source
stringlengths 1
2.05k
⌀ | target
stringlengths 1
11.7k
|
|---|---|
Table 3 summarises measurements of RVs in different groups lines, which are arranged in order of increasing RV amplitude. | Table \ref{mean_rv} summarises measurements of RVs in different groups of lines, which are arranged in order of increasing RV amplitude. |
of The mean stellar velocity is +23 kmss~!. | The mean stellar velocity is +23 $^{-1}$. |
Note that the mean RV is more positive for the "high-temperature" lines, up to +36 kmss~! for the narrow component. | Note that the mean RV is more positive for the ''high-temperature'' lines, up to +36 $^{-1}$ for the narrow component. |
Our search for radial velocity changes was then extended to the other cTTS included in the present study. | Our search for radial velocity changes was then extended to the other cTTS included in the present study. |
As shown in Fig. 11, | As shown in Fig. \ref{anti-ph0}, |
all programme stars show the same type of anti-phase variations as observed for DR | all programme stars show the same type of anti-phase variations as observed for DR |
transitious. providing it with its nune “line-driven wind. | transitions, providing it with its name “line-driven wind”. |
Ou the one hand. mass loss is thought to be a key agent in revealing cliemically processed material at the stellar surface. making it responsible for evolutionary scenarios such as the O > Lunminous Blue Variable (LBV) » WolfRavet (WR) star. » SN sequence (e.g. Conti 1976. Cliosi Maeder 1986. Langer et al. | On the one hand, mass loss is thought to be a key agent in revealing chemically processed material at the stellar surface, making it responsible for evolutionary scenarios such as the O $\rightarrow$ Luminous Blue Variable (LBV) $\rightarrow$ Wolf-Rayet (WR) star $\rightarrow$ SN sequence (e.g. Conti 1976, Chiosi Maeder 1986, Langer et al. |
1991). | 1994). |
Furthermore. it determines the stellar mass before collapse and is thus relevant for the type of compact relmnaut that is left behind G.e. neutron star or black hole). | Furthermore, it determines the stellar mass before collapse and is thus relevant for the type of compact remnant that is left behind (i.e. neutron star or black hole). |
Ou the other haud. the role of mass loss may be equally relevant for the loss of angular 1iomienutuum (e.g. Mevuet Maeder 2003). | On the other hand, the role of mass loss may be equally relevant for the loss of angular momentum (e.g. Meynet Maeder 2003). |
With respect to the latter. it has been suggested that low metallicity (actually low "àirou couteuts: Vink cde Ixoter 2005) leads to less mass aud angular moment loss iu low metallicity euvironmneuts. perliaps resulting in a preference of lone gamma-ray bursts (CORB) in the carly Universe. but however interesting the metallicity dependence aud the long CRB puzzle may be. the temperature dependence of stellar winds and its role in the augular momentum evolution of massive stars has been highlighted more receutly with respect to the possibility of (Wink ct al. | With respect to the latter, it has been suggested that low metallicity (actually low “iron” contents; Vink de Koter 2005) leads to less mass and angular momentum loss in low metallicity environments, perhaps resulting in a preference of long gamma-ray bursts (GRB) in the early Universe, but however interesting the metallicity dependence and the long GRB puzzle may be, the temperature dependence of stellar winds and its role in the angular momentum evolution of massive stars has been highlighted more recently with respect to the possibility of (Vink et al. |
2010). | 2010). |
Caven the crucial role that mass loss plays for luassive star evolution. we discuss the theory of massive star nass loss aud its duplications. with a focus on the imetallicitv (Z) aud effective temperature (Tig) dependence. | Given the crucial role that mass loss plays for massive star evolution, we discuss the theory of massive star mass loss and its implications, with a focus on the metallicity $Z$ ) and effective temperature $T_{\rm eff}$ ) dependence. |
We will see that theopacity plays a domimaut role iu both cases. | We will see that the plays a dominant role in both cases. |
