Issue |
A&A
Volume 506, Number 2, November I 2009
|
|
---|---|---|
Page(s) | 779 - 788 | |
Section | Interstellar and circumstellar matter | |
DOI | https://doi.org/10.1051/0004-6361/200912408 | |
Published online | 27 August 2009 |
A&A 506, 779-788 (2009)
Physical structure and dust reprocessing in a sample of HH jets
,![[*]](/icons/foot_motif.png)
L. Podio1 - S. Medves2 - F. Bacciotti3 - J. Eislöffel4 - T. Ray1
1 - Dublin Institute for Advanced Studies, School of Cosmic Physics,
31 Fitzwilliam Place, Dublin 2, Ireland
2 -
Università di Pisa, Dipartimento di Fisica ``Enrico Fermi'',
Largo Bruno Pontecorvo 3, 56127 Pisa, Italy
3 -
INAF - Osservatorio Astrofisico di Arcetri, Largo Enrico Fermi 5,
50125 Firenze, Italy
4 -
Thüringer Landessternwarte Tautenburg,
Sternwarte 5, 07778 Tautenburg, Germany
Received 30 April 2009 / Accepted 21 July 2009
Abstract
Context. Stellar jets are an essential ingredient of the
star formation process and a wealth of information can be derived from
their characteristic emission-line spectra.
Aims. We investigate the physical structure and dust
reprocessing in the shocks along the beam of a number of classical
Herbig-Haro (HH) jets in the Orion and Lupus molecular clouds (HH 111,
HH 1/2, HH 83, HH 24 M/A/E/C, and Sz68). Parameters
describing plasma conditions, as well as dust content, are derived as a
function of distance from the source and, for HH 111, of gas
velocity.
Methods. Spectral diagnostic techniques are applied to obtain the jet physical conditions (the electron and total density, and
,
the ionisation fraction,
,
and the temperature,
)
from the ratios of selected forbidden lines. The presence of dust
grains is investigated by estimating the gas-phase abundance of calcium
with respect to its solar value.
Results. We find the electron density varies between 0.05-
cm-3, the ionisation fraction
from 0.01-0.7, the temperature ranges between 0.6-
K, and the hydrogen density between 0.01-
cm-3. Interestingly, in the HH 111 jet,
,
,
and
peak in the high velocity interval (HVI) of the strongest working
surfaces, confirming a prediction from shocks models. Calcium turns out
to be depleted with respect to its solar value, but its gas-phase
abundance is higher than estimates for the interstellar medium in
Orion. The depletion is high (up to 80%) along the low-excited
jets, while low or no depletion is measured in those jets which show
higher excitation conditions. Moreover, for HH 111 the depletion
is lower in the HVI of the fastest shock.
Conclusions. Our results confirm the shock structure predicted
by models and indicate that shocks occurring along jets, and presumably
those present in the launch zone, only partially destroy dust grains
and that the efficiency of dust reprocessing strongly depends on shock
velocity. However, the high Ca gas-phase abundance estimated in some of
the knots, is not well understood in terms of existing models of dust
reprocessing in shocks, and indicates that the dust must have been
partially reprocessed in the region where the flow originates.
Key words: ISM: jets and outflows - ISM: Herbig-Haro objects - ISM: dust, extinction - stars: formation
1 Introduction
Although it is widely accepted that jets play an essential role in star formation, there are still many open questions about how they are generated and propagate in the interstellar medium. Their optical spectra are characterised by emission lines from atomic species which are collisionally excited in the shock waves generated by the interaction of the ejected material with the interstellar medium or previously ejected jet material. These lines contain a wealth of information on jet and shock properties. Beyond the jet/shock morphology and kinematics, information on the gas physical conditions can be derived by developing a grid of shock models the predictions of which are compared with the observed line ratios (e.g., Hartigan et al. 1994; Raga & Böhm 1986; Hartigan et al. 1987); or, alternatively, by using spectral diagnostics techniques which are independent of any assumption of the heating mechanism (e.g., Bacciotti & Eislöffel 1999, hereafter referred to as the BE technique). The application of these diagnostic techniques to various datasets allowed derivation of the jet physical/dynamical structure, highlighting the stratification in the shock cooling regions (Nisini et al. 2005; Podio et al. 2006), the variation of the excitation conditions along the jet (Bacciotti & Eislöffel 1999), across it (Bacciotti et al. 2000; Hartigan & Morse 2007) and, in a few cases, with gas velocity (Garcia Lopez et al. 2008; Lavalley-Fouquet et al. 2000; Coffey et al. 2008).
It is also important to estimate the dust content in jets,
a topic that has been poorly studied to date.
Several theoretical studies have explored the dust content and its distribution
in the interstellar medium.
On the one hand both observations and theoretical models show that
the gas-phase abundances of elements like iron (Fe), magnesium (Mg),
silicon (Si), and calcium (Ca)
are considerably depleted (by a factor of the order of 102-104)
with respect to their solar abundances because of the formation of
dust grains (Baldwin et al. 1991; Savage & Sembach 1996; van Dishoek et al. 1993).
On the other hand the chemical composition, the size and the content
of dust grains can evolve as they lose atoms to the gas phase through
shocks waves. The high energy gas-grain and grain-grain collisions
occurring behind the shock front, in fact, may lead to the erosion of the
grain surfaces (sputtering) and/or to the vaporisation and fragmentation
of grains
(Draine 2003; Jones 2000; Jones et al. 1994; Guillet et al. 2009, and references therein).
Total dust destruction is expected in high velocity shocks such as
supernova-generated shock waves (
km s-1) but there are very few
observations about dust reprocessing in slower shocks such as those
occurring along stellar jets (
-80 km s-1in the jet).
A few studies have investigated the gas phase abundance of Fe in HH jets
(Mouri & Taniguchi 2000; Nisini et al. 2002; Beck-Winchatz et al. 1994; Böhm & Matt 2001; Beck-Winchatz et al. 1996) and
only recently, the analysis has been extended to other refractory species such
as Si, Ca, carbon (C),
nickel (Ni), chromium (Cr), and titanium (Ti)
(Nisini et al. 2005; Garcia Lopez et al. 2008; Podio et al. 2006; Nisini et al. 2007).
To further investigate this topic, we study here the distribution of physical parameters and, in particular, the gas phase abundance of calcium in a sample of HH jets (HH 111, HH 1/2, HH 83, HH 24 M/A/C/E, Sz68), as a function of distance from the source, i.e. in the different working surfaces along the jet, and, where possible, of gas velocity. As we will show, these kind of estimates are very important to understand dust reprocessing in shocks and are also useful to derive constraints on the properties of the jet launching region.
This paper is organized as follows: in Sect. 2 we present the observations and the data reduction process and in Sect. 3 we briefly recall the principle of the so-called BE diagnostic technique, which is used to derive gas physical conditions along the jets. In Sect. 4 we present the physical properties derived for HH 111, in two velocity intervals, and for the jet from Sz68, that is analysed here for the first time. Similar results derived for the other jets in the sample, at higher spatial sampling than in previous studies, are presented in the online material. In Sect. 5 we use the derived physical parameters to estimate the presence of dust grains in the jets and we discuss the efficiency of shocks in reprocessing the dust. A comparison of our results with the predictions of theroretical models is then made. Finally, in Sect. 6 we present a summary of our results.
