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In this paper. we present various infrared 2€CDs for O- C-rich. S-type ACD stars using a revised version of the catalog of AGB stars by Sub Kwon (2009). | In this paper, we present various infrared 2CDs for O-rich, C-rich, S-type AGB stars using a revised version of the catalog of AGB stars by Suh Kwon (2009). |
For each object in the new catalog. we cross-identily the AAA27. MSX and PALASS counterparts by finding the nearest source from the position information in the Z/2:19 PSC. | For each object in the new catalog, we cross-identify the $AKARI$, $MSX$ and $2MASS$ counterparts by finding the nearest source from the position information in the $IRAS$ PSC. |
For the large sample of AGB stars. we present various infrared 2€Ds using {δν (PSC). AAA (PSC and BSC). ALS (DSC) and NIC (A anc £ bands: including 244.195 cata at Avs band) data. | For the large sample of AGB stars, we present various infrared 2CDs using $IRAS$ (PSC), $AKARI$ (PSC and BSC), $MSX$ (PSC) and NIR $K$ and $L$ bands; including $2MASS$ data at $K_S$ band) data. |
In this paper. we are concerned about 2CDs for cillerent classes of AGB stars based on the chemistry of the dust shell and/or the central star. | In this paper, we are concerned about 2CDs for different classes of AGB stars based on the chemistry of the dust shell and/or the central star. |
On the 2€Ds. we plot tracks of the theoretical radiative transfer model results with increasing dust shell optical cepths. | On the 2CDs, we plot tracks of the theoretical radiative transfer model results with increasing dust shell optical depths. |
Comparing the observations with the theoretical tracks. we discuss the meaning of the infrared 2C'Ds. | Comparing the observations with the theoretical tracks, we discuss the meaning of the infrared 2CDs. |
For the sample of AGB stars. we use a revised. version of the catalog by Sub Ixwon (2009). | For the sample of AGB stars, we use a revised version of the catalog by Suh Kwon (2009). |
"μον presented a catalog of AGB stars in our Galaxy [rom the sources Listed in the ZRAS PSC compiling the lists of previous works with verifving processes: their catalog was made of 2193 O-rich stars. 14167 C-rich stars. 287 5 stars and 36 silicate carbon stars. | They presented a catalog of AGB stars in our Galaxy from the sources listed in the $IRAS$ PSC compiling the lists of previous works with verifying processes; their catalog was made of 2193 O-rich stars, 1167 C-rich stars, 287 S stars and 36 silicate carbon stars. |
In this paper. we add SiO maser sources to the List of O-rich AGB stars and more sources for C-rich stars and $8 stars from new references and make some corrections to the previous catalog. | In this paper, we add SiO maser sources to the list of O-rich AGB stars and more sources for C-rich stars and S stars from new references and make some corrections to the previous catalog. |
For a general description of the identifving ancl verifving processes. refer to Sub Ixwon (2009). | For a general description of the identifying and verifying processes, refer to Suh Kwon (2009). |
O-rich AGB stars typically show the conspicuous 10 jum and LS pam features in emission or absorption. | O-rich AGB stars typically show the conspicuous 10 $\mu$ m and 18 $\mu$ m features in emission or absorption. |
They suggest the presence of silicate dust. grains in the outer envelopes around them (e.g.. Suh. 1999). | They suggest the presence of silicate dust grains in the outer envelopes around them (e.g., Suh 1999). |
Low mass-loss rate O-rich AGB (LMOA) stars with thin dust envelopes show the 10 pam and TS yom emission features. | Low mass-loss rate O-rich AGB (LMOA) stars with thin dust envelopes show the 10 $\mu$ m and 18 $\mu$ m emission features. |
High. mass-loss rate O-rich AGB (LAIOA) stars with thick dust. envelopes show the absorbing features at the same wavelengths. | High mass-loss rate O-rich AGB (HMOA) stars with thick dust envelopes show the absorbing features at the same wavelengths. |
OLL maser observations identified many OLL/LB. stars. | OH maser observations identified many OH/IR stars. |
Suh Ixwon (2009) listed. 1533 sources from 14 papers as O-rich AGB stars. | Suh Kwon (2009) listed 1533 sources from 14 papers as O-rich AGB stars. |