The theorv goes back to the early 1970s when Lucy Solomon ((1970) sugeested.OO that selective. radiation pressure on spectral lues is capable of driving stellar | The theory goes back to the early 1970s when Lucy Solomon (1970) suggested that selective radiation pressure on spectral lines is capable of driving stellar |
the Antennae galaxies have comparable escape velocities (?).. but ? and ? find that the star clusters in the Autennac show the same “infant mortality” pheuomenon asx those in the Milkv Wav: roughlv of all star clusters dissolve within ~10 Myr of formation. almost certainly because the formation process operates with a low cficiency and removal of a majority of gas leaves the remaining stars unbound. | the Antennae galaxies have comparable escape velocities \citep{whitmore99a}, but \citet{fall05a} and \citet{whitmore07a} find that the star clusters in the Antennae show the same “infant mortality" phenomenon as those in the Milky Way: roughly of all star clusters dissolve within $\sim 10$ Myr of formation, almost certainly because the formation process operates with a low efficiency and removal of a majority of gas leaves the remaining stars unbound. |
Clearly some mechanism must remove eas from these clusters as they form. and we argue below that radiation-driven regious are the most natural explanation. | Clearly some mechanism must remove gas from these clusters as they form, and we argue below that radiation-driven regions are the most natural explanation. |
Our approach to the problem is as follows. | Our approach to the problem is as follows. |
In 2 we derive a condition for when radiation pressure is nmuportaut. aud we then give a solution to the idealized problemi of an region expaudius iuto au züubient iiedium iucludiug radiation pressure effects. | In \ref{derivation} we derive a condition for when radiation pressure is important, and we then give a solution to the idealized problem of an region expanding into an ambient medium including radiation pressure effects. |
Iu 3 woe discuss the contribution of trapped radiation to region dynamics. | In \ref{trapping} we discuss the contribution of trapped radiation to region dynamics. |
In 4. we discuss the relative importance of radiation-driven regions aud supernovac. | In \ref{supernovae} we discuss the relative importance of radiation-driven regions and supernovae. |
Finally. we suumuauizein& 5. | Finally, we summarize in \ref{summary}. |
Consider a source of bolometric huuinositv £ that produces ionizing photous at a rate 5. located at 7—0. | Consider a source of bolometric luminosity $L$ that produces ionizing photons at a rate $S$, located at $r=0$. |
Following ?.. we investigate two cases. which cau be treated in parallel. | Following \citet{matzner02}, we investigate two cases, which can be treated in parallel. |
The first is a region of neutral gas of deusity p=polefry)Me. | The first is a region of neutral gas of density $\rho=\rho_0 (r/r_0)^{-\krho}$. |
The second is a reeion in which the density is p=py(r/ro)"v for «>O aud 0 for e«0. | The second is a region in which the density is $\rho=\rho_0 (r/r_0)^{-\krho}$ for $x>0$ and 0 for $x<0$. |
The former corresponds to the case of an “embedded” region that is completely surrounded by deuse gas. and the latter to a “blister” region iu which the diiviug source is at the edge of a dense cloud. aud the ionized eas can escape frecly, | The former corresponds to the case of an “embedded" region that is completely surrounded by dense gas, and the latter to a “blister" region in which the driving source is at the edge of a dense cloud, and the ionized gas can escape freely. |
We illustrate these two xossible configurations iu Figure 1.. | We illustrate these two possible configurations in Figure \ref{diagram_hiireg}. |
We define ej=13.6 eV as the threshold photon euergyv required to ionize a jeutral hyvdroec4 atom. and for convenience we define CHdk(Sey) o be the ratio of the stars bolometric oower to its jionizing power. counting onlv an cucrey ἐμ CL ioniziue plicXtou. | We define $\epsth = 13.6$ eV as the threshold photon energy required to ionize a neutral hydrogen atom, and for convenience we define $\psi=L/(S\epsth)$ to be the ratio of the star's bolometric power to its ionizing power, counting only an energy $\epsth$ per ionizing photon. |
For massive stars and clusters whose "nnumositv comes mostly frou massive stars. 0671. | For massive stars and clusters whose luminosity comes mostly from massive stars, $\psi\sim 1$. |