2 Observations and data reduction
We obtained optical spectra of a sample of HH jets
(HH 111, HH 1/2, HH 83, HH 24 M/A/C/E, and Sz68).
Our spectra were acquired in January 1998 at the ESO 3.6 m-telescope
by using the EFOSC2 spectrograph, the GR9 grism and a 1
5 slit width.
The slit was aligned parallel to the jet axis with position
angles: 276.7
for HH 111, 324.7
for HH 1/2,
297.7
for HH 83, 328.1
for HH 24 C/E, 219.0
for HH 24 G,
and 334.9
for Sz68.
The spectra cover the wavelength range
from 5600 to 7335
with a spectral resolution
1300
(
km s-1).
The spatial and spectral sampling
scales are 0.157
/pixel and 0.872
/pixel
corresponding to
40 km s-1.
The spectral images were flat-fielded, sky-subtracted, wavelength and
flux calibrated. In addition to the standard data reduction process the
spectral images of the detected lines ([O I]6300, 6363,
[N II]
6548, 6583, [S II]
6716, 6731, and [Ca II]
7291, 7324)
were velocity-calibrated, resampled to the same velocity scale
for all the lines, and corrected for velocity shifts due to atmospheric
differential refraction.
During the observations, in fact, the slit was aligned along the jet emission
through pre-imaging in the [S II]
6731 line.
Because of atmospheric differential refraction, however, the peak of the
emission in the other lines may be off the slit centre inducing a velocity
shift in these lines (see the Appendix of Bacciotti et al. 2002, for a discussion
of the uneven slit illumination effect).
The resampling in velocity and the corrections mentioned above
allowed us to accurately align the lines in velocity despite the low
spectral resolution.
As shown in Hartigan & Morse (2007), the interplay between the slit width and the projected width of the target is crucial in spectroscopic observations. If the slit is narrower than the jet (and the lines are broader than the spectral resolution) the detected lines are spectrally resolved and the velocity resolution is higher for decreasing slit width. If, in contrast, the slit is much wider than the jet (and the velocity dispersion is lower than the spectral resolution) the observations produce emission-line images for each line in the spectrum, but no velocity information can be recovered. The latter is so-called ``slitless spectroscopy'' (see also Hartigan et al. 2004).
Our observations are seeing-limited and the jet widths are spatially
unresolved. Thus the critical parameter to be
compared with the slit width is the
full width half maximum (FWHM) of the seeing.
This varied between 1
,
when observing HH 1/2, HH 83,
HH 24 M/A/C/E, and Sz68, and
1
5 when observing HH 111.
In the case of HH 111, the seeing was comparable to the slit width and the acquired spectra show line profiles which are broader than the spectral resolution. Thus the HH 111 jet is partially resolved in velocity, allowing an analysis of its structure along the jet and in two different velocity intervals, as will be explained in more detail in the next section.
In the other jets, in contrast, the slit is wider than the seeing FWHM. Thus we obtain images of the jet in the different lines. This is evident when comparing spectral images of the detected lines with pre-imaging acquired with the [S II] filter. To avoid confusion between the spectral and spatial information we integrated the line fluxes over the line ``spectral'' profile, suppressing the velocity dimension.
In the next section we explain in more detail the analysis performed in the two different cases.
![]() |
Figure 1:
Position-velocity diagrams of the forbidden line ratios used in
the diagnostic technique for the HH 111 jet.
From top to bottom panel:
the [S II] |
Open with DEXTER |
3 Application of spectral diagnostic
Our analysis relies upon the dependence of the ratios between
optical forbidden lines on the gas physical conditions in the jet, i.e. on
the electron and total hydrogen density,
and
,
the hydrogen ionisation fraction,
,
and the
electron temperature,
(Bacciotti & Eislöffel 1999; Bacciotti et al. 1995).
As a first step, we thus computed for all the jets in our sample the line
ratios used in the diagnostics, i.e. [S II]6731/6716,
[N II]
6548, 6583/[O I]
6300, 6363, and
[O I]
6300, 6363/[S II]
6716, 6731.
For HH 111, which is partially spectrally resolved,
we divided the spectral images of the detected
S+, O0, and N+ lines pixel by pixel, thus obtaining
position-velocity (PV) diagrams of the line ratios
(see Fig. 1).
For the other jets (HH 1/2, HH 83, HH 24 M/A/C/E, and Sz68)
we integrated over the spectral profile
obtaining the variation of the ratios only as a function of the distance
from the source.
The distribution of the ratios already gives a qualitative idea of the
variations of the gas physical conditions along the jet and, for HH 111,
with the gas velocity.
The ratio between the sulphur doublet lines ([S II]
6731/6716)
is a tracer of the gas electron density
up to the critical density of the [S II] lines
(
cm-3), the [N II]/[O I]
increases primarily with
,
and
the [O I]/[S II] ratio is dependent on both
and
.
To obtain a more quantitative estimate of the gas physical properties
we use the BE technique (Bacciotti & Eislöffel 1999), which extracts
diagnostic information from the comparison of observed and calculated ratios.
The method is based on the fact that the observed forbidden lines
are collisionally excited
and that, in low-excitation conditions, and provided that no strong
sources of ionizing photons are present, the ionization fraction of hydrogen is
tightly related to those of nitrogen and oxygen via charge exchange
processes.
This allows one to easily retrieve (and the total hydrogen density)
and
.
As discussed in Bacciotti & Eislöffel (1999) this method assumes that the
lines used in the technique are emitted in a region of similar temperature,
density and ionization fraction.
Bacciotti & Eislöffel (1999) and Nisini et al. (2005) carefully verified this
assumption by using the results from shock models (e.g., Hartigan et al. 1987).
They showed that the considered forbidden lines have their peak emission in the
same region behind shock fronts of different shock velocity, and verified that
the obtained parameters are representative averages.
In this paper we use an improved version of the BE code
in which the values of the collision strength are updated using the results
in Keenan et al. (1996) for S+, Hudson & Bell (2005) for N+,
and Berrington & Burke (1981) and Mendoza (1983) for O0.
This allows us to obtain a better fit of the collision strength for high
temperatures ( between 2-
K)
and thus to trace correctly the plasma physical conditions
in regions of high excitation, i.e.
where the [N II] emission is comparable to that in the [S II] and
[O I] lines.
Examples of the latter are some knots of
the HH 83 and the HH 24 jets where we could not obtain results in our
previous study (Podio et al. 2006).
Note that the diagnostic technique uses
ratios between different species and, thus, the obtained and
may depend on the chosen set of elemental abundances.
However, as stressed in Podio et al. (2006), the relative variations
of the parameters (i.e. in different knots and/or velocity
intervals) do not depend on the
choice of an abundance set (see Fig. 1 of Podio et al. 2006) which, therefore,
can be thought as a model parameter.
