Array for Millimeter-Wave Astronomy (CARMA) CO J = 1−0 observations by Feddersen et al. (2018). They argued that Bruno 193 − an F9IV star at the geometric center – is driving this CO bubble. This bubble is thought to be embedded in the OMC–1 cloud behind the Veil. Based upon the kinematic information, we consider that this CO bubble is not related to the protrusion and that this star is insufficient to ionize the gas, notably given that this star is displaced by 7′ (0.85 pc) from the center of our protrusion. The general morphology, the subcomponents, and the expanding shells are discussed in more detail in Sects. 4.1, 4.2, and 4.3.

thumbnail Fig. 3

Two pv diagrams of the protrusion. Top: selected crosscuts along the green arrows overlaid on the integrated [C II] intensity map. The number of the crosscuts is indicated at the starting point of the cut. Middle and bottom: pv diagrams generated along the magenta crosscuts (cuts 8 and 23, respectively). The pv diagrams with horizontal green lines both show the [C II] emission produced by the FUV-illuminated surface of OMC, and the arcuate white and yellow lines trace the shell expanding at 12 km s−1 and 6 km s−1, respectively. The pv diagrams in Figs. B.9 and B.10 have the same scale on both axes. A 12CO-pv diagram along crosscut 23 is shown in Fig. B.8 for comparison with [C II].

thumbnail Fig. 4

Spectral lines toward the protrusion. Upper panel: Spitzer 8 μm image of the protrusion. Red circles indicate the 12 positions that we used to extract line profiles and have an aperture of 16′′. Lower panel: velocity-resolved spectra of [C II] (colored in gray), 12CO J = 2−1 (blue), and 13CO J = 2−1 (cyan) in the direction of protrusion for the 12 selected positions in the upper panel. The vertical dotted red line at 9 km s−1 marks the approximate velocity of the emission generated by the OMC and the associated star-forming molecular cloud behind the Veil.

thumbnail Fig. 5

Three-color image of the protrusion created using the three components estimated in the [C II] line profile in Fig. 4. Blue emission is the integrated emission between −5 and +3 km s−1, green between +3 and +12 km s−1, and red between +12 and +15 km s−1 of the SOFIA [C II] 158 μm (upper panel) and IRAM 12CO (lower panel) emission maps. White circles show the 12 selected positions from Fig. 4. A movie is available online.

4.4 Kinetic energy and momentum

To identify the driving mechanism of the protrusion, we calculated its momentum and kinetic energy. For this, we followed the same methods as in Pabst et al. (2020). This also allowed us to directly compare our results with the Veil shell (Pabst et al. 2019). To calculate the mass in the limb-brightened shell of the protrusion, we used Herschel PACS (70 μm, 100 μm, and 160 μm) and SPIRE (250 μm and 350 μm) maps. All maps were convolved to the SPIRE 350 μm beam size of 20′′ FWHM, as this resolution is comparable to the spatial resolution of SOFIA [C II]. We converted the units of the SPIRE maps from Jy beam−1 to Jy px−1 using the beam areas given in the HIPE6 manual. The flux densities at each pixel are modeled as a modified blackbody,

Here, Td denotes the effective dust temperature, τ0 the dust optical depth at the reference wavelength (λ0), and β the dust grain opacity index. The reference wavelength (λ0) is 160 μm. The Td and τ160 are free parameters. The dust emissivity index (β) is fixed at 2 in all models (Goicoechea et al. 2015; Kavak et al. 2019; Pabst et al. 2019). Maps of the fitted optical depth and dust temperature are shown in Fig. B.6. The statistical values of the dust temperature, the maximum, minimum, and median, are 50 K, 20 K, and 26 K, respectively. The same statistics for the optical depth at 160 μm are 2 × 10−1, 8 × 10−4, and 2 × 10−3, respectively. Using an average value of the dust optical depth over the protrusion, we calculated the hydrogen column density as (1)

where к160 is the 160 μm dust opacity per H atom7, which is 2.3 × 10−25 cm2/H atom for Rv = 5.5 (Weingartner & Draine 2001). Using these values and the median optical depth, which is 2 × 10−3, we calculated the column density, NH ~ 1.20 × 1022 cm−2 (or a visual extinction of Av = 8 mag). However, we also note that the limb-brightened shell of the protrusion seen in the [C II] map does not appear in the 12CO J=2−1 map, indicating a low column density (Av < 3 mag) – in other words, a thin expanding shell. The high column density we derived reflects a difference in geometry. The dust emission estimate refers to the column density along the line of sight of a limb-brightened shell. Assuming a spherical homogeneous shell with a relative thickness of 0.1 pc, the column density estimates perpendicular to the surface of the protrusion will be about 0.2 times the observed column density, and this confirms the upper limit expected from the absence of CO in the protrusion.

Assuming elliptical geometry, the mass of the limb-brightened shell is given by the surface area, S, times the surface density along the line of sight: M = S NH μ mH. With the dimensions of the ellipse and a thickness of 0.1 pc, the surface area is calculated to be 0.13 pc2, corresponding to a mass in the limb-brightened shell of 18 M. A geometric correction factor (see Appendix A) of 2.5 converts this then into the mass of the [C II] -emitting shell, ~45 M, which is ~3% of the mass estimate of the Veil shell (1500 M; Pabst et al. 2020). Using the mass estimate and the expansion velocity (12 km s−1), we calculated the kinetic energy (Ekin) of the [C II] gas tracing the neutral shell to be ~7 × 1046 erg. Our energy estimate is ~3% of the kinetic energy of the entire expanding Veil shell (Pabst et al. 2020). Also, the momentum of the protrusion would be −540 M km s−1.