We wish to «etermine how the gas moves in response o the radiation flux. aud to understand uuder what cieunistances radiation plavs an important role iu determining gas motions. | We wish to determine how the gas moves in response to the radiation flux, and to understand under what circumstances radiation plays an important role in determining gas motions. |
Following the usual procedure. we approximate that the ionized gas is isothermal at temperature yy. aud has a sound speed ey which is much leger than the sound speed in the neutral eas. | Following the usual procedure, we approximate that the ionized gas is isothermal at temperature $\tii$, and has a sound speed $\cii$ which is much larger than the sound speed in the neutral gas. |
If radiation pressure is negligible. we have the usual 7— solution: the ionized material expands duc to its thermal pressure. | If radiation pressure is negligible, we have the usual \citet{spitzer78} solution: the ionized material expands due to its thermal pressure. |
This expansion sweeps the neutral eas iuto a thin shell. which coutaius most of the mass that was originally inside the radius my of the reeion. | This expansion sweeps the neutral gas into a thin shell, which contains most of the mass that was originally inside the radius $\rii$ of the region. |
If radiation pressure is the dominant force acting ou the eas then a fluid clement at a distance r from the source undergoes a radiative acceleration dyad[fst)Le7(Acre)dp, whore jo(r) is the opacity of the fiui at radius r to photous of frequency pr. and T(r)=iGpcrdre! is the optical depth frou the source ton that point. | If radiation pressure is the dominant force acting on the gas then a fluid element at a distance $r$ from the source undergoes a radiative acceleration $a_{\rm rad} = \int\, \kappa_\nu(r) L e^{-\tau_\nu(r)}/(4 \pi r^2 c)\, d\nu$ , where $\kappa_\nu(r)$ is the opacity of the fluid at radius $r$ to photons of frequency $\nu$ , and $\tau_\nu(r) = \int_{0}^r \kappa_\nu(r') \rho(r') \, dr'$ is the optical depth from the source to that point. |
Photous below the Lyman luit carry roughly half the radiative ποιοτα. aud suce these are absorbed primarily by dust erains. weeTa? aud thus the radiative acceleration is a decreasing function of ;r. | Photons below the Lyman limit carry roughly half the radiative momentum, and since these are absorbed primarily by dust grains, $\kappa_\nu(r) e^{-\tau_\nu(r)}/r^2$ and thus the radiative acceleration is a decreasing function of $r$. |
The other half of the momenta is carried by photons above the Lyman Πατ, which can be absorbed by either ID or dust erains see Appendix A.. | The other half of the momentum is carried by photons above the Lyman limit, which can be absorbed by either H or dust grains -- see Appendix \ref{dustabsorption}. |
If dust absorption dominates. thle radiative acceleration falls with radius as for lower-energv photons. | If dust absorption dominates, then radiative acceleration falls with radius as for lower-energy photons. |
If II absorption domunates. the acceleration is proportional to the recombination rate. which is cither flat (if the eas density is uniformi) or again declines witli radius (if radiative acceleration causes eas to pile up near the shell edge). | If H absorption dominates, the acceleration is proportional to the recombination rate, which is either flat (if the gas density is uniform) or again declines with radius (if radiative acceleration causes gas to pile up near the shell edge). |
Thus. the total radiative acceleration is alwavs largest closest to the source. and again material will be swept iuto a thin shell of racius ry. | Thus, the total radiative acceleration is always largest closest to the source, and again material will be swept into a thin shell of radius $\rii$ . |
The interior of this shell will be optically thin. | The interior of this shell will be optically thin. |
Thus. after a rapid initial expansion- phase. the dynamics of the eas reduce to the problem of computing the dynamics of the thin shell that bounds the reeion. | Thus, after a rapid initial expansion phase, the dynamics of the gas reduce to the problem of computing the dynamics of the thin shell that bounds the region. |