On the other hand the most recent determinations by Asplund et al. (2005) (solar)
and Esteban et al. (2004) (Orion) are in good agreement and give rise to diagnostic
results that differ by at most 15%.
One of the main goal of this paper is the determination of the
calcium gas-phase abundance with respect to solar.
Thus, in contrast with Podio et al. (2006), the solar abundances estimated by
Asplund et al. (2005) are adopted for consistency in all our diagnostic analysis.
In particular, the values of the temperature and the ionization fraction are
larger when using the abundances from Esteban et al. (2004).
The uncertainty over the absolute values is not shown in
Figs. 3,
4, A.1, A.3, and A.4, because
these are intended to be used to analyse the relative variations
of parameters along the jet and in different velocity intervals.
However, we do consider the uncertainty due to abundances choice
when computing theoretical [Ca II]/[S II] ratios to estimate the calcium
gas-phase abundance (see Sect. 5 for more details).
Since the observations are seeing-limited,
i.e. the angular resolution of the data
is much wider than the spatial sampling, we integrated the line fluxes
over the seeing FWHM (
pixels) before
applying the diagnostics.
Thus the de facto spatial sampling in all the
figures is equal to the seeing.
Moreover, we integrated the line fluxes over the spectral profile for the
spectrally unresolved jets (HH 1/2, HH 83, HH 24 M/A/C/E, Sz68) and
over two velocity intervals for HH 111:
low (for v >-100 km s-1) and high (for v <-100 km s-1)
velocity Interval (LVI and HVI, respectively).
For the HH 111 and HH 1 jets we corrected the observed fluxes for
reddening using the values of the visual extinction, ,
estimated by
Nisini et al. (2005) and Podio et al. (2006), the standard dereddening procedure from
Draine (1989) and an interpolation of the extinction law derived by
Rieke & Lebofsky (1985) for the near-IR bands.
For the HH 83 and HH 24 jets we have no estimates of
the visual extinction because the [Fe II]
1.64, 1.32, 1.25 lines were not
detected in our earlier optical/NIR spectra (Podio et al. 2006).
Thus we could not correct the line fluxes for the reddening before
computing the ratios and applying the diagnostics technique.
On the other hand, Bacciotti & Eislöffel (1999) showed that, since the
lines used in this diagnostic are very close in wavelength, any errors induced
by not correcting for extinction are at most 8-10%
for the ionization fraction and 15% for the temperature
as long as the visual extinction is lower than
3.
Once a set of abundances is assumed, the errors affecting the parameters
obtained using the BE technique are due to the signal-to-noise ratio of the
measured fluxes and the uncertainty in the determination of
the .
These are very low (<5%) in the bright knots of the HH 111, HH 1, and
HH 83 jets which have strong emission in all the lines, and larger
(up to 50%) in the fainter knots of the HH 24 and Sz68 jets.
4 The physical structure of the jets
In this section we briefly describe the new information gathered by applying spectral diagnostics (i.e. the BE technique) to our 3.6 m/EFOSC2 spectra of a sample of HH jets.
In Sect. 4.1 we present the physical parameters (,
,
,
)
for HH 111,
for which we have been able to highlight for the first time the
stratification of the gas physical conditions as a function of
velocity.
In Sect. 4.2 we present the parameters obtained for the HH jet from Sz68 which has not been analysed before with the BE method.
The other objects in our sample, already analysed previously (Bacciotti & Eislöffel 1999; Nisini et al. 2005; Podio et al. 2006), are here re-analysed with the improved version of the code and increased spatial sampling, for an accurate determination of Ca gas-phase abundance along the jets (see Sect. 5). The physical properties of these jets are shown in the online material, for reference.
4.1 HH 111: the high and low velocity intervals
![]() |
Figure 2:
Variation of the forbidden line ratios used in
the diagnostic technique for the HH 111 jet as a
function of distance from the source and in two velocity intervals:
the high velocity interval (HVI, red points/solid line, v <-100 km s-1) and
the low velocity interval (LVI, black points/dotted line, v >-100 km s-1).
In contrast, in the terminal bow (knot V) the ratios have been computed
by integrating the fluxes over the line spectral profile thus there is no
information on the velocity intervals.
From top to bottom panel:
intensity profiles of the optical lines, [S II] |
Open with DEXTER |
![]() |
Figure 3:
Variation of the physical parameters for the HH 111 jet as a
function of distance from the source and in two velocity intervals:
the high velocity interval (HVI, red points/solid line, v <-100 km s-1) and
the low velocity interval (LVI, black points/dotted line, v >-100 km s-1).
In contrast, in the terminal bow (knot V) the parameters are derived
by integrating the fluxes over the line spectral profile thus there is no
information on the velocity intervals.
From top to bottom panel:
intensity profiles of the optical lines, the electron density, |
Open with DEXTER |
HH 111 is a well-known parsec-scale jet powered by the young star
IRAS 05491+0247 located in the L1617 cloud in Orion (D=460 pc).
The blueshifted lobe is detected in the optical only from a distance of
20
from the source, because the first part is almost
completely obscured by a dusty molecular envelope.
Then it emerges from the cloud showing
a chain of bright, equally spaced and well-collimated
knots up to a distance of
48
(knots E-L,
following the nomenclature from Reipurth et al. 1997).
These are followed by a series of more widely
spaced and fainter knots up to 80
(M, N, O, P), and then the jet terminates with a strong bow-shock located at
150
(knot V).
The physical properties of HH 111 have been investigated by several authors
using both optical (e.g., Hartigan et al. 1994; Morse et al. 1993a) and near-infrared
(e.g., Nisini et al. 2002) lines.
Reipurth et al. (1997) and Davis et al. (2001), instead,
focused on the gas kinematics,
through analysing [S II] and H2 line profiles, and
highlighted the presence of two velocity components peaking at
-75/-85 and
-15/-30 km s-1.
In a previous work (Podio et al. 2006) a detailed analysis of HH 111
was performed using several spectral tracers, demonstrating
the density and temperature stratification in each spatially
unresolved cooling region.
Nevertheless, none of the previous studies investigated the
variation of the physical parameters as a function of gas velocity, i.e.
in different velocity intervals.
As explained in Sect. 2
the data for HH 111 show partially velocity resolved line profiles.
This is confirmed by the position-velocity (PV) diagrams
of the computed line ratios in Fig. 1, which clearly
show different value of the ratios in the two velocity intervals
in the jet (knots E to L).
A comparison of the values of the ratios integrated over the two considered
velocity intervals (LVI and HVI) is also shown in
Fig. 2.
Since these ratios depend on the gas physical conditions,
they give a qualitative idea of the variations
of electron density, temperature, and ionisation fraction along the jet
and with the gas velocity.
The PV diagrams
of the ratio between the sulphur doublet lines ([S II]6731/6716),
for example, shows that
peaks at the positions of the brightest
knots and at high velocities.
The [N II]/[O I] and [O I]/[S II] ratios, on the other hand,
indicate a peak of the excitation conditions (
,
)
in the high velocity interval of knot L.