Additionally, we estimated the density of the protrusion by using the relation between G0 and 70 μm reported by Goicoechea et al. (2020) for Orion to estimate the mass inside the protrusion. The G0 is given by (2)

where I70 is the 70 μm dust surface brightness in MJy sr−1. The median value of G0 is ~600 toward the protrusion, although a gradient can be seen in the G0 map (see Fig. 6). Using the estimate of G0, we calculated the density of a face-on PDR using Eq. (9.4) of Tielens (2010), which is given in Eq. (3). We isolated the density and expressed it in terms of G0:

(3)

The resulting density map is also shown in Fig. 6. The density decreases in the northwestern direction from the boundary of the Veil to the outer shell of the protrusion. We can check our gas density from the observed [C II] intensity using PDR models. For this purpose, we used the intensity of the [C II] 158 μm line emitted from the surface of an edge-on PDR as a function of the density and radiation field based on the PDR models8 of Kaufman et al. (1999), adopting an average G0 of 600 Habings. This results in an average density of 103 cm−3, in agreement with the estimates in Fig. 6. The density along the limb-brightened shell of the protrusion is comparable with that of the Veil shell (Pabst et al. 2020) and two or three orders of magnitude lower than that of the Orion Bar (Kavak et al. 2019; Pabst et al. 2020). We calculated the mass of the shell to be ~30 M. This is in good agreement with the values calculated box in this section. Lastly, the momentum budget would be between 360 and 540 M0 km s−1 (see Table 1 for a complete list of the characteristics of the protrusion).

thumbnail Fig. 6

Map of the incident radiation field, G0 (left) and the density (right) of the protrusion for a face-on PDR model adopted from Tielens (2010) (see Sect. 4.4 for a more detailed discussion).

Table 1

Comparison of the masses and energetics of the protrusion with the Veil shell.

4.5 Ionizing source

In Sect. 3, we show that the Hα emission follows a similar morphology as the [C II] emission. To understand the origin of ionized gas along the limb-brightened shell, we used the Hα flux to estimate the source of the ionizing photons. We can make an estimate for the extinction associated with the protrusion from the thickness of the shell and the estimated column density of the limb-brightened shell. Adopting a spherical half-shell with a relative thickness of 0.1 pc, we estimated that the column density along the line of sight is 0.2 times the column density derived from the dust emission of the limb-brightened shell, 2 × 1021 H nuclei per cm2. Using the extinction curve of Weingartner & Draine (2001), this corresponds to an extinction at Hα of 1.1 mag. Correcting the observed surface brightness for extinction results in an intrinsic Hα surface brightness of 525 MJy sr−1, or 2.7 × 10−7 erg s−1 cm−2 arcsec−2. We converted the surface brightness into emission measure9 (EM) using Eq. (4). Given a constant temperature of 8500 K obtained from radio recombination line observations (Wilson et al. 1997), (4)

with Hα in units of erg s−1 cm−2 arcsec−2. We derived the EM to be 1.40 × 107 pc cm−6. We then calculated the total number of ionizing photons (ℕLyc) emitted by the star (see Sect. 7.4.1 of Tielens 2010): (5)

where A is surface area in pc2 and the EM is in pc cm−6. We find 1.8 × 1050 photons s−1. We measured the number of ionizing photons over a hole in the wall of the Veil of 1 pc2, which is 1/16 of the total inner surface area of the Veil. In this case, the final number of ionizing photons is 1.1 × 1049 photons s−1. This indicates that the source of the ionizing photons should be an O-type star. The only O star in the Trapezium cluster is θ1 Ori C, the main ionizing star in the Orion Nebula (O'Dell et al. 2017). Therefore, we conclude that the source of the ionized gas in the protrusion should be θ1 Ori C.

Additionally, to find other possible driving stars/sources in the protrusion, we display young stars and protostars identified with Infrared Array Camera on the Spitzer (green circles in Fig. B.2; Megeath et al. 2005, 2012). In addition, we searched the Set of Identifications, Measurements and Bibliography for Astronomical Data (SIMBAD) archive for O, B, and A stars within a 0.5′ circle around the Veil and listed 53 stars in Table B.l (see also Fig. 7). This table consists of the identification and object name of the stars, coordinates in right ascension and declination (in degree units), spectral type, and object type10. The closest star to the protrusion is an A3 star (star 39 in Table B.1)11, which has a luminosity of 14 L and a mass of about 2.0 M. We consider that this star is insufficient to ionize the surrounding gas and cause a protrusion because these types of stars have a low effective temperature (Teff) and ionizing luminosity (Qi). Thus, we find no nearby powerful star that could ionize the gas or locally affect the shell or Veil in the northwest (see Sect. 4.4 for a detailed analysis). Hence, the ionizing photons from the Trapezium cluster must be able to reach this surface almost unimpeded.

thumbnail Fig. 7

SOFIA [C II] map of Orion with O, B, and A stars found in SIM-BAD. The list of stars retrieved from the archive is given in Table B.l. The blue, orange, and red circles are O, B, and A stars, respectively. The light green arrows indicate the positions of the Trapezium stars. Two A stars (Star 21 and 39 in Table B.l), which are the nearest to the protrusion, are also labeled.

5 Discussion

In this section, we examine three different scenarios in order to determine the driving mechanism of the protrusion. To begin, we discuss the role of the winds of the Trapezium stars, which are primarily responsible for the expansion of the whole Veil shell (Pabst et al. 2019). Later on, we examine the initial "dumpiness" of the preexisting low-density regions (Güdel et al. 2008). Finally, we discuss the fossil outflow of Trapezium stars, which may provide directional feedback on the Veil shell, as seen in the massive star-forming region NGC 1333 (Quillen et al. 2005).

5.1 Stellar winds

If the protrusion is driven by stellar winds of the Trapezium stars, in particular θ1 Ori C, as found for the Veil (Pabst et al. 2019), the protrusion itself should expand like the Veil shell. However, despite the fact that the velocity is (slightly) lower than that of the Veil, the protrusion goes far beyond the Veil wall. Alternatively, the stellar winds could originate from another massive star within or near the protrusion. To check this, we superimposed the positions of the known O, B, and A stars on the [C II] map (Fig. 7). There is no massive star within the protrusion. Only two A stars are found near the protrusion. However, the nearest A star (Star 39) does not follow the elongated morphology of the protrusion. The second A star (Star 21 in Fig. 7) is located at a comparable distance to the Trapezium stars. These findings lead us to believe that there is a preexisting structure that is now being overtaken by the expanding Veil shell.

Another way to estimate the role of the winds is to compare them with X-ray observations, in which the hot X-ray-emitting gas is traced inside the Veil. Using X-ray observations of the Veil, Güdel et al. (2008) showed that the X-ray emission from the ionized region indicates a hot plasma heated to a few times 106 K by the shocks created by the stellar winds. In other words, the presence of X-ray-emitting hot gas can be taken as an indication of stellar winds. However, there is no X-ray observation covering the protrusion. It should also be noted that X-ray emission is very susceptible to extinction by foreground material (Güdel et al. 2008). Therefore, X-ray observations may not be the best tool for investigating the effect of stellar winds, at least in our case. Imaging of optical line emission with the Apache Point Observatory will help us to detect the hot plasma (T > 30 000 K) inside the cavity (Bally et al., in prep.).

