# Harvard Astronomy 201b

## ARTICLE: Interpreting Spectral Energy Distributions from Young Stellar Objects

In Journal Club, Journal Club 2013 on April 15, 2013 at 8:02 pm

Posted by: Meredith MacGregor

#### 1. Introduction

When discussing young stellar objects and protoplanetary disks (as well as many other topics in astronomy), astronomers continually throw out the term ‘SED.  In early April, I attended a conference titled ‘Transformational Science with ALMA: From Dust to Rocks to Planets– Formation and Evolution of Planetary Systems.’  And, I can attest to the fact that the term ‘SED’ has come up in a very significant fraction of the contributed talks.  For those not intimately familiar with this sub-field, this rampant use of abbreviations can be confusing, making it difficult to glean any useful take-aways from a talk.  So, in addition to summarizing the article by Robitaille et al. (2007), the goal of this post is to give a bit of an introduction to the terminology and motivation for the growing field of star and planetary system formation.

SEDs: What They Are and Why We Care So Much

The abbreviation ‘SED’ stands for ‘Spectral Energy Distribution.’  If you want to sound like an expert, this should be pronounced exactly as it appears (analogous to the non-acronym ‘said’).  A SED is essentially a graph of flux versus wavelength.  In the context of the Robitaille et al. article, we are most interested in the SEDs for young stars and the envelopes and disks surrounding them.  So, why exactly does the flux from a young stellar object (YSO) vary with wavelength?  As it turns out, different regions of the YSO emit at different wavelengths.  This means that when we observe at different wavelengths, we are actually probing distinct regions of the star and its surrounding disk and envelope. By tracing out the entire SED for a YSO, we can determine what the geometry, structure, and constituents are for that object.

Before we go too far, it is worth taking a moment to clarify the terminology used to classify young stars.  A young stellar object or YSO is typically defined as any star in the earliest stages of development.  YSOs are almost always found in or near clouds of interstellar dust and gas.  The broad YSO class of objects is then divided into two sub-classes: protostars and pre-main sequence stars.  Protostars are heavily embedded in dust and gas and are thus invisible at optical wavelengths.  Astronomers typically use infrared, sub-millimeter, and millimeter telescopes to explore this stage of stellar evolution, where a star acquires the bulk of its material via infall and accretion of surrounding material.  Pre-main sequence (PMS) stars are defined to be low mass stars in the post-protostellar phase of evolution.  These stars have yet to enter the hydrogen-burning phase of their life and are sill surrounded by remnant accretion disks called ‘protoplanetary’ disks.  A more detailed a summary of this classification scheme can be found here.  However, because the evolution of young stars is far from well-understand, astronomers often typically these words interchangeably.

SEDs for pre-main sequence stars are often seen to have a bump or excess in the infrared. This bump is generally interpreted as being due to thermal emission from warm dust and gas surrounding the central star. As an illustrative example, let’s consider a protoplanetary disk around a young star. The image below is taken from a review article by Dullemond & Monnier (2010) and shows a graphical representation of the emission spectrum from such a disk.

A graphical representation of the emission spectrum from a protoplanetary disk and the telescopes that can be used to probe different regions of the disk. Taken from Dullemond & Monnier (2010).

In this case, emission in the near-infrared traces the warmer inner regions of the disk.  As you move into submillimeter wavelengths, you start to probe the outer, cooler regions of the disk.  By modeling the observed SED of a pre-main sequence star, you can determine what components of the source are contributing to the observed flux.  The second figure below is taken from a paper by Guarcello et al. (2010).  The left panel of this figure shows the observed SED (solid line) of a BWE star (in line with other ridiculous astronomy acronyms, this stands from ‘blue stars with excesses’) and the unreddened photospheric flux we would expect to see if the star did not have a disk (dashed line).  The right panel shows a model fit to this data.  The authors describe the observed SED using a model with four distinct components, each represented by a colored line in the figure: (1) emission from the reddened photosphere of the central star, (2) radiation scattered into the line of sight from dust grains in the disk, (3) emission from a collapsing outer envelope, and (4) thermal emission from a circumstellar disk.  The summation of these four component makes up the complete SED for this BWE star.

Left panel: The observed SED (solid line) of a BWE star and the unreddened photospheric flux we would expect to see if the star did not have a disk (dashed line). Right panel: A four component model fit to this data. Taken from Guarcello et al. (2010).

Describing and Classifying the Evolution of a Protostar

As a protostar evolves towards the Zero Age Main Sequence (ZAMS), the system geometry (and thus the SED) will evolve as well. Therefore, the stage of evolution of a protostar is often classified according to both the general shape and the features of the SED.  A graphical overview of the four stages of protostellar evolution are shown below (Andrea Isella’s thesis, 2006).  Class 0 objects are characterized by a very embedded central core in a much larger accreting envelope.  The mass of the central core grows in Class I objects and a flattened circumstellar accretion disk develops.  For Class II objects, the majority of circumstellar material is now found in a disk of gas and dust.  Finally, for Class III objects, the emission from the disk becomes negligible and the SED resembles a pure stellar photosphere.  The distinction between these different classes was initially defined by the slope of the SED (termed the ‘spectral index’) at infrared wavelengths.  Class I sources typically have SEDs that rise in the far- and mid-infrared, while Class II sources have flat or falling SEDs in the mid-infrared.  However, this quantitative distinction is not always clear and is not an effective way to unambiguously distinguish between the different object classes (Protostars and Planets V, page 127-128).

A graphical overview of the four stages of protostar evolution taken from Andrea Isella’s thesis (2006). A typical SED of each class is shown in the left column and a cartoon of the corresponding geometry is shown in the right column.

#### 2. The Article Itself

One common method of fitting SEDs is to assume a given gas and circumstellar dust geometry and set of dust properties (grain size distribution, composition, and opacity), and then use radiative transfer models to predict the resulting SED and find a set of parameters that best reproduce the observations.  However, fitting SEDs by trial and error is a time consuming way to explore a large parameter space.  The problem is even worse if you want to consider thousands of sources.  So, what’s to be done?  Enter Robitaille et al.  In order to attempt to make SED fitting more efficient, they have pre-calculated a large number of radiative transfer models that cover a reasonable amount of parameter space.  Then, for any given source, one can compare the observed SED to this set of models to quickly find the set of parameters that best explains the observations.

Let’s Get Technical

The online version of this fitting tool draws from 20,000 combinations of physical parameters and 10 viewing angles (if you are particularly curious, the online tool is available here).  A brief overview of the parameter space covered is as follows:

• Stellar mass between 0.1 and 50 solar masses
• Stellar ages between $10^3$ and $10^7$ years
• Stellar radii and temperatures (derived directly from stellar mass using evolutionary tracks)
• Disk parameters (disk mass, accretion rate, outer radius, inner radius, flaring power, and scale height) and envelope parameters (the envelope accretion rate, outer radius, inner radius, cavity opening angle, and cavity density) sampled randomly within ranges dictated by the age of the source

However, there are truly a vast number of parameters that could be varied in models of YSOs.  Thus, for simplicity, the authors are forced to make a number of assumptions.  Here are some of the biggest assumptions involved:

1. All stars form via accretion through a disk and an envelope.
2. The gas-to-dust ratio in the disk is 100.
3. The apparent size of the source is not larger than the given aperture.

The last constraint is not required, but it allows a number of model SEDs to be cut out and thus speeds up the process.  Furthermore, the authors make a point of saying that the results can always be scaled to account for varying gas-to-dust ratios, since only the dust is taken into account in the actual radiative transfer calculations.

Does This Method Really Work?

If this tool works well, it should be able to correctly reproduce previous results.  In order to test this out, the authors turn to the Taurus-Auriga star forming region.  They select a sample of 30 sources from Kenyon & Hartman (1995) that are spatially resolved, meaning that there is prior knowledge of their evolutionary stage from direct observations (i.e. there is a known result that astronomers are fairly certain of to compare the model fits against).  When fitting their model SEDs with the observed SEDs for this particular star forming region, the authors throw in a few additional assumptions:

1. All sources are within a distance range of 120 – 160 AU (helps to rule out models that are too faint or too luminous).
2. The foreground interstellar extinction is no more than $A_v$ = 20.
3. None of the sources appeared larger than the apertures used to measure fluxes.

The authors then assign an arbitrary cut-off in chi-squared for acceptable model fits: $\chi^2 - \chi^2_\text{best} < 3$.  Here, $\chi^2_\text{best}$ is the $\chi^2$ of the best-fit model for each source.  Robitaille et al. acknowledge that this cut-off has no statistical justification: ‘Athough this cut-off is arbitrary, it provides a range of acceptable fits to the eye.’  After taking Jim Moran’s Noise and Data Analysis class, I for one would like to see the authors try a Monte Carlo Markov Chain (MCMC) analysis of their 14-dimensional space (for more detail on MCMC methods see this review by Persi Diaconis).  That might make the analysis a bit less ‘by eye’ and more ‘by statistics.’

The upshot of this study is that for the vast majority of the sources considered, the best-fit values obtained by this new SED fitting tool are close to the previously known values. Check.

