CoolWiki - User contributions [en]

File:Ourdraft.ppt

2007-12-16T02:21:00Z

Weehler:

Talk:Working area for AAS jan 2008 poster

2007-12-16T02:19:46Z

Weehler: wiki poster contribution "Characteristics of Young Stars"

== Front Page Information ==
--[[User:Weehler|Weehler]] 18:13, 15 December 2007 (PST)Hi, this is what I've been preparing for the poster. Tell me if I need more detail, less detail, or edit how you see fit--Cindy [[Media:ourdraft.ppt]]
--[[User:Danielle|DanielleYeager]] 07:50, 5 December 2007 (PST) As of now, I have a very short and simple statement on the front page of this area of the wiki. If you don't know what I mean, I mean the article section for 'Working area for AAS jan 2008 poster'... Who all is going to be presenting it? What exactly is the abstract? What information to we currently have that I can post out on the main/article page. [end]

--[[User:Rebull|Rebull]] 14:25, 5 December 2007 (PST) Hi Danielle - I just went ahead and updated it. [end]

--[[User:Danielle|DanielleYeager]] 07:24, 6 December 2007 (PST) Thank you very much! [end]

== Highlights for Wiki AAS Poster ==

--[[User:Danielle|Dani]] 07:19, 5 December 2007 (PST) this is just copied and pasted over from the IC 2118 thread.

....DANIELLEYEAGER 12/03/07 1036 est....

Newest part of the project = Find 'interesting' or 'important' parts of the wiki to highlight on poster. You can post parts you would like to use in this thread. ''[If you don't remember how to post on a particular thread, you click edit to the left of the heading of the thread you want to post. If you want to start a new thread, scroll to the top and click on the + (plus) sign beside the tab for edit. And as always, if you have questions on the wiki, type my name in caps and then your question.]''

Example / Rough Draft of poster.

[[Image:Slide1.JPG]]

[END OF INTRODUCTION, POST BELOW.]
___________________________________________________________________________

.....DANIELLEYEAGER 12/03/07 1253 EST...

*Basics on the fact that you are working on real research, and at the same time it is an informational, so if you have questions, nine times out of ten, the information is right there for you to get and work with. [Main page.]

*A lot of time we work with IR and if someone didn’t know what IR is, we have a section specifically noted for that. [Main Page, section 3]

*Future research project, Galaxy Classifications.

*Main page section 8.

*Users Guide on the bottom of the Main Page.

[END OF DANIELLEYEAGER POST.]

Nick Kelley 12/03/07 1342 est

*Past T-Tauri work.

*Things the wiki allows use to do better.

*Interaction with other students

*Allows real Scientific work to be done by students.

--[[User:Spuck|Spuck]] 11:55, 4 December 2007 (PST) These are great ideas ... keep them coming. Also we have a knew location in the WIKI for posting ideas for the 2008 AAS Wiki poster ideas and work items. SO, please post at the [https://coolwiki.ipac.caltech.edu/index.php/Working_area_for_AAS_jan_2008_poster AAS Wiki Poster Page] in the future. Thanks.

--[[User:Danielle|DanielleYeager]] 10:48, 5 December 2007 (PST)
*Finding cluster members

*The playground

*How can I find already-reduced Spitzer data?

--[[User:Nicholas|NicholasJamesKelley]] 10:40, 6 December 2007 (PST) It would also be very nice if other people contributed to Ideas.

--[[User:Danielle|DanielleYeager]] 10:06, 7 December 2007 (PST) Yes, Nick, it would be nice if other people contributed, but until they do, it's important that we continue to research and post. The most important feature is that a non-user[meaning, someone without a user name and password to the page] can get more information from the site than they more than likely know what to do with and the fact that the authors are talking 'behind the scenes', what we know as the discusion sections, and the site evolves and shifts as the projects progress.

== Highlights for Galaxy Poster ==
Researchers and presenters for the galaxy project include Cale McClintock, Jennifer Butchart, Alexis McCool and Alix Holcomb.
The poster will be presented at the upcoming AAS meeting in Austin Texas.

== Highlights for WZ Sge Poster ==

WZSge is an eclipsing binary star system composed of a white dwarf(serving as the primary star) and a brown dwarf(which serves as a secondary star). This is a close orbital binary star system with only an 82 minute orbital period. What is happening in WZSge is the white dwarf is stripping material away from the brown dwarf, however due to angular momentum the material is not falling directly into the white dwarf it is falling into an accretion disk which surrounds the white dwarf. This accrection disk is so bright that when WZSge is imaged in the optical we are not actually seeing the white dwarf but in fact the accretion disk which surrounds it. In the optical we can see by generating a light curve that the brown dwarf eclipses in front of the white dwarf and its accretion disk. However the current phenomena of WZSge is that an eclipse is visible in the IR (at both 4.5 and 8 microns).

Researchers for this project include:
Steve Howell NOAO WIYN, Don Hoard SSC Harvey Mudd, J.M. Santiago and Jeff Adkins Deer Valley High School, Kimmerly Johnson and Beth Thomas, Matt Walentosky and Tim Spuck Oil City High School.

The poster "WZSge: Dark Matter in Accretion Disks" will be presented at the 211th annual meeting of the American Astronomical Society in Austin Texas this coming January.
--[[User:Weehler|Weehler]] 18:19, 15 December 2007 (PST)

Talk:Working area for AAS jan 2008 poster

2007-12-16T02:13:48Z

Weehler: contribution to the wiki "Characteristics of Young Stars"

== Front Page Information ==
--[[User:Weehler|Weehler]] 18:13, 15 December 2007 (PST)Hi, this is what I've been preparing for the poster. Tell me if I need more detail, less detail, or edit how you see fit--Cindy
--[[User:Danielle|DanielleYeager]] 07:50, 5 December 2007 (PST) As of now, I have a very short and simple statement on the front page of this area of the wiki. If you don't know what I mean, I mean the article section for 'Working area for AAS jan 2008 poster'... Who all is going to be presenting it? What exactly is the abstract? What information to we currently have that I can post out on the main/article page. [end]

--[[User:Rebull|Rebull]] 14:25, 5 December 2007 (PST) Hi Danielle - I just went ahead and updated it. [end]

--[[User:Danielle|DanielleYeager]] 07:24, 6 December 2007 (PST) Thank you very much! [end]

== Highlights for Wiki AAS Poster ==

--[[User:Danielle|Dani]] 07:19, 5 December 2007 (PST) this is just copied and pasted over from the IC 2118 thread.

....DANIELLEYEAGER 12/03/07 1036 est....

Newest part of the project = Find 'interesting' or 'important' parts of the wiki to highlight on poster. You can post parts you would like to use in this thread. ''[If you don't remember how to post on a particular thread, you click edit to the left of the heading of the thread you want to post. If you want to start a new thread, scroll to the top and click on the + (plus) sign beside the tab for edit. And as always, if you have questions on the wiki, type my name in caps and then your question.]''

Example / Rough Draft of poster.

[[Image:Slide1.JPG]]

[END OF INTRODUCTION, POST BELOW.]
___________________________________________________________________________

.....DANIELLEYEAGER 12/03/07 1253 EST...

*Basics on the fact that you are working on real research, and at the same time it is an informational, so if you have questions, nine times out of ten, the information is right there for you to get and work with. [Main page.]

*A lot of time we work with IR and if someone didn’t know what IR is, we have a section specifically noted for that. [Main Page, section 3]

*Future research project, Galaxy Classifications.

*Main page section 8.

*Users Guide on the bottom of the Main Page.

[END OF DANIELLEYEAGER POST.]

Nick Kelley 12/03/07 1342 est

*Past T-Tauri work.

*Things the wiki allows use to do better.

*Interaction with other students

*Allows real Scientific work to be done by students.

--[[User:Spuck|Spuck]] 11:55, 4 December 2007 (PST) These are great ideas ... keep them coming. Also we have a knew location in the WIKI for posting ideas for the 2008 AAS Wiki poster ideas and work items. SO, please post at the [https://coolwiki.ipac.caltech.edu/index.php/Working_area_for_AAS_jan_2008_poster AAS Wiki Poster Page] in the future. Thanks.

--[[User:Danielle|DanielleYeager]] 10:48, 5 December 2007 (PST)
*Finding cluster members

*The playground

*How can I find already-reduced Spitzer data?

--[[User:Nicholas|NicholasJamesKelley]] 10:40, 6 December 2007 (PST) It would also be very nice if other people contributed to Ideas.

--[[User:Danielle|DanielleYeager]] 10:06, 7 December 2007 (PST) Yes, Nick, it would be nice if other people contributed, but until they do, it's important that we continue to research and post. The most important feature is that a non-user[meaning, someone without a user name and password to the page] can get more information from the site than they more than likely know what to do with and the fact that the authors are talking 'behind the scenes', what we know as the discusion sections, and the site evolves and shifts as the projects progress.

== Highlights for Galaxy Poster ==
Researchers and presenters for the galaxy project include Cale McClintock, Jennifer Butchart, Alexis McCool and Alix Holcomb.
The poster will be presented at the upcoming AAS meeting in Austin Texas.

== Highlights for WZ Sge Poster ==

WZSge is an eclipsing binary star system composed of a white dwarf(serving as the primary star) and a brown dwarf(which serves as a secondary star). This is a close orbital binary star system with only an 82 minute orbital period. What is happening in WZSge is the white dwarf is stripping material away from the brown dwarf, however due to angular momentum the material is not falling directly into the white dwarf it is falling into an accretion disk which surrounds the white dwarf. This accrection disk is so bright that when WZSge is imaged in the optical we are not actually seeing the white dwarf but in fact the accretion disk which surrounds it. In the optical we can see by generating a light curve that the brown dwarf eclipses in front of the white dwarf and its accretion disk. However the current phenomena of WZSge is that an eclipse is visible in the IR (at both 4.5 and 8 microns).

Researchers for this project include:
Steve Howell NOAO WIYN, Don Hoard SSC Harvey Mudd, J.M. Santiago and Jeff Adkins Deer Valley High School, Kimmerly Johnson and Beth Thomas, Matt Walentosky and Tim Spuck Oil City High School.

The poster "WZSge: Dark Matter in Accretion Disks" will be presented at the 211th annual meeting of the American Astronomical Society in Austin Texas this coming January.

Units

2007-11-04T18:46:52Z

Weehler: /* Units of Spitzer Images */

=General Units=
Wavelengths in infrared astronomy are commonly expressed in microns = micrometers = µm (or um if you don't have a µ).

*5000 Å =500 nm =0.5 µm =Visible light
*~0.9 to 5 µm =Near-infrared (~smoke particles)
*5 µm to ~30 µm = Mid-infrared (~hair)
*30 µm to ~350 µm = Far-infrared (~salt grain)

Brightnesses or fluxes are most likely to be given in Janskys (Jy) or mJy (milli Jy) or µJy (micro Jy). 1 Jansky = <math>10^{-26}</math> Watts/m^2/Hz.

Jy can be converted to magnitudes which have historically been relatively rarely used in the mid- or far-infrared.

Because the unit is named for [http://en.wikipedia.org/wiki/Karl_Jansky Karl Jansky], the plural of the unit is really Janskys, not Janskies.

==Aside on fluxes and flux densities==

Astronomically, it can be important to understand the difference between luminosity, flux, and flux density. In practice for this stuff, you probably don't need to know the gritty details of this until you are more familiar with the numbers and the jargon.

Colloquially, flux means the rate of something through something else, such as water through a pipe, or traffic on a highway. In physics and astronomy, it means the same thing.

''Flux'' is a measurement of ''energy per unit area per unit time.'' Using our analogies above, this would be the number of cars per lane per second that pass under a bridge on a highway (or grams of water through the cross-sectional area of the pipe per second). In measuring energy from celestial objects, the units of flux are Joules per second per meter squared if you like mks (meters-kilograms-seconds) units, or ergs per second per centimeter squared if you like cgs (centimeters-grams-seconds) units.

''Luminosity'' is a measurement of ''energy per unit of time,'' such as Joules per second if you like mks units, or ergs per second if you like cgs units. This would be, in our analogy, the total number of cars on the highway passing under the bridge per second. (The flux of cars is the luminosity per lane.)

''Flux density'' is a measurement essentially of ''energy per unit area per unit time "per photon".'' In our analogy, this would be the number of RED cars per lane per second that pass under the bridge on the highway. In this analogy, the "per photon" is seen in the red cars. In astronomy, the "per photon" manifests itself as a "per Hz" (unit of frequency) or "per cm" (unit of wavelength). A Jansky is proportional to Watts/m^2/Hz. Recall that Watts are energy per second. So this is energy per second per square meter per Hertz.

Now, just to further confuse things, the units of Spitzer ''mosaics'' are not just Janskys, but Janskys per pixel! To make the numbers easier, they are in MJy/sr, but they could also be in uJy/square arcsecond. Read on for more, including definitions and scale factors!

=Units of Spitzer Images=

Optical data with which you are familiar may be in counts or photons, or possibly (like Hubble data) calibrated to be energies. That, combined with the exposure time of the image, gives you ''flux units''. Spitzer data comes in ''flux (density) per unit (pixel) area'' instead, MegaJanskys per steradian (MJy/sr). 1 MJy = <math>10^{6}</math> Jy, and a sr is a solid angle.

If you've done photometry before, and expect to do it exactly the same way again here, '''it won't work''', because '''this matters'''.

1 square arcsec is <math>2.3504 \times 10^{-11}</math> sr. (1 degree = 60 arcmin = 3600 arcsec.)

If you want to convert the image from MJy/sr to uJy/square arcsec, multiply the image by 23.5045. The units of this number are (uJy/arcsec)/(MJy/sr).

If you want to take a Spitzer image and use your previous routines on it, the most efficient way to do this is probably to take the image in MJy/sr and multiply out the "per sr" part of it so that it is instead in MJy/px. The subtlety in this step is that each Spitzer array has slightly different pixel sizes, and the mosaics that we create have different sizes yet again from the original images. You can make mosaics with whatever size pixels you want, so if you get Spitzer mosaics from more than one astronomer, or more than one Spitzer wavelength, chances are excellent that the pixels will be slightly different sizes. The information on the pixel sizes are in the [[FITS_format|FITS]] header of each image.

The following paragraphs are a high-level summary of what to do for any Spitzer image data you may encounter; see below for a cookbook of the process for one mosaic.

Look in the FITS header of the mosaics for the keywords "CDELT1" and "CDELT2". These keywords are set to be the scale of the rows and columns in degrees per pixel. Using the values of these keywords, and the conversions above, you can figure out the number of square degrees per pixel, the number of square arcsec per pixel, and finally the number of steradians per pixel. Multiply the whole image in MJy/sr by the number of sr/px to get MJy/px.

If you are instead working with the individual BCDs (read this as: the individual little images that went into the big mosaic), you should look for keywords "PXSCAL1" and "PXSCAL2". NOTE that these pixels ARE NOT SQUARE, and this is more important for MIPS data. From here, you now have the same information as the "CDELT1" and "CDELT2" above, so you can follow the same procedure.

=Units of Spitzer Photometry=

==Introduction==

The photometry software that people use at the SSC, called APEX, produces fluxes in microJanskys. The final bandmerged catalog you can get has listed fluxes in microJanskys, as well as magnitudes.

Astronomers use magnitudes in color-color or color-magnitude plots. Astronomers use a variant on fluxes in spectral energy distribution (SED) plots.

==Magnitudes==

A magnitude is really a flux ratio. It is defined as follows, where M's are magnitudes and F's are fluxes:
<math>M_1 - M_2 = 2.5 \times \log \left(\frac{F_2}{F_1}\right)</math> (eqn 1)

The magnitude system (in the optical) was defined to be referenced to Vega. In other words, Vega is defined to be zero magnitude, and you would then define magnitudes of anything else as follows:
<math>M = 2.5 \times \log \left(\frac{F_{\mathrm{Vega}}}{F}\right)</math> (eqn 2)

When they looked at Vega with IRAS, they discovered that it did NOT look like they expected, and in fact it has a large infrared excess! Therefore, infrared magnitudes are defined with respect to what Vega would be, if it did not have an excess.

We have published the zero points (e.g., the "Vega flux") for most of our bandpasses. They are (copied from various places on the web):
*IRAC 1 : 280.9 Jy
*IRAC 2 : 179.7 Jy
*IRAC 3 : 115.0 Jy
*IRAC 4 : 64.13 Jy
*MIPS 1 : 7.14 Jy
*MIPS 2 : 0.775 Jy
*MIPS 3 : 0.159 Jy

Therefore, in order to convert the uJy that apex returns into magnitudes, use the equation 2 above, substituting these so-called "zero-point fluxes" in for "Fvega." Note that the zero-point fluxes are in Janskys and the fluxes returned by APEX are in microJanskys.

You can find the zeropoints for 2MASS magnitudes on the web as well:
*J : 1594 Jy
*H : 1024 Jy
*K : 666.7 Jy

Note that plain magnitudes get fainter (the number gets larger) as the distance of the object increases. BUT, colors (differences in magnitudes) are ratios of fluxes, and therefore independent of distance.

==Spectral Energy Distributions (SEDs)==

SEDs are energy plotted against some measure of the photon -- frequency or wavelength. The reason astronomers do this is to see how much energy is produced by the object as a function of frequency or wavelength.
Now it's really going to get a little hairy! Steel your nerves and plunge onwards... it really all comes down to unit conversion.

1Jy = <math>10^{-23}</math> erg/s/cm^2/Hz (in cgs units rather than mks units, sorry). A Jansky is technically a unit of "flux density." In order to get rid of the "per Hz", you need to multiply the Jy by the frequency of the bandpass center.

Astronomers coming from the longer wavelengths will tend to plot up nu * F(nu) (written as <math>\nu F_{\nu}</math>) against nu, where "nu" (<math>\nu</math>) is the frequency. The units of <math>F_{\nu}</math> are Janskys.

Astronomers coming from the shorter wavelengths will tend to plot up lambda * F(lambda) (written as <math>\lambda F_\lambda</math>), where "lambda" (<math>\lambda</math>) is the wavelength of the light. The units of <math>F_{\lambda}</math> are NOT Janskys.

<math>\lambda \times \nu = c</math>, the speed of light. In order to convert the Janskys into units of <math>F_{\lambda}</math>, you need to take into account the differentials (ah-HA, calculus being used here!), e.g., the fact that
<math>\frac{dF}{d\lambda} = \frac{dF}{d\nu} \frac{d\nu}{d\lambda}</math> and <math>d\nu = \frac{c}{\lambda^2}d\lambda</math>
So you need to multiply the <math>F_{\nu}</math> by <math>c/\lambda^2</math> to convert it into <math>F_{\lambda}</math>.

Additionally, to analyze the Spitzer data, it's often useful to pretend that the contribution from the star is a blackbody. It's not really, but it's awful close, especially in the infrared.

A blackbody's flux density is given by (where T is temperature, and other constants are given below)
<math>B_{\lambda} = \left(\frac{2hc^2/\lambda^5}{\exp(hc/\lambda kT)-1)}\right)</math> (eqn 3)
but of course we want to plot <math>\lambda \times B_{\lambda}</math>:
<math>\lambda B_{\lambda} = \left(\frac{2hc^2/\lambda^4}{\exp(hc/\lambda kT)-1)}\right)</math> (eqn 4)
Values of these constants all in cgs units:
*h = 6.6260755d-27 erg*sec
*c = 2.997924d10 cm/sec
*k = 1.380658d-16 erg/deg

In words, in order to analyze our data, we need to have something that does the following:

#Reads in the fluxes from the files.
#Converts the Spitzer fluxes (and errors) into magnitudes (if necessary).
#Converts the 2MASS magnitudes (and errors) into fluxes (if necessary).
#Makes color-color and color-magnitude plots for stars in our region using magnitudes.
#Makes SED plots for individual objects, but converting numbers first into the right units:
##Creates an array of the wavelengths of each measurement, keeps a copy of the version in microns, and converts to cm.
##For any real measurements, converts the microJanskys into cgs units.
##For any real measurements, converts <math>F_{\nu}</math> into <math>F_{\lambda}</math> by multiplying the <math>F_{\nu}</math> values by the <math>d\nu/d\lambda</math> corresponding to the wavelength of each bandpass.
##For any real measurements, multiplies <math>F_{\lambda}</math> by the lambda corresponding to the wavelength of each bandpass to get <math>\lambda F_{\lambda}</math>.
#For any real measurements, plots the log of the <math>\lambda F_{\lambda}</math> data points (in cgs units) against the log of the lambda data points (in microns, only because that makes it easier to read). Labels the axes (with units)! Plots the error bars on top of the data points (also converted from uJy).
#For any real measurements, for any star with at least 2 fluxes, fits a blackbody to the energies derived from the three 2MASS and first 2 IRAC bands. There are two free parameters in this fit -- the temperature of the blackbody and an additive (in the log) offset related to the distance of the object. If we know the temperature of the star (via a spectral type) and the distance to the object, then we know the values for the temperature and the offset.

Why are we plotting <math>\lambda F_{\lambda}</math> vs. <math>\lambda</math>? Well, only because I think in wavelength, not frequency. I don't know off the top of my head the frequencies of the Spitzer bandpasses, but I do know their wavelengths. Why are we plotting <math>\lambda F_{\lambda}</math> instead of <math>\nu F_{\nu}</math>? Well, only for internal consistency. Since one axis is in wavelength units, it makes sense to have the other axis also in wavelength units.

=Cookbook for image conversion: Method One=

==Step zero. What do you have and what do you need?==
You have an image in MJy/sr(/px). You have the number of degrees per pixel.

You need to convert the image to MJy(/px), a.k.a "get rid of the steradians."

The things that make this hard are:
*The pixel size of the mosaic changes depending on wavelength and where you got the mosaic, so I can't just give you one number to work for all mosaics every time.
*The "pixels" in the above are kind of a funny, hidden unit and the accounting of it works in some unexpected ways, which is why it's in parentheses above (and below).

==Step one. Find out what the size of the pixels are in your images==
For the IRAC-1 mosaic I created for you in July 2006, CDELT1=-0.000339 degrees per pixel, and CDELT2=0.000339 degrees per pixel. (ignore the minus sign; it has something to do with a fits convention.) You can find this out for any given mosaic by looking in the fits header.

==Step two. Find out what the size of the pixels are in square degrees per pixel==

degrees degrees square degrees
0.000339 ------- * 0.000339 ------- = 1.14921e-7 ---------------
pixel pixel (square) pixel

==Step three. Find out what the conversion is between square degrees and sr.==
There are 60 arcminutes in a degree. There are 60 arcseconds in an arcminute.

60 arcminutes 60 arcseconds 3600 arcseconds
------------- * -------------- = ---------------
1 degree 1 arcminutes 1 degree

Square it!

(1 degree)^2 = 1 square degree = (3600 arcsec)^2 = 1.296e7 square arcsec

We look up that 1 square arcsec is 2.3504x10^(-11) sr.

1.297e7 square arcsec 2.3504e-11 sr sr
------------- * ---------------- = 0.000304847 ------------
square degree 1 square arcsec square degree

==Step four. Find out what the size of the pixels are in sr.==

square degrees sr sr
1.14921e-7 --------------- * 0.000304847 ------------ = 3.50333e-11 ----
(square) pixel square degree px

==Step five. Convert the units of the image.==

MJy sr MJy
--------- * 3.50333e-11 ---- = ---
sr px px

So multiply this whole image by 3.50333e-11. The units of the image (and consequently the photometry you get out) are in MJy (MegaJanskys). If you want to get it in Janskys:

1e6 Jy
----
MJy

so multiply the image by 1e6 to get the image into Jy:

MJy 1e6 Jy Jy
--- * ---- = ----
px MJy px

If you want to get it into microJy, there are 1e6 uJy in a Jy, and I'll let you do that one.

=Method 2: Alternative but completely equivalent and possibly more straightforward solution=
==Step zero. What do you have and what do you need?==
You have an image in MJy/sr(/px). You have the number of degrees per pixel.

You need to convert the image to uJy(/px), a.k.a "get rid of the steradians" AND convert to microJanskys to get the numbers to still be reasonable and not very tiny or very large.