However. methanol maser sources at 6.7 Cillz are only associated with massive star formation (\linier et al. | However, methanol maser sources at 6.7 GHz are only associated with massive star formation (Minier et al. |
2003). | 2003). |
More than 550 methanol maser sources have been detected (Pestalozzi. Minier Booth 2005). | More than 550 methanol maser sources have been detected (Pestalozzi, Minier Booth 2005). |
Using the catalog by Pestalozzi et al. ( | Using the catalog by Pestalozzi et al. ( |
2005). we have excluded: 15 methanol maser sources from the ΟΕ maser list in Suh Ixwon (2009). | 2005), we have excluded 18 methanol maser sources from the OH maser list in Suh Kwon (2009). |
Many of the SiO maser sources are O-rich AGB stars. | Many of the SiO maser sources are O-rich AGB stars. |
The catalog of Sub Kwon (2009) did not consider the SiO maser sources. | The catalog of Suh Kwon (2009) did not consider the SiO maser sources. |
By compiling the sources listed in 17 related papers. we have added: 515 SiO maser sources which are identified to be O-rich ACD stars (see Table 1). | By compiling the sources listed in 17 related papers, we have added 815 SiO maser sources which are identified to be O-rich AGB stars (see Table 1). |
655 sources detected. by other methods. including photometrie and spectrometric methods (molecular emission or spectral types) are compiled from 10 related. papers. | 655 sources detected by other methods including photometric and spectrometric methods (molecular emission or spectral types) are compiled from 10 related papers. |
Sec Suh Kwon (2009) for a detailed description. | See Suh Kwon (2009) for a detailed description. |
‘Table 1 shows the revised list of O-rich AGB stars which contains 3003 sources. | Table 1 shows the revised list of O-rich AGB stars which contains 3003 sources. |
Compared to Sub Wkwon (2009). the number has been increased by SLO. | Compared to Suh Kwon (2009), the number has been increased by 810. |
For cach object. we have cross-icentilied the -LALLRZ and MS.N. sources as we describe in Section 3. | For each object, we have cross-identified the $AKARI$ and $MSX$ sources as we describe in Section 3. |
The main components of dust in the envelopes. around carbon stars are believed to be featureless amorphous carbon CXMC) grains and SiC’ erains whieh produce the 11.3. sam | The main components of dust in the envelopes around carbon stars are believed to be featureless amorphous carbon (AMC) grains and SiC grains which produce the 11.3 $\mu$ m |
the age of the Pleiades. but discovered that in order to make their model isochrones a better fit with the members of the Pleiades identitied in ?.. they need to lower the effective temperature of the L-T transition which is set at KK. This suggests that for vounger ages and hence lower gravities. the L-T transition may happen at a lower effective temperature. | the age of the Pleiades, but discovered that in order to make their model isochrones a better fit with the members of the Pleiades identified in \citet{casewell07}, they need to lower the effective temperature of the L-T transition which is set at K. This suggests that for younger ages and hence lower gravities, the L-T transition may happen at a lower effective temperature. |
This is consistent with the analysis of HN Peg B presented by ?.. | This is consistent with the analysis of HN Peg B presented by \citet{leggett08}. |
In IC348. ?. find that for younger objects. a cooler temperature is expected for the same methane colour. | In IC348, \citet{burgess09} find that for younger objects, a cooler temperature is expected for the same methane colour. |
This is consistent with the assumption that the L-T transition. and hence the onset of methane absorption. occurs at a lower temperature for younger (lower gravity) objects. | This is consistent with the assumption that the L-T transition, and hence the onset of methane absorption, occurs at a lower temperature for younger (lower gravity) objects. |
The recent spectroscopic study on the Pleiades L dwarfs by 9. also finds evidence of low gravity features in some early L dwarfs. | The recent spectroscopic study on the Pleiades L dwarfs by \citet{bihain10}, also finds evidence of low gravity features in some early L dwarfs. |