The star θ1 Ori C drives the most powerful wind, 2 × 10−7 M yr−1, with Vwind around 1000 km s−1. The winds of the lower-mass stars are weaker and have lower mass-loss rates. As a rough estimate, the total wind mass-loss rate from main-sequence OB stars is likely to be about 10−6 M yr−1, or about 1 M in 1 Myr. The momentum of the wind will be ~ 1000 M km s−1 in 1 Myr, comparable to that of the massive star outflows during their formation.

The protrusion has a limited lifetime due to the photo-ablation of its walls. Once the massive stars reach the zero-age main sequence and begin to ionize their surroundings, photo-ablation of the inner walls of these cavities will start to fill their interiors with plasma. To first order, the plasma will expand at the speed of sound in ionized gas at V[CII] = 10 km s−1. Using our mass estimations in Table 1, the surface area of the protrusion (0.385 pc2), and the incident flux of Lyman continuum photons (2 × 1049 s−1 for θ1 Ori C), we can estimate the mass-loss rate of the protrusion walls and how long the walls would survive (tsur). The mass-loss rate is given by (6)

where ƒ is a factor of order unity depending on geometry, which is taken to be to recover the Strömgren condition for a spherical HII region. The plasma density (ne) can be calculated assuming that the incident Lyman continuum flux (L(LyC)/(4 π D2)) equals the recombination along a path length (R) (7)

where αB is the Case B recombination coefficient of H; 2.6 × 10−13 cm3 s−1. The number of electron-proton recombinations per unit volume and unit time is equal to ne np αB. Using Eq. (7), we derive a plasma density (ne) of ~ 2 × 103 cm−3. The mass-loss rate from the protrusion walls is 1.8 × 10−4 M yr−1. Therefore, the lifetime of the protrusion (i.e., tsur = M/(dM/dt), where M is the mass of [C II]-emitting protrusion walls) is ~1.6 × 105 yr, which is consistent with the age of the Trapezium stars and the expansion timescale derived in Sect. 4.2, but not with the ages of the 09 to early B stars box the bright Orion Bar (the θ2 Ori A stars), which are more than 106 yr. We argue that the location of the protrusion is an ideal place to break Orion’s Veil and ventilate its hot plasma before a possible supernova occurs (~5 × 106 yr; see also Williams & McKee 1997).

thumbnail Fig. 8

Schematic picture of the protrusion and the fossil bipolar outflow with apparent structures overlaid on the WISE image (see also Fig. B.4). the suspected blue- and red-shifted outflow lobes are shown in blue and red, respectively. The Kelvin-Helmholtz instabilities reported by Berné et al. (2010) are indicated in yellow.

5.2 Blowout of the Veil shell

The wind of θ1 Ori C would produce a spherical bubble only if there were no obstacles blocking the propagation of the wind and the post-shock hot plasma. However, we know that there is a dense cloud in the region of the Kelvin-Helmholtz instabilities (see Figs. 2 and 8) containing CO (Berné et al. 2010) whose surface is affected by radiative feedback (maybe a wind) from the Trapezium cluster. The H-alpha ionization front of the nebula wraps around this structure. An even larger protrusion is visible south of this obstacle, especially in wide-field [S II] image (Fig. B.5). The core of this obstacle is seen in CO; it is the feature marked as the “KH Ripples” in the CARMA Orion CO survey figures (Kong et al. 2018). This also prevents hot plasma from flowing westward and forces it to head south of the Veil (Güdel et al. 2008; Bally 2010). A possible model for the northwestern protrusion is that the plasma driving the Veil shell has found a path of least resistance toward the northwest. Moreover, the south of the obstacle of the protrusion is seen at Vlsr +5 to +8 km s−1 (see the 12CO-RGB map in Fig. 5). Hence, blowout of the nebula around this obstacle is certainly possible The channel movie shows how the western rim of the Veil shell changes morphologically around this obstacle.. However, both the expansion of the protrusion, which is slower than that of the Veil shell, and the bipolar jet-like structures seen toward the Veil shell depicted in Fig. 8 are difficult to reconcile with this scenario.

5.3 Fossil bipolar outflow

The expansion velocity of the protrusion being slightly lower than that of the Veil and its extension beyond the boundary of the Veil suggest that the protrusion is a preexisting structure in the OMC–1 core that is now being overtaken by the Veil bubble. Following Bally et al. (in prep.), we suggest that this preexisting structure is the result of fossil outflow activity in the OMC–1 core created during the accretion phase of the massive protostars in the Trapezium cluster. In Fig. 8 we show probable blue- and red-shifted outflows on the WISE image. Blue-shifted outflows ejected from one or more of the Trapezium stars (Outflow-1, Outflow-2, and Outflow-3) create our protrusion in the northwestern part of the Veil shell. If there is truly a bipolar fossil outflow, we can expect to observe a red-shifted lobe toward the eastern rim of the Veil shell. The suspected red-shifted lobe is also associated with high-velocity 12CO emission (see Fig. B.4). Additionally, the red-shifted lobe appears to have broken the eastern rim, ejecting a tail-like structure at the head of the outflow (see Fig. B.4).

Once the protostellar jet switches off, the cavity blown by this jet will enter the momentum-conserving phase and expand while slowing down. As θ1 Ori C entered its main-sequence phase, its stellar wind started to blow the Veil bubble. The large amount of momentum involved in this kinematic structure could indicate outflow activity associated with the formation of the most massive star. To identify the possible protostellar source(s), we used the bolometric source luminosity and momentum of the outflows from Maud et al. (2015). The momentum of Moutflow red- and blue-shifted lobes are given individually in Fig. 9. The interpretation of the relation in Fig. 9 is that the jet or wind from the most luminous protostar drives the strongest and most powerful outflows. For the scatter in the momentum values, Maud et al. (2015) argue that it is caused either by outflow inclination angles or by multiple outflows driven by sources within dense cores. We also emphasize that this type of outflow activity is generally found in systems with ages of less than a few times 104 yr (Arce et al. 2007) and hence is a clear signature of protostellar activity.