The things that make this hard are:
* The pixel size of the mosaic changes depending on wavelength and where you got the mosaic, so I can't just give you one number to work for all mosaics every time.
* The "pixels" in the above are kind of a funny, hidden unit and the accounting of it works in some unexpected ways.

==Step one. Find out what the size of the pixels are in your images==
For the IRAC-1 mosaic I created for you in July 2006, CDELT1=-0.000339 degrees per pixel, and CDELT2=0.000339 degrees per pixel. (ignore the minus sign; it has something to do with a fits convention.) For any given mosaic, you can find out these values by looking in the fits header.

==Step two. Convert your image from MJy/sr to uJy/square arcsec==
We look up that there are :

[uJy/sq. arcsec]
23.5045 -----------------
[MJy/sr]

So multiply the image by 23.5045 to get it into uJy/square arcsec
==Step three. Find out what the size of the pixels are in square degrees per pixel==

degrees degrees square degrees
0.000339 ------- * 0.000339 ------- = 1.14921e-7 ---------------
pixel pixel (square) pixel

==Step four. Find out how many square arcsec there are in a pixel.==
There are 60 arcminutes in a degree. There are 60 arcseconds in an arcminute.

60 arcminutes 60 arcseconds 3600 arcseconds
------------- * -------------- = ---------------
1 degree 1 arcminutes 1 degree

Square it!

(1 degree)^2 = 1 square degree = (3600 arcsec)^2 = 1.296e7 square arcsec

Convert the pixel size.

square degrees square arcsec square arcsec
1.14921e-7 --------------- * 1.296e7 -------------- = 1.48938 ---------------
(square) pixel square degree (square) px

==Step five. Convert the image==

uJy 1.48938 square arcsec uJy
----------- * --------------- = ------
sq arcsec (*px) (square) px (px)

So multiply the image by 1.48938 to get it into uJy/px.

Units

2007-11-04T18:36:53Z

Weehler: /* Units of Spitzer Images */

=General Units=
Wavelengths in infrared astronomy are commonly expressed in microns = micrometers = µm (or um if you don't have a µ).

*5000 Å =500 nm =0.5 µm =Visible light
*~0.9 to 5 µm =Near-infrared (~smoke particles)
*5 µm to ~30 µm = Mid-infrared (~hair)
*30 µm to ~350 µm = Far-infrared (~salt grain)

Brightnesses or fluxes are most likely to be given in Janskys (Jy) or mJy (milli Jy) or µJy (micro Jy). 1 Jansky = <math>10^{-26}</math> Watts/m^2/Hz.

Jy can be converted to magnitudes which have historically been relatively rarely used in the mid- or far-infrared.

Because the unit is named for [http://en.wikipedia.org/wiki/Karl_Jansky Karl Jansky], the plural of the unit is really Janskys, not Janskies.

==Aside on fluxes and flux densities==

Astronomically, it can be important to understand the difference between luminosity, flux, and flux density. In practice for this stuff, you probably don't need to know the gritty details of this until you are more familiar with the numbers and the jargon.

Colloquially, flux means the rate of something through something else, such as water through a pipe, or traffic on a highway. In physics and astronomy, it means the same thing.

''Flux'' is a measurement of ''energy per unit area per unit time.'' Using our analogies above, this would be the number of cars per lane per second that pass under a bridge on a highway (or grams of water through the cross-sectional area of the pipe per second). In measuring energy from celestial objects, the units of flux are Joules per second per meter squared if you like mks (meters-kilograms-seconds) units, or ergs per second per centimeter squared if you like cgs (centimeters-grams-seconds) units.

''Luminosity'' is a measurement of ''energy per unit of time,'' such as Joules per second if you like mks units, or ergs per second if you like cgs units. This would be, in our analogy, the total number of cars on the highway passing under the bridge per second. (The flux of cars is the luminosity per lane.)

''Flux density'' is a measurement essentially of ''energy per unit area per unit time "per photon".'' In our analogy, this would be the number of RED cars per lane per second that pass under the bridge on the highway. In this analogy, the "per photon" is seen in the red cars. In astronomy, the "per photon" manifests itself as a "per Hz" (unit of frequency) or "per cm" (unit of wavelength). A Jansky is proportional to Watts/m^2/Hz. Recall that Watts are energy per second. So this is energy per second per square meter per Hertz.

Now, just to further confuse things, the units of Spitzer ''mosaics'' are not just Janskys, but Janskys per pixel! To make the numbers easier, they are in MJy/sr, but they could also be in uJy/square arcsecond. Read on for more, including definitions and scale factors!

=Units of Spitzer Images=

Optical data with which you are familiar may be in counts or photons, or possibly (like Hubble data) calibrated to be energies. That, combined with the exposure time of the image, gives you ''flux units''. Spitzer data comes in ''flux (density) per unit (pixel) area'' instead, MegaJanskys per steradian (MJy/sr). 1 MJy = <math>10^{6}</math> Jy, and a sr is a solid angle.

If you've done photometry before, and expect to do it exactly the same way again here, '''it won't work''', because '''this matters'''.

1 square arcsec is <math>2.3504 \times 10^{-11}</math> sr. (1 degree = 60 arcmin = 3600 arcsec.)
'''okay, I'm confused: what is a square arc sec? How do you square an angle measurement?'''
If you want to convert the image from MJy/sr to uJy/square arcsec, multiply the image by 23.5045. The units of this number are (uJy/arcsec)/(MJy/sr).

If you want to take a Spitzer image and use your previous routines on it, the most efficient way to do this is probably to take the image in MJy/sr and multiply out the "per sr" part of it so that it is instead in MJy/px. The subtlety in this step is that each Spitzer array has slightly different pixel sizes, and the mosaics that we create have different sizes yet again from the original images. You can make mosaics with whatever size pixels you want, so if you get Spitzer mosaics from more than one astronomer, or more than one Spitzer wavelength, chances are excellent that the pixels will be slightly different sizes. The information on the pixel sizes are in the [[FITS_format|FITS]] header of each image.

The following paragraphs are a high-level summary of what to do for any Spitzer image data you may encounter; see below for a cookbook of the process for one mosaic.

Look in the FITS header of the mosaics for the keywords "CDELT1" and "CDELT2". These keywords are set to be the scale of the rows and columns in degrees per pixel. Using the values of these keywords, and the conversions above, you can figure out the number of square degrees per pixel, the number of square arcsec per pixel, and finally the number of steradians per pixel. Multiply the whole image in MJy/sr by the number of sr/px to get MJy/px.

If you are instead working with the individual BCDs (read this as: the individual little images that went into the big mosaic), you should look for keywords "PXSCAL1" and "PXSCAL2". NOTE that these pixels ARE NOT SQUARE, and this is more important for MIPS data. From here, you now have the same information as the "CDELT1" and "CDELT2" above, so you can follow the same procedure.

=Units of Spitzer Photometry=

==Introduction==

The photometry software that people use at the SSC, called APEX, produces fluxes in microJanskys. The final bandmerged catalog you can get has listed fluxes in microJanskys, as well as magnitudes.

Astronomers use magnitudes in color-color or color-magnitude plots. Astronomers use a variant on fluxes in spectral energy distribution (SED) plots.

==Magnitudes==

A magnitude is really a flux ratio. It is defined as follows, where M's are magnitudes and F's are fluxes:
<math>M_1 - M_2 = 2.5 \times \log \left(\frac{F_2}{F_1}\right)</math> (eqn 1)

The magnitude system (in the optical) was defined to be referenced to Vega. In other words, Vega is defined to be zero magnitude, and you would then define magnitudes of anything else as follows:
<math>M = 2.5 \times \log \left(\frac{F_{\mathrm{Vega}}}{F}\right)</math> (eqn 2)

When they looked at Vega with IRAS, they discovered that it did NOT look like they expected, and in fact it has a large infrared excess! Therefore, infrared magnitudes are defined with respect to what Vega would be, if it did not have an excess.

We have published the zero points (e.g., the "Vega flux") for most of our bandpasses. They are (copied from various places on the web):
*IRAC 1 : 280.9 Jy
*IRAC 2 : 179.7 Jy
*IRAC 3 : 115.0 Jy
*IRAC 4 : 64.13 Jy
*MIPS 1 : 7.14 Jy
*MIPS 2 : 0.775 Jy
*MIPS 3 : 0.159 Jy

Therefore, in order to convert the uJy that apex returns into magnitudes, use the equation 2 above, substituting these so-called "zero-point fluxes" in for "Fvega." Note that the zero-point fluxes are in Janskys and the fluxes returned by APEX are in microJanskys.

You can find the zeropoints for 2MASS magnitudes on the web as well:
*J : 1594 Jy
*H : 1024 Jy
*K : 666.7 Jy

Note that plain magnitudes get fainter (the number gets larger) as the distance of the object increases. BUT, colors (differences in magnitudes) are ratios of fluxes, and therefore independent of distance.

==Spectral Energy Distributions (SEDs)==

SEDs are energy plotted against some measure of the photon -- frequency or wavelength. The reason astronomers do this is to see how much energy is produced by the object as a function of frequency or wavelength.
Now it's really going to get a little hairy! Steel your nerves and plunge onwards... it really all comes down to unit conversion.

1Jy = <math>10^{-23}</math> erg/s/cm^2/Hz (in cgs units rather than mks units, sorry). A Jansky is technically a unit of "flux density." In order to get rid of the "per Hz", you need to multiply the Jy by the frequency of the bandpass center.

Astronomers coming from the longer wavelengths will tend to plot up nu * F(nu) (written as <math>\nu F_{\nu}</math>) against nu, where "nu" (<math>\nu</math>) is the frequency. The units of <math>F_{\nu}</math> are Janskys.

Astronomers coming from the shorter wavelengths will tend to plot up lambda * F(lambda) (written as <math>\lambda F_\lambda</math>), where "lambda" (<math>\lambda</math>) is the wavelength of the light. The units of <math>F_{\lambda}</math> are NOT Janskys.

<math>\lambda \times \nu = c</math>, the speed of light. In order to convert the Janskys into units of <math>F_{\lambda}</math>, you need to take into account the differentials (ah-HA, calculus being used here!), e.g., the fact that
<math>\frac{dF}{d\lambda} = \frac{dF}{d\nu} \frac{d\nu}{d\lambda}</math> and <math>d\nu = \frac{c}{\lambda^2}d\lambda</math>
So you need to multiply the <math>F_{\nu}</math> by <math>c/\lambda^2</math> to convert it into <math>F_{\lambda}</math>.

Additionally, to analyze the Spitzer data, it's often useful to pretend that the contribution from the star is a blackbody. It's not really, but it's awful close, especially in the infrared.

A blackbody's flux density is given by (where T is temperature, and other constants are given below)
<math>B_{\lambda} = \left(\frac{2hc^2/\lambda^5}{\exp(hc/\lambda kT)-1)}\right)</math> (eqn 3)
but of course we want to plot <math>\lambda \times B_{\lambda}</math>:
<math>\lambda B_{\lambda} = \left(\frac{2hc^2/\lambda^4}{\exp(hc/\lambda kT)-1)}\right)</math> (eqn 4)
Values of these constants all in cgs units:
*h = 6.6260755d-27 erg*sec
*c = 2.997924d10 cm/sec
*k = 1.380658d-16 erg/deg

In words, in order to analyze our data, we need to have something that does the following:

#Reads in the fluxes from the files.
#Converts the Spitzer fluxes (and errors) into magnitudes (if necessary).
#Converts the 2MASS magnitudes (and errors) into fluxes (if necessary).
#Makes color-color and color-magnitude plots for stars in our region using magnitudes.
#Makes SED plots for individual objects, but converting numbers first into the right units:
##Creates an array of the wavelengths of each measurement, keeps a copy of the version in microns, and converts to cm.
##For any real measurements, converts the microJanskys into cgs units.
##For any real measurements, converts <math>F_{\nu}</math> into <math>F_{\lambda}</math> by multiplying the <math>F_{\nu}</math> values by the <math>d\nu/d\lambda</math> corresponding to the wavelength of each bandpass.
##For any real measurements, multiplies <math>F_{\lambda}</math> by the lambda corresponding to the wavelength of each bandpass to get <math>\lambda F_{\lambda}</math>.
#For any real measurements, plots the log of the <math>\lambda F_{\lambda}</math> data points (in cgs units) against the log of the lambda data points (in microns, only because that makes it easier to read). Labels the axes (with units)! Plots the error bars on top of the data points (also converted from uJy).
#For any real measurements, for any star with at least 2 fluxes, fits a blackbody to the energies derived from the three 2MASS and first 2 IRAC bands. There are two free parameters in this fit -- the temperature of the blackbody and an additive (in the log) offset related to the distance of the object. If we know the temperature of the star (via a spectral type) and the distance to the object, then we know the values for the temperature and the offset.

Why are we plotting <math>\lambda F_{\lambda}</math> vs. <math>\lambda</math>? Well, only because I think in wavelength, not frequency. I don't know off the top of my head the frequencies of the Spitzer bandpasses, but I do know their wavelengths. Why are we plotting <math>\lambda F_{\lambda}</math> instead of <math>\nu F_{\nu}</math>? Well, only for internal consistency. Since one axis is in wavelength units, it makes sense to have the other axis also in wavelength units.

=Cookbook for image conversion: Method One=

==Step zero. What do you have and what do you need?==
You have an image in MJy/sr(/px). You have the number of degrees per pixel.

You need to convert the image to MJy(/px), a.k.a "get rid of the steradians."

The things that make this hard are:
*The pixel size of the mosaic changes depending on wavelength and where you got the mosaic, so I can't just give you one number to work for all mosaics every time.
*The "pixels" in the above are kind of a funny, hidden unit and the accounting of it works in some unexpected ways, which is why it's in parentheses above (and below).

==Step one. Find out what the size of the pixels are in your images==
For the IRAC-1 mosaic I created for you in July 2006, CDELT1=-0.000339 degrees per pixel, and CDELT2=0.000339 degrees per pixel. (ignore the minus sign; it has something to do with a fits convention.) You can find this out for any given mosaic by looking in the fits header.

==Step two. Find out what the size of the pixels are in square degrees per pixel==

degrees degrees square degrees
0.000339 ------- * 0.000339 ------- = 1.14921e-7 ---------------
pixel pixel (square) pixel

==Step three. Find out what the conversion is between square degrees and sr.==
There are 60 arcminutes in a degree. There are 60 arcseconds in an arcminute.

60 arcminutes 60 arcseconds 3600 arcseconds
------------- * -------------- = ---------------
1 degree 1 arcminutes 1 degree

Square it!

(1 degree)^2 = 1 square degree = (3600 arcsec)^2 = 1.296e7 square arcsec

We look up that 1 square arcsec is 2.3504x10^(-11) sr.

1.297e7 square arcsec 2.3504e-11 sr sr
------------- * ---------------- = 0.000304847 ------------
square degree 1 square arcsec square degree

==Step four. Find out what the size of the pixels are in sr.==

square degrees sr sr
1.14921e-7 --------------- * 0.000304847 ------------ = 3.50333e-11 ----
(square) pixel square degree px

==Step five. Convert the units of the image.==

MJy sr MJy
--------- * 3.50333e-11 ---- = ---
sr px px

So multiply this whole image by 3.50333e-11. The units of the image (and consequently the photometry you get out) are in MJy (MegaJanskys). If you want to get it in Janskys:

1e6 Jy
----
MJy

so multiply the image by 1e6 to get the image into Jy:

MJy 1e6 Jy Jy
--- * ---- = ----
px MJy px

If you want to get it into microJy, there are 1e6 uJy in a Jy, and I'll let you do that one.

=Method 2: Alternative but completely equivalent and possibly more straightforward solution=
==Step zero. What do you have and what do you need?==
You have an image in MJy/sr(/px). You have the number of degrees per pixel.

You need to convert the image to uJy(/px), a.k.a "get rid of the steradians" AND convert to microJanskys to get the numbers to still be reasonable and not very tiny or very large.

The things that make this hard are:
* The pixel size of the mosaic changes depending on wavelength and where you got the mosaic, so I can't just give you one number to work for all mosaics every time.
* The "pixels" in the above are kind of a funny, hidden unit and the accounting of it works in some unexpected ways.

==Step one. Find out what the size of the pixels are in your images==
For the IRAC-1 mosaic I created for you in July 2006, CDELT1=-0.000339 degrees per pixel, and CDELT2=0.000339 degrees per pixel. (ignore the minus sign; it has something to do with a fits convention.) For any given mosaic, you can find out these values by looking in the fits header.

==Step two. Convert your image from MJy/sr to uJy/square arcsec==
We look up that there are :

[uJy/sq. arcsec]
23.5045 -----------------
[MJy/sr]

So multiply the image by 23.5045 to get it into uJy/square arcsec
==Step three. Find out what the size of the pixels are in square degrees per pixel==

degrees degrees square degrees
0.000339 ------- * 0.000339 ------- = 1.14921e-7 ---------------
pixel pixel (square) pixel

==Step four. Find out how many square arcsec there are in a pixel.==
There are 60 arcminutes in a degree. There are 60 arcseconds in an arcminute.

60 arcminutes 60 arcseconds 3600 arcseconds
------------- * -------------- = ---------------
1 degree 1 arcminutes 1 degree

Square it!

(1 degree)^2 = 1 square degree = (3600 arcsec)^2 = 1.296e7 square arcsec

Convert the pixel size.

square degrees square arcsec square arcsec
1.14921e-7 --------------- * 1.296e7 -------------- = 1.48938 ---------------
(square) pixel square degree (square) px

==Step five. Convert the image==

uJy 1.48938 square arcsec uJy
----------- * --------------- = ------
sq arcsec (*px) (square) px (px)

So multiply the image by 1.48938 to get it into uJy/px.

Units

2007-11-04T18:32:42Z

Weehler: /* Units of Spitzer Images */

=General Units=
Wavelengths in infrared astronomy are commonly expressed in microns = micrometers = µm (or um if you don't have a µ).

*5000 Å =500 nm =0.5 µm =Visible light
*~0.9 to 5 µm =Near-infrared (~smoke particles)
*5 µm to ~30 µm = Mid-infrared (~hair)
*30 µm to ~350 µm = Far-infrared (~salt grain)

Brightnesses or fluxes are most likely to be given in Janskys (Jy) or mJy (milli Jy) or µJy (micro Jy). 1 Jansky = <math>10^{-26}</math> Watts/m^2/Hz.

Jy can be converted to magnitudes which have historically been relatively rarely used in the mid- or far-infrared.

Because the unit is named for [http://en.wikipedia.org/wiki/Karl_Jansky Karl Jansky], the plural of the unit is really Janskys, not Janskies.

==Aside on fluxes and flux densities==

Astronomically, it can be important to understand the difference between luminosity, flux, and flux density. In practice for this stuff, you probably don't need to know the gritty details of this until you are more familiar with the numbers and the jargon.

Colloquially, flux means the rate of something through something else, such as water through a pipe, or traffic on a highway. In physics and astronomy, it means the same thing.

''Flux'' is a measurement of ''energy per unit area per unit time.'' Using our analogies above, this would be the number of cars per lane per second that pass under a bridge on a highway (or grams of water through the cross-sectional area of the pipe per second). In measuring energy from celestial objects, the units of flux are Joules per second per meter squared if you like mks (meters-kilograms-seconds) units, or ergs per second per centimeter squared if you like cgs (centimeters-grams-seconds) units.

''Luminosity'' is a measurement of ''energy per unit of time,'' such as Joules per second if you like mks units, or ergs per second if you like cgs units. This would be, in our analogy, the total number of cars on the highway passing under the bridge per second. (The flux of cars is the luminosity per lane.)

''Flux density'' is a measurement essentially of ''energy per unit area per unit time "per photon".'' In our analogy, this would be the number of RED cars per lane per second that pass under the bridge on the highway. In this analogy, the "per photon" is seen in the red cars. In astronomy, the "per photon" manifests itself as a "per Hz" (unit of frequency) or "per cm" (unit of wavelength). A Jansky is proportional to Watts/m^2/Hz. Recall that Watts are energy per second. So this is energy per second per square meter per Hertz.

Now, just to further confuse things, the units of Spitzer ''mosaics'' are not just Janskys, but Janskys per pixel! To make the numbers easier, they are in MJy/sr, but they could also be in uJy/square arcsecond. Read on for more, including definitions and scale factors!

=Units of Spitzer Images=

Optical data with which you are familiar may be in counts or photons, or possibly (like Hubble data) calibrated to be energies. That, combined with the exposure time of the image, gives you ''flux units''. Spitzer data comes in ''flux (density) per unit (pixel) area'' instead, MegaJanskys per steradian (MJy/sr). 1 MJy = <math>10^{6}</math> Jy, and a sr is a solid angle.

If you've done photometry before, and expect to do it exactly the same way again here, '''it won't work''', because '''this matters'''.

1 square arcsec is <math>2.3504 \times 10^{-11}</math> sr. (1 degree = 60 arcmin = 3600 arcsec.)
okay, I'm confused: what is a square arc sec? How do you square an angle measurement?
If you want to convert the image from MJy/sr to uJy/square arcsec, multiply the image by 23.5045. The units of this number are (uJy/arcsec)/(MJy/sr).

If you want to take a Spitzer image and use your previous routines on it, the most efficient way to do this is probably to take the image in MJy/sr and multiply out the "per sr" part of it so that it is instead in MJy/px. The subtlety in this step is that each Spitzer array has slightly different pixel sizes, and the mosaics that we create have different sizes yet again from the original images. You can make mosaics with whatever size pixels you want, so if you get Spitzer mosaics from more than one astronomer, or more than one Spitzer wavelength, chances are excellent that the pixels will be slightly different sizes. The information on the pixel sizes are in the [[FITS_format|FITS]] header of each image.

The following paragraphs are a high-level summary of what to do for any Spitzer image data you may encounter; see below for a cookbook of the process for one mosaic.

Look in the FITS header of the mosaics for the keywords "CDELT1" and "CDELT2". These keywords are set to be the scale of the rows and columns in degrees per pixel. Using the values of these keywords, and the conversions above, you can figure out the number of square degrees per pixel, the number of square arcsec per pixel, and finally the number of steradians per pixel. Multiply the whole image in MJy/sr by the number of sr/px to get MJy/px.

If you are instead working with the individual BCDs (read this as: the individual little images that went into the big mosaic), you should look for keywords "PXSCAL1" and "PXSCAL2". NOTE that these pixels ARE NOT SQUARE, and this is more important for MIPS data. From here, you now have the same information as the "CDELT1" and "CDELT2" above, so you can follow the same procedure.

=Units of Spitzer Photometry=

==Introduction==

The photometry software that people use at the SSC, called APEX, produces fluxes in microJanskys. The final bandmerged catalog you can get has listed fluxes in microJanskys, as well as magnitudes.

Astronomers use magnitudes in color-color or color-magnitude plots. Astronomers use a variant on fluxes in spectral energy distribution (SED) plots.

==Magnitudes==

A magnitude is really a flux ratio. It is defined as follows, where M's are magnitudes and F's are fluxes:
<math>M_1 - M_2 = 2.5 \times \log \left(\frac{F_2}{F_1}\right)</math> (eqn 1)

The magnitude system (in the optical) was defined to be referenced to Vega. In other words, Vega is defined to be zero magnitude, and you would then define magnitudes of anything else as follows:
<math>M = 2.5 \times \log \left(\frac{F_{\mathrm{Vega}}}{F}\right)</math> (eqn 2)

When they looked at Vega with IRAS, they discovered that it did NOT look like they expected, and in fact it has a large infrared excess! Therefore, infrared magnitudes are defined with respect to what Vega would be, if it did not have an excess.

We have published the zero points (e.g., the "Vega flux") for most of our bandpasses. They are (copied from various places on the web):
*IRAC 1 : 280.9 Jy
*IRAC 2 : 179.7 Jy
*IRAC 3 : 115.0 Jy
*IRAC 4 : 64.13 Jy
*MIPS 1 : 7.14 Jy
*MIPS 2 : 0.775 Jy
*MIPS 3 : 0.159 Jy

Therefore, in order to convert the uJy that apex returns into magnitudes, use the equation 2 above, substituting these so-called "zero-point fluxes" in for "Fvega." Note that the zero-point fluxes are in Janskys and the fluxes returned by APEX are in microJanskys.