Interestingly, HST high-angular resolution images of the jet
clearly show that HH 111 is a chain of shock working surfaces
where knot L has the clearest ``bow'' morphology (Reipurth et al. 1997) and
is the knot with the largest [S II] and H2 line profiles
(Reipurth et al. 1997; Davis et al. 2001) and the largest shock velocity (Hartigan et al. 2001)
in the jet beam.
Our PV diagrams confirm this picture highlighting
the shock velocity structure.
By integrating the line fluxes over the seeing FWHM and over the LVI and HVI
and then by applying the BE technique we obtain
the variations of the physical parameters both along the jet
and in the two velocity intervals for knots E-L (see Fig. 3).
The derived parameters are in agreement with previous estimates
(Podio et al. 2006):
varies between 102 and
cm-3,
goes
from a few percent to approximately 30%,
-
K,
and
ranges between
and
cm-3.
With respect to previous analyses, however,
the increased spatial sampling (more than twice
that of previous studies) and the
velocity information highlight several shocks features,
such as
maxima
in the HVI of the brightest knots, higher values of
in the HVI
along the full jet length, and, finally,
a strong and steep increase of
and
in the HVI of knot L.
Note that very weak emission in the [O III]
4959, 5007 lines was
detected by Morse et al. (1993a) in this knot. The presence of O++ is not accounted for in the adopted diagnostic technique.
Nevertheless, since the emission in these lines is very weak, we
tentatively applied the diagnostic to knot L and
the obtained high values of
and
are
in agreement with the detection of such high-excitation lines.
These results are an important observational validation of the shock
velocity structure predicted by models (Hartigan et al. 1987).
We do not apply the diagnostic to knots M, N, O, and P because of the
very weak emission in the [N II] lines.
Nevertheless, the electron density can be inferred by integrating the
[S II] emission over the knot spatial profile to increase the S/N,
which is low at the position of these faint knots.
We obtain values of about 150 cm-3, 200 cm-3, 250 cm-3, and 150 cm-3
in the HVI of knot M, N, O, and P and values of
100 cm-3 lower in the
LVI for the same knots.
The average values are in agreement with previous results from Podio et al. (2006).
Previous studies did not detect high excitation lines at the position
of the terminal bow shock (knot V) (Morse et al. 1993b),
thus we apply our diagnostic to this knot.
Even if two velocity components were detected in [S II] by
Reipurth et al. (1997) the PVs in Fig. 1
do not show different values of the ratios in the two velocity intervals.
This can be due to the low spectral resolution combined with the much
lower intensity of knot V with respect to other knots in the jet (knot E-L).
Thus we integrate the line fluxes over their
spectral profiles obtaining the trend shown in Fig. 3.
The electron and total density are 400-550 cm-3 and
2-
cm-3, respectively.
The ionisation fraction and the temperature show an unexpected
anti-correlation and varies between
0.23 and
0.07,
and
K and
K moving from the shock front
towards the source.
Note that despite the high shock velocity of knot V (Hartigan et al. 2001) we did
not infer high values of
and
as for knot L.
4.2 The Sz68 jet
![]() |
Figure 4:
Variation of the physical parameters for the Sz68 jet as a
function of the distance from the source.
From top to bottom panel:
intensity profiles of the optical lines, the electron density, |
Open with DEXTER |
The Sz68 jet has never been analysed using spectral diagnostic techniques.
This jet is located in the B228 molecular cloud between the two
emission-line stars Sz68 and Sz69.
The collimated chain of [S II] shock-excited blueshifted
emission, extending up to 34
from Sz68 at a position angle of 135
,
was detected by Heyer & Graham (1989).
They report three knots located at 21
,
28
,
and 34
from the source, the brightest ones being the first knot, A,
and the last one, B.
Comparing the position of the knots in the spectra of Heyer & Graham (1989) with
the positions in our EFOSC2 data (see upper panel of Fig. 4 for
the knots spatial profile), we derive an estimate of the knot proper motions.
This indicates a tangential velocity of
600 km s-1 for knot A,
460 km s-1
for the intermediate knot, and
330 km s-1 for knot B,
which combined with the radial velocity
corrected for the source heliocentric velocity
(
-55 km s-1 for knot A and
-29 km s-1 for knot B,
Heyer & Graham 1989)
gives an upper limit for the inclination angle out of the plane of the sky
of
5
.
We applied the BE technique to the line fluxes measured along the three knots
by integrating over the line spectral profiles.
Since this object is very faint, the error-bars are much larger than for the
other jets.
The inferred electron density is quite low varying between 200 cm-3 and
cm-3 . These values are slightly lower than those
derived by Heyer & Graham (1989).
The ionisation fraction is 0.05-0.25 and the temperature ranges between
0.6-
K, with a maximum at the position of the intermediate knot,
where there is a peak of
,
as well.
Finally the hydrogen total density is 2-
cm-3.
The determination of the jet physical conditions, as well as
the jet tangential velocity
and inclination angle allows us to estimate the mass loss rate as
rJ 2 vJ,
where
is the mean atomic weight,
the proton mass,
the hydrogen density and rJ and vJ, respectively, the jet radius and
velocity.
By assuming a typical jet radius of
100-300 AU we obtain a mass loss rate
between a few 10-8 and a few 10-7
yr-1.
These values are of the same order of magnitude of
found for typical
HH jets (Bacciotti & Eislöffel 1999; Hartigan et al. 1994; Podio et al. 2006).
5 Dust reprocessing in the jets
Refractory species, such as Ca, Fe, Ni, Cr, Ti, Mg, Si
are often depleted in the interstellar
medium because their atoms are locked into dust grains (Savage & Sembach 1996).
The calcium gas-phase abundance, for example,
is estimated to be
in Orion (Baldwin et al. 1991),
i.e. two orders of magnitude lower
than its solar abundance (
,
Asplund et al. 2005),
and up to 4 orders of magnitude lower than solar in
Oph
(Savage & Sembach 1996).
On the other hand, the passage of shocks can partially or completely
destroy the dust releasing the refractory atoms into
the gas phase.
Theoretical models of dust reprocessing in shocks, such as those driven by supernova explosions or those occurring in stellar jets, estimate the dust destruction rate by considering the different processes at work, i.e. thermal and inertial sputtering, photoevaporation and shattering. They show that the relative importance of these processes, and thus the efficiency of the shock in reprocessing dust, strongly depends on a number of parameters such as the shock velocity, the gas pre-shock density, the intensity of the magnetic field, and the size and structure of the dust grains (e.g., Draine 2003; Jones 2000; Jones et al. 1994; Guillet et al. 2009).
From an observational point of view an estimate of the depletion of refractory species with respect to their solar abundance can be used to gauge the amount of dust grains in jets from young stars and thus to evaluate the efficiency of mild shocks in destroying the dust. This in turn can help constrain the location and size of the region around the star from where the jet originates.
In Sect. 5.1 we explain in detail the method used to estimate the abundance of calcium in the gas-phase, while in Sect. 5.2 we present the results obtained and in Sect. 5.3 we compare them with the predictions of theoretical models.