Following ?.. we can solve this problem by writing down the momentum equation for the shell: where Afj, aud Ag, are the shelbs mass aud area. and py and wy are the deusity and velocity of the eas iuuuediately iuterior to it. | Following \citet{matzner02}, we can solve this problem by writing down the momentum equation for the shell: where $\msh$ and $\ash$ are the shell's mass and area, and $\rhoii$ and $\uii$ are the density and velocity of the gas immediately interior to it. |
The shell area aud mass are duy=(L2) aud AL,=(L2)rip3. where PU)=[BABbepule/ru)Ke is the mean density inside the spherical or hemusplerical region of radius r in the initial cloud. and the values in pareutheses refer to the cases for a (spherical. hemispherical) region. | The shell area and mass are $\ash = (4,2)\pi \rii^2$ and $\msh = (4,2)\pi \rii^3 \overline{\rho}/3$, where $\overline{\rho}(r)=[3/(3-\krho)]\rho_0 (r/r_0)^{-\krho}$ is the mean density inside the spherical or hemispherical region of radius $r$ in the initial cloud, and the values in parentheses refer to the cases for a (spherical, hemispherical) region. |
The first teiii on the right-hand side represeuts eas pressure. while the second represcuts radiation pressure. | The first term on the right-hand side represents gas pressure, while the second represents radiation pressure. |
Note that in writing this equation we have implicitly assumed that all the radiatiou force is applied at the thin shell. rather han in the region interior. | Note that in writing this equation we have implicitly assumed that all the radiation force is applied at the thin shell, rather than in the region interior. |
This is certainly a good approximation when radiation pressure is dominant. since. as discetsed above. the interior of the shell will ο cleared. by radiation pressure. aud thus all photous will be absorbde in or near the shell. | This is certainly a good approximation when radiation pressure is dominant, since, as discussed above, the interior of the shell will be cleared by radiation pressure, and thus all photons will be absorbed in or near the shell. |
When gas pressure doniiuates au the shell interior is of uuiforii deusitv. he rate of ux1ueutuinmi deposition bv ionizing photous natches the xxconibinatiou rate. | When gas pressure dominates and the shell interior is of uniform density, the rate of momentum deposition by ionizing photons matches the recombination rate. |
Since this is wuiform. so the mean radius at which momentum is deposited is 3/L of the shell radius. | Since this is uniform, so the mean radius at which momentum is deposited is $3/4$ of the shell radius. |
Moreover. uou-ionizius plotous. which carry half the total momentum. still deposit alltheir iiomientun iu the shell | Moreover, non-ionizing photons, which carry half the total momentum, still deposit alltheir momentum in the shell. |
Thus our approximation that all the moment is deposited in the shell is a good one. | Thus our approximation that all the momentum is deposited in the shell is a good one. |
The quautity frayin equation (1)) represents the | The quantity $f_{\rm trap}$in equation \ref{momeqn}) ) represents the |
the other haud. projection effects could mask the actual orientation of both structures. for instance Tsvetauov&Walsh(1992) proposed that the ionization cone is very inclined (35°) with respect to the plane of the slaw. | the other hand, projection effects could mask the actual orientation of both structures, for instance \citet{Tsvetanov92} proposed that the ionization cone is very inclined $35^\circ$ ) with respect to the plane of the sky. |
In addition. there is also a tougue extending (075 towards the NNW which euds in two πα. blobs that resembles a structure observed in the excitation maps (sce Figure 2)). | In addition, there is also a tongue extending $0\farcs5$ towards the N–NW which ends in two small blobs that resembles a structure observed in the excitation maps (see Figure \ref{fig:X_OIII}) ). |
The presence of dust within the ionization cone has Όσσα reported before oulv in NGC 1068 (Bocketal.2000). | The presence of dust within the ionization cone has been reported before only in NGC 1068 \citep{Bock00}. |