You can find the zeropoints for 2MASS magnitudes on the web as well:
*J : 1594 Jy
*H : 1024 Jy
*K : 666.7 Jy

Note that plain magnitudes get fainter (the number gets larger) as the distance of the object increases. BUT, colors (differences in magnitudes) are ratios of fluxes, and therefore independent of distance.

==Spectral Energy Distributions (SEDs)==

SEDs are energy plotted against some measure of the photon -- frequency or wavelength. The reason astronomers do this is to see how much energy is produced by the object as a function of frequency or wavelength.
Now it's really going to get a little hairy! Steel your nerves and plunge onwards... it really all comes down to unit conversion.

1Jy = <math>10^{-23}</math> erg/s/cm^2/Hz (in cgs units rather than mks units, sorry). A Jansky is technically a unit of "flux density." In order to get rid of the "per Hz", you need to multiply the Jy by the frequency of the bandpass center.

Astronomers coming from the longer wavelengths will tend to plot up nu * F(nu) (written as <math>\nu F_{\nu}</math>) against nu, where "nu" (<math>\nu</math>) is the frequency. The units of <math>F_{\nu}</math> are Janskys.

Astronomers coming from the shorter wavelengths will tend to plot up lambda * F(lambda) (written as <math>\lambda F_\lambda</math>), where "lambda" (<math>\lambda</math>) is the wavelength of the light. The units of <math>F_{\lambda}</math> are NOT Janskys.

<math>\lambda \times \nu = c</math>, the speed of light. In order to convert the Janskys into units of <math>F_{\lambda}</math>, you need to take into account the differentials (ah-HA, calculus being used here!), e.g., the fact that
<math>\frac{dF}{d\lambda} = \frac{dF}{d\nu} \frac{d\nu}{d\lambda}</math> and <math>d\nu = \frac{c}{\lambda^2}d\lambda</math>
So you need to multiply the <math>F_{\nu}</math> by <math>c/\lambda^2</math> to convert it into <math>F_{\lambda}</math>.

Additionally, to analyze the Spitzer data, it's often useful to pretend that the contribution from the star is a blackbody. It's not really, but it's awful close, especially in the infrared.

A blackbody's flux density is given by (where T is temperature, and other constants are given below)
<math>B_{\lambda} = \left(\frac{2hc^2/\lambda^5}{\exp(hc/\lambda kT)-1)}\right)</math> (eqn 3)
but of course we want to plot <math>\lambda \times B_{\lambda}</math>:
<math>\lambda B_{\lambda} = \left(\frac{2hc^2/\lambda^4}{\exp(hc/\lambda kT)-1)}\right)</math> (eqn 4)
Values of these constants all in cgs units:
*h = 6.6260755d-27 erg*sec
*c = 2.997924d10 cm/sec
*k = 1.380658d-16 erg/deg

In words, in order to analyze our data, we need to have something that does the following:

#Reads in the fluxes from the files.
#Converts the Spitzer fluxes (and errors) into magnitudes (if necessary).
#Converts the 2MASS magnitudes (and errors) into fluxes (if necessary).
#Makes color-color and color-magnitude plots for stars in our region using magnitudes.
#Makes SED plots for individual objects, but converting numbers first into the right units:
##Creates an array of the wavelengths of each measurement, keeps a copy of the version in microns, and converts to cm.
##For any real measurements, converts the microJanskys into cgs units.
##For any real measurements, converts <math>F_{\nu}</math> into <math>F_{\lambda}</math> by multiplying the <math>F_{\nu}</math> values by the <math>d\nu/d\lambda</math> corresponding to the wavelength of each bandpass.
##For any real measurements, multiplies <math>F_{\lambda}</math> by the lambda corresponding to the wavelength of each bandpass to get <math>\lambda F_{\lambda}</math>.
#For any real measurements, plots the log of the <math>\lambda F_{\lambda}</math> data points (in cgs units) against the log of the lambda data points (in microns, only because that makes it easier to read). Labels the axes (with units)! Plots the error bars on top of the data points (also converted from uJy).
#For any real measurements, for any star with at least 2 fluxes, fits a blackbody to the energies derived from the three 2MASS and first 2 IRAC bands. There are two free parameters in this fit -- the temperature of the blackbody and an additive (in the log) offset related to the distance of the object. If we know the temperature of the star (via a spectral type) and the distance to the object, then we know the values for the temperature and the offset.

Why are we plotting <math>\lambda F_{\lambda}</math> vs. <math>\lambda</math>? Well, only because I think in wavelength, not frequency. I don't know off the top of my head the frequencies of the Spitzer bandpasses, but I do know their wavelengths. Why are we plotting <math>\lambda F_{\lambda}</math> instead of <math>\nu F_{\nu}</math>? Well, only for internal consistency. Since one axis is in wavelength units, it makes sense to have the other axis also in wavelength units.

=Cookbook for image conversion: Method One=

==Step zero. What do you have and what do you need?==
You have an image in MJy/sr(/px). You have the number of degrees per pixel.

You need to convert the image to MJy(/px), a.k.a "get rid of the steradians."

The things that make this hard are:
*The pixel size of the mosaic changes depending on wavelength and where you got the mosaic, so I can't just give you one number to work for all mosaics every time.
*The "pixels" in the above are kind of a funny, hidden unit and the accounting of it works in some unexpected ways, which is why it's in parentheses above (and below).

==Step one. Find out what the size of the pixels are in your images==
For the IRAC-1 mosaic I created for you in July 2006, CDELT1=-0.000339 degrees per pixel, and CDELT2=0.000339 degrees per pixel. (ignore the minus sign; it has something to do with a fits convention.) You can find this out for any given mosaic by looking in the fits header.

==Step two. Find out what the size of the pixels are in square degrees per pixel==

degrees degrees square degrees
0.000339 ------- * 0.000339 ------- = 1.14921e-7 ---------------
pixel pixel (square) pixel

==Step three. Find out what the conversion is between square degrees and sr.==
There are 60 arcminutes in a degree. There are 60 arcseconds in an arcminute.

60 arcminutes 60 arcseconds 3600 arcseconds
------------- * -------------- = ---------------
1 degree 1 arcminutes 1 degree

Square it!

(1 degree)^2 = 1 square degree = (3600 arcsec)^2 = 1.296e7 square arcsec

We look up that 1 square arcsec is 2.3504x10^(-11) sr.

1.297e7 square arcsec 2.3504e-11 sr sr
------------- * ---------------- = 0.000304847 ------------
square degree 1 square arcsec square degree

==Step four. Find out what the size of the pixels are in sr.==

square degrees sr sr
1.14921e-7 --------------- * 0.000304847 ------------ = 3.50333e-11 ----
(square) pixel square degree px

==Step five. Convert the units of the image.==

MJy sr MJy
--------- * 3.50333e-11 ---- = ---
sr px px

So multiply this whole image by 3.50333e-11. The units of the image (and consequently the photometry you get out) are in MJy (MegaJanskys). If you want to get it in Janskys:

1e6 Jy
----
MJy

so multiply the image by 1e6 to get the image into Jy:

MJy 1e6 Jy Jy
--- * ---- = ----
px MJy px

If you want to get it into microJy, there are 1e6 uJy in a Jy, and I'll let you do that one.

=Method 2: Alternative but completely equivalent and possibly more straightforward solution=
==Step zero. What do you have and what do you need?==
You have an image in MJy/sr(/px). You have the number of degrees per pixel.

You need to convert the image to uJy(/px), a.k.a "get rid of the steradians" AND convert to microJanskys to get the numbers to still be reasonable and not very tiny or very large.

The things that make this hard are:
* The pixel size of the mosaic changes depending on wavelength and where you got the mosaic, so I can't just give you one number to work for all mosaics every time.
* The "pixels" in the above are kind of a funny, hidden unit and the accounting of it works in some unexpected ways.

==Step one. Find out what the size of the pixels are in your images==
For the IRAC-1 mosaic I created for you in July 2006, CDELT1=-0.000339 degrees per pixel, and CDELT2=0.000339 degrees per pixel. (ignore the minus sign; it has something to do with a fits convention.) For any given mosaic, you can find out these values by looking in the fits header.

==Step two. Convert your image from MJy/sr to uJy/square arcsec==
We look up that there are :

[uJy/sq. arcsec]
23.5045 -----------------
[MJy/sr]

So multiply the image by 23.5045 to get it into uJy/square arcsec
==Step three. Find out what the size of the pixels are in square degrees per pixel==

degrees degrees square degrees
0.000339 ------- * 0.000339 ------- = 1.14921e-7 ---------------
pixel pixel (square) pixel

==Step four. Find out how many square arcsec there are in a pixel.==
There are 60 arcminutes in a degree. There are 60 arcseconds in an arcminute.

60 arcminutes 60 arcseconds 3600 arcseconds
------------- * -------------- = ---------------
1 degree 1 arcminutes 1 degree

Square it!

(1 degree)^2 = 1 square degree = (3600 arcsec)^2 = 1.296e7 square arcsec

Convert the pixel size.

square degrees square arcsec square arcsec
1.14921e-7 --------------- * 1.296e7 -------------- = 1.48938 ---------------
(square) pixel square degree (square) px

==Step five. Convert the image==

uJy 1.48938 square arcsec uJy
----------- * --------------- = ------
sq arcsec (*px) (square) px (px)

So multiply the image by 1.48938 to get it into uJy/px.

Studying Young Stars

2007-10-26T23:58:35Z

Weehler: /* More in depth, and SEDs */

=Brief philosophical notes from Luisa=
==Real science vs. textbook science==
*Science (history) as presented in textbooks may seem a never-ending series of right answers. Real science has a lot of dead ends as we struggle to find out what the ‘right answer’ is.
*Science problems in textbooks have well-defined problems, specific methods you’re supposed to use to solve them, and right (exact) answers (1.2 can be wrong when 1.3 is right). Real science is not quite “made up as you go along” but it may feel that way in the coming days. Different people approach the same problem in different ways, and many answers can be right (1.2 and 1.3 can both be right). '''The only way you know it’s the right answer is if you believe that everything you did to get there is right.'''

==Why should anyone care about young stars?==
*Understanding star formation includes understanding how planets form, including planets like Earth.
*Star formation is the "happening field" right now! TONS of new discoveries happening all the time, many driven by Spitzer.
*A friend who is the author of a popular college textbook told me that the chapter that she revises most frequently (particularly recently) in response to new developments is the star formation chapter.
*By doing this project, you are participating in the revolution!

=Young stars in general: Introduction to (low-mass) star formation=
==Overview==

{| cellpadding="1"
| [[Image:starformationcartoon.png]]
|''Cartoon from Greene, American Scientist, Jul-Aug 2001''
|}

Stars begin their lives in a cloud of gas and dust (a nebula). Gravitational forces cause the nebula to start to condense (shrink). (a, b)

As the nebula shrinks, like a spinning skater pulling in her arms, it begins to spin more rapidly. The same physics ("conservation of angular momentum") means that the dust and gas in the nebula doesn't fall straight into the center; it falls onto a disk surrounding the central object, and from the disk, the matter falls onto the central object. The temperature at the center of the shrinking nebula rises due to increasing pressure and friction between the particles. The figure has this stage labeled as a "protostar", but for some astronomers, beginning at this stage, and until the star starts to turn H into He the object is still called a protostar. Since the protostar is still embedded in a thick cloud of gas and dust, it can only be detected in the infrared. (c)

When the protostar enters the next stage, labeled in the figure as the T Tauri stage, it’s still gaining mass and contracting slowly because material is still falling onto it, but it begins to eject gas in two giant gas jets, called bipolar flows. These jets and stellar winds eventually sweep away the envelope of gas still surrounding the protostar. In the surrounding disk protoplanets are beginning to form. (d)

Leftover material in the disk surrounding the star clumps together and undergoes many collisions until most of the material has been swept up by objects orbiting the star, such as planets, asteroids and comets. (e)

The star’s life so far has been governed by the continuous inward pressure of gravity. The gravitational pressure keeps compressing the gas into a smaller and smaller volume, making it hotter and hotter in the core. As soon as the temperature in the core of the protostar becomes great enough, nuclear fusion begins. When this nuclear fusion begins, finally the star has a way to "fight back" against gravity. So much energy is released in this reaction that it enables the star to "push back" with an outward radiation pressure that balances the inward push of gravity. The protostar is now a full-fledged star, fusing hydrogen into helium in its core. (f) The star will stay the same size until it runs out of nuclear fuel in the core (all of the hydrogen has been converted into helium). Then, the pressure from gravity takes over again, pushing in on the star.

'''Movies!''' -- The sequence of star formation described above can also be visualized with some movies created by the Spitzer public affairs group. First [http://ipac.jpl.nasa.gov/web_movies/pa/ssc2004-08v2_qt4.mov form grains in the disk] (note that the grains are distinctly green because they are olivine, like green sand beaches in Hawaii, and the grains get covered in ice). Then [http://ipac.jpl.nasa.gov/web_movies/pa/ssc2004-08v3_qt4.mov form a planet that clears a gap in the disk]. Next,
[http://ipac.jpl.nasa.gov/web_movies/pa/ssc2004-22v2_qt4.mov this is what happens when you form many planets at once] - there are many gaps formed at once. (Where else in the Solar System have you seen this physics before? [http://saturn.jpl.nasa.gov/multimedia/images/raw/raw-images-details.cfm?feiImageID=112884 stumped?]) Of course, it can get crowded in these protoplanetary disks, so [http://ipac.jpl.nasa.gov/web_movies/pa/ssc2004-17v1_qt4.mov occasionally stuff hits each other], creating a second generation of dust. When this happens, [http://ipac.jpl.nasa.gov/web_movies/pa/ssc2004-17v2_qt4.mov the dust gets smeared out into a ring], and there can be many collisions that create transient dust, e.g., dust that comes and goes depending on when you look. Finally, [http://ipac.jpl.nasa.gov/web_movies/pa/ssc2004-22v1_qt4.mov you end up with a fully-fledged planetary system], but you have the remnants of the protoplanetary disk left in the outer reaches of the system. (What do we call this in our Solar System? hint: not the asteroid belt. out farther than that -- see in the movie, we go out well past a "jupiter".)

==More in depth, and SEDs==
The cartoon above is the version of this information that is appropriate for an "educated person from the general public." Now, let's look at this same story again, but using the kinds of plots and information used by professional astronomers.

A spectral energy distribution (SED) is a graph of the energy emitted by an object (any object) as a function of different wavelengths. There are some example SEDs of young stellar objects below. Astronomers originally used the slope of the SEDs for protostars between about 2 and 20 microns quite literally to '''define''' different classes. (This is similar to the process that astronomers such as [http://en.wikipedia.org/wiki/Annie_Jump_Cannon Annie Jump Cannon] followed when they originally classified stars -- astronomers start by putting similar objects together, and through this process, eventually deeper physical understanding follows.) These stages that were defined based on the SED slope are dubbed Class I, II, and III. As we learned more, we created another class called "flat" between Class I and II. The very earliest stages, ones where there are no 2 micron observations at all, are the Class 0s. These classifications more or less match up to the overall sequence of events described by the cartoon above.

{| cellpadding="2"
| [[Image:class0.png]]
| This SED corresponds to the earliest, most embedded phase, called Class 0. The x-axis is wavelength, and the y-axis is energy. It looks like a cold black body; all of the emission comes from the dust and gas in the cloud. The mass of the envelope surrounding the star is more than half a solar mass. The age of the object is about 10,000 (<math>10^4</math>) years. This is thought to be the main accretion phase, where most of the mass of the object is accreted from the cloud.
|}

{| cellpadding="2"
| [[Image:classI.png]]
|This Class I SED is the next stage. See how now there is a warmer blackbody corresponding to the central object, but most of the energy is coming from the dusty cocoon around the star. The "bite" that is at about 10 microns tells us that there are silicates (beach sand) in the dust around the star. About this time is when the rate of accretion slows. The envelope is now about 0.1 Msun. The age of the object is about <math>10^5</math> years.
|}

{| cellpadding="2"
| [[Image:classII.png]]
|This Class II SED is also a kind of object also known as a young T Tauri ("Classical T Tauri"). Most of the energy is coming from the central object (warm blackbody), but there is a little emission from the disk; the disk is optically thick. The disk mass now is very roughly 0.01 Msun. The age of the object is about <math>10^6</math> years. Several of the objects in IC 2118 look like this.
|}

{| cellpadding="2"
| [[Image:classIII.png]]
|This SED is of the last stage, Class III; this is an older T Tauri ("Weak-Lined T Tauri"). The protostar still has a little dust left around it. The disk now is optically thin. The disk mass is very roughly 0.003 Msun. The object is about <math>10^7</math> years old.
Several of the objects in IC2118 look like this.
|}

==Several important notes about this classification scheme==
*As the ages mentioned above suggest, this sequence of Class 0 to I to II to III is often interpreted as an age sequence, so a Class I object is younger than a Class II object, etc. Some of the most recent evidence suggests that maybe the connection to age is not as secure as we have been thinking! So let me just reiterate: the Class of each object is ''defined'' by the slope of the SED. The physical interpretation of the classification definition is degree of embeddedness, e.g., Class 0s are still buried deep within their natal cloud, and Class IIIs have freed themselves. '''The interpretation of the Class as an age may change.'''
*This process strictly only applies to low-mass stars. High-mass stars might very well do this, only faster. We just don’t know yet for sure.
*Class 0s are the the hardest to "catch in the act", from which we infer that they are the shortest lived. Not too long ago, the list of all of the Class 0s known could fit on one page. Spitzer is changing that. Class Is are also being found by Spitzer in abundance.
*Class 0s used to be defined as "undetectable in IR." Even before Spitzer, deeper integrations forced a change in that definition.
*Although the story seems nice and well-defined, even before Spitzer, Class IIs and IIIs have been found at the same ages, e.g., some stars lose their disks very quickly, and some hold on to them for a long time. Now with Spitzer, we're muddying the waters even more.
*A current major question in star formation is the how and why of this process.
*It’s not clear whether Class 0s and Is are found at the same age – until very recently, too few of them were known, and getting an age for them is tough.
*[http://www.spitzer.caltech.edu/Media/releases/ssc2004-17/release.shtml This press release] talks about A stars, which are a little massive for Class 0/I/II/III, but the confusion in disk clearing timescales is vividly displayed there.
*We can be fooled! You can imagine that a Class III that is edge-on might look like a Class II. It could be that some things we think are the youngest protostars are actually just edge-on older things. This is also one of the current burning questions.
*Most people still use the series Class 0-I-II-III to mean a series of youngest to oldest, but it’s important to remember all of these uncertainties.

The details of the shape of the SED can tell us about the disk structure. Dips and wiggles in the SED may suggest, e.g., that there is no (or little) dust near the star, just further out. (see [http://www.spitzer.caltech.edu/Media/releases/ssc2004-08/ssc2004-08c.shtml this graphic from the SSC press release archive].)

=Young stars in general: Finding the cluster members=

Spitzer is so sensitive that it easily sees things at the far reaches of the Universe with only a few seconds' integration. When studying clusters of stars, not just with Spitzer, one of the first major goals is to figure out which objects are truly cluster members and which are not. [[Media:findingclustermembers.pdf| This pdf file]] has a discussion of how to find members of young clusters in general.

Studying Young Stars

2007-10-26T23:47:44Z

Weehler: /* Overview */

=Brief philosophical notes from Luisa=
==Real science vs. textbook science==
*Science (history) as presented in textbooks may seem a never-ending series of right answers. Real science has a lot of dead ends as we struggle to find out what the ‘right answer’ is.
*Science problems in textbooks have well-defined problems, specific methods you’re supposed to use to solve them, and right (exact) answers (1.2 can be wrong when 1.3 is right). Real science is not quite “made up as you go along” but it may feel that way in the coming days. Different people approach the same problem in different ways, and many answers can be right (1.2 and 1.3 can both be right). '''The only way you know it’s the right answer is if you believe that everything you did to get there is right.'''

==Why should anyone care about young stars?==
*Understanding star formation includes understanding how planets form, including planets like Earth.
*Star formation is the "happening field" right now! TONS of new discoveries happening all the time, many driven by Spitzer.
*A friend who is the author of a popular college textbook told me that the chapter that she revises most frequently (particularly recently) in response to new developments is the star formation chapter.
*By doing this project, you are participating in the revolution!

=Young stars in general: Introduction to (low-mass) star formation=
==Overview==

{| cellpadding="1"
| [[Image:starformationcartoon.png]]
|''Cartoon from Greene, American Scientist, Jul-Aug 2001''
|}

Stars begin their lives in a cloud of gas and dust (a nebula). Gravitational forces cause the nebula to start to condense (shrink). (a, b)

As the nebula shrinks, like a spinning skater pulling in her arms, it begins to spin more rapidly. The same physics ("conservation of angular momentum") means that the dust and gas in the nebula doesn't fall straight into the center; it falls onto a disk surrounding the central object, and from the disk, the matter falls onto the central object. The temperature at the center of the shrinking nebula rises due to increasing pressure and friction between the particles. The figure has this stage labeled as a "protostar", but for some astronomers, beginning at this stage, and until the star starts to turn H into He the object is still called a protostar. Since the protostar is still embedded in a thick cloud of gas and dust, it can only be detected in the infrared. (c)

When the protostar enters the next stage, labeled in the figure as the T Tauri stage, it’s still gaining mass and contracting slowly because material is still falling onto it, but it begins to eject gas in two giant gas jets, called bipolar flows. These jets and stellar winds eventually sweep away the envelope of gas still surrounding the protostar. In the surrounding disk protoplanets are beginning to form. (d)

Leftover material in the disk surrounding the star clumps together and undergoes many collisions until most of the material has been swept up by objects orbiting the star, such as planets, asteroids and comets. (e)

The star’s life so far has been governed by the continuous inward pressure of gravity. The gravitational pressure keeps compressing the gas into a smaller and smaller volume, making it hotter and hotter in the core. As soon as the temperature in the core of the protostar becomes great enough, nuclear fusion begins. When this nuclear fusion begins, finally the star has a way to "fight back" against gravity. So much energy is released in this reaction that it enables the star to "push back" with an outward radiation pressure that balances the inward push of gravity. The protostar is now a full-fledged star, fusing hydrogen into helium in its core. (f) The star will stay the same size until it runs out of nuclear fuel in the core (all of the hydrogen has been converted into helium). Then, the pressure from gravity takes over again, pushing in on the star.

'''Movies!''' -- The sequence of star formation described above can also be visualized with some movies created by the Spitzer public affairs group. First [http://ipac.jpl.nasa.gov/web_movies/pa/ssc2004-08v2_qt4.mov form grains in the disk] (note that the grains are distinctly green because they are olivine, like green sand beaches in Hawaii, and the grains get covered in ice). Then [http://ipac.jpl.nasa.gov/web_movies/pa/ssc2004-08v3_qt4.mov form a planet that clears a gap in the disk]. Next,
[http://ipac.jpl.nasa.gov/web_movies/pa/ssc2004-22v2_qt4.mov this is what happens when you form many planets at once] - there are many gaps formed at once. (Where else in the Solar System have you seen this physics before? [http://saturn.jpl.nasa.gov/multimedia/images/raw/raw-images-details.cfm?feiImageID=112884 stumped?]) Of course, it can get crowded in these protoplanetary disks, so [http://ipac.jpl.nasa.gov/web_movies/pa/ssc2004-17v1_qt4.mov occasionally stuff hits each other], creating a second generation of dust. When this happens, [http://ipac.jpl.nasa.gov/web_movies/pa/ssc2004-17v2_qt4.mov the dust gets smeared out into a ring], and there can be many collisions that create transient dust, e.g., dust that comes and goes depending on when you look. Finally, [http://ipac.jpl.nasa.gov/web_movies/pa/ssc2004-22v1_qt4.mov you end up with a fully-fledged planetary system], but you have the remnants of the protoplanetary disk left in the outer reaches of the system. (What do we call this in our Solar System? hint: not the asteroid belt. out farther than that -- see in the movie, we go out well past a "jupiter".)

==More in depth, and SEDs==
The cartoon above is the version of this information that is appropriate for an "educated person from the general public." Now, let's look at this same story again, but using the kinds of plots and information used by professional astronomers.