5.1 Observational estimates of calcium gas-phase abundance in stellar jets: the method
To estimate the depletion of a given refractory specie, in our case calcium, with respect to its solar abundance, we compare ``expected'' and observed ratios between refractory and non-refractory species known to be present in solar abundance in the interstellar medium (e.g. Savage & Sembach 1996).
The ``expected'' ratios can be computed through two different methods:
- (i)
- by using shock models which predict the gas physical conditions and
thus the emission in different lines behind the shock front. This method was
used by Mouri & Taniguchi (2000) and, in part, by Nisini et al. (2002) to estimate
Fe gas-phase abundance in HH jets. They found that Fe is strongly depleted
in these objects, up to 80% with respect to its solar abundance;
- (ii)
- by using the gas physical conditions inferred through spectral diagnostic tools. In this case the estimate does not depend upon assumptions of the shock characteristics (velocity, pre-shock density, magnetic field intensity) but depends on the accuracy of the determinations of the gas parameters. Beck-Winchatz et al. (1994,1996) used this method to estimate Fe and Ni gas-phase abundance in HH jets. They found that Fe is not, or only partially, depleted in HH jets, while Ni is surprisingly over-abundant. The method was then used by Nisini et al. (2005,2002) and Podio et al. (2006) to estimate the gas-phase abundance of both Fe and Ca obtaining a strong depletion (up to 90%) of these species which, however, may vary along the jet and decreases to 0% in some knots.
Following Nisini et al. (2005) and Podio et al. (2006), here we use the
[Ca II]7291/[S II]
6731
ratio to estimate the Ca gas-phase abundance in our sample of HH jets.
These lines have similar excitation temperatures (
K for
[Ca II]
7291 and
K for [S II]
6731) and, despite the
critical density of the calcium line being larger than that of the
sulphur line (
([S II])
cm-3,
([Ca II])
cm-3), at the typical low density of
HH jets the emission of the two lines peaks at the same distance from the shock
front.
In the same region the [S II], [O I], and [N II] lines which are used
to derive the gas physical conditions along the jet are also excited.
Thus the estimated parameters can be used consistently
to model the theoretical [Ca II]/[S II] ratio.
Following Hartigan et al. (2004) we assume that the [Ca II] lines are
collisionally excited and compute the level populations.
As in previous works (Nisini et al. 2005; Podio et al. 2006) we assume that there is no
calcium in the
form of Ca0, because its ionization potential is very low, 6.1 eV.
On the other hand, the ionization potential of Ca II is also quite low,
11.9 eV, i.e. lower than that for hydrogen (
13.6 eV).
In Nisini et al. (2005) and Podio et al. (2006)
we computed the [Ca II]/[S II] ratios by assuming
that all calcium is in the form of Ca+.
This is a reasonable approximation in the case of poorly ionized objects
such as HH 1 and HH 111 (
0.3).
However, the presence of Ca in the form of Ca++ cannot be neglected when
analysing higher excitation jets such as HH 83 and HH 24.
The coefficients for collisional ionization and radiative and dielectronic
recombination of Ca+ and H0 are very similar for temperatures
between
K and
K
(at
K, the recombination rates for Ca+ and H0 are:
(Ca+)
cm3 s-1,
(H0)
cm3 s-1;
while the coefficients for collisional ionization are:
C(Ca+)
cm3 s-1,
C(H0)
cm3 s-1).
Therefore, given that at these temperatures all calcium is in the
form of Ca+ and Ca++, in conditions of ionization equilibrium the
ionization fraction of Ca+ (defined as x(Ca+) = Ca++/Ca) and
H0 are almost the same: x(Ca+)
.
In stellar jets, however, the ionization fraction of hydrogen is never at
equilibrium with the local temperature because,
as explained in Bacciotti & Eislöffel (1999), the gas is moving along the jet and
the recombination timescale at the observed electron densities is of
103 yr, i.e. of the same order of the jet dynamical time.
Because of the similarity of the recombination and collisional ionization
coefficients, however, Ca+ and H0 atoms can be thought to be
undergoing the same processes when moving along the jet.
Thus we assume that, at every position along the jet, the Ca+ ionization
fraction is equal to the hydrogen ionization fraction,
.
Taking Ca and S gas-phase abundances equal to solar ones
(Asplund et al. 2005), we can finally compute the [Ca II]/[S II] theoretical ratios.
Note that,
in contrast with previous works (Nisini et al. 2005; Podio et al. 2006), here we use the
[Ca II]7291/[S II]
6731 ratio, instead of the
[Ca II]
(7291+7324)/[S II]
(6716+6731),
since the [Ca II]
7324 line may be blended with
the [O II]
7318.6, 7319.4 and the [O II]
7329.9, 7330.7 lines.
As explained in Hartigan et al. (2004) the calcium line ratio
[Ca II]
7291/
7324 is expected to be
1.5 both in the case of
collisional excitation and for fluorescent pumping from the
ground state.
Thus by measuring the ratio between the calcium lines we can check for the
presence, if any, of [O II] emission blended with the [Ca II]
7324 line.
We find that this is not the case for the HH 1 and the HH 34 jets
(Nisini et al. 2005; Podio et al. 2006), but may be the case in the HH 111, HH 83 and HH 24
jets which show higher excitation conditions in some of the knots
(see Figs. 3, A.3, and A.4).
These jets, in fact, show a wider [Ca II]
7324 line in some knots
indicating blending with the [O II] lines.
Moreover, in our EFOSC2 spectra, the [Ca II]
7324 line is at the border of the
detector and hence the line flux can be affected by distortion and/or lower S/N.
Thus we decided to use the [Ca II]
7291/[S II]
6731 for all the jets in our
sample.
For the HH 111 and the HH 1 jet we corrected the observed [Ca II]/[S II]
ratios for reddening following the procedure explained in
Sect. 3.
For the HH 83 and the HH 24 jet we have no estimates of the visual extinction,
thus we could not correct the Ca+ and S+ line fluxes for
reddening and the observed ratios (computed assuming
)
are actually upper limits to the real values
(the latter decrease for increasing
).
This means that the calcium depletion inferred from
Figs. 7 and 8 is a lower limit.
The errors affecting the observed ratios are due to the signal-to-noise
over the measured line fluxes.
For HH 1 and HH 111 we consider also the errors due to the uncertainty in
estimating .
The errors on the theoretical ratios are due to the uncertainty of the estimated
physical parameters and are computed by varying ,
,
and
between
their minimum and maximum values.
The influence of these errors is represented in the form of a green stripe in
Figs. 5-8.
Since, as explained in Sect. 3, the choice of the abundance set
may affect the absolute values of
and
,
we also computed the
variations of [Ca II]/[S II] when adopting the alternative abundance set
given by Esteban et al. (2004).
The difference between the [Ca II]/[S II] ratios obtained with the two sets
is illustrated by the width of the beige stripes in
Figs. 5-8.