These spectra show the characteristic triangular shaped bump in the in the // band as seen in young L dwarfs in Trapezium (2. | These spectra show the characteristic triangular shaped bump in the in the $H$ band as seen in young L dwarfs in Trapezium \citep{lucas00}. |
This shape could be due to a reduction in the H» collision -induced absorption and water absorption in a low gravity and dustier environment than is commonly seen in field L dwarfs (?:: 22)). | This shape could be due to a reduction in the $_{2}$ collision -induced absorption and water absorption in a low gravity and dustier environment than is commonly seen in field L dwarfs \citealt*{borysow97}; \citealt{kirkpatrick06,mohanty07}) ). |
This dustier environment is consistent with the L-T transition happening at a lower temperature in the Pleiades. | This dustier environment is consistent with the L-T transition happening at a lower temperature in the Pleiades. |
This youth and dusty atmosphere is also given as a reason for the redder ./.—A colours seen in the Pleiades L dwarfs compared to field dwarfs as the He collision -induced absorption opacity also affects the fy band and depends on the square of the gas density (2). which is lower at lower gravity. making the atmosphere more transparent and hence the flux brighter. | This youth and dusty atmosphere is also given as a reason for the redder $J-K$ colours seen in the Pleiades L dwarfs compared to field dwarfs as the $_{2}$ collision -induced absorption opacity also affects the $K$ band and depends on the square of the gas density \citep{knapp04}, which is lower at lower gravity, making the atmosphere more transparent and hence the flux brighter. |
A similar effect was seen by ?. in Upper Scorpius where numerous candidates with spectral types between M and mid-L were Selected from the UKIDSS GCS using photometry. but spectra revealed them to all have much earlier spectral types than expected. | A similar effect was seen by \citet{lodieu09} in Upper Scorpius where numerous candidates with spectral types between M and mid-L were selected from the UKIDSS GCS using photometry, but spectra revealed them to all have much earlier spectral types than expected. |
This lead the authors to deduce that the ./ἐν colour for young L dwarfs is much redder than older L dwarfs of the same spectra type. | This lead the authors to deduce that the $J-K$ colour for young L dwarfs is much redder than older L dwarfs of the same spectral type. |
To compare the PLZI23. PLZI93 and PLZIHOO with models. we have derived effective temperatures using an almost mode independent method described in the following paragraphs. | To compare the PLZJ23, PLZJ93 and PLZJ100 with models, we have derived effective temperatures using an almost model independent method described in the following paragraphs. |
To do this we must assume that PLZI23. PLZJ93 and PLZI100 are al single objects. and indeed. we have no evidence of their being binaries. | To do this we must assume that PLZJ23, PLZJ93 and PLZJ100 are all single objects, and indeed, we have no evidence of their being binaries. |
With the J. Z. J. df. dy. [3.6] and [4.5] filters measurec magnitudes. it is possible to sum the corresponding fluxes. add a small modelled contribution from the wavelengths not covered. and thus. using the distance to the Pleiades. obtain the bolometric luminosity. | With the $I$, $Z$, $J$, $H$, $K$, [3.6] and [4.5] filters measured magnitudes, it is possible to sum the corresponding fluxes, add a small modelled contribution from the wavelengths not covered, and thus, using the distance to the Pleiades, obtain the bolometric luminosity. |
If we further assume a suitable brown dwarf radius. we can then determine the effective temperature. | If we further assume a suitable brown dwarf radius, we can then determine the effective temperature. |
We have used the distance to the Pleiades ppe: 2)). the AMES-COND spectra aos?1. My and the age of the Pleiades MMyr: 29) to calculate 18 brown dwarf radii. R. We then converted our Vega magnitudes to intensities. using the standard zero magnitude intensities (2).. | We have used the distance to the Pleiades pc; \citealt{an07}) ), the AMES-COND spectra \citep{allard01}, $_{K}$ and the age of the Pleiades Myr; \citealt*{stauffer98}) ) to calculate the brown dwarf radii, R. We then converted our Vega magnitudes to intensities, using the standard zero magnitude intensities \citep*{bessell98}. |