A spectral energy distribution (SED) is a graph of the energy emitted by an object (any object) as a function of different wavelengths. There are some example SEDs of young stellar objects below. Astronomers originally used the slope of the SEDs for protostars between about 2 and 20 microns quite literally to '''define''' different classes. (This is similar to the process that astronomers such as [http://en.wikipedia.org/wiki/Annie_Jump_Cannon Annie Jump Cannon] followed when they originally classified stars -- astronomers start by putting similar objects together, and through this process, eventually deeper physical understanding follows.) These stages that were defined based on the SED slope are dubbed Class I, II, and III. As we learned more, we created another class called "flat" between Class I and II. The very earliest stages, ones where there are no 2 micron observations at all, are the Class 0s. These classifications more or less match up to the overall sequence of events described by the cartoon above.

{| cellpadding="2"
| [[Image:class0.png]]
| This SED corresponds to the earliest, most embedded phase, called Class 0. The x-axis is wavelength, and the y-axis is energy. It looks like a cold black body; all of the emission comes from the dust and gas in the cloud. The mass of the envelope surrounding the star is more than half a solar mass. The age of the object is about 10,000 (<math>10^4</math>) years. This is thoguht to be the main accretion phase, where most of the mass of the object is accreted from the cloud.
|}

{| cellpadding="2"
| [[Image:classI.png]]
|This Class I SED is the next stage. See how now there is a warmer blackbody corresponding to the central object, but most of the energy is coming from the dusty cocoon around the star. The "bite" that is at about 10 microns tells us that there are silicates (beach sand) in the dust around the star. About this time is when the rate of accretion slows. The envelope is now about 0.1 Msun. The age of the object is about <math>10^5</math> years.
|}

{| cellpadding="2"
| [[Image:classII.png]]
|This Class II SED is also a kind of object also known as a young T Tauri ("Classical T Tauri"). Most of the energy is coming from the central object (warm blackbody), but there is a little emission from the disk; the disk is optically thick. The disk mass now is very roughly 0.01 Msun. The age of the object is about <math>10^6</math> years. Several of the objects in IC 2118 look like this.
|}

{| cellpadding="2"
| [[Image:classIII.png]]
|This SED is of the last stage, Class III; this is an older T Tauri ("Weak-Lined T Tauri"). The protostar still has a little dust left around it. The disk now is optically thin. The disk mass is very roughly 0.003 Msun. The object is about <math>10^7</math> years old.
Several of the objects in IC2118 look like this.
|}

==Several important notes about this classification scheme==
*As the ages mentioned above suggest, this sequence of Class 0 to I to II to III is often interpreted as an age sequence, so a Class I object is younger than a Class II object, etc. Some of the most recent evidence suggests that maybe the connection to age is not as secure as we have been thinking! So let me just reiterate: the Class of each object is ''defined'' by the slope of the SED. The physical interpretation of the classification definition is degree of embeddedness, e.g., Class 0s are still buried deep within their natal cloud, and Class IIIs have freed themselves. '''The interpretation of the Class as an age may change.'''
*This process strictly only applies to low-mass stars. High-mass stars might very well do this, only faster. We just don’t know yet for sure.
*Class 0s are the the hardest to "catch in the act", from which we infer that they are the shortest lived. Not too long ago, the list of all of the Class 0s known could fit on one page. Spitzer is changing that. Class Is are also being found by Spitzer in abundance.
*Class 0s used to be defined as "undetectable in IR." Even before Spitzer, deeper integrations forced a change in that definition.
*Although the story seems nice and well-defined, even before Spitzer, Class IIs and IIIs have been found at the same ages, e.g., some stars lose their disks very quickly, and some hold on to them for a long time. Now with Spitzer, we're muddying the waters even more.
*A current major question in star formation is the how and why of this process.
*It’s not clear whether Class 0s and Is are found at the same age – until very recently, too few of them were known, and getting an age for them is tough.
*[http://www.spitzer.caltech.edu/Media/releases/ssc2004-17/release.shtml This press release] talks about A stars, which are a little massive for Class 0/I/II/III, but the confusion in disk clearing timescales is vividly displayed there.
*We can be fooled! You can imagine that a Class III that is edge-on might look like a Class II. It could be that some things we think are the youngest protostars are actually just edge-on older things. This is also one of the current burning questions.
*Most people still use the series Class 0-I-II-III to mean a series of youngest to oldest, but it’s important to remember all of these uncertainties.

The details of the shape of the SED can tell us about the disk structure. Dips and wiggles in the SED may suggest, e.g., that there is no (or little) dust near the star, just further out. (see [http://www.spitzer.caltech.edu/Media/releases/ssc2004-08/ssc2004-08c.shtml this graphic from the SSC press release archive].)

=Young stars in general: Finding the cluster members=

Spitzer is so sensitive that it easily sees things at the far reaches of the Universe with only a few seconds' integration. When studying clusters of stars, not just with Spitzer, one of the first major goals is to figure out which objects are truly cluster members and which are not. [[Media:findingclustermembers.pdf| This pdf file]] has a discussion of how to find members of young clusters in general.

IC 2118 Current Research Activities

2007-08-06T12:01:41Z

Weehler:

=Some more of the basics behind what we are doing!=

[[Image:L1014.jpg|right]]
Here is an image of an object called L1014. This object used to be known as a "starless core." On the left is a picture in visible light, and you can see why people thought it was starless. On the right is a picture from Spitzer clearly revealing a baby star inside (it's the bright red thing)!

When the protostar enters the next stage, labeled in the figure as the T Tauri stage, it’s still gaining mass and contracting slowly because material is still falling onto it, but it begins to eject gas in two giant gas jets, called bipolar flows. These jets and stellar winds eventually sweep away the envelope of gas still surrounding the protostar. In the surrounding disk protoplanets are beginning to form. (d) Astronomers call this phase the T Tauri phase, but some objects with jets are also "Class I"; T Tauris are also either "Class II" (for ones with thick disks) or "Class III" (for ones with thinner disks).

Leftover material in the disk surrounding the star clumps together and undergoes many collisions until most of the material has been swept up by objects orbiting the star, such as planets, asteroids and comets. (e)

The star’s life so far has been governed by the continuous inward pressure of gravity. The gravitational pressure keeps compressing the gas into a smaller and smaller volume, making it hotter and hotter in the core. As soon as the temperature in the core of the protostar becomes great enough, about 1,000,000 K, nuclear fusion begins. Nuclear fusion is the process in which small atomic nuclei combine to make larger atomic nuclei, releasing lots of energy. Inside the star, the gravitational pressure eventually pushes 4 protons and 2 electrons so close together that they fuse together to make helium. When this nuclear fusion begins, finally the star has a way to "fight back" against gravity. So much energy is released in this reaction that it enables the star to "push back" with an outward radiation pressure that balances the inward push of gravity. The protostar is now a full-fledged star, fusing hydrogen into helium in its core. (f) The star will stay the same size until it runs out of nuclear fuel in the core (all of the hydrogen has been converted into helium). Then, the pressure from gravity takes over again, pushing in on the star.

Star formation can be triggered by the collapse of large clouds of gas and dust. Sometimes the clouds collapse all by themselves (due to gravitational forces within the cloud) and sometimes they collapse because they’ve been pushed by the radiation and winds from stars that have already been formed. This seems to be the reason behind the star formation in IC 2118. It seems to have been triggered by some combination of forces from the main Orion Nebula Cluster (the fuzzy patch in the sword) and Rigel.

=Color-Color Plots=
[[Image:irac-colorcolor.png|left]]

The magnitude scale measures the brightness of a star. The system for assigning magnitude numbers was developed in ancient times. The '''brighter''' the star is, the '''lower''' its magnitude.

An infrared color-color diagram is a useful tool in (a) finding the young stars, and (b) making a guess at the age of young stars. In this diagram astronomers compare the differences in magnitude of a star at two wavelengths to the difference in magnitude of the same star in two other wavelengths.

Here is an example of a color-color plot using IRAC colors. By making models of stars, astronomers have determined that a star's position on this graph is related to its age since the star’s position on the graph is related to how warm or cool it is. Stars start out as cool (or red) and become warmer (bluer), therefore they shift positions on the diagram as they age.

--Class 0: these are very young stars, still buried deep inside their cocoon of gas.

--Class I: these are protostars surrounded by an infalling envelope of gas.

--Class II: these protostars, now in the T-Tauri stage, have developed protoplanetary disks.

--Class III: these are nearly fully developed stars, with just a remnant disk (and possibly planets) surrounding them.

==Homework: Your own literature search!==

Search the literature for references to IC 2118. Nearly all the astronomical literature is online at Harvard’s ADS website. (http://adsabs.harvard.edu) The ADS website allows you to search for abstracts using either names (IC2118) or coordinates (Ra/Dec). The website will let you look at older full papers, but only abstracts for recent papers. If there is a recent paper you wish to read, you will need to connect through a local university.

# Find all papers by Maria Kun. Which are refereed, and which are conference proceedings?
# Find all papers involving IC 2118.
# Find Hartmann et. al, 2005, ApJ, 629, 881.
# Search the web for observations in other wavelengths, not just images but photometry as well. Some good places to start are:
#*Gator, at IRSA, provides access to 2MASS, IRAS, and other wavelengths (including some large Spitzer surveys). This site can be accessed at: http://irsa.ipac.caltech.edu/applications/Gator/
#*USNO provides optical magnitudes, if they exist. Make sure to retrieve tables, not images. This site can be accessed at: http://www.nofs.navy.mil/data/FchPix/
#*Skyview provides images at other wavelengths. This site can be accessed at: http://skyview.gsfc.nasa.gov

== [[Writing the Research Paper for IC2118]] ==

== [[Information for Oil City High School meetings]] ==

The next meeting should be at 4 0'clock pm wednesday the 25th at matt walentosky's house. (Specifically for people going to California, so we can plan out events and travel...)

== Proposal for Radio Observations of T-Tauri Candidates in IC2118 ==

'''''SUPPORT INFORMATION - T-Tauri Candidates Emit Radio'''''

'''FROM MATT WALENTOSKY 6/30/2007 @ 2:30 PM - Two Abstracts Below'''

Title: Radio emission from pre-main-sequence stars
Authors: Skinner, Stephen Lee

Affiliation: AA(Colorado Univ., Boulder.)
Publication: Ph.D. Thesis Colorado Univ., Boulder.
Publication Date: 01/1992
Category: Space Radiation
Origin: STI

Abstract
This study focuses on the properties and physical origin of radio continuum emission from pre-main-sequence (PMS) stars. These are young stars, typically less than a few million years old, and are still in a phase of gravitational contraction that will ultimately be halted by the onset of hydrogen burning in their cores. First, I address the question of the origin of centimeter continuum emission in intermediate mass (approx. equal to 3-20 solar mass) PMS stars, the so-called 'Herbig Ae/Be stars'. A high-sensitivity radio survey of 57 such stars was undertaken using the Very Large Array and Australia Telescope, resulting in the detection of twelve stars. These observations provide a homogenous data base consisting of information on source sizes, radio luminosities, variability timescales, circular polarization, and spectral energy distributions in the wavelength range 2-20 cm. Using these data along with previously published spectroscopy, I conclude that centimeter radio emission from Herbig Ae/Be stars is predominantly thermal and in many cases wind-related. An unexpected result of the above program was the serendipitous detection of circularly polarized radio emission in the low mass (approx. equal to 1 solar mass) PMS star Hubble 4, a member of the class of 'weak-lined T Tauri stars' (WTTS). This provides some of the most convincing evidence to date for the existence of ordered magnetic fields in WTTS. In a second observing program, I have searched for evidence of cold (less than or equal to 50 K) circumstellar dust around WTTS, which might exist in the form of remnant disks. Of the sixteen WTTS that were observed in the wavelength range 450-1100 microns using the James Clerk Maxwell Telescope, only V836 Tau was detected. Its spectral energy distribution longward of 10 microns is consistent with that expected for a flat, axisymmetric circumstellar disk of mass approx. equal to 0.04 solar mass (= 42 Jupiter masses). This star may be a rare example of an object in which disk dispersal is underway, but not yet complete.

Title: Centimeter Radio Emission from Low-Mass Weak T Tauri Stars in Taurus-Auriga
Authors: Chiang, Eugene; Phillips, R. B.

Affiliation: AA(MIT), AB(MIT Haystack Observatory)
Publication: American Astronomical Society, 185th AAS Meeting, #48.08; Bulletin of the American Astronomical Society, Vol. 26, p.1388
Publication Date: 12/1994
Origin: AAS
Abstract Copyright:(c) 1994: American Astronomical Society

Abstract
We report on the results of a sensitive survey for lambda 3.6 cm radio emission from low-mass, weak T Tauri (WTT) stars in the Taurus-Auriga cloud complex. The target population consists of stars in the Herbig and Bell Catalog of spectral type K7 or later, and W(Hα ) <= 10 Angstroms. Of the 28 such stars surveyed using the Very Large Array down to detection thresholds of ~ 0.1 mJy, 7 (possibly 8) are observed to emit at strengths ranging from 0.1 to 2 mJy. Five of these young radio stars are newly discovered in our survey: V827 Tau and V710 Tau B are discovered to be relatively strong sources of mJy emission, while IW Tau, UX Tau B, and the possible detection LkHa 332-G1 form a new population of relatively weak emitters. Our radio survey and complementary surveys are pooled, and of 43 WTT stars K7 or later in Tau-Aur, 14 are now known to be radio emitters at lambda 6 and lambda 3.6 cm. Correlations between radio luminosity and other stellar parameters have been attempted but generally yield null results. Wide binarity (component separations in excess of 0''.13, 20 AU) appears unrelated to radio emission, as does spectral type. Furthermore, we find no convincing evidence for the extreme youth of radio stars, contrary to claims in the literature over the past decade. While we do find that radio-loud stars in our sample are formally younger than the radio-quiet stars by about 0.5 Myr, the reality of this relatively small age difference is highly suspect given uncertainties in the placement of these stars on the HR diagram. Moreover, Monte Carlo-type calculations involving distributing the stars on both the HR diagram and local CO gas density cast doubts on any differences between the radio stars and the general WTT population. We conclude that the age effect for low-mass radio WTT stars in Tau-Aur, if real, is much smaller than previous estimations by factors of 4-10. It is also possible centimeter wavelength surveys to date have still not properly described the radio luminosity function of low-mass WTT stars in Tau-Aur, and we urge future observations of these young stars with denser temporal coverage.

'''FROM SPUCK 6/30/2007 @ 2:35pm'''

I'm not sure that we will be able to use the GBT telescope in Green Bank for observations. The beamwidth at 10-15 GHz is probably too big.

Here are the specs for the GBT
Beamwidth (Table 3) Diffraction beamwidth (FWHM)
8 GHz it is 90",
20 GHz it is 36",
50 GHz it is 14"

Here are some questions from Sue Ann Heatherly at NRAO-Green Bank - One question I have is: is the resolving
power of the GBT sufficient to distinguish spatially between individual stars? Does large scale emission exist within the nebula that will confuse your results?

The VLA is probably the only instrument that can be used.

Synthesized Beamwidth (arcsec)depending on configuration
'''Nick Kelley 6/30/2007 3:00PM'''

400 cm 24.0" 80.0" 260.0" 850.0"

90 cm 6.0" 17.0" 56.0" 200.0"

20 cm 1.4 " 3.9" 12.5" 44.0"

6 cm 0.4 " 1.2""" 3.9 " 14.0 "

3.6 cm 0.24 " 0.7" 2.3 " 8.4 "

2 cm 0.14 " 0.4 " 1.2" 3.9"

1.3 cm 0.08 " 0.3" 0.9 "2.8"

0.7 cm 0.05 0.15 0.47 1
restricting my creativity

== T-Tauri/IC2118 Presentation for Astroblast 2007 ==

'''''Oil City High School students will be presenting August 11 at the Oil Region Astronomical Observatory'''''

FROM Nicholas James Kelley

Astroblast is amazin ... see pic

[[Image: astroblast1.JPG|center]]

Is their anyway we can attach the powerpoint here?? Matt W.

FROM SPUCK 7/25/07 - Matt ... we may be able to upload the powerpoint somewhere and then create a link to it. I'll need to check on this. Thanks, Mr. Spuck

== Monitoring T-Tauri Stars using the Perth Obervatory Telescope ==

'''Generating Light Curves'''

Using Perth Telescope

1. Check to see what the current sky conditions are by checking the live camera

http://www.perthobservatory.wa.gov.au//information/po_sky_camera.html

2. Check to see what the current weather conditions are

http://www.perthobservatory.wa.gov.au//information/po_weather.html

3. Weather forecast for Perth - http://www.bom.gov.au/weather/wa/

Log into telescope if conditions look good.

http://perthobservatory.org or http://202.72.190.15

EMAIL tspuck@hotmail.com for telescope access

== Chandra X-Ray Data ==

'''''To date no X-Ray data for IC2118 has been located.''''' If you know of any please contact Tim Spuck at tspuck@hotmail.com.

== Kitt Peak H-alpha Data ==

'''''In January of 2007 students from Oil City High School used the 0.9 Meter Telescope at Kitt peak to image regions of IC2118 in H-alpha.''''

Matt Walentosky, Nick Kelley, Sandy Weiser were the students who went!!!!!!

== USNO U, B, V, R, and I Data ==

'''''U, B, V, R, and I data will be used to generate more accurate SEDs for the T-Tauri Candidates'''''

== Outflows or Jets ==

'''''Visible outflows or jets are strong evidence that a candidate is indeed a T-Tauri Star'''''

== Multi Wavelength Composite Images ==

'''''Images and Methods'''''

FROM SPUCK 6/30/2007 See Image Below- IC 2118 3.6 µm (blue), 5.8 µm (green), 8.0 µm (red) tri-color composite generated using MaxIm DL. (By M. Heath, N. Kelley, P. Morton, M. Walentosky, S. Weiser – Oil City High School, Oil City, PA)

[[Image: spuckim3.JPG|center]]

Talk:IC 2118 Current Research Activities

2007-08-06T12:00:51Z

Weehler:

'''What do (d) and (e) refer to? Where are (a)-(c)? Are there illustrations that go here?--[[User:Weehler|Weehler]] 05:37, 3 August 2007 (PDT)'''

'''Are electrons involved in fusion?--[[User:Weehler|Weehler]] 05:48, 3 August 2007 (PDT)'''

Talk:IC 2118 Current Research Activities

2007-08-06T11:59:31Z

Weehler:

'''What do (d) and (e) refer to? Where are (a)-(c)? Are there illustrations that go here?--[[User:Weehler|Weehler]] 05:37, 3 August 2007 (PDT)'''

IC 2118 Current Research Activities

2007-08-06T11:59:17Z

Weehler:

=Some more of the basics behind what we are doing!=

[[Image:L1014.jpg|right]]
Here is an image of an object called L1014. This object used to be known as a "starless core." On the left is a picture in visible light, and you can see why people thought it was starless. On the right is a picture from Spitzer clearly revealing a baby star inside (it's the bright red thing)!

When the protostar enters the next stage, labeled in the figure as the T Tauri stage, it’s still gaining mass and contracting slowly because material is still falling onto it, but it begins to eject gas in two giant gas jets, called bipolar flows. These jets and stellar winds eventually sweep away the envelope of gas still surrounding the protostar. In the surrounding disk protoplanets are beginning to form. (d) Astronomers call this phase the T Tauri phase, but some objects with jets are also "Class I"; T Tauris are also either "Class II" (for ones with thick disks) or "Class III" (for ones with thinner disks).

Leftover material in the disk surrounding the star clumps together and undergoes many collisions until most of the material has been swept up by objects orbiting the star, such as planets, asteroids and comets. (e)

The star’s life so far has been governed by the continuous inward pressure of gravity. The gravitational pressure keeps compressing the gas into a smaller and smaller volume, making it hotter and hotter in the core. As soon as the temperature in the core of the protostar becomes great enough, about 1,000,000 K, nuclear fusion begins. Nuclear fusion is the process in which small atomic nuclei combine to make larger atomic nuclei, releasing lots of energy. Inside the star, the gravitational pressure eventually pushes 4 protons and 2 electrons '''Are electrons involved in fusion?--[[User:Weehler|Weehler]] 05:48, 3 August 2007 (PDT)'''so close together that they fuse together to make helium. When this nuclear fusion begins, finally the star has a way to "fight back" against gravity. So much energy is released in this reaction that it enables the star to "push back" with an outward radiation pressure that balances the inward push of gravity. The protostar is now a full-fledged star, fusing hydrogen into helium in its core. (f) The star will stay the same size until it runs out of nuclear fuel in the core (all of the hydrogen has been converted into helium). Then, the pressure from gravity takes over again, pushing in on the star.

Star formation can be triggered by the collapse of large clouds of gas and dust. Sometimes the clouds collapse all by themselves (due to gravitational forces within the cloud) and sometimes they collapse because they’ve been pushed by the radiation and winds from stars that have already been formed. This seems to be the reason behind the star formation in IC 2118. It seems to have been triggered by some combination of forces from the main Orion Nebula Cluster (the fuzzy patch in the sword) and Rigel.

=Color-Color Plots=
[[Image:irac-colorcolor.png|left]]

The magnitude scale measures the brightness of a star. The system for assigning magnitude numbers was developed in ancient times. The '''brighter''' the star is, the '''lower''' its magnitude.

An infrared color-color diagram is a useful tool in (a) finding the young stars, and (b) making a guess at the age of young stars. In this diagram astronomers compare the differences in magnitude of a star at two wavelengths to the difference in magnitude of the same star in two other wavelengths.

Here is an example of a color-color plot using IRAC colors. By making models of stars, astronomers have determined that a star's position on this graph is related to its age since the star’s position on the graph is related to how warm or cool it is. Stars start out as cool (or red) and become warmer (bluer), therefore they shift positions on the diagram as they age.

--Class 0: these are very young stars, still buried deep inside their cocoon of gas.

--Class I: these are protostars surrounded by an infalling envelope of gas.

--Class II: these protostars, now in the T-Tauri stage, have developed protoplanetary disks.

--Class III: these are nearly fully developed stars, with just a remnant disk (and possibly planets) surrounding them.

==Homework: Your own literature search!==

Search the literature for references to IC 2118. Nearly all the astronomical literature is online at Harvard’s ADS website. (http://adsabs.harvard.edu) The ADS website allows you to search for abstracts using either names (IC2118) or coordinates (Ra/Dec). The website will let you look at older full papers, but only abstracts for recent papers. If there is a recent paper you wish to read, you will need to connect through a local university.

# Find all papers by Maria Kun. Which are refereed, and which are conference proceedings?
# Find all papers involving IC 2118.
# Find Hartmann et. al, 2005, ApJ, 629, 881.
# Search the web for observations in other wavelengths, not just images but photometry as well. Some good places to start are:
#*Gator, at IRSA, provides access to 2MASS, IRAS, and other wavelengths (including some large Spitzer surveys). This site can be accessed at: http://irsa.ipac.caltech.edu/applications/Gator/
#*USNO provides optical magnitudes, if they exist. Make sure to retrieve tables, not images. This site can be accessed at: http://www.nofs.navy.mil/data/FchPix/
#*Skyview provides images at other wavelengths. This site can be accessed at: http://skyview.gsfc.nasa.gov

== [[Writing the Research Paper for IC2118]] ==

== [[Information for Oil City High School meetings]] ==

The next meeting should be at 4 0'clock pm wednesday the 25th at matt walentosky's house. (Specifically for people going to California, so we can plan out events and travel...)