5.2 Results: Ca depletion along HH jets
![]() |
Figure 5:
Comparison between observed (diamonds) and predicted (dashed lines)
[Ca II] |
Open with DEXTER |
![]() |
Figure 6:
Comparison between observed (diamonds) and predicted (dashed line)
[Ca II] |
Open with DEXTER |
![]() |
Figure 7:
Comparison between observed (diamonds) and predicted (dashed line)
[Ca II] |
Open with DEXTER |
![]() |
Figure 8:
Comparison between observed (diamonds) and predicted (dashed line)
[Ca II] |
Open with DEXTER |
Table 1: Gas-phase abundance of calcium with respect to solar Abundances.
The comparison between the observed and the predicted [Ca II]/[S II] ratios
is shown in Figs. 5-8 and the inferred values of the calcium gas-phase abundance
are summarized in Table 1.
Note that [Ca]
/[Ca]
values in Table 1
and in the bottom panel of Fig. 5 are derived by using the
[Ca II]/[S II] ratios predicted assuming solar elemental abundances from
Asplund et al. (2005) (illustrated by a dotted line in the figures).
Along the HH 111 and HH 1 jets the observed [Ca II]7291/[S II]
6731 ratios
are lower than the predicted ones, i.e. calcium is depleted
with respect to its solar abundance, the depletion being between 20% and 80%.
For these jets the results agree with what has been found in Nisini et al. (2005)
and Podio et al. (2006),
slight differences being explained by the fact that in the latter studies
the presence of calcium in the form of Ca++ was neglected and a
different set of abundances was assumed.
In constrast, along HH 83 dust grains appear to have been destroyed
substantially, as the derived Ca abundance is equal or even larger than solar
in the first knots and is depleted up to 47% at distances larger than
25
from the source.
In the case of HH 83, however, we did not correct the fluxes for
extinction, thus the obtained values of Ca gas-phase abundance are upper
limits.
It can be shown that with an
the Ca gas-phase abundance is between 0.36 and 1.
In any case the low depletion in this jet appears to be consistent with its
high excitation.
Finally, in HH 24 we derive the [Ca II]7291/[S II]
6731 ratio only in
a few of the observed knots because for the fainter ones the Ca+ emission
is below our assumed S/N threshold (equal to three times the background noise).
Figure 8 shows that the depletion
along the jet shows no evident trend: calcium is depleted by 48%-18%
in HH 24 C, by
38%-15% in HH 24 E and by
60%-31% in HH 24 A.
The lack of a trend along the jet reflects the situation found for the
physical parameters (see Fig. A.4) and supports the idea that
HH 24 C, E, and A may not belong to a single jet (Eislöffel & Mundt 1997).
Note that also in the case of HH 24
we did not correct the observed [Ca II]/[S II]
ratios for extinction thus the inferred Ca depletion values are lower limits.
The width of the beige stripes in Figs. 5-8 show the variation of the theoretical [Ca II]/[S II] ratios when adopting abundances from Esteban et al. (2004). Note that also in this case calcium turns out to be largely depleted in most of the knots, although the inferred Ca depletion is lower.
5.3 Discussion
In general, we find an amount of Ca depletion comparable to the Fe depletion estimated in other studies (Mouri & Taniguchi 2000; Nisini et al. 2005; Podio et al. 2006; Nisini et al. 2002; Böhm & Matt 2001; Beck-Winchatz et al. 1996; Garcia Lopez et al. submitted to A&A). Together the results obtained for the two different refractory species indicate that, despite the presence of shocks, there are still dust grains in the jets. However, partial or total destruction of dust grains can be inferred by comparing the level of calcium depletion in the jets with that of the diffuse interstellar medium in Orion (see Table 1).
According to the most recent models of dust reprocessing from
Guillet et al. (2009) the dust destruction rate in J-shocks with velocities lower
than 50 km s-1 and pre-shock densities from 104 to 106 cm-3
is around a few percent.
Previous models by Jones et al. (1994) predict a destruction rate for silicate
grains as high as 50% for shock velocities of 100 km s-1 and low
values of the b parameter (
b=B0 n0-1/2, where B0 is the
intensity of the magnetic field and n0 is the pre-shock density).
Finally, up to 90% of Fe and 60%-70% of Mg and Si can be released in
gaseous form in C-shocks with velocities of only
45 km s-1(May et al. 2000).
It is important to stress that none of the above mentioned models has been
tested with grains containing atoms of Ca and that
these predictions are highly dependent on dust size,
structure, and composition.
Different results, for example, are retrieved when considering
silicate or graphite grains, and if the grains are solid or porous
(Jones et al. 1994; Guillet et al. 2009).
In any case, none of the present models predict total dust grain
destruction
in the mild shocks that occurr in the jet where gas and dust interact
with matter already put into motion by the passage of previous fronts
(
km s-1 in the jet, Hartigan et al. 2001).
This discrepancy between the predictions of shock models and the measured
Ca gas-phase abundance is easily solved if we assume that the dust has been
reprocessed in the disk or circumstellar region
before being extracted and accelerated into jets.
According to recent theoretical models and observations, the accretion
disk in Herbig and Classical T Tauri Stars is populated by dust grains only
ouside of the so-called dust evaporation radius,
(e.g., Isella & Natta 2005).
One can thus infer that
the material in the jet which is coming from a region smaller than
would not contain grains, as these have been
destroyed by the stellar radiation.
The dust evaporation radius is located at
0.1-1.5 AU in Herbig
stars (Eisner et al. 2007) and at
0.05-0.3 AU in T Tauri stars
(Akeson et al. 2005).
Note however there are no estimates of the radius for younger and more embedded
sources such as those giving rise to the HH jets in our sample.
On the other hand, depletion is observed at various levels along the jet,
and to explain the presence of dust grains
we have to assume that part of the material in the jet is coming from a region
of the disk which extends beyond
.
Note that we are not observing ambient dust since this has been swept-out or
destroyed by the passage of major bow-shocks, which have larger
shock velocities (for example
up to
150 km s-1 in the bow of HH 1,
Bally et al. 2002) and the dust reformation time-scale is much longer
than the dynamical timescale of the jet, i.e. 103-104 yr.
Moreover, the rise in gas-phase calcium with distance from the source
along the HH 111 and HH 1 jets, and the lack of deceleration, argues
against significant entrainment of ambient matter along the jet,
as would be suggested by
models of prompt and/or turbulent entrainment (Stahler 1994; Masson & Chernin 1993).
Interestingly, in HH 111, for which we could perform the analysis in
the two velocity intervals, the Ca depletion is at a minimum in
the HVI of the bow-shaped knot L
(see HST images from Reipurth et al. 1997) which is also the one with the
highest shock velocity (
km s-1, Hartigan et al. 2001).
This result suggests that the efficiency of shocks in destroying dust grains
strongly depends on the shock velocity, in agreement with the
predictions of theoretical models (Draine 2003; Jones et al. 1994; Guillet et al. 2009).
Our estimate also indicates that this shock should be able to destroy at
least 40% of the dust grains present at this position.
The proposed link between destruction efficiency and shock velocity is also supported by our finding that calcium is less depleted, or not depleted at all, along HH 83 and HH 24. These jets, in fact, show higher excitation conditions in comparison to HH 1 and HH 111 suggesting that their gas has undergone stronger internal shocks.