. Towards the uucleus of the galaxy there is an excess of reddening which can be attributed to a natural increase in the extinction due to higher dust concentration. | Towards the nucleus of the galaxy there is an excess of reddening which can be attributed to a natural increase in the extinction due to higher dust concentration. |
Assunmiug an intrinsic colour similar to that observed iu the disk of galaxies (IIT 2: 1998). we have estimated a value of Ay-=6.5 mae. which results in Ny=1.2«I?7. | Assuming an intrinsic colour similar to that observed in the disk of galaxies \citep[$\mathrm{I-H} 2; , we have estimated a value of ${A_V = 6.5}$ mag, which results in ${N_H = 1.2\times 10^{22}} \mathrm{cm}^{-2}$. |
N-rav spectral analysis of the observed soft N-rav extended. emission is crucial to determine the excitation nmiechanisni of the plasma. aud its relationship to the optical bicone-like structure (see previous section). | X-ray spectral analysis of the observed soft X-ray extended emission is crucial to determine the excitation mechanism of the plasma, and its relationship to the optical bicone-like structure (see previous section). |
The combination of high spectral resolution aud high spatial resolution data is key to achieve this purpose. | The combination of high spectral resolution and high spatial resolution data is key to achieve this purpose. |
Iun this section we describe iu detail the methodology and main results obtained. | In this section we describe in detail the methodology and main results obtained. |
Iu Section ?? we discuss the origiu of this extended ciission based on the results presented iu this section. | In Section \ref{sec:origin} we discuss the origin of this extended emission based on the results presented in this section. |
The analysis of the spectral counts was performed using the software package XSPEC 12.1019:1996). | The analysis of the spectral counts was performed using the software package XSPEC \citep[version 12.4.0\footnote{http://cxc.heasarc.gsfc.nasa.gov/docs/xanadu/xspec/, }; . |
Soft N-rav cussion in Sevfert ealaxies has been proven o consist of a plethora of enüssion nes plus a sinall raction of continui enmiüsson that can be described with a single flat power-law (Cauüinazzictal.2008) with a fixed spectral iudex of P=1. | Soft X-ray emission in Seyfert galaxies has been proven to consist of a plethora of emission lines plus a small fraction of continuum emission that can be described with a single flat power-law \citep{Guainazzi08} with a fixed spectral index of $\rm{\Gamma = 1}$. |
We obtained the emission ine fluxes of the central 30 aresec region (note that lis iucludes the uucleus and. circumuuclear emission) using the data. | We obtained the emission line fluxes of the central 30 arcsec region (note that this includes the nucleus and circumnuclear emission) using the data. |
We searched for the oeseunce of 37 emission lines of €. O. N. Si. Mg and Fe species bv fitting the spectra of the two RCS cameras ο Gaussian profiles together with a continuum. | We searched for the presence of 37 emission lines of C, O, N, Si, Mg and Fe species by fitting the spectra of the two RGS cameras to Gaussian profiles together with a continuum. |
We usec Cash statistic for this purposes. | We used Cash statistic for this purposes. |
The triplet fits were performed keeping the relative distance between centroids im energv aud the ceutrok energv was left as a free pazameter. | The triplet fits were performed keeping the relative distance between centroids in energy and the centroid energy was left as a free parameter. |
A line was considere detected when the flux was higher than 0 at the lo level. | A line was considered detected when the flux was higher than 0 at the ${\sf \sigma}$ level. |
The resulting RCS spectra aud detected euission mes are presented im Figure D and Table 2.. respectively. | The resulting RGS spectrum and detected emission lines are presented in Figure \ref{fig:RGSspec} and Table \ref{tab:RGS}, respectively. |
Al enerev centroids are consistent with the laboratory value eiven the error bars. | All energy centroids are consistent with the laboratory value given the error bars. |
Cominazzietal.(2008). previously studied the spectra of 5573. | \citet{Guainazzi08} previously studied the spectra of 573. |
Unfortunately. they ouly reported some of the dues. all of them agreeiug with our cussion liue fluxes. | Unfortunately, they only reported some of the lines, all of them agreeing with our emission line fluxes. |