== Proposal for Radio Observations of T-Tauri Candidates in IC2118 ==

'''''SUPPORT INFORMATION - T-Tauri Candidates Emit Radio'''''

'''FROM MATT WALENTOSKY 6/30/2007 @ 2:30 PM - Two Abstracts Below'''

Title: Radio emission from pre-main-sequence stars
Authors: Skinner, Stephen Lee

Affiliation: AA(Colorado Univ., Boulder.)
Publication: Ph.D. Thesis Colorado Univ., Boulder.
Publication Date: 01/1992
Category: Space Radiation
Origin: STI

Abstract
This study focuses on the properties and physical origin of radio continuum emission from pre-main-sequence (PMS) stars. These are young stars, typically less than a few million years old, and are still in a phase of gravitational contraction that will ultimately be halted by the onset of hydrogen burning in their cores. First, I address the question of the origin of centimeter continuum emission in intermediate mass (approx. equal to 3-20 solar mass) PMS stars, the so-called 'Herbig Ae/Be stars'. A high-sensitivity radio survey of 57 such stars was undertaken using the Very Large Array and Australia Telescope, resulting in the detection of twelve stars. These observations provide a homogenous data base consisting of information on source sizes, radio luminosities, variability timescales, circular polarization, and spectral energy distributions in the wavelength range 2-20 cm. Using these data along with previously published spectroscopy, I conclude that centimeter radio emission from Herbig Ae/Be stars is predominantly thermal and in many cases wind-related. An unexpected result of the above program was the serendipitous detection of circularly polarized radio emission in the low mass (approx. equal to 1 solar mass) PMS star Hubble 4, a member of the class of 'weak-lined T Tauri stars' (WTTS). This provides some of the most convincing evidence to date for the existence of ordered magnetic fields in WTTS. In a second observing program, I have searched for evidence of cold (less than or equal to 50 K) circumstellar dust around WTTS, which might exist in the form of remnant disks. Of the sixteen WTTS that were observed in the wavelength range 450-1100 microns using the James Clerk Maxwell Telescope, only V836 Tau was detected. Its spectral energy distribution longward of 10 microns is consistent with that expected for a flat, axisymmetric circumstellar disk of mass approx. equal to 0.04 solar mass (= 42 Jupiter masses). This star may be a rare example of an object in which disk dispersal is underway, but not yet complete.

Title: Centimeter Radio Emission from Low-Mass Weak T Tauri Stars in Taurus-Auriga
Authors: Chiang, Eugene; Phillips, R. B.

Affiliation: AA(MIT), AB(MIT Haystack Observatory)
Publication: American Astronomical Society, 185th AAS Meeting, #48.08; Bulletin of the American Astronomical Society, Vol. 26, p.1388
Publication Date: 12/1994
Origin: AAS
Abstract Copyright:(c) 1994: American Astronomical Society

Abstract
We report on the results of a sensitive survey for lambda 3.6 cm radio emission from low-mass, weak T Tauri (WTT) stars in the Taurus-Auriga cloud complex. The target population consists of stars in the Herbig and Bell Catalog of spectral type K7 or later, and W(Hα ) <= 10 Angstroms. Of the 28 such stars surveyed using the Very Large Array down to detection thresholds of ~ 0.1 mJy, 7 (possibly 8) are observed to emit at strengths ranging from 0.1 to 2 mJy. Five of these young radio stars are newly discovered in our survey: V827 Tau and V710 Tau B are discovered to be relatively strong sources of mJy emission, while IW Tau, UX Tau B, and the possible detection LkHa 332-G1 form a new population of relatively weak emitters. Our radio survey and complementary surveys are pooled, and of 43 WTT stars K7 or later in Tau-Aur, 14 are now known to be radio emitters at lambda 6 and lambda 3.6 cm. Correlations between radio luminosity and other stellar parameters have been attempted but generally yield null results. Wide binarity (component separations in excess of 0''.13, 20 AU) appears unrelated to radio emission, as does spectral type. Furthermore, we find no convincing evidence for the extreme youth of radio stars, contrary to claims in the literature over the past decade. While we do find that radio-loud stars in our sample are formally younger than the radio-quiet stars by about 0.5 Myr, the reality of this relatively small age difference is highly suspect given uncertainties in the placement of these stars on the HR diagram. Moreover, Monte Carlo-type calculations involving distributing the stars on both the HR diagram and local CO gas density cast doubts on any differences between the radio stars and the general WTT population. We conclude that the age effect for low-mass radio WTT stars in Tau-Aur, if real, is much smaller than previous estimations by factors of 4-10. It is also possible centimeter wavelength surveys to date have still not properly described the radio luminosity function of low-mass WTT stars in Tau-Aur, and we urge future observations of these young stars with denser temporal coverage.

'''FROM SPUCK 6/30/2007 @ 2:35pm'''

I'm not sure that we will be able to use the GBT telescope in Green Bank for observations. The beamwidth at 10-15 GHz is probably too big.

Here are the specs for the GBT
Beamwidth (Table 3) Diffraction beamwidth (FWHM)
8 GHz it is 90",
20 GHz it is 36",
50 GHz it is 14"

Here are some questions from Sue Ann Heatherly at NRAO-Green Bank - One question I have is: is the resolving
power of the GBT sufficient to distinguish spatially between individual stars? Does large scale emission exist within the nebula that will confuse your results?

The VLA is probably the only instrument that can be used.

Synthesized Beamwidth (arcsec)depending on configuration
'''Nick Kelley 6/30/2007 3:00PM'''

400 cm 24.0" 80.0" 260.0" 850.0"

90 cm 6.0" 17.0" 56.0" 200.0"

20 cm 1.4 " 3.9" 12.5" 44.0"

6 cm 0.4 " 1.2""" 3.9 " 14.0 "

3.6 cm 0.24 " 0.7" 2.3 " 8.4 "

2 cm 0.14 " 0.4 " 1.2" 3.9"

1.3 cm 0.08 " 0.3" 0.9 "2.8"

0.7 cm 0.05 0.15 0.47 1
restricting my creativity

== T-Tauri/IC2118 Presentation for Astroblast 2007 ==

'''''Oil City High School students will be presenting August 11 at the Oil Region Astronomical Observatory'''''

FROM Nicholas James Kelley

Astroblast is amazin ... see pic

[[Image: astroblast1.JPG|center]]

Is their anyway we can attach the powerpoint here?? Matt W.

FROM SPUCK 7/25/07 - Matt ... we may be able to upload the powerpoint somewhere and then create a link to it. I'll need to check on this. Thanks, Mr. Spuck

== Monitoring T-Tauri Stars using the Perth Obervatory Telescope ==

'''Generating Light Curves'''

Using Perth Telescope

1. Check to see what the current sky conditions are by checking the live camera

http://www.perthobservatory.wa.gov.au//information/po_sky_camera.html

2. Check to see what the current weather conditions are

http://www.perthobservatory.wa.gov.au//information/po_weather.html

3. Weather forecast for Perth - http://www.bom.gov.au/weather/wa/

Log into telescope if conditions look good.

http://perthobservatory.org or http://202.72.190.15

EMAIL tspuck@hotmail.com for telescope access

== Chandra X-Ray Data ==

'''''To date no X-Ray data for IC2118 has been located.''''' If you know of any please contact Tim Spuck at tspuck@hotmail.com.

== Kitt Peak H-alpha Data ==

'''''In January of 2007 students from Oil City High School used the 0.9 Meter Telescope at Kitt peak to image regions of IC2118 in H-alpha.''''

Matt Walentosky, Nick Kelley, Sandy Weiser were the students who went!!!!!!

== USNO U, B, V, R, and I Data ==

'''''U, B, V, R, and I data will be used to generate more accurate SEDs for the T-Tauri Candidates'''''

== Outflows or Jets ==

'''''Visible outflows or jets are strong evidence that a candidate is indeed a T-Tauri Star'''''

== Multi Wavelength Composite Images ==

'''''Images and Methods'''''

FROM SPUCK 6/30/2007 See Image Below- IC 2118 3.6 µm (blue), 5.8 µm (green), 8.0 µm (red) tri-color composite generated using MaxIm DL. (By M. Heath, N. Kelley, P. Morton, M. Walentosky, S. Weiser – Oil City High School, Oil City, PA)

[[Image: spuckim3.JPG|center]]

IC 2118 Current Research Activities

2007-08-06T11:58:26Z

Weehler:

=Some more of the basics behind what we are doing!=

[[Image:L1014.jpg|right]]
Here is an image of an object called L1014. This object used to be known as a "starless core." On the left is a picture in visible light, and you can see why people thought it was starless. On the right is a picture from Spitzer clearly revealing a baby star inside (it's the bright red thing)!

When the protostar enters the next stage, labeled in the figure as the T Tauri stage, it’s still gaining mass and contracting slowly because material is still falling onto it, but it begins to eject gas in two giant gas jets, called bipolar flows. These jets and stellar winds eventually sweep away the envelope of gas still surrounding the protostar. In the surrounding disk protoplanets are beginning to form. (d) Astronomers call this phase the T Tauri phase, but some objects with jets are also "Class I"; T Tauris are also either "Class II" (for ones with thick disks) or "Class III" (for ones with thinner disks).
'''What do (d) and (e) refer to? Where are (a)-(c)? Are there illustrations that go here?--[[User:Weehler|Weehler]] 05:37, 3 August 2007 (PDT)'''
Leftover material in the disk surrounding the star clumps together and undergoes many collisions until most of the material has been swept up by objects orbiting the star, such as planets, asteroids and comets. (e)

The star’s life so far has been governed by the continuous inward pressure of gravity. The gravitational pressure keeps compressing the gas into a smaller and smaller volume, making it hotter and hotter in the core. As soon as the temperature in the core of the protostar becomes great enough, about 1,000,000 K, nuclear fusion begins. Nuclear fusion is the process in which small atomic nuclei combine to make larger atomic nuclei, releasing lots of energy. Inside the star, the gravitational pressure eventually pushes 4 protons and 2 electrons '''Are electrons involved in fusion?--[[User:Weehler|Weehler]] 05:48, 3 August 2007 (PDT)'''so close together that they fuse together to make helium. When this nuclear fusion begins, finally the star has a way to "fight back" against gravity. So much energy is released in this reaction that it enables the star to "push back" with an outward radiation pressure that balances the inward push of gravity. The protostar is now a full-fledged star, fusing hydrogen into helium in its core. (f) The star will stay the same size until it runs out of nuclear fuel in the core (all of the hydrogen has been converted into helium). Then, the pressure from gravity takes over again, pushing in on the star.

Star formation can be triggered by the collapse of large clouds of gas and dust. Sometimes the clouds collapse all by themselves (due to gravitational forces within the cloud) and sometimes they collapse because they’ve been pushed by the radiation and winds from stars that have already been formed. This seems to be the reason behind the star formation in IC 2118. It seems to have been triggered by some combination of forces from the main Orion Nebula Cluster (the fuzzy patch in the sword) and Rigel.

=Color-Color Plots=
[[Image:irac-colorcolor.png|left]]

The magnitude scale measures the brightness of a star. The system for assigning magnitude numbers was developed in ancient times. The '''brighter''' the star is, the '''lower''' its magnitude.

An infrared color-color diagram is a useful tool in (a) finding the young stars, and (b) making a guess at the age of young stars. In this diagram astronomers compare the differences in magnitude of a star at two wavelengths to the difference in magnitude of the same star in two other wavelengths.

Here is an example of a color-color plot using IRAC colors. By making models of stars, astronomers have determined that a star's position on this graph is related to its age since the star’s position on the graph is related to how warm or cool it is. Stars start out as cool (or red) and become warmer (bluer), therefore they shift positions on the diagram as they age.

--Class 0: these are very young stars, still buried deep inside their cocoon of gas.

--Class I: these are protostars surrounded by an infalling envelope of gas.

--Class II: these protostars, now in the T-Tauri stage, have developed protoplanetary disks.

--Class III: these are nearly fully developed stars, with just a remnant disk (and possibly planets) surrounding them.

==Homework: Your own literature search!==

Search the literature for references to IC 2118. Nearly all the astronomical literature is online at Harvard’s ADS website. (http://adsabs.harvard.edu) The ADS website allows you to search for abstracts using either names (IC2118) or coordinates (Ra/Dec). The website will let you look at older full papers, but only abstracts for recent papers. If there is a recent paper you wish to read, you will need to connect through a local university.

# Find all papers by Maria Kun. Which are refereed, and which are conference proceedings?
# Find all papers involving IC 2118.
# Find Hartmann et. al, 2005, ApJ, 629, 881.
# Search the web for observations in other wavelengths, not just images but photometry as well. Some good places to start are:
#*Gator, at IRSA, provides access to 2MASS, IRAS, and other wavelengths (including some large Spitzer surveys). This site can be accessed at: http://irsa.ipac.caltech.edu/applications/Gator/
#*USNO provides optical magnitudes, if they exist. Make sure to retrieve tables, not images. This site can be accessed at: http://www.nofs.navy.mil/data/FchPix/
#*Skyview provides images at other wavelengths. This site can be accessed at: http://skyview.gsfc.nasa.gov

== [[Writing the Research Paper for IC2118]] ==

== [[Information for Oil City High School meetings]] ==

The next meeting should be at 4 0'clock pm wednesday the 25th at matt walentosky's house. (Specifically for people going to California, so we can plan out events and travel...)

== Proposal for Radio Observations of T-Tauri Candidates in IC2118 ==

'''''SUPPORT INFORMATION - T-Tauri Candidates Emit Radio'''''

'''FROM MATT WALENTOSKY 6/30/2007 @ 2:30 PM - Two Abstracts Below'''

Title: Radio emission from pre-main-sequence stars
Authors: Skinner, Stephen Lee

Affiliation: AA(Colorado Univ., Boulder.)
Publication: Ph.D. Thesis Colorado Univ., Boulder.
Publication Date: 01/1992
Category: Space Radiation
Origin: STI

Abstract
This study focuses on the properties and physical origin of radio continuum emission from pre-main-sequence (PMS) stars. These are young stars, typically less than a few million years old, and are still in a phase of gravitational contraction that will ultimately be halted by the onset of hydrogen burning in their cores. First, I address the question of the origin of centimeter continuum emission in intermediate mass (approx. equal to 3-20 solar mass) PMS stars, the so-called 'Herbig Ae/Be stars'. A high-sensitivity radio survey of 57 such stars was undertaken using the Very Large Array and Australia Telescope, resulting in the detection of twelve stars. These observations provide a homogenous data base consisting of information on source sizes, radio luminosities, variability timescales, circular polarization, and spectral energy distributions in the wavelength range 2-20 cm. Using these data along with previously published spectroscopy, I conclude that centimeter radio emission from Herbig Ae/Be stars is predominantly thermal and in many cases wind-related. An unexpected result of the above program was the serendipitous detection of circularly polarized radio emission in the low mass (approx. equal to 1 solar mass) PMS star Hubble 4, a member of the class of 'weak-lined T Tauri stars' (WTTS). This provides some of the most convincing evidence to date for the existence of ordered magnetic fields in WTTS. In a second observing program, I have searched for evidence of cold (less than or equal to 50 K) circumstellar dust around WTTS, which might exist in the form of remnant disks. Of the sixteen WTTS that were observed in the wavelength range 450-1100 microns using the James Clerk Maxwell Telescope, only V836 Tau was detected. Its spectral energy distribution longward of 10 microns is consistent with that expected for a flat, axisymmetric circumstellar disk of mass approx. equal to 0.04 solar mass (= 42 Jupiter masses). This star may be a rare example of an object in which disk dispersal is underway, but not yet complete.

Title: Centimeter Radio Emission from Low-Mass Weak T Tauri Stars in Taurus-Auriga
Authors: Chiang, Eugene; Phillips, R. B.

Affiliation: AA(MIT), AB(MIT Haystack Observatory)
Publication: American Astronomical Society, 185th AAS Meeting, #48.08; Bulletin of the American Astronomical Society, Vol. 26, p.1388
Publication Date: 12/1994
Origin: AAS
Abstract Copyright:(c) 1994: American Astronomical Society

Abstract
We report on the results of a sensitive survey for lambda 3.6 cm radio emission from low-mass, weak T Tauri (WTT) stars in the Taurus-Auriga cloud complex. The target population consists of stars in the Herbig and Bell Catalog of spectral type K7 or later, and W(Hα ) <= 10 Angstroms. Of the 28 such stars surveyed using the Very Large Array down to detection thresholds of ~ 0.1 mJy, 7 (possibly 8) are observed to emit at strengths ranging from 0.1 to 2 mJy. Five of these young radio stars are newly discovered in our survey: V827 Tau and V710 Tau B are discovered to be relatively strong sources of mJy emission, while IW Tau, UX Tau B, and the possible detection LkHa 332-G1 form a new population of relatively weak emitters. Our radio survey and complementary surveys are pooled, and of 43 WTT stars K7 or later in Tau-Aur, 14 are now known to be radio emitters at lambda 6 and lambda 3.6 cm. Correlations between radio luminosity and other stellar parameters have been attempted but generally yield null results. Wide binarity (component separations in excess of 0''.13, 20 AU) appears unrelated to radio emission, as does spectral type. Furthermore, we find no convincing evidence for the extreme youth of radio stars, contrary to claims in the literature over the past decade. While we do find that radio-loud stars in our sample are formally younger than the radio-quiet stars by about 0.5 Myr, the reality of this relatively small age difference is highly suspect given uncertainties in the placement of these stars on the HR diagram. Moreover, Monte Carlo-type calculations involving distributing the stars on both the HR diagram and local CO gas density cast doubts on any differences between the radio stars and the general WTT population. We conclude that the age effect for low-mass radio WTT stars in Tau-Aur, if real, is much smaller than previous estimations by factors of 4-10. It is also possible centimeter wavelength surveys to date have still not properly described the radio luminosity function of low-mass WTT stars in Tau-Aur, and we urge future observations of these young stars with denser temporal coverage.

'''FROM SPUCK 6/30/2007 @ 2:35pm'''

I'm not sure that we will be able to use the GBT telescope in Green Bank for observations. The beamwidth at 10-15 GHz is probably too big.

Here are the specs for the GBT
Beamwidth (Table 3) Diffraction beamwidth (FWHM)
8 GHz it is 90",
20 GHz it is 36",
50 GHz it is 14"

Here are some questions from Sue Ann Heatherly at NRAO-Green Bank - One question I have is: is the resolving
power of the GBT sufficient to distinguish spatially between individual stars? Does large scale emission exist within the nebula that will confuse your results?

The VLA is probably the only instrument that can be used.

Synthesized Beamwidth (arcsec)depending on configuration
'''Nick Kelley 6/30/2007 3:00PM'''

400 cm 24.0" 80.0" 260.0" 850.0"

90 cm 6.0" 17.0" 56.0" 200.0"

20 cm 1.4 " 3.9" 12.5" 44.0"

6 cm 0.4 " 1.2""" 3.9 " 14.0 "

3.6 cm 0.24 " 0.7" 2.3 " 8.4 "

2 cm 0.14 " 0.4 " 1.2" 3.9"

1.3 cm 0.08 " 0.3" 0.9 "2.8"

0.7 cm 0.05 0.15 0.47 1
restricting my creativity

== T-Tauri/IC2118 Presentation for Astroblast 2007 ==

'''''Oil City High School students will be presenting August 11 at the Oil Region Astronomical Observatory'''''

FROM Nicholas James Kelley

Astroblast is amazin ... see pic

[[Image: astroblast1.JPG|center]]

Is their anyway we can attach the powerpoint here?? Matt W.

FROM SPUCK 7/25/07 - Matt ... we may be able to upload the powerpoint somewhere and then create a link to it. I'll need to check on this. Thanks, Mr. Spuck

== Monitoring T-Tauri Stars using the Perth Obervatory Telescope ==

'''Generating Light Curves'''

Using Perth Telescope

1. Check to see what the current sky conditions are by checking the live camera

http://www.perthobservatory.wa.gov.au//information/po_sky_camera.html

2. Check to see what the current weather conditions are

http://www.perthobservatory.wa.gov.au//information/po_weather.html

3. Weather forecast for Perth - http://www.bom.gov.au/weather/wa/

Log into telescope if conditions look good.

http://perthobservatory.org or http://202.72.190.15

EMAIL tspuck@hotmail.com for telescope access

== Chandra X-Ray Data ==

'''''To date no X-Ray data for IC2118 has been located.''''' If you know of any please contact Tim Spuck at tspuck@hotmail.com.

== Kitt Peak H-alpha Data ==

'''''In January of 2007 students from Oil City High School used the 0.9 Meter Telescope at Kitt peak to image regions of IC2118 in H-alpha.''''

Matt Walentosky, Nick Kelley, Sandy Weiser were the students who went!!!!!!

== USNO U, B, V, R, and I Data ==

'''''U, B, V, R, and I data will be used to generate more accurate SEDs for the T-Tauri Candidates'''''

== Outflows or Jets ==

'''''Visible outflows or jets are strong evidence that a candidate is indeed a T-Tauri Star'''''

== Multi Wavelength Composite Images ==

'''''Images and Methods'''''

FROM SPUCK 6/30/2007 See Image Below- IC 2118 3.6 µm (blue), 5.8 µm (green), 8.0 µm (red) tri-color composite generated using MaxIm DL. (By M. Heath, N. Kelley, P. Morton, M. Walentosky, S. Weiser – Oil City High School, Oil City, PA)

[[Image: spuckim3.JPG|center]]

Studying Young Stars

2007-08-06T11:49:52Z

Weehler:

=Brief philosophical notes from Luisa=
==Real science vs. textbook science==
*Science (history) as presented in textbooks may seem a never-ending series of right answers. Real science has a lot of dead ends as we struggle to find out what the ‘right answer’ is.
*Science problems in textbooks have well-defined problems, specific methods you’re supposed to use to solve them, and right (exact) answers (1.2 can be wrong when 1.3 is right). Real science is not quite “made up as you go along” but it may feel that way in the coming days. Different people approach the same problem in different ways, and many answers can be right (1.2 and 1.3 can both be right). '''The only way you know it’s the right answer is if you believe that everything you did to get there is right.'''

==Why should anyone care about young stars?==
*Understanding star formation includes understanding how planets form, including planets like Earth.
*Star formation is the "happening field" right now! TONS of new discoveries happening all the time, many driven by Spitzer.
*A friend who is the author of a popular college textbook told me that the chapter that she revises most frequently (particularly recently) in response to new developments is the star formation chapter.
*By doing this project, you are participating in the revolution!

=Young stars in general: Introduction to (low-mass) star formation=
==Overview==

{| cellpadding="1"
| [[Image:starformationcartoon.png]]
|''Cartoon from Greene, American Scientist, Jul-Aug 2001''
|}

Stars begin their life in a cloud of gas and dust (a nebula). Gravitational forces cause the nebula to start to condense (shrink). (a, b)

As the nebula shrinks, like a spinning skater pulling in her arms, it begins to spin more rapidly. The same physics ("conservation of angular momentum") means that the dust and gas in the nebula doesn't fall straight into the center; it falls onto a disk surrounding the central object, and from the disk, the matter falls onto the central object. The temperature at the center of the shrinking nebula rises due to increasing pressure and friction between the particles. The figure has this stage labeled as a "protostar", but for some astronomers, beginning at this stage, and until the star starts to turn H into He the object is still called a protostar. Since the protostar is still embedded in a thick cloud of gas and dust, it can only be detected in the infrared. (c)

When the protostar enters the next stage, labeled in the figure as the T Tauri stage, it’s still gaining mass and contracting slowly because material is still falling onto it, but it begins to eject gas in two giant gas jets, called bipolar flows. These jets and stellar winds eventually sweep away the envelope of gas still surrounding the protostar. In the surrounding disk protoplanets are beginning to form. (d)

Leftover material in the disk surrounding the star clumps together and undergoes many collisions until most of the material has been swept up by objects orbiting the star, such as planets, asteroids and comets. (e)

The star’s life so far has been governed by the continuous inward pressure of gravity. The gravitational pressure keeps compressing the gas into a smaller and smaller volume, making it hotter and hotter in the core. As soon as the temperature in the core of the protostar becomes great enough, nuclear fusion begins. When this nuclear fusion begins, finally the star has a way to "fight back" against gravity. So much energy is released in this reaction that it enables the star to "push back" with an outward radiation pressure that balances the inward push of gravity. The protostar is now a full-fledged star, fusing hydrogen into helium in its core. (f) The star will stay the same size until it runs out of nuclear fuel in the core (all of the hydrogen has been converted into helium). Then, the pressure from gravity takes over again, pushing in on the star.