6 Summary and conclusions
In this paper we analyse optical spectra of a sample of HH jets located in the Orion and in the Lupus molecular clouds (HH 111, HH 1/2, HH 83, HH 24 M/A/E/C, and Sz68), to investigate primarily the amount of Ca in the gas-phase and, in turn, the dust reprocessing in jets. Towards this aim we first apply spectral diagnostics (Bacciotti & Eislöffel 1999) to derive the physical conditions of the gas in the jets from the line ratios of the detected forbidden lines, i.e. the values of the electron and total density, the ionisation fraction, and the temperature. These parameters are then used to estimate the expected Ca abundance.
We find that in the studied jets the electron density ranges between 50 and
cm-3.
The ionisation fraction varies from a few percent in low excited knots
where the [N II] emission is very faint to 40-70% in high
excitation knots with strong [N II] emission.
The temperature follows the same trend as the ionisation fraction, being
104 K in the weakly ionised knots and
1.5-
K in the
highly ionised ones, in agreement with predictions from shock
models (e.g., Hartigan et al. 1987).
Finally,
the total density derived by combining the retrieved values of
and
varies between
100 and
cm-3.
The most interesting results are obtained for the HH 111 jet for which we derive
the values of the physical parameters separately in two velocity intervals,
Low (v >-100 km s-1) and High (v <-100 km s-1) (LVI and HVI).
This analysis highlights many shock signatures such as maxima in the
HVI of the
brightest knots and a sharp increase of
and
in the HVI of the
strong knot L, which has a clear bow morphology in the HST images of
Reipurth et al. (1997)
and is the knot with the highest shock velocity in the jet beam
(Hartigan et al. 2001).
We also analysed for the first time the faint jet powered by the young source
Sz68. In this case we obtain low values for the electron
and total density (
200-103 cm-3 and
2-
cm-3)
and intermediate
values for the ionisation fraction and the temperature
(
0.05-0.25 and
K).
For this jet we also measured proper motions of
300-600 km s-1
from which, assuming
a jet radius of
100-300 AUs we derived a mass loss rate ranging between
a few 10-8 and a few 10-7
yr-1, i.e. of the same order of magnitude
of
found for typical HH jets (Bacciotti & Eislöffel 1999; Hartigan et al. 1994; Podio et al. 2006).
Using the inferred physical parameters we derive an estimate for the
dust content along the jets by comparing observed and predicted ratios between
calcium and sulphur lines.
The derived Ca gas-phase abundance is lower than solar
([Ca]
/[Ca]
-0.85) in all the jets, with the
exception of HH 83 to which, however, no dereddening has been applied, thus
yielding upper limits to the Ca gas-phase abundance.
In general,
the measured Ca depletion is lower with respect to
the interstellar medium ([Ca]
/[Ca]
in Orion,
Baldwin et al. 1991).
This result suggests that mild shocks occurring along the jets
are partially destroying dust grains, thus releasing calcium atoms into
gaseous form.
Interestingly, a higher degree of depletion is estimated in the low-excited
jets (i.e. HH 111 and HH 1), while little or no depletion is found along
HH 83 and HH 24 which show higher ionization fractions and temperatures.
This seems to indicate that the efficiency of shocks in
reprocessing dust grains strongly increases with shock
velocity, as expected from theoretical models
(e.g., Jones et al. 1994; Guillet et al. 2009).
This idea is also supported by the finding that along HH 111 the
depletion is minimum in the HVI of knot L, which has the largest shock velocity
(Hartigan et al. 2001).
The measured values of Ca gas-phase abundance, however, are too large
to be explained by dust reprocessing in shocks which are slower than
100 km s-1 (Jones et al. 1994; Guillet et al. 2009; May et al. 2000).
This discrepancy between observational results and predictions of
theoretical models may indicate that the material in the jet is
extracted from a wide region of the disk that includes both the inner region
close to the star, inside the so-called dust evaporation radius,
,
where the dust is destroyed by the stellar radiation (Isella & Natta 2005),
and the region beyond
from where the dust is lifted and accelerated
in the flow.
Modeling, however, is required to test this idea.
We are grateful to G. Pineau des Forêts, V. Guillet, and S. Cabrit for useful discussions on dust reprocessing in shocks. We are also grateful to the referee for the attentive comments that allowed us to improve the first version of this paper. We thank the Irish Research Council for Science, Engineering and Technology which funded the work of Linda Podio. This work was partially supported by the European Community's Marie Curie Research and Training Network JETSET (Jet Simulations, Experiments and Theory) under contract MRTN-CT-2004-005592. Tom Ray was partially supported by Science Foundation Ireland through their Research Frontiers Programme.
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Online Material
Appendix A: Physical structure of previously analysed jets
In this appendix we present the results obtained by
applying spectral diagnostics to HH 1 (jet and bow), HH 2, HH 83, and HH 24.
These jets have been already analysed in previous papers with the same
diagnostics, but are re-analysed here with higher spatial sampling
(
)
and an improved version of the diagnostic code
as explained in Sect. 3.
The values presented are used for investigating dust reprocessing
in these jets (see Sect. 5).
A.1 HH 1/2: the jet and its terminal bows
The physical structure of the HH 1 jets and its terminal bows has been previously analysed by Solf & Böhm (1991), Nisini et al. (2005), and Böhm & Solf (1985), Solf et al. (1988), Eislöffel et al. (1994), Bally et al. (2002). In our observations the slit is aligned along the jet and it covers the fainter western/eastern parts of HH 1 and HH 2 respectively, i.e. the knots B and G, which are located to the west of the HH 1 bow apex (knot F), and the knots L, J, G, B, T, and Q located to the east of the bright H and A knots in HH 2.
The results obtained by applying the BE technique to the HH 1 jet are shown in Fig. A.1. The values for the physical parameters are in a good agreement with the ones obtained in Nisini et al. (2005) but here the spatial sampling is two times better.
![]() |
Figure A.1:
Variation of the physical parameters for the HH 1 jet as a
function of distance from the source.
From top to bottom panel:
intensity profiles of the optical lines, the electron density, |
Open with DEXTER |
![]() |
Figure A.2:
Analysis of the excitation conditions in the HH 1 jet and its
terminal bows HH 1 and HH 2 as a
function of distance from the source.
From top to bottom panel:
intensity profiles of the optical lines,
the electron density, |
Open with DEXTER |
We could not apply the diagnostic to the terminal bows. The BE technique, in fact, relies on the assumption of low excitation conditions, i.e. S being totally ionised but there is no S++, and oxygen and nitrogen are ionised at most once (Bacciotti & Eislöffel 1999). This assumption is satisfied for the jet where gas interacts with material already in motion from previous outflow events, but may not be correct when dealing with the terminal bows. This picture is confirmed by proper motions studies by Bally et al. (2002) which indicate shock velocities lower than 30 km s-1 in the jet and velocity jumps of up to 100-200 km s-1 in the bows. Moreover, Böhm & Solf (1985) and Solf et al. (1988) detected many high excitation lines in HH 1 and HH 2 such as O++, S++, and Ar+++.