Using the relation in Fig. 9, we estimate that a massive dense shell with a momentum of ~540 M km s−1 would require a luminosity of 3 × 104 to 3 × 105 L. This corresponds to B0- to 07-type stars (cf. Vacca et al. 1996, for stellar parameters of O and B stars), and several stars in the Trapezium region could be responsible. The most massive star, θ1 Ori C, is likely the culprit. We do notice that there are several other jet-like morphological structures present in the 8 μm and WISE images in the area of the protrusion (Figs. 8 and B.4). Our [C II] observations do not cover these structures, and therefore we have no kinematic information on their expansion. Further (deeper) studies are warranted to determine their kinematics. Here, we recognize that these structures may indicate the presence of multiple protostellar outflows, for example associated with the several Trapezium star clusters. Alternatively, these jet-like structures may reflect intermittent activity of a single, precessing object in a binary of the Trapezium cluster. We note that the trajectories of these jet-like structures do trace back to the Trapezium stars (Fig. 8).

At this point, we speculated that the protrusion was likely created by outflow activity when accretion in a protostar-disk structure was accompanied by a jet or wind in the polar directions. On this basis, it can be argued that the Trapezium stars (specifically θ1 Ori C) should have formed via disk-mediated accretion. This model of massive star formation is supported by recent studies that have found disks (Cesaroni et al. 2017), outflows (López-Sepulcre et al. 2010; Sánchez-Monge et al. 2013), and jets (Sanna et al. 2018; Kavak et al. 2019). If the protrusion is made of fossil outflow cavities, there must be counter-flows corresponding to the red-shifted lobe of the northwestern protrusions from the Trapezium cluster. WISE and 8 μm images show a vague protrusion in the opposite direction of the northwestern protrusion. Given the blue-shifts of the northwestern protrusions, this component should be the red-shifted lobe (the red arrow shows the red-shifted lobe in Fig. 8). However, the [C II] emission is weak at this position, preventing us from studying this red-shifted lobe in this work. We also note that the fossil outflow activity is not related to the explosive outflow or to the H2 fingers seen in near-infrared lines (Bally et al. 2017), as these fingers are still far (~1.5 pc) from the boundary of the Veil shell and are a relatively recent (~500 yr) “explosive” event.

thumbnail Fig. 9

Momentum of outflows from massive young stellar objects as a function of the source luminosity of the cores (Maud et al. 2015). The blue and red symbols indicate the blue- and red-shifted outflow lobe values, respectively, which are joined by a dotted line for each source. The horizontal blue-shaded range indicates the momentum of the protrusion, which is between 360 and 540 M km s−1. The cross at the bottom right shows the uncertainty for both axes.

6 Conclusion

In this study we have investigated the origin of the protrusion in the northwestern part of the Orion Veil shell using velocity-resolved [C II] 158 μm observations. The protrusion, which appears as a half-elliptical cap on the Veil, expands at a velocity of 12 km s−1 with a radius of 1.3 pc. In Sect. 5 we examine three possible mechanisms that could drive the protrusion in the northwest. We propose that the protrusion is formed by extinct or previously active Trapezium star outflows. During the early stages of massive star formation, the outflows of the Trapezium stars weaken the northwestern portion of the Veil. Later on, as the massive stars reach the zero-age main sequence, their radiation ionizes the surrounding gas while their stellar wind starts to blow a bubble filled with hot gas. The EUV and FUV photons can travel virtually unimpeded in the fossil outflow cavity, illuminating the protrusion and lighting it up in Hα emission, PAH emission, and [C II] emission. This suggests that mechanical feedback is the mechanism responsible for the formation of the protrusion rather than radiative feedback. Moreover, in Sect. 3 we also see that the lack of CO detections in the protrusion indicates a low NH, or, in other words, a thin shell in the northwestern Veil. In Sect. 5.1 we also show that the fossil outflow activity could cause breaks in the ionization front of Orion’s Veil because of photo-ablation from the protrusion walls, making the protrusion a suitable place for the Veil to break up.

Moreover, the diagonal pv diagrams parallel to the direction of expansion, in particular cuts 18, 19, and 20 in Fig. 3, show [C II] emission that extends somewhat beyond the protrusion. The densities of the limb-brightened shell are lower (by up to a factor of two) at the head of the protrusion. Outflows, particularly Outflow-3 in Fig. 8, appear to be associated with the chimney-like top of the protrusion, suggesting that the Veil shell has already been pierced there. This location could be a suitable place for the bubble to break. Furthermore, beyond the area mapped in [C II], the outflows and their extended morphology are also seen in the Spitzer 8 μm image and the dust emission maps of Herschel PACS 70 μm and WISE observations. Thus, future, more sensitive [C II] observations could clarify whether or not the Veil is already broken at the location of the protrusion.

If the protrusion is formed by a bipolar fossil outflow, the southeastern cavity carved by the red-shifted lobe of the fossil outflow (see Fig. 8) should exhibit characteristics similar to those of our protrusion. However, unlike our protrusion, the southeastern cavity is not filled with ionized hydrogen plasma, but is associated with high-velocity CO emission. We require further data to determine what protects it from Lyman continuum illumination from the Trapezium or θ2 Ori A. For example, a shock tracer such as the H2 v = 2–1 S(1)/v = 1–0 S(1) line intensity ratio (at 2.24 μm and 2.12 μm, respectively) and/or [Fe ii] 1.644 μm (e.g., observed with Keck or JWST in the near future) tracing the presence of shocked gas toward the southeastern protrusion may reveal the shocked gas within the red-shifted lobe of fossil outflow. Additionally, the EM of the southeastern protrusion could be too low. If this were true, we would witness free-free plasma at the cavity walls. This, however, is not the case. Thus, further effort should be directed toward determining the nature of the protrusion carved by the red-shifted lobe and establishing the involvement of bipolar fossil outflows in the formation of the northwestern protrusion.

In summary, three mechanisms can contribute to the formation, morphology, and expansion of H II regions: (i) pre-cavitation caused by powerful bipolar outflows; (ii) ionization and thermal expansion of the H II region in the early phases; and (iii) stellar winds in the late phases. The protrusions can be readily explained by pre-cavitation caused by collimated bipolar outflow. However, we cannot rule out the possibility of obstacle molecular clouds playing a role. Further research could focus on the outflow kinematics beyond the Veil’s northwestern protrusion and the obstacles at the ankle of the protrusion.

Acknowledgements.