'''Movies!''' -- The sequence of star formation described above can also be visualized with some movies created by the Spitzer public affairs group. First [http://ipac.jpl.nasa.gov/web_movies/pa/ssc2004-08v2_qt4.mov form grains in the disk] (note that the grains are distinctly green because they are olivine, like green sand beaches in Hawaii, and the grains get covered in ice). Then [http://ipac.jpl.nasa.gov/web_movies/pa/ssc2004-08v3_qt4.mov form a planet that clears a gap in the disk]. Next,
[http://ipac.jpl.nasa.gov/web_movies/pa/ssc2004-22v2_qt4.mov this is what happens when you form many planets at once] - there are many gaps formed at once. (Where else in the Solar System have you seen this physics before? [http://saturn.jpl.nasa.gov/multimedia/images/raw/raw-images-details.cfm?feiImageID=112884 stumped?]) Of course, it can get crowded in these protoplanetary disks, so [http://ipac.jpl.nasa.gov/web_movies/pa/ssc2004-17v1_qt4.mov occasionally stuff hits each other], creating a second generation of dust. When this happens, [http://ipac.jpl.nasa.gov/web_movies/pa/ssc2004-17v2_qt4.mov the dust gets smeared out into a ring], and there can be many collisions that create transient dust, e.g., dust that comes and goes depending on when you look. Finally, [http://ipac.jpl.nasa.gov/web_movies/pa/ssc2004-22v1_qt4.mov you end up with a fully-fledged planetary system], but you have the remnants of the protoplanetary disk left in the outer reaches of the system. (What do we call this in our Solar System? hint: not the asteroid belt. out farther than that -- see in the movie, we go out well past a "jupiter".)

==More in depth, and SEDs==
The cartoon above is the version of this information that is appropriate for an "educated person from the general public." Now, let's look at this same story again, but using the kinds of plots and information used by professional astronomers.

A spectral energy distribution (SED) is a graph of the energy emitted by an object (any object) as a function of different wavelengths. There are some example SEDs of young stellar objects below. Astronomers originally used the slope of the SEDs for protostars between about 2 and 20 microns quite literally to '''define''' different classes. (This is similar to the process that astronomers such as [http://en.wikipedia.org/wiki/Annie_Jump_Cannon Annie Jump Cannon] followed when they originally classified stars -- astronomers start by putting similar objects together, and through this process, eventually deeper physical understanding follows.) These stages that were defined based on the SED slope are dubbed Class I, II, and III. As we learned more, we created another class called "flat" between Class I and II. The very earliest stages, ones where there are no 2 micron observations at all, are the Class 0s. These classifications more or less match up to the overall sequence of events described by the cartoon above.

{| cellpadding="2"
| [[Image:class0.png]]
| This SED corresponds to the earliest, most embedded phase, called Class 0. The x-axis is wavelength, and the y-axis is energy. It looks like a cold black body; all of the emission comes from the dust and gas in the cloud. The mass of the envelope surrounding the star is more than half a solar mass. The age of the object is about 10,000 (<math>10^4</math>) years. This is thoguht to be the main accretion phase, where most of the mass of the object is accreted from the cloud.
|}

{| cellpadding="2"
| [[Image:classI.png]]
|This Class I SED is the next stage. See how now there is a warmer blackbody corresponding to the central object, but most of the energy is coming from the dusty cocoon around the star. The "bite" that is at about 10 microns tells us that there are silicates (beach sand) in the dust around the star. About this time is when the rate of accretion slows. The envelope is now about 0.1 Msun. The age of the object is about <math>10^5</math> years.
|}

{| cellpadding="2"
| [[Image:classII.png]]
|This Class II SED is also a kind of object also known as a young T Tauri ("Classical T Tauri"). Most of the energy is coming from the central object (warm blackbody), but there is a little emission from the disk; the disk is optically thick. The disk mass now is very roughly 0.01 Msun. The age of the object is about <math>10^6</math> years. Several of the objects in IC 2118 look like this.
|}

{| cellpadding="2"
| [[Image:classIII.png]]
|This SED is of the last stage, Class III; this is an older T Tauri ("Weak-Lined T Tauri"). The protostar still has a little dust left around it. The disk now is optically thin. The disk mass is very roughly 0.003 Msun. The object is about <math>10^7</math> years old.
Several of the objects in IC2118 look like this.
|}

==Several important notes about this classification scheme==
*As the ages mentioned above suggest, this sequence of Class 0 to I to II to III is often interpreted as an age sequence, so a Class I object is younger than a Class II object, etc. Some of the most recent evidence suggests that maybe the connection to age is not as secure as we have been thinking! So let me just reiterate: the Class of each object is ''defined'' by the slope of the SED. The physical interpretation of the classification definition is degree of embeddedness, e.g., Class 0s are still buried deep within their natal cloud, and Class IIIs have freed themselves. '''The interpretation of the Class as an age may change.'''
*This process strictly only applies to low-mass stars. High-mass stars might very well do this, only faster. We just don’t know yet for sure.
*Class 0s are the the hardest to "catch in the act", from which we infer that they are the shortest lived. Not too long ago, the list of all of the Class 0s known could fit on one page. Spitzer is changing that. Class Is are also being found by Spitzer in abundance.
*Class 0s used to be defined as "undetectable in IR." Even before Spitzer, deeper integrations forced a change in that definition.
*Although the story seems nice and well-defined, even before Spitzer, Class IIs and IIIs have been found at the same ages, e.g., some stars lose their disks very quickly, and some hold on to them for a long time. Now with Spitzer, we're muddying the waters even more.
*A current major question in star formation is the how and why of this process.
*It’s not clear whether Class 0s and Is are found at the same age – until very recently, too few of them were known, and getting an age for them is tough.
*[http://www.spitzer.caltech.edu/Media/releases/ssc2004-17/release.shtml This press release] talks about A stars, which are a little massive for Class 0/I/II/III, but the confusion in disk clearing timescales is vividly displayed there.
*We can be fooled! You can imagine that a Class III that is edge-on might look like a Class II. It could be that some things we think are the youngest protostars are actually just edge-on older things. This is also one of the current burning questions.
*Most people still use the series Class 0-I-II-III to mean a series of youngest to oldest, but it’s important to remember all of these uncertainties.

The details of the shape of the SED can tell us about the disk structure. Dips and wiggles in the SED may suggest, e.g., that there is no (or little) dust near the star, just further out. (see [http://www.spitzer.caltech.edu/Media/releases/ssc2004-08/ssc2004-08c.shtml this graphic from the SSC press release archive].)

=Young stars in general: Finding the cluster members=

Spitzer is so sensitive that it easily sees things at the far reaches of the Universe with only a few seconds' integration. When studying clusters of stars, not just with Spitzer, one of the first major goals is to figure out which objects are truly cluster members and which are not. [[Media:findingclustermembers.pdf| This pdf file]] has a discussion of how to find members of young clusters in general.

Talk:Studying Young Stars

2007-08-06T11:48:48Z

Weehler:

Real science v. textbook science bullett 2:
''''''this would make an excellent link to the previous comment I made on bullet 1 of Photometry--[[User:Weehler|Weehler]] 06:37, 3 August 2007 (PDT)''''''

Studying Young Stars

2007-08-03T13:37:12Z

Weehler: possible link

=Brief philosophical notes from Luisa=
==Real science vs. textbook science==
*Science (history) as presented in textbooks may seem a never-ending series of right answers. Real science has a lot of dead ends as we struggle to find out what the ‘right answer’ is.
*Science problems in textbooks have well-defined problems, specific methods you’re supposed to use to solve them, and right (exact) answers (1.2 can be wrong when 1.3 is right). Real science is not quite “made up as you go along” but it may feel that way in the coming days. Different people approach the same problem in different ways, and many answers can be right (1.2 and 1.3 can both be right). '''The only way you know it’s the right answer is if you believe that everything you did to get there is right.'''''''''this would make an excellent link to the previous comment I made on bullet 1 of Photometry--[[User:Weehler|Weehler]] 06:37, 3 August 2007 (PDT)''''''

==Why should anyone care about young stars?==
*Understanding star formation includes understanding how planets form, including planets like Earth.
*Star formation is the "happening field" right now! TONS of new discoveries happening all the time, many driven by Spitzer.
*A friend who is the author of a popular college textbook told me that the chapter that she revises most frequently (particularly recently) in response to new developments is the star formation chapter.
*By doing this project, you are participating in the revolution!

=Young stars in general: Introduction to (low-mass) star formation=
==Overview==

{| cellpadding="1"
| [[Image:starformationcartoon.png]]
|''Cartoon from Greene, American Scientist, Jul-Aug 2001''
|}

Stars begin their life in a cloud of gas and dust (a nebula). Gravitational forces cause the nebula to start to condense (shrink). (a, b)

As the nebula shrinks, like a spinning skater pulling in her arms, it begins to spin more rapidly. The same physics ("conservation of angular momentum") means that the dust and gas in the nebula doesn't fall straight into the center; it falls onto a disk surrounding the central object, and from the disk, the matter falls onto the central object. The temperature at the center of the shrinking nebula rises due to increasing pressure and friction between the particles. The figure has this stage labeled as a "protostar", but for some astronomers, beginning at this stage, and until the star starts to turn H into He the object is still called a protostar. Since the protostar is still embedded in a thick cloud of gas and dust, it can only be detected in the infrared. (c)

When the protostar enters the next stage, labeled in the figure as the T Tauri stage, it’s still gaining mass and contracting slowly because material is still falling onto it, but it begins to eject gas in two giant gas jets, called bipolar flows. These jets and stellar winds eventually sweep away the envelope of gas still surrounding the protostar. In the surrounding disk protoplanets are beginning to form. (d)

Leftover material in the disk surrounding the star clumps together and undergoes many collisions until most of the material has been swept up by objects orbiting the star, such as planets, asteroids and comets. (e)

The star’s life so far has been governed by the continuous inward pressure of gravity. The gravitational pressure keeps compressing the gas into a smaller and smaller volume, making it hotter and hotter in the core. As soon as the temperature in the core of the protostar becomes great enough, nuclear fusion begins. When this nuclear fusion begins, finally the star has a way to "fight back" against gravity. So much energy is released in this reaction that it enables the star to "push back" with an outward radiation pressure that balances the inward push of gravity. The protostar is now a full-fledged star, fusing hydrogen into helium in its core. (f) The star will stay the same size until it runs out of nuclear fuel in the core (all of the hydrogen has been converted into helium). Then, the pressure from gravity takes over again, pushing in on the star.

'''Movies!''' -- The sequence of star formation described above can also be visualized with some movies created by the Spitzer public affairs group. First [http://ipac.jpl.nasa.gov/web_movies/pa/ssc2004-08v2_qt4.mov form grains in the disk] (note that the grains are distinctly green because they are olivine, like green sand beaches in Hawaii, and the grains get covered in ice). Then [http://ipac.jpl.nasa.gov/web_movies/pa/ssc2004-08v3_qt4.mov form a planet that clears a gap in the disk]. Next,
[http://ipac.jpl.nasa.gov/web_movies/pa/ssc2004-22v2_qt4.mov this is what happens when you form many planets at once] - there are many gaps formed at once. (Where else in the Solar System have you seen this physics before? [http://saturn.jpl.nasa.gov/multimedia/images/raw/raw-images-details.cfm?feiImageID=112884 stumped?]) Of course, it can get crowded in these protoplanetary disks, so [http://ipac.jpl.nasa.gov/web_movies/pa/ssc2004-17v1_qt4.mov occasionally stuff hits each other], creating a second generation of dust. When this happens, [http://ipac.jpl.nasa.gov/web_movies/pa/ssc2004-17v2_qt4.mov the dust gets smeared out into a ring], and there can be many collisions that create transient dust, e.g., dust that comes and goes depending on when you look. Finally, [http://ipac.jpl.nasa.gov/web_movies/pa/ssc2004-22v1_qt4.mov you end up with a fully-fledged planetary system], but you have the remnants of the protoplanetary disk left in the outer reaches of the system. (What do we call this in our Solar System? hint: not the asteroid belt. out farther than that -- see in the movie, we go out well past a "jupiter".)

==More in depth, and SEDs==
The cartoon above is the version of this information that is appropriate for an "educated person from the general public." Now, let's look at this same story again, but using the kinds of plots and information used by professional astronomers.

A spectral energy distribution (SED) is a graph of the energy emitted by an object (any object) as a function of different wavelengths. There are some example SEDs of young stellar objects below. Astronomers originally used the slope of the SEDs for protostars between about 2 and 20 microns quite literally to '''define''' different classes. (This is similar to the process that astronomers such as [http://en.wikipedia.org/wiki/Annie_Jump_Cannon Annie Jump Cannon] followed when they originally classified stars -- astronomers start by putting similar objects together, and through this process, eventually deeper physical understanding follows.) These stages that were defined based on the SED slope are dubbed Class I, II, and III. As we learned more, we created another class called "flat" between Class I and II. The very earliest stages, ones where there are no 2 micron observations at all, are the Class 0s. These classifications more or less match up to the overall sequence of events described by the cartoon above.

{| cellpadding="2"
| [[Image:class0.png]]
| This SED corresponds to the earliest, most embedded phase, called Class 0. The x-axis is wavelength, and the y-axis is energy. It looks like a cold black body; all of the emission comes from the dust and gas in the cloud. The mass of the envelope surrounding the star is more than half a solar mass. The age of the object is about 10,000 (<math>10^4</math>) years. This is thoguht to be the main accretion phase, where most of the mass of the object is accreted from the cloud.
|}

{| cellpadding="2"
| [[Image:classI.png]]
|This Class I SED is the next stage. See how now there is a warmer blackbody corresponding to the central object, but most of the energy is coming from the dusty cocoon around the star. The "bite" that is at about 10 microns tells us that there are silicates (beach sand) in the dust around the star. About this time is when the rate of accretion slows. The envelope is now about 0.1 Msun. The age of the object is about <math>10^5</math> years.
|}

{| cellpadding="2"
| [[Image:classII.png]]
|This Class II SED is also a kind of object also known as a young T Tauri ("Classical T Tauri"). Most of the energy is coming from the central object (warm blackbody), but there is a little emission from the disk; the disk is optically thick. The disk mass now is very roughly 0.01 Msun. The age of the object is about <math>10^6</math> years. Several of the objects in IC 2118 look like this.
|}

{| cellpadding="2"
| [[Image:classIII.png]]
|This SED is of the last stage, Class III; this is an older T Tauri ("Weak-Lined T Tauri"). The protostar still has a little dust left around it. The disk now is optically thin. The disk mass is very roughly 0.003 Msun. The object is about <math>10^7</math> years old.
Several of the objects in IC2118 look like this.
|}

==Several important notes about this classification scheme==
*As the ages mentioned above suggest, this sequence of Class 0 to I to II to III is often interpreted as an age sequence, so a Class I object is younger than a Class II object, etc. Some of the most recent evidence suggests that maybe the connection to age is not as secure as we have been thinking! So let me just reiterate: the Class of each object is ''defined'' by the slope of the SED. The physical interpretation of the classification definition is degree of embeddedness, e.g., Class 0s are still buried deep within their natal cloud, and Class IIIs have freed themselves. '''The interpretation of the Class as an age may change.'''
*This process strictly only applies to low-mass stars. High-mass stars might very well do this, only faster. We just don’t know yet for sure.
*Class 0s are the the hardest to "catch in the act", from which we infer that they are the shortest lived. Not too long ago, the list of all of the Class 0s known could fit on one page. Spitzer is changing that. Class Is are also being found by Spitzer in abundance.
*Class 0s used to be defined as "undetectable in IR." Even before Spitzer, deeper integrations forced a change in that definition.
*Although the story seems nice and well-defined, even before Spitzer, Class IIs and IIIs have been found at the same ages, e.g., some stars lose their disks very quickly, and some hold on to them for a long time. Now with Spitzer, we're muddying the waters even more.
*A current major question in star formation is the how and why of this process.
*It’s not clear whether Class 0s and Is are found at the same age – until very recently, too few of them were known, and getting an age for them is tough.
*[http://www.spitzer.caltech.edu/Media/releases/ssc2004-17/release.shtml This press release] talks about A stars, which are a little massive for Class 0/I/II/III, but the confusion in disk clearing timescales is vividly displayed there.
*We can be fooled! You can imagine that a Class III that is edge-on might look like a Class II. It could be that some things we think are the youngest protostars are actually just edge-on older things. This is also one of the current burning questions.
*Most people still use the series Class 0-I-II-III to mean a series of youngest to oldest, but it’s important to remember all of these uncertainties.

The details of the shape of the SED can tell us about the disk structure. Dips and wiggles in the SED may suggest, e.g., that there is no (or little) dust near the star, just further out. (see [http://www.spitzer.caltech.edu/Media/releases/ssc2004-08/ssc2004-08c.shtml this graphic from the SSC press release archive].)

=Young stars in general: Finding the cluster members=

Spitzer is so sensitive that it easily sees things at the far reaches of the Universe with only a few seconds' integration. When studying clusters of stars, not just with Spitzer, one of the first major goals is to figure out which objects are truly cluster members and which are not. [[Media:findingclustermembers.pdf| This pdf file]] has a discussion of how to find members of young clusters in general.

Mosaics

2007-08-03T13:11:14Z

Weehler: comment

=What is a mosaic?='''really good!--[[User:Weehler|Weehler]] 06:11, 3 August 2007 (PDT)'''

A mosaic is a larger picture composed of many smaller pictures. This way, you can get an image of a region much larger than the field of view of a single frame.

Here is an example taken from real life. We'd like to take a picture of some friendly people standing in front of the 200-inch Hale Telescope at Palomar Observatory. But our camera has a relatively small field of view, and we can either take a picture of the people, or the telescope, but not both at the same time. So we take two pictures:

[[Image:pal200top.jpg]]

[[Image:pal200bottom.jpg]]

Then, after downloading the images from our digital camera, we use the computer to combine them together into one larger picture, covering the entire region we're interested in.

[[Image:pal200mos.jpg]]

Note here too that the mosaicking process compensated for distortion in the camera -- the edges of the image are no longer a straight line.

Real astronomical mosaickers do this too -- they combine smaller images into a larger image, and compensate for distortion in the camera. The distortion is just (usually) much smaller than what is seen here.

=Spitzer mosaics=

Most of Spitzer's cameras have a field of view that is 5 arcminutes on a side. But, by design, Spitzer is really, really good at covering huge areas. So by taking many, many individual frames, we can combine them all together into huge maps -- I've worked on maps that are more than 10 square degrees, containing about 21,000 individual frames!

In Spitzer jargon, an individual frame that has just been taken is called a '''Data Collection Event (DCE)''', or "raw data." After it is processed through the SSC's pipeline, it is called '''Basic Calibrated Data (BCD)'''. Many BCDs are combined into a mosaic, or a '''post-BCD product'''.

Photometry

2007-08-03T13:09:30Z

Weehler: "yanked out of Varoujan's"

yanked out of my intro talk. needs to be reformulated, fixed, integrated properly.

*Every astronomer does things a little differently. (It’s kind of amazing that any two astronomers working on the same object and the same wavelength ever get the same answer.) Everyone does what they think is right. '''Perhaps we should add something to the effect that, That doesn't mean they're making it up; it's just the creative nature of the science process that you truly are starting from scratch.? --Don't want people to think ANYTHING is okay.--[[User:Weehler|Weehler]] 05:58, 3 August 2007 (PDT)'''
*The shape of the point-source pattern is called the “point spread function (PSF),” or sometimes the “point response function (PRF).” (Technically these two things are subtly different, but never mind that for now.) The PSF is HUGE, and there is a lot of flux surprisingly far from the star that needs to be included. '''include a definition of flux here? or is it enough that they can access the definition from the Units section?--[[User:Weehler|Weehler]] 06:02, 3 August 2007 (PDT)'''
*One can measure fluxes from point sources (like stars) in two ways: aperture photometry or PSF fitting.
*Aperture photometry measures all of the flux within a (usually circular) aperture centered on the star, minus the flux in an annulus around the aperture. This is quick, but not necessarily accurate, and can lead to large errors (especially if the background is complicated), and is essentially impossible in crowded fields. One must take into account fractional pixels within the aperture (which matters particularly when the units are MJy/sr''' why does it matter particularly with these units?--[[User:Weehler|Weehler]] 06:04, 3 August 2007 (PDT)'''). Usually one needs to apply an aperture correction to correct for the ‘missing’ flux outside the aperture.
*PSF fitting takes the basic shape of the PSF and matches it to the point source, thereby taking into account the fluxes at large distances from the star, as well as ignoring complicated structure in the background and other nearby point sources.
*MOPEX does both aperture and PSF fitting, and understands the units of Spitzer images. In practice with MOPEX, one needs to use aperture photometry for the brightest stars and PSF fitting for the rest. And, as of July 2007, one ought to use MOPEX aperture photometry for IRAC, and either aperture or PRF fitting for MIPS - this has to do with how well-sampled the PRFs are in the various channels.

yanked out of Varoujan's GAVRT ''''''what is GAVRT? I don't know about any of this : ) --[[User:Weehler|Weehler]] 06:09, 3 August 2007 (PDT)''''''stuff. needs to be reformulated, fixed, integrated properly.

Photometry is the measurement of the brightness of an object. The brightness that is recorded on an electronic detector (or any kind of detector) is a combination of the brightness of the source plus the brightness of the background that the source is on.

In the case of the center of the galaxy in our project, a great deal of the "background" is in fact from the surrounding host galaxy. But what we are interested in is only the light from the center and not from the rest of the galaxy. To account for this, we will determine how much light is coming from where the center of the AGN is, and then compare it to how much light is coming from near the center of the AGN. We assume that the "background" near the AGN is the same as the background right on the AGN, so by subtracting the two, we are left with only the brightness of the center of the AGN.

Photometry

2007-08-03T13:04:58Z

Weehler: bullet 4

yanked out of my intro talk. needs to be reformulated, fixed, integrated properly.

*Every astronomer does things a little differently. (It’s kind of amazing that any two astronomers working on the same object and the same wavelength ever get the same answer.) Everyone does what they think is right. '''Perhaps we should add something to the effect that, That doesn't mean they're making it up; it's just the creative nature of the science process that you truly are starting from scratch.? --Don't want people to think ANYTHING is okay.--[[User:Weehler|Weehler]] 05:58, 3 August 2007 (PDT)'''
*The shape of the point-source pattern is called the “point spread function (PSF),” or sometimes the “point response function (PRF).” (Technically these two things are subtly different, but never mind that for now.) The PSF is HUGE, and there is a lot of flux surprisingly far from the star that needs to be included. '''include a definition of flux here? or is it enough that they can access the definition from the Units section?--[[User:Weehler|Weehler]] 06:02, 3 August 2007 (PDT)'''
*One can measure fluxes from point sources (like stars) in two ways: aperture photometry or PSF fitting.
*Aperture photometry measures all of the flux within a (usually circular) aperture centered on the star, minus the flux in an annulus around the aperture. This is quick, but not necessarily accurate, and can lead to large errors (especially if the background is complicated), and is essentially impossible in crowded fields. One must take into account fractional pixels within the aperture (which matters particularly when the units are MJy/sr''' why does it matter particularly with these units?--[[User:Weehler|Weehler]] 06:04, 3 August 2007 (PDT)'''). Usually one needs to apply an aperture correction to correct for the ‘missing’ flux outside the aperture.
*PSF fitting takes the basic shape of the PSF and matches it to the point source, thereby taking into account the fluxes at large distances from the star, as well as ignoring complicated structure in the background and other nearby point sources.
*MOPEX does both aperture and PSF fitting, and understands the units of Spitzer images. In practice with MOPEX, one needs to use aperture photometry for the brightest stars and PSF fitting for the rest. And, as of July 2007, one ought to use MOPEX aperture photometry for IRAC, and either aperture or PRF fitting for MIPS - this has to do with how well-sampled the PRFs are in the various channels.

yanked out of Varoujan's GAVRT stuff. needs to be reformulated, fixed, integrated properly.

Photometry is the measurement of the brightness of an object. The brightness that is recorded on an electronic detector (or any kind of detector) is a combination of the brightness of the source plus the brightness of the background that the source is on.