The ost intense emission lines comprisue the spectrum are: ο VI Ly. O VII G). O VII (f). O VIII Lvo. O VIL. O VITRRC. Fe NVIT 3d-2p. and Ne IN (1). | The most intense emission lines comprising the spectrum are: C VI $\rm{\beta}$, O VII (r), O VII (f), O VIII $\rm{\alpha}$, O VII $\rm{\gamma}$, O VII RRC, Fe XVII 3d-2p, and Ne IX (r). |
The fit of data with a thermal nodel produces poor results below 2 keV (4? ~16). | The fit of data with a thermal model produces poor results below 2 keV $\sf{\chi^{2}}$ $\sim$ 16). |
lustead. a model composed of multiple emission lines was tried. | Instead, a model composed of multiple emission lines was tried. |
Taking advautage of the fit. we πρόοδος, that the intensity of the lines iu the ow vesolition spectra fit do not exceed the RCS ueasurenieuts. | Taking advantage of the fit, we imposed that the intensity of the lines in the low resolution spectra fit do not exceed the RGS measurements. |
This is acceptable because the crosscalibrations betweeu EPIC aud RCS3ustriuenuts shows a πολλαο constant im the range of 0.9 to 1.0 (seePlu-cluskyetσαal. 2008). | This is acceptable because the cross-calibrations between EPIC and RGS instruments shows a normalization constant in the range of 0.9 to 1.0 \citep[see][]{Plucinsky08}. |
. The assumed Caussian width is 100 eV. Note that for EPIC (aud also Chandra)) data we do not question the existence of the emission lines detected ou the but we use them as a template. | The assumed Gaussian width is 100 eV. Note that for EPIC (and also ) data we do not question the existence of the emission lines detected on the but we use them as a template. |
Triplets were fitted using the total flux of all components of the He-like lines O VIL. N VI. and Ne IX. | Triplets were fitted using the total flux of all components of the He-like lines O VII, N VI, and Ne IX. |
The coutinmun emission was fitted to a power-law to be consistent with the high spectral resolution analysis. | The continuum emission was fitted to a power-law to be consistent with the high spectral resolution analysis. |
However. this fit has poor statistics (\?> 2). | However, this fit has poor statistics $\rm{\chi^{2}_{r} > 2}$ ). |
Five lines were added at energies above 0.95 keV in order to achieve an acceptable fit ο.= 0).8): FoNNN at 0.97 keV (τισ, NoXN Lyo at 1.02 keV. (4225.62). TIX Πο 6 at 1.10 keV (4223.72). XXI. triplet at ~1.33 keV (\2=0.82). and Si XIII triplet at 1.81 keV(\2=0.80). | Five lines were added at energies above 0.95 keV in order to achieve an acceptable fit $\rm{\chi^{2}_{r} = 0.8}$ ): XX at 0.97 keV $\rm{\chi_{r}^{2}}$ =5.64), X $\rm{\alpha}$ at 1.02 keV $\rm{\chi_{r}^{2}}$ =5.62), IX He $\rm{\delta}$ at 1.16 keV $\rm{\chi_{r}^{2}}$ =3.72), XI triplet at $\rm{\sim}$ 1.33 keV $\rm{\chi_{r}^{2}}$ =0.82), and Si XIII triplet at 1.84 $\rm{\chi_{r}^{2}}$ =0.80). |
The final ft is shown in Figure 5.. | The final fit is shown in Figure \ref{fig:xspecXMM}. |
The low spectral resolution spectrum shows the following intense emission lines: VV Tes. VWI Ly. NVVII Lvo. OVVII triplet. OVVIII Lvo. O VILIITes. OVVIT RRC. NXNVII. 3d-2p. aud triplet. XX Lya. HIN Te 8. and Mg NI triplet. | The low spectral resolution spectrum shows the following intense emission lines: V $\rm{\gamma}$, VI $\rm{\beta}$, VII $\rm{\alpha}$, VII triplet, VIII $\rm{\alpha}$, O $\rm{\gamma}$, VII RRC, XVII 3d-2p, and IX triplet, X $\rm{\alpha}$, IX He $\rm{\delta}$, and Mg XI triplet. |
In HIN.the best- modelthe flux of the OVIII RRC feature appears uceelieible aud the adjaceut line NNVIT 3d2p is preseut. | In the best-fit modelthe flux of the OVIII RRC feature appears negligible and the adjacent line XVII 3d2p is present. |
This is in contrast to what happens iu the uuclear spectrum (see Section ?? and Table 3)). | This is in contrast to what happens in the nuclear spectrum (see Section \ref{sec:chandraspec} and Table \ref{tab:low}) ). |
Iu order to check the compatibility of the results of both | In order to check the compatibility of the results of both |
is - 0.09 dex kpc!. to be compared with - 0.044 and -0.04 dex Κροτ. for M33 and the Milky Way respectively. | is - 0.09 dex $^{-1}$, to be compared with - 0.044 and -0.04 dex $^{-1}$ for M33 and the Milky Way respectively. |