We want to thank Martin Vogelaar (Groningen) for his help for solving Python programming problems and Anthony G. A. Brown (Leiden) for retrieving the list of O–, B–, and A– stars from the SIMBAD database. We also thank Marc William Pound and Mark Wolfire for their help on the PDR Toolbox. Studies of interstellar dust and gas at Leiden Observatory are supported by a Spinoza award from the Dutch Science agency, NWO. JRG thanks the Spanish MICINN for funding support under grant PID2019-106110GB-I00 partially based on IRAM 30 m telescope observations. IRAM is supported by INSU/CNRS (France), MPG (Germany), and IGN (Spain). This study was based on observations made with the NASA/DLR Stratospheric Observatory for Infrared Astronomy (SOFIA). SOFIA is jointly operated by the Universities Space Research Association Inc. (USRA), under NASA contract NAS2-97001, and the Deutsches SOFIA Institut (DSI), under DLR contract 50 OK 0901 to the University of Stuttgart. upGREAT is a development by the MPI für Radioastronomie and the KOSMA/Universität zu Köln, in cooperation with the DLR Institut für Optische Sensorsysteme. We acknowledge the work, during the C+upGREAT square degree survey of Orion, of the USRA and NASA staff of the Armstrong Flight Research Center in Palmdale, the Ames Research Center in Mountain View (California), and the Deutsches SOFIA Institut.

Appendix A Geometric correction factor

The limb-brightened shell observed in different tracers is seen as an arc of emission. If we assume that the emission seen in the dust tracer, the [C II] line, or the CO line is proportional to the total volume, then we need some geometry to figure out what the enhancement factor, fv, is that scales the volume of the limb-brightened part to that of the full shell. We considered two concentric nested ellipsoids with major diameters 2Co and 2Ci and minor diameters 2Bo and 2Bi. The protrusion is half of this ellipsoid (see Fig. A.1). If the cap height is h, the cap volume is given by (A.1)

The base surface area of the cap is (A.2)

The volume of the cylinder is (A.3)

The total volume of the outer ellipsoid is (A.4)

The volume of the rim is then (A.5)

We compared this to the volume in between the two nested ellipsoids, (A.6)

thumbnail Fig. A.1

Shell geometry.

As only half the ellipsoid protrudes out of the Veil, we should divide all of these volumes by two. However, as we are really interested in Vrim/V, these factors of two cancel each other out. Now we have to express h in the sizes of the inner and outer ellipsoid. The base area of the cap is equal to the surface area of the inner spheroid: (A.7)

Thus, h can be found from (A.8)

For Bo = 0.5 pc, Co = 1.3 pc, Bi = 0.4 pc, and Ci = 1.2 pc, we find that the height of the cap would be 0.244 pc. Using this, we estimate that the volume of the [C II] -emitting limb-brightened rim is 2.5% of the total volume of the half ellipse in Fig A.l. In this case, the mass in the limb-brightened shell would be 45 M, which is in good agreement with the mass estimation of 30 M based on the PDR models. Finally, the mass of the limb-brightened shell is between 30 and 45 M.

Appendix B Additional maps

Figures B.1 and B.2 show Spitzer 8 μm and Ha maps, respectively. In Fig. B.1, [C II] traces 8 μm closely.

thumbnail Fig. B.1

Spitzer 8 μm image, which outlines the PDR surfaces. The blue contours show the integrated SOFIA [C II] line emission.

thumbnail Fig. B.2

Hα image, which traces the ionized gas in the protrusion. The blue contours show the integrated SOFIA [C II] line emission. Green circles show the young stars and protostars surveyed by Megeath et al. (2005, 2012).

thumbnail Fig. B.3

WISE image of the Orion Nebula provided by the University of Berkeley. Blue represents emission at 3.4 μm, and cyan (blue-green) represents emission at 4.6 μm, both of which come mainly from hot stars. Relatively cooler objects, such as PAHs (i.e., the dust of the nebulae), appear green and red. Green and red represent 12 μm and 22 μm emission, respectively, which trace very small grains. The field of view of the original image is 3° × 3°, but we trimmed the image to only show the outflow beyond the protrusion. The original file can be retrieved at http://wise.ssl.berkeley.edu/gallery_OrionNebula.html.

thumbnail Fig. B.4

High-velocity CO emission. Left: Red- and blue-shifted lobes of suspected fossil OMC-1 outflow on a WISE image, including CO emission (red; Bally et al., in prep.). Right: Red-shifted lobe of suspected fossil OMC-1 outflow on [C II] (blue emission) and 12CO emission (red emission). In both panels, CO emission is integrated between +10 and +13 km s−1.

thumbnail Fig. B.5

Orion [S ii] image, which has been taken with the Víctor M. Blanco 4-Meter Telescope in Chile, showing the giant outflow structures.

thumbnail Fig. B.6

Temperature map (upper left) and optical depth at 160 μm map (τ160) (upper right) of dust emission, which traces the mass of the shell. The bottom panel shows an example of spectral energy distributions from the bottom left of the dust temperature map.

thumbnail Fig. B.7

Channel map of [C II] emission from VLSR from −3.2 to +13.6 km s−1 overlaid with 12CO J = 2–1 observations (white contours). The contour levels are [3, 6, 10, 15, 20] K km s−1. The velocity resolution of both maps is smoothed to 0.5 km s−1. Globule #10, which is a bright CO emission with a velocity of −0.8 km s−1, is the CO globule reported in Orion’s Veil (see also Fig. B.8; Goicoechea et al. 2020).

thumbnail Fig. B.8

pv diagram of 12CO J = 2–1 along crosscut 23 in Fig. 3. Globule #10, which is a bright CO emission at 200″, is the CO globule reported by Goicoechea et al. (2020).

Table B.1

List of O, B, and A stars within a 0.5′ circle centered at the Veil shell.

thumbnail Fig. B.9

[C II] pv diagrams from the protrusion sliced along the expansion direction (i.e., cuts 1 to 15 in Fig. 3). All diagram have the same scales as in Fig. 3.

thumbnail Fig. B.10

[C II] pv diagrams from the protrusion sliced along the expansion direction (i.e., cuts 16 to 30 in Fig. 3). All diagram have the same scales as in Fig. 3.