In the case of the center of the galaxy in our project, a great deal of the "background" is in fact from the surrounding host galaxy. But what we are interested in is only the light from the center and not from the rest of the galaxy. To account for this, we will determine how much light is coming from where the center of the AGN is, and then compare it to how much light is coming from near the center of the AGN. We assume that the "background" near the AGN is the same as the background right on the AGN, so by subtracting the two, we are left with only the brightness of the center of the AGN.

Photometry

2007-08-03T13:02:05Z

Weehler: bullet 2

yanked out of my intro talk. needs to be reformulated, fixed, integrated properly.

*Every astronomer does things a little differently. (It’s kind of amazing that any two astronomers working on the same object and the same wavelength ever get the same answer.) Everyone does what they think is right. '''Perhaps we should add something to the effect that, That doesn't mean they're making it up; it's just the creative nature of the science process that you truly are starting from scratch.? --Don't want people to think ANYTHING is okay.--[[User:Weehler|Weehler]] 05:58, 3 August 2007 (PDT)'''
*The shape of the point-source pattern is called the “point spread function (PSF),” or sometimes the “point response function (PRF).” (Technically these two things are subtly different, but never mind that for now.) The PSF is HUGE, and there is a lot of flux surprisingly far from the star that needs to be included. '''include a definition of flux here? or is it enough that they can access the definition from the Units section?--[[User:Weehler|Weehler]] 06:02, 3 August 2007 (PDT)'''
*One can measure fluxes from point sources (like stars) in two ways: aperture photometry or PSF fitting.
*Aperture photometry measures all of the flux within a (usually circular) aperture centered on the star, minus the flux in an annulus around the aperture. This is quick, but not necessarily accurate, and can lead to large errors (especially if the background is complicated), and is essentially impossible in crowded fields. One must take into account fractional pixels within the aperture (which matters particularly when the units are MJy/sr). Usually one needs to apply an aperture correction to correct for the ‘missing’ flux outside the aperture.
*PSF fitting takes the basic shape of the PSF and matches it to the point source, thereby taking into account the fluxes at large distances from the star, as well as ignoring complicated structure in the background and other nearby point sources.
*MOPEX does both aperture and PSF fitting, and understands the units of Spitzer images. In practice with MOPEX, one needs to use aperture photometry for the brightest stars and PSF fitting for the rest. And, as of July 2007, one ought to use MOPEX aperture photometry for IRAC, and either aperture or PRF fitting for MIPS - this has to do with how well-sampled the PRFs are in the various channels.

yanked out of Varoujan's GAVRT stuff. needs to be reformulated, fixed, integrated properly.

Photometry is the measurement of the brightness of an object. The brightness that is recorded on an electronic detector (or any kind of detector) is a combination of the brightness of the source plus the brightness of the background that the source is on.

In the case of the center of the galaxy in our project, a great deal of the "background" is in fact from the surrounding host galaxy. But what we are interested in is only the light from the center and not from the rest of the galaxy. To account for this, we will determine how much light is coming from where the center of the AGN is, and then compare it to how much light is coming from near the center of the AGN. We assume that the "background" near the AGN is the same as the background right on the AGN, so by subtracting the two, we are left with only the brightness of the center of the AGN.

Photometry

2007-08-03T12:59:52Z

Weehler:

yanked out of my intro talk. needs to be reformulated, fixed, integrated properly.

*Every astronomer does things a little differently. (It’s kind of amazing that any two astronomers working on the same object and the same wavelength ever get the same answer.) Everyone does what they think is right. '''Perhaps we should add something to the effect that, That doesn't mean they're making it up; it's just the creative nature of the science process that you truly are starting from scratch.? --Don't want people to think ANYTHING is okay.--[[User:Weehler|Weehler]] 05:58, 3 August 2007 (PDT)'''
*The shape of the point-source pattern is called the “point spread function (PSF),” or sometimes the “point response function (PRF).” (Technically these two things are subtly different, but never mind that for now.) The PSF is HUGE, and there is a lot of flux surprisingly far from the star that needs to be included. include a definition of flux here? or is it enough that they can access definitionfrom
*One can measure fluxes from point sources (like stars) in two ways: aperture photometry or PSF fitting.
*Aperture photometry measures all of the flux within a (usually circular) aperture centered on the star, minus the flux in an annulus around the aperture. This is quick, but not necessarily accurate, and can lead to large errors (especially if the background is complicated), and is essentially impossible in crowded fields. One must take into account fractional pixels within the aperture (which matters particularly when the units are MJy/sr). Usually one needs to apply an aperture correction to correct for the ‘missing’ flux outside the aperture.
*PSF fitting takes the basic shape of the PSF and matches it to the point source, thereby taking into account the fluxes at large distances from the star, as well as ignoring complicated structure in the background and other nearby point sources.
*MOPEX does both aperture and PSF fitting, and understands the units of Spitzer images. In practice with MOPEX, one needs to use aperture photometry for the brightest stars and PSF fitting for the rest. And, as of July 2007, one ought to use MOPEX aperture photometry for IRAC, and either aperture or PRF fitting for MIPS - this has to do with how well-sampled the PRFs are in the various channels.

yanked out of Varoujan's GAVRT stuff. needs to be reformulated, fixed, integrated properly.

Photometry is the measurement of the brightness of an object. The brightness that is recorded on an electronic detector (or any kind of detector) is a combination of the brightness of the source plus the brightness of the background that the source is on.

In the case of the center of the galaxy in our project, a great deal of the "background" is in fact from the surrounding host galaxy. But what we are interested in is only the light from the center and not from the rest of the galaxy. To account for this, we will determine how much light is coming from where the center of the AGN is, and then compare it to how much light is coming from near the center of the AGN. We assume that the "background" near the AGN is the same as the background right on the AGN, so by subtracting the two, we are left with only the brightness of the center of the AGN.

Photometry

2007-08-03T12:58:15Z

Weehler: bullet 1

yanked out of my intro talk. needs to be reformulated, fixed, integrated properly.

*Every astronomer does things a little differently. (It’s kind of amazing that any two astronomers working on the same object and the same wavelength ever get the same answer.) Everyone does what they think is right. '''Perhaps we should add something to the effect that, That doesn't mean they're making it up; it's just the creative nature of the science process that you truly are starting from scratch.? --Don't want people to think ANYTHING is okay.--[[User:Weehler|Weehler]] 05:58, 3 August 2007 (PDT)'''
*The shape of the point-source pattern is called the “point spread function (PSF),” or sometimes the “point response function (PRF).” (Technically these two things are subtly different, but never mind that for now.) The PSF is HUGE, and there is a lot of flux surprisingly far from the star that needs to be included.
*One can measure fluxes from point sources (like stars) in two ways: aperture photometry or PSF fitting.
*Aperture photometry measures all of the flux within a (usually circular) aperture centered on the star, minus the flux in an annulus around the aperture. This is quick, but not necessarily accurate, and can lead to large errors (especially if the background is complicated), and is essentially impossible in crowded fields. One must take into account fractional pixels within the aperture (which matters particularly when the units are MJy/sr). Usually one needs to apply an aperture correction to correct for the ‘missing’ flux outside the aperture.
*PSF fitting takes the basic shape of the PSF and matches it to the point source, thereby taking into account the fluxes at large distances from the star, as well as ignoring complicated structure in the background and other nearby point sources.
*MOPEX does both aperture and PSF fitting, and understands the units of Spitzer images. In practice with MOPEX, one needs to use aperture photometry for the brightest stars and PSF fitting for the rest. And, as of July 2007, one ought to use MOPEX aperture photometry for IRAC, and either aperture or PRF fitting for MIPS - this has to do with how well-sampled the PRFs are in the various channels.

yanked out of Varoujan's GAVRT stuff. needs to be reformulated, fixed, integrated properly.

Photometry is the measurement of the brightness of an object. The brightness that is recorded on an electronic detector (or any kind of detector) is a combination of the brightness of the source plus the brightness of the background that the source is on.

In the case of the center of the galaxy in our project, a great deal of the "background" is in fact from the surrounding host galaxy. But what we are interested in is only the light from the center and not from the rest of the galaxy. To account for this, we will determine how much light is coming from where the center of the AGN is, and then compare it to how much light is coming from near the center of the AGN. We assume that the "background" near the AGN is the same as the background right on the AGN, so by subtracting the two, we are left with only the brightness of the center of the AGN.

IC 2118 Current Research Activities

2007-08-03T12:50:00Z

Weehler: para 1

=Some more of the basics behind what we are doing!=

[[Image:L1014.jpg|right]]
Here is an image of an object called L1014. This object used to be known as a "starless core." On the left is a picture in visible light, and you can see why people thought it was starless. On the right is a picture from Spitzer clearly revealing a baby star inside (it's the bright red thing)!'''really good illustration!--[[User:Weehler|Weehler]] 05:50, 3 August 2007 (PDT)'''

When the protostar enters the next stage, labeled in the figure as the T Tauri stage, it’s still gaining mass and contracting slowly because material is still falling onto it, but it begins to eject gas in two giant gas jets, called bipolar flows. These jets and stellar winds eventually sweep away the envelope of gas still surrounding the protostar. In the surrounding disk protoplanets are beginning to form. (d) Astronomers call this phase the T Tauri phase, but some objects with jets are also "Class I"; T Tauris are also either "Class II" (for ones with thick disks) or "Class III" (for ones with thinner disks).
'''What do (d) and (e) refer to? Where are (a)-(c)? Are there illustrations that go here?--[[User:Weehler|Weehler]] 05:37, 3 August 2007 (PDT)'''
Leftover material in the disk surrounding the star clumps together and undergoes many collisions until most of the material has been swept up by objects orbiting the star, such as planets, asteroids and comets. (e)

The star’s life so far has been governed by the continuous inward pressure of gravity. The gravitational pressure keeps compressing the gas into a smaller and smaller volume, making it hotter and hotter in the core. As soon as the temperature in the core of the protostar becomes great enough, about 1,000,000 K, nuclear fusion begins. Nuclear fusion is the process in which small atomic nuclei combine to make larger atomic nuclei, releasing lots of energy. Inside the star, the gravitational pressure eventually pushes 4 protons and 2 electrons '''Are electrons involved in fusion?--[[User:Weehler|Weehler]] 05:48, 3 August 2007 (PDT)'''so close together that they fuse together to make helium. When this nuclear fusion begins, finally the star has a way to "fight back" against gravity. So much energy is released in this reaction that it enables the star to "push back" with an outward radiation pressure that balances the inward push of gravity. The protostar is now a full-fledged star, fusing hydrogen into helium in its core. (f) The star will stay the same size until it runs out of nuclear fuel in the core (all of the hydrogen has been converted into helium). Then, the pressure from gravity takes over again, pushing in on the star.

Star formation can be triggered by the collapse of large clouds of gas and dust. Sometimes the clouds collapse all by themselves (due to gravitational forces within the cloud) and sometimes they collapse because they’ve been pushed by the radiation and winds from stars that have already been formed. This seems to be the reason behind the star formation in IC 2118. It seems to have been triggered by some combination of forces from the main Orion Nebula Cluster (the fuzzy patch in the sword) and Rigel.

=Color-Color Plots=
[[Image:irac-colorcolor.png|left]]

The magnitude scale measures the brightness of a star. The system for assigning magnitude numbers was developed in ancient times. The '''brighter''' the star is, the '''lower''' its magnitude.

An infrared color-color diagram is a useful tool in (a) finding the young stars, and (b) making a guess at the age of young stars. In this diagram astronomers compare the differences in magnitude of a star at two wavelengths to the difference in magnitude of the same star in two other wavelengths.

Here is an example of a color-color plot using IRAC colors. By making models of stars, astronomers have determined that a star's position on this graph is related to its age since the star’s position on the graph is related to how warm or cool it is. Stars start out as cool (or red) and become warmer (bluer), therefore they shift positions on the diagram as they age.

--Class 0: these are very young stars, still buried deep inside their cocoon of gas.

--Class I: these are protostars surrounded by an infalling envelope of gas.

--Class II: these protostars, now in the T-Tauri stage, have developed protoplanetary disks.

--Class III: these are nearly fully developed stars, with just a remnant disk (and possibly planets) surrounding them.

==Homework: Your own literature search!==

Search the literature for references to IC 2118. Nearly all the astronomical literature is online at Harvard’s ADS website. (http://adsabs.harvard.edu) The ADS website allows you to search for abstracts using either names (IC2118) or coordinates (Ra/Dec). The website will let you look at older full papers, but only abstracts for recent papers. If there is a recent paper you wish to read, you will need to connect through a local university.

# Find all papers by Maria Kun. Which are refereed, and which are conference proceedings?
# Find all papers involving IC 2118.
# Find Hartmann et. al, 2005, ApJ, 629, 881.
# Search the web for observations in other wavelengths, not just images but photometry as well. Some good places to start are:
#*Gator, at IRSA, provides access to 2MASS, IRAS, and other wavelengths (including some large Spitzer surveys). This site can be accessed at: http://irsa.ipac.caltech.edu/applications/Gator/
#*USNO provides optical magnitudes, if they exist. Make sure to retrieve tables, not images. This site can be accessed at: http://www.nofs.navy.mil/data/FchPix/
#*Skyview provides images at other wavelengths. This site can be accessed at: http://skyview.gsfc.nasa.gov

== [[Writing the Research Paper for IC2118]] ==

== [[Information for Oil City High School meetings]] ==

The next meeting should be at 4 0'clock pm wednesday the 25th at matt walentosky's house. (Specifically for people going to California, so we can plan out events and travel...)

== Proposal for Radio Observations of T-Tauri Candidates in IC2118 ==

'''''SUPPORT INFORMATION - T-Tauri Candidates Emit Radio'''''

'''FROM MATT WALENTOSKY 6/30/2007 @ 2:30 PM - Two Abstracts Below'''

Title: Radio emission from pre-main-sequence stars
Authors: Skinner, Stephen Lee

Affiliation: AA(Colorado Univ., Boulder.)
Publication: Ph.D. Thesis Colorado Univ., Boulder.
Publication Date: 01/1992
Category: Space Radiation
Origin: STI

Abstract
This study focuses on the properties and physical origin of radio continuum emission from pre-main-sequence (PMS) stars. These are young stars, typically less than a few million years old, and are still in a phase of gravitational contraction that will ultimately be halted by the onset of hydrogen burning in their cores. First, I address the question of the origin of centimeter continuum emission in intermediate mass (approx. equal to 3-20 solar mass) PMS stars, the so-called 'Herbig Ae/Be stars'. A high-sensitivity radio survey of 57 such stars was undertaken using the Very Large Array and Australia Telescope, resulting in the detection of twelve stars. These observations provide a homogenous data base consisting of information on source sizes, radio luminosities, variability timescales, circular polarization, and spectral energy distributions in the wavelength range 2-20 cm. Using these data along with previously published spectroscopy, I conclude that centimeter radio emission from Herbig Ae/Be stars is predominantly thermal and in many cases wind-related. An unexpected result of the above program was the serendipitous detection of circularly polarized radio emission in the low mass (approx. equal to 1 solar mass) PMS star Hubble 4, a member of the class of 'weak-lined T Tauri stars' (WTTS). This provides some of the most convincing evidence to date for the existence of ordered magnetic fields in WTTS. In a second observing program, I have searched for evidence of cold (less than or equal to 50 K) circumstellar dust around WTTS, which might exist in the form of remnant disks. Of the sixteen WTTS that were observed in the wavelength range 450-1100 microns using the James Clerk Maxwell Telescope, only V836 Tau was detected. Its spectral energy distribution longward of 10 microns is consistent with that expected for a flat, axisymmetric circumstellar disk of mass approx. equal to 0.04 solar mass (= 42 Jupiter masses). This star may be a rare example of an object in which disk dispersal is underway, but not yet complete.

Title: Centimeter Radio Emission from Low-Mass Weak T Tauri Stars in Taurus-Auriga
Authors: Chiang, Eugene; Phillips, R. B.

Affiliation: AA(MIT), AB(MIT Haystack Observatory)
Publication: American Astronomical Society, 185th AAS Meeting, #48.08; Bulletin of the American Astronomical Society, Vol. 26, p.1388
Publication Date: 12/1994
Origin: AAS
Abstract Copyright:(c) 1994: American Astronomical Society

Abstract
We report on the results of a sensitive survey for lambda 3.6 cm radio emission from low-mass, weak T Tauri (WTT) stars in the Taurus-Auriga cloud complex. The target population consists of stars in the Herbig and Bell Catalog of spectral type K7 or later, and W(Hα ) <= 10 Angstroms. Of the 28 such stars surveyed using the Very Large Array down to detection thresholds of ~ 0.1 mJy, 7 (possibly 8) are observed to emit at strengths ranging from 0.1 to 2 mJy. Five of these young radio stars are newly discovered in our survey: V827 Tau and V710 Tau B are discovered to be relatively strong sources of mJy emission, while IW Tau, UX Tau B, and the possible detection LkHa 332-G1 form a new population of relatively weak emitters. Our radio survey and complementary surveys are pooled, and of 43 WTT stars K7 or later in Tau-Aur, 14 are now known to be radio emitters at lambda 6 and lambda 3.6 cm. Correlations between radio luminosity and other stellar parameters have been attempted but generally yield null results. Wide binarity (component separations in excess of 0''.13, 20 AU) appears unrelated to radio emission, as does spectral type. Furthermore, we find no convincing evidence for the extreme youth of radio stars, contrary to claims in the literature over the past decade. While we do find that radio-loud stars in our sample are formally younger than the radio-quiet stars by about 0.5 Myr, the reality of this relatively small age difference is highly suspect given uncertainties in the placement of these stars on the HR diagram. Moreover, Monte Carlo-type calculations involving distributing the stars on both the HR diagram and local CO gas density cast doubts on any differences between the radio stars and the general WTT population. We conclude that the age effect for low-mass radio WTT stars in Tau-Aur, if real, is much smaller than previous estimations by factors of 4-10. It is also possible centimeter wavelength surveys to date have still not properly described the radio luminosity function of low-mass WTT stars in Tau-Aur, and we urge future observations of these young stars with denser temporal coverage.

'''FROM SPUCK 6/30/2007 @ 2:35pm'''

I'm not sure that we will be able to use the GBT telescope in Green Bank for observations. The beamwidth at 10-15 GHz is probably too big.

Here are the specs for the GBT
Beamwidth (Table 3) Diffraction beamwidth (FWHM)
8 GHz it is 90",
20 GHz it is 36",
50 GHz it is 14"

Here are some questions from Sue Ann Heatherly at NRAO-Green Bank - One question I have is: is the resolving
power of the GBT sufficient to distinguish spatially between individual stars? Does large scale emission exist within the nebula that will confuse your results?

The VLA is probably the only instrument that can be used.

Synthesized Beamwidth (arcsec)depending on configuration
'''Nick Kelley 6/30/2007 3:00PM'''

400 cm 24.0" 80.0" 260.0" 850.0"

90 cm 6.0" 17.0" 56.0" 200.0"

20 cm 1.4 " 3.9" 12.5" 44.0"

6 cm 0.4 " 1.2""" 3.9 " 14.0 "

3.6 cm 0.24 " 0.7" 2.3 " 8.4 "

2 cm 0.14 " 0.4 " 1.2" 3.9"

1.3 cm 0.08 " 0.3" 0.9 "2.8"

0.7 cm 0.05 0.15 0.47 1
restricting my creativity

== T-Tauri/IC2118 Presentation for Astroblast 2007 ==

'''''Oil City High School students will be presenting August 11 at the Oil Region Astronomical Observatory'''''

FROM Nicholas James Kelley

Astroblast is amazin ... see pic

[[Image: astroblast1.JPG|center]]

Is their anyway we can attach the powerpoint here?? Matt W.

FROM SPUCK 7/25/07 - Matt ... we may be able to upload the powerpoint somewhere and then create a link to it. I'll need to check on this. Thanks, Mr. Spuck

== Monitoring T-Tauri Stars using the Perth Obervatory Telescope ==

'''Generating Light Curves'''

Using Perth Telescope

1. Check to see what the current sky conditions are by checking the live camera

http://www.perthobservatory.wa.gov.au//information/po_sky_camera.html

2. Check to see what the current weather conditions are

http://www.perthobservatory.wa.gov.au//information/po_weather.html

3. Weather forecast for Perth - http://www.bom.gov.au/weather/wa/

Log into telescope if conditions look good.

http://perthobservatory.org or http://202.72.190.15

EMAIL tspuck@hotmail.com for telescope access

== Chandra X-Ray Data ==

'''''To date no X-Ray data for IC2118 has been located.''''' If you know of any please contact Tim Spuck at tspuck@hotmail.com.

== Kitt Peak H-alpha Data ==

'''''In January of 2007 students from Oil City High School used the 0.9 Meter Telescope at Kitt peak to image regions of IC2118 in H-alpha.''''

Matt Walentosky, Nick Kelley, Sandy Weiser were the students who went!!!!!!

== USNO U, B, V, R, and I Data ==

'''''U, B, V, R, and I data will be used to generate more accurate SEDs for the T-Tauri Candidates'''''

== Outflows or Jets ==

'''''Visible outflows or jets are strong evidence that a candidate is indeed a T-Tauri Star'''''

== Multi Wavelength Composite Images ==

'''''Images and Methods'''''

FROM SPUCK 6/30/2007 See Image Below- IC 2118 3.6 µm (blue), 5.8 µm (green), 8.0 µm (red) tri-color composite generated using MaxIm DL. (By M. Heath, N. Kelley, P. Morton, M. Walentosky, S. Weiser – Oil City High School, Oil City, PA)

[[Image: spuckim3.JPG|center]]

IC 2118 Current Research Activities

2007-08-03T12:48:49Z

Weehler: paragraph 3

=Some more of the basics behind what we are doing!=

[[Image:L1014.jpg|right]]
Here is an image of an object called L1014. This object used to be known as a "starless core." On the left is a picture in visible light, and you can see why people thought it was starless. On the right is a picture from Spitzer clearly revealing a baby star inside (it's the bright red thing)!

When the protostar enters the next stage, labeled in the figure as the T Tauri stage, it’s still gaining mass and contracting slowly because material is still falling onto it, but it begins to eject gas in two giant gas jets, called bipolar flows. These jets and stellar winds eventually sweep away the envelope of gas still surrounding the protostar. In the surrounding disk protoplanets are beginning to form. (d) Astronomers call this phase the T Tauri phase, but some objects with jets are also "Class I"; T Tauris are also either "Class II" (for ones with thick disks) or "Class III" (for ones with thinner disks).
'''What do (d) and (e) refer to? Where are (a)-(c)? Are there illustrations that go here?--[[User:Weehler|Weehler]] 05:37, 3 August 2007 (PDT)'''
Leftover material in the disk surrounding the star clumps together and undergoes many collisions until most of the material has been swept up by objects orbiting the star, such as planets, asteroids and comets. (e)

The star’s life so far has been governed by the continuous inward pressure of gravity. The gravitational pressure keeps compressing the gas into a smaller and smaller volume, making it hotter and hotter in the core. As soon as the temperature in the core of the protostar becomes great enough, about 1,000,000 K, nuclear fusion begins. Nuclear fusion is the process in which small atomic nuclei combine to make larger atomic nuclei, releasing lots of energy. Inside the star, the gravitational pressure eventually pushes 4 protons and 2 electrons '''Are electrons involved in fusion?--[[User:Weehler|Weehler]] 05:48, 3 August 2007 (PDT)'''so close together that they fuse together to make helium. When this nuclear fusion begins, finally the star has a way to "fight back" against gravity. So much energy is released in this reaction that it enables the star to "push back" with an outward radiation pressure that balances the inward push of gravity. The protostar is now a full-fledged star, fusing hydrogen into helium in its core. (f) The star will stay the same size until it runs out of nuclear fuel in the core (all of the hydrogen has been converted into helium). Then, the pressure from gravity takes over again, pushing in on the star.

Star formation can be triggered by the collapse of large clouds of gas and dust. Sometimes the clouds collapse all by themselves (due to gravitational forces within the cloud) and sometimes they collapse because they’ve been pushed by the radiation and winds from stars that have already been formed. This seems to be the reason behind the star formation in IC 2118. It seems to have been triggered by some combination of forces from the main Orion Nebula Cluster (the fuzzy patch in the sword) and Rigel.