These were then multiplied by the known filter widths to obtain fluxes. | These were then multiplied by the known filter widths to obtain fluxes. |
These fluxes were then summed and multiplied by πα” where d is 133 pe (the distance to the Pleiades) to get most of the bolometric luminosity. | These fluxes were then summed and multiplied by $\pi$ $^{2}$ where d is 133 pc (the distance to the Pleiades) to get most of the bolometric luminosity. |
The above paragraph deseribes the principle of the method and also provides a first Tip estimate. using the method below. | The above paragraph describes the principle of the method and also provides a first $_{\rm eff}$ estimate, using the method below. |
However. the Vega spectrum is very different from that of a brown dwarf. | However, the Vega spectrum is very different from that of a brown dwarf. |
Thus for each filter we integrate the Vega spectrum over the filter transmission profile. | Thus for each filter we integrate the Vega spectrum over the filter transmission profile. |
We also integrate an AMES-COND spectrum with the first estimate μι over the filter profile. | We also integrate an AMES-COND spectrum with the first estimate $_{\rm eff}$ over the filter profile. |
Comparing the two gives the normalisation factor of the AMES- spectrum. | Comparing the two gives the normalisation factor of the AMES-COND spectrum. |
Finally. we integrate the AMES-COND spectrum over the wavelength region where the transmission exceeds 5 per cent to calculate the total filter flux. | Finally, we integrate the AMES-COND spectrum over the wavelength region where the transmission exceeds 5 per cent to calculate the total filter flux. |
As above. then summing the | As above, then summing the |
ihe NASA/IPAC Extragalactic Database (NED). which is operated bv the Jet Propulsion Laboratory. California Institute of Technology. under contract with NASA.IST... | the NASA/IPAC Extragalactic Database (NED), which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with NASA.. |
inclination. angle our imaging code (see Section. 4)) has clilliculty determining i£ features are located in the Northern or Southern hemisphere so we have limited the inclination anele range to + 107. | inclination angle our imaging code (see Section \ref{Sec_ima}) ) has difficulty determining if features are located in the Northern or Southern hemisphere so we have limited the inclination angle range to $\pm$ . |
Table 2. lists the basic stellar parameters of HD. 121943. that we have determined. for this παν. | Table \ref{Tab_param} lists the basic stellar parameters of HD 141943 that we have determined for this study. |
As mentioned. there are no error measurements eiven for the Llillenbrandctal.(2008) measurements of ellective temperature and luminosity. so the estimations of the errors in the stellar parameters of LID 141943 may. well be uncerestimiates. | As mentioned, there are no error measurements given for the \citet{HillenbrandLA:2008} measurements of effective temperature and luminosity, so the estimations of the errors in the stellar parameters of HD 141943 may well be underestimates. |
From the values for photospheric temperature. and unspottec Luminosity we can determine the mass anc age of LD 141943 from the evolutionary models of(2000).. see Fig. 1.. | From the values for photospheric temperature and unspotted luminosity we can determine the mass and age of HD 141943 from the evolutionary models of, see Fig. \ref{Fig_Siess}. |
Phis makes HD 141943 a 1.3 star of age ~17 Alves. in reasonable agreement with the values from Cutispotoctal.(2003).. and thus it can be classified as a PAIS star. | This makes HD 141943 a $\sim$ 1.3 star of age $\sim$ 17 Myrs, in reasonable agreement with the values from \citet{CutispotoG:2003}, and thus it can be classified as a PMS star. |
Nordstróm.ctal.(2004) took 7. radial velocity measurements of HD. 141943 over a 2979 day period (epochs and measurements were not given). | \citet{NordstromB:2004} took 7 radial velocity measurements of HD 141943 over a 2979 day period (epochs and measurements were not given). |