References

  1. Abel, N. P., Ferland, G. J., & O'Dell, C. R. 2019, ApJ, 881, 130 [NASA ADS] [CrossRef] [Google Scholar]
  2. Anderson, L. D., Bania, T. M., Balser, D. S., et al. 2014, ApJS, 212, 1 [Google Scholar]
  3. André, P., Men'shchikov, A., Bontemps, S., et al. 2010, A&A, 518, L102 [Google Scholar]
  4. Arce, H. G., Shepherd, D., Gueth, F., et al. 2007, in Protostars and Planets V, eds. B. Reipurth, D. Jewitt, & K. Keil, 245 [Google Scholar]
  5. Bally, J. 2010, Nature, 466, 928 [Google Scholar]
  6. Bally, J., Ginsburg, A., Arce, H., et al. 2017, ApJ, 837, 60 [Google Scholar]
  7. Bennett, C. L., Kogut, A., Hinshaw, G., et al. 1994, ApJ, 436, 423 [NASA ADS] [CrossRef] [Google Scholar]
  8. Berné, O., Marcelino, N., & Cernicharo, J. 2010, Nature, 466, 947 [CrossRef] [Google Scholar]
  9. Beuther, H., Schilke, P., Sridharan, T. K., et al. 2002, A&A, 383, 892 [Google Scholar]
  10. Castor, J., McCray, R., & Weaver, R. 1975, ApJ, 200, L107 [NASA ADS] [CrossRef] [Google Scholar]
  11. Cesaroni, R., Sanchez-Monge, A., Beltran, M. T., et al. 2017, A&A, 602, A59 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  12. Ceverino, D., & Klypin, A. 2009, ApJ, 695, 292 [Google Scholar]
  13. Churchwell, E. 2002, ARA&A, 40, 27 [Google Scholar]
  14. Churchwell, E., Povich, M. S., Allen, D., et al. 2006, ApJ, 649, 759 [Google Scholar]
  15. Da Rio, N., Robberto, M., Soderblom, D. R., et al. 2009, ApJS, 183, 261 [NASA ADS] [CrossRef] [Google Scholar]
  16. Feddersen, J. R., Arce, H. G., Kong, S., et al. 2018, ApJ, 862, 121 [NASA ADS] [CrossRef] [Google Scholar]
  17. Fierlinger, K. M., Burkert, A., Ntormousi, E., et al. 2016, MNRAS, 456, 710 [NASA ADS] [CrossRef] [Google Scholar]
  18. Fujimoto, Y., Chevance, M., Haydon, D. T., Krumholz, M. R., & Kruijssen, J. M. D. 2019, MNRAS, 487, 1717 [NASA ADS] [CrossRef] [Google Scholar]
  19. Genzel, R., Tacconi, L. J., Lutz, D., et al. 2015, ApJ, 800, 20 [Google Scholar]
  20. Goicoechea, J. R., Teyssier, D., Etxaluze, M., et al. 2015, ApJ, 812, 75 [Google Scholar]
  21. Goicoechea, J. R., Pabst, C. H. M., Kabanovic, S., et al. 2020, A&A, 639, A1 [Google Scholar]
  22. Grenier, I. A., Casandjian, J.-M., & Terrier, R. 2005, Science, 307, 1292 [Google Scholar]
  23. Griffin, M. J., Abergel, A., Abreu, A., et al. 2010, A&A, 518, L3 [EDP Sciences] [Google Scholar]
  24. Großschedl, J. E., Alves, J., Meingast, S., et al. 2018, A&A, 619, A106 [Google Scholar]
  25. Güdel, M., Briggs, K. R., Montmerle, T., et al. 2008, Science, 319, 309 [Google Scholar]
  26. Guzmán, V., Pety, J., Gratier, P., et al. 2012, A&A, 543, L1 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  27. Haid, S., Walch, S., Seifried, D., et al. 2018, MNRAS, 478, 4799 [Google Scholar]
  28. Hollenbach, D. J., & Tielens, A. G. G. M. 1997, ARA&A, 35, 179 [Google Scholar]
  29. Hosokawa, T., & Omukai, K. 2009, ApJ, 691, 823 [Google Scholar]
  30. Kaufman, M. J., Wolfire, M. G., Hollenbach, D. J., & Luhman, M. L. 1999, ApJ, 527, 795 [Google Scholar]
  31. Kavak, Ü., van der Tak, F. F. S., Tielens, A. G. G. M., & Shipman, R. F. 2019, A&A, 631, A117 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  32. Kavak, Ü., Sánchez-Monge, A., López-Sepulcre, A., et al. 2021, A&A, 645, A29 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  33. Kereš, D., Katz, N., Davé, R., Fardal, M., & Weinberg, D. H. 2009, MNRAS, 396, 2332 [CrossRef] [Google Scholar]
  34. Kong, S., Arce, H. G., Feddersen, J. R., et al. 2018, ApJS, 236, 25 [NASA ADS] [CrossRef] [Google Scholar]
  35. Kuiper, R., Klahr, H., Beuther, H., & Henning, T. 2011, ApJ, 732, 20 [Google Scholar]
  36. Kurtz, S. 2005, in Massive Star Birth: A Crossroads of Astrophysics, eds. R. Cesaroni, M. Felli, E. Churchwell, & M. Walmsley, 227, 111 [Google Scholar]
  37. Lopez, L. A., Krumholz, M. R., Bolatto, A. D., Prochaska, J. X., & Ramirez-Ruiz, E. 2011, ApJ, 731, 91 [Google Scholar]
  38. Lopez, L. A., Krumholz, M. R., Bolatto, A. D., et al. 2014, ApJ, 795, 121 [Google Scholar]
  39. López-Sepulcre, A., Cesaroni, R., & Walmsley, C. M. 2010, A&A, 517, A66 [Google Scholar]
  40. Maud, L. T., Moore, T. J. T., Lumsden, S. L., et al. 2015, MNRAS, 453, 645 [Google Scholar]
  41. Megeath, S. T., Flaherty, K. M., Hora, J., et al. 2005, in Massive Star Birth: A Crossroads of Astrophysics, eds. R. Cesaroni, M. Felli, E. Churchwell, & M. Walmsley, IAU Symp., 227, 383 [NASA ADS] [CrossRef] [Google Scholar]
  42. Megeath, S. T., Gutermuth, R., Muzerolle, J., et al. 2012, AJ, 144, 192 [Google Scholar]
  43. Menten, K. M., Reid, M. J., Forbrich, J., & Brunthaler, A. 2007, A&A, 474, 515 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  44. Motte, F., Bontemps, S., & Louvet, F. 2018, ARA&A, 56, 41 [Google Scholar]
  45. Naab, T., & Ostriker, J. P. 2017, ARA&A, 55, 59 [Google Scholar]
  46. Newman, R. C., & Axford, W. I. 1968, ApJ, 153, 595 [NASA ADS] [CrossRef] [Google Scholar]
  47. O'Dell, C. R. 2018, MNRAS, 478, 1017 [NASA ADS] [CrossRef] [Google Scholar]
  48. O'Dell, C. R., Kollatschny, W., & Ferland, G. J. 2017, ApJ, 837, 151 [CrossRef] [Google Scholar]
  49. Olivier, G. M., Lopez, L. A., Rosen, A. L., et al. 2021, ApJ, 908, 68 [Google Scholar]
  50. Pabst, C., Higgins, R., Goicoechea, J. R., et al. 2019, Nature, 565, 618 [NASA ADS] [CrossRef] [Google Scholar]
  51. Pabst, C. H. M., Goicoechea, J. R., Teyssier, D., et al. 2020, A&A, 639, A2 [Google Scholar]
  52. Pineda, J. L., Langer, W. D., & Goldsmith, P. F. 2014, A&A, 570, A121 [Google Scholar]
  53. Poglitsch, A., Waelkens, C., Geis, N., et al. 2010, A&A, 518, L2 [Google Scholar]
  54. Purser, S. J. D., Lumsden, S. L., Hoare, M. G., & Cunningham, N. 2018, MNRAS, 475, 2 [NASA ADS] [CrossRef] [Google Scholar]
  55. Quillen, A. C., Thorndike, S. L., Cunningham, A., et al. 2005, ApJ, 632, 941 [NASA ADS] [CrossRef] [Google Scholar]
  56. Risacher, C., Güsten, R., Stutzki, J., et al. 2016, A&A, 595, A34 [Google Scholar]
  57. Risacher, C., Güsten, R., Stutzki, J., et al. 2018, J. Astron. Instrum., 7, 1840014 [Google Scholar]
  58. Sánchez-Monge, Á., López-Sepulcre, A., Cesaroni, R., et al. 2013, A&A, 557, A94 [Google Scholar]
  59. Sanna, A., Moscadelli, L., Goddi, C., Krishnan, V., & Massi, F. 2018, A&A, 619, A107 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  60. Schneider, N., Simon, R., Guevara, C., et al. 2020, PASP, 132, 104301 [NASA ADS] [CrossRef] [Google Scholar]
  61. Spitzer, L. 1978, Physical Processes in the Interstellar Medium [Google Scholar]
  62. Tan, J. C., Beltrán, M. T., Caselli, P., et al. 2014, in Protostars and Planets VI, eds. H. Beuther, R. S. Klessen, C. P. Dullemond, & T. Henning, 149 [Google Scholar]
  63. Tielens, A. G. G. M. 2010, The Physics and Chemistry of the Interstellar Medium (Cambridge, UK: Cambridge University Press) [Google Scholar]
  64. Tielens, A. G. G. M., & Hollenbach, D. 1985, ApJ, 291, 722 [Google Scholar]
  65. Vacca, W. D., Garmany, C. D., & Shull, J. M. 1996, ApJ, 460, 914 [NASA ADS] [CrossRef] [Google Scholar]
  66. Walch, S. K., Whitworth, A. P., Bisbas, T., Wünsch, R., & Hubber, D. 2012, MNRAS, 427, 625 [NASA ADS] [CrossRef] [Google Scholar]
  67. Weaver, R., McCray, R., Castor, J., Shapiro, P., & Moore, R. 1977, ApJ, 218, 377 [Google Scholar]
  68. Weingartner, J. C., & Draine, B. T. 2001, ApJ, 548, 296 [Google Scholar]
  69. Williams, J. P., & McKee, C. F. 1997, ApJ, 476, 166 [NASA ADS] [CrossRef] [Google Scholar]
  70. Wilson, T. L., Filges, L., Codella, C., Reich, W., & Reich, P. 1997, A&A, 327, 1177 [NASA ADS] [Google Scholar]
  71. Wolfire, M. G., McKee, C. F., Hollenbach, D., & Tielens, A. G. G. M. 2003, ApJ, 587, 278 [Google Scholar]
  72. Young, E. T., Becklin, E. E., Marcum, P. M., et al. 2012, ApJ, 749, L17 [NASA ADS] [CrossRef] [Google Scholar]