It is also worth mentioning here, that there are many other sets of SED modeling codes.  One set of codes of particular note are those written by Paola D’Alessio (D’Alessio et al., 1998; D’Alessio et al., 1999; D’Alessio et al, 2001).  These codes were the most frequently used in the results presented at the ALMA conference I attended.  The distinct change in the D’Alessio models is that they solve for the detailed hydrostatic vertical disk structure in order to account for observations of ‘flared’ disks around T Tauri stars (flaring refers to an increase in disk thickness at larger radii).

But, Wait! There are Caveats!

Although the overall conclusion is that this method fits SEDs with reasonable accuracy, there are a number of caveats that are raised.  First of all, the models tend to overestimate the mid-IR fluxes for DM Tau and GM Aur (two sources known to have inner regions cleared of dust).  The authors explain that this is most likely due to the fact that their models for central holes assume that there is no dust remaining in the hole.  In reality, there is most likely a small amount of dust that remains.  Second, the models do not currently account for the possibility of young binary systems and circumbinary disks (relevant for CoKu Tau 1).

The paper also addresses estimating parameters such as stellar temperature, disk mass, and accretion rate from SED fits.  And, yes, you guessed it, these calculations raise several more issues.  For very young objects, it is difficult to disentangle the envelope and disk, making it very challenging to estimate a total disk mass.  To make these complications clearer, the set of two plots below from the paper show calculated values for the disk mass plotted against the accepted values from the literature.  It is easily seen that the disk masses for the embedded sources are the most dissimilar from the literature values.

Two plots from Robitaille et al. (2007) that show calculated values for the disk mass plotted against the accepted values from the literature. It is easily seen that the disk masses for the embedded sources (right) are the most dissimilar from the literature values.

Furthermore, even if the disk can be isolated, the dust mass in the disk is affected by the choice of dust opacity.  That’s a pretty big caveat!  A whole debate was started at the ALMA conference over exactly this issue and the authors have simply stated the problem and swept it under the rug in just one sentence.  In 2009, David Hogg conducted a survey of the rho Ophiucus region and used the models of Robitaille et al. (2007) to determine the best-fit dust opacity index, $\beta$ for this group of sources.   Hogg found that $\beta$ actually decreases for Class II protostars, a possible indication of the presence of larger grains in the disk.  Robitaille et al. also mention that the calculated accretion rates from SED fitting are systematically larger than what is presented in the literature.  The authors conclude that future models should include disk emission inside the dust destruction radius, the radius inside which it is too hot for dust to survive.  A great example of the complications that arise from a disk with a central hole can be seen in LkCa 15 (Espaillat et al., 2010Andrews et al., 2011).   The figure below shows the observed and simulated SEDs for the source (left) as well as the millimeter image (right).  The double ansae (bright peaks or ‘handles‘ apparent on either side of the disk) seen in the millimeter contours are indicative of a disk with a central cavity.

Left: The observed and simulated SEDs for Lk Ca 15. The sharp peak seen at 10 microns is due to silicate grains within the inner hole of the disk. Right: The millimeter image of the disk. (Espaillat et al., 2010; Andrews et al., 2011)

In this case, a population of sub-micron sized dust within the hole is needed in order to produce the observed silicate feature at 10 microns.  Furthermore, an inner ring is required to produce the strong near-IR excess shortward of 10 microns.  A cartoon image of the predicted disk geometry is shown below.  To make things even more complicated, the flux at shorter wavelengths appears to vary inversely with the flux at longer wavelengths over time (Espaillat et al., 2010).  This phenomenon is explained by changing the height of the inner disk wall over time.

A cartoon image of the predicted disk geometry for Lk Ca 15 showing the outer ring, silicate grains within the hole, and the inner ring. (Espaillat et al., 2010)

Finally, Robitaille et al. discuss how well parameters are constrained given different combinations of data points for two example sources: AA Tau and IRAS 04361+2547.  In both sources, if only IRAC (Infrared Array Camera on the Spitzer Space Telescope) fluxes obtained between 3.6 and 8 microns are used, the stellar mass, stellar temperature, disk mass, disk accretion rate, and envelope accretion rate are all poorly constrained.  Things are particularly bad for AA Tau in this scenario, where only using IRAC data results in ~5% of all SED models meeting the imposed goodness of fit criterion (yikes!).  Adding in optical data to the mix helps to rule out models that have low central source temperatures and disk accretion rates.  Adding data at wavelengths longer than ~ 20 microns helps to constrain the evolutionary stage of the YSO, because that is where any infrared excess is most apparent.  And, adding submillimeter data helps to pin down the disk mass, since the emission at these longer wavelengths is dominated by the dust.  This just goes to show how necessary it is to obtain multi-wavelength data if we really want to understand YSOs, disks, and the like.

#### 3. Sources:

Where to look if you want to read more about anything mentioned here…

## ARTICLE: An L1551 Extravaganza: Three Articles

In Journal Club 2013 on April 1, 2013 at 11:55 am

Wide-Field Near-Infrared Imaging of the L1551 Dark Cloud by Masahiko Hayashi and Tae-Soo Pyo

Observations of CO in L1551 – Evidence for stellar wind driven shocks by Ronald L. Snell, Robert B. Loren & Richard L. Plambeck

Multiple Bipolar Molecular Outflows from the L1551 IRS5 Protostellar System by Po-Feng Wu, Shigehisa Takakuwa, and Jeremy Lim

Summary by Fernando Becerra, Lauren Woolsey, and Walker Lu

## Introduction

### Young Stellar Objects and Outflows

In the early stages of star formation, Young Stellar Objects (YSOs) produce outflows that perturb the surrounding medium, including their parental gas cloud. The current picture of star formation indicates that once gravity has overcome pressure support, a central protostar is formed surrounded by an infalling and self-supported gas disk. In this context outflows are powered by the release of gravitational potential energy liberated by matter accreting onto the protostar. Outflows are highly energetic and often spatially extended phenomena, and are observable over a wide range of wavelengths from x-ray to the radio. Early studies of molecular outflows (predominantly traced by CO emission lines, e.g. Snell et al. 1980, see below) have shown that most of their momentum is deposited in the surrounding medium and so provide a mass loss history of the protostar. In contrast, the optical and near-infrared (NIR) emission trace active hot shocked gas in the flow.

### Interactions with the surrounding medium: Herbig-Haro objects, bow shocks and knots

When outflows interact with the medium surrounding a protostar, emission can often be produced. One example of this is emission from Herbig-Haro (HH) objects, which can be defined as “small nebulae in star-forming regions as manifestations of outflow activity from newborn stars”. The most common pictures show a HH object as a well-collimated jet ending in a symmetric bow shock. Bow shocks are regions where the jet accelerates the ambient material. The shock strength should be greatest at the apex of the bow, where the shock is normal to the outflow, and should decline in the wings, where the shocks become increasingly oblique. Another interesting feature we can distinguish are knots. Their origin is still unknown but a few theories have been developed over the years. They can formed due to the protostar producing bursts of emission periodically in time, or producing emission of varying intensity. They can also form due to interactions between the jet and the surrounding Interstellar Medium (ISM), or due to different regions of the jet having different velocities.

## An exceptional case: The L1551 region

The L1551 system is an example of a region in which multiple protostars exhibiting outflows are seen, along with several HH objects and knots. This system has been catalogued for over fifty years (Lyons 1962), but ongoing studies of the star formation and dynamical processes continue to the present day (e.g. Hayashi and Pyo 2009; Wu et al. 2009). L1551 is a dark cloud with a diameter of ~20′ (~1 pc) located at the south end of the Taurus molecular cloud complex. The dark cloud is associated with many young stellar objects. These YSOs show various outflow activities and characteristics such as optical and radio jets, Herbig-Haro objects, molecular outflows, and infrared reflection nebulae. We will start by giving a broad view of the region based on Hayashi and Pyo 2009, and then we will focus on a subregion called L1551 IRS 5 following Snell et al. 1980 and Wu et al. 2009.

## Paper I: An overview of the L1551 Region (Hayashi and Pyo 2009)

The L1551 region is very rich in YSOs, outflows and their interaction with the ISM. The most prominent of the YSOs in this region are HL Tau, XZ Tau, LkHα 358, HH 30, L1551 NE, and L1551 IRS 5 (see Fig. 1), arrayed roughly north to south and concentrated in the densest part (diameter ~10′) of the cloud. The authors based their study on observations using two narrowband filters [Fe II] ($\lambda_c$ = 1.6444 μm, $\Delta\lambda$ = 0.026 μm), $H_2$ ($\lambda_c$ = 2.116 μm, $\Delta\lambda$ = 0.021 μm) and two broad-band filters: $H$ ($\lambda_c$ = 1.64 μm, $\Delta\lambda$ = 0.28 μm) $K_s$ ($\lambda_c$ = 2.14 μm, $\Delta\lambda$ = 0.31 μm). The choice of [Fe II] and $H_2$ is motivated by previous studies suggesting that the [Fe II] line has higher velocity than the $H_2$, and thus arises in jet ejecta directly accelerated near the central object, while $H_2$ emission may originate in shocked regions. In the particular case of bow shocks, regions of higher excitation near the apex are traced by [Fe II], while $H_2$ is preferentially found along bow wings. The broadband filters were chosen for comparison with NIR narrowband filters and comparison with previous studies. The total sky coverage was 168 arcmin2, focused on 4 regions of the densest part of the L1551 dark cloud, including HL/XZ Tau, HH30, L1551 IRS5, some HH objects to the west, L1551 NE, and part of HH 262 (see Fig. 1).