Based on the fact that the standard. error of cach measurement was less than the standard: deviation ofthe measurements they determined | Based on the fact that the standard error of each measurement was less than the standard deviation ofthe measurements they determined |
modes. at 1.0pm.8.5prm.17.5jim and. jam JCG0b.. | modes, at $\sim7.0\mic, 8.5\mic, 17.5\mic$ and \\citep[e.g.][]{C60b}. |
Fig. | Fig. |
4. also shows evidence for a feature aboTg. and subtraction of the 2005 spectrum. from the 2007 spectrum leaves a feature with central wavelength 7.01+0.01 ((see ‘Table 1)): there is no evidence for the gum feature. | \ref{IRS1} also shows evidence for a feature at $\sim7$, and subtraction of the 2005 spectrum from the 2007 spectrum leaves a feature with central wavelength $7.01\pm0.01$ (see Table \ref{fluxes}) ); there is no evidence for the ' feature. |
The yam" feature we observe in iis actually atyam... quite dillerent from the expected value of [for gas phase Coo (rerumetal.1991). | The ' feature we observe in is actually at, quite different from the expected value of for gas phase $_{60}$ \citep{C60b}. |
.. However.solid Ci has a feature at. 17.3 citep*C60a.C60c.. closer to the [feature inOph. | However, $_{60}$ has a feature at \\citep*{C60a,C60c}, closer to the feature in. |
. We should therefore consider whether the features in Fig. | We should therefore consider whether the features in Fig. |
6 arise in gaseous or solid. Cou. | \ref{C60} arise in gaseous or solid $_{60}$. |
The fux ratios of the putative Coo. features. enable an estimate of the vibrational temperature. Zi. if. the Coo is in gaseous form. | The flux ratios of the putative $_{60}$ features enable an estimate of the vibrational temperature $T_{\rm vib}$ if the $_{60}$ is in gaseous form. |
Using Einstein cocllicicnts from Alitzner&Campbell (1995: included in Table 1)). values of Lan~520+50 Ix are obtained: however the jm flux seems underestimated bv a factor 2. | Using Einstein coefficients from \citeauthor{mitzner} (1995; included in Table \ref{fluxes}) ), values of $T_{\rm vib}\sim520\pm50$ K are obtained; however the ' flux seems underestimated by a factor $\sim2$. |
A similar value (~670 Ix) is obtained assuming that the energy ofa ~10 eV photon absorbed. by a Coo molecule is equally. distributed amongst the available vibrational modes. | A similar value $\sim670$ K) is obtained assuming that the energy of a $\sim10$ eV photon absorbed by a $_{60}$ molecule is equally distributed amongst the available vibrational modes. |
However. at. this temperature the feature would have a [lux ~1.5.10DL Wom 7. far greater than observed. | However, at this temperature the ' feature would have a flux $\sim1.5\times10^{-15}$ W $^{-2}$, far greater than observed. |
We also note that laboratory measurements on solid. ου (Ixràtschmoeretal.1990a) suggest that the feature is rather weaker than the other three. | We also note that laboratory measurements on solid $_{60}$ \citep{C60a} suggest that the ' feature is rather weaker than the other three. |
Therefore on the basis of (1) the wavelength of the m feature and (i) the weakness of the pim feature. we conclude that the Coo in iis most likely in solid form: if so this is the first astrophysical detection of solid. Cao. | Therefore on the basis of (i) the wavelength of the ' feature and (ii) the weakness of the ' feature, we conclude that the $_{60}$ in is most likely in solid form; if so this is the first astrophysical detection of solid $_{60}$. |
The absorption cross-section of Cao has been measured bv Yagietal.(2009)... from which we estimate the Planck mean absorption cross-section per Cao molecule (averaged over the emission of the B star) to be ~710.74 mt. | The absorption cross-section of $_{60}$ has been measured by \cite{yagi}, from which we estimate the Planck mean absorption cross-section per $_{60}$ molecule (averaged over the emission of the B star) to be $\sim7\times10^{-21}$ $^2$. |
Lf (ef. | If (cf. |