Even if in our case the slit covers the lateral part of the bows
the excitation conditions may still be too high.
This is indicated by the detection of the high excitation [Ar III]7135.8
line in our spectra and by the fact that the [N II] lines,
which are faint in the jet, show strong emission in the bows,
comparable to the [S II] lines (see the upper panel of Fig. A.2).
In order to compare the physical conditions in the jet and in the bows
we computed the electron density, ,
and the [N II]/[O I] ratio along all the slit length.
The electron density does not show large differences.
It varies between 0.05-
cm-3 along all the jet.
While in the jet
is decreasing with distance from the source,
both in the HH 1 and in HH 2 bows, on the contrary,
there is a decreasing trend going from the shock apex toward the source.
This is expected behind a shock front and confirms the results found by
Böhm & Solf (1985) and Solf et al. (1988).
Our values of
are lower than those inferred by Böhm & Solf (1985), however,
because of the different slit alignment (on the brightest spots at the apex
of HH 1 and HH 2 for Böhm & Solf 1985, and on the bows wings in our case)
and show that the density in the two bows is maximum at the apex and
fades towards the wings.
The [N II]/[O I] ratio is a good indicator of excitation conditions and,
in particular, of the ionisation state of the gas.
Figure A.2 shows that [N II]/[O I] is <1 in the jet
and in knot L, while it is >1 in the HH 1 and HH 2 bows, indicating
that the excitation level is much higher in the bows, and a higher value of
is expected.
A.2 HH 83: the physical structure of the inner knots
The physical structure of the HH 83 jet has been already derived in Podio et al. (2006). Thanks to the high S/N of these data, however, we obtained a sampling which is four times larger in comparison to previous results. Moreover the good quality of the data, which allowed us to properly subtract the continuum emission from the reflection nebula Re 17 (Rolph et al. 1990), and the use of the improved diagnostic code allowed us to estimate the gas physical conditions in the inner part of the jet, where emission from the [N II] lines is comparable to the [S II] and [O I] emission.
The derived parameters, shown in Fig. A.3, indicate that the excitation conditions are very high
in the knots close to the source (knots from A to D)
with
400-700 cm-3, and high values of the ionisation fraction
and temperature (
0.4-0.7, and
1.5-
K).
![]() |
Figure A.3:
Variation of the physical parameters for the HH 83 jet as a
function of distance from the source.
From top to bottom panel:
intensity profiles of the optical lines, the electron density, |
Open with DEXTER |
A.3 The HH 24 jets
The physical parameters along the jets HH 24 C, E, and A, have already been
derived in Bacciotti & Eislöffel (1999) and Podio et al. (2006).
In our observations the slit has been aligned along the axis HH 24 M-A-E and
thus only partially covers the knots of the HH 24 C jet up to 40
(knot C6).
This is why the line profiles in the upper panel of
Fig. A.4 show fainter emission in the knots of group C,
contrary to what was found in Podio et al. (2006), where the slit was aligned
along the HH 24 C jet.
The variation of the physical parameters obtained by applying the BE technique
is shown in Fig. A.4.
The sampling is improved by around one third with respect to previous analyses
(Bacciotti & Eislöffel 1999; Podio et al. 2006)
allowing us to highlight the different excitation conditions in the
various groups of knots detected in the HH 24 complex
(HH 24 A, HH 24 M/E, and HH 24 C)
and supporting the idea that these knots may belong to different jets.
![]() |
Figure A.4:
Variation of the physical parameters for the HH 24 C/E jet as a
function of distance from the source.
From top to bottom panel:
intensity profiles of the optical lines, the electron density, |
Open with DEXTER |
Footnotes
- ... jets
- Based on observations collected at the European SouthernObservatory, La Silla, Chile (ESO programmes 60.C-0398(A).
- ...
- Appendix A is only available in electronic form at http://www.aanda.org
All Tables
Table 1: Gas-phase abundance of calcium with respect to solar Abundances.
All Figures
![]() |
Figure 1:
Position-velocity diagrams of the forbidden line ratios used in
the diagnostic technique for the HH 111 jet.
From top to bottom panel:
the [S II] |
Open with DEXTER | |
In the text |
![]() |
Figure 2:
Variation of the forbidden line ratios used in
the diagnostic technique for the HH 111 jet as a
function of distance from the source and in two velocity intervals:
the high velocity interval (HVI, red points/solid line, v <-100 km s-1) and
the low velocity interval (LVI, black points/dotted line, v >-100 km s-1).
In contrast, in the terminal bow (knot V) the ratios have been computed
by integrating the fluxes over the line spectral profile thus there is no
information on the velocity intervals.
From top to bottom panel:
intensity profiles of the optical lines, [S II] |
Open with DEXTER | |
In the text |
![]() |
Figure 3:
Variation of the physical parameters for the HH 111 jet as a
function of distance from the source and in two velocity intervals:
the high velocity interval (HVI, red points/solid line, v <-100 km s-1) and
the low velocity interval (LVI, black points/dotted line, v >-100 km s-1).
In contrast, in the terminal bow (knot V) the parameters are derived
by integrating the fluxes over the line spectral profile thus there is no
information on the velocity intervals.
From top to bottom panel:
intensity profiles of the optical lines, the electron density, |
Open with DEXTER | |
In the text |
![]() |
Figure 4:
Variation of the physical parameters for the Sz68 jet as a
function of the distance from the source.
From top to bottom panel:
intensity profiles of the optical lines, the electron density, |
Open with DEXTER | |
In the text |
![]() |
Figure 5:
Comparison between observed (diamonds) and predicted (dashed lines)
[Ca II] |
Open with DEXTER | |
In the text |
![]() |
Figure 6:
Comparison between observed (diamonds) and predicted (dashed line)
[Ca II] |
Open with DEXTER | |
In the text |
![]() |
Figure 7:
Comparison between observed (diamonds) and predicted (dashed line)
[Ca II] |
Open with DEXTER | |
In the text |
![]() |
Figure 8:
Comparison between observed (diamonds) and predicted (dashed line)
[Ca II] |
Open with DEXTER | |
In the text |
![]() |
Figure A.1:
Variation of the physical parameters for the HH 1 jet as a
function of distance from the source.
From top to bottom panel:
intensity profiles of the optical lines, the electron density, |
Open with DEXTER | |
In the text |
![]() |
Figure A.2:
Analysis of the excitation conditions in the HH 1 jet and its
terminal bows HH 1 and HH 2 as a
function of distance from the source.
From top to bottom panel:
intensity profiles of the optical lines,
the electron density, |
Open with DEXTER | |
In the text |
![]() |
Figure A.3:
Variation of the physical parameters for the HH 83 jet as a
function of distance from the source.
From top to bottom panel:
intensity profiles of the optical lines, the electron density, |
Open with DEXTER | |
In the text |
![]() |
Figure A.4:
Variation of the physical parameters for the HH 24 C/E jet as a
function of distance from the source.
From top to bottom panel:
intensity profiles of the optical lines, the electron density, |
Open with DEXTER | |
In the text |
Copyright ESO 2009
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