Figure 1: An overview of L1551 (Figure 1 of Hayashi and Pyo 2009)

HL/XZ Region

Some of the features the authors identify in this region are:

• A faint [Fe II] jet emanating from HL Tau to its northeast and southwest. The $H_2$ emission is hard to identify in the northeast part, but significant $H_2$ emission blobs are detected in the southwest part (denoted “H2 jet” in Fig. 2)
• A diffuse feature is also distinguished to the north-northeast of XZ Tau, which may be related to the outflow from one member of the XZ Tau binary.
• A continuum arc from HL Tau to the north and then bending to the east (“cont arc” in Fig. 2) is also identified. This arc may be a dust density discontinuity where enhanced scattering is observed. Although it is not clear if this arc is related to activities at HL Tau or XZ Tau.
• Another arc feature to the south from HL Tau curving to the southeast can be identified. Two $H_2$ features are located in the arc and indicated by arrows in Fig. 2. This may be shocked regions in the density discontinuity.
• Other $H_2$ features can be distinguished: “A” (interpreted as a limb-brightened edge of the XZ Tau counter-outflow) and “B”, “C”, “a” (blobs driven by the LkH$\alpha$ 358 outflow and interacting with the southern outflow bubble of XZ Tau).

Figure 2: HL/XZ Region (Figure 2 of Hayashi and Pyo 2009)

HH 30 Region

HH 30 is a Herbig-Haro (HH) object including its central star, which is embedded in an almost edge-on flared disk. Although this object doesn’t have clear signs of large-scale [Fe II] or $H_2$ emission (see Fig. 3), a spectacular jet was detected in the [S II] emission line in previous studies. Despite that, the authors identify two faint small-scale features based on the [Fe II] frame: one to the northeast (corresponding to the brightest part of the [S II] jet) and one to the south-southeast (corresponding to a reflection nebula)

Figure 3: HH 30 Region (Figure 3 of Hayashi and Pyo 2009)

L1551 NE

L1551 NE is a deeply embedded object associated with a fan-shaped infrared reflection nebula opening toward the west-southwest seen in the broad-band Ks continuum emission. It has an opening angle of $60^o$. The most important features in this region are:

• A needle-like feature connecting L1551 NE and HP2 is distinguished from the continuum-substracted [Fe II] image, associated with an [Fe II] jet emanating from L1551 NE.
• A diffuse red patch at the southwest end of the nebula (denoted as HP1) is dominated by $H_2$ emission
• Five isolated compact features are detected in the far-side reflection nebula: HP3 and HP3E ([Fe II] emission), HP4 (both [Fe II] and $H_2$ emission) and HP5 and HP6 ($H_2$ emission). All of them are aligned on a straight line that is extrapolated from the jet connecting NE and HP2, naturally assigned to features on the counter-jet.
• Comparing this data to previous observations in [S II] and radio we can deduce radial velocities of 160-190 km/s for HP2, and 140-190 km/s for HP4 and HP5. With radial velocities in the range 100-130 km/s for these knots, the inclination of the jet axis is estimated to be $45^o$$60^o$.

L1551 IRS-5

L1551 IRS 5 is a protostellar binary system with a spectacular molecular outflow (Snell et al. 1980; see below) and a pair of jets emanating from each of the binary protostars. A conspicuous fan-shaped infrared reflection nebula is seen in Fig. 4, widening from IRS 5 toward the southwest. At the center of this nebula, the two [Fe II] jets appear as two filaments elongated from IRS 5 to its west-southwest; the northern jet is the brighter of the two. Knots A, B and C located farther west and west-southwest of PHK3 (associated with $H_2$ line emission) have significant [Fe II] emission.

Figure 4: A close-up of IRS-5 (Figure 5 of Hayashi and Pyo 2009)

A counter-jet only seen in the [Fe II] frame can be distinguished to the northeast of IRS 5. Considering its good alignment with the northern jet, it can be interpreted as the receding part of the jet. Based on brightness comparison between the both jets, and transforming H-band extinction to visual extinction the authors deduce a total visual extinction of Av=20-30 mag. Besides the counter-jet, the authors also detect the northern and southern edge of the reflection nebula that delineate the receding-side outflow cone of IRS5.

A brief summary of the HH objects detected in the IRS5 region:

• HH29: Consistent with a bow shock, its [Fe II] emission features are compact, while the $H_2$ emission is diffuse. Both emissions are relatively separate.
• HH260: Consisted with a bow shock with compact [Fe II] emission knot located at the apex of a parabolic $H_2$ emission feature.
• HP7: Its [Fe II] and $H_2$ emission suggest it is also a bow shock driven by an outflow either from L1551 IRS5 or NE.
• HH264: It is a prominent $H_2$ emission loop located in the overlapping molecular outflow lobes of L1551 IRS5 and NE. Its velocity gradients are consistent with the slower material surrounding a high-velocity (~ -200 km/s in radial velocity) wind axis from L1551 IRS 5 (or that from L1551 NE)
• HH 102: Loop feature dominated by $H_2$ emission (and no [Fe II] emission) similar to HH264. Considering that the major axes of the two elliptical features are consisted with extrapolated axis of the HL Tau jet, it is suggested that they might be holes with wakes on L1551 IRS5 and/or NE outflow lobe(s) that were bored by the collimated flow from HL Tau.

### Comparison of Observations

Near-infrared [Fe II] and $H_2$ emission show different spatial distributions in most of the objects analyzed here. On one hand the [Fe II] emission is confined in narrow jets or relatively compact knots. On the other hand, the $H_2$ emission is generally diffuse or extended compared with the [Fe II] emission, with none of the $H_2$ features showing the well collimated morphology as seen in [Fe II].
These differences can be understood based on the conditions that produce different combinations of [Fe II] and $H_2$ emission:

• Case of spatially associated [Fe II] and $H_2$ emissions: Generally requires fast dissociative J shocks (Hollenbach & McKee 1989; Smith 1994; Reipurth et al. 2000).
• Case of a strong $H_2$ emission without detectable [Fe II] emission: Better explained by non-dissociative C shocks

The interpretation of differences in [Fe II] and $H_2$ emission as a result of distinct types shocks is supported by observational evidence showing that the [Fe II] emission usually has a much higher radial velocity than the $H_2$ emission. In the case of HH 29, HH 260 and HP 7 the [Fe II] emission arises in the bow tips where the shock velocity is fast (~50 km/s) and dissociative whereas $H_2$ emission occurs along the trailing edges where the shock is slower (~20 km/s)

## Paper II: Landmark Observations of Snell et al. 1980

One of the original papers in the study of L1551 was written by Snell, Loren and Plambeck (1980). In this paper, the authors use 12CO to map what they find to be a double-lobed structure extending from the infrared source IRS-5 (see Figures 1, 4). This system is also associated with several Herbig-Haro objects, which are small dense patches that are created in the first few thousands of years after a star is formed. This star is consistent with a B star reddened by 20 magnitudes of dust extinction along the line of sight, through a distance of 160 pc (Snell 1979). By studying these outflows, we are able to better understand the evolution of YSOs.

### Observations

Snell et al. (1980) made their observations using the 4.9 meter antenna at Millimeter Wave Observatory in Texas. Specifically, they considered the J = 1-0 and J = 2-1 transitions of $^{12}$CO and $^{13}$CO. Additionally, they made J = 1-0 observations with the NRAO 11 meter antenna. They found asymmetries in the spectral lines, shown below in Figure 5. To the northeast of IRS-5, the high-velocity side of the line has a broad feature, and the southwest of IRS-5 presents a similar broad feature on the low-velocity side of the spectral line. No such features were found to the NW, SE, or in the central position of IRS-5.

Figure 5: 12CO and 13CO 1-0 transition lines; top is NE of central source, bottom is SW of source (Figure 4 of Snell et al. 1980)

The J = 2-1 $^{12}$CO transition is enhanced relative to the J = 1-0 transition of $^{12}$CO, suggesting that the $^{12}$CO emission is not optically thick. If the emission was optically thick, the J = 1-0 line would be the expected dominant transition as it is a lower level transition. The observations also suggest an excitation temperature for the 2-1 transition of T$_{ex}$ ~ 8-35 K. This would only relate to the gas temperature if the environment is in local thermal equilibrium, but it does set a rough minimum temperature. The $^{13}$CO emission for the 1-0 transition is roughly 40 times weaker than the same transition for $^{12}$CO, which further suggests both isotopes are optically thin in this region (if the $^{12}$CO is already optically thin, the weaker transition means $^{13}$CO is even more so). The geometry of the asymmetries in the line profiles seen to the NE and SW combined with the distance to L1551 suggest lobes that extend out 0.5 pc in both directions.