Section 2)) the B star is situated at ~7.2;1044 m from the inner boundary. of the dust shell. the temperature of a Cao grain of radius e is ~200(a/0.03pm)! K. While the apparent absence of Coy in the HUS spectrum immecdiately. after eclipse in 2005. and its presence in 2007. is sugeestive. it is dillieult to argue that the eclipse is in any wav connected with the presence of Coo in the spectrum. especially as it is the giant that is eclipsed: it is likely therefore that the appearance of Coy in 2007 is unconnected with the eclipse of Fig. 3.. | Section \ref{binary}) ) the B star is situated at $\sim7.2\times10^{11}$ m from the inner boundary of the dust shell, the temperature of a $_{60}$ grain of radius $a$ is $\sim200\,(a/0.03\mic)^{1/4}$ K. While the apparent absence of $_{60}$ in the IRS spectrum immediately after eclipse in 2005, and its presence in 2007, is suggestive, it is difficult to argue that the eclipse is in any way connected with the presence of $_{60}$ in the spectrum, especially as it is the giant that is eclipsed: it is likely therefore that the appearance of $_{60}$ in 2007 is unconnected with the eclipse of Fig. \ref{LC}. |
We can make an estimate of the mass of Caso. using the combined Dux in the Cao features. | We can make an estimate of the mass of $_{60}$ using the combined flux in the $_{60}$ features. |
Assuming (cf. | Assuming (cf. |
Section 2)) the B star is situated at ~7.2.101 m from the inner boundary. of the dust shell. and. using the Planck mean absorption cross-section above. the absorbed power per Coo particle is —S1 ο ‘Phe emitted. power (assuming is ~30107 NV. so —ALS1075 Coy particles (Le. ~ y). in solid form. are required. | Section \ref{binary}) ) the B star is situated at $\sim7.2\times10^{11}$ m from the inner boundary of the dust shell, and using the Planck mean absorption cross-section above, the absorbed power per $_{60}$ particle is $\sim8.1\times10^{-18}$ W. The emitted power \citep[assuming a distance of 2~kpc for \XX;][]{evansetal93} is $\sim3.9\times10^{26}$ W, so $\sim4.8\times10^{43}$ $_{60}$ particles (i.e. $\sim2.9\times10^{-11}$ ), in solid form, are required. |
This suggests that the number of Cyg molecules is 0.03 the number of PALL molecules. | This suggests that the number of $_{60}$ molecules is $\sim0.03$ the number of PAH molecules. |
excellent. | excellent. |
Zoccalietal. themselves founc a BOO agreement between thei observatioua valttes of AVPMPIB and the models of Cassisi&Saaris(1997) asstnilg a alplia-capture overabuudanuce a /Fe] = 0.30 for [Fe/H] « —1.0. | \citeauthor{zoccali} themselves found a good agreement between their observational values of $\Delta
V^\mathrm{bump} _\mathrm{HB}$ and the models of \citet{cas3}
assuming an alpha-capture overabundance $\alpha$ /Fe] = 0.30 for [Fe/H] $<$ $-$ 1.0. |
Our resul show. similarly. a good agreement between observatious and theoretical values of Vip calculated witha median alpha-capti'e overabundauce [a /Fe — 0.15. | Our results show, similarly, a good agreement between observations and theoretical values of $V_\mathrm{bump}$ calculated with a median alpha-capture overabundance $\alpha$ /Fe] = 0.45. |
loreover. ihe ‘esults of our Monte. Calo uncertainty analysis W.how that differences. between heoretical aud observed values of the bUup maguiude lie witli1 the theoretical uncertaluties in stellar evolutjon as well as within the oervational uncertainties in the [n]«;lobular cluster cisalice uodulus aud the metallicity scale. | Moreover, the results of our Monte Carlo uncertainty analysis show that differences between theoretical and observed values of the bump magnitude lie within the theoretical uncertainties in stellar evolution as well as within the observational uncertainties in the globular cluster distance modulus and the metallicity scale. |
It is clear that there is currently. uo cdiscrepaucy between the V-baud bunp magniude as observed in Gaactic globular clusters aud as calculated iu staudal stellar evolu101 moclels. | It is clear that there is currently no discrepancy between the $V$ -band bump magnitude as observed in Galactic globular clusters and as calculated in standard stellar evolution models. |
To analyze the iuclivicual influence of each continuously varying stellar evolution paralneter onMya. al 1120 Monte €‘arlo realigatious are plotted on a graph of versus paraljeter value. | To analyze the individual influence of each continuously varying stellar evolution parameter on, all 1120 Monte Carlo realizations are plotted on a graph of versus parameter value. |