1

The GLIMPSE survey has revealed the omnipresence of bubbles in the ISM (Churchwell et al. 2006). Most of these bubbles reflect the presence of H II regions (Anderson et al. 2014). The H II region bubble morphology may be driven by the thermal expansion of gas ionized by a central star, by the activity of a stellar wind during the main-sequence phase, or by fossil cavities created by now extinct energy and momentum sources, such as protostellar outflows. Throughout the paper, we presume that the Orion Nebula is mainly driven by stellar winds from the Trapezium stars (Pabst et al. 2019).

2

SOFIA is a Boeing 747SP aircraft modified to carry a 2.7-m telescope (Young et al. 2012).

3

FEEDBACK is a SOFIA legacy program dedicated to studying the interaction of massive stars with their environment. It performs a survey of 11 Galactic high-mass star-forming regions in the 158 μm (1.9 THz) line of [C II] and the 63 μm (4.7 THz) line of [O I].

4

We used 414 pc, provided by Menten et al. (2007), as the distance. The OMC does show a substantial distance gradient (Großschedl et al. 2018), but it is on a much larger scale and not relevant for our paper.

5

The WISE image of the Extended Orion Nebula can be retrieved via: http://wise.ssl.berkeley.edu/gallery_OrionNebula.html

6

The software package for Berschel Interactive Processing Environment (HIPE) is designed to work with the Berschel data and find the data products, conduct an interactive analysis, plot data, and perform data manipulation.

9

The EM is defined as , where ne is the electron density and L is the total path length in the ionized gas.

10

More information on object type is available at http://simbad.u-strasbg.fr/simbad/sim-display?data=otypes

11

In the Gaia DR2 survey, Großschedl et al. (2018) reported a parallax of 2.4792 ± 0.0374 mas, which corresponds to 404 ± 6.1 pc, indicating that Star 39 could be associated with the Orion Nebula.

Movie

Movie 1 associated with Fig. 5 (Kavaketal2022_BreakingVeil_OrionCII_channelmaps) (Access here)

All Tables

Table 1

Comparison of the masses and energetics of the protrusion with the Veil shell.

Table B.1

List of O, B, and A stars within a 0.5′ circle centered at the Veil shell.

All Figures

thumbnail Fig. 1

Integrated intensity [C II] 158 μm map. Left: integrated (between −5 and +14 km s−1) intensity [C II] 158 μm map of the OMC observed by the upGREAT receiver on board SOFIA. The positions of NGC 1977, Trapezium stars, M42, M43, and the Orion Bar PDR are labeled. The green box shows the extracted region from the map, including the area of interest for this study, the protrusion. Right: close-up view of the protrusion. The bright ridge of emission is the edge of the expanding Veil shell. Faint emission extends well beyond this shell – the protrusion of which has a multi-component structure. More details about the kinematics and components of the protrusion are given in Sect. 3.