### Interpretations

Column density

The authors make a rough estimate of the column density of the gas in these broad velocity features by making the following assumptions:

• the $^{12}$CO emission observed is optically thin
• the excitation temperature is 15 K
• the ratio of CO to H2 is a constant $5 \times 10^{-5}$

With these assumptions, the authors find a column density of $10^{20} cm^{-2}$. This is much lower than the region’s extinction measurement of $A_{v} =$ 20 magnitudes by Snell (1979), as the outflow is sweeping out material around the star(s).

Stellar wind and bow shocks

The model of the wind that Snell et al. (1980) suggest is a bimodal wind that sweeps out material in two “bubble-like” lobes, creating a dense shell and possible ionization front that shocks the gas (More on shocks). The physical proximity of the Herbig-Haro (HH) objects in the southwest lobe coming from IRS-5 suggests a causal relationship. Previous work found that the optical spectra of the HH objects resemble spectra expected of shocked gas (Dopita 1978; Raymond 1979).

There is evidence that the CO lobes are the result of a strong stellar wind, the authors clarify this with the schematic shown in Figure 6. They suggest that the wind is creating a bow shock and a shell of swept-up material (More on outflows in star-forming regions). The broad velocity features on the CO emission line wings reach up to 15 km/s, suggesting the shell is moving out at that speed. The Herbig-Haro objects HH29 and HH102 have radial velocities of approximately 50 km/s in the same direction as the SW lobe is expanding (Strom, Grasdalen and Strom 1974). Additionally, Cudworth and Herbig (1979) measured the transverse velocities of HH28 and HH29, and found that the objects were moving at a speed of 150 to 170 km/s away from IRS-5. To have reached these velocities, the HH objects must have been accelerated, most likely by a strong stellar wind at speeds above 200 km/s. The bimodal outflow suggests a thick accretion disk around the young star.

Figure 6: Schematic drawing of stellar outflow (Figure 5 of Snell et al. 1980)

Mass-loss rate

The average density in the region away from the central core is $10^{3} {\rm ~cm}^{-3}$ (Snell 1979), so the extent and density of the shell implies a swept-up mass of 0.3 to 0.7 solar masses. With the measured velocity of ~15 km/s assumed to be constant during the lifetime of the shell and at the measured distance of 0.5 pc from the star, the shell was created 30,000 years ago. With this age, the authors determined a mass loss using the lower end of the assumed swept-up mass and the observed volume of the shell. They found a mass-loss rate of $8 \times 10^{-7} {\rm ~M}_{Sun}{\rm ~yr}^{-1}$, which can be compared to other stars using a chart like that shown in Figure 7. This is not meant to present constraints on the processes that produce the mass loss in the IRS-5 system, but rather to simply provide context for the stellar wind observed. The low-mass main sequence star(s) that will eventually arise from the IRS-5 system will be characterized by much lower mass loss rates, and studies of mass loss rates from other YSOs suggest that this source is at the high end of the range of expected rates.

Figure 7: A representative plot of the different types of stellar wind, presented by Steve Cranmer in lectures for Ay201a, Fall 2012

Snell et al. (1980) suggest observational tests of this wind-driven shock model:

• H2 emission from directly behind the shock
• FIR emission from dust swept up in the shell; this is a possibly significant source of cooling
• radio emission from the ionized gas in the wind itself near to IRS-5 with the VLA or similar; an upper limit of 21 mJy at 6 cm for this region was determined by Gilmore (1978), which suggests the wind is completely ionized

The results of some newer observations that support this wind model are presented in the following section.

## Paper III: A new look at IRS-5 by Wu et al. 2009

Wu et al. (2009) focus on the outflows in L1551 IRS5, the same region studied by Snell et al. (1980), but at a higher angular resolution (~3 arcsec; <1000 AU) and much smaller field of view (~1 arcmin; ~0.05 pc or 10,000 AU). Using the sub-millimeter array (SMA) right after it was formally dedicated, the authors detected CO(2-1) line and millimeter continuum in this low-mass star formation system. The mm continuum, which comes mostly from thermal dust emission, is used to estimate the dust mass. The CO(2-1) spectral line is used to trace the outflows around the binary or triple protostellar system at the center, revealing complex kinematics that suggest the presence of three possible bipolar outflows. The authors construct a cone-shaped outflow cavity model to explain the X-shaped component, and a precessing outflow or a winding outflow due to orbital motion model to explain the S-shaped component. The third component, the compact central one, is interpreted as newly entrained material by high-velocity jets.

### Important concepts

There are several concepts related to radio interferometry that merit some discussion:

1. Extended emission filtered out by the interferometer
This is the well-known ‘missing flux problem’ unique in interferometry. There is a maximum scale over which structure cannot be detected by an interferometer, and this scale is set by the minimum projected baseline (i.e. projected distance between a pair of antennas) in the array. In the channel maps (Fig. 3 of Wu et al. 2009), there is a big gap between 5.8 km/s and 7.1 km/s. It does not indicate that there is no CO gas at these velocities, but rather is the result of very extended and homogenous CO distribution which is unfortunately filtered out. This effect applies to all channels.

2. Visibility vs. image
The data directly obtained from the interferometers are called visibility data, in which the amplitude and phase are stored for each baseline. The amplitude, as it literally means, measures the flux; and the phase derives the relative location with respect to the phase center (a reference position on the antenna). We need to convolve the visibility data with a point spreading function, also called ‘dirty beam’, to get the image we need. Mathematically, visibility and image are related through a Fourier transform. For more information, see this online course.

3. Channel maps and P-V diagram
In radio observations, velocity of spectral lines plays an important role by providing kinematic information inside ISM. Normally, the radio data has (at least) three dimensions, two in spatial (e.g. R.A. and Decl.) and one in frequency or velocity. Velocity itself can be used to identity outflows, turbulence or infall by the analysis of line profile, or it can be combined with spatial distribution of emission, if the spatial resolution allowed as in this paper, in the form of channel maps or P-V diagram, to show the three-dimensional structure. In terms of outflows, we expect to see gas at velocities much different from the systematic velocity, and a symmetric pattern both in red- and blue-shifted sides will be even more persuasive. For the efforts in visualizing the 3-d datacube in a fancier way, see this astrobite.

### The classical image of low-mass star formation

Fig. 8 shows a diagram of a simplified four-step model of star formation (Fig. 7 of Shu et al. 1987). First, the dense gas collapses to form a core; second, a disk forms because of conservation of angular momentum; third, a pair of outflows emerge along the rotational axis; finally, a stellar system comes into place. During this process, bipolar outflows are naturally formed when the wind breaks through surrounding gas. Therefore, bipolar outflows are useful tools to indirectly probe the properties of protostellar systems.

Figure 8: The formation of a low-mass star (see Shu et al. 1987)

### Protostar candidates in L1551 IRS5

Two protostellar components as well as their own ionized jets and circumstellar disks have been found in this source. In addition, Lim & Takakuwa (2006) found a third protostellar candidate, as seen in Fig. 9. In this paper the authors investigated the possible connection between these protostars and the outflows.

Figure 9: Two 7mm continuum peaks in the north and south represent the binary system in IRS 5. Arrows show the direction of jets from each of the protostars. A third protostellar candidate is found close to the northern protostar, marked by a white cross. (see Lim and Takakuwa 2006)

### Three outflow components

Based on CO(2-1) emission, the authors found three distinct structures. Here the identification was not only based on morphology, but also on the velocity (see Fig. 5 and Fig. 9 of Wu et al. 2009). In other words, it is based on information in the 3-d datacube, as shown in 3 dimensions by the visualization below.

Figure 10: 3-D Datacube of the observations. Arrows mark the outflows identified in this paper. Red/blue colors indicate the red/blue-shifted components. The solid arrows mark the X-shaped component; the broad arrows are the S-shaped component; the narrow arrows are the compact central component. There is an axes indicator at the lower-left corner (x – R.A., y – Decl., z- velocity). (visualization by Walker Lu)

• The X-shaped component

The first one is an X-shaped structure, with its morphology and velocity shown in the paper. Four arms comprise an hour-glass like structure, with an opening angle of ~90 degree. The northwest and southwest arms are blue-shifted with respect to the systematic velocity, and the northeast and southeast arms are red-shifted. This velocity trend is the same with the large-scale bipolar outflow (Snell et al. 1980, Fig. 6 above). However the two blue-shifted arms, i.e. the NW and SW arms, are far from perfect: the SW arm is barely seen, while the NW arm consists of two components, and presents a different velocity pattern. This component coincides well the U-shaped infrared emission found to the SW of IRS5 (see Hayashi & Pyo 2009, or Fig. 4 above).

Figure 11: Components of the Outflow. Coordinates are offset from the pointing center. (Figure 7 of Wu et al. 2009)

• The S-shaped component

The second component is an S-shaped structure. It extends along the symmetry axis of the X-shaped component as well as the large scale outflow, but in an opposite orientation. As it literally means, this component is twisted like an ‘S’, although the western arm is not so prominent.

• The compact central component

The third component is a compact, high-velocity outflow very close to the protostars. The authors fitted a 2-d Gaussian structure to this component and made a integrated blue/red-shifted intensity map, which shows a likely outflow feature in the same projected orientation with the X-shaped component and the large scale outflow.