TI ence of ool a giveu parameter value is then ¢haracterized with a straight line fit. | The dependence of on a given parameter value is then characterized with a straight line fit. |
Specifically. the parameter ange is divded into 20 bi llaining 56 Monte Carlo realizatious each. and for each bin the neciai value of aud the confidence lit etermined. | Specifically, the parameter range is divided into 20 bins containing 56 Monte Carlo realizations each, and for each bin the median value of and the confidence limits are determined. |
A straielt line is fit to the meclian points of each jn using the simje least-serares method with each bi1 weiglited equally. | A straight line is fit to the median points of each bin using the simple least-squares method with each bin weighted equally. |
Straight lines are also fit o the upper aud ower dence poiuts in each ju. | Straight lines are also fit to the upper and lower confidence points in each bin. |
Table 3 presentis tie. slope of the ]near dlepeuclence for tlie most sigit stellar evolution pa'aineters as well as the otal change in ulnup uaguitucde as the j»arameter varies across its range of τιicertaiuty. either between the endpoints of a uiform cistribution or )elwee he confience limits of a ωςο distribution. | Table \ref{table3} presents the slope of the linear dependence for the most significant stellar evolution parameters as well as the total change in bump magnitude as the parameter varies across its range of uncertainty, either between the endpoints of a uniform distribution or between the confidence limits of a Gaussian distribution. |
The impact of srface boundary. cotditious is explored by compajug the Monte Carlo distribution of ulnup nag[n]utuces obtained using the Eddiugtou P(r relation to that using the Wrishua-Swamy(1966) ‘elation. aud Table 3 records the difference between tie mecdiau bump magniπο obtainec with eac of the two treatments. | The impact of surface boundary conditions is explored by comparing the Monte Carlo distribution of bump magnitudes obtained using the Eddington $T(\tau)$ relation to that using the \citet{krish} relation, and Table \ref{table3} records the difference between the median bump magnitudes obtained with each of the two treatments. |
The three most sigtificant sources of uncertainty iL the bump maguicde are the alpha-capt abtdance. the naixine length. aux the low-temperattο opacities. | The three most significant sources of uncertainty in the bump magnitude are the alpha-capture abundance, the mixing length, and the low-temperature opacities. |
All three of these parameters have a strouger intluence at higher inetallic‘ities. which explains why the overall uice‘tainty in the bump magiude increases with meallicity. | All three of these parameters have a stronger influence at higher metallicities, which explains why the overall uncertainty in the bump magnitude increases with metallicity. |
Other sigulicaut sources of uncertainty are the surface boundary coiditious. convective overshoot. the PC4j)—HN+ reaction rate. aid the helium diffusion coeficients. The impact o “these )araiueters ¢oes not vary signilicautly wit1 [Fe/H]. and so the values give1 in Table: are averaged across a| five metallicities. | Other significant sources of uncertainty are the surface boundary conditions, convective overshoot, the $\mathrm{^{13}C} + p \rightarrow \mathrm{^{14}N} + \gamma$ reaction rate, and the helium diffusion coefficients, The impact of these parameters does not vary significantly with [Fe/H], and so the values given in Table \ref{table3} are averaged across all five metallicities. |
All of the other stellar evolutiou pa‘ameters explored iu tre Mone Carlo simulation are uegligible. impactug the bump umagulitude a the evel of 0.05 iuag or less. | All of the other stellar evolution parameters explored in the Monte Carlo simulation are negligible, impacting the bump magnitude at the level of 0.05 mag or less. |
Cassisi&Salaris(1997) aud Cassisi.«eelIuuocenti&Salaris(1997) have alreacy studied the üunpact of a ew of these stellar evolution parameters ou theoretical bump maguituces. | \citet{cas3} and \citet{cas1}
have already studied the impact of a few of these stellar evolution parameters on theoretical bump magnitudes. |
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