Indication of steeper gradients. and sizable chemical evolution. are thus the defining characteristics of M81 compared with other spirals. | Indication of steeper gradients, and sizable chemical evolution, are thus the defining characteristics of M81 compared with other spirals. |
The oxygen enrichment indicates that this galaxy has suffered outflow to a lesser extent than the comparison galaxies. and steeper oxygen gradients are compatible with this explanation. | The oxygen enrichment indicates that this galaxy has suffered outflow to a lesser extent than the comparison galaxies, and steeper oxygen gradients are compatible with this explanation. |
It is evident from the above analysis that the metallicity gradient slopes in the galaxies examined do not depend on the average galactic metallicity: M33 is metal poor (its metallicity is similar to that of the LMC. Leisy Dennefeld 2006) and Μδ| is closer in metal contents to the Galaxy. | It is evident from the above analysis that the metallicity gradient slopes in the galaxies examined do not depend on the average galactic metallicity: M33 is metal poor (its metallicity is similar to that of the LMC, Leisy Dennefeld 2006) and M81 is closer in metal contents to the Galaxy. |
Since M81 is also closer to the Galaxy in stellar mass content (see Table 7 for a direet comparison of the main properties of these two galaxies). it makes sense to compare these two galaxies directly. to explore the relations of the metallicity gradients to the characteristics of the galaxy disks and to their evolution. | Since M81 is also closer to the Galaxy in stellar mass content (see Table 7 for a direct comparison of the main properties of these two galaxies), it makes sense to compare these two galaxies directly, to explore the relations of the metallicity gradients to the characteristics of the galaxy disks and to their evolution. |
The observations of PNe tell us that the gradient of M8I is steeper than that of our Galaxy in old stellar populations: the trend seems to persist in the young stellar population. but this needs further confirmation. | The observations of PNe tell us that the gradient of M81 is steeper than that of our Galaxy in old stellar populations; the trend seems to persist in the young stellar population, but this needs further confirmation. |
From a theoretical point of view. there are two main reasons that could influence the gradient slopes. (1) the rotational velocity of the galaxy and (2) the different situation of these two galaxies within their environment. | From a theoretical point of view, there are two main reasons that could influence the gradient slopes, (1) the rotational velocity of the galaxy and (2) the different situation of these two galaxies within their environment. |
Molla Diaz (2005) noted that metallicity gradients depend on the rotational velocity of the galaxy (thus presumably on its total mass) and on its morphological type: more massive galaxies tend to evolve faster and have flatter gradients than lower mass galaxies: for a given rotation velocity. gradients are steeper for late type than for early type galaxies. | Mollá Diaz (2005) noted that metallicity gradients depend on the rotational velocity of the galaxy (thus presumably on its total mass), and on its morphological type: more massive galaxies tend to evolve faster and have flatter gradients than lower mass galaxies; for a given rotation velocity, gradients are steeper for late type than for early type galaxies. |
Both effects would imply a steeper gradient for the Galaxy than for M81. contrary to the observations. | Both effects would imply a steeper gradient for the Galaxy than for M81, contrary to the observations. |
Some other factor must be at play. | Some other factor must be at play. |
It is likely that the situation of M81 within its group of galaxies influences its metallicity gradient. | It is likely that the situation of M81 within its group of galaxies influences its metallicity gradient. |
The tidal ffeatures near M81 are consistent with a large-scale redistribution of gas in this galaxy. | The tidal features near M81 are consistent with a large-scale redistribution of gas in this galaxy. |
Subsets and Splits
No community queries yet
The top public SQL queries from the community will appear here once available.