### Modeling the outflows

• A Cone-shaped cavity model for the X-shaped component

The authors then move on to construct outflow models for these components. For the X-shaped component, a cone-shaped outflow cavity model is proposed (see Fig. 12 and compare with Fig. 11). By carefully selecting the opening angle of the cone and position angle of the axis, plus assuming a Hubble-like radial expansion, this model can reproduce the X-shaped morphology and the velocity pattern. The origin of this cone is related to a high-velocity and well collimated wind, followed by a low-velocity and wide-angle wind that excavates the cone-shaped cavity. Therefore, what we see as X-shaped structure is actually the inner walls of the cavity. However, this model cannot incorporate the NW arm into the picture.

Figure 12: Models from the paper (Figure 10 of Wu et al. 2009)

• Entrained material model of the compact central component

For the compact central component, the authors argue that it is material that has been entrained recently by the jets. After comparing the momentum of this component and that of the optical jets, they found that the jets are able to drive this component. Moreover, the X-ray emission around this component indicates strong shocks, which could be produced by the high-velocity jets as well. Finally, the possibility of infalling motion instead of outflow is excluded, because the velocity gradient is larger than the typical value and rotation should dominate over infall at this scale.

• Three models of the X-shaped component

For the S-shaped component, the authors present the boldest idea in this paper. They present three possible explanations at one go, then analyze their pros and cons in four pages. But before we proceed, do we need to consider other possibilities, such as that this S-shaped component is not interconnected at all, but instead contains two separate outflows from different objects, or that the western arm of the component is not actually part of the outflow but some background or foreground emission? Although the model can reproduce the velocity pattern, it has so many adjustable parameters that we could use it to reproduce anything we like (which reminds me of the story of ‘drawing an elephant with four parameters‘). Anyway, let’s consider the three explanations individually.

1. Near and far sides of outflow cavity walls? This possibility is excluded because it cannot be incorporated into the cone-shaped cavity model aforementioned, and cannot explain the S-shaped morphology.
2. A precessing outflow? The outflow should be driven by a jet. Then if the jet is precessing, the outflow will also be twisted. The authors considered two scenarios: a straight jet and a bent jet, and found with finely tuned precessing angle and period, the bent jet model can best reproduce the velocity pattern along the symmetry axis. Therefore, the orbital motion between the third protostellar component, which is thought to be the driving source of this jet, and northern protostellar component, is proposed to cause the precession of the jet thus the outflow. A schematic image is shown in Fig 12.
3. A winding outflow due to orbital motion? The difference between this explanation and the previous one is that in a precessing outflow the driving source itself is swung by its companion protostar, so the outflow is point symmetric, while in a winding outflow the driving source is unaffected but the outflow is dragged by the gravity of its companion as it proceeds, so it has mirror symmetry with respect to the circumstellar disk plain. Again, if we fine-tune the parameters in this model, we can reproduce the velocity pattern.

Figure 13: The best-fit model for the S-shaped component, a bent precessing jet. Note the velocity patterns for red/blue lobes are not symmetric (Figure 11 of Wu et al. 2009)

A problem here, however, is although either the precessing jet or the winding outflow model is assumed to be symmetric, the authors use asymmetric velocity patterns to fit the two arms of the S-shaped component (see Fig. 12 and 13 in the paper). In the winding outflow model for instance, in order to best fit the observed velocities, the authors fit the eastern arm starting at a net velocity of 2 km/s at the center, while they fit the western arm starting at ~1.2 km/s. This means the two arms start at different velocities at the center.

### Discussion

The nature of X-shaped and S-shaped structures interpreted in this paper is based on the analysis of kinematics and comparison with toy models. However, the robustness of their conclusion suffers from several questions: for example, how to explain the uniqueness of the NW arm in the X-shaped structure? Is the X-shaped structure really a bipolar outflow system, or just two crossing outflows? Why is the compact central component filtered out around the systematic velocity? Is the S-shaped structure really a twisted outflow, or it is two outflow lobes from two separated protostars?

All these questions might be caused by the missing flux problem discussed above. Observations from a single-dish telescope could be  combined with the interferometric data to: 1) find the front and back walls of the outflow cavity, given sufficient sensitivity, to confirm that the X-shaped component is interconnected; 2) detect the extended structure around the systematic velocity, thus verify the nature of the compact central component; 3) recover at least part of the flux in the SW arm of the X-shaped component and the west arm of the S-shaped component, and better constrain the models.

## Conclusions

Using radio and infrared observations, these three papers together provide a integrated view of jets and outflows around YSOs in L1551. The near infrared observations of Hayashi & Pyo (2009) searched for [Fe II] and H2 features introduced by shocks, and found quite different configurations among the YSOs in this region. Some have complicated IR emission, such as HL/XZ Tau, while others like L1551 NE and IRS5 have well-collimated jets traced by [Fe II]. Among them, L1551 IRS5 is particularly interesting because it shows two parallel jets. The pilot work of Snell et al. (1980) revealed a bipolar molecular outflow traced by 12CO from IRS 5, which is interpreted to be created by a strong stellar wind from the young star. High angular resolution observation by Wu et al. (2009) confirms this outflow component, as well as the presence of another two bipolar outflows originating from the binary, or triple system in IRS 5. All these observations show us that jets and outflows are essential in star formation, no only by transporting the angular momentum so that YSOs can continue accreting, but also by stirring up ambient gas and feeding turbulence into the ISM, which might determine the core mass function as mentioned in Alves et al. 2007.

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## ARTICLE: Dark Nebulae, Globules, and Protostars

In Journal Club, Journal Club 2013 on February 19, 2013 at 10:45 pm

### Dark Nebulae, Globules, and Protostars by Bart Bok (1977)

Summary by George Miller

### Introduction

In Bart Bok’s 1977 paper Dark nebulae, globules, and protostars (Bok 1977), largely based on a lecture given upon acceptance of the Astroomical Society of the Pacific’s Bruce Medal, he presents two fundamentally different pictures of star formation. The first, constituting the majority of the paper’s discussion, occurs in large Bok globules which are compact, rounded and remarkably well-defined regions of high-extinction ranging from 3′ to 20′.  The globules show a strong molecular hydrogen and dust component and relatively little signs of higher neutral HI concentrations than its surroundings. In contrast, Bok briefly examines star formation in the Magellanic Clouds which show a vast amount of neutral atomic hydrogen and a comparatively small amount of cosmic dust. In this review, I will summarize a number of key points made by Bok, as well as provide additional information and modern developments since the paper’s original publishing.

### Large Bok Globules

#### A history of observations

In 1908, Barnard drew attention to “a number of very black, small, sharply defined spots or holes” in observations of the emission nebula Messier 8 (Barnard 1908).  39 years later Bok published extensive observations of 16 “globules” present in M8 as well others in $\eta$ Carinae, Sagittarius, Ophiuchus and elsewhere, making initial estimates of their distance, diameter and extinction (Bok & Reilly 1947). He further claimed that these newly coined “globules” were gravitationally contracting clouds present just prior to star formation, comparing them to an “insect’s cocoon” (Bok 1948). As we will see, this bold prediction was confirmed over 40 years later to be correct. Today there over 250 globules known within roughly 500 pc of our sun and, as Bok claims in his 1977 paper, identifying more distant sources is difficult due to their small angular diameter and large number of foreground stars.  There are currently four chief methods of measuring the column density within Bok Globules: extinction mappings of background stars, mm/sub-mm dust continuum emission, absorption measurements of galactic mid-IR background emission, and mapping molecular tracers.  See Figure 1 for a depiction of the first three of these methods.  At the time Bok published his paper in 1977, only extinction mapping and molecular tracer methods were readily available, thus I will primarily discuss these two.  For a more in depth discussion, see Goodman et. al. 2009 and the subsequent AST201b Journal Club review.

Figure 1.  Three methods of determining column density of starless molecular cores or Bok globules. (a) K-band image of Barnard 68 and plot of the $A_K$ as a function of radius from the core.  This method measures the H–K excess, uses the extinction law to convert into $A_V$, and then correlated to the $H_2$ column density from UV line measurements, parameterized by f. (b) 1.2-mm dust continuum emission map and ﬂux versus radius for L1544.  $\kappa_{\nu}$ is the dust opacity per unit gas mass, ρ is the dust density, and m the hydrogen mass (corrected for He). (c) 7-μm ISOCAM image and opacity versus radius for ρ Oph D.  In this method the absorbing opacity is related to the hydrogen column via the dust absorption cross section, $\sigma_{\lambda}$.  Figure taken from Bergin & Tafalla 2007.