In the text
thumbnail Fig. 2

Images of Orion's protrusion at different wavelengths and angular resolutions. The observed transition or frequency is given for each panel. The [C II],12CO (2−1), and 13CO (2−1) observations are integrated between −5 and +14 km s−1.

In the text
thumbnail Fig. 3

Two pv diagrams of the protrusion. Top: selected crosscuts along the green arrows overlaid on the integrated [C II] intensity map. The number of the crosscuts is indicated at the starting point of the cut. Middle and bottom: pv diagrams generated along the magenta crosscuts (cuts 8 and 23, respectively). The pv diagrams with horizontal green lines both show the [C II] emission produced by the FUV-illuminated surface of OMC, and the arcuate white and yellow lines trace the shell expanding at 12 km s−1 and 6 km s−1, respectively. The pv diagrams in Figs. B.9 and B.10 have the same scale on both axes. A 12CO-pv diagram along crosscut 23 is shown in Fig. B.8 for comparison with [C II].

In the text
thumbnail Fig. 4

Spectral lines toward the protrusion. Upper panel: Spitzer 8 μm image of the protrusion. Red circles indicate the 12 positions that we used to extract line profiles and have an aperture of 16′′. Lower panel: velocity-resolved spectra of [C II] (colored in gray), 12CO J = 2−1 (blue), and 13CO J = 2−1 (cyan) in the direction of protrusion for the 12 selected positions in the upper panel. The vertical dotted red line at 9 km s−1 marks the approximate velocity of the emission generated by the OMC and the associated star-forming molecular cloud behind the Veil.

In the text
thumbnail Fig. 5

Three-color image of the protrusion created using the three components estimated in the [C II] line profile in Fig. 4. Blue emission is the integrated emission between −5 and +3 km s−1, green between +3 and +12 km s−1, and red between +12 and +15 km s−1 of the SOFIA [C II] 158 μm (upper panel) and IRAM 12CO (lower panel) emission maps. White circles show the 12 selected positions from Fig. 4. A movie is available online.

In the text
thumbnail Fig. 6

Map of the incident radiation field, G0 (left) and the density (right) of the protrusion for a face-on PDR model adopted from Tielens (2010) (see Sect. 4.4 for a more detailed discussion).

In the text
thumbnail Fig. 7

SOFIA [C II] map of Orion with O, B, and A stars found in SIM-BAD. The list of stars retrieved from the archive is given in Table B.l. The blue, orange, and red circles are O, B, and A stars, respectively. The light green arrows indicate the positions of the Trapezium stars. Two A stars (Star 21 and 39 in Table B.l), which are the nearest to the protrusion, are also labeled.

In the text
thumbnail Fig. 8

Schematic picture of the protrusion and the fossil bipolar outflow with apparent structures overlaid on the WISE image (see also Fig. B.4). the suspected blue- and red-shifted outflow lobes are shown in blue and red, respectively. The Kelvin-Helmholtz instabilities reported by Berné et al. (2010) are indicated in yellow.

In the text
thumbnail Fig. 9

Momentum of outflows from massive young stellar objects as a function of the source luminosity of the cores (Maud et al. 2015). The blue and red symbols indicate the blue- and red-shifted outflow lobe values, respectively, which are joined by a dotted line for each source. The horizontal blue-shaded range indicates the momentum of the protrusion, which is between 360 and 540 M km s−1. The cross at the bottom right shows the uncertainty for both axes.

In the text
thumbnail Fig. A.1

Shell geometry.

In the text
thumbnail Fig. B.1

Spitzer 8 μm image, which outlines the PDR surfaces. The blue contours show the integrated SOFIA [C II] line emission.

In the text
thumbnail Fig. B.2

Hα image, which traces the ionized gas in the protrusion. The blue contours show the integrated SOFIA [C II] line emission. Green circles show the young stars and protostars surveyed by Megeath et al. (2005, 2012).

In the text
thumbnail Fig. B.3

WISE image of the Orion Nebula provided by the University of Berkeley. Blue represents emission at 3.4 μm, and cyan (blue-green) represents emission at 4.6 μm, both of which come mainly from hot stars. Relatively cooler objects, such as PAHs (i.e., the dust of the nebulae), appear green and red. Green and red represent 12 μm and 22 μm emission, respectively, which trace very small grains. The field of view of the original image is 3° × 3°, but we trimmed the image to only show the outflow beyond the protrusion. The original file can be retrieved at http://wise.ssl.berkeley.edu/gallery_OrionNebula.html.

In the text
thumbnail Fig. B.4

High-velocity CO emission. Left: Red- and blue-shifted lobes of suspected fossil OMC-1 outflow on a WISE image, including CO emission (red; Bally et al., in prep.). Right: Red-shifted lobe of suspected fossil OMC-1 outflow on [C II] (blue emission) and 12CO emission (red emission). In both panels, CO emission is integrated between +10 and +13 km s−1.

In the text
thumbnail Fig. B.5

Orion [S ii] image, which has been taken with the Víctor M. Blanco 4-Meter Telescope in Chile, showing the giant outflow structures.

In the text
thumbnail Fig. B.6

Temperature map (upper left) and optical depth at 160 μm map (τ160) (upper right) of dust emission, which traces the mass of the shell. The bottom panel shows an example of spectral energy distributions from the bottom left of the dust temperature map.

In the text
thumbnail Fig. B.7

Channel map of [C II] emission from VLSR from −3.2 to +13.6 km s−1 overlaid with 12CO J = 2–1 observations (white contours). The contour levels are [3, 6, 10, 15, 20] K km s−1. The velocity resolution of both maps is smoothed to 0.5 km s−1. Globule #10, which is a bright CO emission with a velocity of −0.8 km s−1, is the CO globule reported in Orion’s Veil (see also Fig. B.8; Goicoechea et al. 2020).

In the text
thumbnail Fig. B.8

pv diagram of 12CO J = 2–1 along crosscut 23 in Fig. 3. Globule #10, which is a bright CO emission at 200″, is the CO globule reported by Goicoechea et al. (2020).

In the text
thumbnail Fig. B.9

[C II] pv diagrams from the protrusion sliced along the expansion direction (i.e., cuts 1 to 15 in Fig. 3). All diagram have the same scales as in Fig. 3.

In the text
thumbnail Fig. B.10

[C II] pv diagrams from the protrusion sliced along the expansion direction (i.e., cuts 16 to 30 in Fig. 3). All diagram have the same scales as in Fig. 3.

In the text

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