=Color-Color Plots=
[[Image:irac-colorcolor.png|left]]

The magnitude scale measures the brightness of a star. The system for assigning magnitude numbers was developed in ancient times. The '''brighter''' the star is, the '''lower''' its magnitude.

An infrared color-color diagram is a useful tool in (a) finding the young stars, and (b) making a guess at the age of young stars. In this diagram astronomers compare the differences in magnitude of a star at two wavelengths to the difference in magnitude of the same star in two other wavelengths.

Here is an example of a color-color plot using IRAC colors. By making models of stars, astronomers have determined that a star's position on this graph is related to its age since the star’s position on the graph is related to how warm or cool it is. Stars start out as cool (or red) and become warmer (bluer), therefore they shift positions on the diagram as they age.

--Class 0: these are very young stars, still buried deep inside their cocoon of gas.

--Class I: these are protostars surrounded by an infalling envelope of gas.

--Class II: these protostars, now in the T-Tauri stage, have developed protoplanetary disks.

--Class III: these are nearly fully developed stars, with just a remnant disk (and possibly planets) surrounding them.

==Homework: Your own literature search!==

Search the literature for references to IC 2118. Nearly all the astronomical literature is online at Harvard’s ADS website. (http://adsabs.harvard.edu) The ADS website allows you to search for abstracts using either names (IC2118) or coordinates (Ra/Dec). The website will let you look at older full papers, but only abstracts for recent papers. If there is a recent paper you wish to read, you will need to connect through a local university.

# Find all papers by Maria Kun. Which are refereed, and which are conference proceedings?
# Find all papers involving IC 2118.
# Find Hartmann et. al, 2005, ApJ, 629, 881.
# Search the web for observations in other wavelengths, not just images but photometry as well. Some good places to start are:
#*Gator, at IRSA, provides access to 2MASS, IRAS, and other wavelengths (including some large Spitzer surveys). This site can be accessed at: http://irsa.ipac.caltech.edu/applications/Gator/
#*USNO provides optical magnitudes, if they exist. Make sure to retrieve tables, not images. This site can be accessed at: http://www.nofs.navy.mil/data/FchPix/
#*Skyview provides images at other wavelengths. This site can be accessed at: http://skyview.gsfc.nasa.gov

== [[Writing the Research Paper for IC2118]] ==

== [[Information for Oil City High School meetings]] ==

The next meeting should be at 4 0'clock pm wednesday the 25th at matt walentosky's house. (Specifically for people going to California, so we can plan out events and travel...)

== Proposal for Radio Observations of T-Tauri Candidates in IC2118 ==

'''''SUPPORT INFORMATION - T-Tauri Candidates Emit Radio'''''

'''FROM MATT WALENTOSKY 6/30/2007 @ 2:30 PM - Two Abstracts Below'''

Title: Radio emission from pre-main-sequence stars
Authors: Skinner, Stephen Lee

Affiliation: AA(Colorado Univ., Boulder.)
Publication: Ph.D. Thesis Colorado Univ., Boulder.
Publication Date: 01/1992
Category: Space Radiation
Origin: STI

Abstract
This study focuses on the properties and physical origin of radio continuum emission from pre-main-sequence (PMS) stars. These are young stars, typically less than a few million years old, and are still in a phase of gravitational contraction that will ultimately be halted by the onset of hydrogen burning in their cores. First, I address the question of the origin of centimeter continuum emission in intermediate mass (approx. equal to 3-20 solar mass) PMS stars, the so-called 'Herbig Ae/Be stars'. A high-sensitivity radio survey of 57 such stars was undertaken using the Very Large Array and Australia Telescope, resulting in the detection of twelve stars. These observations provide a homogenous data base consisting of information on source sizes, radio luminosities, variability timescales, circular polarization, and spectral energy distributions in the wavelength range 2-20 cm. Using these data along with previously published spectroscopy, I conclude that centimeter radio emission from Herbig Ae/Be stars is predominantly thermal and in many cases wind-related. An unexpected result of the above program was the serendipitous detection of circularly polarized radio emission in the low mass (approx. equal to 1 solar mass) PMS star Hubble 4, a member of the class of 'weak-lined T Tauri stars' (WTTS). This provides some of the most convincing evidence to date for the existence of ordered magnetic fields in WTTS. In a second observing program, I have searched for evidence of cold (less than or equal to 50 K) circumstellar dust around WTTS, which might exist in the form of remnant disks. Of the sixteen WTTS that were observed in the wavelength range 450-1100 microns using the James Clerk Maxwell Telescope, only V836 Tau was detected. Its spectral energy distribution longward of 10 microns is consistent with that expected for a flat, axisymmetric circumstellar disk of mass approx. equal to 0.04 solar mass (= 42 Jupiter masses). This star may be a rare example of an object in which disk dispersal is underway, but not yet complete.

Title: Centimeter Radio Emission from Low-Mass Weak T Tauri Stars in Taurus-Auriga
Authors: Chiang, Eugene; Phillips, R. B.

Affiliation: AA(MIT), AB(MIT Haystack Observatory)
Publication: American Astronomical Society, 185th AAS Meeting, #48.08; Bulletin of the American Astronomical Society, Vol. 26, p.1388
Publication Date: 12/1994
Origin: AAS
Abstract Copyright:(c) 1994: American Astronomical Society

Abstract
We report on the results of a sensitive survey for lambda 3.6 cm radio emission from low-mass, weak T Tauri (WTT) stars in the Taurus-Auriga cloud complex. The target population consists of stars in the Herbig and Bell Catalog of spectral type K7 or later, and W(Hα ) <= 10 Angstroms. Of the 28 such stars surveyed using the Very Large Array down to detection thresholds of ~ 0.1 mJy, 7 (possibly 8) are observed to emit at strengths ranging from 0.1 to 2 mJy. Five of these young radio stars are newly discovered in our survey: V827 Tau and V710 Tau B are discovered to be relatively strong sources of mJy emission, while IW Tau, UX Tau B, and the possible detection LkHa 332-G1 form a new population of relatively weak emitters. Our radio survey and complementary surveys are pooled, and of 43 WTT stars K7 or later in Tau-Aur, 14 are now known to be radio emitters at lambda 6 and lambda 3.6 cm. Correlations between radio luminosity and other stellar parameters have been attempted but generally yield null results. Wide binarity (component separations in excess of 0''.13, 20 AU) appears unrelated to radio emission, as does spectral type. Furthermore, we find no convincing evidence for the extreme youth of radio stars, contrary to claims in the literature over the past decade. While we do find that radio-loud stars in our sample are formally younger than the radio-quiet stars by about 0.5 Myr, the reality of this relatively small age difference is highly suspect given uncertainties in the placement of these stars on the HR diagram. Moreover, Monte Carlo-type calculations involving distributing the stars on both the HR diagram and local CO gas density cast doubts on any differences between the radio stars and the general WTT population. We conclude that the age effect for low-mass radio WTT stars in Tau-Aur, if real, is much smaller than previous estimations by factors of 4-10. It is also possible centimeter wavelength surveys to date have still not properly described the radio luminosity function of low-mass WTT stars in Tau-Aur, and we urge future observations of these young stars with denser temporal coverage.

'''FROM SPUCK 6/30/2007 @ 2:35pm'''

I'm not sure that we will be able to use the GBT telescope in Green Bank for observations. The beamwidth at 10-15 GHz is probably too big.

Here are the specs for the GBT
Beamwidth (Table 3) Diffraction beamwidth (FWHM)
8 GHz it is 90",
20 GHz it is 36",
50 GHz it is 14"

Here are some questions from Sue Ann Heatherly at NRAO-Green Bank - One question I have is: is the resolving
power of the GBT sufficient to distinguish spatially between individual stars? Does large scale emission exist within the nebula that will confuse your results?

The VLA is probably the only instrument that can be used.

Synthesized Beamwidth (arcsec)depending on configuration
'''Nick Kelley 6/30/2007 3:00PM'''

400 cm 24.0" 80.0" 260.0" 850.0"

90 cm 6.0" 17.0" 56.0" 200.0"

20 cm 1.4 " 3.9" 12.5" 44.0"

6 cm 0.4 " 1.2""" 3.9 " 14.0 "

3.6 cm 0.24 " 0.7" 2.3 " 8.4 "

2 cm 0.14 " 0.4 " 1.2" 3.9"

1.3 cm 0.08 " 0.3" 0.9 "2.8"

0.7 cm 0.05 0.15 0.47 1
restricting my creativity

== T-Tauri/IC2118 Presentation for Astroblast 2007 ==

'''''Oil City High School students will be presenting August 11 at the Oil Region Astronomical Observatory'''''

FROM Nicholas James Kelley

Astroblast is amazin ... see pic

[[Image: astroblast1.JPG|center]]

Is their anyway we can attach the powerpoint here?? Matt W.

FROM SPUCK 7/25/07 - Matt ... we may be able to upload the powerpoint somewhere and then create a link to it. I'll need to check on this. Thanks, Mr. Spuck

== Monitoring T-Tauri Stars using the Perth Obervatory Telescope ==

'''Generating Light Curves'''

Using Perth Telescope

1. Check to see what the current sky conditions are by checking the live camera

http://www.perthobservatory.wa.gov.au//information/po_sky_camera.html

2. Check to see what the current weather conditions are

http://www.perthobservatory.wa.gov.au//information/po_weather.html

3. Weather forecast for Perth - http://www.bom.gov.au/weather/wa/

Log into telescope if conditions look good.

http://perthobservatory.org or http://202.72.190.15

EMAIL tspuck@hotmail.com for telescope access

== Chandra X-Ray Data ==

'''''To date no X-Ray data for IC2118 has been located.''''' If you know of any please contact Tim Spuck at tspuck@hotmail.com.

== Kitt Peak H-alpha Data ==

'''''In January of 2007 students from Oil City High School used the 0.9 Meter Telescope at Kitt peak to image regions of IC2118 in H-alpha.''''

Matt Walentosky, Nick Kelley, Sandy Weiser were the students who went!!!!!!

== USNO U, B, V, R, and I Data ==

'''''U, B, V, R, and I data will be used to generate more accurate SEDs for the T-Tauri Candidates'''''

== Outflows or Jets ==

'''''Visible outflows or jets are strong evidence that a candidate is indeed a T-Tauri Star'''''

== Multi Wavelength Composite Images ==

'''''Images and Methods'''''

FROM SPUCK 6/30/2007 See Image Below- IC 2118 3.6 µm (blue), 5.8 µm (green), 8.0 µm (red) tri-color composite generated using MaxIm DL. (By M. Heath, N. Kelley, P. Morton, M. Walentosky, S. Weiser – Oil City High School, Oil City, PA)

[[Image: spuckim3.JPG|center]]

IC 2118 Current Research Activities

2007-08-03T12:37:56Z

Weehler: confused about references to (a)-(e) in paragraph 2

=Some more of the basics behind what we are doing!=

[[Image:L1014.jpg|right]]
Here is an image of an object called L1014. This object used to be known as a "starless core." On the left is a picture in visible light, and you can see why people thought it was starless. On the right is a picture from Spitzer clearly revealing a baby star inside (it's the bright red thing)!

When the protostar enters the next stage, labeled in the figure as the T Tauri stage, it’s still gaining mass and contracting slowly because material is still falling onto it, but it begins to eject gas in two giant gas jets, called bipolar flows. These jets and stellar winds eventually sweep away the envelope of gas still surrounding the protostar. In the surrounding disk protoplanets are beginning to form. (d) Astronomers call this phase the T Tauri phase, but some objects with jets are also "Class I"; T Tauris are also either "Class II" (for ones with thick disks) or "Class III" (for ones with thinner disks).
'''What do (d) and (e) refer to? Where are (a)-(c)? Are there illustrations that go here?--[[User:Weehler|Weehler]] 05:37, 3 August 2007 (PDT)'''
Leftover material in the disk surrounding the star clumps together and undergoes many collisions until most of the material has been swept up by objects orbiting the star, such as planets, asteroids and comets. (e)

The star’s life so far has been governed by the continuous inward pressure of gravity. The gravitational pressure keeps compressing the gas into a smaller and smaller volume, making it hotter and hotter in the core. As soon as the temperature in the core of the protostar becomes great enough, about 1,000,000 K, nuclear fusion begins. Nuclear fusion is the process in which small atomic nuclei combine to make larger atomic nuclei, releasing lots of energy. Inside the star, the gravitational pressure eventually pushes 4 protons and 2 electrons so close together that they fuse together to make helium. When this nuclear fusion begins, finally the star has a way to "fight back" against gravity. So much energy is released in this reaction that it enables the star to "push back" with an outward radiation pressure that balances the inward push of gravity. The protostar is now a full-fledged star, fusing hydrogen into helium in its core. (f) The star will stay the same size until it runs out of nuclear fuel in the core (all of the hydrogen has been converted into helium). Then, the pressure from gravity takes over again, pushing in on the star.

Star formation can be triggered by the collapse of large clouds of gas and dust. Sometimes the clouds collapse all by themselves (due to gravitational forces within the cloud) and sometimes they collapse because they’ve been pushed by the radiation and winds from stars that have already been formed. This seems to be the reason behind the star formation in IC 2118. It seems to have been triggered by some combination of forces from the main Orion Nebula Cluster (the fuzzy patch in the sword) and Rigel.

=Color-Color Plots=
[[Image:irac-colorcolor.png|left]]

The magnitude scale measures the brightness of a star. The system for assigning magnitude numbers was developed in ancient times. The '''brighter''' the star is, the '''lower''' its magnitude.

An infrared color-color diagram is a useful tool in (a) finding the young stars, and (b) making a guess at the age of young stars. In this diagram astronomers compare the differences in magnitude of a star at two wavelengths to the difference in magnitude of the same star in two other wavelengths.

Here is an example of a color-color plot using IRAC colors. By making models of stars, astronomers have determined that a star's position on this graph is related to its age since the star’s position on the graph is related to how warm or cool it is. Stars start out as cool (or red) and become warmer (bluer), therefore they shift positions on the diagram as they age.

--Class 0: these are very young stars, still buried deep inside their cocoon of gas.

--Class I: these are protostars surrounded by an infalling envelope of gas.

--Class II: these protostars, now in the T-Tauri stage, have developed protoplanetary disks.

--Class III: these are nearly fully developed stars, with just a remnant disk (and possibly planets) surrounding them.

==Homework: Your own literature search!==

Search the literature for references to IC 2118. Nearly all the astronomical literature is online at Harvard’s ADS website. (http://adsabs.harvard.edu) The ADS website allows you to search for abstracts using either names (IC2118) or coordinates (Ra/Dec). The website will let you look at older full papers, but only abstracts for recent papers. If there is a recent paper you wish to read, you will need to connect through a local university.

# Find all papers by Maria Kun. Which are refereed, and which are conference proceedings?
# Find all papers involving IC 2118.
# Find Hartmann et. al, 2005, ApJ, 629, 881.
# Search the web for observations in other wavelengths, not just images but photometry as well. Some good places to start are:
#*Gator, at IRSA, provides access to 2MASS, IRAS, and other wavelengths (including some large Spitzer surveys). This site can be accessed at: http://irsa.ipac.caltech.edu/applications/Gator/
#*USNO provides optical magnitudes, if they exist. Make sure to retrieve tables, not images. This site can be accessed at: http://www.nofs.navy.mil/data/FchPix/
#*Skyview provides images at other wavelengths. This site can be accessed at: http://skyview.gsfc.nasa.gov

== [[Writing the Research Paper for IC2118]] ==

== [[Information for Oil City High School meetings]] ==

The next meeting should be at 4 0'clock pm wednesday the 25th at matt walentosky's house. (Specifically for people going to California, so we can plan out events and travel...)

== Proposal for Radio Observations of T-Tauri Candidates in IC2118 ==

'''''SUPPORT INFORMATION - T-Tauri Candidates Emit Radio'''''

'''FROM MATT WALENTOSKY 6/30/2007 @ 2:30 PM - Two Abstracts Below'''

Title: Radio emission from pre-main-sequence stars
Authors: Skinner, Stephen Lee

Affiliation: AA(Colorado Univ., Boulder.)
Publication: Ph.D. Thesis Colorado Univ., Boulder.
Publication Date: 01/1992
Category: Space Radiation
Origin: STI

Abstract
This study focuses on the properties and physical origin of radio continuum emission from pre-main-sequence (PMS) stars. These are young stars, typically less than a few million years old, and are still in a phase of gravitational contraction that will ultimately be halted by the onset of hydrogen burning in their cores. First, I address the question of the origin of centimeter continuum emission in intermediate mass (approx. equal to 3-20 solar mass) PMS stars, the so-called 'Herbig Ae/Be stars'. A high-sensitivity radio survey of 57 such stars was undertaken using the Very Large Array and Australia Telescope, resulting in the detection of twelve stars. These observations provide a homogenous data base consisting of information on source sizes, radio luminosities, variability timescales, circular polarization, and spectral energy distributions in the wavelength range 2-20 cm. Using these data along with previously published spectroscopy, I conclude that centimeter radio emission from Herbig Ae/Be stars is predominantly thermal and in many cases wind-related. An unexpected result of the above program was the serendipitous detection of circularly polarized radio emission in the low mass (approx. equal to 1 solar mass) PMS star Hubble 4, a member of the class of 'weak-lined T Tauri stars' (WTTS). This provides some of the most convincing evidence to date for the existence of ordered magnetic fields in WTTS. In a second observing program, I have searched for evidence of cold (less than or equal to 50 K) circumstellar dust around WTTS, which might exist in the form of remnant disks. Of the sixteen WTTS that were observed in the wavelength range 450-1100 microns using the James Clerk Maxwell Telescope, only V836 Tau was detected. Its spectral energy distribution longward of 10 microns is consistent with that expected for a flat, axisymmetric circumstellar disk of mass approx. equal to 0.04 solar mass (= 42 Jupiter masses). This star may be a rare example of an object in which disk dispersal is underway, but not yet complete.

Title: Centimeter Radio Emission from Low-Mass Weak T Tauri Stars in Taurus-Auriga
Authors: Chiang, Eugene; Phillips, R. B.

Affiliation: AA(MIT), AB(MIT Haystack Observatory)
Publication: American Astronomical Society, 185th AAS Meeting, #48.08; Bulletin of the American Astronomical Society, Vol. 26, p.1388
Publication Date: 12/1994
Origin: AAS
Abstract Copyright:(c) 1994: American Astronomical Society

Abstract
We report on the results of a sensitive survey for lambda 3.6 cm radio emission from low-mass, weak T Tauri (WTT) stars in the Taurus-Auriga cloud complex. The target population consists of stars in the Herbig and Bell Catalog of spectral type K7 or later, and W(Hα ) <= 10 Angstroms. Of the 28 such stars surveyed using the Very Large Array down to detection thresholds of ~ 0.1 mJy, 7 (possibly 8) are observed to emit at strengths ranging from 0.1 to 2 mJy. Five of these young radio stars are newly discovered in our survey: V827 Tau and V710 Tau B are discovered to be relatively strong sources of mJy emission, while IW Tau, UX Tau B, and the possible detection LkHa 332-G1 form a new population of relatively weak emitters. Our radio survey and complementary surveys are pooled, and of 43 WTT stars K7 or later in Tau-Aur, 14 are now known to be radio emitters at lambda 6 and lambda 3.6 cm. Correlations between radio luminosity and other stellar parameters have been attempted but generally yield null results. Wide binarity (component separations in excess of 0''.13, 20 AU) appears unrelated to radio emission, as does spectral type. Furthermore, we find no convincing evidence for the extreme youth of radio stars, contrary to claims in the literature over the past decade. While we do find that radio-loud stars in our sample are formally younger than the radio-quiet stars by about 0.5 Myr, the reality of this relatively small age difference is highly suspect given uncertainties in the placement of these stars on the HR diagram. Moreover, Monte Carlo-type calculations involving distributing the stars on both the HR diagram and local CO gas density cast doubts on any differences between the radio stars and the general WTT population. We conclude that the age effect for low-mass radio WTT stars in Tau-Aur, if real, is much smaller than previous estimations by factors of 4-10. It is also possible centimeter wavelength surveys to date have still not properly described the radio luminosity function of low-mass WTT stars in Tau-Aur, and we urge future observations of these young stars with denser temporal coverage.

'''FROM SPUCK 6/30/2007 @ 2:35pm'''

I'm not sure that we will be able to use the GBT telescope in Green Bank for observations. The beamwidth at 10-15 GHz is probably too big.

Here are the specs for the GBT
Beamwidth (Table 3) Diffraction beamwidth (FWHM)
8 GHz it is 90",
20 GHz it is 36",
50 GHz it is 14"

Here are some questions from Sue Ann Heatherly at NRAO-Green Bank - One question I have is: is the resolving
power of the GBT sufficient to distinguish spatially between individual stars? Does large scale emission exist within the nebula that will confuse your results?

The VLA is probably the only instrument that can be used.

Synthesized Beamwidth (arcsec)depending on configuration
'''Nick Kelley 6/30/2007 3:00PM'''

400 cm 24.0" 80.0" 260.0" 850.0"

90 cm 6.0" 17.0" 56.0" 200.0"

20 cm 1.4 " 3.9" 12.5" 44.0"

6 cm 0.4 " 1.2""" 3.9 " 14.0 "

3.6 cm 0.24 " 0.7" 2.3 " 8.4 "

2 cm 0.14 " 0.4 " 1.2" 3.9"

1.3 cm 0.08 " 0.3" 0.9 "2.8"

0.7 cm 0.05 0.15 0.47 1
restricting my creativity

== T-Tauri/IC2118 Presentation for Astroblast 2007 ==

'''''Oil City High School students will be presenting August 11 at the Oil Region Astronomical Observatory'''''

FROM Nicholas James Kelley

Astroblast is amazin ... see pic

[[Image: astroblast1.JPG|center]]

Is their anyway we can attach the powerpoint here?? Matt W.

FROM SPUCK 7/25/07 - Matt ... we may be able to upload the powerpoint somewhere and then create a link to it. I'll need to check on this. Thanks, Mr. Spuck

== Monitoring T-Tauri Stars using the Perth Obervatory Telescope ==

'''Generating Light Curves'''

Using Perth Telescope

1. Check to see what the current sky conditions are by checking the live camera

http://www.perthobservatory.wa.gov.au//information/po_sky_camera.html

2. Check to see what the current weather conditions are

http://www.perthobservatory.wa.gov.au//information/po_weather.html

3. Weather forecast for Perth - http://www.bom.gov.au/weather/wa/

Log into telescope if conditions look good.

http://perthobservatory.org or http://202.72.190.15

EMAIL tspuck@hotmail.com for telescope access

== Chandra X-Ray Data ==

'''''To date no X-Ray data for IC2118 has been located.''''' If you know of any please contact Tim Spuck at tspuck@hotmail.com.

== Kitt Peak H-alpha Data ==

'''''In January of 2007 students from Oil City High School used the 0.9 Meter Telescope at Kitt peak to image regions of IC2118 in H-alpha.''''

Matt Walentosky, Nick Kelley, Sandy Weiser were the students who went!!!!!!

== USNO U, B, V, R, and I Data ==

'''''U, B, V, R, and I data will be used to generate more accurate SEDs for the T-Tauri Candidates'''''

== Outflows or Jets ==

'''''Visible outflows or jets are strong evidence that a candidate is indeed a T-Tauri Star'''''

== Multi Wavelength Composite Images ==

'''''Images and Methods'''''

FROM SPUCK 6/30/2007 See Image Below- IC 2118 3.6 µm (blue), 5.8 µm (green), 8.0 µm (red) tri-color composite generated using MaxIm DL. (By M. Heath, N. Kelley, P. Morton, M. Walentosky, S. Weiser – Oil City High School, Oil City, PA)

[[Image: spuckim3.JPG|center]]

User:Weehler

2007-06-08T14:24:58Z

Weehler: suggestions

Luisa,
This information is so cool. As usual, though, I need more help. Can the captions include what the different pictures are showing me about sensitivity, resolution, etc? In other words, can the differences i'm supposed to notice btwn ISO, IRAF, and SIRTF be explained? Does this question make sense?
Cindy--[[User:Weehler|Weehler]] 07:24, 8 June 2007 (PDT)