#### Measuring photometric extinction

Measuring the photometric absorption, and thus yielding a minimum dust mass, for these globules is itself an arduous process. For globules with $A_v<10$   mag, optical observations with large telescopes can be used to penetrate through the globules and observe the background stars.  Here $A_{\lambda} \equiv m_{\lambda}-m_{\lambda, 0} = 2.5 \, log(\frac{F_{\lambda,0}}{F_{\lambda}})$.  Thus an extinction value of $A_v=10$ mag means the flux is decreased by a factor of $10^4$.  By using proper statistics of the typical star counts and magnitudes seen within a nearby unobstructed field of view, extinction measurements can be made for various regions.  It is important to note that the smaller an area one tries to measure an extinction of, the greater the statistical error (due to a smaller number of background stars).  This is one of the key limitations of extinction mappings.  For the denser cores or more opaque globules with $10 < A_V < 20$ mag, observations in the near infrared are needed (which is relatively simple by today’s standards but not so during Bok’s time). This is further complicated due to imprecisely defined BVRI photometric standard sequences for fainter stars, a problem still present today with various highly-sensitive space telescopes such as the HST. Bok mentions two methods. In the past a Racine (or Pickering) prism was used to produce fainter companion images of known standards, yet as discussed by Christian & Racine 1983 this method can produce important systematic errors. The second, and more widely used, method is to pick an easily accessible progression of faint stars and calibrate all subsequent photographic plates (or ccd images) from this. See Saha et. al. 2005 for a discussion of this problem in regards to the Hubble Space Telescope.

Obtaining an accurate photometric extinction for various regions within the globule is valuable as it leads an estimate of the dust density. Bok reports here from his previous Nature paper (Bok et. al. 1977) that the extinction $A_v$ within the Coalsack Globule 2 varies inversely as the square of distance, thus also implying the dust density varies inversely as the cube of distance from the core.  Modern extinction mappings, as seen in Figure 1(a) of Barnard 68,  show that at as one approaches the central core the extinction vs. distance relation actually flattens out nearly to $r^{-1}$.  This result was a key discovery, for the Bonnor-Ebert (BE) isothermal sphere model predicts a softer power law at small radii.  In his paper, Bok remarks “The smooth density gradient seems to show that Globule 2 is […] an object that reminds one of the polytropic models of stars studied at the turn of the century by Lane (1870) and Emden (1907)”.  It is truly incredible how accurate this assessment was.  The Bonnor-Ebert sphere is a model derived from the Lane-Emden equation for an isothermal, self-gravitating sphere which remains in hydrostatic equilibrium.  Figure 2 displays a modern extinction mapping of Barnard 68 along with the corresponding BE sphere model, showing that the two agree remarkably well.  There are, however, a number of detractors from the BE model applied to Bok globules.  The most obvious is that globules are rarely spherical, implying that some other non-symmetric pressure must be present.  Furthermore, the density gradient between a globule’s core and outer regions often exceeds 14 ($\xi_{max} > 6.5$) as required for a stable BE sphere (Alves, Lada & Lada 2001).

Figure 2.  Dust extinction mapping for Barnard 68 plotted against an isothermal Bonnor-Ebert sphere model.  Figure taken from Alves, Lada & Lada 2001.

#### Using CO as a tracer

Important tracer molecules, such as CO, are used to study the abundance of $H_2$, temperatures and kinematics of these globules. Because the more common $^{12}CO$ isotope tends to become optically thick and saturate in regions of higher column density such as globules, the strength of $^{13}CO$ emission is usually used to indicate the density of H2.  The conversion factor of $N_{H_2} = 5.0 \pm 2.5 \times 10^5 \times N_{13},$ from Dickman 1978 has changed little in over three decades. The column density of $H_2$, combined with its known mass and radius of the globule, can then be used to estimate the globule’s total mass. Furthermore, the correlation of $^{13}CO$ density with photometric extinction, $A_v = 3.7 \times 10^{-16} \times N_{13},$ is another strong indication that $^{13}CO$ emission is an accurate tracer for H$_2$ and dust. Further studies using $C^{17}O$ and $C^{18}O$ have also been used to trace even higher densities when even $^{13}CO$ can become optically thick(Frerking et. al. 1982).  As an example, Figure 3 shows molecular lines from the central region of the high-mass star forming region G24.78+0.08.  In the upper panel we can see the difference between the optically thick $^{12}CO$ and thin $C^{18}O$.  The $^{12}CO$ line shows obvious self-absorption peaks associated with an optically thick regime, and one clearly can not make a Gaussian fit to determine the line intensity.   $^{12}CO$, due to the small dipole moment of its $J=1 \rightarrow 0$ transition and thus ability to thermalize at relatively low densities, is also used to measure the gas temperature within globules. These temperatures usually range from 7K to 15K. Finally, the width of CO lines are used to measure the velocity dispersion within the globule. As Bok states, most velocities range from 0.8 to 1.2 km/s. This motion is often complex and measured excess line-widths beyond their thermal values are usually attributed to turbulence (Bergin & Tafalla 2007). Importantly, the line-width vs. size relationship within molecular clouds first discovered by Barnard 1981 does not extend to their denser cores (which have similar velocity motions as Bok globules).  Instead, a “coherence” radius is seen where the non-thermal component of a linewidth is approximately constant (Goodman et. al. 1998).  In the end, as Bok surmises, the subsonic nature of this turbulence implies it plays a small role compared to thermal motions.

Figure 3.  Spectra taken from the core of the high-mass star forming region G24.78+0.08.  The solid line corresponds to $^{12}CO (1\rightarrow 0)$, $^{12}CO (2\rightarrow 1)$, and $C^{32}S (3\rightarrow 2)$, the dashed line to $^{13}CO (1\rightarrow 0)$$^{13}CO (2\rightarrow 1)$, and $C^{34}S (3\rightarrow 2)$ and the dotted line to $C^{18}O (1\rightarrow 0)$.  From the top panel, one can clearly see the difference between the optically thick, saturated $^{12}CO (1\rightarrow 0)$ line and the optically thin $C^{18}O (1\rightarrow 0)$ transition.  Figure taken from Cesaroni et. al. 2003.

#### The current status of Bok globules

Today, the majority of stars are thought to originate within giant molecular clouds or larger dark cloud complexes, with only a few percent coming from Bok globules. However, the relative simplicity of these globules still make them important objects for studying star formation. While an intense debate rages today regarding the influence of turbulence, magnetic fields, and other factors on star formation in GMCs, these factors are far less important than simple gravitational contraction within Bok globules. The first list of candidate protostars within Bok globules, obtained by co-adding IRAS images, was published in 1990 with the apropos title “Star formation in small globules – Bart Bok was correct” (Yun & Clemens 1990).  To conduct the search, Yun & Clemens first fit a single-temperature modified blackbody model the the IRAS 60 and 100 μm images (after filtering out uncorrelated background emission) to obtain dust temperature and optical depth values.  This result was then used as a map to search for spatially correlated 12 and 25 μm point sources (see Figure 4.).  More evidence of protostar outflows (Yun & Clemens 1992), Herbig-Haro objects due to young-star jets (Reipurth et al. 1992) and the initial stages of protostar collapse (Zhou et. al. 1993) have also been detected within Bok Globules. Over 60 years after Bok’s pronouncement that these globules were “insect cocoons” encompassing the final stages of protostar formation, his hypothesis remains remarkably accurate and validated. It is truly “pleasant indeed that globules are there for all to observe!”

Figure 4.    (a) Contour map of the dust temperature $T_{60/100}$ of the Bok Globule CB60 derived from 60 and 100 μm IRAS images.  (b) 12 μm IRAS image of CB60 after subtracting background emission using median-filtering.  This source is thought to be a young stellar object or protostar located within the globule.  The other 12 μm field sources seen in (b) are thought not to be associated with the globule. Figure taken from Yun & Clemens 1990.

### Magellanic Cloud Star Formation

At the end of his paper, Bok makes a 180 degree turn and discusses the presence of young stars and blue globulars within the Magellanic Clouds. These star formation regions stand in stark contrast to the previously discussed Bok globules; they contain a rich amount of HI and comparatively small traces of dust, they are far larger and more massive, and they form large clusters of stars as opposed to more isolated systems. Much more is known of the star-formation history in the MCs since Bok published this 1977 paper. The youngest star populations in the MCs are found in giant and supergiant shell structures which form filamentary structures throughout the cloud. These shells are thought to form from supernova, ionizing radiation and stellar wind from massive stars which is then swept into the cool, ambient molecular clouds. Further gravitational, thermal and fluid instabilities fragment and coalesce these shells into denser star-forming regions and lead to shell-shell interactions (Dawson et. al. 2013). The initial onset of this new ($\sim$ 125 Myr) star formation is thought to be due to close encounters between the MCs, and is confirmed by large-scale kinematic models (Glatt et al. 2010).

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## ARTICLE: The Physical State of Interstellar Hydrogen

In Journal Club 2013 on February 12, 2013 at 9:57 pm

The Physical State of Interstellar Hydrogen by Bengt Strömgren (1939)

Summary by Anjali Tripathi

Abstract

In 1939, Bengt Strömgren published an analytic formulation for the spatial extent of ionization around early type stars.  Motivated by new H-alpha observations of sharply bound “diffuse nebulosities,” Strömgren was able to characterize these ionized regions and their thin boundaries in terms of the ionizing star’s properties and abundances of interstellar gas.  Strömgren’s work on these regions, which have come to be eponymously known as Strömgren spheres, has found longstanding use in the study of HII regions, as it provides a simple analytic approach to recover the idealized properties of such systems.

Background: Atomic Physics in Astronomy & New Observations

Danish astronomer Bengt Strömgren (1908-87) was born into a family of astronomers and educated during a period of rapid development in our understanding of the atom and modern physics.  These developments were felt strongly in Copenhagen where Strömgren studied and worked for much of his life.  At the invitation of Otto Struve, then director of Yerkes Observatory, Strömgren visited the University of Chicago from 1936 to 1938, where he encountered luminaries from across astrophysics, including Chandrasekhar and Kuiper.  With Struve and Kuiper, Strömgren worked to understand how photoionization could explain observations of a shell of gas around an F star, part of the eclipsing binary $\epsilon$ Aurigae (Kuiper, Struve and Strömgren, 1937).  This work laid out the analytic framework for a bounded region of ionized gas around a star, which provided the theoretical foundation for Strömgren’s later work on HII regions.

The observational basis for Strömgren’s 1939 paper came from new spectroscopic measurements taken by Otto Struve.  Using the new 150-Foot Nebular Spectrograph (Struve et al, 1938) perched on a slope at McDonald Observatory, pictured below, Struve and collaborators were able to resolve sharply bound extended regions “enveloped in diffuse nebulosities” in the Balmer H-alpha emission line (Struve and Elvey, 1938).  This emission line results from recombination when electrons transition from the n = 3 to n = 2 energy level of hydrogen, after the gas was initially ionized by UV radiation from O and B stars.  Comparing these observations to those of the central parts of the Orion Nebula led the authors to estimate that the number density of hydrogen with electrons in the n=3 state is $N_3 = 3 \times 10^{-21} cm^{-3}$, assuming a uniform concentration of stars and neglecting self-absorption (Struve and Elvey, 1938).  From his earlier work on $\epsilon$ Aurigae, Strömgren had an analytic framework with which to understand these observations.

Instrument used to resolve HII Regions in H-alpha (Struve et al, 1938)

Putting it together – Strömgren’s analysis

To understand the new observations quantitatively, Strömgren worked out the size of these emission nebulae by finding the extent of the ionized gas around the central star.  As in his paper with Kuiper and Struve, Strömgren considered only neutral and ionized hydrogen, assumed charge neutrality, and used the Saha equation with additional terms:

${N'' N_e \over N'} = \underbrace{{(2 \pi m_e)^{3/2} \over h^3} {2q'' \over q'} (kT)^{3/2}e^{-I/kT}}_\text{Saha} \cdot \underbrace{\sqrt{T_{el} \over T}}_\text{Temperature correction} \cdot \underbrace{R^2 \over 4 s^2}_\text{Geometrical Dilution}\cdot \underbrace{e^{-\tau_u}}_\text{Absorption}\\ N': \text{Neutral hydrogen (HI) number density}\\ N'':\text{Ionized hydrogen (HII) number density}\\ N_e:\text{Electron number density, }N_e = N''\text{ by charge neutrality}\\ x: \text{Ionization fraction}, x = N''/(N'+N'')$

Here, the multiplicative factor of $\sqrt{T_{el} \over T}$ corrects for the difference between the stellar temperature($T$) and the electron temperature($T_{el}$) at a distance $s$ away from the star.  The dilution factor ${R^2 \over 4 s^2}$, where $R$ is the stellar radius and $s$ is the distance from the star, accounts for the decrease in stellar flux with increasing distance.  The factor of $e^{-\tau_u}$, where $\tau_u$ is the optical depth, accounts for the reduction in the ionizing radiation due to absorption.  Taken together, this equation encapsulates the physics of a star whose photons ionize surrounding gas.  This ionization rate is balanced by the rate of recombination of ions and electrons to reform neutral hydrogen.  As a result, close to the star where there is abundant energetic flux, the gas is fully ionized, but further from the star, the gas is primarily neutral.  Strömgren’s formulation allowed him to calculate the location of the transition from ionized to neutral gas and to find the striking result that the transition region between the two is incredibly sharp, as plotted below.

Plot of ionization fraction vs. distance for an HII Region (Values from Table 2 of Strömgren, 1939)

Strömgren found that the gas remains almost completely ionized until a critical distance $s_0$, where the ionization fraction sharply drops and the gas becomes neutral due to absorption.  This critical distance has become known as the Strömgren radius, considered to be the radius of an idealized, spherical HII region.  The distance over which the ionization fraction drops from 1 to 0 is small (~0.01 pc), corresponding to one mean free path of an ionizing photon, compared to the Strömgren radius(~100pc).  Thus Strömgren’s analytic work provided an explanation for sharply bound ionized regions with thin transition zones separating the ionized gas from the exterior neutral gas.

Strömgren also demonstrated how the critical distance depends on the total number density $N$, the stellar effective temperature $T$, and the stellar radius $R$:

$\log{s_0} = -6.17 + {1 \over 3} \log{ \left( {2q'' \over q'} \sqrt{T_{el} \over T} \right)} - {1 \over 3} \log{a_u} - {1 \over 3} \frac{5040K}{T} I + {1 \over 2} \log{T} + {2 \over 3} \log{R} - {2 \over 3} \log{N},$

where $a_u$ is the absorption coefficient for the ionizing radiation per hydrogen atom (here assumed to be frequency independent) and $s_0$ is given in parsecs.  From this relation, we can see that for a given stellar radius and a fixed number density, $s_0 \propto T^{1/2}$, so that hotter, earlier type stars have larger ionized regions.  Plugging in numbers, Strömgren found that for a total number density of $3~cm^{-3}$, a cluster of 10 O7 stars would have a critical radius of 100-150 parsecs, in agreement with estimates made by the Struve and Elvey observations.

To estimate the hydrogen number density from the H-alpha observations, Strömgren also considered the excitation of the n=3 energy levels of hydrogen.  Weighing the relative importance of various mechanisms for excitation – free electron capture, Lyman-line absorption, Balmer-line absorption, and collisions – Strömgren found that their effects on the number densities of the excited states and electron number densities were comparable.  As a result, he estimated from Struve’s and Elvey’s $N_3$ that the number density of hydrogen is 2-3 $cm^{-3}$.

Strömgren’s analysis of ionized regions around stars and neutral hydrogen in “normal regions” matched earlier theoretical work by Eddington into the nature of the ISM (Eddington, 1937).  “With great diffidence, having not yet rid myself of the tradition that ‘atoms are physics, but molecules are chemistry’,” Eddington wrote that “presumably a considerable part” of the ISM is molecular.  As a result, Strömgren outlined how his analysis for ionization regions could be modified to consider regions of molecular hydrogen dissociating, presciently leaving room for the later discovery of an abundance of molecular hydrogen.  Instead of the ionization of atomic hydrogen, Strömgren worked with the dissociation of molecular hydrogen in this analysis.   Given that the energy required to dissociate the bond of molecular hydrogen is less than that required to ionize atomic hydrogen, Stromgren’s analysis gives a model of a star surrounded by ionized atoms, which is surrounded by a sharp, thin transition region of atomic hydrogen, around which molecular hydrogen remains.

In addition to HI and HII, Strömgren also considered the ionization of other atoms and transitions.  For example, Strömgren noted that if the helium abundance was smaller than that of hydrogen, then most all of the helium will be ionized out to the boundary of the hydrogen ionization region.  From similar calculations and considering the observations of Struve and Elvey, Strömgren was able to provide an estimate of the abundance of OII, a ratio of $10^{-2}-10^{-3}$ oxygen atoms to each hydrogen atom.

Strömgren Spheres Today

Strömgren’s idealized formulation for ionized regions around early type stars was well received initially and  has continued to influence thinking about HII regions in the decades since.  The simplicity of Strömgren’s model and its assumptions, however, have been recognized and addressed over time.  Amongst these are concerns about the assumption of a uniformly dense medium around the star.  Optical and radio observations, however, have revealed that the surrounding nebula can have clumps and voids – far from being uniformly dense (Osterbrock and Flather, 1959).  To address this, calculations of the nebula’s density can include a ‘filling factor’ term.  Studies of the Orion Nebula (M42), pictured below, have provided examples of just such clumpiness.  M42 has also been used to study another related limitation of Strömgren’s model – the assumption of a central star surrounded by spherical symmetry.

Orion Nebula, infrared image from WISE. Credit: NASA/JPL/Caltech

Consideration of the geometry of Strömgren spheres has been augmented by blister models of the 1970s whereby a star ionizes surrounding gas but the star is at the surface or edge of a giant molecular cloud (GMC), rather than at the center of it.  As a result, ionized gas breaks out of the GMC, like a popping blister, which in turn can prompt “champagne flows” of ionized gas leaching into the surrounding medium.  In a review article of Strömgren’s work, Odell (1999) states that due to observational selection effects, many HII regions observed in the optical actually are more akin to blister regions, rather than Strömgren spheres, since Strömgren spheres formed at the center or heart of a GMC may be obscured so much that they are observable only at radio wavelengths.

In spite of its simplifying assumptions, Strömgren’s work remains relevant today.   Given its abundance, hydrogen dominates the physical processes of emission nebulae and, thus, Strömgren’s idealized model provides a good first approximation for the ionization structure, even though more species are involved than just atomic hydrogen.  Today we can enhance our understanding of these HII regions using computer codes, such as CLOUDY, to calculate the ionization states of various atoms and molecules.  We can also computationally model the  hydrodynamics  of shocks radiating outwards from the star and use spectral synthesis codes to produce mock spectra.  From these models and the accumulated wealth of observations over time, we have come to accept that dense clouds of molecular gas, dominated with molecular hydrogen, are the sites of star formation.  Young O and B-type stars form out of clumps in these clouds and their ionizing radiation will develop into an emission nebula with ionized atomic hydrogen, sharply bound from the surrounding neutral cloud.  As the stars age and the shocks race onwards, the HII regions will evolve.  What remains, however, is Strömgren’s work which provides a simple analytic basis for understanding the complex physics of HII regions.

Strömgren, “The Physical State of Interstellar Hydrogen”, ApJ (1939)

Kuiper et al, “The Interpretation of $\epsilon$  Aurigae”, ApJ (1937)

Struve et al, “The 150-Foot Nebular Spectrograph of the McDonald Observatory”, ApJ (1938)

Eddington, “Interstellar Matter”, The Observatory (1937)

Osterbrock and Flather, “Electron Densities in the Orion Nebula. II”, ApJ (1959)

O’Dell, “Strömgren Spheres”, ApJ (1999)

## ARTICLE: On the Dark Markings in the Sky

In Journal Club, Journal Club 2013 on February 8, 2013 at 2:46 pm

On the Dark Markings in the Sky by Edward E. Barnard (1919)

Summary by Hope Chen

#### Abstract

By examining photographic plates of various regions on the sky, Edward E. Barnard concluded in this paper that what he called “dark markings” were in fact due to the obscuration of nearby nebulae in most cases. This result had a significant impact on the debate regarding the size and the dimension of the Milky Way and also the research of the interstellar medium, particularly work by Vesto Slipher, Heber Curtis and Robert Trumpler. The publication of  Photographic Atlas of Selected Regions of the Milky Way after Barnard’s death, which included many of the regions mentioned in the paper, further provided a new method of doing astronomy research. In this paper and the Atlas, we are also able to see a paradigm very different from that of today.

It is now well-known that the interstellar medium causes extinction of light from background stars. However, think of a time when the infrared imaging was impossible, and the word “photon” meant nothing but a suspicious idea. Back in such a time in the second decade of the twentieth century, Edward Edison Barnard, by looking at hundreds of photographic plates, proposed an insightful idea that “starless” patches of the sky were dark because they are obscured by nearby nebulae. This idea not only built the foundation of the modern concept of the interstellar medium, but also helped astronomers figure out that the Universe extended so much farther beyond the Milky Way.

#### Young Astronomer and His Obsession of the Sky

In 1919, E. E. Barnard published this paper and raised the idea that what he called “dark markings” are mostly obscuration from nebulae close to us. The journey, however, started long before the publication of this paper. Born in Nashville, Tennessee in 1857, Barnard was not able to receive much formal education owing to poverty. His first interest, which became important for his later career, was in photography. He started working as a photographer’s assistant at the age of nine, and the work continued throughout most of his teenage years. He then developed an interest in astronomy, or rather, “star-gazing,” and would go watch the sky almost every night with his own telescope. He took courses in natural sciences at Vanderbilt University and started his professional career as an astronomer at the Lick Observatory in 1888. He helped build the Bruce Photographic Telescope at the Lick Observatory and there he started taking pictures of the sky on photographic plates. He then moved on to his career at the Yerkes Observatory at Chicago University and worked there until his death in 1922. (Introduction of the Atlas, Ref. 2)

Fig. 1 One of the many plates in the Atlas including the region around Rho Ophiuchii, which was constantly mentioned in many of Barnard’s works. (Ref. 2)

Fig. 1 is one of the many plates taken at the Yerkes Observatory. It shows the region near Rho Ophiuchii, which was a region constantly and repetitively visited by Barnard and his telescope. Barnard noted in his description of this plate, “the [luminous] nebula itself is a beautiful object. With its outlying connections and the dark spot in which it is placed and the vacant lanes running to the East from it, … it gives every evidence that it obscures the stars beyond it.” Numerous similar comments spread throughout his descriptions of various regions covered in A Photographic Atlas of Selected Regions of the Milky Way (hereafter, the Atlas). Then finally in his 1919 paper, he concluded, “To me these are all conclusive evidence that masses of obscuring matter exist in space and are readily shown on photographs with the ordinary portrait lenses,” although “what the nature of this matter may be is quite another thing.” The publication of these plates in the Atlas (unfortunately after his death, put together by Miss Mary R. Calvert, who was Barnard’s assistant at the Yerkes Observatory and helped publish many of Barnard’s works after his death) also provided a new way of conducting astronomical research just as the World Wide Telescope does today. The Atlas for the first time allowed researchers to examine the image and the astronomical coordinates along with lists of dominant objects at the same time.

Except quoting Vesto Slipher’s work on spectrometry measurements of these nebulae, most of the evidences in Barnard’s paper seemed rather qualitative than quantitative. So, as of today’s standard, was the “evidence” really conclusive? Again, the question cannot be answered without knowing the limits of astronomical research at the time. Besides an immature understanding of the underlying physics, astronomers in the beginning of the twentieth century were limited by the lack of tools on both the observation and analysis fronts. Photographic plates as those in the Atlas were pretty much the most advanced imaging technique at the time, on which even a quantitative description of “brightness” was not easy, not to mention an estimation of the extinction of these “dark markings.” However, this being said, a very meaningful and somewhat “quantitative” assumption was drawn in Barnard’s paper: the field stars were more or less uniformly distributed. Barnard came to this assumption by looking at many different places, both in the galactic plane and off the plane, and observing the densities of field stars in these regions. Although numbers were not given in the paper, this was inherently similar to a star count study. Eventually, this assumption lead to what Barnard thought as the conclusive evidence of these dark markings being obscuring nebulae instead of “vacancies.” Considering the many technical limits at the time, while the paper might not seem to be scientific in today’s standard, this paper did pose a “conclusion” which was strong enough to sustain many of the more quantitative following examinations.

#### The “Great Debate”

Almost at the same time, perviously mentioned Vesto Slipher (working at the Lowell Observatory) began taking spectroscopic measurements of various clouds and tried to understand the constituents of these clouds. Although limited by the wavelength range and the knowledge of different radiative processes (the molecular transition line emission used largely in the research of the interstellar medium today was not observed until half a century later in 1970, by Robert Wilson, who, on a side note, also discovered the Cosmic Microwave Background), Slipher was able to determine the velocities of clusters by measuring the Doppler shifts and concluded that many of these clusters move at a faster rate than the escape velocity of the Milky Way (Fig. 2). This result, coupled with Barnard’s view of intervening nebulae, revolutionized the notion of the Universe in the 1920s.

Fig. 2 The velocity measurements from spectroscopic observations done by Vesto Slipher. (Ref. 3)

On April 26, 1920 (and in much of the 1920s), the “Great Debate” took place between Harlow Shapley (the Director of Harvard College Observatory at the time) and Curtis Heber (the Lick Observatory, 1902 to 1920). The general debate concerned the dimension of the Universe and the Milky Way, but the basic issue was simply whether distant “spiral nebulae” were small and lay within the Milky Way or whether they were large and independent galaxies. Besides the distance and the velocity measurements, which suffered from large uncertainties due to the technique available at the time, Curtis Heber was able to “win” the debate by claiming that dark lanes in the “Great Andromeda Nebula” resemble local dark clouds as those observed by Barnard (Fig. 3, taken in 1899). The result of the debate then sparked a large amount of work on “extragalactic astronomy” in the next two decades and was treated as the beginning of this particular research field.

Fig. 3 The photographic plate of the “Great Andromeda Nebula” taken in 1988 by Isaac Roberts.

#### The Paper Finally Has a Plot

Then after the first three decades of the twentieth century, astronomers were finally equipped with a relatively more correct view of the Universe, the idea of photons and quantum theory. In 1930, Robert J. Trumpler (the namesake of the Trumpler Award) published his paper about reddening and reconfirmed the existence of local “dark nebulae.” Fig. 4 shows the famous plot in his paper which showed discrepancies between diameter distances and photometric distances of clusters. In the same paper, Trumpler also tried to categorize effects of the ISM on light from background stars, including what he called “selective absorption” or reddening as it is known today. This paper, together with many of Trumpler’s other papers, is one of the first systematic research results in understanding the properties of Barnard’s dark nebulae, which are now known under various names such as clouds, clumps, and filaments, in the interstellar medium.

Fig. 4 Trumpler’s measurements of diameter distances v. photometric distances for various clusters.

#### Moral of the Story

As Alyssa said in class, it is often more beneficial than we thought to understand what astronomers knew and didn’t know at different periods of time and how we came to know what we see as common sense today, not only in the historically interesting sense but also in the sense of better understanding of various ideas. In this paper, Barnard demonstrated a paradigm which we may call unscientific today but made a huge leap into what later became the modern research field of the interstellar medium.